{
    "0806/0806.0496_arXiv.txt": {
        "abstract": "We discuss the possibility to identify anisotropic and/or inhomogeneous cosmological models using type Ia supernova data. A search for correlations in current type Ia peak magnitudes over a large range of angular scales yields a null result.  However, the same analysis limited to supernovae at low redshift, shows a feeble anticorrelation at the 2$\\sigma$ level at angular scales $\\theta \\approx 40^{\\circ}$. Upcoming data from, e.g., the SNLS (Supernova Legacy Survey) and the SDSS-II (SDSS: Sloan Digital Sky Survey) supernova searches will improve our limits on the size of -- or possibly detect -- possible correlations also at high redshift at the per cent level in the near future. With data from the proposed SNAP (SuperNova Acceleration Probe) satellite, we will be able to detect the induced correlations from gravitational lensing on type Ia peak magnitudes on scales less than a degree. ",
        "introduction": "Apart from the identification of dark matter, the most dominant question in cosmology today is that of what is responsible for the apparent acceleration of the universal expansion. Assuming that the cause is a dominant energy component with negative pressure, current efforts are focused on establishing whether this dark energy (DE) component can be described using a cosmological constant (CC), i.e. an energy component with constant density and a fixed equation of state (EOS) given by $p = \\omega\\rho$ where $w=-1$, or whether dark energy is dynamical (DDE) with a varying EOS, $w=w(z)$.  Two main techniques are employed in these investigations: probing cosmological distances and probing the growth of cosmological structures, both of which are sensitive to the energy content of the universe. Distances are most robustly probed via type Ia supernovae (SNe~Ia) \\cite{1998AJ....116.1009R,1999ApJ...517..565P,2006A&A...447...31A,2007ApJ...666..694W}, the baryon acoustic oscillation (BAO) peak \\cite{2005MNRAS.362..505C,2005ApJ...633..560E,2007MNRAS.381.1053P} and the last scattering surface of the cosmic microwave background (CMB) \\cite{2008arXiv0803.0547K}. Structure growth is mainly probed by large galaxy surveys such as the 2dF \\cite{2003MNRAS.346...78H} and the SDSS \\cite{2004ApJ...607..655P,2007ApJ...657..645P} but also via galaxy cluster counts \\cite{2006ApJ...653..954D} and weak lensing \\cite{2006ApJ...647..116H}. Despite some claims to the contrary, there is a consensus that current data are perfectly consistent with having approximately 70\\% of the energy density in the universe in the form of a CC \\cite{2007ApJ...666..716D}.  However, cosmological distances only depend on the dark energy EOS in the form of a double integral and it is thus very difficult to investigate the time evolution of the EOS parameter, $w(z)$. Though not demanding it, current data thus still allow for large deviations of $w(z)$ from the CC value of $w=-1$ \\cite{2007ApJ...666..716D}. Even if the limits on $w(z)$ will improve in the future, detecting such an evolution remains notoriously difficult, especially if the deviation from the CC value turns out to be small.  It is thus of interest to look for alternative ways to detect or constrain any behaviour of the DE that differs from that of a CC. One such possibility is to study not temporal evolution, but rather spatial variations in DE properties. The CC value $w=-1$ is special in the sense that it not only gives a constant energy density in time, but also that it prevents any spatial clustering of DE. Any deviations from this value inevitably cause the DE to cluster, a clustering which -- if detected -- would refute the CC as the dominant energy component in the universe.  The clustering properties of DDE in different models and scenarios have been investigated in several recent articles (see, e.g., \\cite{2007arXiv0709.2227M} and references therein). For values close to $w=-1$, such clustering is expected to be weak and take place mainly on very large scales, larger than the current Hubble radius and thus inaccessible to observations. On smaller scales, it can be shown that the DE density will tend to be anticorrelated with the matter density \\cite{2007arXiv0709.2227M,2007PhRvD..75f3507D}. Although, again, the amount of clustering is expected to be small (on the order of $10^{-5}[1+w]$), it is nevertheless important to get observational confirmation that this is indeed the case, even if it is not expected from purely theoretical considerations. On a more speculative note, also in models where the apparent acceleration of the universe is explained in terms of the back-reaction of inhomogeneities in the universe, we expect to have spatial variations in, e.g., cosmological distances \\cite{2008PhRvD..77b3003M}. Furthermore, cosmological observations can be used to put constraints on the degree of centricity in spherically symmetric inhomogeneous models \\cite{2007JCAP...02...19E,2007PhRvD..75b3506A} and on anisotropic DE models \\cite{2007arXiv0707.0279K,2006MNRAS.371.1373M,2008PhRvD..77b3534R,2007MNRAS.382..793M,2006PhRvD..74d3505T}.  One way to probe models with inhomogeneous DE or back-reaction is to look for anisotropies in the observed peak magnitudes of SNe~Ia. In general, we expect cosmological distances to be correlated on scales similar to the clustering scale of the DE or the matter inhomogeneities responsible for the back-reaction effect. The detection of such correlations is demanding because of the intrinsic variation and observational uncertainties in SN~Ia peak magnitudes.  Since systematic effects connected to, e.g., the observational properties of different telescopes, calibration issues and details of the light-curve fitting procedure can induce correlations in SN~Ia magnitudes, it is important to minimize these effects by using as homogeneous a data set as possible in the analysis. Also, physical effects such as peculiar motions \\cite{2007ApJ...661..650H}, gravitational lensing \\cite{2006PhRvL..96b1301C} and dust extinction \\cite{2007ApJ...657...71Z} will introduce correlations in the SN Ia data and need to be controlled and, if possible, corrected for.  In Kolatt \\& Lahav \\cite{2001MNRAS.323..859K}, an early data set consisting of a total of 79 SNe~Ia combined from the Supernova Cosmology Project \\cite{1999ApJ...517..565P} and the High-$z$ Supernova Search Team \\cite{1998AJ....116.1009R}  was used to look for directional variations in the best fit cosmological parameters. The result is consistent with the expected statistical variations in a homogeneous universe. Gupta \\textit{et al} \\cite{2007astro.ph..1683G} used a more homogeneous set of 157 SNe~Ia \\cite{2004ApJ...607..665R} to look for directional variations in the statistical scatter around the best fit cosmological model, again with a result consistent with null variations. A slightly different technique was used in Bochner \\cite{2007astro.ph..2730B} to look for anisotropic scatter in 172 SNe~Ia \\cite{2003ApJ...594....1T} with similar results. Although current investigations have been mainly inconclusive, Schwartz \\& Weinhorst \\cite{2007A&A...474..717S} find an offset in the best fit calibration of low $z$ SNe~Ia, between the north and south equatorial hemispheres at the 95\\% confidence level. Whether this hint of a north/south asymmetry is due to a statistical coincidence, observational systematic differences, or correlated peculiar motions or has a cosmological origin is not clear at the moment.  In this paper we devise a general methodology for detecting angular correlations in SN~Ia magnitude residuals. We apply this method to current SN~Ia data as well as simulated upcoming data sets.  In section~\\ref{method}, we describe the method for analysing SN~Ia magnitude residuals for angular correlations and define the detection limit for such correlations. Two data sets from the literature are analysed for angular correlations in section~\\ref{analyse}. In section~\\ref{future}, we present different SN Ia surveys and investigate the detection limits for each of them. In section~\\ref{toymodel}, we assume a toy model correlation function and illustrate how future data can put limits on such a correlation. The paper is concluded in section~\\ref{conclusions}. ",
        "conclusions": "The cosmological community is hard at work trying to constrain the behaviour of DE. The most pressing question is whether DE is a CC or something dynamical. A detection of temporal or spatial variations of DE would answer this question and refute the assertion of the CC being the dominant energy component in the universe. At the same time, there are alternative attempts to explain the apparent acceleration of the universe without invoking DE at all: instead as originating in the large scale inhomogeneities in the matter distribution. Irrespective of their origin, large scale inhomogeneities would manifest themselves as anisotropies in the observed magnitudes of SNe~Ia.  In this paper, we have devised a methodology for detecting angular correlations in SN~Ia magnitude residuals. The methodology was applied to two recently found data sets (astier06 and davis07), neither of which show any signs of correlations at angular scales $0^\\circ<\\theta<180^\\circ$. The uncertainties on the measured correlations, $C(\\theta)$, are approximately 10\\% and 5\\% for the astier06 and davis07 data sets respectively, using an angular resolution of $\\sim 15^\\circ$.  Note, however, that the two data sets are not independent, since they overlap partially for the high redshift sample, and include almost the same set of nearby SNe. The main difference between the data sets is that different assumptions on SN~Ia properties, such as intrinsic colour and dust extinction in the host galaxy, are used when deriving the SN~Ia peak magnitudes. Such assumptions are responsible for some of the most important systematic uncertainties in SN~Ia cosmology. A comparison between the results for the two data sets allows us to keep the impact of these uncertainties on the measured correlation function under control. Because systematic uncertainties are typically different at low and high redshifts, and we expect effects from DE inhomogeneities to grow with redshift, we have applied the same analysis to the two sub-samples separately. At low $z$, we found an anticorrelation at the 2$\\sigma$ level at angular scales $\\theta \\approx 40^{\\circ}$. Due to the weak statistical significance of this detection, and the low redshifts, we cannot draw any conclusions about the possible implications for DE inhomogeneity. It is worth noticing that the nearby SNe analysed are a collection of very inhomogeneous data samples, collected by different observers at various telescopes and thus, there could be systematic uncertainties involved which have not been included in the error bars (see, e.g., discussion in \\cite{2003A&A...404..901N}). We note that we see no signs of anisotropy on angular scales $\\theta \\sim 180^{\\circ}$, corresponding to the north/south asymmetry detected in \\cite{2007A&A...474..717S}. Our data sample, however, is only partially overlapping with that work, since we have only included SNe in the Hubble flow ($z \\geq 0.015$) in our analysis in order to mitigate the effects from correlations arising from peculiar motions.  Gravitational lensing correlations will only be significant on small scales, where we expect the effects from DE clustering to be negligible. Using data from the proposed SNAP satellite, we should be able to detect the lensing correlations in SN~Ia magnitudes with an angular resolution of $<0.3^\\circ$.  Correlations induced by intervening galactic dust will also mainly take place on very small angular scales, whereas correlations from dust in the Milky Way could potentially be a problem for an all-sky survey. Most surveys, however, observe in directions in the sky where the Milky Way dust extinction is well measured and understood.  More SN data are needed if we are to detect presumably weak correlations and better constrain inhomogeneous/anisotropic models. Our simulations of future data illustrate how the survey geometries govern on what scales, and to what precision, we will be able to detect possible angular correlations. We find that using data from the soon to be completed SDSS-II and SNLS surveys, we will typically be able to detect correlations in the magnitude residuals at the per cent level.  Any claim of a deviation of DE properties from those of a CC will be subject to  intense scrutiny from the cosmological community, and will need to be backed up by independent evidence in order to be generally accepted. It is therefore of utmost importance to pursue the study of DE along multiple paths. Although observationally demanding, the search for spatial variations in DE properties is a useful complement to studies aiming at constraining the time evolution of DE. Specifically, for SNe~Ia, we expect systematic effects connected to the time evolution of DE, such as temporal evolution of SN~Ia properties, to be of less importance when studying spatial variations of DE. Since systematic effects constitute the limiting factor for future SN~Ia surveys, we expect interesting results for the spatial correlations of SN magnitudes in the near future, irrespective of whether such correlations are found or not.  \\ack MB acknowledges support from the HEAC Centre funded by the Swedish Research Council. EM acknowledges support for this study by the Swedish Research Council and from the Anna-Greta and Holger Crafoord fund. The authors would like to thank Ariel Goobar for useful discussions."
    },
    "0806/0806.2988_arXiv.txt": {
        "abstract": "Photometry and long-slit spectroscopy are presented for 14 S0 and spiral galaxies of the Fornax, Eridanus and Pegasus cluster, and NGC~7582 group. The structural parameters of the galaxies are derived from the $R-$band images by performing a two-dimensional photometric decomposition of the surface-brightness distribution. This is assumed to be the sum of the contribution of a bulge and disc component characterized by elliptical and concentric isophotes with constant (but possibly different) ellipticity and position angles. The rotation curves and velocity dispersion profiles are measured from the spectra obtained along the major axis of galaxies. The radial profiles of the \\Hb , Mg, and Fe line-strength indices are presented too. Correlations between the central values of \\Mgd , \\Fe , \\Hb , and $\\sigma$ are found. The age, metallicity and $\\alpha/$Fe enhancement of the stellar population in the center and at the radius where bulge and disc give the same contribution to the total surface brightness are obtained using stellar population models with variable element abundance ratios. Three classes of bulges are identified. The youngest bulges ($\\sim2$ Gyr) with ongoing star formation, intermediate-age bulges (4--8 Gyr) have solar metallicity, and old bulges ($\\sim10$ Gyr) have high metallicity. Most of the sample bulges display solar $\\alpha/$Fe enhancement, no gradient in age, and a negative gradient of metallicity.  The presence of negative gradient in the metallicity radial profile favors a scenario with bulge formation via dissipative collapse. This implies strong inside-out formation that should give rise to a negative gradient in the $\\alpha$/Fe enhancement too. But, no gradient is measured in the \\aFe\\/ radial profiles for all the galaxies, except for NGC~1366. In this galaxy there is a kinematically-decoupled component, which is younger than the rest of host bulge. It possibly formed by enriched material probably acquired via interaction or minor merging. The bulge of NGC~1292 is the most reliable pseudobulge of our sample. The properties of its stellar population are consistent with a slow buildup within a scenario of secular evolution. ",
        "introduction": "\\label{sec:introduction}  The relative importance of the dissipative collapse \\citep{eglbsa62,sandage90,gilwys98}, major and minor merging events \\citep{kauffmann96,coletal00,aguetal01}, and redistribution of disc material due to the presence of a bar or environmental effects \\citep{korken04} drives the variety of properties observed in bulges.  The bulges of lenticulars and early-type spirals are similar to low-luminosity elliptical galaxies. Their photometric and kinematic properties satisfy the same fundamental plane correlation (FP) found for ellipticals \\citep{bendetal92,bendetal93,bursetal97,aguetal05a}. The surface-brightness radial profile of big bulges is well described by the de Vaucouleurs law \\citep{andretal95,caroetal98,molletal01} {\\bf even if this could drastically change by taking into account the small-scale inner structures smoothed by the seeing in the ground-base observations \\citep{balceetal03}.} Some of them are rotationally-flattened oblate spheroids with little or no anisotropy \\citep{kormetal82,davill83,capetal06}. But, the intrinsic shape of a large fraction of early-type bulges is triaxial, as shown by the isophotal twisting \\citep{lindblad56,zarlo86}, misalignment with respect to disc \\citep{bertetal91,mendetal07}, and non-circular gas motions \\citep{bertetal89,gerhetal89,berman01,corsetal03,coccetal04}. The bulk of their stellar population formed between redshift 3 and 5 ($\\simeq$12 Gyr) in a short time-scale \\citep{bernetal98,mehletal03,thometal05}. The enrichment of interstellar medium is strongly related to the time delay between SNII and SNIa, which contributed most of the $\\alpha$ elements and Iron, respectively \\citep{wofago92,thmabe03}.  On the contrary, the bulges of late-type spiral galaxies are reminiscent of discs. They are flat components \\citep{fatpel03} with exponential surface-brightness radial profiles \\citep{AndSan94,dejong96,macartetal03} and rotate as fast as discs \\citep{korm93,kormetal02}. Moreover, the stellar population in late-type bulges is younger than in early-type bulges \\citep{tragetal99,goud99,thda06}. They appear to have lower metallicity \\citep{gandetal07} and lower $\\alpha/$Fe enhancement with respect to early type galaxies \\citep{procetal02,peleetal07,afasil05}.  In the current paradigm, early-type bulges were formed by rapid collapse and merging events while late-type bulges have been slowly assembled by internal and environmental secular processes \\citep{korken04}. But many questions are still open.  For instance, the monolithic collapse can not explain the presence in bulges of kinematically-decoupled components \\citep{pizzetal02,krajaf03,emsetal04,mcdetal06}. Moreover, the environment plays a role in defining the properties of galaxies \\citep[e.g.,][]{dresler80,cozietal01,clemetal06,brouetal07}. Recent studies of early-type galaxies in different environments \\citep{beui02,thometal05,thda06} have shown that age, metallicity, and $\\alpha/$Fe enhancement are more correlated with the total mass of the galaxy than local environment.  To investigate the formation and evolution of the bulges, there are two possible approaches: going backward in redshift and look to evolution of galaxies through cosmic times or analyze in detail nearby galaxies to understand the properties of their stellar population in terms of dominant mechanism at the epochs of star formation and mass assembly. In this work, we present a photometric and spectroscopic study of the bulge dominated region of a sample of spiral galaxies in the Fornax and Pegasus clusters. Our aim is to estimate the age and metallicity of the stellar population and the efficiency and timescale of the last episode of star formation to disentangle between early rapid assembly and late slow growing.  The galaxy sample is presented in Sect. \\ref{sec:sample}. The photometric observations are described in Sect. \\ref{sec:observation_photometry}. The structural parameters of the bulge and disc of the sample galaxies are derived by analyzing their two-dimensional surface brightness distribution in Sect. \\ref{sec:decomposition}. The spectroscopic observations are described in Sect. \\ref{sec:observation_spectroscopy}.  The stellar kinematics and line-strength indices are measured from long-slit spectra in Sect. \\ref{sec:kinematics}. The central values of the line-strength indices are derived in Sect. \\ref{sec:linestreng_cent}. They are used to estimate the age, metallicity, and $\\alpha/$Fe-enhancement of the stellar population of the bulge in Sect. \\ref{agemet_cent} . Their gradients in the bulge dominated region are discussed in Sect. \\ref{sec:agemetalpha_grad}. The identification of pseudobulges hosted by sample galaxies is performed in Sect. \\ref{slowfastrotator}. Finally, conclusions are given in Sect. \\ref{conclusion}. ",
        "conclusions": "\\label{conclusion}  The structural parameters and properties of the stellar population of the bulges of sample of 14 S0 and spiral galaxies of the Fornax, Eridanus and Pegasus cluster, and NGC~7582 group were investigated to constrain the dominant mechanism at the epoch of their assembly.  \\begin{itemize}  \\item The bulge and disc parameters of the sample galaxies were derived performing a two-dimensional photometric decomposition of their $R-$band images. The surface-brightness distribution of the galaxy was assumed to be the sum of the contribution of a S\\'ersic bulge and an exponential disk. The two components were characterized by elliptical and concentric isophotes with constant (but possibly different) ellipticity and position angles. Most of the bulges have a S\\'ersic index $n\\leq2$ and for few of them the apparent flattening of the bulge is similar to that of the disc. According to \\citet{korken04} the disc-like flattening and radial profile are the photometric signature of the pseudobulge.  \\item The central values of velocity dispersion $\\sigma$ and \\Mgb, \\Mgd, \\Hb, \\Fe, and \\MgFe\\/ line-strength indices were derived from the major-axis spectra. Correlations between \\Mgd , \\Fe , \\Hb , and $\\sigma$ were found. The \\Mgd$-\\sigma$ and \\Hb$-\\sigma$, correlations are steeper than those found for early-type galaxies \\citep[e.g.,][]{bernetal98,jorgen99,kuntetal00,mehletal03}.  The \\Fe$-\\sigma$ correlation is consistent with previous findings for spiral bulges \\citep{idiaetal96,prugetal01,procetal02}.  \\item The mean ages, total metallicities, and total $\\alpha/$Fe enhancements in the center of the sample bulges were derived by using the stellar population models by \\citet{thmabe03}. The youngest bulges have an average age of 2 Gyr. They are characterized by ongoing star formation.  The stellar population of intermediate-age bulges is 4 to 8 Gyr old. It has solar metallicity (\\ZH$\\;=0.0$ dex).  The older bulges have a narrow distribution in age around 10 Gyr and high metallicity (\\ZH$\\;=0.30$ dex). Most of the sample bulges display solar $\\alpha/$Fe enhancements. A few have a central super-solar enhancement (\\aFe$\\;=0.3$).  \\item There is no correlation between age, metallicity, and $\\alpha/$Fe enhancement of bulges with the membership of the host galaxy to different cluster. There is a correlation with the velocity dispersion. The more massive bulges of our sample galaxies are older, more metal rich and characterized by a fast star formation. Since we did not find any correlation with galaxy morphology we exclude a strong interplay between the bulge and disc components.  \\item Most of the sample galaxies show no gradient in age and a negative gradient of metallicity. This is in agreement with the earlier findings by \\citet{mehletal03} and \\citet{sancetal06s} for the early-type galaxies, and by \\citet{jabletal07} for bulges. The presence of negative gradient in the metallicity radial profile favors a scenario with bulge formation via dissipative collapse. This implies strong inside-out formation that should give rise to a negative gradient in the $\\alpha$/Fe enhancement too \\citep{fesi02}. But, no gradient was measured in the \\aFe\\/ radial profiles for all the galaxies, except for NGC~1366 and NGC~7531. Moreover, the correlation between the central value and gradient of metallicity can not be built by pure dissipative collapse \\citep{besh99,cobari99} and suggests that mergers or acquisition events need to be invoked during the bulge assembly.  \\item The peculiar gradients observed for the stellar population of the bulges of NGC~1366 and NGC~7531 suggest they host a substructure.  Very interestingly, in NGC~1366 we found the presence of a kinematically-decoupled component. It is younger than the host bulge and formed by enriched material probably acquired via interaction or minor merging.  \\item According to the prescriptions by \\citet{korken04} the bulge of NGC~1292 is a pseudobulge. The properties of its stellar population are consistent with a slow buildup within a scenario of secular evolution. Indeed, the bulge of NGC~1292 has a intermediate age (3 Gyr) and low metal content (\\ZH$=-0.7$ dex). The $\\alpha/$Fe enhancement is the lowest in our sample (\\aFe$=-0.12$ dex) suggesting a prolonged star formation history. The presence of emission lines in the spectrum is a signature of ongoing star formation.  \\end{itemize}"
    },
    "0806/0806.0944_arXiv.txt": {
        "abstract": "\\leftskip1.0cm \\rightskip1.0cm We present and compare the predictions of various cosmic-ray Monte Carlo models for the energy  ($dE/d\\eta$) and particle ($dN/d\\eta$) flows in p-p, p-Pb and Pb-Pb collisions at $\\sqrtsnn =$~14, 8.8, and 5.5 TeV respectively, in the range covered by forward LHC detectors like CASTOR or TOTEM (5.2$<|\\eta|<$6.6) and ZDC or LHCf ($|\\eta|\\gtrsim$8.1 for neutrals). ",
        "introduction": "The origin and nature of cosmic rays (CRs) with energies between $10^{15}$~eV and the Greisen-Zatsepin-Kuzmin (GZK) cutoff at about $10^{20}$\\,eV, recently measured by the HiRes~\\cite{hires} and Auger~\\cite{auger}  experiments, remains a central open question in high-energy astrophysics. One key to solving this question is the determination of the elemental composition of cosmic rays in this energy range. The candidate particles, ranging from protons to nuclei as massive as iron, generate ``extended air-showers'' (EAS) in interactions with air nuclei when entering the Earth's atmosphere. Due to their low observed flux (Fig.~\\ref{fig:CRs}, left), only indirect (yet complementary) measurements are possible % using the atmosphere as ``calorimeter''. The first method relies on measuring the fluorescence light emitted by air molecules excited by the cascade of secondaries. % The second one relies on the use of either scintillators or water \\v{C}erenkov tanks to sample the shower at ground. Recent Auger results are consistent with the hypothesis that the highest energy CRs are protons (source correlated with Active Galactic Nuclei)~\\cite{auger2}. However, fluorescence-based measurements of the shower maximum as a function of primary energy % favour a mixed composition (Fig.~\\ref{fig:CRs}, right). %  \\begin{figure}[htbp] \\centering{\\includegraphics[width=7.cm]{CR-flux-HEP-scaled2p5-5c.eps} \\includegraphics[width=7.cm]{augerHiresFlysEye.eps}} \\caption{Left: Measured cosmic ray energy spectrum compared to the center-of-mass energy of various hadron colliders. Right: Measurements of CRs shower maximum in the atmosphere as a function of primary energy for several experiments along with MC predictions for proton and iron primaries.} \\label{fig:CRs} \\end{figure}  The determination of the primary energy (from the surface detectors alone) and mass (from either method) relies on hadronic Monte Carlo (MC) codes which describe the interactions of the primary cosmic-ray in the upper atmosphere. The bulk of the primary particle production is dominated by forward and soft QCD interactions, modeled commonly in Regge-Gribov-based~\\cite{Gribov:1968fc} approaches with parameters constrained by the existing collider data ($E_{lab}\\lesssim$~10$^{15}$~eV). When extrapolated to energies around the GZK-cutoff, the current MCs predict energy and multiplicity flows differing by factors as large as three, with significant inconsistencies in the forward region. % The coming energy frontier for hadron collisions will be reached by the Large Hadron Collider (LHC). The LHC will boast a full compliment of detectors in almost the full range of pseudorapidity (Fig.~\\ref{fig:LHC}). Measurement of forward particle production in p-p, p-Pb, and Pb-Pb collisions at LHC energies ($E_{lab} \\approx 10^{17}$ eV) will thus provide strong constraints on these models and allow for more reliable determinations of the CR energy and composition at the highest energies. ",
        "conclusions": ""
    },
    "0806/0806.0343_arXiv.txt": {
        "abstract": "We present luminosity and surface-brightness distributions of 40\\,111 galaxies with $K$-band photometry from the United Kingdom Infrared Telescope (UKIRT) Infrared Deep Sky Survey (UKIDSS) Large Area Survey (LAS), Data Release 3 and optical photometry from Data Release 5 of the Sloan Digital Sky Survey (SDSS). Various features and limitations of the new UKIDSS data are examined, such as a problem affecting Petrosian magnitudes of extended sources. Selection limits in $K$- and $r$-band magnitude, $K$-band surface brightness and $K$-band radius are included explicitly in the $1/V_\\mathrm{max}$ estimate of the space density and luminosity function. The bivariate brightness distribution in $K$-band absolute magnitude and surface brightness is presented and found to display a clear luminosity--surface brightness correlation that flattens at high luminosity and broadens at low luminosity, consistent with similar analyses at optical wavelengths. Best fitting Schechter function parameters for the $K$-band luminosity function are found to be $M^*-5\\log h=-23.19 \\pm 0.04$, $\\alpha=-0.81 \\pm 0.04$ and $\\phi^*=(0.0166 \\pm 0.0008)h^3$\\,Mpc$^{-3}$, although the Schechter function provides a poor fit to the data at high and low luminosity, while the luminosity density in the $K$ band is found to be $j = (6.305 \\pm 0.067) \\times 10^8$\\,L$_\\odot\\,h$\\,Mpc$^{-3}$. However, we caution that there are various known sources of incompleteness and uncertainty in our results. Using mass-to-light ratios determined from the optical colours we estimate the stellar mass function, finding good agreement with previous results. Possible improvements are discussed that could be implemented when extending this analysis to the full LAS. ",
        "introduction": "It is possible to learn much about a population's history by taking a census of the present-day population. With a large and increasing number of deep, large-area surveys taking place, a census of the low-redshift galaxy population may be undertaken. Deep imaging allows many different properties to be studied simultaneously.  Much can be learned about galaxy formation and evolution by investigating these properties: how they correlate with each other, and the sub-populations that exist.  The advantages of working at low redshift ($z \\sim 0.1$) are as follows: (1) a more complete sample may be studied, including galaxies with low luminosity or low surface brightness, (2) the more luminous galaxies may be studied in more depth, investigating morphology and structure as well as luminosity and colour and (3) the evolution and selection effects that plague high-redshift surveys are less of a problem.  The advantages (in principle) of building such a census on near-infrared (NIR) observations are well known. First, mass-to-light ($M/L$) ratios in the NIR are largely insensitive to galaxy or stellar type, certainly much less than in the optical \\citep{BelldJ2001}. This means that the NIR light is a good tracer of the total stellar mass in a galaxy. Moreover, the range of $M/L$ ratios is much smaller in the NIR, so uncertainties in the stellar mass are much smaller. Not only does this mean that a survey limited in NIR magnitude will be approximately limited in `apparent stellar mass', but also that morphological measurements in the NIR, for example the S\\'ersic index and the half-light radius, will reflect the distribution of stellar mass within the galaxy, whereas such measures in the optical will be significantly biased by the presence of young stellar populations. The $K$-band galaxy luminosity function (LF) is, for these reasons, a convenient quantity for numerical or semi-analytic models to predict: see e.g.\\ \\citet{Croton...2006}, \\citet{Bower...2006}, \\citet{deLuciaB2007} and \\citet*{BertonedLT2007}.  A second advantage is that the $K$-corrections in the $K$ band are also relatively independent of galaxy type \\mbox{\\citep{Mannucci...2001}}, leading to smaller uncertainties in the absolute magnitudes.  A third advantage is that dust is much less of a problem in the NIR than in the optical.  This means that, whereas optical measurements of galaxy properties are affected by dust obscuration, and therefore strongly dependent on the inclination of the galaxy, producing a smoothing of the galaxy LF, this is not such a problem in the NIR \\citep{Driver...2007b,Maller...2009}.  However, the main disadvantage of the NIR (for ground-based telescopes) is the sky brightness, which is around 13.5\\,mag\\,arcsec$^{-2}$ in $K$ for the data used here \\citep{Dye...2006}.  There have been several studies of the low-redshift $K$-band LF using Two-Micron All-Sky Survey (2MASS) imaging combined with various redshift surveys \\citep{Cole...2001,Kochanek...2001b,Bell...2003c,Eke...2005,Jones...2006}. Other estimates have been made by \\citet*{MobasherSE1993}, \\citet{Szokoly...1998} and \\citet{Loveday2000}, using optically selected samples, and by \\citet{Glazebrook...1995b}, \\citet{Gardner...1997} and \\citet{Huang...2003}, using samples selected in the $K$ band. Table \\ref{tbl:samplesize} shows the sample size of these $K$-band LF estimates.  An accurate detemination of the local $K$-band LF also has great value as a baseline for comparison with studies at higher redshift \\citep[e.g.][]{Cirasuolo...2007}. The principal uncertainty remaining at low redshift is connected with the low-luminosity end of the LF.  Moreover, there has been significant discussion about possible low-surface-brightness incompleteness in 2MASS \\citep{Andreon2002a}, which would affect the low-luminosity end of the LF.  \\begin{table} \\caption[Sample sizes of $K$-band galaxy LFs]{\\label{tbl:samplesize} Sample sizes of $K$-band galaxy luminosity functions.} \\vspace{.2in} \\centering \\begin{tabular}{lr} \\hline \\hline Paper & Number of galaxies in sample \\\\ \\hline \\citet{MobasherSE1993} & 181 \\\\ \\citet{Glazebrook...1995b} & 124 \\\\ \\citet{Gardner...1997} & 510 \\\\ \\citet{Szokoly...1998} & 175 \\\\ \\citet{Loveday2000} & 345 \\\\ \\citet{Kochanek...2001b} & 3878 \\\\ \\citet{Cole...2001} & 5683 \\\\ \\citet{Huang...2003} & 1056 \\\\ \\citet{Bell...2003c} & 6282 \\\\ \\citet{Eke...2005} & 15\\,644 \\\\ \\citet{Jones...2006} & 60\\,869 \\\\ This work & 40\\,111 \\\\ \\hline \\end{tabular} \\end{table}  Here we present the first statistical study of galaxies in the UKIRT Infrared Deep Sky Survey (UKIDSS) Large Area Survey, while leaving detailed investigation of surface brightness completeness and the very faint-end of the LF to future work.  This paper is organized as follows. Details about the sample used are found in Section 2. The method of analysing the data is described in Section 3. In Section 4, the bivariate brightness distribution (BBD) and LF are presented. The stellar mass function (SMF) is estimated in Section 5. There is a discussion in Section 6, followed by the conclusions in Section 7.  For ease of comparison with previous results, a flat cosmological model with $\\Omega_\\mathrm{M}=0.3$ and $\\Omega_\\Lambda=0.7$ is used, with $H_0 = 100h$\\,km\\,s$^{-1}$\\,Mpc$^{-1}$. AB magnitudes are used for Sloan Digital Sky Survey (SDSS) magnitudes and Vega magnitudes for $K$-band quantities.  For reference, AB and Vega magnitudes are related in the $r$ band by $r_\\mathrm{AB} = r_\\mathrm{Vega} + 0.146$ and in the $K$ band by $K_\\mathrm{AB} = K_\\mathrm{Vega} + 1.900$ \\citep{Hewett...2006}. ",
        "conclusions": "We have presented the first statistical analysis of galaxies from the UKIDSS LAS\\@.  The $1/V_\\mathrm{max}$ space density estimator has been used in a four-dimensional form to produce results for the $K$- and $r$-band LF and the SMF consistent with previous findings, with a form similar to a Schechter function with an almost flat faint-end slope.  We have presented the first $K$-band BBD in $K$-band absolute magnitude and effective surface brightness. This shows similar trends to the optical BBD: a correlation between luminosity and surface brightness, with a broadening of the surface brightness distribution at low luminosity and a flattening of the luminosity--surface brightness relation at high luminosity.  The multiple limits on the survey have been taken into account. For example, limits in $K$-band and $r$-band magnitude, $K$-band Petrosian radius and $K$-band surface brightness have been used to estimate the volume within which each galaxy would have been visible.  When the sample is subdivided according to colour, we have found a clear distinction between two populations, one of red, high-luminosity and high-surface-brightness galaxies, and the other of blue, low-luminosity and low-surface-brightness galaxies. This agrees well with previous investigations of the bimodality of the galaxy population.  We have considered various sources of incompleteness of the data and find that in order to obtain better results for low luminosity galaxies with UKIDSS, various selection effects need further attention.  A stable multivariate estimator for the space density needs to be developed that takes consideration of large-scale structure.  Incompleteness at low surface brightness also needs to be addressed, perhaps using source extraction with elliptical apertures, and the colour-dependent selection effects could be addressed by a $K$-band limited redshift survey to extend the SDSS spectroscopic sample to a greater depth in $K$. Physical properties of galaxies could be investigated further with good NIR and optical--NIR colours and/or extrapolated model magnitudes."
    },
    "0806/0806.2807_arXiv.txt": {
        "abstract": "White dwarfs are the end product of the lifes of intermediate- and low-mass stars and their evolution is described as a simple cooling process.  Recently, it has been possible to determine with an unprecedented precision their luminosity function, that is, the number of stars per unit volume and luminosity interval. We show here that the shape of the bright branch of this function is only sensitive to the averaged cooling rate of white dwarfs and we propose to use this property to check the possible existence of axions, a proposed but not yet detected weakly interacting particle.  Our results indicate that the inclusion of the emission of axions in the evolutionary models of white dwarfs noticeably improves the agreement between the theoretical calculations and the observational white dwarf luminosity function. The best fit is obtained for $m_a \\cos^2 \\beta \\approx 5$~meV, where $m_a$ is the mass of the axion and $\\cos^2 \\beta$ is a free parameter. We also show that values larger than 10 meV are clearly excluded.  The existing theoretical and observational uncertainties do not yet allow the confirmation of the existence of axions, but our results clearly show that if their mass is of the order of few meV, the white dwarf luminosity function is sensitive enough to detect their existence. ",
        "introduction": "One solution to the strong CP problem of quantum chromodynamics is the Peccei-Quinn symmetry (Peccei \\& Quinn 1977a, 1977b). This symmetry is spontaneously broken at an energy scale that gives rise to the formation of a light pseudo-scalar particle named the ``axion'' \\citep{wei78,wil78}.  This scale of energies is not defined by the theory but it has to be well above the electroweak scale to ensure that the coupling between axions and matter is weak enough to account for the lack of a positive detection up to now. The mass of axions and the energy scale are related by $m_a\\approx0.6(10^7~{\\rm GeV}/f_a)$~eV. Astrophysical and cosmological arguments \\citep{raf07} have been used to constrain this mass to the range $10^{-2}~{\\rm eV}\\geq m_a \\geq 10^{-4}~{\\rm eV}$.  For this mass range, axions can escape from stars and act as a sink of energy.  White dwarfs are the final evolutionary phase of low- and intermediate-mass stars ($M\\leq10\\pm2~M_{\\sun}$).  Since they are degenerate objects, they cannot obtain energy from thermonuclear reactions and their evolution can be described just as a gravothermal process of cooling.  Therefore, if axions exist, the properties of these stars would be noticeably perturbed. Furthermore, white dwarfs have a relatively simple structure:  a degenerate core that contains the bulk of the mass and acts as an energy reservoir and a partially degenerate envelope that controls the energy outflow.  The vast majority of white dwarfs have masses in the range $0.4\\leq M/M_{\\odot}\\leq 1.05$ --- although these figures are still uncertain --- and have a core made of a mixture of carbon and oxygen.  All of them are surrounded by a thin helium layer, with a mass ranging from $10^{-2}$ to $10^{-4}~M_{\\odot}$ which, in turn, is surrounded by an even thinner layer of hydrogen with a mass between $10^{-4}$ and $10^{-15}~M_{\\sun}$, although about 25\\% of white dwarfs do not have hydrogen atmospheres.  White dwarfs displaying hydrogen in their spectra are called DA and the remaining ones are known as non-DAs. Because of the different opacities, DA white dwarfs cool more slowly than the non-DA ones.  The standard theory of white dwarf cooling can be summarized as follows \\citep{ise98}. When the luminosity is large, $M_{\\rm bol}<8$, the evolution is dominated by neutrino emission.  In this phase the main uncertainties come from our poor knowledge of the initial conditions.  Fortunately, it has been shown that all the initial thermal structures converge toward a unique one \\citep{dan89}. For smaller luminosities, $8\\leq M_{\\rm bol}\\leq12$, the main source of energy is of gravothermal origin. In this phase, the Coulomb plasma coupling parameter is not large and the cooling can be accurately described.  Furthermore, the energy flux through the envelope is controlled by a thick nondegenerate or partially degenerate layer with an opacity dominated by hydrogen, when present, and helium, and it is weakly dependent on the metal content since metals sink towards the base of the envelope by gravitationally induced diffusion.  Below these luminosities, white dwarfs evolve into a region of densities and temperatures where the plasma crystallizes.  When this happens, two additional sources of energy appear. The first one is the release of latent heat during crystallization. The second one is the release of gravitational energy induced by phase separation of the different chemical species (Garc\\'{\\i}a-Berro et al.~1988a, 1988b; Isern et al.~1997, 2000).  When the bulk of the star is solid the white dwarf enters into the Debye cooling phase and the only important source of energy comes from the compression of the outer layers.  These late phases of cooling are not yet well understood \\citep{ise98}. ",
        "conclusions": "Figure \\ref{axion3} displays several luminosity functions obtained using different axion masses, adopting a constant star formation rate and an age of the Galactic disk of 11 Gyr.  As already mentioned, it is important to realize that the bright branch of the luminosity function is not sensitive to these last assumptions.  All the luminosity functions have been normalized to the luminosity bin at $\\log (L/L_{\\sun})\\simeq-3$ or, equivalently, $M_{\\rm bol} \\simeq 12.2$.  The best-fit model --- namely that which minimizes the $\\chi^2$ test in the region $-1>\\log (L/L_{\\sun})>-3$ (that is, $ 7.2 < M_{\\rm bol} < 12.2$), which is the region where both the observational data and the theoretical models are reliable --- is obtained for $m_a \\cos^2 \\beta\\approx 5$ meV and solutions with $m_a \\cos^2 \\beta > 10$ meV are clearly excluded.  Figure \\ref{axion4} displays the behavior of $\\chi^2$ as a function of the mass of the axion in our fiducial case (solid line) and in the case in which we use the initial-final mass relationship of \\cite{wood}, which is marginally compatible with the white dwarf mass distribution \\citep{cat08}.  In both cases the behavior of the luminosity function is similar and gives similar values for the mass of the axion once one takes into account the present uncertainties.  It is also important to notice that the largest contribution to the lack of accuracy comes from the brightest bins of the luminosity function, which have large error bars.  \\begin{figure} \\vspace{8cm} \\special{psfile=fig4.eps hscale=57 vscale=57 voffset=-100 hoffset=-35} \\caption{Value of $\\chi^2$ as a function  of the mass of the axion for the  case in  which the  initial--final mass  relationship of \\cite{cat08}  (solid line)  and that  of  \\cite{wood} (dotted line) are used.} \\label{axion4} \\end{figure}  This result is completely compatible with the previously existing constraints \\citep{raf07}.  Furthermore, these values are also compatible with the bounds imposed by the drift of the pulsational period of the ZZ Ceti star G117$-$B15A \\citep{ise92,cor01,bis08}.  It is worthwhile to mention here that axions with $m_a \\approx 5$ meV would change the expected period drift of variable DB white dwarfs --- which have values between $\\dot P \\sim 10^{-13}$ and $10^{-14}$ s~s$^{-1}$ \\citep{cor04} --- by a factor 2, the exact value depending on the adopted temperature of the stellar core and that the detection of such a drift would provide a strong additional argument in favor of the existence of axions.  The results presented here are not a definite proof of the existence of axions, since there are still some observational and theoretical uncertainties. However, the calculations reported here show that the hot branch of the white dwarf luminosity function is a powerful tool to test the existence of weakly interacting massive particles because it is only sensitive to the averaged cooling rate of white dwarfs and not to the details of the star formation rate or the initial mass function, as shown in \\S 2.  Moreover, our results are problably the first evidence that the shape of the white dwarf luminosity function could be affected by the emission of axions and that this change can be measured. If this is indeed the case the mass of the axion would be of about 5 meV. In addition, we have derived an upper bound to the mass of the axions of 10 meV, which is compatible with other recent determinations \\citep{bis08}.  It is worth mentioning that this result is relevant for other research fields and, in particular, for cosmology.  Specifically, assuming $\\cos^2\\beta = 1$, the contribution of axions to dark matter would be of the order of 0.2\\% \\citep{raf07}.  In view of this, we consider it of the largest importance to improve the observational determination of the white dwarf luminosity function, especially in the region of the hottest white dwarfs (see Fig.~\\ref{axion3}).  Thus, the extension of the SDSS is of the maximum interest not only for astronomers and cosmologists, but also for particle physicists.  However, not only observational efforts are needed, since in order to obtain a reliable determination of the mass of the axion it is also important to decrease the uncertainties in the plasmon neutrino emission rates at the relevant temperature range.  Furthermore, it would be also convenient to intensify the study of the drift of the pulsational periods of variable white dwarfs in order to obtain additional independent evidence."
    },
    "0806/0806.4610_arXiv.txt": {
        "abstract": "Infrared selection is a potentially powerful way to identify heavily obscured AGN missed in even the deepest X-ray surveys.  Using a 24~\\micron-selected sample in GOODS-S, we test the reliability and completeness of three infrared AGN selection methods: (1) IRAC color-color selection, (2) IRAC power-law selection, and (3) IR-excess selection; we also evaluate a number of infrared excess approaches. We find that the vast majority of non-power-law IRAC color-selected AGN candidates in GOODS-S have colors consistent with those of star-forming galaxies. Contamination by star-forming galaxies is most prevalent at low 24~\\micron\\ flux densities ($\\sim 100$~\\microjy) and high redshifts ($z\\sim 2$), but the fraction of potential contaminants is still high ($\\sim 50\\%$) at 500~\\microjy, the highest flux density probed reliably by our survey.  AGN candidates selected via a simple, physically-motivated power-law criterion (\\plagas), however, appear to be reliable.  We confirm that the infrared excess methods successfully identify a number of AGN, but we also find that such samples may be significantly contaminated by star-forming galaxies.  Adding only the secure \\textit{Spitzer}-selected \\plaga, color-selected, IR-excess, and radio/IR-selected AGN candidates to the deepest X-ray--selected AGN samples directly increases the number of known X-ray AGN (84) by $54-77\\%$, and implies an increase to the number of 24~\\micron-detected AGN of $71-94\\%$. Finally, we show that the fraction of MIR sources dominated by an AGN decreases with decreasing MIR flux density, but only down to $f_{\\rm 24~\\micron} = 300$~\\microjy.  Below this limit, the AGN fraction levels out, indicating that a non-negligible fraction ($\\sim 10\\%$) of faint 24~\\micron\\ sources (the majority of which are missed in the X-ray) are powered not by star formation, but by the central engine.  The fraction of all AGN (regardless of their MIR properties) exceeds 15\\% at all 24~\\micron\\ flux densities. ",
        "introduction": "Identifying complete and reliable samples of AGN has become a necessity for extragalactic surveys, whether the goal be the selection of AGN candidates or the removal of AGN ``contaminants''.  Only when armed with complete samples of AGN will we be able to determine the role of obscured accretion in the build-up of the present day black hole mass function, or accurately characterize the star-formation history of the universe.  Complete AGN samples are also required to test proposed evolutionary theories in which black hole formation and star-formation are intimately linked by merger and feedback processes (e.g. Hopkins et al. 2006), ultimately producing the correlation between black hole mass and bulge velocity dispersion (Ferrarese \\& Merritt 2000; Gebhardt et al. 2000). Unfortunately, the varied luminosities, accretion rates, orientations, and intrinsic obscurations of AGN prevent any one selection technique from reliably identifying all of them.  For instance, while current UV, optical, and X-ray surveys are capable of detecting unobscured AGN, they miss many of the obscured AGN and nearly all of the Compton-thick AGN thought to dominate AGN number counts at both low and high redshift (e.g. Gilli et al. 2007; Daddi et al. 2007a,b).  Likewise, only 10-15\\% of AGN are radio-loud, making radio surveys relatively incomplete.  The Multiband Imaging Photometer (MIPS; Rieke et al. 2004) and Infrared Array Camera (IRAC; Fazio et al. 2004) instruments aboard \\textit{Spitzer} have provided sensitive surveys in multiple mid-IR bands. Infrared selection with MIPS and IRAC data is being used widely to select AGN candidates independently of their optical and/or X-ray properties.  In addition to identifying AGN in fields with little or no X-ray data, infrared selection criteria are capable of identifying heavily obscured AGN missed in even the deepest X-ray fields (e.g. Donley et al. 2007). As such, IR selection has the potential to complement traditional AGN selection methods and to yield a more complete census of AGN activity.  In this paper, we critically review the following infrared selection criteria: (1) IRAC color-color selection, (2) IRAC power-law selection, and (3) IR-excess selection.  The first selection method employs color cuts in two representations of IRAC 4-band mid-infrared (MIR) color-color space (Lacy et al. 2004; Stern et al. 2005), the second identifies AGN whose IRAC SEDs are well-fit by a power-law (Alonso-Herrero et al. 2006, Donley et al. 2007), and the third selects red galaxies with large infrared to UV/optical flux ratios (Daddi et al. 2007a, Dey et al. 2008, Fiore et al. 2008, Polletta et al. 2008).  The first two criteria are based on the same principle: the hot dust near an AGN's central engine reprocesses absorbed UV, optical, and X-ray emission into short-wavelength MIR emission, filling in the gap between the stellar emission that peaks near 1.5~\\micron\\ and the long-wavelength dust emission features that dominate the SEDs of star-forming galaxies.  The color-color and power-law selection criteria, however, differ in the range of mid-IR characteristics they include as possible AGN indicators.  The third selection method identifies sources in which heavy obscuration with reemission in the infrared diminishes the optical emission and/or enhances the infrared emission.  This paper utilizes improved spectral templates with a sample of infrared color-selected, power-law galaxies (\\plagas), and IR-excess galaxies in the ultra-deep GOODS-S field to test the reliability and completeness of these selection techniques over a wide range of sample properties.  From this analysis, we then quantify the contribution of these approaches plus \\textit{Spitzer} identification of radio-intermediate and radio-loud AGN to the X-ray--selected AGN population.  The paper is organized as follows.  In \\S2, we describe the selection of the color-selected, \\plaga, and IR-excess samples. The construction of high-reliability photometric redshifts is described in \\S3, as are the overall redshift properties of the sample.  In \\S4, we briefly discuss the X-ray properties of our MIPS-selected sample.  The infrared color selection criteria of Lacy et al. (2004) and Stern et al. (2005) are discussed in \\S5, where we compare and contrast the two selection criteria, investigate the behavior in color space of the star-forming templates, determine the redshift and flux dependencies of the color selection techniques, and investigate the properties of the most secure color-selected AGN candidates.  In \\S6, we discuss the \\plaga\\ selection criteria, and in \\S7 we investigate the IR-excess sources. Finally, in \\S8, we discuss the overall statistics of AGN revealed by IR-based methods compared with X-ray--selected samples.  A summary is given in \\S9.  Throughout the paper, we assume the following cosmology: ($\\Omega_{\\rm m}$,$\\Omega_{\\rm \\Lambda},H_0$)=(0.3, 0.7, 72~km~s$^{-1}$~Mpc$^{-1}$). ",
        "conclusions": "Infrared selection of AGN is a powerful technique.  Using new accurate star-forming and AGN templates along with a flux-limited MIPS-selected sample drawn from the GOODS-S field, we critically review three MIR selection criteria: (1) the IRAC color cuts of Lacy et al. (2004) and Stern et al. (2005), (2) the power-law galaxy (\\plaga) selection technique of Alonso-Herrero et al. (2006) and Donley et al. (2007), and (3) the IR-excess selection criteria of Daddi et al. (2007a,b), Dey et al.  (2008), Fiore et al. (2008), and Polletta et al. (2008). From this analysis, we then quantify the contribution of \\textit{Spitzer}-selected AGN to the X-ray selected AGN population. The main conclusions of this paper are as follows:  \\begin{itemize}  \\item The majority of non-power-law IRAC color-selected AGN candidates have IR colors consistent with those of redshift-appropriate star-forming templates. In comparison, the majority of \\plaga\\ AGN candidates lie outside of the star-forming contours.  PLG selection recovers the majority of high-quality AGN candidates.  \\item The reliability of AGN IRAC color-color selection improves with increasing flux as high-redshift star-forming galaxies fall out of the sample.  Nevertheless, the fraction of potential star-forming contaminants is still high ($\\sim 50\\%$) at the highest fluxes probed by our survey ($f_{\\rm 24~\\micron} \\sim 500$~\\microjy).  \\item A comparison of the 24~\\micron\\ to 3.6~\\micron\\ colors of the X-ray non-detected \\plagas\\ to those of AGN and star-forming templates suggests that the X-ray non-detected \\plagas, like their X-ray--detected counterparts, have more hot dust emission than can be explained by star-formation alone.  \\item An analysis of the Daddi et al. (2007) IR-excess sources in our MIPS sample indicates that while these sources may be Compton-thick AGN, it is also possible that they are low-luminosity, Compton-thin AGN and/or luminous, highly-reddened star-forming galaxies.  \\item An X-ray stacking analysis of the sources selected via the Fiore et al. (2008) criteria indicate that $\\sim 42\\%$ of the sources are consistent with being obscured AGN, and that the remaining 58\\% may be star-forming galaxies.  \\item Adding secure \\textit{Spitzer}-selected power-law, color-color, radio/IR, and IR-excess AGN candidates to the deepest X-ray samples directly increases the number of known AGN by $\\sim 54-77\\%$, and implies a total increase of $71-94$\\%.  This fraction excludes the full contributions from the Daddi et al. and Fiore et al. AGN candidates, whose nature is still uncertain.  \\item The fraction of MIR sources dominated by an AGN decreases with decreasing flux density, but only down to a 24~\\micron\\ flux density of $\\sim 300$~\\microjy.  Below this limit, the AGN fraction levels out at $\\sim 10\\%$. This indicates that a non-negligible fraction of faint 24~\\micron\\ sources are primarily powered not by star-formation, but by the central engine.  In addition, the majority of AGN at low 24~\\micron\\ flux densities are missed in the X-ray, indicating that X-ray emission alone cannot be used to identify AGN, especially amongst faint IR samples.  \\end{itemize}  \\vspace*{0cm}"
    },
    "0806/0806.1490_arXiv.txt": {
        "abstract": "We present results obtained from a series of observations of the supernova remnant RX~J1713.7$-$3946 by the {\\it Suzaku} satellite  which  cover about two-thirds of the remnant surface.  Hard X-rays have been detected from each pointing up to $\\sim 40~{\\rm keV}$. The hard X-ray spectra are described by power-law functions with photon indices of $\\sim 3.0$, which are larger than those in the energy region below 10~keV. The combination of the spatially-integrated XIS and HXD spectra clearly reveals a spectral cutoff in the  X-ray spectrum which  is linked to the maximum energy of accelerated electrons  emitting synchrotron radiation. The broad-band coverage of {\\it Suzaku} observations  from 0.4 keV to 40 keV   allows  us to derive,  for the first time, the energy spectrum of parent  electrons in the cutoff region. The inferred cutoff energy in the spatially-integrated X-ray spectrum indicates that the electron acceleration in the remnant proceeds  close to the  Bohm-diffusion limit. We discuss  implications of the spectral and morphological properties of {\\it Suzaku} data in the context of the origin of   nonthermal emission.  The {\\it Suzaku}  X-ray and the H.E.S.S.  TeV gamma-ray  data together hardly can be  explained within a  pure leptonic scenario, unless we introduce an additional  component of relativistic  electrons with softer energy spectrum. Moreover, the leptonic models  require  very weak magnetic field  which does not agree with the recently discovered filamentary structure and  short-term variability  features of the X-ray emitting region. The  hadronic  models with strong magnetic field provide perfect fits to the  observed X-ray and TeV gamma-ray spectra through the synchrotron radiation of electrons and {\\it p-p} interactions of  protons,   but require special  arrangements of model parameters  to explain the lack of thermal  component of X-ray emission.  For the morphology studies, we compare the X-ray an TeV gama-ray surface brightness maps using the {\\it Suzaku} XIS and the H.E.S.S.  data. We confirm  the previously reported strong correlation between X-ray and  TeV gamma-ray emission  components. At the same time the {\\it Suzaku}  data  reveal a deviation  from the general tendency,  namely,   the X-ray emission in the western rim regions  appears  brighter  than expected  from the average X-ray to gamma-ray  ratio. ",
        "introduction": "Supernova remnants (SNRs) have long been considered to be likely acceleration sites of cosmic-ray particles below the energy of the {\\it knee}, $\\sim 10^{15}~{\\rm eV}$. The energy supply to explain the energy density of comic rays is satisfied if $\\sim 1$--10\\% of the energy of each supernova is transferred to accelerated particles. Also, the well developed theory of diffusive shock acceleration nicely explains the universal power-law spectrum of cosmic rays (e.g. \\citealt{BE87,malkov01}). Although synchrotron emission detected in the radio band supports this idea observationally, no evidence of acceleration to TeV energy had been observed until recently. During the last decade, such evidence was revealed through observations of X-rays and TeV gamma rays from several shell-type SNRs. \\cite{koyama95} discovered synchrotron X-rays from the shell of SN~1006, which indicates electrons are accelerated up to multi-TeV energies. This finding was followed by detections of synchrotron X-rays from other SNRs, including RX~J1713.7$-$3946 (e.g. \\citealt{koyama97,slane01}). Further evidence for multi-TeV particles (electrons and/or protons) has been provided by discovery of TeV gamma rays from some SNRs, such as Cassiopeia~A \\citep{aha01} or RX~J1713.7$-$3946 \\citep{muraishi00}, although their spectral parameters and morphologies were not well determined due to the limited sensitivity of TeV observatories. Subsequently, high quality morphological and spectral studies have been performed by H.E.S.S. (e.g. \\citealt{aha04,aha06,aha07}). These pioneering measurements by the H.E.S.S. telescope, together with the high resolution X-ray data, have enabled direct comparison of X-ray and TeV gamma-ray data.  The shell-type SNR RX~J1713.7$-$3946 (also known as G347.3$-$0.5), is one of the best-studied SNRs from which both non-thermal X-rays and TeV gamma rays are detected. This SNR was discovered in soft X-rays during the {\\it ROSAT} All-Sky Survey \\citep{pfe96}. The {\\it ASCA} satellite, with wider energy coverage than that of {\\it ROSAT}, revealed that the X-ray spectrum is featureless and can be best interpreted as synchrotron emission from very high energy electrons in the TeV regime \\citep{koyama97,slane99}. The X-ray spectrum was well fitted with a power-law function of photon index $\\Gamma = 2.2$--$2.4$ and  interstellar absorption column density $N_{\\rm H} = 0.6$--$0.8 \\times 10^{22}~{\\rm cm}^{-2}$ without any observable evidence for a thermal emission component. Subsequent observations by {\\it Chandra} and {\\it XMM-Newton} have unveiled structure with a complex network of bright filaments and knots, in the western part of the SNR \\citep{uchi03,laz04,cassam04,hiraga05}.  TeV gamma-ray emission from RX~J1713.7$-$3946 was first reported by the CANGAROO collaboration in 1998 \\citep{muraishi00}, and confirmed by the subsequent observations with CANGAROO-II in 2000 and 2001 \\citep{enomoto02}. Later, the H.E.S.S. collaboration obtained  a resolved  image of the source in TeV gamma rays \\citep{aha04} showing that the gamma-ray emission from RX~J1713.7$-$3946 arises mainly in the shell. These observations revealed a striking correlation between the X-ray and the gamma-ray images, which indicates a strong connection between the physical processes responsible for X-ray and TeV gamma-ray emission components \\citep{aha06}.  Based on the spectral and morphological information, they discussed two possible gamma-ray emission scenarios, one where gamma rays are generated by inverse Compton scattering of accelerated electrons with diffuse radiation  fields  (the so-called leptonic scenario) and the other where the decay of  secondary $\\pi^{0}$-mesons is  responsible for gamma rays (hadronic scenario). The later observations with H.E.S.S.  revealed that the flux  extends to 30~TeV and, likely, beyond,  which implies particle acceleration up to energies  well above 100~TeV for either model \\citep{aha07} .  Most recently, our X-ray observations using {\\it Chandra} and {\\it Suzaku} have provided important clues for understanding the acceleration process in the SNR. From a series of observations of the northwest part of the SNR with {\\it Chandra} in 2000, 2005 and 2006, we discovered that compact regions of the northwest (NW) shell are variable in flux on a one-year time scale \\citep{uchi07}. The fast variability was interpreted as one-year scale acceleration and synchrotron cooling of electrons with amplified magnetic fields of order of 1~mG. Such a large magnetic field in compact regions strongly favors $\\pi^0$-decay emission as the origin of TeV gamma rays. Also, thanks to the wide-band coverage of {\\it Suzaku} and its low background level, we were able to measure a hard X-ray spectrum up to 40~keV from the southwest portion of RX~J1713.7$-$3946 with a clear indication of a high-energy cutoff in the synchrotron spectrum (\\citealt{takahashi07}, hereafter Paper I). Combined with the upper limit on a shock speed of $4500~{\\rm km}~{\\rm s}^{-1}$ placed by {\\it Chandra}, the cutoff energy determined by the  {\\it Suzaku} observation of the southwest part  indicates that particle acceleration within the SNR shock is so efficient that it approaches the theoretical limit corresponding to  the so-called Bohm diffusion regime (e.g. \\citealt{malkov01}).   In this paper, we present results of mapping observations of RX~J1713.7$-$3946 with {\\it Suzaku}, which covers about two-thirds of the SNR region with 11 pointings. The low background level of the Hard X-ray Detector (HXD) enables us to detect hard X-ray emission up to $\\sim 40~{\\rm keV}$ from each of the pointings. At the same time, its small field-of-view (FoV) of $\\sim 25^{\\prime} \\times 25^{\\prime}$ FWHM gives us information about the spatial distribution of hard X-ray emission and spectral differences from region to region. Thanks to its low instrumental background and large effective area, the other detector system aboard {\\it Suzaku}, the X-ray Imaging Spectrometer (XIS), also uncovers new observational facts, such as spectral features below 10~keV and the morphology of relatively dim regions left unclear in previous studies by {\\it ASCA}, {\\it Chandra}, and {\\it XMM-Newton}. By combining the XIS and HXD spectra summed over the data from all the pointings, we show a wide-band X-ray spectrum (0.4--40~keV) with quite high statistics, with which we investigate not only the existence of a cutoff, but also its shape. We then compare the cutoff shape obtained with theoretical predictions.  In Section 2, we describe our {\\it Suzaku} observations and the data reduction procedures. Analysis and results of HXD and XIS data are shown in \\S3.1 and \\S3.2, respectively. We present the wide-band spectrum by connecting the XIS and the HXD data in \\S3.3. A detailed study regarding the cutoff structure is also given there. Section 4 is devoted to multi-wavelength spectral and morphological studies. The results obtained are discussed in the following section, and the results are finally summarized. Throughout this paper, we assume that the  distance to RX~J1713.7$-$3946 is close  to  1~kpc as proposed by \\cite{koyama97} based on the $N_{\\rm H}$ value. A similar distance has been claimed based on the NANTEN CO data \\citep{fukui03,moriguchi05}. The typical age of the remnant for such a distance is estimated of order of 1000~yr, which can be an indication of  association of RX~J1713.7$-$3946  with an explosion in A.D. 393 as proposed by \\cite{wang97}. ",
        "conclusions": "\\subsection{Cutoff in the Synchrotron Spectrum} We conducted a series of {\\it Suzaku} observations which covers about two-thirds of the surface of SNR RX~J1713.7$-$3946. Through the data analysis, we successfully detected signals up to $\\sim 40$~keV from each of the pointings. The HXD spectra above 10~keV are significantly steeper than those obtained from the XIS below 10~keV, suggesting that a spectral cutoff is common throughout the remnant. By combining the XIS and HXD spectra, we obtained a wide-band spectrum with high statistics, which clearly shows a cutoff around 10~keV.  Taking advantage of the high photon statistics, we performed a detailed study of the cutoff shape and compared it with a recent theoretical prediction by \\cite{zira07}. A sharp cutoff of the accelerated electron spectrum is needed to reproduce the cutoff shape in the synchrotron spectrum detected with {\\it Suzaku}.  The spectrum of electrons derived from {\\it Suzaku} data is in good agreement with the analytical model of  \\cite{zira07}.   The cutoff energy in the spectrum of synchrotron radiation contains an important information about the efficiency of diffusive shock acceleration. For  acceleration in the Bohm diffusion regime and when energy losses of electrons are dominated by synchrotron cooling,  the cutoff energy,  $\\varepsilon_0$ in equation (\\ref{za07_sync_spec})  is expressed as \\citep{zira07} \\begin{eqnarray} \\varepsilon_0 = 0.55 \\left( \\frac{v_s}{3000~{\\rm km}~{\\rm s}^{-1}} \\right)^2 \\eta^{-1}~{\\rm keV}, \\end{eqnarray} where $v_s$ is the shock speed and $\\eta~(\\ge1)$ is the so-called ``gyrofactor''. The case of  $\\eta = 1$ corresponds to the  ``Bohm limit'', and implies  high level of turbulence $\\delta B \\sim B$. The {\\it Suzaku} spectrum is characterized  by the best-fit  parameter $\\varepsilon_0 = 0.67~{\\rm keV}$ which gives $v_s =  3300 \\eta^{1/2}~{\\rm km}~{\\rm s}^{-1}$. Here, we assume that the shock speed $v_s$ is uniform throughout the remnant, which is supported by the fact that the outer boundary of the X-ray morphology is nearly circular. The upper-limit of the shock speed $v_s \\le 4500~{\\rm km}~{\\rm s}^{-1}$ derived from  the   {\\it Chandra} data \\citep{uchi07}, results in   $\\eta \\le 1.8$. This is a strong  evidence of acceleration of electrons in the regime close to the Bohm limit. Note that a similar result was obtained for the SW rim of the remnant in \\cite{uchi07}. Here we confirm this conclusion for a larger area of the remnant with higher statistics.  \\subsection{Multi-wavelength Spectrum}  While there is little doubt in the  synchrotron origin of broad-band X-ray emission measured by {\\it Suzaku},  the X-ray spectrum alone does not give preference to the strength of the magnetic field in the region of production of synchrotron radiation. Formally, the field can be as small as $10~\\mu{\\rm  G}$ and as large as $100~\\mu{\\rm G}$. Meanwhile, the strength of the magnetic field has dramatic impact on the origin of TeV gamma-rays. The so-called leptonic or inverse Compton models require magnetic field between $10~\\mu{\\rm G}$  and $15~\\mu{\\rm G}$.  Even so, it is difficult to achieve, at least within a simple one-zone model,  a  satisfactory explanation of both X-ray and TeV gamma-ray spectral features, unless we invoke an extremely high diffuse radiation field of optical  photons  to enhance the IC gamma-radiation below 1~TeV (see  Figure~\\ref{fig:SED_leptonic_opt}). A more realistic approach  for  explanation of  the broad-band TeV gamma-ray spectrum  within IC models  can be realized under the assumption of existence an additional,  low-energy electron component in the shell (see Figure~\\ref{fig:SED_leptonic_2ele}). Even so, the most serious  problem for IC models remains the requirement of low magnetic field in the gamma-ray production region,  in contrast to large magnetic field required to explain the fast  variability of X-ray emission on small scales. Formally, one may assume that gamma-rays are mainly produced in \"voids\", i.e. in regions  with very low magnetic field. This would imply quite inhomogeneous distribution of the magnetic field in the shell. One the other hand, the  observed strong X-ray and TeV correlation within the IC models can be explained  only in the case of homogeneous distribution of  magnetic field.  The large-scale magnetic fields on parsec scales with an average  strength larger than $\\geq 15~\\mu{\\rm G}$  make the IC gamma-ray production inefficient, and thus give preference to the so-called hadronic models of gamma-rays produced at interactions of accelerated protons with the ambient gas via production and decay of secondary $\\pi^0$-mesons. What concerns X-rays, they are produced,  as in leptonic models, by synchrotron radiation  of directly accelerated electrons. This is demonstrated in  Figure~\\ref{fig:SED} for  very strong magnetic field, $B=200~\\mu{\\rm G}$. Note that  while comparing the model predictions with measurements in the radio band, one should take into account that the radio points shown correspond to measurements of NW rim, while the X-ray and gamma-ray points are for the entire remnant. If the ratio of the radio flux from the NW rim to that from the whole remnant is not much different  from  the corresponding ratio in X-rays, the flux from the whole SNR should be significantly larger. This would reduce the difference between the measurements and predictions. In any case,  the radio flux can be significantly reduced assuming somewhat smaller magnetic field or harder electron spectrum.  Indeed, in Figure \\ref{fig:SED_100uG} we show model calculations performed for a magnetic field  $B=100~\\mu{\\rm G}$. While the synchrotron X-ray flux is described perfectly as before (in Figure~\\ref{fig:SED}), the radio flux is  by a factor of four lower; at 1.4~GHz it is 34~Jy, which is close  to the latest estimates of radio flux from the whole remnant based on observations with ATCA and the 30-m radio telescopes of IRA (F. Acero et al. 2008, in preparation; G. Dubner 2008, private communications).   \\begin{figure} \\epsscale{1.0} \\plotone{./f20.eps} \\caption{The SED of RX~J1713.7$-$3946 from radio to X-ray with a synchrotron model curve calculated by assuming $B = 100~\\mu{\\rm G}$. } \\label{fig:SED_100uG} \\end{figure}   The  radio flux  can be suppressed  even for the ambient field larger than $100~\\mu {\\rm G}$,  provided that the electron injection spectrum is harder than $E^{-2}$. Figure~\\ref{fig:SED_17} demonstrates this  possibility, where we assume an electron/proton index of $s = 1.7$  which corresponds to a compression ratio of $\\sigma = 5.3$. Note that $\\sigma$ can exceed the adiabatic upper limit of 4 as described by \\cite{bere06}.  Note that the value of $s = 1.7$ is consistent with the conclusion of  \\cite{villa07}  based on semi-analytical derivation of the  parent proton  spectrum  from the H.E.S.S. data. In model calculations shown  in Figure~\\ref{fig:SED_17}, the spectrum of protons requires  an ``early''  exponential cutoff at $E_{p0} = 25.0~{\\rm TeV}$.  Note that formally the spectral index $s=1.7$ implies  shock acceleration in non-linear regime which in fact  predicts some deviation from pure power-law distribution  of accelerated particles (see e.g. \\citealt{Ellison07,bere06}). This  would lead to further reduction of the radio flux.  The convection of low energy electrons could be another reason for low radio flux. Note that the escape of electrons through convection has strong impact only on low energy electrons; because of fast synchrotron cooling,  the effect of escape is negligible for multi-TeV electrons. A significant quantity of low-energy electrons can escape from the  shell of the SNR before emitting radio photons. Therefore, the radio flux can be reduced while the X-ray flux will remain unchanged.  \\begin{figure} \\epsscale{1.0} \\plotone{./f21.eps} \\caption{The SED of RX~J1713.7$-$3946 with a hadronic model when the index of the electron/proton spectrum is 1.7. The magnetic field is $200~\\mu{\\rm G}$. The line styles for the IC spectra are the same as those in Figure~\\ref{fig:SED} the total energy of electrons is $W_e = 6.0 \\times 10^{45}~{\\rm erg}$. The total proton energy is $W_p = 1.6 \\times 10^{50}~(n/1~{\\rm cm}^{-3})^{-1}~{\\rm erg}$. } \\label{fig:SED_17} \\end{figure}  For calculations  shown in Figure~\\ref{fig:SED}, the total energy of electrons is estimated as $W_e =  3.1 \\times 10^{46}~(d/1~{\\rm kpc})^2~{\\rm erg}$,  and the energy for protons as $W_p = 2.7 \\times 10^{50}~(n/1~{\\rm cm}^{-3})^{-1}(d/1~{\\rm kpc})^2~{\\rm erg}$. In the case of harder energy spectra with power-law index $s=1.7$ corresponding to Figure~\\ref{fig:SED_17},  one has $W_e =  6.0 \\times 10^{45}~(d/1~{\\rm kpc})^2~{\\rm erg}$,  and $W_p = 1.6  \\times 10^{50}~(n/1~{\\rm cm}^{-3})^{-1}(d/1~{\\rm kpc})^2~{\\rm erg}$, respectively.  The proton/electron ratio in either case is very small, $K_{ep} \\leq 10^{-4} (n/1~{\\rm cm}^{-3})$. This value is significantly smaller  than that for directly observed local cosmic rays ($K_{ep} \\sim 0.01$),  unless a large ambient matter density of $n \\sim 100~{\\rm cm}^{-3}$ is assumed. \\cite{katz08} and \\cite{butt08} argued that the hadronic scenario for this SNR has difficulties because the $K_{ep}$ value should be consistent with local cosmic rays and other SNRs. However, the $K_{ep}$ value of one SNR at a fixed age does not necessarily need to agree with the local cosmic ray value. The low-energy electrons are likely produced in later stages of SNR evolution, when the $K_{ep}$ value can be different from the present value. A comparison with other SNRs should be performed with care, as well.  Cutoff energy in the gamma-ray spectrum should give an important hint whether SNRs are sources of cosmic rays below the {\\it knee}, if gamma rays observed by H.E.S.S. have hadronic origins. \\cite{plaga08} argued that the cutoff energy of the H.E.S.S. spectrum of RX~J1713$-$3946 is around 18~TeV, which can be translated to proton energy more than 10 times below the energy of the {\\it knee}. Indeed our multi-wavelength study requires a cutoff in the proton spectrum around 100~TeV or even less for hard acceleration spectra of protons. However, one should take into account that even in the case of effective acceleration the highest energy protons beyond 100 TeV escape the source in a quite short time scales, and hence do not contribute to the gamma-ray production at the present epoch \\citep{ptuskin05, gabici07}.  A  unique feature of RX~J 1713.7-3946 is the lack of thermal X-ray emission. Recently, \\cite{katz08} and \\cite{butt08}  interpreted this fact as an argument against the hadronic model for TeV gamma rays. Generally it is true that plasma in young supernova remnants is heated to high temperatures observed via thermal X-ray emission of hot electrons. However, one should take into account that we deal with a unique object, and the lack of thermal X-ray emission cannot {\\it a priori} be invoked as an argument against the hadronic origin  of the observed TeV gamma rays.  It is important to note that in SNR shocks the formation of high plasma temperatures with $kT_i=3/16~m_i v_s^2$ is relevant only to protons (ions), and that a high ion temperature does not automatically (from first principles) mean a high electron temperature. In fact, the only known heating process of thermal electrons is Coulomb collisions between electrons and protons (ions), which, however, has too long time scale to establish electron-proton equipartition. On the other hand, we do know from X-ray observations that the electrons in young SNRs are heated to keV temperatures.  This can be explained by assuming that a hypothetical mechanism, most likely related to the energy exchange through excited plasma waves, is responsible for effective electron heating in SNRs. As long as the nature of this mechanism in collisionless shocks remains unknown, one cannot predict, even qualitatively, the specifics of its operation on a source by source basis.  We indeed deal with two interesting facts. First, many young SNRs, like Tycho and Cassiopeia A, with intense thermal X-ray emission and intense nonthermal radio emission emit little (or do not emit at all) TeV gamma rays. On the other hand, RX~J1713.7$-$3946, with lack of (or rather very low)  thermal X-ray emission and with relatively weak nonthermal radio emission, is a source of powerful TeV radiation. These two facts can be treated as a hint for low efficiency of establishing equipartition in the thermal plasma in very effective TeV particle accelerators like RX~J1713.7$-$3946. It is interesting to note in this regard, that such a tendency is found also for another effective TeV accelerator --- SN 1006  \\citep{vink03}. Whether the reduction of the exchange rate between different particle species in thermal plasma has a link to the particle acceleration in high Mach number shocks, as proposed by \\cite{vink03}, is a very interesting question to be explored in future deep theoretical and phenomenological studies. In this regard, RX~J1713.7$-$3946 can serve as a key ``template'' source for such studies.  \\cite{hughes00} also discussed low electron temperature based on {\\it Chandra} observations of a young SNR in the Small Magellanic Could, 1E~0102.2$-$7219. They measured a blast-wave velocity of $\\sim 6000~{\\rm km}~{\\rm s}^{-1}$ from the expansion rate and predicted the electron temperature of $kT_e > 2.5~{\\rm keV}$ by considering Coulomb heating. However, the electron temperature derived from their spectral analysis is 0.5~keV, which is far below the prediction. According to their discussion, not only electron heating but also ion heating is suppressed and substantial fraction of energy may be going into cosmic-ray production due to the non-linear effects in the shock.  The nonlinear shock acceleration in this object can convert a significant, up to  $f  \\sim  0.5$  fraction of the  kinetic energy of explosion into relativistic particles. Correspondingly the fraction of  available energy which goes to the heating of the ambient  plasma will be reduced  $1-f \\sim 0.5$. Yet, conservative estimates show that plasma in RX~J1713.7$-$3946 can be heated to quite high temperatures even the heating of  electrons and  protons proceeds only through the Coulomb exchange. This question recently has studied by \\cite{Ellison07} for a standard SNR of age $t_{\\rm SNR}=500$~yr and  energy $E_{\\rm SN}=10^{51}$~erg. In particular, it has been shown that in the case of effective diffusive shock acceleration and the plasma density $n=0.1~{\\rm cm}^{-3}$, the ratio of synchrotron luminosity to thermal (bremsstrahlung) luminosity can be as large as 100. This implies that in the case of RX~J1713.7$-$3946,  from which thermal X-ray emission is not observed, the plasma density cannot  significantly exceed $0.1~{\\rm cm}^{-3}$ (the luminosity of thermal bremsstrahlung is proportional to $n^2$).  The low density of order of $0.1~{\\rm cm}^{-3}$  reduces the parameter space for hadronic models but does not exclude it. Indeed, as mentioned above,  the TeV gamma-ray flux  of RX~J1713.7$-$3946 can be explained  by interactions of protons if the parameter $A=(W_p/10^{50}~{\\rm erg}) (n/1~{\\rm cm}^{-3})  (d/1~{\\rm kpc})^{-2}$ exceeds 1.5   to 3, depending on the spectrum of protons. Assuming that  more than $30\\%$ of the explosion energy of this SNR is released in accelerated protons, and that  the plasma in the  gamma-ray production region is compressed by a factor of few, we find that the  located of the source at a distance of about 1 kpc would marginally support the hadronic model.  While  closer location of the source would make the model requirements quite viable and flexible, the location of  the source beyond 1~kpc hardly can be accommodated within a standard shock acceleration scenario. If the SNR is closer, the distance of $d \\gtrsim 0.5~{\\rm kpc}$ seems reasonable considering the upper limit on the shock speed of $4500(d/1~{\\rm kpc})~{\\rm km}~{\\rm s}^{-1}$ by \\cite{uchi07}.  Finally we should mention the model  suggested by \\cite{Malkov05} which can naturally explain both the low synchrotron flux at radio frequencies and lack of thermal X-ray emission of  RX~J1713.7$-$3946. The standard scenarios of gamma-ray production in SNRs assume that radiation is produced in downstream where the density of both  relativistic particles and thermal plasma is higher than in upstream.  However,  in the  cases when the shock is expanding into a low-density  wind bubble and approaching  cold dense material, e.g. swept-up shell or  surrounding molecular clouds, the gamma-radiation is contributed predominantly from upstream. While the energy distribution of accelerated particles downstream is coordinate-independent in both linear and nonlinear regimes, the  particle distribution upstream  is coordinate-dependent. Because of energy-dependent diffusion coefficient, the high-energy particles diffuse ahead of low-energy  particles, thus a dense material adjacent upstream will ``see'' relativistic particles (protons and electrons) with low-energy exponential cutoff,  $E_{\\rm min}$,  which depends on the location of the  dense regions. This implies that the effective production of TeV  gamma-rays (from {\\it p-p} interactions) and X-rays (from synchrotron radiation of TeV electrons) will be not accompanied by low energy (GeV) gamma-rays and synchrotron radio emission. Obviously, this model is not constraint by lack of thermal emission.  \\subsection{Morphology} In addition to the spectral information, the comparison of X-ray and TeV gamma-ray images presented in Figure~\\ref{fig:plot_keV_vs_TeV} helps us to draw the physical picture of RX~J1713.7$-$3946. Let us first discuss the tight correlation observed in most parts of the remnant. Within the hypothesis of hadronic origin of gamma-rays, the gamma-ray flux is proportional to the number densities of the ambient matter and relativistic protons, while the X-ray flux is proportional to number density of electrons. If the matter distribution significantly varies throughout the SNR, we need fine parameter tuning among the matter distribution, the electron injection rate, and the proton injection rate, in order to produce the tight correlation. Therefore, more natural explanation is that the matter density is uniform and the injection rate of the electrons and that of the protons are proportional to each other.  The X-ray flux excess along the NW and SW rims provides a unique probe of recent acceleration activity. Let us consider a \"toy\" model and compare its predictions with the observational results. In the toy model, the injection rate of electrons and protons keeps constant but increases by a factor of 1.5 in the last 10~yr only at the NW and SW rims. What is important here is the difference of cooling time between electrons and protons. The synchrotron cooling time of an electron emitting synchrotron photons with energy of $\\varepsilon$ is given as equation (\\ref{eq:sync_cooling_time}), while the cooling time of protons due to $p$-$p$ interactions is expressed as \\begin{eqnarray} t_{pp} = 5.3 \\times 10^7 \\left( \\frac{n}{1~{\\rm cm}^{-3}} \\right)^{-1}~{\\rm yr} \\end{eqnarray} and is almost energy-independent. For the magnetic field $B=100 \\  \\rm \\mu G$, the cooling times of  electrons emitting 2~keV and 5~keV X-rays are $12~{\\rm yr}$ and $7.6~{\\rm yr}$, respectively. For any reasonable density of  ambient gas, the cooling time of protons is much longer than the lifetime of this SNR.  While the synchrotron X-rays we observe at present are emitted by electrons accelerated during the last $\\sim 10~{\\rm yr}$. the  flux of  $\\pi^0$-decay gamma rays is provided by protons accelerated throughout the lifetime of the SNR. Figure~\\ref{fig:plot_keV_vs_TeV_toymodel} shows the scatter plot of $F_{\\rm keV}$ and $F_{\\rm TeV}$ expected from the toy model, which shows a similar distribution to the observational results in Figure~\\ref{fig:plot_keV_vs_TeV}. The recent active acceleration increases the X-ray flux while keeping the gamma-ray flux almost unchanged.   \\begin{figure} \\epsscale{1.0} \\plotone{./f22.eps} \\caption{The same plot as Figure~\\ref{fig:plot_keV_vs_TeV} but results obtained by assuming the \"toy\" model described in the text.} \\label{fig:plot_keV_vs_TeV_toymodel} \\end{figure}"
    },
    "0806/0806.3947_arXiv.txt": {
        "abstract": "The physical nature of the X-ray/radio correlation of AGN is still an unsolved question. High angular resolution observations are necessary to disentangle the associated energy dynamics into nuclear and stellar components. We present MERLIN/EVN 18~cm observations of 13 X-raying AGN. The sample consists of Seyfert 1, Narrow Line Seyfert 1, and LINER-like galaxies. We find that for all objects the radio emission is unresolved and that the radio luminosities and brightness temperatures are too high for star formation to play an important role. This indicates that the radio emission in these sources is closely connected to processes that occur in the vicinity of the central massive black hole, also where the X-ray emission is believed to originate in. ",
        "introduction": "Among current problems in AGN research are the feedback between the central engine and the host galaxy environment \\cite{page_submillimeter_2001}. Related to this is the observation that only about 10\\% of the local galaxy population display Seyfert characteristics \\cite{ho_search_1997}, which raises questions like: (i) How is the nuclear engine fueled? (ii) What is the relationship between Seyfert activity and circumnuclear star formation? In order to address these questions it is important to disentangle the energy dynamics around the nucleus at high angular resolution. The NIR-to-X-ray spectral energy distributions (SEDs) of Seyfert 1 galaxies are quite flat \\cite{2004ASPC..311...37W}. The X-rays are believed to originate close to the accretion disk in a kind of hot corona \\cite{1978Natur.272..706S}. Only a small fraction of X-rays can be attributed to a stellar component in AGN \\cite{2005ApJ...631..707P}. On the other hand, extended radio emission is most likely produced via Synchrotron processes in radio jets or supernovae \\cite{1984ARA&A..22..319B,1992ARA&A..30..575C}. The core radio emission generally exhibits a flat, self-absorbed spectrum, whereas the spectral slope of radio jets is steep. Furthermore, there appears to be a tight correlation between X-ray and radio flux, both for AGN and star-forming galaxies \\cite{2007A&A...467..519P}. Such a correlation suggests a common physical nature that links the emission mechanisms of both phenomena. It is, however, not clear what this link is, since both types of radiation are believed to originate on different physical scales (accretion  disk vs. jet). One solution for flat spectrum point sources is possible provided that the emitting region is hot and optically thin \\cite{blundell_origin_2007}. In this case, optically thin bremsstrahlung from a slow, dense disk wind can contribute significantly to the observed levels of nuclear emission. In this case, radio and X-ray emission originate from the same regions and the correlation can be attributed to a common disk origin.  The spatial distribution of the radio emission can, therefore, tell us something about the origin and relative importance of the radio emission. Extended but collimated emission is linked to a jet, diffuse emission is usually attributed to star formation, and compact nuclear emission to processes close to the supermassive black hole (SMBH).  \\begin{figure}[t!] \\begin{center} \\includegraphics[width=12cm]{zuther_j_fig_1.pdf} \\end{center} \\caption{Absolute $i$-band PSF magnitude as function of redshift. The dashed, horizontal line marks the demarcation between Seyfert galaxies and QSOs. The color coding of the symbols corresponds to the $(g-i)_\\mathrm{PSF}$ color of the targets. The dotted circles indicate sources that are members of galaxy clusters. Numbers are target identifiers \\citep[cf][]{zuther08}.} \\label{fig:MagZDist} \\end{figure} ",
        "conclusions": "We have presented initial results from a radio-interferometric study of X-ray bright AGN. The detection rate is high and most of the sources are unresolved on the MERLIN (500~pc) scale. The three EVN observations are unresolved on the 40~pc scale and contain all the flux of the MERLIN measurements. Our observations show that the radio emission in these kind of sources is compact and likely related to the nuclear accretion processes, even considering the various spectral classes of objects and the connection of part of the X-rays to hot cluster gas in three cases.  Future observations at other radio frequencies can provide further spectral information that can be incorporated in models of the nuclear emission processes \\citep[e.g.,][]{blundell_origin_2007}. AO-assisted near-infrared follow-up studies, furthermore, have the power to complement the radio studies on scales comparable to that of MERLIN and can shed light on possible small-scale jet/interstellar medium interaction in these sources \\citep[e.g.,][and references therein]{2007A&A...466..451Z}.  \\ack This work is partly funded by the Deutsche Forschungsgemeinschaft via grant SFB 494. This work has benefited from research funding from the European Community's sixth Framework Programme under RadioNet R113CT 2003 5058187. We are grateful to Tom Muxlow and Giuseppe Cimo for their support on the data reduction."
    },
    "0806/0806.2085_arXiv.txt": {
        "abstract": "s{We show that the dragging of axis directions of local inertial frames by a weighted average of the energy currents in the universe (Mach's postulate) is exact for all linear perturbations of all Friedmann-Robertson-Walker universes and for all types of matter.} ",
        "introduction": " ",
        "conclusions": ""
    },
    "0806/0806.3610_arXiv.txt": {
        "abstract": "In this paper, we use the outer-galactic HI scale height data as well as the observed rotation curve as constraints to determine the halo density distribution of the Andromeda galaxy (M31). We model the galaxy as a gravitationally-coupled system of stars and gas, responding to the external force-field of a known Hernquist bulge and the dark matter halo, the density profile of the latter being characterized by four free parameters. The parameter space of the halo is optimized so as to match the observed HI thickness distribution as well as the rotation curve on an equal footing, unlike the previous studies of M31 which were based on rotation curves alone. We show that an oblate halo, with an isothermal density profile, provides the best fit to the observed data. This gives a central density of 0.011 M$_{\\odot}$ pc$^{-3}$, a core radius of 21 kpc, and an axis ratio of 0.4. The main result from this work is the flattened dark matter halo for M31, which is required to match the outer galactic HI scale height data. Interestingly, such flattened halos lie at the most oblate end of the distribution of halo shapes found in recent cosmological simulations.\\\\ ",
        "introduction": "It is well-known that the dark matter halo plays an important role in the dynamics of galaxies, especially in the outer regions (Binney \\& Tremaine 1987). Since a galactic disk is rotationally supported, the  rotation curve serves as a useful tracer of the gravitational potential in the plane of the galaxy. The observed rotation curve is routinely used to deduce the mass distribution in a galaxy and hence its dark matter content (e.g., Begeman 1987, Kent 1986, 1987, Geehan et al. 2006). The thickness of the gas layer, on the other hand, depends on the vertical gravitational force and traces the potential perpendicular to the mid-plane (e.g., Narayan \\& Jog 2002 a). In this work, we use the rotation curves as well as the radial distribution of the thickness of the HI gas layer in the outer galaxy to study the shape and density profile of the dark matter halo in M31. In  a disk plus bulge plus halo model of an external  galaxy, the disk and the bulge can be mostly studied observationally. Therefore, the rotation curve and the vertical HI scale height data effectively complement each other to determine the dark matter halo distribution of a galaxy uniquely.  In the past, the idea of studying the dark matter halo properties by using the outer galactic HI flaring data has been used to explore tha halos of NGC 4244 (Olling 1996), NGC 891 (Becquaert \\& Combes 1997), and the Galaxy (Olling \\& Merrifield 2000, 2001). However, the HI scale height distribution was mainly used to constrain the oblateness of the halo, and not its other parameters such as the power-law index. In some cases, the gas gravity and even the stellar gravity was ignored (Becquaert \\& Combes 1997) in determining the net galactic potential, and hence the gas scale height distribution.  These issues were taken care of in determining the Galactic halo parameters by Narayan et al (2005). Using the gravitationally-coupled, 3-component Galactic disk model (Narayan \\& Jog 2002 b), various density profiles of the halo were investigated, and an attempt was made to obtain the halo parameters, which provided the best fit (in the least square sense) to the observed HI scale height distribution. Finally, conformity with the shape of the observed rotation curves was used to remove the degeneracies in the best-fit values obtained by the first constraint. Also, unlike some of the previous models, the self-gravity of the gas was included in the analysis. From their study Narayan et al. (2005) concluded that a spherical halo, with a density falling off more rapidly than an isothermal halo,  provides the best fit to the available data. This study was based on the HI scale height data then available upto 24 kpc from Wouterloot et al. (1990). Kalberla et al. (2007) have confirmed this by using their recent extended HI scale height data upto 40 kpc, and have also included a dark matter ring which they claim is needed to explain the observed HI scale height distribution in the Galaxy.  In this paper, we apply the above approach to investigate the dark matter halo properties of the Andromeda galaxy (M31 or NGC 224). Here we use both the rotation curve and the HI scale height data as rigorous constraints simultaneously and scan the entire parameter space systematically so as to obtain the best-fit halo parameters. In addition to the various density profiles, we also try to fit various shapes of the halo which was not done by Narayan et al. (2005). Earlier studies on M31  (Widrow et al. 2003, Widrow \\& Dubinski 2005, Geehan et al. 2006, Seigar et al. 2006, Tamm et al. 2007) were mostly aimed at developing a complete mass model (disk plus bulge plus halo), based on comparisons made with the available structural and kinematical data (surface brightness profiles, bulge-velocity dispersion relations, rotation curves), which assumed a spherical-shaped halo. On the other hand, we have studied the dark matter halo profile, and  show it to be flattened.  In Section 2, we describe our model, and in  Section 3 discuss the numerical calculations involved, and the input parameters used. In Section 4, we present the results and analysis of the numerical results; followed by the discussion, and conclusions in Sections 5 and 6 respectively. ",
        "conclusions": "We have used both the observed rotation curve and the outer galactic HI scale height data to constrain the dark matter halo profile of M31. We have systematically explored various shapes and power-law indices for the density distributions for the halo to fit the observed data. Our galactic disk model consists of coupled stars and HI gas, where the gas-gravity is taken into account on an equal footing with the stellar gravity. We find that an oblate isothermal halo with a central density of 0.011 M$_{\\odot}$ pc$^{-3}$, a core radius of 21 kpc best fits the observations. The axis ratio for the best-fit results is 0.4. This is in a sharp contrast to the spherical halo used to model M31 in the literature so far. The rotation curve constraint alone is usually used which determines the mass within a radius but cannot uniquely determine the shape of the halo. The present work highlights the fact that using the two simultaneous and complementary constraints of rotation curve and the HI scale height data in the outer galactic region, allows one to identify the shape as well as the density distribution of the dark matter halo in spiral galaxies.  We stress that the availability of more data points for the HI scale heights in the outer galaxy beyond 5-6 disk scale lengths, and an accurate determination of the HI gas velocity dispersion, would provide a tighter  constraint for the shape and the density profile of the dark matter halo. In fact, having such data for other galaxies would allow the above method to be applied to a systematic study of the dark matter halo properties in different galaxies."
    },
    "0806/0806.1938_arXiv.txt": {
        "abstract": "The origin of the rapid quasi-periodic variabilities observed in a number of accreting black hole X-ray binaries is not understood. It has been suggested that these variabilities are associated with diskoseismic oscillation modes of the black hole accretion disk. In particular, in a disk with no magnetic field, the so-called g-modes (inertial oscillations) can be self-trapped at the inner region of the disk due to general relativistic effects. Real accretion disks, however, are expected to be turbulent and contain appreciable magnetic fields. We show in this paper that even a weak magnetic field (with the magnetic energy much less than the thermal energy) can modify or ``destroy'' the self-trapping zone of disk g-modes, rendering their existence questionable in realistic black hole accretion disks. The so-called corrugation modes (c-modes) are also strongly affected when the poloidal field approaches equal-partition. On the other hand, acoustic oscillations (p-modes), which do not have vertical structure, are not affected qualitatively by the magnetic field, and therefore may survive in a turbulent, magnetic disk. ",
        "introduction": "In recent years, quasi-periodic variability has been observed from a number of Galactic compact X-ray binary systems. Of particular interest are several accreting black hole (BH) binaries which show pairs of quasi-periodic oscillations (QPOs) of fixed frequencies having ratios close to $2:3$ (for example, GRO J1655-40 shows $f=300,~450$~Hz; see Remillard \\& McClintock 2006 for a review). The origin of these QPOs is not understood. The fixed frequencies and frequency ratio led to the suggestion that these QPOs involve certain nonlinear resonant phenomenon in the disk (e.g., coupling between the radial and vertical epicyclic oscillations of the disk fluid element; Kluzniak \\& Abramowicz 2002), but so far no fluid dynamical model producing these resonances has been developed (see Rebusco 2008 and references therein). Alternatively, it has been suggested that these QPOs may arise from acoustic oscillations in an accretion torus (Rezzolla et al. 2003; Lee, Abramowicz \\& Kluziniak 2004), with the oscillation frequencies determined by the (unknown) radial extent of the torus. Perhaps the theoretically most developed model for QPOs is the relativistic diskoseismic oscillation model (Kato \\& Fukue 1980; Okazaki et al.~1987; Nowak \\& Wagoner 1991; see Wagoner 1999; Kato 2001 for reviews), in which general relativistic (GR) effect produces trapped g-mode (also called inertial mode or inertial-gravity mode) oscillations in the inner region of the disk. Various theoretical issues related to this model have been studied, such as the role of corotational wave absorption (Kato 2003; Li, Goodman \\& Narayan 2003; Silbergleit \\& Wagoner 2007) and resonant mode excitations due to global disk deformation (Kato 2008; Ferreira \\& Ogilvie 2008).  The studies of the oscillation modes of disks/torii, such as those mentioned above, usually assume that the unperturbed flow is laminar and has no magnetic field. Real accretion disks, on the other hand, are highly turbulent due to the nonlinear development of magnetorotational instability (MRI) (see Balbus \\& Hawley 1998 for a review).  The question therefore arises as to how the MRI-driven turbulence affects the oscillation modes obtained from hydrodynamical models and to what extent these trapped modes remain ``valid'' in a realistic situation. Arras, Blaes \\& Turner (2006) attempted to address this issue by carrying out MHD simulations in the shearing-box geometry. They showed that axisymmetric standing sound waves give rise to distinct peaks in the temporal power spectrum, while inertial waves do not. The discrete frequencies obtained by them were due to the imposed periodic boundary conditions adopted in the simulations, and not due to any relativistic effect.  Arras et al. suggested that their result posses a serious problem for QPO models based on g-modes (inertial waves). Recently, Reynolds \\& Miller (2008) reported on the results of global simulations of BH accretion disks (using Paczynski-Wiita pseudo-Newtonian potential) and showed that, while axisymmetric g-mode oscillations manifest in the hydrodynamic disk with no magnetic field, they disappear in the magnetic disk where MHD turbulence develops.  In this paper, we study analytically the effects of magnetic fields on the relativistic diskoseismic modes in accretion disks around BHs. We consider both poloidal and toroidal fields and use local analysis of the full MHD equations to examine how the magnetic field changes the radial wave propagation diagrams for various modes.  We show that the trapping region of g-modes can be easily ``destroyed'' even when the disk field strength is such that the associated Alfv\\'{e}n speed is much smaller than the sound speed. On the other hand, the propagation characteristics of p-modes (acoustic oscillations) and c-modes are largely unchanged. Note that since we assume laminar flows for our unperturbed disks, we do not directly address the effects of turbulence on disk modes. However, we believe that our work is relevant to this issue, since magnetic fields naturally arise in a turbulent disk.  We summarize the basic MHD equations in \\S 2 and review the properties of diskoseismic modes important for our analysis in \\S 3. We examine in \\S 4 and \\S5 the effects of poloidal field and toroidal field on those modes, respectively, and discuss the implications of our result in \\S 6. ",
        "conclusions": "In this paper we have studied the effects of both poloidal and toroidal magnetic fields on the diskoseismic modes in BH accretion disks. Previous works by Kato, Wagoner and others have been based on hydrodynamic disks with no magnetic field. The key finding of our paper is that the g-mode self-trapping zone (which arises from GR effect) disappears when the disk contains even a small poloidal magnetic field, corresponding to $v_{Az}/c_s =0.01-0.1$ (see Fig.~2-3; $v_{Az}$ is the Alfv\\'{e}n speed and $c_s$ is the sound speed). It is well-known that the combination of a weak poloridal field and differential rotation gives rise to MRI, making real astrophysical disks turbulent. Earlier numerical simulations indicated that the magnetic field grows as MRI develops, until it saturates at $v_{Az}/c_s \\sim 0.1-1$, with the toroidal field stronger than the poloidal field by a factor of a few (see, e.g., Hawley et al. 1996; Balbus \\& Hawley 1998). Recent simulations showed that the turbulent state depends strongly on the net magnetic flux through the disk (e.g., Fromang \\& Papaloizou 2007; Simon, Hawley \\& Beckwith 2008). In any case, it is likely that the magnetic field in a turbulent disk is large enough to ``destroy'' the g-mode self-trapping zone.  Thus, the g-mode properties (including the frequencies and excitations) derived from hydrodynamical models are unlikely to be applicable to real BH accretion disks. The disappearance of the g-mode trapping zone might also explain why Arras et al. (2006) and Reynolds \\& Miller (2008) did not see any global g-modes in their MHD simulations.  As mentioned in \\S1, g-mode oscillations have been considered a promising candidate to explain QPOs in BH X-ray binaries. Theoretically, these modes are appealing because in hydrodynamic disks their existence depends on general relativistic effect and does not require special disk boundary conditions. Our analytical results presented in this paper, together with recent numerical simulations (Arras et al. 2006; Reynolds \\& Miller 2008), suggest that magnetic fields and turbulence associated with real accretion disks can change this picture significantly.  While g-modes can be easily modified or ``destroyed'' by magnetic fields, our analysis showed that p-modes are not affected qualitatively. The magnetic field simply changes the sound speed to the fast magnetosonic wave speed and leaves the p-mode propagation diagram unchanged. We also showed that a weak poloidal field ($v_{Az}/c_s \\ll 1$) does not affect the c-mode propagation zone, although a stronger field modifies it. Our results therefore suggests that global p-mode oscillation is robust and may exist in real BH accretion disks, provided that partial wave reflection at the disk inner edge can be achieved.\\footnote{Kato (2001) has discussed why such reflection may be possible.} Of particular interest is the non-axisymmetric p-modes, since they may be excited by instabilites associated with corotation resonance (Tsang \\& Lai 2008a, b)."
    },
    "0806/0806.4389_arXiv.txt": {
        "abstract": "We analyze the phenomenology of a prolonged  early epoch of matter domination by an unstable but very long-lived massive particle. This new matter domination era can help to relax some of the requirements on the primordial inflation. Its main effect is the huge entropy production produced by the decays of such particle that can dilute any possible unwanted relic, as the gravitino in supersymmetric models, and thus relax the constraints on the inflationary reheating temperature. A natural candidate for such heavy, long-lived particle already present in the Standard Model of the electroweak interactions would be a heavy right-handed neutrino. In this case, we show that its decays can also generate the observed baryon asymmetry with right-handed neutrino masses well above the bound from gravitino overproduction. ",
        "introduction": "Inflation was introduced in the 80s \\cite{inflation} as a solution to several problems of the big bang cosmology. Perhaps the main problem was the large-scale smoothness problem. Why different patches of the Universe, that were not in causal contact in the radiation last scattering era, have approximately the same  temperature today \\cite{Hinshaw:2008kr}?. In inflationary models, an epoch of exponential expansion inflates a small patch of the Universe in causal contact to contain all the observable Universe today. Simultaneously if the temperature after inflation is low enough, inflation helps also to dilute unwanted relics from higher scales and reduces the flatness problem.  It is usually assumed that some kind of inflation starts already at the Planck scale to avoid the Universe collapse in a few Planck times if $\\Omega>1$ or (for any $\\Omega$) to prevent the invasion of the surrounding inhomogeneity to our homogeneous patch before inflation. On the other hand, the scales observable today   in the cosmic microwave background left the horizon at an energy $V^{1/4} \\lsim 6 \\times 10^{16} $ GeV \\cite{Alabidi:2005qi}, or 60 $e$-foldings before the end of inflation. So that, inflation must end below this scale. After the end of inflation comes an era of reheating when the inflaton field oscillates around its minimum and decays to ordinary particles. The final reheating temperature where we recover ordinary big bang cosmology can take any value from $V^{1/4}$ above to scales as low as 1 MeV. However, the required dilution of unwanted relics, as GUT monopoles or gravitinos in supersymmetric models, forces the reheating temperature to be well below the GUT scale or even below $T_{RH} \\leq 10^{8}$ GeV in SUSY models.  In this letter, we propose a simple and economic mechanism that helps solving some of these problems and reduces the requirements on the primordial inflationary mechanism without further additions to the particle spectrum of the Standard Model {\\bf with right-handed neutrinos} \\footnote{From now on, we call the Standard Model with right-handed neutrinos simply ``Standard Model''}. After an initial inflationary epoch (still necessary to reproduce the observed correlation on temperature fluctuations at large scales) we assume our Universe is radiation dominated for a short period and then enters a matter domination era due to the existence of a heavy long-lived unstable particle that decays to radiation well-before nucleosynthesis, when we  connect with usual cosmology. In the Standard Model, as we will show, this role could be played by a heavy right-handed neutrino. In the literature, it is well-known that late time entropy release can help to ameliorate some of the problems of standard cosmology. However, so far most of these works have only considered moduli fields in supersymmetric theories (see for instance \\cite{Nagano:1998aa,Kawasaki:1999na,Kohri:2004qu,Pradler:2006hh,Kitano:2008tk}) and their real presence in nature could be considered more speculative than the existence of right-handed neutrinos.  As we show below, this matter domination mechanism, naturally embedded in the SM, is able to help primordial inflation in several aspects. A long period of matter domination can reduce mildly the number of e-folds before the end of inflation at which observable perturbations were generated, relaxing this way flatness conditions on the inflationary potential. Moreover, the large entropy production in the decay of this particle  completely  dilutes any unwanted relics, eliminating the constraint on the inflation reheating temperature. In this sense our matter domination epoch has the same advantages as thermal inflation \\cite{Lyth:1995hj}, without resorting to yet another scalar field and /or scalar potential. ",
        "conclusions": "In this work, we have shown that an early epoch of matter domination by a long-lived massive particle can help to solve some of the problems of primordial inflation. We have seen that the large entropy production generated by the decays of such particle can dilute unwanted relics from higher temperatures, relaxing the  constraints on the inflationary reheating temperature. In supersymmetric theories this mechanism can help to solve the gravitino problem. Moreover, a long period of matter domination reduces the number of $e$-foldings before the end of inflation at which the observable cosmological perturbations were generated. In the Standard Model a natural candidate for such heavy, long-lived particle is a heavy right-handed neutrino. For low enough mass of the lightest left-handed neutrino and neutrino Yukawa mixings sufficiently close to the PNMS mixing matrix, the right-handed neutrino dominates the energy density of the universe for a long time and generates a large amount of entropy in its decay.  In this case, we show that its decays can also generate the observed baryon asymmetry for right-handed neutrino masses well above the bound from gravitino overproduction."
    },
    "0806/0806.1279_arXiv.txt": {
        "abstract": "GRB 060124 is the first event that both prompt and afterglow emission were observed simultaneously by the three \\emph{Swift} instruments. Its main peak also triggered Konus-Wind and HETE-II. Therefore, investigation on both the temporal and spectral properties of the prompt emission can be extended to X-ray bands. We perform a detailed analysis on the two well identified pulses of this burst, and find that the pulses are narrower at higher energies, and both X-rays and gamma-rays follow the same $w - E$ relation for an individual pulse. However, there is no a universal power-law index of the $w - E$ relation among pulses. We find also that the rise-to-decay ratio $r/d$ seems not to evolve with $E$ and the $r/d$ values are well consistent with that observed in typical GRBs. The broadband spectral energy distribution also suggest that the X-rays are consistent with the spectral behavior of the gamma-rays. These results indicates that the X-ray emission tracks the gamma-ray emission and the emissions in the two energy bands are likely to be originated from the same physical mechanism. ",
        "introduction": "The \\emph{Swift} satellite (Gehrels et al., 2004) was successfully launched on 20th November 2004. It is a multi-wavelength observatory, covering the gamma-ray, X-ray and UV/optical bands. Thanks to its rapid repointing capability, the mission has revolutionized the Gamma-Ray Burst (GRB) observations in many aspects (for recent reviews, see M\\'{e}sz\\'{a}ros, 2006; Fox and M\\'{e}sz\\'{a}ros, 2006; Zhang, 2007). The prompt slewing capability of the X-Ray Telescope (XRT, Burrows et al., 2005) and UV-Optical Telescope (UVOT, Roming et al., 2005) allows the satellite to swiftly catch very early X-ray and UV/optical signals following the GRB prompt emission detected by the Burst Alert Telescope (BAT, Barthelmy et al., 2005).   In the pre-\\emph{Swift} era, the temporal and spectral behaviors of GRB prompt emission have been studied extensively. It is found that a pulse of the prompt gamma-rays at lower energy bands tend to be wider, which is roughly depicted as $w\\propto E^{-0.4}$ ($w - E$ relation; Fishman et al., 1992; Link et al., 1993; Fenimore et al., 1995; Norris et al., 1996, 2005; Piro et al., 1998; Costa, 1999; Nemiroff, 2000; Feroci et al., 2001; Crew et al., 2003; Peng et al., 2006; Zhang et al., 2007b). The photons at lower energy bands also lag behind that of the photons at higher energy bands (the so-called spectral lag; Cheng et al., 1995; Norris et al., 1996, 2005; Norris, Marani and Bonnell, 2000; Wu and Fenimore, 2000; Chen et al., 2005; Yi et al., 2006; Zhang et al., 2006b, 2006c; Peng et al., 2007).  Although many attempts have been made to explain the $w - E$ relation and the spectral lag behavior(e.g. Fenimore et al. 1995; Cohen et al., 1997; Chiang, 1998; Dermer, 1998; Kazanas et al. 1998; Piran, 1999; Wang et al., 2000; Nemiroff, 2000; Qin et al., 2004, 2005; Shen et al., 2005; Lu et al., 2006; Zhang et al., 2007b; Dado, Dar and De R\\'{u}jula, 2007), its nature remains a matter of active debate in the community. On the other hand, it is unclear whether this correlation can be extended to X-ray bands. The broadband observations showed that the X-ray emission of some GRBs have unusual properties. in't zand et al. (1999) found that the prompt X-ray emission of GRB 980519, measured by BeppoSAX, undergoes a strong soft-to-hard-to-soft evolution. An exceptionally intense gamma-ray burst, GRB 030329, was detected and localized by the instruments on board the High Energy Transient Explorer satellite (HETE). It's lightcurve has a distinct, bright, soft X-ray component(Vanderspek et al. 2004). A thermal emission component is identified from the XRT data of a nearby XRF 060218 (Campana et al. 2006), but its non-thermal X-rays are from the same emission component as the gamma-rays (Liang et al., 2006). Vetere et al. (2006) analyzed the X-ray temporal and spectral characteristics of ten GRBs detected by the WFCs on board BeppoSAX and argued that there exist two components (slow and fast) in the X-ray emission. These facts suggest that the physics of these prompt X-rays are also very uncertain.  With benefit from the Swift satellite, the prompt X-rays of some bursts were observed, e.g. GRB 050117 (Hill et al., 2006), GRB050713A (Morris et al., 2007), GRB050820A (Cenko et al., 2006), GRB 060124 (Romano et al., 2006), GRB 060218 (Campana et al., 2006; Liang et al., 2006); GRB 061121 (Page et al., 2007) and GRB 070129( Godet et al., 2007; Krimm et al., 2007). These panchromatic observations unveiled the unprecedented spectral and temporal information of GRB prompt emission. This also makes it possible to measure their temporal structures and to examine whether the well-known $w - E$ relation can be extended to X-ray bands for these bursts. Except for GRB 060124 and 060218, these bursts have very complicated temporal structures that consist of a series of overlapping pulses. Liang et al. (2006) analyzed the non-thermal emission of GRB 060218 from the gamma-ray to X-ray bands and obtained that $w\\propto E^{-0.31\\pm0.03}$, which roughly satisfies the $w- E$ relation and the relation between spectral lag and luminosity derived from typical GRBs (Fenimore et al. 1995; Norris et al., 1996, 2000, 2005), although it has the longest pulse duration and spectral lag observed to date among the observed GRBs. They suggested that the prompt X-rays and gamma-rays are from the same component. In this paper we present a detailed analysis on the prompt emission of GRB 060124 to examine whether the $w-E$ relation can be extended to X-rays. In Section 2, we present the data reduction. Results are given in Section 3. Our conclusions are presented in Section 4. ",
        "conclusions": "We have analyzed the temporal properties of the two well-identified pulses of GRB 060124 from X-ray (0.2-10 keV) to gamma-ray (15-1160 keV) energy bands. We find that the pulse width $w$ is energy-dependent for the two pulses in eight energy bands (0.2-1, 1-4, 4-10, 15-25, 25-50, 50-100, 100-350 and 300-1160 keV). The pulses are found narrower at higher energies, and both X-rays and gamma-rays follow the same $w - E$ relation for an individual pulse. However, we find that there is no a universal power-law index of the $w - E$ relation among pulses. We also find that the rise-to-decay ratio $r/d$ seems not to evolve with $E$ and the $r/d$ values are well consistent with that observed in typical GRBs. The peak fluxes for the two pulses become larger at higher energies. These results indicates that the X-ray emission tracks the gamma-ray emission and the emissions in the two energy bands are likely to be originated from the same physical mechanism.   One remarkable advance from {\\em Swift} is that the on-board XRT has established a large sample of X-ray lightcurves from tens of seconds to days (Zhang et al., 2006a; Nousek et al., 2006; O'Brien et al., 2006). The physical mechanisms of these X-rays are of great uncertain and is on debate in the GRB community (see review by Zhang, 2007). It is possible that the mechanisms are diverse (e.g., Zhang et al., 2006a; Zhang et al., 2007a, Liang et al., 2007, 2008). As we show here that some X-rays are possibly from the same emission component (see also Liang et al., 2006 for GRB 060218). However, Vetere et al. (2006) analyzed the temporal and spectral characteristics of X-rays for ten GRBs detected by the WFCs on board BeppoSAX and argued that there exist two components (slow and fast) in the X-ray emission. This feature actually is also seen in GRB 030329 (Vanderspek et al., 2004). Both temporal and spectral properties are critical to discriminate these components.   \\textbf"
    },
    "0806/0806.4953_arXiv.txt": {
        "abstract": "We present results from detailed general relativistic simulations of stellar core collapse to a proto-neutron star, using two different microphysical nonzero-temperature nuclear equations of state as well as an approximate description of deleptonization during the collapse phase. Investigating a wide variety of rotation rates and profiles as well as masses of the progenitor stars and both equations of state, we confirm in this very general setup the recent finding that a generic gravitational wave burst signal is associated with core bounce, already known as type~I in the literature. The previously suggested type~II (or ``multiple-bounce'') waveform morphology does not occur. Despite this reduction to a single waveform type, we demonstrate that it is still possible to constrain the progenitor and postbounce rotation based on a combination of the maximum signal amplitude and the peak frequency of the emitted gravitational wave burst. Our models include to sufficient accuracy the currently known necessary physics for the collapse and bounce phase of core-collapse supernovae, yielding accurate and reliable gravitational wave signal templates for gravitational wave data analysis. In addition, we assess the possiblity of nonaxisymmetric instabilities in rotating nascent proto-neutron stars. We find strong evidence that in an iron core-collapse event the postbounce core cannot reach sufficiently rapid rotation to become subject to a classical bar-mode instability. However, many of our postbounce core models exhibit sufficiently rapid and differential rotation to become subject to the recently discovered dynamical instability at low rotation rates. ",
        "introduction": "\\label{section:introduction}  The final event in the life of a massive star is the catastrophic collapse of its central, electron-degenerate core composed of iron-peak nuclei. When silicon shell burning pushes the iron core over its effective Chandrasekhar mass, collapse is initiated by a combination of electron capture and photo-disintegration of heavy nuclei, both leading to a depletion of central pressure support. Massive stars in the approximate mass range of about $ 10 $ to $ 100 $ solar masses ($ M_\\odot $) experience such a collapse phase until their homologously contracting~\\cite{goldreich_80_a, yahil_83_a} inner core reaches densities near and above nuclear saturation density where the nuclear equation of state (EoS) stiffens, leading to an almost instantaneous rebound of the inner core (core bounce) into the still supersonically infalling outer core. The hydrodynamic supernova shock is born, travels outward in radius and mass, but rapidly loses its kinetic energy to the dissociation of infalling iron-group nuclei and to neutrinos that deleptonize the immediate postshock material and stream off from these regions quasi-freely. The shock stalls, turns into an accretion shock and must be revived to produce the observable explosion associated with a core-collapse supernova. Mechanisms of shock revival are still under debate (a recent review is presented in~\\cite{janka_07_a}, but see also~\\cite{burrows_07_a, burrows_07_b, bruenn_06_a}) and may involve heating of the postshock region by neutrinos, multi-dimensional hydrodynamic instabilities of the accretion shock, in the postshock region, and/or in the proto-neutron star, rotation, magnetic fields, and nuclear burning. If the shock is not revived, black-hole formation (on a timescale of $ \\sim 1\\mbox{\\,--\\,}2 \\mathrm{\\ s} $~\\cite{liebendoerfer_04_a}) is inevitable and the stellar collapse event may remain undetected by conventional astronomy or, perhaps, appear as a gamma-ray burst if the progenitor star has a compact enough envelope and sufficiently rapid rotation in its central regions~\\cite{woosley_06_a, dessart_08_a}.  Conventional astronomy can constrain core-collapse supernova theory and the supernova explosion mechanism via secondary observables only, e.g., the explosion energy, ejecta morphology, nucleosynthesis yields, residue neutron star or black hole mass and proper motion, and pulsar magnetic fields. Neutrinos and gravitational waves, on the other hand, are emitted deep inside the supernova core and travel to observers on Earth practically unscathed by intervening material. They can act as messengers to provide first-hand and live dynamical information on the intricate multi-dimensional dynamics of the proto-neutron star and postshock region and may constrain directly the core-collapse supernova mechanism. Importantly, core-collapse events that do not produce the canonical observational astronomical signature or whose observational display is shrouded from view can still be observed in neutrinos and gravitational waves if occurring sufficiently close to Earth.  Gravitational waves, in contrast to neutrinos, have not yet been observed directly, but an international array of gravitational wave observatories (see, e.g., \\cite{aufmuth_05}) is active and taking data. Since gravitational waves from astrophysical sources are expected to be weak, their detection is notoriously difficult and involves extensive signal processing and detailed analysis of the detector output. Chances for the detection of an astrophysical event of gravitational wave emission are significantly enhanced if accurate theoretical knowledge of the expected gravitational wave signature from such an event is at hand.  Theoretical predictions of the gravitational wave signature from a core-collapse supernova are complicated, since the emission mechanisms are very diverse. While the prospective gravitational wave burst signal from the collapse, bounce, and the very early postbounce phase is present only when the core rotates~\\cite{mueller_82_a, moenchmeyer_91_a, zwerger_97_a, dimmelmeier_02_a, kotake_03_a, ott_04_a, ott_07_a, dimmelmeier_07_a}, gravitational wave signals with sizeable amplitudes can also be expected from convective motions at postbounce times, instabilities of the standing accretion shock, anisotropic neutrino emission, excitation of various oscillations in the proto-neutron star, or nonaxisymmetric rotational instabilities~\\cite{rampp_98_a, mueller_04_a, shibata_05_a, ott_05_a, ott_06_a, ott_07_a}.  In the observational search for gravitational waves from merging black hole or neutron star binaries, powerful data analysis algorithms such as matched filtering are applicable, as the waveform from the inspiral phase can be modeled with high accuracy (see, e.g., \\cite{blanchet_06_a}) and gravitational wave data analysts already have access to robust template waveforms that depend only on a limited number of macroscopic parameters. In contrast, the complete gravitational wave signature of a core-collapse supernova cannot be predicted with template-level accuracy as the postbounce dynamics involve chaotic processes (turbulence, [magneto-] hydrodynamic instabilities) that are sensitive not only to a multitude of precollapse parameters, but also to small-scale perturbations of any of the hydrodynamic variables.  While the complete supernova gravitational wave signature may remain inaccessible to template-based data analysis, a number of individual constituent emission processes, in particular those involving coherent global bulk dynamics and/or rotation, allow, in principle, for accurate and robust waveform predictions that may be applied to template-based searches. Rotating core collapse and core bounce as well as pulsations or nonaxisymmetric rotational deformations of a proto-neutron star constitute this group of processes. Among them, rotating collapse and bounce is the historically most extensively studied case (see, e.g., \\cite{ott_06_b} for a historical review) and may be the most promising for becoming robustly predictable in its gravitational wave emission. Yet, to date, the gravitational wave signal from rotating stellar core collapse and bounce has not been predicted with the desired accuracy and robustness.  These deficiencies of previous simulations result from the fact that the physically realistic modeling of core collapse requires a general relativistic description of consistently coupled gravity and hydrodynamics in conjunction with a microphysical treatment of the sub- and supernuclear EoS, electron capture on heavy nuclei and free protons, and neutrino radiation transport. Only very few multi-dimensional general relativistic codes have recently begun to approach these requirements~\\cite{ott_07_a, dimmelmeier_07_a}. In addition, the properties of the EoS around and above nuclear density are not very well constrained by theory or experiments. The same applies to the rotation rate and angular velocity profile of the progenitor core, which are also not directly accessible by observation and very difficult to model numerically in stellar evolution codes. Furthermore, variations with progenitor structure and mass are to be expected. Therefore, the influence of rotation and progenitor structure on the collapse and bounce dynamics and thus the gravitational wave burst signal must be investigated by extensive and computationally expensive parameter studies.  Previous parameter studies have considered a large variety of rotation rates and progenitor core configurations, but generally ignored important microphysical aspects and/or the influence of general relativity. M\\\"onchmeyer et al.~\\cite{moenchmeyer_91_a} performed axisymmetric Newtonian calculations with progenitor models from stellar evolutionary studies. They employed the microphysical nuclear EoS of Hillebrandt and Wolff~\\cite{hillebrandt_85_a} and included deleptonization via a neutrino leakage scheme and electron capture on free protons. Capture on heavy nuclei was neglected, which resulted in a too high electron fraction $ Y_e $ at core bounce and a consequently overestimated inner core mass~\\cite{yahil_83_a, van_riper_81_a}. In that study a limited set of four calculations was computed and two qualitatively and quantitatively different types of gravitational wave burst signals were identified. Their morphology can be classified alongside with the collapse and bounce dynamics: \\emph{Type~I} signals are emitted when the collapse of the quasi-homologously contracting inner core is not strongly influenced by rotation, but stopped by a \\emph{pressure-dominated bounce} due to the stiffening of the EoS near nuclear density $ \\rho_\\mathrm{nuc} \\approx 2 \\times 10^{14} \\mathrm{\\ g\\ cm}^{-3} $, where the adiabatic index $ \\gamma_\\mathrm{eos} $ rises above $ 4 / 3 $. This leads to a bounce with a maximum core density $ \\rho_\\mathrm{max} \\ge \\rho_\\mathrm{nuc} $. \\emph{Type~II} signals occur when centrifugal forces, which grow during contraction owing to angular momentum conservation, are sufficiently strong to halt the collapse, resulting in consecutive (typically multiple) \\emph{centrifugal bounces} with intermediate coherent re-expansion of the inner core, seen as density drops by sometimes more than an order of magnitude; thus here $ \\rho_\\mathrm{max} < \\rho_\\mathrm{nuc} $ after bounce. Type~I and~II dynamics and waveforms were also found in the more recent Newtonian studies by Kotake et al.~\\cite{kotake_03_a}, who employed a more complete leakage/capture scheme, but still obtained too high $ Y_e $ at bounce, and by Ott et al.~\\cite{ott_04_a}, who performed an extensive parameter study and for the first time also considered variations in progenitor star structure, but neglected deleptonization during collapse.  Zwerger and M\\\"uller~\\cite{zwerger_97_a} carried out an extensive two-dimensional Newtonian study of rotating collapse of idealized polytropes in rotational equilibrium~\\cite{komatsu_89_a} with a simplified \\emph{hybrid} EoS, consisting of a polytropic and a thermal component~\\cite{janka_93_a}. Electron capture during collapse was mimicked by an instantaneous lowering of the adiabatic index $ \\gamma_\\mathrm{eos} $ from its initial value of $ 4 / 3 $ to trigger the onset of collapse. At $ \\rho_\\mathrm{nuc} $, the adiabatic index was raised to $ \\gtrsim 2 $ to qualitatively model the stiffening of the nuclear EoS. Zwerger and M\\\"uller also obtained the previously suggested signal types and introduced \\emph{type~III} signals that appear in a pressure-dominated bounce when the inner core has a very small mass due to very efficient electron capture (approximated in~\\cite{zwerger_97_a} via a $ \\gamma_\\mathrm{eos} \\lesssim 1.29 $ in their hybrid EoS). Obergaulinger et al.~\\cite{obergaulinger_06_a} also employed the hybrid EoS, but included magnetic fields. They introduced the additional dynamics/signal \\emph{type~IV}, which occurs only in the case of very strong precollapse core magnetization. They found that weak to moderate core magnetization in agreement with predictions from stellar evolution theory (see, e.g., \\cite{heger_05_a}) has little effect on the collapse and bounce dynamics and the resulting gravitational wave signal. This finding is in agreement with~\\cite{kotake_04_a} (see also~\\cite{burrows_07_b, cerda_07_a}), where magneto-rotational collapse simulations were performed, a smaller model set was considered, but the neutrino leakage scheme of~\\cite{kotake_03_a} was employed and it was made use of two different microphysical EoSs to study the EoS dependence of the collapse dynamics and gravitational wave signal.  The first extensive set of general relativistic simulations of rotating iron core collapse to a proto-neutron star were presented by Dimmelmeier et al.~\\cite{dimmelmeier_02_a}, who employed an analytic hybrid EoS and polytropic precollapse models in rotational equilibrium as initial data (but see also the pioneering early work of~\\cite{nakamura_81_a}). These simulations were subsequently confirmed in~\\cite{shibata_04_a, cerda_05_a, ott_06_b, obergaulinger_06_b}. Dimmelmeier et al.\\ studied a subset of the models in~\\cite{zwerger_97_a} in the same parameter space of rotation rate and degree of differential rotation, and found that general relativistic effects counteract centrifugal support and shift the occurrence of type~II dynamics and wave signals to a higher precollapse rotation rate at a fixed degree of differential rotation.  Recently, new general relativistic simulations of rotating core collapse in two and three dimensions were carried out by Ott et al.~\\cite{ott_06_b, ott_07_a, ott_07_b} who included the microphysical EoS of Shen et al.~\\cite{shen_98_a}, precollapse models from stellar evolutionary calculations as well as an approximate deleptonization scheme~\\cite{liebendoerfer_05_a}. The results of these calculations indicate that the gravitational wave burst signal associated with rotating core collapse is \\emph{exclusively of type~I}. In addition, the simulations showed that rotating stellar iron cores stay axisymmetric throughout collapse and bounce, and only at postbounce times develop nonaxisymmetric features.  In a general relativistic two-dimensional follow-up study, Dimmelmeier et al.~\\cite{dimmelmeier_07_a, dimmelmeier_07_b} considerably extended the number of models and comprehensively explored a wide parameter space of precollapse rotational configurations. Even for this more general setup they found gravitational wave signals solely of type~I form, although for rapid precollapse rotation some of their models experience a core bounce due to centrifugal forces only, which however is always a \\emph{single centrifugal bounce} rather than the multiple ones observed in earlier work (see, e.g., \\cite{zwerger_97_a, dimmelmeier_02_a, ott_04_a}). They identified the physical conditions that lead to the emergence of this generic gravitational wave signal type and quantified their relative influence. These results strongly suggest that the waveform of the gravitational wave burst signal from the collapse of rotating iron cores in a core-collapse event is much more generic than previously anticipated.  In this work, we extend the above study of the gravitational wave signal from rotating core collapse and consider not only variations in the precollapse rotational configuration, but also in progenitor structure and nuclear EoS. In this way, we carry out the to-date largest and most complete parameter study of rotating stellar core collapse that includes all the (known) necessary physics to produce reliable predictions of the gravitational wave signal associated with rotating collapse and bounce. All our computed gravitational wave signals are made available to the detector data analysis community in a freely accessible waveform catalog~\\cite{wave_catalog}.  We perform a large number of two-dimensional simulations with our general relativistic core-collapse code \\coconut\\ and employ $ 11.2 $, $ 15.0 $, $ 20.0 $, and $ 40.0 \\, M_\\odot $ (masses at zero-age main sequence) precollapse stellar models from the stellar evolutionary studies of Heger et al.~\\cite{heger_00_a, heger_05_a}. In addition to the EoS by Shen et al.~\\cite{shen_98_a} used in our previous studies, we also calculate models with the EoS by Lattimer and Swesty~\\cite{lattimer_91_a}. We describe in detail and explain comprehensively the qualitative and quantitative aspects of the collapse and bounce dynamics and the resultant gravitational wave signal. We lay out the individual effects of general relativity, deleptonization, precollapse stellar structure and rotational configuration, and nuclear EoS on the gravitational wave signature from rotating core collapse. We study the prospects for nonaxisymmetric rotational instabilities in our postbounce cores, which could lead to an enhancement of the gravitational wave signature. Furthermore, we set our model gravitational radiation waveforms in context with present and future detector technology and assess their detectability.  This paper is organized as follows: In Section~\\ref{section:physical_model} we introduce our treatment of the general relativistic spacetime curvature and hydrodynamics equations. Furthermore, we introduce our variants of the two microphysical EoS we employ, the scheme for deleptonization and neutrino pressure contributions, our precollapse model set, and the gravitational wave extraction technique employed. Section~\\ref{section:numerical_methods} discusses the numerical methods used in the \\coconut\\ code and the computational grid setup for the simulations presented in this paper. In Section~\\ref{section:collapse_dynamics} we present the collapse dynamics and waveform morphology of our simulated models, while in Section~\\ref{section:density_structure_and_waveform} we investigate the stratification of the postbounce core and its impact on the gravitational wave signal. The detection prospects for the gravitational wave burst from core bounce are discussed in Section~\\ref{section:detection_prospects}, while the rotational configuration of the proto-neutron star and its susceptibility to nonaxisymmetric rotational instabilities are examined in Section~\\ref{section:rotation_rate}. Finally, in Section~\\ref{section:summary}, we summarize and discuss our results.  Throughout the paper we use a spacelike signature $ (-, +, +, +) $ and units in which $ c = G = 1 $. Greek indices run from 0 to 3, Latin indices from 1 to 3, and we adopt the standard Einstein summation convention. ",
        "conclusions": "\\label{section:summary}  In this article we have presented results from a comprehensive set of collapse simulations of rotating stellar iron cores to proto-neutron stars, using the axisymmetric general relativistic hydrodynamics code \\coconut. Our simulations treat all the relevant physics of the collapse phase to good approximation. They include precollapse iron core profiles from stellar evolutionary calculations, a highly efficient approximate treatment of deleptonization, a microphysical finite-temperature EoS, as well as neutrino pressure contributions. Magnetic fields are not included, since their relevance in the collapse and early postbounce phases is very likely negligible in cores with realistic precollapse fields~\\cite{heger_05_a, burrows_07_b, cerda_07_a, obergaulinger_06_b}.  The focus of our study is on procuring accurate and reliable waveforms of the gravitational wave burst signal associated with core bounce and on understanding the dependence of the signal characteristics on progenitor star mass, precollapse rotational setup, and nuclear EoS. To this end we have performed the to-date most extensive parameter study of this scenario, covering with more than 100 model calculations the parameter space spanned by (1) progenitor mass and model profile (zero-age main sequence masses from $ 11.2 $ to $ 40\\,M_\\odot $, presupernova models with and without rotation), (2) rotational configuration (slow and uniform to rapid and differential rotation), and (3) nuclear EoS prescription (from relatively soft to relatively stiff). Importantly, the parameter space encompasses and even goes beyond all precollapse rotational configurations that are deemed realistic in the context of collapsing massive stars.  A central result of this work is the finding that the gravitational wave burst from core bounce exhibits a generic waveform shape known as type~I in the literature~\\cite{moenchmeyer_91_a, zwerger_97_a}, independent of the model parameters. The multiple centrifugal bounce dynamics and the corresponding type~II waveform found in previous, technically less complete studies (see, e.g., \\cite{moenchmeyer_91_a, zwerger_97_a, dimmelmeier_02_a, ott_04_a}) do not occur in our models.  We have demonstrated that all models with precollapse core angular velocities $ \\Omega_\\mathrm{c,i} $ below $ \\sim 5 \\mathrm{\\ rad\\ s}^{-1} $ (corresponding periods longer than about $ 1 \\mathrm{\\ s} $) reach nuclear densities and experience a core bounce predominantly due to nuclear pressure effects. More rapidly rotating cores develop sufficient rotational support during collapse to undergo either a mixture of centrifugal and pressure-dominated bounce or a \\emph{single} centrifugal bounce at subnuclear densities. Centrifugal hang-up much below nuclear density or multiple, damped harmonic oscillator-like centrifugal bounces do not occur. Therefore, these models also exhibit a type~I waveform. The detailed analysis of the collapse dynamics presented in this paper reveals that the combined effects of general relativity and deleptonization lead to an increased destabilization of the collapsing core, result in a relatively small radius and mass $ M_\\mathrm{ic,b} $ of the sonically-connected inner core at bounce (but not small enough to show the type~III waveform associated with rapid collapse found in some previous simplistic models), and diminish the dynamical importance of centrifugal forces during collapse.  The key parameter which determines the peak amplitude $ |h|_\\mathrm{max} $ of the gravitational wave burst has turned out to be the precollapse central angular velocity $ \\Omega_\\mathrm{c,i} $. Slowly rotating cores with $ \\Omega_\\mathrm{c,i} \\lesssim 1 \\mathrm{\\ rad\\ s}^{-1} $ produce feeble peak amplitudes on the order of $ 10^{-22} $ at a distance of $ 10 \\mathrm{\\ kpc} $. More rapidly rotating cores with $ 1 \\mathrm{\\ rad\\ s}^{-1} \\lesssim \\Omega_\\mathrm{c,i} \\lesssim 6 \\mathrm{\\ rad\\ s}^{-1} $ develop stronger quadrupole deformations and have a rotationally-increased mass $ M_\\mathrm{ic,b} $ at bounce, resulting in sizeable peak amplitudes in the range of $ 5 \\times 10^{-22} \\lesssim |h|_\\mathrm{max} \\lesssim 10^{-20} $. The peaks of the waveform spectrum of such cores cluster in frequency space in the interval of $ 650\\mbox{\\,--\\,}800 \\mathrm{\\ Hz} $. At larger $ \\Omega_\\mathrm{c,i} $, centrifugal effects become strong, significantly decelerate collapse and bounce, and even lead to a purely centrifugal bounce in a subset of models. This results in a general decrease of $|h|_\\mathrm{max}$ and a shift of the waveform's spectral peak to frequencies below $ \\sim 400 \\mathrm{\\ Hz} $ at high $ \\Omega_\\mathrm{c,i} $.  We have also shown that, in addition to $ \\Omega_\\mathrm{c,i} $, the precollapse core mass in combination with the electron fraction sets the mass $ M_\\mathrm{ic,b} $ of the inner core at bounce, is an important quantity influencing the strength of the gravitational burst. Since more massive progenitors generally (though with notable non-monotonicity in the mass range from about $ 18 $ to $ 23\\,M_\\odot $) form larger iron cores, we observe in our model series a general trend to bigger $ M_\\mathrm{ic,b} $ and larger $ |h|_\\mathrm{max} $ with increasing progenitor mass if all other parameters are kept constant. For instance, the $ 40\\,M_\\odot $ progenitor yields values of $ |h|_\\mathrm{max} $ which are up to 4 times larger than for the lower-entropy $ 11.2\\,M_\\odot $ counterpart with the same rotational configuration.  The variations in the degree of differential rotation considered in this study have only a minor impact on the collapse dynamics and burst waveform amplitude. Increasing differential rotation at fixed $ \\Omega_\\mathrm{c,i} $ generally lowers the centrifugal support of outer core regions. However, since the dynamically most relevant inner core at bounce consists of only $ \\sim 0.5\\mbox{\\,--\\,}1\\,M_\\odot $ located within about $ 1000 \\mathrm{\\ km} $ at the onset of collapse, the effects of differential rotation on the gravitational wave burst are small.  Our results further indicate that the nuclear EoS has little influence on the gravitational wave burst signal. For a given precollapse configuration, a softer nuclear EoS yields higher densities at bounce and postbounce times with shorter variation timescales of the quadrupole moment, but also leads to greater inner core compactness. In our simulations, the two effects generally cancel, leading to no systematic trend in the peak waveform amplitude $ |h|_\\mathrm{max} $ with the EoS. The peak of the waveform spectrum, however, shifts to higher frequencies in the case of a softer EoS. For the models considered here, this frequency shift amounts to $ \\sim 5.5\\% $ on average for models undergoing pressure-dominated bounce. It is significantly smaller for models bouncing at subnuclear densities under the influence of centrifugal effects.  If situated within our Galaxy, a large fraction of our models are comfortably detectable by current gravitational wave detectors with a signal-to-noise ratio of up to 6 in the most optimistic case (which is obtained for the most rapidly rotating models that still undergo pressure-dominated core bounce). Advanced detectors could observe them easily out to $ \\sim 100 \\mathrm{\\ kpc} $ and up to several $ 10 \\mathrm{\\ Mpc} $ for third-generation detectors.  While such a gravitational wave signal may per se be detectable, the extraction of detailed physical information from the signal (i.e., solving the ``inversion problem'') from the signal will be a formidable task. The very generic morphology of the burst waveforms and the clustering in frequency space of most models make it seem unlikely that a pure waveform-template-based inversion (as, e.g., carried out in~\\cite{summerscales_08_a} using the waveforms of~\\cite{ott_04_a}) can be successful for determining key physical parameters to significant precision. Our results, however, suggest that based on $ |h|_\\mathrm{max} $ and the peak frequency $ f_\\mathrm{max} $ of the waveform spectrum alone, it should be possible to discriminate between purely pressure-dominated bounce (small to large $ |h|_\\mathrm{max} $ at frequencies $ f_\\mathrm{max} $ significantly above $ 500 \\mathrm{\\ Hz} $) and centrifugal bounce (large $ |h|_\\mathrm{max} $ at frequencies $ f_\\mathrm{max} $ significantly below $ 500 \\mathrm{\\ Hz} $). Furthermore, we find that for not too rapid rotation $ |h|_\\mathrm{max} $ can be directly used to extract the rotation rate $ \\beta_\\mathrm{b} $ at bounce to good precision.  Making use of the extensive set of postbounce rotational configurations obtained with our simulations, we have also studied the prospects for the development of nonaxisymmetric rotational instabilities in nascent proto-neutron stars. We find that the rotational barrier imposed by centrifugal forces prohibits the spin-up to rotation rates necessary for the classical dynamical bar-mode instability at high values of $ \\beta $. We find, however, that a large subset of our postbounce models exhibits sufficiently differential and rapid rotation to become subject to the recently discovered low-$ \\beta $ instability. Still, three-dimensional simulations as in~\\cite{rampp_98_a, ott_05_a, ott_07_a, ott_07_b} will be necessary to provide conclusive tests of our predictions. Furthermore, the interaction and competition of the low-$ \\beta $ instability and other instabilities operating on the shear energy of differential rotation, for instance the magneto-rotational instability (see, e.g., \\cite{balbus_91_a, cerda_07_a}), remain to be studied.  Finally, we point out that this study may be regarded as part -- with the presently highest level of sophistication -- of a multi-decade effort of our groups~\\cite{mueller_82_a, moenchmeyer_91_a, zwerger_97_a, dimmelmeier_02_a, ott_04_a, dimmelmeier_07_a, ott_07_a} to provide reliable estimates for the gravitational wave burst emission associated with rotating core collapse and core bounce. The waveforms presented here are for the first time not only accurate (i.e., numerically converged), but reliable and robust, since our calculations take into account all the necessary physics, including general relativity, deleptonization and a microphysical EoS. All waveforms are available for download in various formats in a publicly accessible waveform catalog~\\cite{wave_catalog}.  We point out that the gravitational wave emission process considered in this work operates at measurable strength only if the progenitor core is rotating a lot more rapidly than expected for ordinary iron cores (see, e.g., \\cite{ott_06_c, heger_05_a}). In slowly rotating core-collapse supernovae, turbulent convective overturn, instabilities of the accretion shock, and, possibly, proto-neutron star pulsations are likely to be the dominant emission processes of gravitational waves. The characteristics of these emission processes are not as well understood and will require more extensive and precise modeling to provide accurate estimates of the complete gravitational wave signature of core-collapse supernovae."
    },
    "0806/0806.3268_arXiv.txt": {
        "abstract": "Hot, underdense bubbles powered by active galactic nuclei (AGN) are likely to play a key role in halting catastrophic cooling in the centers of cool-core galaxy clusters.  We present three-dimensional simulations that capture the evolution of such bubbles, using an adaptive-mesh hydrodynamic code, FLASH3, to which we have added a subgrid model of turbulence and mixing.  While pure-hydro simulations indicate that AGN bubbles are disrupted into resolution-dependent pockets of  underdense gas, proper modeling of subgrid turbulence indicates that this a poor approximation to a turbulent cascade that continues far beyond the resolution limit.  Instead, Rayleigh-Taylor instabilities act to effectively mix the heated region with its surroundings, while at the same time preserving it as a coherent structure, consistent with observations.  Thus bubbles are transformed into hot clouds  of mixed material as they move outwards in the hydrostatic intracluster medium (ICM), much as large airbursts lead to a distinctive  ``mushroom cloud'' structure as they rise in the hydrostatic atmosphere of Earth.  Properly capturing the evolution of such clouds has important implications for many ICM properties.  In particular, it significantly  changes the impact of AGN on the distribution of entropy and metals in cool-core clusters such as Perseus. ",
        "introduction": "The X-ray and abundance  profiles  of the hot, diffuse medium in galaxy clusters are observed to  be bimodal (Fukazawa \\etal 2000; Matsushita \\etal 2002;  Schmidt \\etal 2002; Churazov \\etal 2003; De Grandi \\etal 2004). Strong intracluster medium (ICM) abundance gradients are associated with cool-core clusters with a peak in their central X-ray surface brightness distributions, and nearly uniform abundance profiles are associated with non cool-core clusters.  Furthermore, these metallicity distributions are much broader than the associated galaxy light, indicating that significant mixing has occurred (\\eg Churazov \\etal  2003; Chandran 2005; David \\& Nulsen 2008).  At the same time, the nature of cool-core clusters remains uncertain. Although strong X-ray emission indicates that the central gas is cooling rapidly, the deficit of star formation and  $<$ 1 keV gas (\\eg Fabian 1994; Tamura \\etal 2001) means that radiative cooling must be balanced by an unknown energy source.  Currently, the most successful model for achieving this balance relies on heating from a central AGN,  yet the details of this process are poorly understood (\\eg Loken \\etal 1995; Br\\\" uggen \\& Kaiser 2002; Reynolds \\etal 2002; Brighenti \\& Mathews 2006; Brunetti \\& Lazarian 2007).  While AGN are observed to drive large bubbles filled with relativistic particles (\\eg Boehringer \\etal 1993; Carilli \\etal 1994; McNamara \\etal 2000; Blanton \\etal 2001; Finoguenov \\& Jones 2001; Nulsen \\etal 2005),  the  synchrotron radiation emitted by the electrons in these regions fades after about $10^8$ years, becoming extremely difficult to detect. Moreover, the  corresponding depressions in the X-ray surface brightness are only visible near the center of the cluster, where the contrast is large. Thus, it is unclear how far these structures rise in the cluster.  Furthermore, AGN have been observed to induce shocks and/or sonic motions in the ICM that are believed to eventually dissipate their energy into this gas (Fabian \\etal 2003; Kraft \\etal 2003; Ruszkowski \\etal 2004; McNamara \\etal 2005), although the impact of the resulting heating is difficult to quantify observationally.  The presence of AGN-heated cavities in galaxy clusters has raised a number of questions.  These buoyant bubbles are unstable  to the Rayleigh-Taylor (RT) instability,  which occurs whenever a fluid  is accelerated or supported against gravity by a fluid of lower density.  In any ideal hydrodynamic model, the cavities must be inflated supersonically or else they would be destroyed by RT instabilities faster than they are inflated (Reynolds \\etal 2005). Curiously, the strong and hot ICM shocks that are expected around the active cavities are absent, and, instead, many cavities are surrounded by shells of gas that is cooler than the ambient ICM (Nulsen \\etal 2002). Secondly, these cavities appear to be intact even after inferred ages of several $10^8$ yrs, as for example the outer cavities in Perseus (Nulsen \\etal 2002). However, hydrodynamic simulations fail to reproduce the observed morphology as the RT and other instabilities shred the bubbles in a relatively short time (Robinson \\etal 2004; Reynolds \\etal 2005, Br\\\"uggen \\etal 2005a), although this time can  be extended somewhat by a more detailed treatment of bubble inflation (Pizzolato \\& Soker 2006; Sternberg \\etal 2008; Sternberg \\& Soker 2008a,b).  A consequence of the evolution of such cavities is the production of turbulence, which is likely to be pervasive in the ICM and play a crucial role in the structure of cool-core clusters (\\eg Schuecker \\etal 2004).  Turbulence occurs in any case in which the Reynolds number, $Re$, is greater than $\\approx 1000$, where  $Re \\equiv d \\, v /\\nu,$  $d$ is the characteristic scale of the flow instability, $v$ is its characteristic velocity, and $\\nu$ is the kinetic viscosity  of the fluid.   While there have been some suggestions that the ICM may have a non-negligible viscosity (Ruszkowski \\etal 2004; Reynolds \\etal 2005), this quantity remains unknown  because the physics of such dilute and magnetized plasmas is poorly constrained. In particular,  even minute magnetic fields lead to small proton gyroradii that suppress viscosity efficiently. However, it has been pointed out that the exponential divergence of neighboring field lines in a tangled magnetic field may lead to only modest suppression (\\eg Narayan \\& Medvedev 2001). On the other hand, Rebusco \\etal (2005) showed that the turbulent  scales and velocities required to spread metals in the Perseus cluster closely correspond to those necessary to balance cooling, and similar results have been recently obtained for several other clusters (Rebusco \\etal 2006).  This turbulence has important implications for  particle acceleration and the mixture and transport of gas energy and  heavy elements.   Observationally, there are several circumstantial clues about the nature of turbulence in the ICM, such as the pressure distribution in the Coma cluster (Schuecker \\etal 2004), the lack of resonant scattering in the 6.7 iron K$\\alpha$ line in the Perseus cluster (Churazov \\etal 2004) and the Faraday rotation map of the Hydra cluster (Vogt \\& Ensslin 2005, see also Iapichino \\& Niemeyer 2008). Turbulence has also been invoked to explain the non-thermal emission in clusters (\\eg Brunetti \\& Lazarian 2007), and Doppler shifts from such bulk motions are likely to be directly detectable with the next generation of X-ray observatories, such as {\\em Constellation-X}, with an envisaged spectral resolution of 1-2 eV (\\eg Br\\\"uggen \\etal 2005b).  Finally, our understanding of AGN-driven turbulence in cool-core clusters is complicated both by a variety of other possible sources of turbulence and competing physical effects.  Clusters form through the accretion of smaller structures  and this infall can generate turbulence (see \\eg Takizawa 2005). Episodes of active merging are also expected to produce turbulence, as seen in simulations, (\\eg  Norman \\& Bryan 1999; Ricker \\& Sarazin 2001). Within clusters, the motion of galaxies can also produce wakes that are likely to generate turbulent motions (Bregman \\& David 1989; Kim 2007), and  microphysical processes, often described in terms of conduction and viscosity, may also play important roles (\\eg Narayan \\& Medvedev 2001; Voigt \\& Fabian 2004; Ruszkowski \\etal 2004; Sternberg \\etal 2007).  In summary, cool-core clusters are a mystery that is carefully observed, poorly understood, and closely tied up with AGN-driven turbulence.  It is with this in mind that we have carried out detailed simulations of AGN-driven turbulence in a cool-core cluster using the adaptive mesh refinement code, FLASH3.  While the direct simulation of turbulence is extremely challenging, computationally expensive, and dependent on resolution  (\\eg  Glimm \\etal 2001), its behavior can be approximated to a good degree of accuracy by adopting a sub-grid approach. In this case the flow is decomposed into mean and fluctuating components, which provides a systematic framework for deriving a set of turbulence equations.  There are several types of models that are able to describe such fluctuations arising from  the RT instability as well as the Richtmyer-Meshkov instability, which occurs when a shock hits a medium of varying acoustic impedance. The simplest such model consists of ordinary differential equations for the mixing region (\\eg Alon 1995; Chen \\etal 1996; Ramshaw 1998), describing for example, the amplitude of the bubble by balancing inertia, buoyancy, and drag forces.  While these yield the right growth rates, they fail when there are multiple interfaces  and are not readily extended to two and more dimensions. Although these problems can be addressed with multifluid models (Youngs 1989), such models are complicated, numerically expensive, and sometimes unstable.  A second class of models evolves the turbulent kinetic energy per unit mass and its dissipation rate. Such ``two-equation turbulence models'' developed for unstable shear flows postulate a turbulent viscosity, a Reynolds stress, and a dissipation term (\\eg Llor 2003). However, the usual Reynolds stress terms must be modified in the presence of shocks, and modeling the RT and RM instabilities requires a buoyancy term that depends on the amplitude and the wavelength.  Recently, DiMonte \\& Tipton (2006, hereafter DT06), described a sub-grid model that is especially suited to capturing the buoyancy-driven turbulent evolution of AGN bubbles. The model captures the self-similar growth of the RT and RM instabilities by augmenting the mean hydrodynamics equations with evolution equations for the turbulent kinetic energy per unit mass and the scale length of the dominant eddies. The equations are based on buoyancy-drag models for RT and RM flows, but constructed with local parameters so that they can be applied to multidimensional flows with multiple materials.  The model is self-similar, conserves energy, preserves Galilean invariance, and works in the presence of shocks, and although it contains several unknown coefficients, these are determined by comparisons with analytic solutions, numerical simulations, and experiments.  Here we implement the DT06 model into FLASH3 to: 1.) examine the impact of turbulence on the morphology and stability of AGN-driven bubbles as observed in nearby clusters (\\eg Fabian 2006); 2.) quantify turbulence, comparing it to present indirect entropy and metal distribution constraints (\\eg Inogamov \\& Sunyaev 2003) and making predictions for radial profiles,  as directly measurable  by future linewidth studies.  Our goal is to focus on better understanding the basic case of purely AGN-driven turbulence in an inviscid, unmagnetized ICMm and to this end, we follow the model of Roediger \\etal (2007; hereafter R07), in which the ICM is described by a spherically-symmetric profile fit to X-ray observations of the Perseus cluster, and feedback is modeled by periodically injecting energy into the center of this distribution.  The structure of this work is as follows.  In \\S2 we give an overview of the FLASH3 code and our implementation the DT06 subgrid-turbulence model within it.  In \\S3 we present tests of our implementation against analytic solutions.  In \\S4 we discuss our modeling of the galaxy cluster and energy input by the central AGN. In \\S5 we present the results from our simulations and discuss their observational consequences.  Our conclusions are summarized in \\S 6.  \\vspace{2cm} ",
        "conclusions": "A wide range of observations suggest that AGN-feedback  plays a key role in the evolution of cool-core galaxy clusters.  At the same time, theoretical studies have pointed out some of the many physical mechanisms that may be important in this evolution, including viscosity, magnetohydrodynamic effects, heat conduction, and cosmic-ray heating.  While any or all of these may operate in nature, none of them can be fully understood without first accurately capturing the underlying hydrodynamic evolution of the ICM.  It is with this basic goal in mind that we have used a subgrid turbulence model to study AGN heating in  an inviscid fluid, neglecting other effects such as magnetic fields and heat-conduction.  Within this context, our study has been focused on two key issues: the growth of instabilities and turbulence caused by AGN-heated bubbles and the role of these instabilities in determining the evolution of the bubbles and the surrounding medium.  Clearly, there  are other sources of turbulence that might add  similar  levels of turbulent energy into the ICM, such as mergers of large subclusters or motions of galaxies within a cluster.  Likewise, the DT06 subgrid turbulence model that we have employed does not include all processes that generate turbulence, such as the shear-driven Kelvin-Helmholtz instability.  However, our tests show that it is effective in capturing the growth of the extremely important RT and RM instabilities.  Thus although many aspects of cluster evolution remain uncertain and beyond the scope of this work, there are a number of robust conclusions that we can make about the role of these two instabilities in shaping the evolution and impact of AGN-heated cavities in clusters.  In particular we find that:  \\begin{itemize}  \\item  Many of the RT and RM unstable modes that drive the evolution of the bubbles evolve on scales that are far below the resolution limits of current simulations. The superposition of these unstable modes  smears out the interface between the bubbles and the ambient medium, transforming them into clouds of mixed material that stay intact and expand as they rise in the  stratified cluster medium. This mixing can explain the coherent X-ray cavities detected in clusters of galaxies. The subgrid-turbulence model also greatly reduces the sensitivity of our results on the resolution of the computational grid and the detailed choice of initial conditions.   \\item  Within the clouds, turbulent motions quickly attain typical velocities of $\\approx 200$ km s$^{-1},$ roughly $10\\%$ of the internal sound speed  and $20\\%$ of the sound speed of the surrounding ICM.  Similarly, the scale of the turbulent eddies rises swiftly but does not exceed the $\\approx 30$ kpc scale of clouds.  A typical turbulent diffusivity is then $\\approx 500$ km/s kpc, which is in excellent agreement with the diffusion coefficient inferred by the abundance profiles studies of Perseus by Rebusco \\etal (2005).  \\item Subgrid turbulence is likely to enhance metal transport significantly.  In our fiducial single-bubble evacuated pure-hydro run, metal transport is halted at $\\approx 50$ kpc by bubble disruption, which occurs when RT instabilities have shredded the evacuated region into resolution-limited cavities.  Yet, in the subgrid-turbulence run, small-scale fluctuations act to keep the cloud coherent for much longer, thus causing it to distribute metals out to larger radii.  \\item In the case where the bubbles produce shock waves, the turbulent velocities and length scales are systematically enhanced with respect to the evacuated-bubble run, due to the RM instability.  While such $\\approx 300$ km/s velocities are roughly 50\\% greater than in the runs with evacuated bubbles, they are still well below the $\\approx 500$ km/s radial velocities of the clouds. However they are large enough to be probed by studying the emission lines of heavy ions with the future {\\em Constellation-X} satellite, which will have an envisaged spectral resolution of 1-2 eV.  \\item Calculating the turbulent kinetic energy produced by the rising bubbles, we find that this is only $\\approx 1\\%$ of  the total energy available for the bubbles to  heat the cluster.  Hence, we do not expect the energy of the turbulent motions themselves to play a major role in the heating of the cool-core regions of the ICM.  Rather, the main role of turbulence is to increase the efficiency with which the thermal energy of the rising clouds is mixed into the surrounding  gas.   \\item   Finally, runs that include radiative cooling and multiple episodes of AGN-feedback indicate that the  impact of turbulence continues to increase with successive generations of heating.   This is because turbulence driven by previous feedback events remains behind, enhancing the mixing of subsequent bubbles.  While in our pure-hydro runs, the bubbles deposit most of their internal energy  at the resolution-dependent radius at which they are disrupted, the turbulent motions captured by the subgrid model lead to more gradual heating, correspondingly shallower temperature and entropy profiles, and shallower metal gradients.  \\end{itemize}   In summary, the properties, evolution, and appearance of AGN-blown bubbles in clusters are substantially different when one properly accounts for turbulent motions on scales well below the limits of current pure-hydro simulations.  Although accounting for these motions may not provide the ultimate solution to many of the mysteries surrounding galaxy clusters, it is a crucial step forward in the modeling of feedback and the interaction between AGN and the ICM. Subgrid models such as the one developed in DT06 provide us with tools for capturing this physics.  In fact, the numerical methodology presented here is likely to have applications in other areas of astrophysics where hydrodynamic modeling of RT instabilities is crucial, such as supernovae, supernova remnants, and galactic winds. It is clear that while many physical processes may play important roles in clusters and other environments, these can only be understood when carefully disentangled from the impact of subgrid turbulence."
    },
    "0806/0806.1437_arXiv.txt": {
        "abstract": " ",
        "introduction": "Revealing the nature of dark energy is fundamentally important not only for astrophysics but also for particle physics. Constraints on the dark energy from astronomical observations is very influential for them. Baryon acoustic oscillations (BAO) in the galaxy power spectrum provide a strong constraint on the dark energy using its acoustic scale as a standard ruler. Large galaxy surveys such as the Sloan Digital Sky Survey and two degree field already provide the constraint and future larger surveys are currently planned to detect the BAO more accurately\\cite{ei05,co05,pe07,ok07}. Hence, an accurate theoretical model of the BAO is crucial, and many authors have been investigating the BAO using numerical simulation\\cite{se05,abfl07,h07,sss07,sss08,t08,sba08} and the perturbation theory (including the renormalized perturbation theory)\\cite{jk06,cs06,cs08,m07,mc07,th08,is07,mp07,n07}.  Previously, several authors investigated the third-order density perturbation and derived the one-loop correction to the linear power spectrum in the EdS model\\cite{j81,j83,ggrw86,ss91,mss92,jb94}. Similarly for the cosmological constant model, Bernardeau (1994) presented the third-order perturbation solution (see also Refs. \\citen{bjcp92,b94,bchj95,clmm95,m95}). They found that the dependence of the cosmological model on the second- and third-order perturbations is very small, if the scale factor in the EdS model is replaced with the linear growth factor.\\footnote{Martel \\& Freudling (1991) and Scoccimarro et al. (1998) showed that this assumption is valid if $f \\equiv d \\ln D_1/ d \\ln a = \\Omega_M^{1/2}$. However, since $f \\approx \\Omega_M^{0.6}$, the approximation is not accurate.} However, since the theoretical model of the BAO should archive the subpercent accuracy to provide a strong constraint on the dark energy, it is useful to reinvestigate this topic to accurately check the above assumption. In this study, we calculate the third-order density perturbation, newly including the dark energy with the time-varying equation of state, and derive the one-loop power spectrum analytically for the first time. We compare our results with the approximate results based on the EdS model in detail, and discuss the effect of the dark energy on the power spectrum near the baryon acoustic scale.   Throughout this paper, we use $\\delta$ as the density fluctuation, $\\theta ~(=\\nabla \\cdot {\\boldsymbol{v}})$ as the divergence of the peculiar velocity field, and $\\tau=a(t) dt$ as the conformal time. $\\Omega_M$, $\\Omega_K$ and $\\Omega_X$ are the density parameter for the matter, the curvature and the dark energy at present. $w(a)$ is the equation of state of the dark energy. The Hubble expansion rate is $H^2(a)= H_0^2 \\left[ \\Omega_M a^{-3} + \\Omega_K a^{-2} + \\Omega_X \\exp \\left[ 3 \\int_a^1 da^\\prime \\left( 1+w(a^\\prime) \\right)/a^\\prime \\right] \\right]$. ",
        "conclusions": "We investigate the third-order density perturbation and the one-loop power spectrum in the dark-energy cosmological model. We present analytical solutions and a fitting formula with the general time-varying equation of state for the first time. It turns out that the cosmological dependence is very weak, for example, less than $1 \\%$ for $k<0.4h/$Mpc for the power spectrum. However, our results may be useful in some cases when one needs a very highly accurate theoretical model of the BAO or in the study of the nonlinear evolution on a smaller scale ($>0.4h$/Mpc)."
    },
    "0806/0806.0412.txt": {
        "abstract": "{}{Neutral hydrogen clouds are found in  the Milky Way and Andromeda halo both as large complexes and smaller isolated clouds.  Here we present a  search for \\hi clouds in the halo of M33, the third spiral galaxy of the Local Group.}{We have used two complementary data sets:  a 3$^{\\circ} \\times 3^{\\circ}$ map of the area provided by the Arecibo Legacy Fast ALFA (ALFALFA) survey and deeper pointed observations carried out with the Arecibo telescope in two fields  that permit sampling of the north eastern and south-western edges of the \\hi disc. } {The total amount of \\hi around M33 detected by our survey is $\\sim 10^7$ M$_{\\odot}$. At least 50\\% of this  mass is made of  \\hi clouds that are related both in space and velocity to the galaxy. We discuss several scenarios for the origin of these clouds focusing on the two most interesting ones: $(a)$ dark-matter dominated gaseous satellites, $(b)$ debris from filaments flowing into M33 from the intergalactic medium or generated by a previous interaction with M31. Both scenarios seem to fit with the observed cloud properties. Some structures are found at anomalous velocities, particularly an extended \\hi complex previously detected by Thilker et al. (2002). Even though the ALFALFA observations seem to indicate that this cloud is possibly  connected to M33 by a faint gas bridge, we cannot firmly establish its extragalactic nature or its relation to M33.} {Taking into account that the clouds associated with M33 are likely to be highly ionised by the extragalactic UV radiation, we predict that the total gas mass associated with them is $\\ge 5\\times 10^7$~M$_\\odot$. If the gas is steadily falling towards the M33 disc it can provide the fuel needed to sustain a current star formation rate of  0.5 \\msun yr$^{-1}$. } ",
        "introduction": "The  origin of \\hi High Velocity Clouds (HVCs) is a long-standing problem since their discovery around the Milky Way about 40 years ago (Muller et al. 1963, Oort 1966). Different scenarios have been addressed to explain the nature of these clouds, and here we briefly summarise the ones that are of particular interest to the interpretation of the data presented herein.  \\begin{itemize}  \\item The \"galactic fountain\" model predicts that the clouds have a local origin and they form in the halo from hot gas ejected by supernova explosions in the disc. As the gas flows into the halo it radiatively cools down  and then  falls back onto the disc appearing as high-velocity moving  gas (Shapiro \\& Field 1976, de Avillez 2000).  \\item Other interpretations suggest that at least a fraction of the observed HVCs have an extragalactic origin. One possibility is that such clouds are the remnant of gas stripped by previous interactions with smaller neighbors. The most direct evidence of such a process is given by the Magellanic Stream between the Milky Way and the Magellanic clouds (Mathewson et al. 1974) which contains about 2 $\\times 10^8$ \\msun of \\hi (at a distance of 55 kpc) (Putman et al. 2003).  \\item Another hypothesis for the origin of the HVCs assumes that \\hi clouds may constitute the gaseous counterparts of the \"missing\" dark satellites predicted by galaxy formation theories (Blitz et al. 1999, Braun \\& Burton 2000). The structure formation scenario in the Lambda Cold Dark Matter ($\\Lambda$CDM) paradigm predicts a larger number of dark matter dominated satellites around massive galaxies, such as  the Milky Way or Andromeda, compared to what is currently observed (Klypin et al. 1999). Although  new very faint Local Group (LG) dwarfs have been discovered in the last few years with the Sloan Digital Sky Survey (SDSS) and other wide field optical surveys (see Simon \\& Geha 2007 and references therein), a discrepancy between theory and observations still exists, unless one assumes that such dark halos have not been able to form stars.  They should then contain mainly ionised and neutral hydrogen bound in their dark matter potential wells. Therefore HVCs (particularly the compact HVC class) have been considered in the last few years as possible tracers of the missing dark matter halos.   \\item Hot intergalactic medium condensing into galaxies is another possibility to accrete gas and fuel star formation. At low redshift, low mass  galaxies can get part of their gas through a cold accretion mode (Binney 1977, Katz \\& White 1993, Fardal et al. 2001, Murali et al. 2002, Keres et al. 2005). In this case the gas is not heated to the virial temperature of the halo but accretes at $T\\sim 10^4$ K and establishes radiative equilibrium with extragalactic ionising radiation. In Cold Dark Matter (CDM) models this gas resides in filamentary structures and its accretion rate depends on the environment, being higher for isolated galaxies in low density regions. The gas might enter the virial radius with high speed, of order 100-300~km~s$^{-1}$ and it might  be shocked close to the galaxy disc where it radiates most of its excess energy. Alternatively the gas might convert part of its infall velocity into rotational velocity (see Keres et al. 2005 for a more detailed discussion). In the case of M33 it is unclear whether any of the gas is accreted in this way,  being a galaxy of intermediate mass, and the detection of any warm neutral gas in the halo will help quantify this possibility. This scenario is not very different from the original Oort suggestion (Oort 1970), that HVCs around galaxies are related to residual gas left over from their formation and  gradually accreted by the host galaxy.   \\item Finally, in the case of galaxies orbiting a more massive one, such as M33 which is a satellite of M31, gas may be removed from the disc when the galaxy is close to  the pericenter, and then, as the distance increases, the gas might fall back onto the disc. The orientation of the M33 gas warp in the direction of M31 and the shift of the center of the orbits in its outermost parts (Corbelli \\& Schneider 1997) might indicate that the M31 disturbances are not negligible on the gaseous disc of its smaller companion, even though the smooth appearance of the M33 stellar distribution constrains  the history of their interaction (see Loeb et al. 2005). Numerical simulations (Jiang \\& Binney 1999) show that the warp itself might be generated by a slow gas-accretion process since gas infall reorients the outer parts of the halo.  \\end{itemize}  Searching for gas in the halo of %HVCs and \\hi complexes in nearby galaxies can help to better understand the properties of HVCs  and to discriminate between these different scenarios. A population of \\hi clouds possibly located in the vicinity of M31  out to a projected distance of 50 kpc has been recently discovered (Thilker et al. 2004, Westmeier et al. 2005). Some of the \\hi features  appear to have a tidal origin, being located  in proximity of the giant stellar stream of M31 and of its satellite NGC 205, while others are spatially isolated with radial velocities that differ from any known dwarf companion.  \\begin{figure*} \\begin{center} \\includegraphics[width=10.cm]{Fig01.ps} %\\newline \\end{center} \\label{map} \\caption{The distribution of all the \\hi clouds detected around the disc of M33 within the ALFALFA cube. The density contours range from N$_{HI} = 2 \\times 10^{18}$ \\cmsq, to $3 \\times 10^{19}$ \\cmsq. {\\em Type 1} clouds with $V > -180$ \\kms are shown in light grey, {\\em Type 1} clouds with $V < -180$ \\kms are in dark grey, while {\\em Type 2} clouds are plotted in black. The extent of the M33 \\hi disc is shown with two contours at a column density of 5 $\\times 10^{19}$ \\cmsq and $1 \\times 10^{21}$ \\cmsq.} \\end{figure*}  \\begin{figure*} %\\begin{center} \\includegraphics[width=8.cm]{Fig02a.ps} \\includegraphics[width=8.cm]{Fig02b.ps} \\caption{{\\em Left}: The distribution of {\\em Type 1} clouds which have radial velocities compatible with the rotation of the M33 disc  ($-80$ \\kms $<V< -320$ \\kms). {\\em Right}: The distribution of {\\em Type 2} clouds whose radial velocities are below -320 \\kms . The extent of the M33 \\hi disc is shown with two contours at a column density of 5 $\\times 10^{19}$ \\cmsq and $1 \\times 10^{21}$ \\cmsq.} %\\end{center} \\label{map_vel} \\end{figure*}    Contrary to the Milky Way and Andromeda, M33 is fairly isolated, and its recent evolution  seems to have been rather smooth since there is no evidence for  recent mergers; the disc does not appear to be disturbed, there is no prominent bulge, and there are no traces of substructure in the stellar population of its outer regions (Ferguson et al. 2006). At the same time, M33 is a blue galaxy with a much higher star formation rate per unit surface compared to the more massive companion M31.  Finding evidence of extraplanar gas, or gas in the very outer disc, can help constrain the evolutionary scenarios for this galaxy. Recent results of a chemical evolution model of M33 point out that, to sustain its star forming activity and to explain its shallow metallicity gradient, M33 needs to accrete gas at a rate of about 1 \\msun yr$^{-1}$ (Magrini et al. 2007). The 21-cm survey around M31 by Westmeier et al. (2005) was also extended to M33, and it led to the discovery of only one HVC. Moreover, the presence of an arch of carbon stars at the edge of the optical disc  (Block et al. 2004) supports an inside-out formation scenario and suggests that M33 is still in the process of accreting matter which sporadically enhances star formation.    Therefore to search for additional HVCs and to inspect the \\hi content of the halo of M33   we have used two complementary sets of data: 1) a $3^{\\circ} \\times 3^{\\circ}$ 21-cm map of the M33 region (centred on the optical position of the galaxy) extracted by the Arecibo Legacy Fast ALFA survey  (ALFALFA; Giovanelli et al. 2005) carried out with the Arecibo telescope; 2) a  map of two $3^{\\circ}.5 \\times 1^{\\circ}.5$ regions in the south-east and north-west part of the galaxy enabling a study of the edge of the \\hi disc. The latter data have also been taken with the Arecibo telescope but with an integration time of 6$^{m}$ per pointing ($\\sim$10 times longer than ALFALFA) and  higher spectral resolution (5 times ALFALFA).  The main thrust of this paper will be that of testing hypotheses that place the newly discovered clouds at the distance of M33. We describe their main properties and we discuss the possible origins of these complexes  and the implications  for the evolution of the gaseous disc.   The paper is organised as follows: in Section 2 we describe the data sets we used for this study, in Section 3 we give a brief overview of the results, in Sections 4, 5 and  6 we describe  the clouds detected in both data sets, in Section 7 we present the global properties of the clouds, in Section 8 we discuss some possible formation scenarios and we estimate the ionisation fractions and total gas masses of these clouds, and in Section 9 we present our final conclusions. ",
        "conclusions": "\\hi clouds seem to be a common feature in the halos of  spiral discs at least in the LG. The Milky Way and Andromeda have been known to contain a population of high velocity clouds (HVCs) and here we show that a similar gaseous component is detected in the environment of M33.  At the distance of 840 kpc, the \\hi masses of the clouds  range  between 10$^4$ and few times 10$^6$ solar masses.  They are found within a projected distance of $\\sim$ 20 kpc and they appear to be distributed (in projection)  along the major axis of the outer disc  which is warped towards the direction of the Andromeda galaxy. The total \\hi mass associated with {\\em Type 1} clouds is 5 $\\times 10^6$ \\msun and it goes up to $7 \\times 10^6$ \\msun if one includes also {\\em Type 2} objects. %, which is  about 10$\\%$ of the total \\hi mass of M33.    The origin of high velocity clouds is still under debate; in this paper we have discussed some possible scenarios for the formation of HVCs in M33 and how they fit with their observed properties.  We have  explored the possibility that these clouds might be associated with a population of dark-matter dominated satellites. From ST02 approximately 25 dark mini-halos would be expected within 50 kpc from M33 with total masses around $10^8$ \\msun. This is comparable to the total number of clouds we have detected (19 excluding those with anomalous velocity). From the comparison  to their mini-halo models we find that this scenario would be  consistent with the observed line widths and radii. Given the overall low \\hi column densities (below $10^{19}$ \\cmsq)  a large fraction of the gas would be ionised. The ionised component could be up to 50 times larger than the neutral one, with the ionised envelope extending for more than 3 kpc. We derive that the total gas mass associated with all {\\em Type 1} clouds would be of the order of $10^8$ \\msun.  Assuming that all the gas locked in the halo is falling back towards the disc at an average velocity of 100 \\kms from an average distance of 20 kpc, we derive a gas accretion rate of $\\lesssim$ 0.8 \\msun yr$^{-1}$.  We have also considered other scenarios in which  \\hi clouds are  condensations within a hot intergalactic medium (cosmic filaments) or they are either tidal debris of a previous interaction between M33 and M31 or remnants of gas stripped from a dwarf satellite at earlier times. Quantifying the total gas mass (\\hi and H{\\sc ii}) within the M33 halo in these cases is more difficult and we can only set a lower limit of $5 \\times 10^7$ \\msun.    Higher angular resolution observations of the M33 population of \\hi clouds will be useful to probe their inner structure, and to search for a higher column density core and internal velocity gradients.  At smaller radial velocities ($V < -320$ \\kms) we have resolved a \\hi complex that had been serendipitously discovered by Thilker et al. (2002) with Westerbork observations of M33.  The complex is  at about one degree from the centre of the galaxy ($\\sim$ 15 kpc), but it is  characterised by a large velocity relative to M33 ($\\sim 200$ \\kms) which complicates a firm  association to the galaxy. Two main substructures can be distinguished from the ALFALFA data and there appear to be faint clumps of gas pointing towards the center of M33, which may give an indication for a possible interaction with the galaxy (see  Figure 2).  We cannot exclude that this is a cloud not gravitationally bound to M33 and higher sensitivity observations of the region  are needed to test this hypothesis. On the other hand, if this is a local \\hi complex, a search for similar clouds in this part of the sky would help to disentangle the origin of this object."
    },
    "0806/0806.4738_arXiv.txt": {
        "abstract": "A cellular automaton model of pulsar glitches is described, based on the superfluid vortex unpinning paradigm.  Recent analyses of pulsar glitch data suggest that glitches result from scale-invariant avalanches \\citep{Melatos07a}, which are consistent with a self-organized critical system (SOCS).  A cellular automaton provides a computationally efficient means of modelling the collective behaviour of up to $10^{16}$ vortices in the pulsar interior, whilst ensuring that the dominant aspects of the microphysics are not lost.  The automaton generates avalanche distributions that are qualitatively consistent with a SOCS and with glitch data.  The probability density functions of glitch sizes and durations are power laws, and the probability density function of waiting times between successive glitches is Poissonian, consistent with statistically independent events.  The output of the model depends on the physical and computational paramters used.  The fitted power law exponents for the glitch sizes ($a$) and durations ($b$) decreases as the strength of the vortex pinning increases.  Similarly the exponents increase as the fraction of vortices that are pinned decreases.  For the physical and computational parameters considered, one finds $-4.3\\leq a \\leq -2.0$ and $-5.5\\leq b\\leq -2.2$, and mean glitching rates in the range $0.0023\\leq\\lambda\\leq0.13\\rm$ in units of inverse time.  \\newline\\newline ",
        "introduction": "\\label{sec:intro}  It has recently been shown that glitches in individual pulsars obey avalanche statistics \\citep{Melatos07a}.   Such statistical behaviour, in conjunction with the wide dynamical range of glitches, suggests that the underlying physics is that of a self-organized critical system (SOCS) \\citep{Pines85,Bak96,Jensen98,Melatos07a}.  SOCS are characterised by transitions between metastable states via scale invariant avalanches \\citep{Jensen98}.  The fractional increase in spin frequency spans seven decades across the entire glitch population ($10^{-11}\\leq\\delta\\nu / \\nu\\leq 10^{-4}$), and up to four decades in any individual pulsar.  A natural way to model a scale invariant, avalanching system is to construct a cellular automaton that is driven slowly to a threshold of instability \\citep{Field}.  Such models have been explored in detail in the context of earthquakes \\citep{Sornette}, granular assemblies \\citep{Bak88,Wiesenfeld,Bak96,Frette96,Pruessner03}, solar flares \\citep{Lu}, and superconducting vortices \\citep{Jensen,Field,Bassler,Linder}.  Theories of pulsar glitches centre around the mass unpinning model first proposed by \\cite{Anderson} and extended by many others \\citep{Alpar84,Pines85,Alpar89,Link02}.  The neutron superfluid in the stellar interior is threaded by many ($\\sim 10^{16}$) vortices, approximately one percent of which are pinned to the stellar crust at grain boundaries and/or nuclear lattice sites \\citep{Alpar89}.  As the pulsar crust spins down electromagnetically, a lag builds up between the velocity of the pinned vortex lines (corotating with the crust) and the superfluid.  When the transverse Magnus force (directly proportional to the lag) surpasses a threshold value (equal to the strength of the pinning force), a catastrophic unpinning of vortices occurs, transferring angular momentum to the crust.  In order for this mechanism to generate glitches on the scale observed, it requires up to $10^{12}$ vortices to unpin simultaneously, exhibiting a high level of collective, non-local behaviour.  The mass unpinning model spawned much activity in quantifying the microphysics that would lead to such macroscopic behaviour \\citep{Pines85,Jones97,Jones98a,Jones98b,DeB99,Elgaroy01,Donati03,Avogadro}.  In particular, the strength of the pinning force has recently been discussed in detail using a self-consistent, semi-analytical model for the vortex-nucleus interaction \\citep{Avogadro}, yielding the unexpected result that the strongest pinning occurs in regions of the crust that are of high and low density, rather than intermediate density \\citep{Donati06}.  An interest has also been taken in the evolution and morphology of the neutron star crust, with particular attention paid to impurities and defects, which determine the density of pinning centres available to the vortices, and their relative strengths \\citep{DeB98}.  \\cite{Link02} looked at the role that precession, if present, plays in the unpinning process, and discussed the dependence of the pinnning force on various microphysical parameters of the superfluid vortices and of the nuclear crustal lattice.  Several other theories of pulsar glitches, similarly focused on the microphysics of the superfluid in the stellar interior and crust, are currently under consideration.  These include thermally driven glitches arising from sudden changes in the frictional coupling within the neutron star \\citep{Link96}, superfluid turbulence \\citep{Peralta05,Peralta06a,Peralta06b,Melatos07b}, superfluid two-stream instabilities \\citep{Andersson04,Mastrano05}, crust cracking \\citep{Ruderman69,Bak96,Alpar96,Middleditch06} and mass accretion \\citep{Morley96a,Morley96b}.  Regardless of the model one chooses to examine, there is a common need to synthesise the microphysics with the global, \\emph{collective} dynamics in order to accurately describe the observational data.  In order to produce a realistic model of the global dynamics, some subtle aspects of the microphysics need to be sacrificed.  Cellular automaton models of complex systems aim to employ only the dominant aspects of the microphysics to generate replicable and efficient automaton rules.  Such an approach was attempted by \\cite{Morley96a} using a cellular automaton platelet model of mass accretion onto the stellar crust.  Although the microphysical details differ from those treated in this paper, the underlying philosophy of using a simple model to describe large scale behaviour is identical to ours.  We are similarly encouraged by the marked parallels between superfluid vortices and those that populate hard, type II superconductors.  The agreement between cellular automaton models of flux creep and vortex avalanches in superconductors and experimental data is excellent \\citep{Linder,Jensen,Field,Bassler,Nicodemi01a,Nicodemi01b}, in particular with respect to the measurement of a power law over several decades in the avalanche (cf. glitch) size.  A general, first principles theory of self-organized criticality does not yet exist for any SOCS, let alone for pulsar glitches.  In this paper we look for empirical agreement between the output of our cellular automaton and the gross qualitative features in pulsar glitch data.  A detailed quantitative comparison between a first principles theory and data is not possible at this stage.  In this paper we present a cellular automaton model of pulsar glitches based on the mass unpinning model.  After describing the key features of a SOCS we revisit the observational pulsar glitch data in \\S\\ref{sec:SOCS}.  \\S\\ref{sec:unpinning} reviews the vortex unpinning paradigm and recent advances in the understanding of neutron star structure which are relevant to the paradigm.  \\S\\ref{sec:model} describes in detail the cellular automaton model and justifies physically the automaton rules.  In \\S\\ref{sec:results} the statistical behaviour of the model is explored with respect to driver parameters, internal physical parameters, and variations in the automaton rules.  Further discussion of the model's validity and a summary of our findings are contained in \\S\\ref{sec:conc}. ",
        "conclusions": "\\label{sec:conc}  In this paper we present a cellular automaton model for glitches in neutron stars, based on the vortex unpinning paradigm \\citep{Anderson,Alpar84}.  Despite the gross idealisations in such a discrete model, the resulting statistical behaviour agrees with recent analyses of radio timing data \\citep{Wang00,Wong01,Melatos07a}.  Of particular interest are the power law distributions of glitch sizes and durations produced by the model, providing further evidence that glitches are a scale invariant phenomenon.  For this reason, as observations improve in resolution and lengthen in duration, we expect to discover smaller and larger glitches than have been seen so far.  Moreover, the exponential waiting-time distribution generated by the model supports the idea that the each vortex avalanche is a statistically independent event.  The values of the power law exponents and Poisson glitching rates depend on the particular parameter set. Our fits return size and duration exponents in the ranges $2.0\\leq a\\leq 4.3$ and $2.2\\leq b\\leq 5.5$ respectively, and mean glitching rates in the range $0.0023\\leq\\lambda\\leq0.13$.  Although we do not have a first-principles theory against which to compare our results, we claim that our cellular automaton is consistent with a SOCS.  The operation of the cellular automaton relies heavily on the assumption that vortices are inhomogeneously distributed.  \\cite{Alpar96} suggested that crust cracking creates capacitive regions where many vortices pin simultaneously.  These regions do not permit a vortex current, so pinned vortices in these regions do not participate in outward vortex creep.  Crust cracking is a nonuniform process, suggesting that the regions themselves are inhomogeneously distributed.  Our model conforms qualitatively with this picture.  Of particular interest is the claim by \\cite{Alpar95} that the neutron star crust is divided coarsely into depleted and capacitive regions. If verified, this claim suggests that the level of coarse graining necessary to render our cellular automaton computationally tractable matches the level of inhomogeneity that is realistically presesnt in a neutron star.  Of course, the opposite viewpoint is also viable:  the crust forms a large, defect-free crystal, allowing the vortices to occupy a contiguous sequence of pinning positions and preventing the type of inhomogeneity represented by our initial conditions \\citep{Jones98a}.  Importantly, such a scenario challenges \\textit{any} model based on collective unpinning of numerous vortices, not just our cellular automaton.  \\cite{DeB98} concluded that the crust is relatively free of microcrystalline structures and vacancies, with impurities, or point defects, the most likely disruptions to the crystal lattice (faults and dislocations were not considered).  Alternatively, if pinning occurs mostly at grain boundaries and cracks, as originally postulated by \\cite{Anderson}, then inhomogeneity on macroscopic scales is possible.   Irrespective of the density and distribution of pinning sites in reality, the model assumes for simplicity that vortex pinning occurs throughout the pulsar.  Several recent studies have shown that pinning forces sufficient to oppose the Magnus force develop only in specific zones of the crust, e.g. regions of high and low, rather than intermediate, denisty \\citep{Avogadro}.  Restricting the simulation to these optimal pinning zones obviates the need to allow for a range of pinning strengths if the zones are thin layers. A model in which pinning does not occur everywhere in the star leads to a different computational topology, e.g. an annulus, or several nested annuli.  However, we choose to be conservative in this introductory paper by not imposing a grid topology a priori, in order to see if a preferred topology emerges by itself.  In fact, it does; we find that a narrow active annulus, situated at the critical radius, participates in the avalanches.  However, it is important to experiment with other topologies in future work, because the active annulus we find in this paper may still be an artifact. If the active region lies near the rotation axis, an automaton executed in an annulus is free of the shortcomings introduced by the `permanently unpinned' regions mentioned in \\S\\ref{subsub:frac}.  To use the full circular grid in a meaningful way, we must include pinning of proton and neutron vortices in the core of the neutron star, together with pinning of neutron vortices in the inner crust.  In this case, the additional physics of entrainment \\citep{Link96,Ruderman98,Sedrakian95,Andersson04} and induced electric currents implies a more detailed set of automaton rules, outside the scope of the model presented here.  It should also be noted that if the superfluid vortices are tangled \\citep{Peralta06a,Peralta06b}, rather than forming a regular Abrikosov array, these two-dimensional annular and circular topologies are clearly too restrictive.  The locations of pinning zones also determine the long-term fate of a pulsar as it spins down.  If pinning occurs everywhere, spin down increases the proportion of vortices with $|\\Delta\\mathbf{v}|>|\\Delta\\mathbf{v}_{\\rm{pin}}|$ which are always supercritical (outside the critical circle) and do not participate in avalanche dynamics as they never pin.  If pinning is restricted to one or more annuli, the glitch behaviour of the pulsar changes significantly when the critical circle moves into and out of the pinning zones.  Clearly, our cellular automaton does not include the response of the pulsar to glitches (ie. the spin up of the pulsar crust and the subsequent `relaxation').  Nor do we account for the finite time-scale on which the unpinned superfluid couples to the crust, governed by the Hall-Vinen-Bekarevich-Khalatnikov equations \\citep{Peralta05,Peralta06a}.  We make the approximation that this time-scale is considerably less than the time-scale of our big time steps, motivated by glitch timing data (where the post-glitch recovery phase typically ends well before the next glitch; cf. Vela).  To elicit avalanches from our model, we require fine tuning in both the physical and computational parameters, such that $|\\Delta\\mathbf{v}_{\\rm{max}}|\\approx|\\Delta\\mathbf{v}_{\\rm{pin}}|$.  This condition ensures that there are enough vortex bundles that switch between becoming sub- and supercritical as time passes.  We achieve this by choosing $\\nu$, $\\epsilon$ and $\\Delta\\mathbf{v}_{\\rm{pin}}$ to place the critical circle near the surface of the star.  We emphasize that for a larger (smaller) star, we would resize the critical circle proportinally, for computational rather than physical reasons.  Fine tuning is also required in the ratio of vortex bundles to grid cells $B$.  For $B\\gg 1$, the automaton output is no longer consistent with a SOCS.  In conclusion, we present an empirical cellular automaton model of pulsar glitches based on the mass vortex unpinning paradigm.  We find that for certain physical and computational parameters the model produces dynamics that are consistent with a SOCS and with radio timing data from pulsars.  There exists no general, first-principles theory of SOC , let alone of pulsar glitches, so many of our results are empirical.  In particular, ther is no way known at present to predict theoretically the size and duration distribution exponents, and the mean rate of the waiting-time distribution.  We do demonstrate that the basic physical principles governing inter-vortex interactions can produce the type of \\textit{collective} behaviour necessary to explain pulsar glitches within the mass unpinning paradigm, especially the puzzle of how so many vortices can unpin in sympathy during a glitch, and why their number varies so much from glitch to glitch.  We thank Carlos Peralta for sharing his up-to-date glitch catalogue, and Stuart Wyithe for advice on statistical methods.  LW acknowledges the support of an Australian Postgraduate Award."
    },
    "0806/0806.0431_arXiv.txt": {
        "abstract": "We present the discovery of PSR~J1410$-$6132, a 50-ms pulsar found during a high-frequency survey of the Galactic plane, using a 7-beam 6.3-GHz receiver on the 64-m Parkes radio telescope.  The pulsar lies within the error box of the unidentified EGRET source 3EG J1410$-$6147, has a characteristic age of 26 kyr and a spin-down energy of 10$^{37}$~erg s$^{-1}$. It has a very high dispersion measure of 960 $\\pm10$ cm$^{-3}$ pc and the largest rotation measure of any pulsar, RM=$+2400 \\pm30$ rad m$^{-2}$. The pulsar is very scatter-broadened at frequencies of 1.4 GHz and below, making pulsed emission almost impossible to detect. Assuming a distance of 15~kpc, the pulsar's spin-down energy and a $\\gamma$-ray efficiency factor of $\\sim$10 per cent is sufficient to power the $\\gamma$-ray source. We therefore believe we have identified the nature of 3EG J1410$-$6147. This new discovery suggests that deep targeted high-frequency surveys of inner-galaxy EGRET sources could uncover further young, energetic pulsars. ",
        "introduction": "\\label{intro}  The Galactic population of pulsars is still poorly known in the inner parts of our Galaxy. Studies of the Galactic structure (e.g.~Bahcall 1986, Gilmore et al.~1989)\\nocite{bah86,gwk89} and studies of the radial distribution of pulsars (e.g.~Lyne et al.~1985)\\nocite{lmt85} however, suggest that a large number of pulsars await discovery in the inner spiral arms. Unfortunately, selection effects imposed by the interstellar medium make the discovery of such pulsars difficult. First, the interaction of the broadband pulsed radio signal with the ionised component of the interstellar medium causes the effect of dispersion, delaying signals observed at lower frequencies relative to their higher frequency counterparts (Hewish et al.~1968)\\nocite{hbp+68}.  The dispersive delay is inversely proportional to the square of the observing frequency, the constant of proportionality being the column density of the free electrons between Earth and the pulsar, known as the dispersion measure (DM). Without a correction for this effect, the recorded pulse becomes smeared and eventually undetectable. Secondly, the pulse also becomes broadened due to interstellar multi-path scattering as the free electron distribution in the turbulent interstellar medium is inhomogeneous \\cite{r77}.  Scattering is especially severe at low frequencies, since the scattering time is proportional to $\\nu^{\\beta}$ where $\\nu$ is the observational frequency and $\\beta\\ga-4$ (L\\\"ohmer et al.~2004)\\nocite{lmg+04}.  Unlike dispersion, scattering cannot be removed by instrumental means. Both effects are particularly severe for high DM pulsars. Another factor that mitigates against low frequency surveys is that at low Galactic latitudes, the background (synchrotron) radiation has an effective temperature, T$_{sky}$, which dominates T$_{sys}$ at low frequencies with a dependence which varies as approximately $\\nu^{-2.6}$.  Any pulsar survey therefore needs to balance the above effects against the fact that pulsars are intrinsically brighter at low frequencies and the larger telescope beam at lower frequencies reduces the survey time for a given area. Population studies take these selection effects into account and a recent population study by Lorimer et al.~(2006)\\nocite{lfl+06} suggests that, using the electron density model of Cordes \\& Lazio (2002)\\nocite{cl02}, a dearth of pulsars exist in the inner Galaxy.  As uncertainties in the electron model remain (cf.~Kramer et al.~2003, Lorimer et al.~2006)\\nocite{kbm+03} a high-frequency survey is the only direct way to shed light on the population of pulsars in the inner Galaxy, as was previously demonstrated by the pioneering use of high frequencies in pulsar surveys by Clifton \\& Lyne (1986) \\nocite{cl86}. This is particularly true for the studies of young pulsars which are especially important for our understanding of birthrates derived in these population studies. Indeed, for young pulsars still residing near their birth place in the Galactic plane, the aforementioned selection effects are extremely severe as they are typically spinning fast. These considerations were the motivation to conduct a survey of the inner Galactic plane at an unusually high frequency of 6.3 GHz, of which the first results are presented here.  The new survey utilises a seven beam receiver that was built in collaboration between Jodrell Bank Observatory and the Australia Telescope National Facility, operating at a wavelength around 5\\,cm, to search the Galaxy for emission of methanol masers. Exploiting this ``methanol multi-beam'' (MMB) receiver (each beam with a width of 0.11 deg) allows us to rapidly cover the inner Galactic plane at a central observing frequency of 6306\\,MHz, with a bandwidth of 576\\,MHz, which is indeed much higher than for usual pulsar surveys. As a result, we expect the survey to discover a sample of new pulsars that is dominated by young objects.  We report here the first discovery of the MMB survey which, in accordance with the expectations, is a young, energetic pulsar. In Section 2 we briefly outline the parameters of the MMB survey and in Section 3 discuss the pulsar's parameters and its discovery. Section 4 discuss the likely association between the pulsar and the unidentified $\\gamma$-ray point source, 3EG J1410$-$6147. ",
        "conclusions": "A survey with the Parkes telescope at the high frequency of 6.3~GHz has been carried out along a thin strip of the Galactic plane. We report here on the discovery of a young, highly energetic pulsar located in the error box of the $\\gamma$-ray source 3EG J1410$-$6147. The parameters of the pulsar indicate that the association is highly likely, although confirmation awaits the reduction of the $\\gamma-$error box and/or the detection of pulsations with AGILE or GLAST."
    },
    "0806/0806.1956_arXiv.txt": {
        "abstract": "We seek to compute the fraction of the volume of the Universe filled by expanding cocoons of the cosmological population of radio galaxies over the Hubble time as well as the magnetic field infused by them, in order to assess their importance in the cosmic evolution of the Universe. Using N-body $\\Lambda$CDM simulations, radio galaxies distributed according to the observed radio luminosity function are allowed to evolve in a cosmological volume as using well defined prescriptions for their expansion. We find that the radio galaxies permeate $10 - 30\\%$ of the total volume with $\\sim 10^{-8}$ G magnetic field by the present epoch. ",
        "introduction": "\\label{sec-intro}  Radio galaxies (RGs) are believed to have significant impact on the formation and evolution of large scale structures in the Universe. The cosmological population of expanding RGs and quasars can permeate large volumes of the intergalactic medium (IGM) and hence impact a considerable fraction of the filamentary protogalactic structures and could contribute substantially toward magnetization and metal enrichment of the Universe \\citep[e.g.,][]{GKW01, kronberg01, FL01, barai04, GWB04, LG05}.  The expansion of shocked and overpressured radio cocoons in a two-phase IGM are argued to compress the cold clouds and trigger star (perhaps even dwarf galaxy) formation \\citep{deYoung89, rees89, daly90, chokshi97, natarajan98, vanBreugel04, silk05}, as supported by recent observations of jet-induced star formation \\citep[e.g.,][]{dopita07, reuland07}. At the same time, some works indicate that RG expansion inhibits star formation by expelling (and heating) the IGM gas % \\citep[e.g.,][]{rawlings04, schawinski06, fujita08}. Other studies \\citep[e.g.,][]{nath02, vernaleo07, mcNamara07} reveal that the RGs heat up the ICM in galaxy clusters.  A key step to quantify the large scale impact RGs have is to address the question that how much of the volume of the Universe do the cosmological population of radio cocoons occupy over the Hubble time, which we seek to answer in the present work.  {From} rough calculations, \\citet{GKW01} argued that the expanding lobes of the generations of RGs can pervade up to 0.5 of the WHIM (warm/hot intergalactic medium) component in the Universe over $z \\sim 1-3$. \\citet{barai06} performed Monte Carlo simulations to construct virtual radio surveys \\citep{BRW, wang08}, and \\citet{barai07} estimated the cumulative volume filling factor to be $\\sim0.05$. These results were expressed as a fraction of the WHIM volume, adopted from the numerical simulations of \\citet{cen99}.  In this work we perform self-consistent cosmological simulations to compute the fractional volume of the Universe occupied by RGs, a more rigorous approach than previous attempts. Finding such volume filling fractions is important to probe in more detail the cosmological impact of RGs. We also perform preliminary estimates of the magnetic field infused in the filled volumes. The simulation method and model are described in \\S\\ref{sec-model}, and the results and discussion are in \\S\\ref{sec-results}. ",
        "conclusions": "\\label{sec-results}  Figure~\\ref{fig1} shows the redshift evolution of two single RGs in the simulation; here we discuss only the one with $\\tau_{\\rm RG} = 100$ Myr ({\\it red} curves). At the end of the active phase ($z = 5.57$) its cocoon is overpressured by a factor of $\\sim650$. So it continues to expand while its pressure falls faster because of adiabatic losses. Finally when $p_c$ falls to a level to match the external pressure it does not expand anymore. From $z = 1.85$ its comoving radius remains constant at $1.4 h^{-1}$ Mpc in the passive Hubble phase.   \\begin{figure} \\includegraphics[width = 3.2 in]{f1.eps} \\caption{ Characteristic quantities for the evolution of two single RGs: {\\it red} curves for one with $\\tau_{\\rm RG} = 100$ Myr and $Q_0 = 1.1 \\times 10^{44}$ erg/s, {\\it blue} curves for another with $\\tau_{\\rm RG} = 500$ Myr and $Q_0 = 1.9 \\times 10^{43}$ erg/s. Upper panel: Comoving size ($R_h$ during active-jet, and $R_c$ during spherical expansion) ({\\it red-solid} and {\\it blue-solid} curves with y-axis labels at top-right), and mean ambient gas density within the cocoon volume ($\\overline{\\rho_x}$) ({\\it red-dashed} and {\\it blue-dashed} curves with y-axis labels at top-left). Lower panel: Cocoon pressure $p_c$ ({\\it red-solid} and {\\it blue-solid} curves with y-axis labels at bottom-right), and overpressure factor of the cocoon w.r.t. external medium $p_c / p_x$ ({\\it red-dashed} and {\\it blue-dashed} curves with y-axis labels at bottom-left). The vertical lines separate the expansion phases of the RGs: active-AGN, post-AGN overpressured and the final passive Hubble evolution.} \\label{fig1} \\end{figure}  This illustrates that the radio cocoons are persistently overpressured for a substantial period of time even after the AGN has stopped activity, and hence continue to expand into the ambient medium. In Figure~\\ref{fig1}, after a active life of $100$ Myr, the RG remains overpressured for $\\sim3000$ Myr. Such results are in accord with other studies \\citep[e.g.,][]{yamada99, kronberg01}. In our simulations $\\sim 20 - 50 \\%$ (depending on the active lifetime) of the sources became spherical in shape during the active-AGN phase.  In order to prevent overcounting of the volume due to overlap of RGs, we count the mesh cells in the simulation box which occur inside the volume of one or more RG cocoons. The total number of these filled cells, $N_{\\rm RG}$, give the total volume of the box occupied by RGs. We express the total volume filled as a fraction of volumes of various overdensities in the box, $N_{\\rho} = N (\\rho > {\\cal C} \\overline{\\rho})$, where $\\overline{\\rho} = (1+z)^3 \\Omega_M 3H_0^2 / (8 \\pi G)$ is the mean matter density of a spatially flat Universe (the box) at an epoch $z$. So $N_{\\rho}$ gives the number of cells which are at a density ${\\cal C}$ times the mean density. We find $N_{\\rho}$ for ${\\cal C} = 0, 1, 2, 3, 5, 7$; ${\\cal C} = 0$ gives the total volume of the box, since then $N_{\\rho} = N (\\rho > 0) = 512^3$ is the total number of cells in the box.  Figure~\\ref{fig2} shows the redshift evolution of the volume filling factors for different active lifetimes. By the present epoch, $0.08$ of the entire Universe is filled by RGs of active lifetime $10$ Myr, the fraction going up to $0.26$ for 100 Myr, and $0.32$ for 500 Myr. With $\\tau_{\\rm RG} = 100$ or $500$ Myr, RGs fill up all of the the overdense regions ($\\rho > \\overline{\\rho}$, or higher) by $z=0.3-0.4$. With $\\tau_{\\rm RG} = 500$ Myr, RGs always fill up regions with $\\rho > 5 \\overline{\\rho}$ or higher at all epochs. The case with $\\tau_{\\rm RG} \\propto 1/\\sqrt{Q_0}$ fills up $0.24$ of the Universe by $z=0$, and give volume filling fractions similar to the values with a constant lifetime of 100 Myr.   \\begin{figure} \\includegraphics[width = 3.2 in]{f2.eps} \\caption{Volume filled by RGs ($N_{\\rm RG}$) as a fraction of total volume of the simulation box {\\it (black)}, and as a fraction of volumes of various overdensities: $N (\\rho > \\overline{\\rho})$ {\\it (blue)}, $N (\\rho > 2 \\overline{\\rho})$ {\\it (red)}, $N (\\rho > 3 \\overline{\\rho})$ {\\it (violet)}, $N (\\rho > 5 \\overline{\\rho})$ {\\it (orange)}, $N (\\rho > 7 \\overline{\\rho})$ {\\it (turquoise)}. The panels from top to bottom are for active RG lifetimes of $\\tau_{\\rm RG} = 10, 100, 500$ Myr, and for $\\tau_{\\rm RG} \\propto 1/\\sqrt{Q_0}$.} \\label{fig2} \\end{figure}  It is the overdense cosmic regions which gravitationally collapse to form stars and galaxies. So evidently the RGs have a profound impact on the protogalactic regions of the Universe. The precise effect on star formation is still open to debate (\\S\\ref{sec-intro}), with possible RG influence on both triggering and suppressing star formation in different regions of the Universe depending on the exact ambient conditions.  Our volume filling factors of $10-30\\%$ are between the values that \\citet{GKW01} ($50\\%$) and \\citet{barai07} ($\\lesssim 5\\%$) obtained as a fraction of the volume of the WHIM component of the Universe. Our results, based on self-consistent cosmological simulations, give a more reliable estimate of the fractional volume of the Universe occupied by RGs. The volumes obtained by \\citet{LG05} (100\\% filling by $z\\sim1$) are much higher, since they consider the whole AGN population.  We perform preliminary estimates of the energy density and magnetic field in the volumes of the Universe filled by radio cocoons. The cocoon energy density behaves similar to the cocoon pressure evolving adiabatically (\\S\\ref{sec-dead-AGN}) $u_E = 3 p_c$. Assuming equipartition of energy between magnetic field of strength $B_c$ and relativistic particles inside the cocoon, the magnetic energy density is $u_B = u_E / 2 = B_c^2 / (8 \\pi)$. The mean thermal energy density of the ambient medium inside the RG volume is $\\overline{u_{T,x}} = 3 \\overline{\\rho_x} k T_x / (2 \\mu)$. We define the volume weighted average of a physical quantity ${\\cal A}$ as $\\langle {\\cal A} \\rangle (z) \\equiv \\sum ({\\cal A} V_{\\rm RG}) / \\sum V_{\\rm RG}$, where the summation is over all RGs existing in the simulation box at that epoch.   \\begin{figure} \\includegraphics[width = 3.2 in]{f3.eps} \\caption{% The volume weighted average of the total energy density inside cocoon volumes $\\langle u_E \\rangle$ (top), ratio of the magnetic energy density to the mean external thermal energy density $\\langle u_B/\\overline{u_{T,x}} \\rangle$ (middle), and the equipartition magnetic field within RG filled volumes $\\langle B_c \\rangle$ (bottom). The color of a curve indicate its lifetime: 10 Myr ({\\it red}), 100 Myr ({\\it violet}), 500 Myr ({\\it blue}), and $\\tau_{\\rm RG} \\propto 1/\\sqrt{Q_0}$ ({\\it orange}).} \\label{fig3} \\end{figure}  Figure~\\ref{fig3} shows the redshift evolution of $\\langle u_E \\rangle$, $\\langle u_B/\\overline{u_{T,x}} \\rangle$ and $\\langle B_c \\rangle$. The energy densities and magnetic field decrease with redshift as the filled volumes get bigger. The ratio $\\langle u_B/\\overline{u_{T,x}} \\rangle$, giving the importance of cocoon magnetic energy over external thermal energy, has a trend similar to that deduced by \\citet{FL01}. We find that, by the present, $u_B$ is comparable to $\\overline{u_{T,x}}$ or greater by factors of few, implying that substantial magnetic energies are infused into the IGM by the expanding radio cocoons. At $z = 0$, a magnetic field of $\\sim 10^{-8}$ G permeates the filled volumes, consistent with the results of \\citet{GKW01} and \\citet{ryu98}. At a given redshift, the energy density and magnetic field are larger for higher source lifetimes. The results for $\\tau_{\\rm RG} \\propto 1/\\sqrt{Q_0}$ are intermediate between those of 10 and 100 Myr.   We conclude that using our N-body cosmological simulations, the expanding population of RGs pervade $10-30\\%$ of the volume of the Universe by the present, and occupy $100\\%$ of the overdense regions by $z \\sim 0.3$. A magnetic field of $\\sim 10^{-8}$ G is infused in the filled volumes at $z = 0$."
    },
    "0806/0806.4097_arXiv.txt": {
        "abstract": "{In the past few years, a new class of High Mass X-Ray Binaries (HMXRB) has been claimed to exist, the Supergiant Fast X-ray Transients (SFXT). These are X-ray binary systems with a compact companion orbiting a supergiant star which show very short and bright outbursts in a series of activity periods overimposed on longer quiescent periods. Only very recently the first attempts to model the behaviour of these sources have been published, some of them within the framework of accretion from clumpy stellar winds. } {Our goal is to analyze the properties of \\object{XTE~J1739-302}/\\object{IGR~J17391-3021} within the context of the clumpy structure of the supergiant wind. } {We have used {\\it INTEGRAL}~ and {\\it RXTE}/PCA observations in order to obtain broad band (1\\,--\\,200 keV) spectra and light curves of \\object{XTE~J1739-302} and investigate its X-ray spectrum and temporal variability. } {We have found that \\object{XTE~J1739-302} follows a much more complex behaviour than expected. Far from presenting a regular variability pattern,   \\object{XTE~J1739-302} shows periods of high, intermediate, and low flaring activity. } {} ",
        "introduction": "Wind-fed Supergiant X-Ray Binaries (SGXRBs) display high energy emission arising from the accretion of material in the wind of an OB supergiant onto the compact component of the system (a neutron star -NS- or black hole in orbit around the supergiant).  SGXRBs are persistent X-ray sources, displaying an X-ray luminosity $L_{{\\rm X}}\\sim10^{36}\\:{\\rm erg}\\,{\\rm s}^{-1}$. Because of the physical characteristics of wind accretion, their emission is variable on short timescales, with frequent flares, but relatively stable on the long term (for example, the long-term {\\it RXTE}/ASM lightcurve of Vela X-1, averaged and smoothed with a running window of 30\\,d length, shows variations by only a factor of $\\sim4$;  \\citealt{ribo06}). If the orbit is eccentric, the luminosity is  modulated on the orbital period of the system \\citep[e.g.,][]{leahy02}. Stronger short flares, with a fast rise and a typical timescale of the order of a few hours, have been observed from several systems, such as Vela X-1 \\citep{lau95,krivonos03} or 4U~1907+09 \\citep{fritz06}.  Recently, thanks to the improved sensitivity of high energy missions, many new SGXRBs have been discovered, leading to the suggestion of new classes of X-ray sources. On the one hand, there is a number of highly absorbed SGXRBs, invisible to previous missions due to high absorption in the softer X-ray bands \\citep[e.g.,][]{chaty05}. On the other hand, Supergiant Fast X-ray Transients (SFXTs) display fast outbursts, with a typical duration of a few hours, but stay in quiescence most of the time \\citep{smith06,neg06a,sgue06}. Unlike in classical SGXRBs, the X-ray luminosity of SFXTs goes down below the sensitivity limit of the INTErnational Gamma-Ray Astrophysics Laboratory ({\\it INTEGRAL}) and they remain undetectable for long time spans. They can only be observed during an outburst or flare, for a short time. Though several models have been proposed for this difference in behaviours, it seems to be a natural consequence of the clumpy nature of OB star winds \\citep{walter07,neg08}. Sidoli et al. (\\citeyear{sidoli07}) propose an alternative hypothesis, based on observations of \\object{IGR~J11215$-$5952}, in which the observed flaring activity is due to the interaction of the compact object with an extended equatorial decretion disc around the supergiant star.  Although the definition of SFXTs as a putative new class of objects was only possible when the optical counterparts to these systems started to be identified, {\\it INTEGRAL} has contributed decisively to the characterization of the high energy behaviour of these sources. So far, $\\sim 12$ SFXTs or related objects have been detected by {\\it INTEGRAL} \\citep{walter07,sgue06}. Among them, the best characterized system is  \\object{XTE J1739-302} = \\object{IGR~J17391-3021}, generally taken as the prototype of the class \\citep{smith06,neg06b}.  \\object{XTE~J1739-302} was discovered by {\\it RXTE} during a short outburst in 1997 \\citep{smith98}, when it was detected only for a period of a few hours. The source spectrum was well described by bremsstrahlung emission with a source temperature $kT\\sim21$~keV. No indications of any periodicity shorter than 300~s could be found. During 2003, the source was detected by {\\it INTEGRAL}/ISGRI \\citep{lutovinov05}. Again, a bremsstrahlung model with $kT\\sim22$~keV fitted the source spectrum well and no evidence of periodicity could be found.  A total of 6 outbursts were detected by {\\it INTEGRAL} up to 2005 \\citep{sgue05}. The mean duration of these outbursts is of the order of 5 hours and they are all highly structured. For a plot of {\\it INTEGRAL}/ISGRI detections in the 20\\,--\\,40 keV energy range during the period 2003\\,--\\,2005, see Fig.~\\ref{gps_lc}. The optical counterpart  was identified thanks to a {\\it Chandra} localization as an O8\\,Iab(f) supergiant at a distance of $\\approx2.3$~kpc \\citep{neg06b}.  In this work, we present a detailed analysis of {\\it INTEGRAL} data for \\object{XTE~J1739-302} obtained mostly through the Galactic Plane Scans (GPS, public data), the Galactic Center Deep Exposure (GCDE, public data), and through three long exposures of the Galactic Center, taken as part of the {\\it INTEGRAL} Key Programme (KP) observations. Section~\\ref{sec:obs} will be devoted to the description of the data and the analysis techniques used, including the presentation of results. An interpretation of these results in the context of our current understanding of SFXTs and models of accretion from clumpy winds \\citep{walter07,neg08} will be presented in Section~\\ref{sec:disc}, followed by our conclusions. ",
        "conclusions": "\\object{XTE~J1739-302}  has been observed with the {\\it INTEGRAL} observatory during a long time span. In the period 2003\\,--\\,2005, as part of the {\\it INTEGRAL} core program observations of the Galactic Center, the source showed moderate activity, with an average outburst frequency below 1~outburst~per day. During the deep exposures of the Galactic Center taken within the frame of our first Key Programme observations, in the 2006\\,--\\,2007 period, \\object{XTE~J1739-302} showed a higher level of activity, with a mean outburst frequency of $\\sim$3~outburst~per day. Surprisingly, during the last observations in 2007 (within the second run of our Key Programme) the source showed an unusually low activity state, and no outburst was detected with a flux above 43~mCrab. The behaviour during the first two periods can be explained well within the framework of the clumpy wind models. These models will not only explain the observed properties of the source but will also predict an orbital periodicity around $\\sim$8~d and suggests the presence of a NS as the compact companion. To explain the low activity observed during KP2 period, geometrical considerations related to the eccentricity of the orbit, or a drop in the mass loss from the supergiant companion, need to be invoked in order to maintain consistency with the proposed model. A continuous monitoring of the system will allow disentangling which of the three observed behaviours (non dectability, moderate activity, and high activity) is representing the {\\it normal} state of the source, will lead to the careful determination of the model parameters, and will permit to constrain the geometrical parameters of the system.  An independent constraint or measure of the orbital period of \\object{XTE~J1739-302} would help to support or discard the clumpy wind model as the explanation to the behaviour of \\object{XTE~J1739-302}. This model would, then, be the first consistent attempt to explain the observational properties of the SFXT and the classical wind-fed supergiant systems all together. The long spans of high activity seem incompatible with the model proposed by Sidoli et al. (2007), where outbursts happen once or twice each orbital cycle. Moreover, the observations presented here represent strong evidence against a coherent periodicity in the recurrence of the outbursts, which is a requirement of the model. It must be stressed that this model was specifically designed for \\object{IGR~J11215$-$5952}, a system presenting periodic outbursts, and its applicability to SFXTs is just a hypothesis, even if it turns to be appropriate for this particular source."
    },
    "0806/0806.3354_arXiv.txt": {
        "abstract": "Fully relativistic calculations of radiative rates and electron impact excitation cross sections for Fe\\,{\\sc x} are used to derive theoretical emission-line ratios involving transitions in the 174--366\\,\\AA\\ wavelength range. A comparison of these with solar active region observations obtained during the 1989 and 1995 flights of the Solar Extreme-ultraviolet Research Telescope and Spectrograph (SERTS) reveals generally very good agreement between theory and experiment. Several Fe\\,{\\sc x} emission features are detected for the first time in SERTS spectra, while the 3s$^{2}$3p$^{5}$ $^{2}$P$_{3/2}$--3s$^{2}$3p$^{4}$($^{1}$S)3d $^{2}$D$_{3/2}$ transition at 195.32\\,\\AA\\ is identified for the first time (to our knowledge) in an astronomical source. The most useful Fe\\,{\\sc x} electron density (N$_{e}$) diagnostic line ratios are assessed to be 175.27/174.53 and 175.27/177.24, which both involve lines close in wavelength and free from blends, vary by factors of 13 between N$_{e}$ = 10$^{8}$ and 10$^{11}$\\,cm$^{-3}$, and yet show little temperature sensitivity. Should these lines not be available, then the 257.25/345.74 ratio may be employed to determine N$_{e}$, although this requires an accurate evaluation of the instrument intensity calibration over a relatively large wavelength range. However, if the weak 324.73\\,\\AA\\ line of Fe\\,{\\sc x} is reliably detected, the use of 324.73/345.74 or 257.25/324.73 is recommended over 257.25/345.74. Electron densities deduced from 175.27/174.53 and 175.27/177.24 for the stars Procyon and $\\alpha$~Cen, using observations from the Extreme-Ultraviolet Explorer (EUVE) satellite, are found to be consistent and in agreement with the values of N$_{e}$ determined from other diagnostic ratios in the EUVE spectra. A comparison of several theoretical extreme-ultraviolet Fe\\,{\\sc x} line ratios with experimental values for a $\\theta$-pinch, for which the plasma parameters have been independently determined, reveals reasonable agreement between theory and observation, providing some independent support for the accuracy of the adopted atomic data. ",
        "introduction": "Emission lines arising from transitions in Fe\\,{\\sc x} have been widely detected in solar extreme-ultraviolet (EUV) spectra (see, for example, Dere 1978; Thomas \\& Neupert 1994). Jordan (1965) first proposed the use of Fe\\,{\\sc x} lines to determine the electron density in the solar corona, while Jordan (1966) employed EUV transitions of Fe\\,{\\sc x} and other Fe ions to derive both the electron density and coronal Fe abundance. Since then, several authors have undertaken analyses of the solar EUV spectrum of Fe\\,{\\sc x}, including for example Nussbaumer (1976) and Bhatia \\& Doschek (1995). To date, the most complete study is probably that of Del Zanna, Berrington \\& Mason (2004), which also provides an excellent review of previous work on Fe\\,{\\sc x}. The theoretical line ratios calculated by Del Zanna et al. employ radiative rates generated with the {\\sc superstructure} code by either Bhatia \\& Doschek (1995) or themselves. For electron impact excitation rates, they use results for transitions among the lowest 31 fine-structure levels of Fe\\,{\\sc x} calculated with the Breit-Pauli {\\sc rmatrx} code, either by Pelan \\& Berrington (2001) or once again by themselves.  Recently, Aggarwal \\& Keenan (2004, 2005) have used the fully relativistic {\\sc grasp} and Dirac {\\sc rmatrx} codes to calculate radiative rates and electron impact excitation cross sections, respectively, for transitions among the energetically lowest 90 fine-structure levels of Fe\\,{\\sc x}. In this paper we use these results, plus additional atomic data presented here, to generate theoretical emission-line ratios for Fe\\,{\\sc x}, and compare these with high resolution spectra from the Solar Extreme-ultraviolet Research Telescope and Spectrograph (SERTS). Our aims are threefold, namely to (i) assess the importance of blending in the SERTS observations, (ii) detect new Fe\\,{\\sc x} emission lines, and (iii) identify the best Fe\\,{\\sc x} line ratios for use as electron density diagnostics.  This work is of particular relevance due to the recent launch of the {\\em Hinode} mission, which has on board the EUV Imaging Spectrometer (EIS), covering the 170--211\\,\\AA\\ and 246--292\\,\\AA\\ wavelength ranges (Culhane et al. 2007), similar to the SERTS spectral coverage of 170--225\\,\\AA\\ and 235--450\\,\\AA\\ (Thomas \\& Neupert 1994). It is clearly important that emission lines observed by the EIS are fully assessed, and the best diagnostics identified. SERTS provides the ideal testbed for this, due to its larger wavelength coverage, allowing more lines from the same species to be detected and compared with theoretical predictions. Furthermore, the best SERTS spectral resolution is about 0.03\\,\\AA\\ [full width at half-maximum (FWHM)], obtained for the 170--225\\,\\AA\\ wavelength range observed in second-order (Brosius, Davila \\& Thomas 1998a), which is a full factor of two better than the 0.065--0.075\\,\\AA\\ resolution available from EIS (Young et al. 2007). Hence the SERTS data sets should allow emission features to be resolved and assessed which are blended in EIS spectra. Indeed, this is illustrated by our previous work on Fe\\,{\\sc xiii} (Keenan et al. 2007), where we resolved the 203.79 and 203.83\\,\\AA\\ features which are blended in EIS observations (Young et al. 2007).   \\section[]{Observational data}  SERTS has had a total of 10 successful flights over the period 1982--2000, each featuring a somewhat different set of optical components and observing targets.  We have selected data from two of these for analysis here, as they provide the most complete sequence of well-observed Fe\\,{\\sc x} emission lines.  The flight on 1989 May 5 (henceforth referred to as SERTS--89) carried a standard gold-coated toroidal diffraction grating, and was the first to observe a strong solar active region, NOAA 5464. It detected hundreds of first-order emission lines in the 235--450\\,\\AA\\ wavelength range, plus dozens of features spanning 170--225\\,\\AA, which appeared in second-order among the 340--450\\,\\AA\\ first-order lines (Thomas \\& Neupert 1994). The spectrum was recorded on Kodak 101--07 emulsion, at a spectral resolution of 50--80\\,m\\AA\\ (FWHM) in first-order, and a spatial resolution of approximately 7 arcsec (FWHM).  This combination of high spectral resolution, strong signal, and broad wavelength coverage remains unsurpassed even today as the best available data set for investigating imaged solar emission features over the full wavelength range of 235--450\\,\\AA, and hence are used in the present paper. Subsequent SERTS flights employed either a multilayer-coated diffraction grating or CCD-detector arrays which provided important technical advantages, but which also restricted the spectral bandpass that could be covered, making their data less suitable for the present study (see Keenan et al. 2007 and references therein).  The one exception was the flight on 1995 May 15 (henceforth referred to as SERTS--95) which incorporated a multilayer-coated toroidal diffraction grating that enhanced its sensitivity to second-order features in the 170--225\\,\\AA\\ wavelength range (Brosius et al. 1998a). It observed active region NOAA 7870, again using Kodak 101--07 emulsion, and had an improved spatial resolution of approximately 5 arcsec.  Its unique shortwave multilayer coating led to the detection of many second-order emission lines not seen on previous SERTS flights (Thomas \\& Neupert 1994; Brosius et al. 1996), and furthermore obtained the highest spectral resolution (0.03\\,\\AA\\ FHWM) ever achieved for spatially resolved active region spectra in this wavelength range.  The SERTS--95 data therefore provide the best observations for investigating Fe\\,{\\sc x} emission lines in the 170--225\\,\\AA\\ region, and hence are employed in this paper.  Further details of the SERTS--95 observations, and the wavelength and absolute flux calibration procedures employed in the data reduction, may be found in Brosius et al. (1998a,b). Similar information for the SERTS--89 spectrum is available from Thomas \\& Neupert (1994). We note that although the relative intensity calibration curves in both cases involved fitting to calculations of insensitive line-ratio pairs, Fe\\,{\\sc x} lines made up less than 10\\%\\ of the points utilized in those procedures. Furthermore, detailed comparisons of line ratio calculations for many ions with SERTS measurements have revealed generally very good agreement between theory and observation, providing support for the SERTS intensity calibrations (see, for example, Keenan et al. 2007 and references therein). Thus, even without fully independent intensity calibrations, the SERTS results still provide a valid test for the new Fe\\,{\\sc x} calculations presented here. It should also be noted that the SERTS--95 first-order calibration curve was extrapolated from its fitted range of 238--336\\,\\AA\\ to cover its complete bandpass. Hence there is some extra uncertainty in the intensities of lines with wavelengths greater than 340\\,\\AA\\ measured from this flight. However, no such problem exists with the SERTS--89 measurements of these lines.  We have searched for Fe\\,{\\sc x} emission lines in the SERTS--89 and SERTS--95 spectra using the detections of Thomas \\& Neupert (1994) and Brosius et al. (1998b), supplemented with those from other sources, including the NIST database,\\footnote{http://physics.nist.gov/PhysRefData/} the latest version (V5.2) of the {\\sc chianti} database (Dere et al. 1997; Landi et al. 2006), the Atomic Line List of van Hoof,\\footnote{http://www.pa.uky.edu/$\\sim$peter/atomic/} and in particular the excellent summary of line identifications by Del Zanna et al. (2004). The latter provides not only a comprehensive list of wavelengths for well-observed Fe\\,{\\sc x} transitions, but also indicates alternative wavelengths where previous identifications are not consistent with their conclusions. In Table 1 we list the Fe\\,{\\sc x} transitions found in the SERTS--89 and SERTS--95 spectra, along with their measured wavelengths. We also indicate possible blending features or alternative identifications as suggested by Thomas \\& Neupert or Brosius et al. in their original line lists.  \\begin{table*} \\centering \\begin{minipage}{140mm} \\caption{Fe\\,{\\sc x} line identifications in the SERTS 1989 and 1995 active region spectra.} \\begin{tabular}{cll} \\hline Wavelength (\\AA)    &   Transition  & Note$^{a}$ \\\\ \\hline 174.53 & 3s$^{2}$3p$^{5}$ $^{2}$P$_{3/2}$--3s$^{2}$3p$^{4}$($^{3}$P)3d $^{2}$D$_{5/2}$ \\\\ 175.27 & 3s$^{2}$3p$^{5}$ $^{2}$P$_{1/2}$--3s$^{2}$3p$^{4}$($^{3}$P)3d $^{2}$D$_{3/2}$ \\\\ 175.48 & 3s$^{2}$3p$^{5}$ $^{2}$P$_{3/2}$--3s$^{2}$3p$^{4}$($^{3}$P)3d $^{2}$P$_{1/2}$ \\\\ 177.24 & 3s$^{2}$3p$^{5}$ $^{2}$P$_{3/2}$--3s$^{2}$3p$^{4}$($^{3}$P)3d $^{2}$P$_{3/2}$ \\\\ 180.38 & 3s$^{2}$3p$^{5}$ $^{2}$P$_{1/2}$--3s$^{2}$3p$^{4}$($^{3}$P)3d $^{2}$P$_{1/2}$ & Blended with Fe\\,{\\sc xi} 180.38 + Fe\\,{\\sc xvi} 360.76 (first-order). \\\\ 184.53 & 3s$^{2}$3p$^{5}$ $^{2}$P$_{3/2}$--3s$^{2}$3p$^{4}$($^{1}$D)3d $^{2}$S$_{1/2}$ \\\\ 190.05 & 3s$^{2}$3p$^{5}$ $^{2}$P$_{1/2}$--3s$^{2}$3p$^{4}$($^{1}$D)3d $^{2}$S$_{1/2}$ \\\\ 193.72 & 3s$^{2}$3p$^{5}$ $^{2}$P$_{3/2}$--3s$^{2}$3p$^{4}$($^{1}$S)3d $^{2}$D$_{5/2}$ \\\\ 195.32 & 3s$^{2}$3p$^{5}$ $^{2}$P$_{3/2}$--3s$^{2}$3p$^{4}$($^{1}$S)3d $^{2}$D$_{3/2}$ \\\\ 201.56 & 3s$^{2}$3p$^{5}$ $^{2}$P$_{1/2}$--3s$^{2}$3p$^{4}$($^{1}$S)3d $^{2}$D$_{3/2}$ \\\\ 220.26 & 3s$^{2}$3p$^{5}$ $^{2}$P$_{3/2}$--3s$^{2}$3p$^{4}$($^{3}$P)3d $^{2}$F$_{5/2}$ \\\\ 256.43 & 3s$^{2}$3p$^{5}$ $^{2}$P$_{3/2}$--3s$^{2}$3p$^{4}$($^{3}$P)3d $^{4}$D$_{3/2}$ \\\\ 257.25 & 3s$^{2}$3p$^{5}$ $^{2}$P$_{3/2}$--3s$^{2}$3p$^{4}$($^{3}$P)3d $^{4}$D$_{5/2,7/2}$ \\\\ 324.73 & 3s$^{2}$3p$^{4}$($^{3}$P)3d $^{4}$D$_{7/2}$--3s3p$^{5}$($^{3}$P)3d $^{4}$F$_{9/2}$ \\\\ 337.24 & 3s$^{2}$3p$^{4}$($^{3}$P)3d $^{2}$F$_{7/2}$--3s3p$^{5}$($^{3}$P)3d $^{2}$F$_{7/2}$ & Listed as Ar\\,{\\sc viii} by Thomas \\& Neupert (1994). \\\\ 345.74 & 3s$^{2}$3p$^{5}$ $^{2}$P$_{3/2}$--3s3p$^{6}$ $^{2}$S$_{1/2}$ \\\\ 365.57 & 3s$^{2}$3p$^{5}$ $^{2}$P$_{1/2}$--3s3p$^{6}$ $^{2}$S$_{1/2}$ \\\\ \\hline \\end{tabular}  $^{a}$From Brosius et al. (1998b) or Thomas \\& Neupert (1994). \\end{minipage} \\end{table*}  Intensities and line widths (FWHM) of the Fe\\,{\\sc x} features are given in Tables 2 and 3 for the SERTS--89 and SERTS--95 active regions, respectively, along with the associated 1$\\sigma$ errors. These were measured using modified versions of the Gaussian fitting routines employed by Thomas \\& Neupert (1994), as discussed by Keenan et al. (2007). As a consequence, the intensities, FWHM values and their uncertainties listed in Tables 2 and 3 are somewhat different from those originally reported in Thomas \\& Neupert and Brosius et al. (1998b). Also, a uniform factor of 1.24 has been applied here to all SERTS--89 intensities, reflecting a more recent re-evaluation of its absolute calibration scale. Even so, in all directly comparable cases, the resulting line intensity values usually differ only slightly from those previously obtained. For the SERTS--95 data set, several of the stronger first-order Fe\\,{\\sc x} lines could also be detected, and their measurements are therefore included in Table 3 along with the second-order results. However, for the SERTS--89 spectrum only the first-order Fe\\,{\\sc x} features could be reliably identified.  \\begin{table} \\centering \\caption{Fe\\,{\\sc x} line intensities and widths from the SERTS 1989 active region spectrum.} \\begin{tabular}{ccc} \\hline Wavelength    &   Intensity & Line width \\\\ (\\AA) & (erg\\,cm$^{-2}$\\,s$^{-1}$\\,sr$^{-1}$) & (m\\AA) \\\\ \\hline 256.43 & 174.8 $\\pm$ 62.5 & 94 $\\pm$ 28 \\\\ 257.25 & 166.2 $\\pm$ 38.3 & 86 $\\pm$ 12 \\\\ 324.73 & 12.2 $\\pm$ 3.2 & 54 $\\pm$ 8 \\\\ 337.24 & 16.5 $\\pm$ 7.9 & 73 $\\pm$ 21 \\\\ 345.74 & 93.6 $\\pm$ 11.9 & 98 $\\pm$ 4 \\\\ 365.57 & 53.1 $\\pm$ 7.3 & 113 $\\pm$ 7 \\\\ \\hline \\end{tabular} \\end{table}  \\begin{table} \\centering \\caption{Fe\\,{\\sc x} line intensities and widths from the SERTS 1995 active region spectrum.} \\begin{tabular}{ccc} \\hline Wavelength    &   Intensity & Line width \\\\ (\\AA) & (erg\\,cm$^{-2}$\\,s$^{-1}$\\,sr$^{-1}$) & (m\\AA) \\\\ \\hline 174.53 & 655.4 $\\pm$ 82.0 & 41 $\\pm$ 3 \\\\ 175.27 & 133.5 $\\pm$ 30.5 & 27 $\\pm$ 5 \\\\ 175.48 & 79.4 $\\pm$ 20.9 & 41 $\\pm$ 9 \\\\ 177.24 & 335.2 $\\pm$ 41.9 & 39 $\\pm$ 3 \\\\ 180.38 & 3345.1 $\\pm$ 373.0 & 48 $\\pm$ 3 \\\\ 184.53 & 163.6 $\\pm$ 19.6 & 33 $\\pm$ 3 \\\\ 190.05 & 47.6 $\\pm$ 8.4 & 41 $\\pm$ 5 \\\\ 193.72 & 6.1 $\\pm$ 2.9 & 16 $\\pm$ 6 \\\\ 195.32 & 6.8 $\\pm$ 2.4 & 25 $\\pm$ 3 \\\\ 201.56 & 71.4 $\\pm$ 14.0 & 78 $\\pm$ 12 \\\\ 220.26 & 32.9 $\\pm$ 21.0 & 67 $\\pm$ 20 \\\\ 256.43 & 223.7 $\\pm$ 68.0 & 62 $\\pm$ 16 \\\\ 257.25 & 159.2 $\\pm$ 48.5 & 44 $\\pm$ 10 \\\\ 345.74 & 92.7 $\\pm$ 37.1 & 50 $\\pm$ 15 \\\\ 365.57 & 35.7 $\\pm$ 16.6 & 53 $\\pm$ 21 \\\\ \\hline \\end{tabular} \\end{table}  In Figs 1--5 we plot portions of the SERTS--89 and SERTS--95 spectra containing Fe\\,{\\sc x} transitions, focusing on emission lines which have not previously been identified in SERTS data sets. We note that several of these features have line intensities and widths comparable to the noise fluctuations. In these instances, the reality of the feature was confirmed by a visual inspection of the original SERTS film. However, the lines are weak and clearly further observations to strengthen these identifications would be desirable.  \\begin{figure} \\epsfig{file=fe10_serts_fig1.eps,angle=0,width=8.5cm} \\caption{Plot of the SERTS 1995 active region spectrum in the 175.0--175.7\\,\\AA\\ wavelength range. The profile fit to the Fe\\,{\\sc x} 175.27 and Fe\\,{\\sc ix}/{\\sc x} 175.48\\,\\AA\\ features is shown by a solid line.} \\end{figure}  \\begin{figure} \\epsfig{file=fe10_serts_fig2.eps,angle=0,width=8.5cm} \\caption{Plot of the SERTS 1995 active region spectrum in the 193.6--194.1\\,\\AA\\ wavelength range. The profile fit to the Fe\\,{\\sc x} 193.72\\,\\AA\\ feature is shown by a solid line. Also clearly visible in the figure are the Ca\\,{\\sc xiv} 193.87\\,\\AA\\ and Ni\\,{\\sc xvi} 194.03\\,\\AA\\ transitions.} \\end{figure}  \\begin{figure} \\epsfig{file=fe10_serts_fig3.eps,angle=0,width=8.5cm} \\caption{Plot of the SERTS 1995 active region spectrum in the 195.2--195.7\\,\\AA\\ wavelength range. The profile fit to the Fe\\,{\\sc x} 195.32\\,\\AA\\ feature is shown by a solid line. Also clearly visible is the wing of the strong Fe\\,{\\sc xii} 195.12\\,\\AA\\ transition.} \\end{figure}  \\begin{figure} \\epsfig{file=fe10_serts_fig4.eps,angle=0,width=8.5cm} \\caption{Plot of the SERTS 1995 active region spectrum in the 219.4--221.0\\,\\AA\\ wavelength range. The profile fit to the Fe\\,{\\sc x} 220.26\\,\\AA\\ feature is shown by a solid line. Also clearly visible in the figure is the Fe\\,{\\sc xiv} 220.09\\,\\AA\\ transition.} \\end{figure}  \\begin{figure} \\epsfig{file=fe10_serts_fig5.eps,angle=0,width=8.5cm} \\caption{Plot of the SERTS 1989 active region spectrum in the 324.1--325.2\\,\\AA\\ wavelength range. The profile fit to the Fe\\,{\\sc x} 324.73\\,\\AA\\ feature is shown by a solid line.} \\end{figure} ",
        "conclusions": "Our comparison of theoretical Fe\\,{\\sc x} emission-line intensity ratios with solar active region spectra from the SERTS 1989 and 1995 flights reveals generally very good agreement between theory and experiment, with several features identified for the first time in the SERTS data sets, including Fe\\,{\\sc x} 193.72, 220.26 and 324.73\\,\\AA, plus Fe\\,{\\sc xi} 201.56\\,\\AA\\ and the Fe\\,{\\sc ix}/{\\sc x} blend at 175.48\\,\\AA. In addition, the 195.32\\,\\AA\\ transition of Fe\\,{\\sc x} is detected for the first time (to our knowledge) in an astronomical source.  We find that the ratios 175.27/174.53 and 175.27/177.24 provide the best Fe\\,{\\sc x} electron density diagnostics, as they involve lines which are strong and free from blends, are close in wavelength and the ratios are highly N$_{e}$--sensitive. Should these lines not be available, then the 257.25/345.74 ratio may be employed as a diagnostic, although this requires an accurate determination of the instrument intensity calibration over a relatively large wavelength range. However, if the weak 324.73\\,\\AA\\ line is reliably detected, then the use of 324.73/345.74 or 257.25/324.73 is recommended in preference to 257.25/345.74."
    },
    "0806/0806.4742_arXiv.txt": {
        "abstract": "We re-examine the evidence for the existence of ultra-massive ($M>1.1M_\\odot$) white dwarfs based on gravitational redshift of white dwarfs in common proper motion binaries or in clusters, on parallax measurements, on orbital solutions, and, finally, on the analysis of hydrogen line profiles. We conclude that the best evidence is largely based on the analysis of Balmer line profiles although the companion to the A8V star HR~8210 is a compelling case made initially using the large binary mass function and confirmed by an analysis of the Lyman line spectrum. The confirmation and identification of high-mass white dwarfs, more particularly non-DA white dwarfs, using parallax measurements may prove critical in establishing the population fraction of these objects and in constraining the high-end of empirical initial-mass to final-mass relations. The existence of a substantial population of ultra-massive white dwarfs supports the concept of a steeper initial-mass to final-mass relations linking $6\\ M_\\odot$ progenitors with $\\ga 1.1 \\ M_\\odot$ white dwarfs. ",
        "introduction": "Ultra-massive white dwarfs, generally defined as having masses exceeding $1.1\\ M_\\odot$, remained relatively rare until deep spectroscopic surveys uncovered these intrinsically faint objects. The hot hydrogen-rich (DA) white dwarf GD~50 \\citep{ber1991}, the carbon-rich white dwarf G35-26 \\citep{the1990}, and two DA white dwarfs (PG~1658+441 and PG~0136+251) from the Palomar-Green survey \\citep{sch1992} were rare examples of this phenomenon. Interestingly, spectroscopic follow-up of extreme ultraviolet (EUV) surveys of the local ($d \\la 100$ pc), hot white dwarf population managed to identify many new ultra-massive white dwarfs on the basis of a {\\it large, spectroscopically determined surface gravity} ($\\log{g}\\ga 9$). Based on the {\\it Extreme Ultraviolet Explorer} ({\\it EUVE}) and ROSAT {\\it Wide Field Camera} ({\\it WFC}) surveys, \\citet{ven1996, ven1997b}, \\citet{ven1999}, \\citet{fin1997}, and \\citet{ mar1997} added a dozen new objects to the population. More recently, a re-analysis of the Palomar-Green (PG) sample of DA white dwarfs \\citep{lie2005} and a study of the white dwarf mass distribution in the Sloan Digital Sky Survey \\citep[SDSS; ][]{kep2007} seem to generate similar yields of ultra-massive white dwarfs.  Are these objects products of single star evolution or double degenerate mergers? \\citet{wei2000} reviews theoretical arguments in favour of an upper mass limit for white dwarf stars larger than the canonical upper limit of $1.1\\ M_\\odot$, and possibly as large as $\\approx 1.3\\ M_\\odot$. A limit of $1.1\\ M_\\odot$ is generally assumed because carbon ignition in high mass cores would lead to thermonuclear runaway or core collapse. However, the effect of mass loss may alter the scenario and lead to the formation of a massive oxygen-neon-magnesium white dwarf \\citep{nom1984}. \\citet{gar1997} and \\citet{rit1996} successfully evolved 9 and 10 $M_\\odot$ stars, which, following an off-centre carbon ignition in partial electron degenerate conditions, generated oxygen-neon core white dwarfs with carbon-oxygen shells, and total masses of $1.15$ and $1.26\\ M_\\odot$, respectively. On the other hand, the merger scenario proposed for the origin of Type Ia supernovae \\citep{yun1994} also generates ultra-massive white dwarfs. However, the merger process itself and the fate of these objects are uncertain \\citep{seg1997}.  We re-examine the evidence for the existence of ultra-massive white dwarfs. Mass measurements are based on gravitational redshift measurements (\\S 2), radius (parallax) measurements (\\S 3), orbital parameters (\\S 4), and surface gravity measurements (\\S 5). The gravitational redshift measurements ($\\propto M/R$), radius measurements, and surface gravity measurements ($\\propto M/R^2$) are converted into mass measurements by adopting mass-radius relations for a variety of model structures. We adopted the models of \\citet{ben1999} with a C/O core and zero metallicity, a helium mantle ($\\log{q}({\\rm He}) = -2$), and with a {\\it thick} ($\\log{q}({\\rm H})=-4$) hydrogen layer to describe hydrogen-rich DA white dwarfs or without a hydrogen layer to describe non-DA white dwarfs. In \\S 5, we also present evidence that deep surveys of high-proper motion white dwarfs are now extending the ultra-massive population toward lower temperatures, and it is estimated that $\\approx 10$\\% of white dwarfs have masses in excess of  $1.1\\ M_\\odot$. We summarize and conclude in \\S 6. ",
        "conclusions": "We critically reviewed current evidence for the existence of ultra-massive white dwarfs. First, we demonstrated that the high-mass white dwarfs listed in \\citet{sil2001} and presented as evidence for an ultra-massive white dwarf population in CPMB are in fact white dwarfs with normal masses ($\\approx 0.6$--$0.7\\ M_\\odot$), or, in the case of G275-B16A, a low-metallicity F star and possible halo member. The absence of any reliable candidates from this sample is puzzling considering the large number of objects identified using other methods. Improved radial velocity measurements of an enlarged sample of CPMB would help provide accurate gravitational redshifts and deliver the expected number of ultra-massive white dwarfs.  The study of the peculiar DBAZ GD~362 exposes potential difficulties in surface gravity, hence mass measurements based on hydrogen line profiles. Although spectroscopically evanescent at an effective temperature of 10\\,000 K , helium was found to be the dominant constituent in the atmosphere of that star. \\citet{zuc2007} demonstrated that the reduced hydrogen abundance had the consequence of lowering the surface gravity measurement to almost a normal level. Although helium lines are considerably stronger at temperatures in excess of 15\\,000 K, it is possible that some cooler ultramassive white dwarfs are in fact helium dominated but with a normal mass.  Next, we examined the evidence based on parallax measurements and found that the small number of measurements, and the white dwarf masses inferred from these measurements are in good agreement with the spectroscopic masses. Table~\\ref{tbl-2} lists spectroscopically-identified ultra-massive white dwarfs for which parallax measurements are desirable. Indeed, there remain some questions concerning the validity of our approaches, \\citet{hum1988} or \\citet{ing1939}, for Balmer and Lyman line merging at high density and temperature. The effect of perturbers on upper energy levels is essentially calibrated using normal gravity white dwarfs ($\\log{g} = 8$) and this calibration may not apply well at higher gravities potentially causing a systematic shift in mass measurements. However, a case-by-case review of the few spectroscopically identified high-mass white dwarfs with parallax measurements shows good agreement between the two methods. Table 2 lists the absolute visual magnitude and predicted parallax for each star based on parameters provided in the listed references. In addition to the calculated masses that we determined using the mass-radius relations assuming a CO core, we have also calculated mass estimates using mass-radius relations for white dwarfs with an ONe core \\citep{alt2005}, which may be more appropriate for massive white dwarfs. The mass-radius relations for ONe cores, predict masses that are systematically $\\approx 0.02\\ M_\\odot$ lower than those predicted by CO mass-radius relations. For masses larger than $1.3\\ M_\\odot$, we used the mass-radius relations of \\citet{ham1961}.  We also confirmed the high mass for the hot white dwarf in the binary HR~8210. A high mass was initially implied by the binary mass function \\citep{lan1993}. We fitted the Lyman line spectrum of the white dwarf and constrained the mass to be $M= 1.08 - 1.24\\ M_\\odot$ in agreement, in the upper mass range, with the binary parameters.  Finally, we show that the EUV-selected population of white dwarf stars is composed of $\\approx 10$\\% objects with masses in excess of $1.1\\ M_\\odot$. A similar yield was obtained by \\citet{lie2005} based on the PG survey and by \\citet{kep2007} based on SDSS {\\it after large corrections} were applied due to the magnitude-limited nature of the samples collected. In particular, it should be noted that only seven objects out of 347 from the Palomar-Green sample, or a fraction of 2\\%, met the criterion. By applying $V/V_{\\rm max}$ corrections due to incompleteness at fainter magnitudes, the estimated fraction was re-evaluated at 10\\% in agreement with the yield directly measured in the EUV selection.  The origin of ultra-massive white dwarfs remains uncertain. Initial-mass to final-mass relations \\citep{cat2008a} indicate that main sequence stars with masses in excess of $\\approx 6\\ M_\\odot$ generate white dwarfs with masses in excess $1\\ M_\\odot$, a situation best illustrated by the massive white dwarf (LB~1497) member of the Pleaides. By re-evaluating available cluster data \\citet{cat2008a} revised the final masses upward, and managed to reproduce the high-mass peak in both SDSS and PG empirical mass distributions. It is therefore possible that white dwarfs with masses in excess of $1.1\\ M_\\odot$ are the products of single star evolution and that the binary merger scenario may only apply to a minority of peculiar objects such as the fast rotating magnetic white dwarf WD~0325$-$857 \\citep[see][and references therein]{ven2003}.  The existence of a substantial population of ultra-massive white dwarfs supports the concept of a steeper initial-mass to final-mass relations linking $6\\ M_\\odot$ progenitors with $\\ga 1.1 \\ M_\\odot$ white dwarfs as proposed by \\citet{cat2008a}. Ultra-massive white dwarfs in close binaries are also likely Type Ia supernova progenitors \\citep{par2007}.  \\begin{table*} \\centering \\begin{minipage}{110mm} \\caption{NLTT and EUV spectroscopic sample of ultra-massive white dwarfs\\label{tbl-2}} \\begin{tabular}{clccccc} \\hline WD & Other names & Mass$^a$    & $M_V$ & Mass$^b$ & Predicted $\\pi$ & Ref. \\\\ &             & ($M_\\odot$) & (mag) & ($M_\\odot$) & (mas) & \\\\ \\hline 0457$-$004 & NLTT~14307          & $1.24\\pm0.02$ & 14.09 & $1.22\\pm0.02$ & 58.8 & 1\\\\ 1236$-$495 & NLTT~31372          & $1.11\\pm0.02$ & 13.13 & $1.09\\pm0.02$ & 66.7 & 2\\\\ 1653$+$256 & NLTT~43827          & $1.31\\pm0.01$ & 14.29 & $1.28\\pm0.01$ & 30.3 & 1 \\\\ 1729$+$371$^c$ & NLTT~44986          & $1.26\\pm0.03$ & 14.54 & $1.24\\pm0.03$ & 45.5 & 1\\\\ 2159$-$754 & NLTT~52728          & $1.17^{+0.06}_{-0.07}$ & 14.30 & $1.14^{+0.05}_{-0.06}$ & 71.4 & 2 \\\\ 0001$+$433 & EUVEJ0003$+$435     & $1.18\\pm0.04$ & 11.19 & $1.16^{+0.04}_{-0.05}$ &  7.2 & 3\\\\ 0136$+$251 & EUVEJ0138$+$253, PG & $1.13^{+0.04}_{-0.05}$ & 10.97 & $1.10\\pm0.05$ &  9.9 & 3\\\\ 0325$-$857 & EUVEJ0317$-$855     & $1.32\\pm0.03$ & 12.64$^d$ & $1.32\\pm0.03$ & 37.0$^d$ & 2, 4\\\\ 0340$+$103 & 1RXSJ0343$+$1030    & $1.22^{+0.04}_{-0.05}$ & 11.01 & $1.20^{+0.05}_{-0.06}$ &  8.0 & 3\\\\ 0346$-$011 & EUVEJ0348$-$009, GD~50 & $1.23\\pm0.02$ & 11.49 & $1.21\\pm0.02$ & 31.3 & 3\\\\ 0440$-$038 & EUVEJ0443$-$037     & $1.14\\pm0.07$ & 10.57 & $1.11\\pm0.07$ &  5.4 & 3\\\\ 0554$-$165 & 1RXSJ0557$-$1635    & $1.30^{+0.07}_{-0.12}$ & 11.86 & $1.28^{+0.09}_{-0.12}$ &  5.4 & 3\\\\ 0630$+$200 & 1RXSJ0633$+$2001    & $1.22^{+0.04}_{-0.05}$ & 11.11 & $1.20^{+0.05}_{-0.06}$ &  6.9 & 3\\\\ 0652$-$563 & EUVEJ0653$-$564     & $1.16^{+0.05}_{-0.06}$ & 11.40 & $1.13\\pm0.06$ & 10.0 & 3 \\\\ 0821$-$252 & 1RXSJ0823$-$2525    & $1.27\\pm0.05$ & 11.83 & $1.25^{+0.04}_{-0.05}$ & 12.2 & 3\\\\ 0914$-$195 & EUVEJ0916$-$197     & $1.27^{+0.06}_{-0.07}$ & 11.47 & $1.25^{+0.06}_{-0.08}$ &  6.8 & 3\\\\ 1022$-$301 & 1RXSJ1024$-$3021    & $1.21\\pm0.04$ & 11.63 & $1.19\\pm0.05$ & 13.3 & 3\\\\ 1529$-$772 & EUVEJ1535$-$774     & $1.26^{+0.03}_{-0.04}$ & 11.54 & $1.24^{+0.03}_{-0.04}$ & 10.6 & 3\\\\ 1658$+$441 & EUVEJ1659$+$440, PG & $1.34^{+0.06}_{-0.14}$ & 12.83 & $1.33^{+0.07}_{-0.16}$ & 43.5 & 3\\\\ 1711$+$668 & 1RXSJ1711$+$6645    & $1.22^{+0.03}_{-0.05}$ & 11.25 & $1.20^{+0.04}_{-0.05}$ &  6.8 & 3\\\\ 1724$-$359 & EUVEJ1727$-$360     & $1.20^{+0.02}_{-0.04}$ & 11.66 & $1.17\\pm0.03$ & 17.2 & 3\\\\ 1740$-$706 & EUVEJ1746$-$706     & $1.17\\pm0.04$ & 11.01 & $1.14\\pm0.04$ &  7.9 & 3\\\\ \\hline \\end{tabular}\\\\ $^a$ Using the mass-radius relations for a CO core \\citep{ben1999}. For masses above $1.2 M_\\odot$, mass-radius relations of \\citet{ham1961} were used. \\\\ $^b$ Using the mass-radius relations for a ONe core \\citep{alt2005}. For masses above $1.3 M_\\odot$, mass-radius relations of \\citet{ham1961} were used.\\\\ $^c$ Based on the model fit of Kawka \\& Vennes (2006) assuming $He/H = 0$. $^d$ Based on the distance modulus of the DA companion LB~9802 \\citep{kaw2007}. \\\\ References: (1) \\citet{kaw2006}; (2) \\citet{kaw2007}; (3) \\citet{ven2008}; (4) \\citet{ven2003}. \\end{minipage} \\end{table*}"
    },
    "0806/0806.4604_arXiv.txt": {
        "abstract": "We  derive stellar  masses,  ages and  star  formation histories  of massive early-type  galaxies in the  z=1.237 RDCS1252.9-2927 cluster and compare  them with those  measured in a  similarly mass-selected sample  of field contemporaries  drawn from  the GOODS  South Field. Robust  estimates of these  parameters are  obtained by  comparing a large  grid of  composite stellar  population models  with  8-9 band photometry in the rest-frame NUV,  optical and IR, thus sampling the entire  relevant  domain  of   emission  of  the  different  stellar populations.  Additionally, we present new, deep $U$-band photometry of both  fields, giving  access to the  critical FUV  rest-frame, in order to constrain empirically  the dependence on the environment of the  most recent  star  formation processes.   We  also analyze  the morphological properties  of both samples to  examine the dependence of their scaling  relations on their mass and  environment.  We find that early-type galaxies, both in the cluster and in the field, show analogous  optical morphologies,  follow comparable  mass  vs.  size relation, have congruent average  surface stellar mass densities and lie  on the  same Kormendy  relation.  We  also that  a  fraction of early-type galaxies  in the field employ  longer timescales, $\\tau$, to assemble  their mass than their cluster  contemporaries. Hence we conclude that, while the  formation epoch of early-type only depends on their mass, the environment does regulate the timescales of their star  formation  histories.   Our  deep  $U$-band  imaging  strongly supports  this conclusions.  I  shows that  cluster galaxies  are at least  0.5 mag fainter  than their  field contemporaries  of similar mass and optical-to-infrared colors,  implying that the last episode of star formation must have happened more recently in the field than in the cluster. ",
        "introduction": "The  description of galaxy  formation and  evolution becomes  far more complicated if one  considers that a galaxy is  not an isolated island universe. % Many  observations have  shown  that  galaxies are  in  fact parts  of groups, clusters  and super-clusters,  and that their  properties are correlated with  the environment in  which they live.  Several authors have shown  that the  local and the  large-scale environments  play an important  role in  determining many  galaxy properties,  such  as star formation rate, gas  content and morphology \\citep{Kodama01,Balogh02}. Several mechanisms have been proposed by theorists to account for these effects,  such as ram  pressure stripping,  mergers and  tidal effects \\citep{Gunn72, Dressler97, Moore96,Moore98,Moore99}.  More than half of all stars in the local Universe are found in massive spheroids (e.g., \\citet{Bell03}).  To derive the assembly  history of the bulk of the stellar mass in the Universe it is thus fundamental to understand how and when the  early-type  galaxies  (hereafter,  ETGs)  built  up  their  mass. Several  studies  at $1<  z  < 1.5$  indicate  that  the most  massive galaxies in the  field ($M_{stars} > 10^{11} M\\odot$)  may also be the oldest             at             a            given             epoch \\citep{Cimatti04,Fontana04,Saracco04,Treu05,Juneau05,Tanaka05}.  Other studies have  shown that the massive early-type  cluster galaxies have evolved  mainly passively  since $z  \\sim 1.0$  and that,  since then, field   galaxies  have   evolved   as  slowly   as  cluster   galaxies \\citep{Bernardi98,    vandokkum96,   vandokkum01,    Treu99,   Treu01, Kochanek00, vandokkumstanford03,  Strazzullo06, Depropris07}.  These studies  imply an  epoch  of massive  early-type  galaxy formation  at redshifts well  beyond $z  = 2$.  At  earlier epochs, an  abundance of dusty  and star forming  galaxies are  found, which  appear to  be the progenitors of massive ETG \\citep{Adelberger05}.  Contrary to what is observed   in  low   redshift   clusters,  observations   of  $z   >2$ proto-clusters  have  shown that  the  quiescent  red sequence,  which traces  the   passively  evolving  ellipticals,  has   yet  to  appear \\citep{Kurk04}.  This supports a paradigm where a more rapid evolution in denser environments is occurring a $z \\sim 2$.  More recent studies \\citep{Kodama07, Zirm08} reveal the  appearance of the red sequence of galaxies by $2< z <3$, although with a large scatter.  Thus, galaxy cluster samples at $1.0 <  z < 2.0$ provide a key link to the more  active epoch at  $z>2$ where proto-clusters  around powerful high  redshift  radio galaxies  are  not  yet  populated by  passively evolving ETG  \\citep{Kurk04}. Cluster galaxies have  evolved and were more luminous and bluer  at high redshift (e.g.  \\citet{vandokkum98}). Other observational results, e.g. \\citet{Labbe05}, \\citet{Papovich06}, \\citet{Kriek06}, suggest that  by $z \\sim 2$ we  are entering the star formation  epoch  of massive  ETGs,  leading  us  to ask  whether  the properties of  ETG at intermediate redshift are  consistent with this interpretation.  On the basis of  these studies, the evidence seems not only to indicate an epoch of massive early-type galaxy formation at $z >  2$ but  also that  in the  range $1.0  < z  < 2$,  the environmental effects start to become more substantial.  In  fact, the most  distant clusters  known to  date (all  at $z<1.5$) provide   the   strongest   test   of  model   predictions.   Relevant investigations include the observation of the aftermath of an off-axis merger in  XMMXCS2215.9-1738 \\citep{Hilton07}  at z=1.45, a  tight red sequence  of  ETG  at   $0.8<  z  <  1.3$  \\citep{Rosati04,  Blake03, Blakeslee06, Lidman04,  Mei06a, Mei06b}, a  slowly evolving $K$-band luminosity  function  at  odds  with  hierarchical  merging  scenarios \\citep{Toft04,  Strazzullo06},   and  a  tight   and  slowly  evolving Fundamental  Plane (hereafter,  FP) out  to z=1.25  \\citep{Hol05} have been found.  Intriguingly,  \\citet{Steidel05} have found that galaxies in a proto-cluster environment at $z=2.3$ have mean stellar masses and inferred ages that are $\\sim 2$ times larger than identically selected galaxies outside of the structure.  A long-standing prediction of hierarchical  models is that ETG in the field are younger for a given  mass than those in cluster cores, since galaxy    formation    is    accelerated   in    dense    environments \\citep{Diaferio01, Delucia04}. Studies at low redshift, using chemical abundance  indicators  \\citep{Bernardi06} or  the  analysis of  fossil record  data  via line  strength  indices \\citep{Thomas05,  Clemens06, Sanchez06} suggest  that star formation in  low density environments was delayed  by 1-2 Gyr.  FP studies at  $z \\simeq 1$ have  shown that massive  ETG in the field  and in  clusters  ($M_{*} >  3 \\cdot  10^{10} M_{\\odot}$)  share the  same FP  evolution ($M/L$  vs.  $z$)  and have approximately similar ages (within  $\\sim 0.4$ Gyr) and star formation histories (e.g., \\citet{vandokkum07}, \\citet{dSA06}).  FP studies have also shown that $M/L$ ratios  of massive cluster and field ETG evolve slowly and regularly and that there is evidence that low-mass galaxies evolve       faster      than      high-mass       galaxies      (e.g, \\citet{Hol05,Jorgensen06,Treu05,vdW05,dSA05}).   This  so-called  {\\it downsizing}  effect  is at  odds  with  earlier semi-analytic  model predictions  \\citep{Baugh96,   Kauffmann98,  Somerville04}  (see  also \\citet{Renzini06}), although  it can be reconciled in  the most recent versions  based  on  $\\Lambda  CDM$  cosmogony  \\citep{Delucia06}  and assuming that dry-mergers between  non star-forming ETG may occur and build up  the massive early-type  galaxy population \\citep{Khochfar03, Bell05}.  The majority of the above mentioned studies have focused on rest-frame optical and/or infrared spectrophotometric data.  However, the optical spectrum  remains largely unaffected  by moderate  amounts of  past or recent  star formation.   More  recently, deep  optical surveys  (e.g. FIRES \\citep{Franx03}, GOODS \\citep{Giava04}, COMBO-17 \\citep{Wolf04}, MUSYC  \\citep{Gawiser06}) have  provided access  to the  rest-frame UV spectrum beyond $z\\sim 0.5$, enabling more in-depth studies of the ETG population.   Studies of  local  ETG have  revealed  the existence  of relatively young stellar populations.  Such fossil record observations of  absorption  line-strengths  \\citep{Trager00, Thomas04}  find  that stellar  populations younger than  $\\sim 5$Gyr  (i.e. which  must have formed between  $z \\sim  1$ and the  present-day) are common  in ETGs. Furthermore,  a  significant  fraction  of  z  $\\leq  0.1$  ETG  show relatively blue NUV-optical  colors \\citep{Kaviraj05}, within extended disks  \\citep{Kauffmann07}, indicating  star-formation  over the  past Gyr.  The inferred recent  star-formation amounts to only $\\sim1$\\% of the total stellar mass.  In a more recent study, \\citet{Kaviraj07} have also shown  that a significant  fraction of $0.5  < z < 1.0$  ETG show relatively blue NUV colors,  indicating star-formation over the past 1 Gyr.  At  slightly higher redshift,  spectroscopic studies at  $z \\sim 1.2$  have  shown  that  both  the brightest  ETG  of  RDCS1252.9-2927 \\citep{Demarco07}   and   the  massive   ETG   in   the  GDDS   fields \\citep{Leborgne06}  show  evidence  for  recent  (i.e.,  within  1Gyr) star-formation on  the basis  of prominent post-starburst  features in galaxy spectra (e.g., $H_{\\delta}$ absorption line).  However, since little is known about the dependence on the environment of  the recent  star-formation rates  (hereafter, SFR)  in ETG  at $z \\simeq  1.2$, in  this  work,  we complement  our  modeling of  galaxy Spectral  Energy Distributions  (hereafter,  SEDs), from  NUV to  NIR, presenting result  of deep observations  of ETG in cluster  and field obtained with  {\\it VLT}/VIMOS in  the $U$-band filter,  that directly probes the  FUV regime  at the  redshift of our  samples, as  shown in Fig~\\ref{sed_model}.  By studying stellar population ages  at $z=1.2$, we provide a key test of the  paradigm of  an accelerated evolution  in the  highest density environments. For galaxies observed at $z \\leq 1$ most of their difference could have been smoothed  out  by  billion  years  of mostly  passive  evolution.   By comparing stellar masses, ages  and inferred star formation histories of cluster  ellipticals with  their field contemporaries,  we directly test  the  prediction  that  the cluster  environment  should  display accelerated evolution, resulting in larger masses and ages.  The primary observational goal of this  work is to use {\\it HST}/ACS in     the     rest-frame     Near-UV     (hereafter,     NUV)     and optical\\footnote{wavelength  ranges  which  are  known to  be  stellar population age sensitive},and  {\\it VLT}/ISAAC, {\\it Spitzer}/IRAC in the  rest-frame near-IR  (hereafter, NIR)\\footnote{a  wavelength range which is known to be strongly correlated with the underlying stellar mass  \\citep{Gavazzi96}}  to  measure  stellar population  ages  and masses   for    ETG   in   the    z=1.237   RDCS1252.9-2927   cluster \\citep{Rosati04}  and  compare them  to  those  measured on  similarly selected  sample of field  contemporaries drawn  from the  GOODS South Field.   This allows  us  to directly  analyze  the  entire relevant spectral energy distribution  of the different stellar populations, enabling us to improve constraints on galaxy ages, masses and star formation histories in both environments at $z \\simeq 1.2$.  We note that in an accompanying paper, \\citet{Gobat08}, we also compare the  coadded spectroscopic  information available  on both  samples of ETG with a large grid of composite stellar population models.  We also analyze  the morphological properties of ETG  in the field and in  the  cluster.  By  studying  scaling  relations  in this  relevant redshift range we can trace back where the majority of stars formed as a function of the environment and stellar mass.\\\\  The structure  of this  paper is as  follows.  The description  of our data-sets, cataloging and sample selection is described in \\S 2. In \\S 3 we describe our methods in  deriving galaxy sizes and morphologies as well as  inferring ages and  masses from stellar  population analysis. The  results  of  our  study are  discussed in  \\S  4,  while  in  \\S \\hspace{0.05cm} 5 we summarize our conclusions.  We assume a $\\Omega_{\\Lambda} = 0.73$, $\\Omega_{m} = 0.27$ and $H_{0} = 71\\ \\rm{km} \\cdot \\rm{s}^{-1} \\cdot \\rm{Mpc}^{-1}$ flat universe \\citep{Spergel03}, and use magnitudes in the AB system throughout this work. ",
        "conclusions": "\\subsection{The dependence of ETG scaling relations on environment}  The FP  is known to be a  powerful tool for studying  the evolution of ETGs \\citep{Djorgo87}.  In as similar  way to the small scatter of the color-magnitude  relation  \\citep{Bower92}, the  tightness  of the  FP \\citep{Jorgensen96,  Bernardi03}  constrains  the homogeneity  of  the ETG  stellar population.   Because of its  dependence on galaxy  luminosity,  the FP  is  sensitive  to  recent star  formation episodes.  One of its projections shows  a tight relation between  the effective radius, $R_{e}$, and the  mean surface brightness $<\\mu_{e}>$ measured inside  $R_{e}$, also  known  as  the \\citet[][hereafter  KR]{kor77} relation: \\begin{equation} <\\mu_{e}> = \\alpha + \\beta log (R_{e}) , \\end{equation} where  the slope  $\\beta  \\simeq 3$  is  found to  be  costant out  to $z\\simeq  0.65$  \\citep{LaBarbera03},  while  the  value  of  $\\alpha$ depends  on the  photometric  band adopted  to  derive the  structural parameters.   Here  we   adopt  the  KR  as  one   of  the  tools  for investigating  the  structural  properties   of  ETG  with  the  aim  of understanding the  role of the  environment in shaping ETG  of similar masses and optical-to-infrared colors.  In Fig.~\\ref{Kormendy}, we find very  similar KRs for the two samples. Both  the derived  zero points  and slopes  are consistent  within the errors.   These relations  show that  at the  effective  radius, large (massive) galaxies  are fainter than small galaxies  regardless of the environment.   This in  turn indicates  that large  galaxies  are less dense than small  galaxies in both the cluster and  field at $z \\simeq 1.2$.  For comparison,  we overplot  the  KR at  $z \\sim  0$ found  by \\citet{LaBarbera03}  (dotted-dashed  red  line),  K-corrected  to  our rest-frame $B$-band.  Our galaxies are  brighter by 1-2mag than at low redshift, a discrepancy that other studies at $1.0 \\lesssim z \\lesssim 1.4$ have  also found  difficult to explain  with the  assumption that galaxies undergo only a pure luminosity evolution with redshift (e.g., \\citet{Longhetti07}).   In fact,  our  galaxies show  an evolution  of $<\\mu_{e}>$ which exceeds $\\sim 2$  times the one expected in the case of   pure   luminosity   evolution   ($\\sim  1$mag).    According   to Eq.~\\ref{eq:mu_e},  the other  quantity affecting  $<\\mu_{e}>$  is the effective radius.  Therefore to recover this discrepancy we can assume that, as a function of redshift, ETG undergo a size evolution as well: the effective radius  of ETG should increase by  at least factor $\\sim 1.5$ from  z$\\simeq 1.2$ to z$\\sim 0$  both in the cluster  and in the field environment.  Recent  studies of  the  dependence  on environment  of  the size  vs. stellar  mass   relation  \\citep{Trujillo04,  McIntosh05,  Trujillo06, Trujillo07} found that the  bulk of galaxies with comparable stellar masses to  ours were at least  a factor 2 smaller  at higher redshifts than  locally.  This  is  qualitatively consistent  with the  observed trend  in our  data,  (see Fig.~\\ref{mass_size})  when  the sizes  and masses our  samples are  compared with the  local relation for  ETG in SDSS \\citep[dotted-dashed red line;][]{Shen03}.  We find no-dependence on the environment of the $R_{e}$ vs.  $M_{*}$ relation, implying that cluster  and  field ETGs  must  undergo  similar  luminosity and  size evolution  to match  the typical  values found  for the  ETG  at lower redshifts.  To  explain how compact  galaxies observed at  $z>1$ could possibly  end-up  on  the  local  relation,  a  possible  evolutionary mechanism  that grows  stellar mass  and  size has  been suggested:  a dissipationless   (``dry'')  merging   of   gas-poor  systems   (e.g., \\citep{Ciotti01, Nipoti03, Khochfar03,  Boylan06} that is efficient in increasing  the size of  the objects,  while remaining  inefficient at forming new stars.  In   the  local  universe,   the  SFR   per  stellar   mass  (specific star-formation   rate,   SSFR)   correlates   strongly   with   galaxy concentration, effective  radius and the average  surface stellar mass density \\citep[$\\sigma_{50}$;][]{Kauffmann03,Brinchmann04}. A striking similarity of  cluster and  field galaxies at  $z \\simeq$1.2  is again shown in  Fig.~\\ref{andrew} where we plot, $\\sigma_{50}$, \\begin{equation} \\sigma_{50} = \\frac{0.5 M_{\\star}}{\\pi R_{e}^{2}} , \\end{equation} versus the  stellar mass,  and compare them  with similar data  in the literature  drawn  from \\citet{Zirm07}  at  $z  \\sim 2.5$.   Quiescent Distant Red  Galaxies (qDRGs) are drawn as  red ellipses, star-forming DRGs (sDRGs) are  drawn as open red stars,  while blue stars indicates Lyman Break  Galaxies (LBGs)  from the same  work.  While some  of our galaxies are almost as  dense as the \\citet{Zirm07} and \\citep{Toft07} quiescent  distant  red galaxies  (qDRGs),  both  our samples  (filled circles) overlap the region occupied by other $1.0 \\lesssim z \\lesssim 1.5$  galaxy  samples  \\citep{Trujillo06, Daddi05,  vdW06,  Rettura06} (open red  squares, open red  circles, open black circles,  open black squares, respectively).   As a comparison  we also overplot  the local relation for ETG in SDSS (red dotted-dashed line) calculated from the mass-size   relation  of   \\citet{Shen03}.   It   is  very   clear  in Fig.~\\ref{andrew} that  the bulk of  our galaxies in both  samples have much larger  densities that their local counterparts.   To account for this  effect  in  the  context  of  a  plausible  formation  scenario, semi-analytical modeling  (e.g. \\citet{Khochfar06}) suggests  that ETGs formed in gas-rich mergers can  result in very dense stellar cores, as the gas is  driven to the center of  this ``wet'' (dissipative) merger where it  very efficiently produces massive  starburts.  Galaxies that merge in the  early universe are likely to  be gas-rich.  Consequently the dense  nature of this objects  could be the result  of much denser conditions of  the universe  at the time  of their formation.  We note that  our   finding that there is no dependence  on  the   environment  of  the $\\sigma_{50}$  vs,  $M_{*}$  relation  at z$\\simeq  1.2$  can  provide an important datum for models of galaxy formation.   \\subsection{The dependence of ETG ages and star formation histories on their environment}  As we  apply the method  described in section  \\S 3.2, we are  able to directly compare  the distribution of star-formation  weighted ages in the  field  and  in  the  cluster.   As shown  in  the  top  panel  of Fig.~\\ref{timing},  we  find  the  overall  relative  distribution  of cluster and field ETG ages to be very similar.  This result implies no significant delay  in relative  formation epochs is  found for  ETG in either environments.   We find that  $\\sim 80$\\% of massive  ETGs have ages in the range $3.5 \\pm 1.0$ Gyr in both cluster and field.  To investigate  the  dependence of  this  result on  the actual  stellar  population synthesis  code  adopted,  we compare  our current results (based on BC03 models) with those obtained with a set of  similar  dust-free CSPs  templates  built with  \\citet{Maraston05} models,   adopting   the    same   exponentially-declining   SFHs   of Eq.~\\ref{eq:sfhs},     and    assuming    solar     metallicity    and \\citet{Salpeter55}  IMF.  The  result  of the  analysis  based on  M05 models  is shown  in the  bottom panel  of Fig.~\\ref{timing}  where we still find the cluster and field relative age distributions to be very similar, despite the fact that the contribution of the TP-AGB stars in these models  are implemented  in a different  way.  However,  here we find that $\\sim 60$\\% of galaxies have ages in the range $3.5 \\pm 1.0$ Gyr; M05 models favor slightly younger ages ($\\sim1-2$ Gyrs) for $\\sim $20\\% of ETG  in both environments.  This effect  can be explained by the fact that, at about 1  Gyr, M05 models account for a larger amount of  2 $\\mu m$  flux than  BC03 models  of the  same age,  resulting in significantly redder  color at  younger ages, thus  can favor  $\\overline t \\sim$ 1-2Gyrs best-fits in a few cases.   To  summarize,  we  find   that,  regardless  of  the  actual  stellar population synthesis code adopted, cluster galaxies ages have the same relative  distribution as their  field contemporaries:  no significant delay in  their formation  epochs is found,  within the  errors ($\\sim 0.5Gyr$).   This result  is  at  variance with  some  versions of  the hierarchical    model    of    galaxy    formation    and    evolution \\citep{Diaferio01,   Delucia06}  and   with   fossil  record   studies \\citep{Thomas05, Clemens06}, but in remarkably good agreement with the ones  derived  by   \\citet{vandokkum07}  and  \\citet{dSA06}  from  the evolution of the $M/L$ ratio.  It should be noted that similar results are found by other works using independent methods and data-sets.  In the top-left panel of Fig.~\\ref{quartetto} we plot for both samples each galaxy  age, $\\overline  t$, as a  function of stellar  mass.  We note  that  the  age  of   ETG  increases  with  galaxy  mass  in  all environments,  which   is  in  agreement  with   the  so--called  {\\it downsizing}  scenario  of  galaxy formation  \\citep{Cowie96}.   This effect   can  also   be   seen   in  the   the   top-right  panel   of Fig.~\\ref{quartetto},  where we  plot  our galaxies'  lookback time  to formation as  a function of  their stellar mass, in  both environments. Our result  is in  agreement with the  one obtained with  an independent method  and data-set by \\citep{dSA06} and  based on z$\\sim 1$ ETG ages estimated from the FP parameters (see their Fig. 3).  Despite of the fact that cluster and field galaxy formation epochs are found to be similar, it could still be possible that the timescales of their SFH  are significantly different.   Firstly, the data  show that the distribution of cluster  and field optical colors is significantly different. As  a function  of the stellar  mass, cluster  galaxies are found to  lie on a very  tight red-sequence, while those  in the field populate     the    color-sequence     with    a     larger    scatter (Fig.~\\ref{UMB_mass}). Secondly,  in \\citet{Gobat08} we  find that the averaged  spectrum  of the  cluster  galaxies  has  a more  pronounced 4000\\AA \\mbox{} break than that of the field sample.  Both these pieces of observational evidence find a natural explanation in the framework of our  modeling.  As shown in the bottom-right panel of Fig.~\\ref{quartetto}, as  a function of stellar mass,  we find that field  ETG  span a  larger  range  of  timescales than  their  cluster contemporaries, which are formed with the shortest $\\tau$ at any given mass.   According  to  our  models,  cluster ETG  are  found  to  have experienced more  similar star-formation  histories.  As shown  in the bottom-left  panel   of  Fig.~\\ref{quartetto},  cluster   ETG  form  a color-age sequence with much smaller scatter than the field ones.  As  discussed in  \\S 2.1,  we  recall that  our field  sample is  more deficient in lower mass objects than the cluster sample because of the different depths of  spectroscopic observations. However, even if  the field  sample were  corrected for  completeness,  this would likely result in a larger fraction of field ETG at low mass, which are  the ones that we found with  longer $\\tau$. Hence this would amplify the  difference between the typical timescales of the two samples, and so not affect our conclusions.  \\subsection{The dependence of ETG FUV magnitudes on their environment}  The rest-frame FUV ($\\sim \\lambda ~  1700 \\AA$) SED is a crucial range where hot ($> 9000 K$),  massive ($M > 2 M_{\\odot}$), short-living ($< 1Gyr$) stars emit most of their light.  Thus it is a wavelength domain which is very sesnsitive to current or recent star formation.  Most of the light from  ETG is emitted in the optical  and the NIR rest-frame. However, the FUV can be used  as a good tracer of the residual current star-formation  and to  trace  back,  within the  last  Gyr, the  most recent episode of star-formation. About 100Myr after star formation ceases,  an ETG  spectrum  becomes dimmer  and  redder. Therefore,  the fainter the  rest-frame UV emission  is, the earlier  the star  formation must  have stopped.   Over time,  the  galaxy spectrum fades  and  slowly reddens  as  the  4000  $\\AA$ break  becomes  more pronounced.   Here we have used the  {\\it  VLT}/VIMOS  $U$-band observations  described  in  \\S   2.2  to  empirically  constrain  the dependence  on  the environment  of  the  most  recent star  formation processes in z$\\simeq 1.2$ ETGs.  However,  when analyzing  UV rest-frame  fluxes of  massive ETG  it is important to recall  that core helium burning stars  on the horizontal branch (HB) are known to  produce a ``UV upturn'' feature \\citep{Yi97, Yi99}.    This   effect    can,   in   principle,   complicate   the disentanglement of the contributions to the UV spectrum of the evolved and young stellar populations.  However, the onset of the HB typically takes 9  Gyr, meaning that, by  z=1.2 (when the Universe  is only 5Gyr old), not enough cosmic time  has elapsed for this population of stars to appear.   Hence the UV flux  seen in our sample  ETG must originate only from young stars.  In Fig.~\\ref{uband} we show  $U$-band magnitudes (1'' radius aperture; rest-frame $\\sim  1700 \\AA $)  as a function  of stellar mass  for the ETG  detected  in  the  field  (filled blue  circles).   Solid  lines represent the  1$\\sigma$ limit magnitudes  of both data-sets  (in blue for  CDFS,  in  red  CL1252).  Dashed lines  represent  the  3$\\sigma$ limits.  As already  pointed  out  in \\S2.2,  the  combined effect  of shorter  total exposure times  and higher  galactic extinction  at the location  of CL1252,  directly translates  into a  $\\sim  1mag$ deeper $U$-band photometry  for the  CDFS.  A  large fraction (75  \\%) of  field ETG  are ($>3 \\sigma$) detected  in the  deep CDFS  images.   The observed  magnitude of  the median stack of these detections is $U$=27.46 mag (blue dotted line of Fig.~\\ref{uband}),   corresponding   to   a  SFR=0.47   $M_{\\odot}/yr$ \\citep{Sawicki06}.  An image of the  median $U$-band stack of the ETGs detected  in the  field, is  displayed  in the  bottom-left corner  of Fig.~\\ref{uband}.  Since none of the CL1252 ETG is actually detected in our $U$-band data, we use their median stack, shown in the middle of Fig.~\\ref{uband}, to provide a robust  upper limit of $U>27.3$ mag  for the CL1252 early-type population, which corresponds to SFR$< 0.55 M_{\\odot}/yr$.  We  note that  the non  detection  of the  CL1252 ETG  cannot only  be attributed to the  shallower $U$-band data for the  cluster.  To prove this we have  simulated how the CDFS detected ETG  would appear in the 1252  data  (more  details  can  be found  in  \\citet{Nonino08}).   We randomly placed the  20 $U$-band detected CDFS ETG  (0.4 mag dimmed to match the  relative difference in  galactic extinction) in  the CL1252 maximally exposed  region, avoiding objects detected  in the $U$-band, and repeated this  step 30 times. Hence we  generated 32 median stacks of  18 simulated  galaxies  each,  picking up  at  random amongst  the cutouts.   Aperture photometry  (1'' radius)  on these  simulated {\\it CDFS@1252} stacks results in a median value of $U$=27.8 mag which is in agreeement with the dimmed inputs of the simulations. In the  bottom-right corner of  Fig.~\\ref{uband} we show one  of these stacks, which would be clearly  detected in 1252.  Comparing this last value to the upper limit we measured in the cluster data, we can state with confidence that cluster ETG are intrinsically fainter by at least 0.5 mag  in the observed  $U$-band than their field  contemporaries of similar mass and optical-to-infrared colors.  This observational  evidence corroborates  the results of  our stellar population  synthesis analysis described  in the  previous subsection. In our proposed  scenario, a generally shorter {\\it  timescale} of the star  formation  process among  the  cluster  galaxy population  would naturally result  in a generally fainter  observed $U$-band magnitudes compared to the field population at z$\\simeq 1.2$."
    },
    "0806/0806.4972_arXiv.txt": {
        "abstract": "We study the gravitationally confined detonation (GCD) model of Type Ia supernovae through the detonation phase and into homologous expansion.  In the GCD model, a detonation is triggered by the surface flow due to single point, off-center flame ignition in carbon-oxygen white dwarfs. The simulations are unique in terms of the degree to which non-idealized physics is used to treat the reactive flow, including weak reaction rates and a time dependent treatment of material in nuclear statistical equilibrium (NSE). Careful attention is paid to accurately calculating the final composition of material which is burned to NSE and frozen out in the rapid expansion following the passage of a detonation wave over the high density core of the white dwarf; and an efficient method for nucleosynthesis post-processing is developed which obviates the need for costly network calculations along tracer particle thermodynamic trajectories.  Observational diagnostics are presented for the explosion models, including abundance stratifications and integrated yields. We find that for all of the ignition conditions studied here, a self regulating process comprised of neutronization and stellar expansion results in final \\iso{Ni}{56} masses of $\\sim$1.1\\msun. But, more energetic models result in larger total NSE and stable Fe peak yields. The total yield of intermediate mass elements is $\\sim0.1$\\msun and the explosion energies are all around 1.5$\\times$10$^{51}$ ergs. The explosion models are briefly compared to the inferred properties of recent Type Ia supernova observations. The potential for surface detonation models to produce lower luminosity (lower \\iso{Ni}{56} mass) supernovae is discussed. ",
        "introduction": "\\par The currently favored model for Type Ia supernovae (SNe Ia) is the thermonuclear incineration of a white dwarf (WD) which has accreted mass to near the Chandrasekhar limit from a binary companion \\citep[e.g.,][]{branch1995,hillebrandt2000}. The enormous luminosity and homogeneity in the properties of the light curves of SNe Ia make them exceptionally good standard candles and as such have shown that the expansion rate of the universe is accelerating and provided intriguing evidence for a cosmological constant \\citep{riess1998}.  \\par Despite the success in using SNe Ia as cosmological probes and identifying a plausible astrophysical progenitor site for the explosions, a detailed understanding of the explosion mechanisms itself remains elusive. Several uncertainties stand in the way of a definitive solution to the SNe Ia problem. On the one hand, the conditions under which the thermonuclear runaway commences remains poorly understood so that the initial number and distribution of flamelets that seed the runaway is still a free parameter. On the other hand, although significant progress has been made in simulating flame fronts in multi-dimensional stellar models \\citep{gamezo2005,schmidt2006b,roepke2007a,townsley2007,jordan2007}, the challenge associated with modeling an unresolved turbulent deflagration \\citep[e.g.][]{schmidt2006a} with limited computational resources injects an additional degree of uncertainty into the outcome of a model for any given choice of initial conditions.        \\par In this paper we describe progress on our ongoing effort to improve the simulation of SNe Ia in multi-dimensions, including methods to perform detailed nucleosynthesis post-processing in a computationally efficient manner. We extend the study of the GCD model for a single ignition point slightly offset a range of distances from the center of the star, as described in \\citet{townsley2007}, through the detonation phase and into homologous expansion. In \\S\\ref{sec:numerics} we describe the treatment of the reactive-hydrodynamics problem used in our simulation code. In \\S\\ref{sec:def} we review the relevant properties of the deflagration phase for single point flame ignition.  In \\S\\ref{sec:det} we examine in some detail the initiation of the detonation, the properties of the detonation wave which disrupts the star, and the resultant remnant morphology. In \\S\\ref{sec:yields} we discuss in detail the nucleosynthetic yields for the explosions studied and decribe the methodology used to efficiently calculate iron peak yields from the simulation data. We conclude with a summary of the salient features of the explosion models in light of observed Type Ia supernvoae. ",
        "conclusions": "\\par We have studied the final outcomes for a range of single point flame ignition models of thermonuclear supernovae within the computational framework developed at the FLASH center (\\S\\ref{sec:numerics} and \\citet{fryxell2000,calder2007,townsley2007,seitenzahl2008a}). For the first time in this work, we have extended the 3-stage reactive ash model for nuclear burning described in \\citet{townsley2007} to study the ignition and propagation of the detonation mode of burning. As a result, our explosion models are unique in terms of the degree to which non-idealized nuclear physics are employed, including a non-static, tabularized treatment of the nuclear statistical equilibrium (NSE) state and the inclusion of contemporary weak reaction rates \\citep{seitenzahl2008a}. In addition, we have demonstrated here, by reaction network post-processing of recorded Lagrangian histories, that the 3-stage reactive ash model provides a suitable reproduction of fluid density and temperature histories to allow detailed nucleosynthesis, including self-consistent neutronization. Using these techniques, we have followed the progression of a thermonuclear flame (deflagration) from a single ignition point which is varied to successively larger distances from the center of a carbon oxygen white dwarf, and we have described in detail the resulting surface flows and detonation which ensue.  \\par Detonations arise within a colliding surface flow for all models which are ignited at a radial location which exceeds $\\sim$ 20 km in our 2D simulations. Flames ignited closer to the stellar center release enough nuclear energy to signifianctly expand the stellar core to a degree that it stalls the surface flow, thus preventing a strong collision region and detonation.  The nuclear binding energy released in these stalled surface flow models, however, is not enough to gravitationally unbind the star and they remain viable candidates for a pulsational detonation upon recollapse \\citep{khoklov1991,arnett1994}. Models which detonate release $\\sim 2\\times10^{51}$ erg in nuclear binding energy, resulting in a supernova-like explosion with total energy $E_{\\rm tot}\\sim1.5\\times 10^{51}$ erg.  \\par  In all of the models in our parameter study which produce supernova-like explosions, detonation initiates within a jet-like flow which forms in the converging surface flow. This is in agreement with the results presented in \\citet{kasen2007}. However, we do not find that the detonation initiates through a shock to detonation transition (SDT) as suggested by these authors, but instead find that the detonation occurs through a gradient mechanism. The initiation of the detonation takes place within the compressed gas which lies ahead of the high velocity jet, and ahead of the internal shock which forms within the jet (see \\S\\ref{sec:det} and Figure~\\ref{fig:jet-slice}). The focusing of the surface flow and the formation of the jet is also present in 3D simulations (\\S\\ref{sec:det}, and \\citealt{jordan2007,roepke2007a}), and is therefore not an artifact of the 2D axisymmetric geometry used.  \\par Within a few seconds after the detonation wave disrupts the stellar core, homologous expansion is beginning to be established. By t$\\sim$ 4 s from flame ignition more than 90\\% of the total energy is in the kinetic energy of expansion, and the expansion velocity has acquired a linear dependence on radius. The final remnant posesses both global and small scale asymmetries which will influence the observational signature. When the remnant enters the homologous expansion phase it is characterized by a smooth, layered inner core surrounded by a low density, flocculent layer of deflagration ash which was dumped onto the surface of the star prior to detonation. The smooth inner core of the remnant has a global north-south asymmetry due to off-center ignition and surface detonation, which is well characterized by circular isodensity contours which are progressively off-center at higher densities (see \\S~\\ref{subsec:remnant} and Figure~\\ref{fig:shape}).   \\par We have analyzed in detail the nucleosynthesis of material burned to NSE in the detonation. These results have been generated from the multi-dimensional simulation data using a newly developed post-processing method which takes advantage of the uniqueness of the NSE state and systematic properties of detonation waves. The method, presented in \\S\\ref{sec:yields}, obviates the need for computationally prohibitive network calculations along each of the millions of particle trajectories which are necessary for good mass resolution in 3D explosion models. This work addresses only material which has relaxed to NSE, which forms, by mass, nearly all of the yield from the 2-dimensional explosion models of this study. Extending the method to include detailed isotopic yields for material incompletely relaxed to NSE (incomplete carbon, oxygen, and silicon burning) is being developed and will be described in a forthcoming publication (C. Meakin, in prep). Nucleosynthesis of material processed in a deflagration instead of detonation burning mode can be processed with a similarly parameterized method, though requiring more parameters, if it reaches NSE.  This leaves only the relative minority of tracks in partially burned deflagration material to be processed directly (only a few percent of all the trajectories).  \\par Larger offsets of the ignition point lead to less stellar expansion prior to detonation and therefore the production of more NSE material.  However, we find that the amount of \\iso{Ni}{56} produced stays roughly fixed at $\\sim 1.1 M_\\odot$ for all of our 2-dimensional explosion models which extend down to a central density of $\\rho_c\\sim4\\times 10^{8}$ g/cm$^{3}$ at the time of detonation. This regulation is due to the enhanced neutronization at the higher densities characteristic of the less-expanded cases. Higher density cores at the time of detonation result in more neutronization, and therefore a larger fractional yield of stable Fe-peak isotopes (e.g., \\iso{Fe}{54} and \\iso{Fe}{58}). The isotopic distribution we find in the Fe-peak is very similar to that found for pure deflagration models, but is characterized by a lower degree of neutronization.  Less neutronization is a result of the lower densities under which the burning proceeds in our surface detonation models compared to pure deflagrations, due to the pre-detonation expansion. Between 0.06 and 0.14\\msun\\ of intermediate mass elements are produced at high velocities. Regions in which more than half of the mass is in the form of IMEs lie above an expansion velocity of 14,000 km/s for all of the 2-dimensional detonation models calculated.   \\par We successfully reproduced the relationship between the central density and mass-density distribution in the pre-detonation expanded star by superposing on the hydrostatic star the lowest order radial mode calculated in a linear approximation. We find that much smaller \\iso{Ni}{56} yields are expected in cores which undergo more expansion prior to detonation (see Figure~\\ref{fig:mnse-rhoc}).  This degree of expansion appears to be achievable in 3-dimensional simulations which relax the constraints on axisymmetry of the ignition conditions necessary for 2-dimensional simulations.  Thus it is expected that more realistic simulations, which include the pre-ignition convection field and its effect on the growing flame bubble, will be characterized by such larger expansions. However, further analysis of such simulations, which will be the subject of future papers, is required.   \\par Future work on elucidating the SNe Ia explosion mechanism which is being pursued at the FLASH center involves the following. (1) We are extending our survey of the mapping between flame ignition conditions and final outcomes within the computational framework developed at the FLASH center, including multi-point ignition conditions and 3D models. (2) A simulation pipeline is being constructed to generate synthetic observational diagnostics for the explosion models, including light curves and spectra, which will allow a more direct comparison between the systematic properties of the single degenerate Type Ia model and observational data."
    },
    "0806/0806.4233_arXiv.txt": {
        "abstract": "We present results from multifrequency polarimetry of NRAO 140 using the Very Long Baseline Array. These observations allow us to reveal the distributions of both the polarization position angle and the Faraday rotation measure (RM). These distributions are powerful tools to discern the projected and line-of-sight components of the magnetic field, respectively. We find a systematic gradient in the RM distribution, with its sign being opposite at either side of the jet with respect to the jet axis. The sign of the RM changes only with the direction of the magnetic field component along the line of sight, so this can be explained by the existence of helical magnetic components associated with the jet itself. We derive two constraints for the pitch angle of the helical magnetic field from the distributions of the RM and the projected magnetic field; the RM distribution indicates that the helical fields are tightly wound, while that of the projected magnetic field suggests they are loosely wound around the jet axis. This inconsistency may be explained if the Faraday rotator is not cospatial with the emitting region. Our results may point toward a physical picture in which an ultra-relativistic jet (``spine'') with a loosely wound helical magnetic field is surrounded by a sub-relativistic wind layer (``sheath'') with a tightly wound helical magnetic field. ",
        "introduction": "Jets from active galactic nuclei (AGNs) maintain well-collimated structure and can travel for more than 100 kpc. It has come to be known that the bulk Lorentz factors of these jets reach $\\Gamma =$ 30 in some cases (Kellermann et al. 2004). The mechanisms of formation, acceleration, and collimation of AGN jets are, however, still unclear. Magnetohydrodynamic (MHD) mechanisms are frequently invoked to model these jets. Magnetically driven outflows powered by a spinning black hole (Blandford \\& Znajek 1977) or an accretion disk (Blandford \\& Payne 1982) have been widely discussed in terms of both the acceleration and collimation. MHD outflows originating in the AGN core, particularly those with a strong toroidal field encircling the collimated flow, can exhibit efficient acceleration and collimation. This continues over a scale of $10^3 - 10^4$ Schwarzschild radii, and the terminal Lorentz factors of MHD outflows have reached values of 10-100 in recent numerical studies (Vlahakis \\& K\\\"onigl 2004; Fendt \\& Ouyed 2004; McKinney 2006). These theoretical pictures are qualitatively consistent with some VLBI observations (e.g., Junor et al. 1999; Horiuchi et al. 2006). In the theoretical models, the toroidal magnetic field plays an important role in the acceleration (through the magnetic pressure gradient force) and the collimation (through the magnetic tension force). Thus, the detection of helically twisted magnetic components associated with such jets is a crucial key to confirming the MHD models with observations.  It is important to reveal the structure and strength of the magnetic fields in jets, although the available observational methods are limited. One powerful option is polarimetry. In the optically thin synchrotron emission from a nonrelativistic plasma, the polarization position angle (P.A.) is perpendicular to the projected direction of the magnetic field (the component perpendicular to the line of sight). In addition, the fractional polarization is closely related to the relative degree of ordering in the magnetic field. For this purpose, very long baseline polarimetry (VLBP) has been applied since the technique's early development. It has been reported that the P.A.'s tend to be parallel with respect to the jet axis for BL Lacertae objects whereas they tend to be perpendicular for quasars (Cawthorne et al. 1993). Jorstad et al. (2007), using multifrequency polarization observations, found a similar dichotomy but also suggested that it is not simply based on the optical classification of an object. The trend in BL Lac objects is sometimes explained by the compression due to strong shocks in the jets (Laing 1980; Hughes et al. 1989). Another possible explanation is that it is due to a toroidal magnetic field (Gabuzda et al. 2000) or the toroidal component of a helical magnetic field (Asada et al. 2002). However, the situation is not this simple, because the observed direction of the P.A. is not generally orthogonal to the projected direction of the magnetic field in the case of a relativistically moving, optically thin jet (Lyutikov et al. 2005). Neither is it simple to discriminate between these possibilities, since we have not taken advantage of the information derivable from the line-of-sight component of the magnetic field.  Further possibilities for examining the structure of such magnetic fields have been provided by the capability of full polarimetry with the Very Long Baseline Array (VLBA). The line-of-sight component of the magnetic field can be probed by the distribution of the Faraday rotation measure (RM), since this is related to the electron density $n_{e}$ and the magnetic field component parallel to the line of sight ${\\it B}_{\\rm LOS}$ as RM $\\sim \\int_{\\rm LOS} n_{e} {\\it B}_{\\rm LOS} dr$, where $\\int_{\\rm LOS} dr$ represents integration along the line of sight. By combining the distributions of both the RM and the projected magnetic field, one can more reliably investigate the three-dimensional structure of the magnetic fields in parsec-scale jets.  The first inference of the presence of a helical magnetic field structure from the RM distribution was for the 3C 273 jet, based on 5-8 GHz VLBA observations (Asada et al. 2002). Those observations revealed a gradient of the RM across the jet, which can be interpreted as indicating the presence of the toroidal component of a helical field (Asada et al. 2002). In the simple case where we are seeing such a field from the side (a viewing angle of 90$^{\\circ}$), the sign of the line-of-sight component of the magnetic field will differ on the two sides of the jet as the field reverses direction. If the viewing angle decreases, the antisymmetric distribution will remain, with only an additional offset in the absolute value of the RM. Therefore, an RM gradient is expected in the presence of a helical magnetic field for arbitrary viewing angles, except 0$^{\\circ}$. After this initial work, a similar RM gradient in 3C 273 was confirmed by independent observations (Zavala \\&Taylor 2005; Attridge et al. 2005; Asada et al. 2008). The same kinds of gradients have been reported for several BL Lac objects as well (Gabuzda et al. 2004).  In this paper we report observations of a bright quasar NRAO 140, which show an RM gradient across its jet, implying the presence of helical magnetic components. NRAO 140 is a distant quasar with a redshift, $z$, of 1. 263. If we assume ${\\it H}_{0}$ = 73 km s$^{-1}$ Mpc$^{-1}$ and ${\\it q}_{0}$ = 0. 5, 1 mas corresponds to 8. 4 pc. A value of 1 mas yr$^{-1}$ corresponds to 27. 4 {\\it c}, and a superluminal motion of 11. 0 {\\it c} was reported with VLBA observations at 15 GHz (Kellermann et al. 2004).  The paper is organized as follows: In \\S 2, we describe the details of the VLBA polarimetry observations and calibrations. In  \\S 3,  we present polarization images of NRAO 140 jet and summarize several properties of the projected magnetic field and the RM. Discussion and conclusions are given in \\S\\S 4 and 5. ",
        "conclusions": "In order to discuss the three-dimensional configuration of the magnetic fields in parsec-scale quasar jets, we have performed multifrequency VLBA polarimetry toward NRAO 140. We revealed the distributions of both the projected component of the magnetic field and the rotation measure. We find a systematic gradient in the RM distribution similar to that seen in the 3C 273 jet. Furthermore, the sign of the RM is opposite on either side of the gradient. This is presumably due to a change of direction of the magnetic field and is naturally explained by a helical field structure. Using the properties of the RM gradient and the viewing angle, we derive a constraint for the pitch angle of the helical magnetic field in the layer or the sheath of the jet of $\\chi$ $<$ 10.4$^{\\circ}$. On the other hand, from the direction of the projected magnetic field and the viewing angle, we derive another constraint for the pitch angle in the spine of the jet of $\\chi$ $>$ 79.8$^{\\circ}$. Therefore, we expect that in NRAO 140 an ultra-relativistic jet (or ``spine'') with a loosely wound helical magnetic field is surrounded by a layer (or ``sheath'') with a tightly wound helical magnetic field.   {\\bf Acknowledgments:} The authors thank E. S. Perlman for a critical reading of this paper and valuable comments as a referee. The VLBA and VLA are operated by the National Radio Astronomy Observatory (NRAO), a facility of the National Science Foundation, operated under cooperative agreement by Associated Universities, Inc."
    },
    "0806/0806.3413_arXiv.txt": {
        "abstract": "Using data from the 2MASS All-Sky Point Source Catalogue, we have extended our census of nearby ultracool dwarfs to cover the full celestial sphere above Galactic latitute 15$^o$. Starting with an initial catalogue of 2,139,484 sources, we have winnowed the sample to 467 candidate late-type M or L dwarfs within 20 parsecs of the Sun. Fifty-four of those sources already have spectroscopic observations confirming them as late-type dwarfs. We present optical spectroscopy of 376 of the remaining 413 sources, and identify 44 as ultracool dwarfs with spectroscopic distances less than 20 parsecs. Twenty-five of the 37 sources that lack optical data have near-infrared spectroscopy. Combining the present sample with our previous results and data from the literature, we catalogue 94 L dwarf systems within 20 parsecs. We discuss the distribution of activity, as measured by H$\\alpha$ emission, in this volume-limited sample. We have coupled the present ultracool catalogue with data for stars in the northern 8-parsec sample and recent (incomplete) statistics for T dwarfs to provide a snapshot of the current 20-parsec census as a function of spectral type. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0806/0806.4005_arXiv.txt": {
        "abstract": "We present results from spatially resolved spectral analyses of the northeastern (NE) rim of the Cygnus Loop supernova remnant (SNR) based on two {\\it Chandra} observations.  One pointing includes northern outermost abundance-enhanced regions discovered by recent {\\it Suzaku} observations, while the other pointing is located on regions with ``normal'' abundances in the NE rim of the Cygnus Loop.  The superior spatial resolving power of {\\it Chandra} allows us to reveal that the abundance-enhanced region is concentrated in a $\\sim$200$^{\\prime\\prime}$-thickness region behind the shock front. We confirm absolute metal abundances (i.e., relative to H) as well as abundance ratios between metals are consistent with those of the solar values within a factor of $\\sim$2.  Also, we find that the emission measure in the region gradually decreases toward the shock front. These features are in contrast with those of the ejecta fragments around the Vela SNR, which leads us to believe that the abundance enhancements are not likely due to metal-rich ejecta.  We suggest that the origin of the plasma in this region is the interstellar medium (ISM).  In the ``normal'' abundance regions, we confirm that abundances are depleted to the solar values by a factor of $\\sim$5 that is not expected in the ISM around the Cygnus Loop.  Introduction of non-thermal emission in our model fitting can not naturally resolve the abundance-depletion problem. The origin of the depletion still remains as an open question. ",
        "introduction": "The Cygnus Loop is one of the nearest (540 pc: Blair et al.\\ 2005) supernova remnants (SNRs). The angular dimensions are $2^\\circ.8\\times3^\\circ.5$ (Leahy et al.\\ 1997) and the age is considered to be $\\sim$10000 yrs.  The low column density of the foreground material makes it an ideal target to study UV and soft X-ray emission from the remnant. The X-ray boundary in the northeastern (NE) rim of the Cygnus Loop is associated with Balmer-dominated filaments which mark current locations of the blast wave (Chevalier 1978).  Balmer-dominated filaments in this region have been well studied by optical as well as UV observations (e.g., Raymond et al.\\ 1983; Sankrit et al.\\ 2000; Ghavamian et al.\\ 2001).  Hester et al.\\ (1994) studied a Balmer-dominated filament in detail and reported that it was recently (in past 1000 yrs) decelerated from $\\sim$400 km\\,sec$^{-1}$ to $\\sim$180 km\\,sec$^{-1}$.  The rapid deceleration of the shock velocity was considered to be a result of the blast wave hitting the wall of a cavity which surrounded the supernova precursor.  The cavity wall/blast wave interaction was also suggested by X-ray observations.  For example, Levenson et al.\\ (1999) revealed that a soft spatially thin ($<5'$) shell surrounded almost the entire rim of the Loop based on a {\\it ROSAT} PSPC hardness map.   They concluded that the soft shell occurred where the cavity wall decelerated the blast wave.  {\\it ASCA} observations of the NE rim of the Cygnus Loop revealed low abundances relative to the solar values by a factor of $\\sim$5 (Miyata \\& Tsunemi 1999).  A follow-up {\\it Suzaku} observation of the region confirmed the depleted abundances there (Miyata et al.\\ 2007).  {\\it Chandra} observation of the southwestern rim of the Cygnus Loop revealed that oxygen abundance there is also depleted by the same factor as that observed in the NE rim (Leahy 2004).  The low abundances were considered to be a common feature in the rim of the Cygnus Loop.  However, recent {\\it Suzaku} observations of the northern outermost region in NE rim of the Cygnus Loop revealed enhanced abundances: a white dashed polygon in Fig.~\\ref{fig:HRI image} showed high metal abundances relative to those in the other region: C$\\sim$0.7, N$\\sim$0.7, O$\\sim$0.4, Ne$\\sim$0.6, Mg$\\sim$0.3, and Fe$\\sim$0.3 times the solar values (Katsuda et al.\\ 2008). Neither a circumstellar medium, fragments of ejecta, nor abundance inhomogeneities of the local interstellar medium around the Cygnus Loop can explain the relatively enhanced abundance in the region. The origin of the abundance inhomogeneity in the NE rim remained as an open question.  Here, we present results from {\\it Chandra} observations of the NE rim of the Cygnus Loop.  Utilizing the high spatial resolving power of {\\it Chandra}, we reveal detailed spatial distributions of metal abundances, emission measure, and thermodynamic parameters from our spatially resolved spectral analysis.  In this paper, we attempt to reveal the origin of the abundance enhancements in the northern outermost region of the NE rim as well as to understand the cause of the abundance depletion in the rest of the NE rim. ",
        "conclusions": "We analyzed archival {\\it Chandra} data of the Cygnus Loop NE rim where abundance inhomogeneities were found by recent {\\it Suzaku} observations (Katsuda et al.\\ 2008).  Thanks to the superior spatial resolving power of {\\it Chandra}, we were able to carry out detailed spatially resolved spectral analyses for the region. We revealed that the abundance-enhanced region was concentrated in a $\\sim$200$^{\\prime\\prime}$-thickness region behind the shock front and confirmed that the values of abundances were consistent with the solar values by a factor of $\\sim$2.  Also, we found that the emission measure decreased toward the shock front.  These features showed stark contrast with those seen in the Vela shrapnels, indicating that the abundance enhancements in the NE rim of the Cygnus Loop were not likely due to fragments of ejecta.  We suggested that the plasma in the abundance-enhanced region originated from the ISM, whereas the plasma in the rest of the NE rim (i.e., abundance-depleted region) originated from the cavity wall or the gas in the cavity."
    },
    "0806/0806.0163.txt": {
        "abstract": "In an effort to explain the short-timescale variability of the broad, double-peaked profiles of some active galactic nuclei, we constructed stochastically perturbed accretion disk models and calculated H$\\alpha$ line profile series as the bright spots rotate, shear and decay. We determined the dependence of the properties of the line profile variability on the spot properties. We compared the variability of the line profile from the models to the observed variability of the H$\\alpha$ line of Arp~102B and 3C~390.3. We find that spots need to be concentrated in the outer parts of the line emitting region to reproduce the observed variability properties for Arp~102B. This rules out spot production by star/disk collisions and favors a scenario where the radius of marginal self-gravity is within the line emitting region, creating a sharp increase in the radial spot distribution in the outer parts. In the case of 3C~390.3, all the families of models that we tested can reproduce the observed variability for a suitable choice of model parameters. ",
        "introduction": "\\subsection{Background and Motivation}  Double-peaked broad Balmer emission lines are found in 20\\% of radio loud active galactic nuclei (AGNs) at $ z<0.4 $ (Eracleous \\& Halpern 1994; 2003) and 4\\% of the Sloan digital Sky Survey (SDSS) quasars at $z<0.33$ \\citep{Sal03}. A number of models have been suggested for the origin of Double-peaked emission lines (DPEL), but most of them face significant challenges when compared to observations. The velocity drift of the two peaks and the single-peaked profile of high-ionization emission lines \\citep{Eracal97} are not consistent with the binary black hole scenario suggested by Gaskell (1983). The observed response of the line profiles to changes in the ionizing flux \\citep{Dietal98, Obrien98} disagrees with predictions from the bipolar outflow model proposed by Zheng et al. (1990). The disagreement between the observed radio lobe geometry and the outflow geometry required to fit the line profile is also a challenge for this model. Wanders et al. (1995) suggested that the double-peaked profile could be produced by a collection of anisotropically illuminated clouds. Even though this scenario agrees with most observations, it has theoretical difficulties since these clouds would be destroyed rapidly by collisions or drag. Finally, Chen, Halpern, \\& Filippenko (1989) and Chen \\& Halpern (1989) suggested that the DPEL are emitted in a zone in an accretion disk over a narrow range of radii (i.e. $r_{\\rm out}/r_{\\rm in}< 10$). In this scenario, the high ionization emission lines are produced at the base of a disk-wind \\citep{CS87}. Thus, the profiles of the high-ionization lines are single-peaked because of radiative transfer effects in the accelerating outflow \\citep{MC97}. The limitation to this model is that the accretion disk structure required to explain the variability of the line profiles cannot be axisymmetric. However one does not expect a perfectly axisymmetric accretion disk so this is not a major challenge. The model of DPEL production in the accretion is the one that holds up best to close scrutiny, thus we adopt it as our working model.  DPEL profiles are observed to vary on timescales of months to years, i.e. on timescales of the order of the dynamical time or longer (e.g., Veilleux \\& Zheng 1991; Zheng et al. 1991; Marziani et al. 1993; Romano et al 1998; Gilbert et al. 1999; Sergeev et al. 2000; Shapovalova et al. 2001; Storchi-Bergmann et al. 2003; Gezari et al. 2007). The line profile variability  does not appear to be correlated to changes in the line and/or continuum flux so it likely traces changes in the accretion disk structure. The long-timescale variablity of the DPEL profile of some objects has been successfully modeled by the precession of a non-axisymmetric accretion disk such as an elliptical disk or a disk with a spiral arm \\citep[and references therein]{Gezarith, Storchietal03, Shap01, Gilbert}. These models, however, fail to explain the long timescale variability of some objects and the short-timescale variability of all objects \\citep{Lewis05}.  Other attemps at explaining the line profile variability through perturbations of the disk structure have introduced bright spots over an axisymmetric accretion disk. Newman et al. (1997) succesfully modeled the variation of the H$\\alpha$ peak intensity ratio of Arp~102B with a single spot rotating within the disk, but Gezari et al. (2007) were not able to apply the same model to the same object at a different time period. It is noteworthy that a similar bright spot model was recently used by Turner et al. (2006) to explain the variability of the Fe K$\\alpha$ line profile of Mrk 766 in the X-ray band. Sergeev et al. (2000) tried a different approach; they modeled the variations of Arp~102B with a collection of 1500 clouds in a disk and were able to reproduce the root mean square (r.m.s.) spectrum and a quasi-sinusoidal variation of the excess flux. This model however introduced phase effects and an edge-on disk in disagreement with the inclination obtained from fitting the total line profile \\citep{CH89, Chenal89, Gezarith}. A  collection of clouds orbiting the central black hole can also explain the power-law shape of the X-ray power spectra observed in AGNs \\citep{Abram91, BO95}. Finally, Pariev \\& Bromley (1998) modeled the Fe K$\\alpha$ line profile resulting from a turbulent disk, which was created by adding a random frequency shift to each pixel in the accretion disk. Motivated by the ideas of Sergeev et al. (2000), we investigate in this paper whether a stochastically perturbed disk model can explain the short and long timescale variation of the DPEL H$\\alpha$ profiles of the two best-monitored objects: Arp~102B and 3C~390.3.  \\subsection{The Physical origin of Bright Spots}  Many processes can lead to density or temperature inhomogeneities, hence brightness perturbations, in the accretion disk: self gravity, disk-star collisions, and baroclinic vorticity. The properties of the inhomogeneities (size, radial distribution, lifetime, resistance to Keplerian shear) depend on their specific formation process, as we outline below.  \\begin{itemize} \\item A disk will fragment under the effect of self-gravity when the Jeans instability sets in. The radius of marginal self-gravity depends on the properties of the disk and the central object and can be as low as 500~$r_{\\rm g}$ ($r_{\\rm g}\\equiv GM/c^2$ is the gravitational radius) for objects with an Eddington ratio of $L/L_{\\rm Edd}\\sim 10^{-3}$ \\citep{GT04}. The radius of marginal self-gravity can be even smaller with more extreme disk properties than those assumed by these authors, such as a Shakura-Sunyaev viscosity parameter less than 0.3 \\citep{SS73}. The size of the clumps will be smaller than the local Jeans radius $r_{\\rm J}=c_{\\rm s}/\\sqrt{G\\rho}\\propto\\xi^{-3/20}$, with a typical size of $10~r_{\\rm g}$ at 1000~$r_{\\rm g}$ for a $10^8~M_{\\odot}$ black hole. The clumps produced in this manner do not shear with differential rotation and they have a high density (and hence brightness) that varies very little over time. The radial distribution of the clumps should be quite uniform past the radius of marginal self-gravity.  \\item Disk-star collisions are thought to occur as frequently as once per day \\citep{Zurekal95}. Such collisions increase the local disk temperature, and hence create a bright spot. These spots will shear with Keplerian rotation and will fade away as the material cools down. The typical size of a spot at the time of formation is at least equal to the size of the star, which is very small in units of $r_{\\rm g}$. As shown by Zurek et al. (1995), the radial distribution of the spots created by star collisions will increase with decreasing radius since collisions are more frequent at smaller radii due to the increased Keplerian frequency of the star about the black hole: $\\nu_{\\rm collison}(r)\\propto r^{p-1/2}$ with $0 < p < 1/2$.  \\item The radial temperature gradient in the accretion disk combined with the Keplerian differential rotation of the disk material will lead to baroclinic vorticity \\citep{PSJ07}. Gas drag will then cause the disk material to spiral to the center of the vortex increasing the local density, hence the brightness, of the vortex \\citep{BM05}. The lifetime and resistance of the spots to shearing is still a subject of debate. The simulations of Petersen et al. (2007) create long-lived, non-decaying, non-shearing vortices. Barranco \\& Marcus (2005) also find the same long-lived vortices when the vorticity is larger than the shearing; but they also find short-lived, shearing vortices when the vorticity is not high enough or when the gas speed is supersonic. Simulations by Johnson \\& Gammie (2005) lead to long-lived, non-shearing vortices that slowly decay as $t^{-1/2}$. All of these models produce structures whose size is comparable to the scale height of the disk, which scales as $r^{9/8}$. The distribution of the spots produced by vorticity is weakly dependent on the radius. \\end{itemize}  We investigate the effects of these different bright spot properties on the variability of DPEL profiles. In \\S 2, we present a method to quantitatively characterize the observed line profile variability and apply it to two AGNs: Arp~102B and 3C~390.3. In \\S 3, we describe the mathematical formalism for computing line profiles in our model. In \\S 4, we explore the parameter space of the model and compare the results to the observations. In \\S 5, we discuss the results and their implications . ",
        "conclusions": "Stochastically perturbed disk models produce variability of the H$\\alpha$ line profile whose properties depend on the model parameters. Allowing for decay of the spots steepens the observed periodograms of the line parameters. This is because the decay introduces a new, longer timescale in the system, thus shifting power to lower frequencies The radial distribution of the spots also influences the power-law indices of the line parameter periodograms: steeper periodograms are produced when spots are concentrated towards the outer parts of the line-emitting region. This is a result of the fact that in such a distribution more spots have longer orbital periods, introducing more power at lower frequencies, thus steepening the periodogram. Shearing does not significantly change the power-law indices of the periodograms because shearing and decaying happen simultaneously and the effect of decaying of a spot is more prominent that that of shearing. Because shearing does not modify the variability properties of the line profile, one cannot constrainr it from observations of the line profile variability. The fractional rms amplitude of the line parameters increases with increasing contrast and size of the spots and does not depend strongly on the number of spots, their radial distribution, or their shear and decay properties. The fractional rms amplitude thus provides another constraint on the contrast of the perturbation of the emissivity function.  %All three families of the stochastically clumpy disk model produce significant power peaks in the line profile parameters power spectra and in the 2-D periodograms at seemingly random frequencies (and velocities). This is a positive output of this model since global models that were previously used to explain line profile variability fail at this. We also note that the two families of models that include fading of the spots have a higher probability of producing power peaks at the same frequencies as the observed values. {\\it why?}  The perturbed disk model can reproduce the observed variability properties of Arp~102B and 3C~390.3 with different levels of success. Most of the observed line parameter periodograms of Arp~102B are too steep to be successfully reproduced. The best models have decaying spots, mostly concentrated towards the outer parts of the line emitting region. Such models produce power-law indices of the peak flux ratio and blue peak position periodograms that are comparable to those of Arp~102B, but all the other line parameter power-law indices are too high by up to 0.4. The maximum radial spot distribution index that we tested is +2. A model with an even steeper radial distribution might produce steep enough periodograms to match the observations. Such a distribution rules out the production of spots by collisions with stars since this process skews the spot distribution towards the inner part of the accretion disk. Spot production by vortices yields a priori a uniform spot distribution. Spot production by self gravity leads to spots only past the self gravity radius so this could skew the spot distribution towards the outer part of the line emitting region. These spots would however neither shear nor decay so one would need an even steeper radial distribution index to achieve agreement with the observed values of Arp~102B. A sharp increase from 0 to a large number of spots past the radius of marginal self-gravity might mimic such a steep radial distribution index. Using the properties of Arp~102B determined by Lewis \\& Eracleous(2006) and the formulae for the radius of marginal self-gravity ($r_{\\rm sg}$) given by Goodman \\& Tan (2004) and Hur\\'e (1998), we estimate that $400<r_{\\rm sg}/r_{\\rm g}<900$ by varying the disk parameters within their plausible range. This estimate falls within the range of the line emitting region determined by fitting the average spectrum. To test whether a sharp boundary in the spot distribution steepens the power-law indices of the periodograms, we ran one simulation of the first family of models with no spots below 750 $r_{\\rm g}$ and 100 spots uniformly distributed above that radius. The spots were large (15 $r_{\\rm g}$) and bright ($\\langle C_i\\rangle =20$) as one would expect from self-gravitating clumps. The resulting line parameter periodograms had power-law indices varying between $-0.17$ and $-0.72$, which is steeper than what a smooth radial spot distribution produces and closer to the values observed for Arp~102B. Moreover, these self-gravitating spots would have the necessary size and contrast to match the observed fractional rms amplitude of Arp~102B. This result bolsters the hypothesis that spots inhabit the outer parts of the line-emitting disk of Arp~102B.  There could also be other factors that steepen the line parameter periodograms. We assumed constant inner and outer radii of the line emitting region, but we tested the effect of a {\\it fast} modulation of these radiii and found that this produces steeper periodograms. The inner and outer radii could vary due a change in the disk ionization structure after a variation of the ionizing continuum from the central source. Developing a physically self-consistent model to account for this effect requires one to calculate the variation of  the whole disk structure with ionizing flux and is beyond the scope of this paper.  The line parameter periodograms of 3C~390.3 are not as steep as those of Arp~120B and can be repoduced by a wide range of model parameters. A radial distribution of spots skewed towards the outer parts of the line emitting region is not absolutely necessary for 3C~390.3 for two reasons. First, the steepening of the line parameter periodograms with increasing radial spot distribution index is not as marked as for Arp~102B. This might be a result of the fact that the line emitting region of 3C~390.3 has a much larger outer radius than that of Arp~102B. Hence, there is naturally a larger number of spots at large radii. This suggests that it is the relative number of spots at large radii compared to lower radii rather than the absolute number of spots at large radii that influences the power-law indices of the periodograms. Also since the observed power-law indices of the line parameter periodograms of 3C~390.3 are higher than those of Arp~102B, we do not need to select the model families and parameters producing steep power spectra in order to reproduce the properties of 3C~390.3. This also means that no particular model for the origin of spots in the accretion disk of 3C~390.3 can be ruled out or favored. We note that $250<r_{\\rm sg}/r_{\\rm g}<400$ for 3C~390.3, which is below $\\xi_{\\rm min}$ and strengthens our assumption of a perturbed accretion disk at all radii.  The behavior of the collapsed periodograms from the 2-D periodogram is a greater challenge to this model. The difference between the two power-law indices agree with the observed difference, but the values of the indices are too high to be reconciled with the observations.  The results of our line profile variability study are quite different between Arp~102B and 3C~390.3 even though their black hole masses are comparable. The main driver of this difference is the higher accretion rate of 3C~390.3 ($0.15~{\\rm M}_\\odot~{\\rm yr}^{-1}$ vs. $1.7\\times 10^{-3}~{\\rm M}_\\odot~{\\rm yr}^{-1}$ for Arp~102B), which makes that disk more unstable to self-gravity. It is encouraging that the perturbed accretion disk model provides a unified solution to explain the profile variability of two AGNs that are so different. In the context of the production of clumps by self-gravity, AGNs with similar black hole masses and even higher Eddington ratios than 3C~390.3 are expected to be unstable at all radii and the variability of their profile would then be well explained by our model. AGNs with smaller black hole masses might have their radius of marginal self-gravity in the line-emitting region and have a variability pattern similar to Arp~102B. In order to test this theory and confirm our result, we need to apply this method to a larger sample of AGNs covering a wider range of black hole masses and accretion rates, such as those presented in Gezari et al. (2007) and Lewis (2005)."
    },
    "0806/0806.0626_arXiv.txt": {
        "abstract": "{\\it Chandra} X-ray observations routinely resolve tens to hundreds of low-mass X-ray binaries (LMXBs) per galaxy in nearby massive early-type galaxies. These studies have raised important issues regarding the behavior of this population of remnants of the once massive stars in early-type galaxies, namely the connection between LMXBs and globular clusters (GCs) and the nature of the LMXB luminosity function (LF). In this paper, we combine five epochs of {\\it Chandra} observations and one central field {\\it Hubble Space Telescope} Advance Camera for Surveys observation of NGC~4697, one of the nearest, optically luminous elliptical (E6) galaxies, to probe the GC-LMXB connection and LMXB-LF down to a detection/completeness limit of $0.6/1.4\\times 10^{37} {\\rm \\, ergs \\, s}^{-1}$. We detect 158 sources, present their luminosities and hardness ratios, and associate 34 LMXBs with GCs. We confirm that GCs with higher encounter rates ($\\Gamma_h$) and redder colors (higher metallicity $Z$) are more likely to contain GCs, and find that the expected number of LMXBs per GC is proportional to $\\Gamma_{h}^{0.79^{+0.18}_{-0.15}} \\, (Z/Z_\\odot)^{0.50^{+0.20}_{-0.18}}$, consistent with fainter X-ray sources in Galactic GCs and LMXBs in Virgo early-type galaxies. Approximately $11\\pm2/8\\pm2$\\% of GCs in NGC~4697 contain an LMXB at the detection/completeness limit. We propose that the larger proportion of metal-rich GCs in NGC~4697 compared to the Milky Way explains why these fractions are much higher than those of the Milky Way at similar luminosities. We confirm that a broken power-law is the best fit to the LMXB-LF, although we cannot rule out a cutoff power-law, and argue that this raises the possibility that there is no universal form for the LMXB-LF in early-type galaxies. We find marginal evidence for different LFs of LMXBs in GCs and the field  and different spectra of GC-LMXBs and Field-LMXBs. ",
        "introduction": "\\label{sec:n4697x_intro}  Observations with the {\\it Einstein Observatory} revealed that early-type galaxies can be luminous X-ray sources \\citep{FJT1985}. Prior to the launch of the {\\it Chandra X-ray Observatory}, observations of X-ray faint galaxies, galaxies with relatively low X-ray-to-optical luminosity ratios, indicated the presence of two distinct spectral components: a soft ($\\sim 0.2 {\\rm \\, keV}$) component \\citep{FKT1994,P1994,KFM+1996} and a hard ($\\sim 5$--$10 {\\rm \\, keV}$) component \\citep{MKA+1997}. The soft component was attributed to hot interstellar gas and the hard component to low-mass X-ray binaries (LMXBs; \\citealt{KFT1992}). Starting with the \\textit{Chandra X-ray Observatory} observation of the X-ray faint elliptical, \\object{NGC 4697} \\citep[hereafter Papers I and II]{SIB2000,SIB2001}, this picture has been confirmed; the majority of X-ray emission in X-ray faint early-type galaxies has been resolved into X-ray point sources, whose properties are consistent with LMXBs. With its ability to resolve LMXBs and accurately measure their positions, {\\it Chandra} has raised at least two major issues regarding the behavior of the LMXBs as a population, namely the connection between LMXBs and globular clusters (GCs) and the nature of the LMXB luminosity function (LF).  From {\\it Chandra}, a high percentage $(\\sim 20-70\\%)$ of LMXBs have positions coincident with globular clusters \\citep[GCs,][]{ALM2001,SKI+2003}; only $\\sim 10\\%$ of Milky Way LMXBs are found in GCs. This raised the question as to where LMXBs in early-type galaxies formed. One early suggestion was that all LMXBs in early-type galaxies formed in GCs, with LMXBs observed to be in the field of the galaxy (Field-LMXBs) having escaped from GCs through supernovae kick velocities, stellar dynamical processes, or the dissolution of the GC due to tidal effects \\citep{WSK2002}. Later analyses have suggested that a significant fraction of Field-LMXBs were formed in situ \\citep{J2005,I2005}, although there is evidence that more Field-LMXBs in lenticular galaxies, as opposed to elliptical galaxies, may have originated in GCs and later escaped into the field \\citep{I2005}. The extent to which LMXBs can be used to probe GC formation and evolution, as well as LMXB formation, depends critically on understanding the GC/LMXB connection. Given that the primary binary formation scenario in GCs is thought to involve dynamical interactions, as opposed to the predominantly primordial nature of binaries formed in the field, these different populations should trace different LMXB formation histories and may trace different star formation histories in early-type galaxies.  {\\it Chandra} has also explored the LF of LMXBs. At the bright end of the LF of NGC~4697, a possible break near the Eddington limit of a $1.4 \\, M_{\\sun}$ neutron star (NS) was found in \\citetalias{SIB2000}. This break was argued to be due to the presence of two LMXB populations, a black hole (BH) population at the bright end, and a predominantly NS population at the faint end. \\citet{BD2004} argued that such a break may be due to ultracompact binaries. Although \\citet{KF2003} argue that no break was required in NGC~4697 after correcting for incompleteness, they do find that a break near $5\\times10^{38} {\\rm \\, ergs \\, s}^{-1}$ gives an improved fit compared to a single power-law for a uniformly selected, incompleteness corrected sample of 14 early-type galaxies. At the faint end, \\citet{VG2005} found that the LF in \\object[NGC 5128]{Cen A} flattens significantly below $5\\times10^{37} {\\rm \\, ergs \\, s}^{-1}$ and follows the $dN/dL \\propto L^{-1}$ law in agreement with the behavior found for LMXBs in the Milky Way and the bulge of M31. One interpretation of this break in the Milky Way LF is that there are two populations of short $(\\lesssim 20 {\\rm \\, hr})$ period LMXBs: those where magnetized stellar winds dominate mass transfer, and those where gravitational radiation drives the accretion \\citep{PK2005}. Measuring and understanding the LF of LMXBs in early-type galaxies has clear implications in our understanding of the number and type of binary systems that make up the zoo of LMXBs.  As more LMXBs are detected and characterized, more examples of extragalactic LMXB candidates with extreme behaviors, such as supersoft sources (SSs) and ultraluminous X-ray sources (ULXs) are being identified by {\\it Chandra} and studied in greater detail. The SSs have very soft X-ray spectra $(\\lesssim 75 {\\rm \\, eV})$ similar to SSs in our Galaxy and M31 that are generally believed to be accreting white dwarfs (WDs); however, the bolometric luminosities of extragalactic SSs can exceed the Eddington luminosity for a Chandrasekhar mass WD. One hypothesis is that these sources contain intermediate-mass ($\\sim 10^2$--$10^3 \\, M_{\\sun}$) accreting BHs \\citep{SGS+2002}. A great deal of attention has been spent on ULXs, which we define as $L_{X} > 2 \\times 10^{39} {\\rm \\, ergs \\, s}^{-1}$ (0.3--10.0 keV). However, except possibly for LMXB candidates in \\object{NGC 720} \\citep{JCB+2003}, \\object{NGC 1399} \\citep{ALM2001}, \\object{NGC 1600} \\citep{SSC2004}, and NGC~4482 \\citep{MKZ+2007}, ULXs are generally not found within old stellar systems beyond the number expected from unrelated foreground or background sources \\citep{IBA2004}.  By stacking multi-epoch {\\it Chandra} observations of an early-type galaxy, more, fainter LMXB candidates can be studied, and brighter LMXB candidates can be studied in greater detail. In this Paper, we report on multi-epoch {\\it Chandra} X-ray observations of NGC~4697, one of the nearest \\citep[$11.3 {\\rm \\, Mpc}$; see footnote 18 of][]{JCB+2005} optically luminous ($M_B < -20$) elliptical (E6) galaxy. (There is a weak disk; however, it not comparable to those seen in lenticular galaxies \\citep{PDI+1990}). We adopted the 2MASS Point Source Catalog \\citep{SCS+2006} position of R.A.\\ $= 12^{\\rm h} 48^{\\rm m} 35\\fs90$ and Dec.\\ $= -5\\arcdeg48\\arcmin02\\farcs6$ as the location of the center of NGC~4697 and the Third Reference Catalogue of Bright Galaxies (RC3) optical photometry values for the effective radius ($r_{\\rm eff}=72\\farcs0$), position angle ($PA=70\\arcdeg$, measured from north to east), and ellipticity ($e=0.354$), which assumes a de Vaucouleurs profile \\citep{VVC+1992}. From these values, we calculate the optical photometry's effective semi-major distance ($a_{\\rm eff} = 89\\farcs6$). This galaxy lies $\\sim 5 {\\rm \\, Mpc}$ in front of the bulk of the galaxies in the Virgo cluster, and is $18\\fdg7$ south of \\object{M87}, the galaxy at the dynamical center of the Virgo cluster. We also report on the GC/LMXB connection as determined from a joint observation of the central region by the {\\it Hubble Space Telescope} Advance Camera for Surveys \\citep[{\\it HST}-ACS;][]{FBB+1998}.  In \\S~\\ref{sec:n4697x_obs}, we discuss the observations and data reduction of NGC~4697. The X-ray image and the detection of X-ray sources are discussed in \\S~\\ref{sec:n4697x_image} and \\S~\\ref{sec:n4697x_detections}. In \\S~\\ref{sec:n4697x_opt_ids}, we discuss the optical counterparts from ground-based and {\\it HST} observations. We examine the GC/LMXB connection in detail in \\S~\\ref{sec:n4697x_gclmxb}. The analyses of luminosity, hardness ratios, and spectra are considered in \\S~\\ref{sec:n4697x_src_lum}--\\ref{sec:n4697x_src_spectra}. Finally, we summarize our conclusions in \\S~\\ref{sec:n4697x_conclusion}. We concentrate on the variability of the X-ray sources in our companion paper (Sivakoff et al.\\ 2008b, hereafter Paper V). Unless otherwise noted, all errors refer to $1 \\sigma $ confidence intervals, count rates are in the $0.3$--$6 {\\rm \\, keV}$ band, and fluxes and luminosities are in the $0.3$--$10 {\\rm \\, keV}$ band, with absorption effects removed. ",
        "conclusions": "\\label{sec:n4697x_conclusion} Multi-epoch {\\it Chandra} observations reveal a wealth of information on LMXBs in NGC~4697, the nearest optically luminous elliptical galaxy. We detect 158 sources, 126 of which have their count rates determined at $\\ge 3 \\sigma$. Ten sources have optical counterparts in ground-based catalogs, including a known AGN (Source 117) and the foreground star \\object{BD-05 3573} (Source 156). With our {\\it Hubble} observations of the galaxy center, we find 36 additional optical counterparts. Most importantly, we identify 34 LMXBs clearly associated with GCs.  We confirm that GCs that are optically brighter ($4.5\\sigma$) and redder ($3.0\\sigma$) are more likely to contain GCs. We find that GCs with larger encounter rates are also more likely to contain GCs ($5.5\\sigma$). When we fit the expected number ($\\lambda_t$) of LMXBs in a GC, we find $\\lambda_t \\propto \\Gamma_{h}^{0.74^{+0.14}_{-0.13}} \\, (Z/Z_\\odot)^{0.50^{+0.18}_{-0.16}}$, where $\\Gamma_{h}$ is the encounter rate and $Z/Z_\\odot$ is the metallicity of the GC. Our results agree well with those found for fainter X-ray sources in Galactic GCs \\citep{PLA+2003} and LMXBs in M87 \\citep{JCF+2004}. These results are also consistent with \\citetalias{SJS+2007}; however, our NGC~4697 data set is included in that analysis.  We detect sources with X-ray luminosities $> 6\\times10^{36} {\\rm \\, ergs \\, s}^{-1}$. The fraction of LMXBs associated with GCs, $f_{X,{\\rm GC}}$, is $38.4^{+6.1}_{-5.7}\\%$ and does not appear to depend on X-ray luminosity. We find $10.7^{+2.1}_{-1.8}\\%$ of GCs contain an LMXB at the detection limit, although we note that it is likely that the percentage of GCs with an active LMXB is even higher due to the X-ray flux limit of the current observations. Furthermore, our X-ray detections are not complete at the detection limit. At the luminosity limit of our Analysis Sample ($1.4\\times10^{37} {\\rm \\, ergs \\, s}^{-1}$), which is $>89\\%$ complete, $8.1^{+1.9}_{-1.6}\\%$ of GCs contain an LMXB. This is the third deepest probe of the GC/LMXB connection in an early-type galaxy to date. [At $3.4 {\\rm \\, Mpc}$, fainter luminosities can be more easily reached with \\object{Cen A}; however, studies of the GC-LMXB connection in Cen~A are made more difficult by its larger angular extent on the sky and recent star-formation \\citep{I1998}]. Deep observations of NGC~3379 probe deeper in luminosity; however, only nine of its GCs contain LMXBs.] At this same limit, there have been two ($1.3^{+1.7}_{-0.9}\\%$) Galactic GCs containing LMXBs (NGC 6440 and NGC 6624) over the history of X-ray astronomy. The discrepancy between the Milky Way and NGC 4697 may be explained by their different GC metallicity distributions (a 3:1 metal-poor to metal rich ratio for Galactic GCs as compared to a 1:1 ratio for NGC 4697 GCs). Since metal-rich GCs ($[{\\rm Fe/H}] > -0.75$) are about 3 times as likely to contain LMXBs, NGC 4697 is predicted to have about twice the percentage of GCs with LMXBs as the Milky Way. This correction eliminates most of the discrepancy between the two galaxies.  We have determined the X-ray luminosity functions from each individual observation, from the combination of our five observations, and the LF of the non-variable sources. There is no statistically significant difference in the LFs of the different observations. This result is critical because it validates using single-epoch observations to measure LFs. While we clearly rule out a single power-law LF, we cannot definitively rule out cutoff power-law models with slopes of $\\alpha = 1.5\\pm0.2$ and cutoff luminosities of $(6^{+4}_{-3}) \\times 10^{38} {\\rm \\, erg\\,s}^{-1}$. Broken power-law models (eq.~[\\ref{eq:n4697x_lfb}]) provide the best fits to our LFs. We adopt our fit of the instantaneous LF as our best-fit, with $N_{0,b} = 3.1\\pm1.5$, $\\alpha_{\\rm l} = 0.83\\pm0.52$, $\\alpha_{h} = 2.38\\pm0.33$, and $L_{\\rm b} = ( 10.8\\pm2.9 ) \\times 10^{37} {\\rm \\, ergs \\, s}^{-1}$. We note that \\cite{KF2004} found evidence for a possible break in the LF of LMXBs in early-type galaxies at slightly larger luminosities; however recent deep observations of NGC~3379 and NGC~4278 have not found strong evidence for such a break \\citep{KFK+2006}. This raises the possibility that there is no universal form for the LF of LMXBs in early-type galaxies.  We find marginal evidence (significant at the $2.1\\sigma$ level) that a larger number of LMXBs above the Eddington limit for a hydrogen accreting $1.4 {\\rm \\, M_\\odot}$ NS tend to be found in GCs than in the field. Although this is consistent with results in \\citet{ALM2001} in NGC 1399, we believe this result needs to be tested with a larger sample. One possible \\citep[e.g.,][]{KMZ2007} explanation is that multiple LMXBs might exist in some GCs. We predict that this effect is small, only $\\sim4$ of the GCs are expected to have multiple LMXBs with total X-ray luminosity above $1.4 \\times 10^{37} {\\rm \\, ergs \\, s}^{-1}$, which corresponds to $\\sim 16\\%$ of the GCs above that X-ray luminosity. An alternative explanation is that any possible discrepancy in the LFs occurs at lower luminosities. Such discrepancies have been seen for the bulge of M31 and the Milky Way \\citep{VG2007} and the elliptical galaxy NGC~3379 \\citep{FBZ+2008}; however, this effect is most evident at luminosities below those probed by our observations of NGC~4697.  Our spectrum of sources in the inner effective semi-major axis is best-fit by bremsstrahlung emission with Galactic absorption ($kT = 9.1^{+1.3}_{-1.1} {\\rm \\, keV}$, $N_{H} = 2.14 \\times 10^{20} {\\rm \\, cm}^2$). Both hardness ratios and spectral analysis indicate that the spectra of X-ray sources at large radii ($a \\gtrsim 180\\arcsec$) differ from those at small radii. We believe that this effect is due to an increasing dominance of unrelated foreground and background sources, particularly background AGNs. The spectra of GC-LMXBs and Field-LMXBs appear to differ (significant at the $3.3\\sigma$ level). The GC-LMXBs are better fit with higher temperatures and greater absorbing columns compared to the Field-LMXBs. Similarly, we find a difference (significant at the $3.5\\sigma$ level) between X-ray fainter and brighter LMXBs. The fainter LMXBs tend to have smaller absorbing columns, while the brighter LMXBs tend to have a small excess in absorption, which may be due to having more accreting material. We find that spectra of LMXBs in metal-poor GCs have harder temperature and lower absorbing columns than those in metal-rich GCs; however, this marginal result is only significant at the $1.9\\sigma$ level. In all cases, the sources with a smaller absorption column tend to be fit with sub-Galactic absorbing columns. This indicates that the spectral model (folded in with the calibration) underpredicts the soft emission in the spectra. To probe the cause of the spectral differences, we require a more physically accurate model and better understanding of the calibration at low energies.  Among the spectral fits to individual bright sources, Source 117 (a known AGN) and Source 134 had spectral fits with large absorbing columns. We predict Source 134 is likely to be an AGN. The spectral fit and variability in Source 156 (the foreground star \\object{BD-05 3573}) are consistent with it being an FK Comae star, a chromospherically active giant."
    },
    "0806/0806.2235_arXiv.txt": {
        "abstract": "% We discuss the formation, evolution and observational parameters of the population of short-period ($\\raise-.5ex\\hbox{$\\buildrel<\\over\\sim$}10$\\,hr) low-mass black-hole binaries (LMBHB). Their evolution is determined by the orbital angular momentum loss and/or nuclear evolution of the donors.  All observed semidetached LMBHB are observed as soft X-ray transients (SXTs). The absence of observed short-period stable luminous X-ray sources with black holes and low-mass optical components suggests that upon RLOF by the donor, the angular-momentum losses are substantially reduced. The model with reduced angular-momentum loss reasonably well reproduces the masses and effective temperatures of the observed secondaries of SXTs. Theoretical mass-transfer rates in SXTs are consistent with those deduced from observations only if the accretion discs in LMBHB are truncated. The population of short-period LMBHB is formed mainly by systems which at RLOF had unevolved or slightly evolved donors (abundance of hydrogen in the center $X_c\\, \\raise-.5ex\\hbox{$\\buildrel>\\over\\sim\\,$} 0.35$). Our models suggest that a very high efficiency of common envelopes ejection is necessary to form LMBHB. ",
        "introduction": "Currently, ten Galactic dynamically-confirmed black-hole candidate X-ray binaries with K/M spectral type secondaries and orbital periods $\\raise-.5ex\\hbox{$\\buildrel<\\over\\sim$}$1 day have been observed \\citep{2006csxs.book..157M}. All these objects are SXTs. Their X-ray luminosity may vary by 5--8 orders of magnitude between quiescence and the peak of the outburst. Recurrence times spread from a about a year to tens or years and could be even longer. Their variability is interpreted as resulting from a thermal-viscous instability of irradiated accretion discs around black holes in semidetached binaries \\citep[see][and references therein]{2001NewAR..45..449L}. The estimated number of low-mass black-hole binaries (LMBHB) in the Galaxy ranges from several hundred to several thousand \\citep{1997ApJ...491..312C,1998A&A...333..583R}. The list of short-period black-hole SXTs and some of their observed and inferred parameters are rendered in Table 1. Observational data presented in Table 1 is based on the survey of the literature \\citep[see ][]{2008arXiv0802.4375Y}.   Below, we consider the formation, evolution and some observational properties of  LMBHBs.  \\begin{table}[!ht] \\caption{Known short-period black-hole SXTs. } \\smallskip \\begin{center} {\\small \\begin{tabular}{lllccl} \\tableline \\noalign{\\smallskip} No. & Name  & $P_{\\rm orb}$,  & Sp & $\\langle \\dot{M}_{\\rm recc} \\rangle$ & $~~\\dot{M}_{\\rm in}$, \\\\ &&hr&&${\\rm M_{\\odot}~yr^{-1}}$&${\\rm M_{\\odot}~yr^{-1}}$ \\\\ \\hline 1. & XTE J1118+480 (KV UMa)   & 4.10  &  K5-M1  & $1.9 \\times 10^{-12}$   & $3.0 \\times 10^{-10}$  \\\\ 2. & GRO J0422+32 (V518 Per)  & 5.09  & M2-M4   & $1.3 \\times 10^{-11}$   & $1.9 \\times 10^{-10}$  \\\\ 3. & GRS 1009-45 (MM Vel)     & 6.84  & K7-M0.5 & $4.4 \\times 10^{-11}$   & $1.8 \\times 10^{-10}$  \\\\ 4. & XTE J1650-500            & 7.68  & K4      &  $2.0 \\times 10^{-11}$  & $2.8 \\times 10^{-10}$  \\\\ 5. & A0620-00 (V616 Mon)      & 7.75  & K3-K7   & $3.3 \\times 10^{-11}$   & $2.8 \\times 10^{-10}$  \\\\ 6. & GS 2000+25 (QZ Vul)      & 8.28  & K3-K6   &  $2.0 \\times 10^{-10}$  & $5.5 \\times 10^{-10}$  \\\\ 7. & XTE J1859+226 (V406 Vul) & 9.12  & G5-K0   & $4.1 \\times 10^{-10}$   & $1.0 \\times 10^{-9}$  \\\\ 8. & GRS 1124-68 (GU Mus)     & 10.39 & K3-K7   & $3.4 \\times 10^{-10}$   & $3.5 \\times 10^{-10}$  \\\\ 9. & H 1705-25 (V2107 Oph)    & 12.50 & K3-K7   & $5.5 \\times 10^{-11}$   & $4.1 \\times 10^{-10}$  \\\\ 10. &  4U 1543-47  (IL Lup)   & 27.0  &  A2     & $4.2 \\times 10^{-10}$   & $1.3 \\times 10^{-9}$  \\\\ \\noalign{\\smallskip} \\tableline \\end{tabular} } \\end{center} \\vspace{-2mm} {\\small Note:  $\\langle \\dot{M}_{\\rm recc} \\rangle$ -- mass-transfer rate estimate based on recurrence times, $\\dot{M}_{\\rm in}$ -- upper limit to mass-transfer rate based on assumption of maximal truncation of accretion discs in SXTs. }  \\label{tab:pte} \\end{table} ",
        "conclusions": "We calculated models of the Galactic population of short-period low-mass black-hole binaries which are identified with soft X-ray transients. We found that using  the values of the common-envelope parameter $\\alpha_{ce}\\lambda$ between $\\approx (0.5 - 2)$ and assuming a strongly reduced magnetic braking it is possible to reproduce, (within the uncertainty of observations) the number of the LMBHBs in the Galaxy and  the effective temperatures and masses of the donors in these systems (as inferred from the spectra of the stars). The above mentioned values of $\\alpha_{ce}\\lambda$ imply that the common-envelope expulsion in the progenitors of SXTs has to be very efficient and that sources other than orbital energy may be required in this process.  In our model, all short-period LMBHB systems are transient in agreement with observations.  Model mass-transfer rates in LMBHBs are consistent with the upper limits derived from observations under assumption that accretion discs in SXTs are leaky."
    },
    "0806/0806.0971_arXiv.txt": {
        "abstract": "Nonlinear effects are crucial in order to compute the cosmological matter power spectrum to the accuracy required by future generation surveys. Here, a new approach is presented, in which the power spectrum, the bispectrum and higher order correlations,  are obtained -- at any redshift and for any momentum scale -- by integrating a system of differential equations. The method is similar to the familiar BBGKY hierarchy. Truncating at the level of the trispectrum, the solution of the equations corresponds  to the summation of an infinite class of perturbative corrections. Compared to other resummation frameworks, the scheme discussed here is particularly suited to cosmologies other than $\\Lambda$CDM, such as those based on modifications of gravity and those containing massive neutrinos. As a first application, we compute the Baryonic Acoustic Oscillation feature of the power spectrum, and compare the results with perturbation theory, the halo model, and N-body simulations. The density-velocity and velocity-velocity power spectra are also computed, showing that they are much less contaminated by nonlinearities than the density-density one. The approach can be seen as a particular formulation of the renormalization group,  in which time is the flow parameter. ",
        "introduction": "Future generation galaxy surveys are going to measure the statistical properties of matter distribution to an unprecedented accuracy, providing informations on fundamental questions such as the nature of Dark Energy (DE) \\cite{EHT98,SE03} and the absolute scale of neutrino masses (see, for instance, \\cite{LP06}, and references therein). In particular, the location and amplitude of Baryon Acoustic Oscillations (BAO), wiggles in the matter power-spectrum produced by the coupling of baryons to radiation by Thomson scattering in the early universe, for wavenumbers in the range $k\\simeq 0.05 - 0.25 \\;h \\mathrm{Mpc}^{-1}$, have the potential to constrain the expansion history of the Universe and the nature of the DE. BAO's have recently been detected both in the 2dF and SDSS surveys data \\cite{Eis05,Hue05,Pad06,BCB06}, and are going to be measured in the near future in a series of high-redshift surveys (see, for instance \\cite{Hill04,GED05}).  A reliable comparison between theoretical models and observations requires going beyond the linear order in perturbation theory \\cite{JK06,JK08,ABFL07,MP07a, MP07b,RPTBAO}. The more established way to deal with nonlinearities is by means of $N$-body simulations (for recent applications to the matter power spectrum (PS) see for instance \\cite{HSW07,SSEW08,TR08}). However, in order to gain the required sensitivity, very large volumes and high resolutions are required, with the consequence that, due to time limitations, only the `vanilla' type $\\Lambda$CDM cosmologies have been investigated so far. Fitting functions for the nonlinear PS, such as those based on the Halo model e.g. refs.~\\cite{PD96,S03}) are uncapable to reach the required level of accuracy \\cite{HSW07}.  An alternative, semi-analytical approach is represented by perturbation theory (PT) (for a review, see \\cite{PT}), which has recently experienced a renewed interest, mainly motivated by two reasons. First, next generation galaxy surveys are going to measure the PS  at large redshift, where the fluctuations are still in the linear regime and 1-loop PT is expected to work \\cite{JK06,JK08}. Second, techniques for the resummation of some classes of perturbative corrections {\\em to all orders} have been recently developed. These techniques, based on field theory tools such as Feynman diagrams and the renormalization group (RG) \\cite{RPTa, RPTb, MP07a, MP07b} have been shown to be able to render a PS in agreement with that from $N$-body simulations down to $z=0$ \\cite{MP07a, MP07b, RPTBAO}. For related work, see \\cite{Valageas03,Valageas06,McD06,Izumi07,Valageas07,Matsubara07,Taruya2007,BV08}.  While these semi-analytic models have proved themselves a viable alternative to $N$-body simulations (at least for what concerns the Dark Matter (DM) PS), they still suffer from some limitations. First of all, these methods are formulated for an Einstein-deSitter (EdS, $\\Omega_m=1$) cosmology, and further extended to more general ones ({\\it e.g.} $\\Lambda$CDM) by replacing the EdS linear growth factor for the growing mode, $D^{EdS}_+=a$, with that of the other cosmology, $D_+(a)$, where $a$ is the scale factor. While this procedure is correct at the linear order, it introduces some inaccuracies at higher orders as soon as the condition $\\Omega_m/f_+^2=1$ (with $f_+=d\\ln D_+/d\\ln a$) fails \\cite{PTreview}. The problem is that in schemes such as \\cite{RPTa, RPTb, MP07a, MP07b} there is no way to independently assess the validity of such approximation. Moreover, and more importantly, such approaches cannot be trivially extended to cosmologies in which the linear growth factor is also scale dependent, $D_+=D_+(k,\\,a)$, as is the case, for instance, when massive neutrinos \\cite{LP06} contribute to the DM.  On a more technical ground, while the leading corrections for the PS  where correctly identified and computed in \\cite{MP07a, MP07b, RPTBAO}, for the bispectrum (BS) the task turns out to be much more involved (for a computation in a different approach, see \\cite{ValBisp}). Finally, a systematic scheme of approximations was not univocally identified in such frameworks, and the independent assessment of the size of the corrections one is neglecting is not a straightforward task.  In this paper we present an alternative approach to the resummation of infinite classes of PT corrections. The method is based on a straightforward application of the `equations of motion', {\\it i.e.} the continuity, Euler, and Poisson equations, to the computation of correlation functions (PS, BS, etc.) for the density and the velocity fields. The correlation functions evolve in time according to a (truncated) system of differential equations, quite similar to the familiar BBGKY hierarchy ones \\cite{PeebBook}. Compared to the approaches discussed in  \\cite{RPTa, RPTb, MP07a, MP07b}, the one we are going to present is much more free from field theory technicalities and, as such, will be more easily implemented by people without a specific field theoretical background. The correlation functions obey a system of integro-differential equations, closed by a well defined approximation procedure, whose solution gives the PS, BS, etc. at different times and for any momentum configuration.  Remarkably, as we will discuss, the present formulation is much more flexible than the previous ones, and can be straightforwardly applied to more general cosmologies, as those based on modifications of gravity and those containing massive neutrinos \\cite{LMPR08}.  As a first illustration of the possible applications of the method, we compute the PS in the BAO region. We discuss the simplest non-trivial ({\\it i.e.} beyond 1-loop) approximation level and show which class of PT corrections it corresponds to. We consider two cosmologies, a `vanilla' $\\Lambda$CDM, and one in which DE is a quintessence with equation of state $w=-0.8$. For the  $\\Lambda$CDM case, we compare our results with those obtained by $N$-body simulations \\cite{JK06} and by other approaches, and for both cases we assess the error made by not properly treating the decaying mode. We find that it can be as large as $\\simeq 1\\%$ only for $z$ close to zero and scales smaller than those relevant for the BAO.  The density-velocity and velocity-velocity power spectra are also computed, showing that they are much less contaminated by nonlinearities than the density-density one.  The paper is organized as follows. In Section \\ref{EOM} we discuss the fluid equations for a large class of cosmologies and derive, starting from those, the evolution equations for the PS's and the BS's. We also state the only approximation made in this paper, consisting in neglecting the effect of the evolution of the connected four point functions, the trispectrum. In Section \\ref{ANSOL} we derive a formal solution of the system of equations  in our approximation and show to which class of PT corrections it corresponds to. In Section \\ref{COMPARE} we compare the present approach to those of refs. \\cite{RPTa, RPTb, MP07a, MP07b}, and also with the one, closer in spirit, of ref.~\\cite{McD06}. In Section \\ref{NUMERIKA} we give and discuss our numerical results and compare it to other approaches. Finally, in Section \\ref{CONCLUSION} we summarize our results and discuss future lines of work. ",
        "conclusions": "\\label{CONCLUSION} Compared to previous semi-analytic approaches to the resummation of perturbative corrections, the scheme presented in this paper has some clear advantages. First of all, it can be easily extended to a large class of  cosmologies, among which the case of massive neutrinos and of scalar-tensor modifications of gravity are the most notable ones. In practice, the approach in its present form can be used for all models in which the nonlinear terms in the continuity and in the Euler equations are the same as for the EdS case. The evolution can be seen as a series of time steps (in the limit of vanishing step length), during which the fields evolve with the --possibly scale-dependent-- linear growth factors. At the end of each step, the growing and decaying mode are mixed by the universal nonlinear terms. Therefore, the different cosmologies enter only via the different background evolution and the different linear growth factor, all the information being contained in the ${\\bf \\Omega}$ matrix of eq.~(\\ref{bigomega}).  The second advantage is the possibility of formulating a clean and systematic scheme of approximations. The next level, compared to the one considered here, is obviously given by the inclusion of the trispectrum in the system of equations, which will correspond to including vertex corrections. Judging from the present agreement between our results and $N$-body simulations in the BAO region, one would infer that such improved approximation will mainly affect higher momentum scales.  For a more practical point of view, though the present approach is completely compatible with those of refs.~\\cite{RPTa,RPTb,MP07a,MP07b} (see the discussion in Section \\ref{COMPARE}), here field theory tools such as Feynman diagrams, generating functionals, and RG flows, are not necessary in order to compute, for instance, the PS and the BS. What is needed, at the end, is just a code to solve the differential equations in eq.(\\ref{syst}). Therefore, a field theoretic background is not a prerequisite to use this method.  The results we obtain are in good agreement with the N-body simulations of ref. \\cite{JK06}, up to $k$ of order $0.3 - 0.35 \\,\\mathrm{h/Mpc}$, showing that the method is able to describe (from first principles) the non-linearities due to mode-mode coupling in the BAO region. We also showed that the velocity-density and the velocity-velocity PS's are less contaminated by the mode-mode coupling and, since they should also be much less contaminated by bias, their possible role in the future BAO measurements should be studied in more detail.  Finally, we showed that the role of the decaying mode can be safely ignored for $z \\agt1$, and can lead to percent-level effects only at lower redshifts and scales higher than those relevant for the BAO's.  The importance of nonlinear effects for the determination of the upper bound on the neutrino mass scale was emphasized in ref. \\cite{STT08}, where 1-loop PT was used to compute the PS. A better nonlinear approach is clearly desirable to derive firmer neutrino mass bounds \\cite{LMPR08}, possibly including the running of the trispectrum, in order to reach higher momentum scales.  \\appendix"
    },
    "0806/0806.0695_arXiv.txt": {
        "abstract": "These lecture notes provide an introduction to mm/submm extragalactic astronomy, focused on AGN studies, with the final goal of preparing students to their future exploitation of the ALMA capabilities. I first provide an overview of the current results obtained through mm/submm observations of galaxies and AGNs, both local and at high redshift. Then I summarize the main mm/submm facilities that are currently available. ALMA is then presented with a general description and by providing some details on its observing capabilities. Finally, I discuss some of the scientific goals that will be achievable with ALMA in extragalactic astronomy, and for AGN studies in particular. ",
        "introduction": "\\label{sec_intro}  The Atacama Large Millimeter Array (ALMA) is one of the largest ground-based astronomy projects of the next decade, which will revolutionize several fields of astronomy. A large community of scientists is expected to use ALMA to tackle several outstanding questions in astrophysics. However, mm/submm astronomy is often considered a field restricted to experts. In the case of students and young scientists in particular, the limited familiarity with mm/submm facilities and observations may prevent them to fully exploit the ALMA capabilities in the future. These lecture notes are aimed at providing students and young researches some background on mm/submm extragalactic astronomy, with a focus on the investigation of AGNs. I will first provide a quick overview of the current results obtained through extragalactic mm/submm observations, by focusing on AGNs (\\S\\ref{sec_mm_astronomy}). I will then summarize the currently available (and forthcoming) mm-submm facilities (\\S\\ref{sec_current_facilities}). Then I will shortly describe ALMA and summarize its observing capabilities (\\S\\ref{sec_alma}). Finally, I will discuss some of the ALMA prospects for extragalactic studies, and in particular for AGNs, both in the local universe and at cosmological distances (\\S\\ref{sec_alma_prospects}). These lecture notes are far from being exhaustive; several scientific cases will not be discussed at all; the main goal of these notes is only to provide an introduction to mm/submm extragalactic astronomy and to highlight some scientific cases that ALMA will be able to tackle. ",
        "conclusions": ""
    },
    "0806/0806.0376_arXiv.txt": {
        "abstract": "Using deep Chandra and optical spectroscopic observations, we investigate an intriguing, young massive group, RXJ1648.7+6109, at $z=0.376$, and we combine these observations with previous measurements to fit the scaling relations of intermediate-redshift groups and poor clusters.  RXJ1648 appears to be in an early stage of formation; while it follows X-ray scaling relations, its X-ray emission is highly elongated and it lacks a central, dominant BCG.  Instead, RXJ1648 contains a central string of seven bright galaxies, which have a smaller velocity dispersion, are on average brighter, and have less star formation (lower EW([OII]) and EW($H_{\\delta}$)) than other group galaxies. The 4-5 brightest galaxies in this string should sink to the center and merge through dynamical friction by $z=0$, forming a BCG consistent with a system of RXJ1648's mass even if 5-50\\% of the light is lost to an intracluster light component (ICL).  The $L_X-T_X$ relation for intermediate-redshift groups/poor clusters is very similar to the low-redshift cluster relation and consistent with the low-redshift group relation.  In contrast, the $L_X-\\sigma_v$ and $\\sigma_v-T_X$ relations reveal that intermediate-redshift groups/poor clusters have significantly lower velocity dispersions for their X-ray properties compared to low-redshift systems, however the intermediate-redshift relations are currently limited to a small range in luminosity. ",
        "introduction": "Groups and poor clusters of galaxies are the building blocks of larger scale structures, and they are particularly important environments for diagnosing the origin of the non-gravitational heating of the intracluster medium (ICM)(Balogh et al.~2006) and for galaxy evolution (Zabludoff \\& Mulchaey 1998).  For example, at low redshift non-gravitational heating appears to be proportionally more important in the group regime, leading to a steepening of the $L_X-T$ relation for groups versus clusters and excess group entropy above self-similar expectations (e.g.~Helsdon \\& Ponman 2000; Ponman et al.~2003).  The evolution in group X-ray scaling relations with redshift can constrain models of non-gravitational heating (Balogh et al.~2006).  In addition, most galaxies in the local universe lie in groups (e.g.~Turner \\& Gott 1976), making them important environments in the study of galaxy evolution.  We have only recently begun to study the evolution of this important environment with redshift (e.g.~Mulchaey et al. 2006; Jeltema et al. 2006,2007; Willis et al. 2005; Pacaud et al. 2007; Wilman et al. 2005a,b; Gerke et al. 2007) and very few groups/poor clusters at even moderate redshifts ($z>0.2$) have both significant X-ray and optical data (Mulchaey et al. 2006; Jeltema et al. 2006,2007; Willis et al. 2005; Gastaldello et al. 2007). Current X-ray observations at intermediate-redshifts ($0.2<z<0.6$) are limited to systems with $kT \\sim 1$ keV or above, and therefore probe a transitional regime of galaxy associations, which we refer to as massive groups or poor clusters, between what are typically referred to as clusters and lower mass groups of galaxies.  Our group has undertaken a program to study in depth X-ray luminous intermediate-redshift groups/poor clusters selected from the ROSAT Deep Cluster Survey (Rosati et al. 1998).  Spectroscopic confirmation and HST imaging of a sample of nine groups was presented in Mulchaey et al. (2006; Paper I).  Jeltema et al. (2006; Paper II) investigated the X-ray properties of six of these groups based on observations taken with XMM-Newton, and Jeltema et al. (2007; Paper III) presented deeper spectroscopic follow-up of seven groups with the Gemini-North and Keck telescopes.  One of the most interesting results of this study was the discovery that unlike X-ray luminous groups at low redshift, most of these systems either lacked a dominant central galaxy or contained multi-component brightest galaxies indicative of a recent merger (Papers I and III).  In other ways these intermediate-redshift groups are similar to low-redshift groups and clusters; they have X-ray luminosities and temperatures consistent with the low-redshift $L_X-T_X$ relation (Paper II), and they contain high fractions of elliptical galaxies and low fractions of star-forming galaxies (Paper III). These observations are consistent with a picture in which BCGs form/grow late with respect to the collapse of the system itself (De Lucia \\& Blaizot 2007; Dubinski 1998).  In fact, the groups in our sample appear to span a range of evolutionary states.  Of the six groups with XMM-Newton observations, two lack a brightest galaxy at the X-ray center and have elongated X-ray emission, two have bright central galaxies composed of three or more luminous nuclei but round X-ray emission, and two have bright galaxies at the X-ray center with no galaxies of similar luminosity nearby.  In this paper, we investigate in detail one particularly intriguing system in our sample, RXJ1648.7+6109 at $z=0.376$.  Based on previous observations, this massive group appears to be in an early state of evolution.  Instead of a single, dominant galaxy it contains a string of seven galaxies within 200 kpc of the X-ray center, five of which are brighter than $L_{\\ast}$ (Paper III).  RXJ1648's X-ray emission appears to be elongated, but the XMM-Newton observation was almost completely contaminated by background flares leading to large uncertainties in its X-ray temperature and spatial distribution (Paper II).  Here we report on a new $\\sim$ 100 ks Chandra observation of RXJ1648, which allows us to determine its average X-ray properties and spatial structure (\\S 3.1).  Combining these new X-ray measurements with the other groups in our sample and intermediate-redshift groups/poor clusters from the literature, we derive the $L_X-T_X$, $L_X-\\sigma_v$, and $\\sigma_v-T_X$ relations for poor clusters and groups at intermediate redshifts (\\S 3.3).  In addition, we investigate in greater detail the galaxy properties in RXJ1648, including the magnitude, star formation, and velocity distributions of the member galaxies (\\S 3.2) and the future formation of a BCG through the mergers of the bright galaxies in the central string (\\S 4.1).  Throughout the paper, we assume a cosmology of $H_0=70h_{70}$ km s$^{-1}$ Mpc$^{-1}$, $\\Omega_m=0.27$, and $\\Lambda = 0.73$. ",
        "conclusions": "The X-ray luminosity and temperature of RXJ1648 ($L_{X,bol}(r<r_{500}) = 3.2^{+1.7}_{-0.8} \\times 10^{43}$ ergs s$^{-1}$, $kT = 2.1^{+0.8}_{-0.6}$ keV) indicate that that it is a massive, collapsed group.  In addition, the average group properties agree well with the intermediate-redshift group/poor cluster scaling relations.  However, the X-ray emission is very elliptical/elongated, and RXJ1648 lacks a dominant BCG at the X-ray center, indicating that this group is still in the process of formation.  Instead, RXJ1648 contains a string of seven galaxies, five of which are brighter than $L_{\\ast}$, within a projected radius of 200 kpc (excluding one galaxy with a velocity offset $>900$ km s$^{-1}$).  The X-ray emission and its elongation trace this central string of galaxies.  These central galaxies have a smaller velocity dispersion, are on average brighter, and have less star formation (lower [OII] and $H_{\\delta}$) than other group galaxies.  The total group velocity dispersion of $417^{+118}_{-86}$ km s$^{-1}$ also indicates a massive group and matches well the X-ray luminosity and temperature; however, the velocity distribution indicates a double Gaussian or large tails, with one system having $\\sigma \\sim 200$ km s$^{-1}$ (dominated by the central string) within a larger system.  Overall, the observations support the picture that RXJ1648 is a recently formed massive group.  Several of the bright galaxies appear to have sunk towards the center of the system, but have not yet merged to form a dominant, central BCG.  The groups in our X-ray selected, intermediate-redshift sample represent an evolutionary sequence in BCG formation, with RXJ1648 representing one of the youngest systems before a dominant BCG has formed.  Two other RDCS groups, both at $z=0.23$, contain multi-component dominant central galaxies with at least three nuclei, indicating a past merger, and round, relaxed X-ray morphologies (Paper III; Paper I).  These groups represent what RXJ1648 might look like in a few Gyrs.  The formation of a central, dominant BCG through mergers of bright group galaxies appears to occur after group collapse and the appearance of luminous X-ray emission.  Below we discuss more quantitatively the prospects for the formation of a central BCG in RXJ1648.  \\subsection{ BCG Formation }  RXJ1648 has a total mass of $M_{200} = 4^{+3}_{-2} \\times 10^{14} M_{\\odot}$, where the mass is estimated from the X-ray temperature and the surface brightness profile assuming hydrostatic equilibrium.  For a group of this mass, we expect a BCG with an absolute R-band magnitude between $-23.5$ and $-24$ (Brough et al. 2008; Hansen et al. 2008).  The brightest observed galaxy in RXJ1648 is about a magnitude fainter than this prediction with $M_R = -22.8$.  Will the central string of galaxies in RXJ1648 merge, and if so will the resulting galaxy have a luminosity consistent with the total group mass?  First, we consider the time it will take for the galaxies in the string to sink to the group center through dynamical friction.  Using the Faber-Jackson relation (Bernardi et al. 2005), we estimate that the five super-$L_{\\ast}$ galaxies within $r<200$ kpc have stellar velocity dispersion between 265 and 350 km s$^{-1}$.  Following the methodology in Zabludoff \\& Mulchaey (1998), we find that these galaxies have predicted tidal radii of $\\sim 80$ kpc ($r_T \\approx r_c \\sigma_{gal}/2\\sigma_{grp}$ where $r_c$ is the group core radius and $\\sigma_{gal}$ and $\\sigma_{grp}$ are the galaxy stellar velocity dispersion and the group velocity dispersion, respectively; Merritt 1984) and masses between $6 \\times 10^{11} M_{\\odot}$ and $14 \\times 10^{11} M_{\\odot}$ ($M \\approx r_T\\sigma_{gal}^2/2G$).  If we instead conservatively assume $M/L \\approx 5 M_{\\odot}/L_{\\odot}$, we find galaxy masses about a factor of three lower, but these lead to a similar range of dynamical friction times.  We calculate the time it will take for each galaxy to fall to the center of the group through dynamical friction (eq.[7-18] and eq.[7-13b] in Binny \\& Tremaine 1987) approximating the group as a singular isothermal sphere with the total group velocity dispersion (eq.[4-123] in Binny \\& Tremaine 1987) and assuming that the galaxies in the string are currently orbiting with a velocity of $\\sim 200$ km s$^{-1}$ at their projected distances.  We find that the five super-$L_{\\ast}$ galaxies will reach the center in between 1-5 Gyrs (regardless of our definition of galaxy mass).  The group redshift corresponds to 4.1 Gyrs ago, in which time four of the galaxies should merge.  The two sub-$L_{\\ast}$ galaxies in the string will not reach the center for $\\sim 9$ Gyrs.  These calculations represent a fairly simple estimate of the dynamical friction timescale ignoring, for example, the gravitational affect of the galaxies on each other, but they do show that it is likely that some of these galaxies will merge by the present epoch.  If the five bright string galaxies merged and no stars were lost in the merging process the resulting BCG would have $M_R = -24.23$, a bit brighter then predicted for a group with the mass of RXJ1648.  If only the four brightest galaxies merge, the BCG magnitude would still be $M_R = -24.06$.  Recently, various groups have argued based on comparisons between simulations and observations that a significant fraction of stars (30-80\\%) in galaxy mergers within clusters are lost to the intracluster light (ICL) (Conroy, Wechsler, \\& Kravtsov 2007; Monaco et al.~2006).  Assuming the observed $L_{BCG}-M_{200}$ relation (Brough et al. 2008; Hansen et al. 2008), we similarly find that 20-50\\% of the stars should be lost in the mergers of the five bright string galaxies in order to produce a BCG with the right luminosity for the group mass.  However, if RXJ1648's mass grows with time or fewer galaxies merge, than no loss of stars to the ICL is necessary.  We note that if the bright, central galaxies in RXJ1648 merge, these mergers are expected to be dry.  Of the five super-$L_{\\ast}$ galaxies in the central region all are elliptical or S0 galaxies with no detected [OII] emission except for the post-starburst, Sb galaxy noted in \\S3.2.  As noted in Paper III, one of these galaxies is in fact a merging system composed of two elliptical galaxies."
    },
    "0806/0806.0012_arXiv.txt": {
        "abstract": "\\noindent  Spectral indices are useful tools for quantifying the strengths of features in moderate-resolution spectra and relating them to intrinsic stellar parameters. This paper focuses on the 4300$\\mbox{\\AA}$ CH G-band, a classic example of a feature interpreted through use of spectral indices. G-band index definitions, as applied to globular clusters of different metallicity, abound in the literature,  and transformations between the various systems, or comparisons between different authors' work, are difficult and not always useful. We present a method for formulating an optimized G-band index, using a large grid of synthetic spectra. To make our new index a reliable measure of carbon abundance, we minimize its dependence on [N/Fe] and simultaneously maximize its sensitivity to [C/Fe]. We present a definition for the new index $S_{2}(CH)$, along with estimates of the errors inherent in using it for [C/Fe] determination, and conclude that it is valid for use with spectra of bright globular cluster red giants over a large range in [Fe/H], [C/Fe], and [N/Fe]. ",
        "introduction": "Spectroscopic indices are a measure of the strength of an absorption or emission feature, usually reported as the magnitude difference between the integrated flux in the region of the feature in question (the ''science band'') and one or two nearby continuum regions (the ''comparison bands''). Indices and narrow-band photometry have a long history of usefulness in terms of physical information returned per time spent observing. They can be used to study a wide range of stellar properties: Ca II H\\&K surveys quickly identify extremely metal-poor stars (e.g., Beers et al 1985, 1999\\nocite{B85}\\nocite{B99}), the Mg I index is an indicator of surface gravity in cool stars (e.g., Morrison et al. 2000\\nocite{M00}), and Balmer indices are diagnostics of stellar age and metallicity when seen in galaxy spectra (e.g., Kauffmann et al. 2003\\nocite{K03}). Because they require far less observing time than high-resolution spectroscopy, the telescope-limited astronomer frequently uses these low-resolution methods in a first pass to identify interesting objects to be followed up later in more detail.  Narrow-band colors are often used to calculate physical quantities: an empirical calibration can be established between the observed quantities and the intrinsic properties in a well-studied sample, and that calibration can then be applied to derive the intrinsic properties of a much larger set of stars, as long as the stars in the larger set are similar enough in temperature, gravity, and composition to the well-studied calibration sample. The \\citet{A99} photometric temperature calibration is a well-known example of this process. The observables can also be used directly as proxies; for example, DDO and Str\\\"{o}mgren colors, which are known to correlate with temperature, metallicity, and surface gravity, can be used for dwarf/giant separation, or to do rough sorting by metallicity (e.g., Grundahl et al. 2002\\nocite{GBN02}). The so-called Lick indices can be measured from integrated spectra of galaxies, and provide information about unresolved stellar populations (e.g., Worthey 2004\\nocite{W04}).  Molecular bandstrength indices measured from low- or moderate-resolution spectroscopy are often the end result themselves: comparing CN versus CH bandstrengths of globular cluster red giants, one finds that CH bandstrength falls with rising luminosity \\citep{C82} (the same is true in the halo field: see Gratton et al. 2000\\nocite{GSCB00}), while at a given luminosity the CH-weaker (and therefore lower carbon) stars tend to have the stronger CN bands, and therefore must be nitrogen-enhanced. These trends imply that red giant atmospheres must be experiencing ongoing CN(O)-cycle processing (e.g., McClure \\& Norris 1977 \\nocite{MN77}, Shetrone et al. 1999\\nocite{S99}). Much research into the behavior of C and N elemental abundances in globular cluster red giants (e.g., Briley \\& Cohen 2001\\nocite{BC01}, Norris et al. 1984\\nocite{N84}, Norris \\& Zinn 1977\\nocite{NZ77}, Zinn 1973\\nocite{Z73}) has relied heavily on the use of the 4300 $\\mbox{\\AA}$ CH bandstrength for [C/Fe] determination. This is true of Population II field giants as well (e.g., Rossi et al. 2005\\nocite{R05}). However, there has been little uniformity in the CH indices used by various researchers. Some G-band indices that have been used to date are $S(CH)$ \\citep{MSB08}: \\begin{displaymath} S(CH) = -2.5 \\log \\frac{\\int_{4280}^{4320} I_{\\lambda}d\\lambda}{\\int_{4050}^{4100} I_{\\lambda}d\\lambda + \\int_{4330}^{4350} I_{\\lambda}d\\lambda } \\end{displaymath} which was defined for use with the low-metallicity globular cluster M53  ([Fe/H]=-1.84);  \\noindent $s_{CH}$ \\citep{BS93}: \\begin{displaymath} s_{CH} = -2.5 \\log \\frac{\\int_{4280}^{4320} I_{\\lambda}d\\lambda}{\\int_{4220}^{4280} I_{\\lambda}d\\lambda} \\end{displaymath} which was defined for use with the moderate-metallicity globular cluster M13  ([Fe/H]=-1.54);  \\noindent $m_{CH}$ \\citep{T83}: \\begin{displaymath} m_{CH} = -2.5 \\log \\frac{(\\frac{1}{50\\mbox{\\AA}}){\\int_{4270}^{4320}F_{\\nu} d\\lambda}}{(\\frac{1}{110\\mbox{\\AA}}){\\int_{4020}^{4130}F_{\\nu} d\\lambda}+(\\frac{1}{90\\mbox{\\AA}}){\\int_{4430}^{4520}F_{\\nu} d\\lambda}} \\end{displaymath} which was defined for use with red giants in metal-poor globular clusters M92 and M15 ([Fe/H] $\\simeq -2.3$);  \\noindent $GP$ \\citep{R05}: \\begin{displaymath} GP = -2.5 \\log \\frac{\\int_{4297.5}^{4312.5}I_{\\lambda}}{\\int_{4247}^{4267}I_{\\lambda}d\\lambda + \\int_{4363}^{4372}I_{\\lambda}d\\lambda} \\end{displaymath} defined for a study of carbon-enhanced metal-poor field stars;  \\noindent and $CH(4300)$ \\citep{H03}: \\begin{displaymath} CH(4300) = -2.5 \\log \\frac{\\int_{4280}^{4320} I_{\\lambda}d\\lambda}{\\int_{4250}^{4280} I_{\\lambda}d\\lambda + \\int_{4320}^{4340} I_{\\lambda}d\\lambda } \\end{displaymath} which was defined for use with main-sequence stars of the globular cluster 47 Tucanae. In the above equations, $I_{\\lambda}$ is the flux per Angstrom, recorded in data such as \\citet{MSB08} in units of ADU per pixel, and $F_{\\nu}$ is the flux per unit frequency, calibrated to true energy units in the study of \\citet{T83}. ",
        "conclusions": "It is our hope that this new G-band index can provide a standard tool for low-resolution spectroscopic studies of bright red giants in globular clusters. It is simultaneously more sensitive to [C/Fe] and less sensitive to [N/Fe] than similar indices used in the literature to date, and would allow for useful comparison and combination of independent researchers' data in multiple globular clusters. The optimization method used here could be generally applicable to any spectral region where there are a few main parameters controlling the strength of an absorption or emission feature. Our characterizations of the dependence of $S_{2}(CH)$ on [Fe/H], [C/Fe], and [N/Fe], as well as the error associated with converting $S_{2}(CH)$ to [C/Fe], apply over a wide abundance range, but they are only strictly true for stars with $M_{V}=-1.5$, since the model atmospheres and synthetic spectra were only calculated for that particular value. While our current model calibration of $S_{2}(CH)$ should lead to a reliable measure of [C/Fe] for a moderate range in $M_{V}$, perhaps $-1.0 \\geq M_{V} \\geq -2.0$, it should not be applied to stars outside this magnitude range without recalibration. Since the index was designed specifically for bright red giants of Population II metallicities, we do not recommend using it to study [C/Fe] in stars belonging to significantly different populations.     \\newpage"
    },
    "0806/0806.1014_arXiv.txt": {
        "abstract": "{As part of an ongoing series of deep GMRT surveys we have observed the Spitzer extragalactic First Look Survey field, producing the deepest wide-field 610-MHz survey published to date.  We reach an rms noise level of 30~$\\mu$Jy~beam$^{-1}$ before primary beam correction, with a resolution of $\\sim$6~arcsec over an area of $\\sim$4~deg$^{2}$.  By combining these observations with the existing 1.4-GHz VLA survey produced by Condon et al.\\ (2003), along with infrared data in up to seven wavebands from the Spitzer Space Telescope, optical photometry from SDSS and a range of spectroscopic redshift surveys, we are able to study the relationship between radio luminosity and star formation rate in star-forming galaxies up to $z \\sim 1$.  The large amount of multi-wavelength data available allows $k$-corrections to be performed in the radio due to the knowledge of the radio spectral index, and in the infrared through the use of a semi-empirical radiative transfer model which models star-forming regions, warm dust surrounding these regions, and diffuse interstellar dust, taking into account the star formation rate, star formation history and hydrogen column density within each galaxy.  A strong correlation is seen between radio luminosity and the infrared-derived star formation rates, which is best fit by a slightly non-linear power-law.  We look for cosmic evolution in the comparative radio brightness of star-forming galaxies by searching for deviations away from the global relationship.  Any such deviation would indicate a systematic variation in one or more of the properties controlling synchrotron radiation, in particular an increase in the magnetic field strengths of star-forming galaxies over time.  The data shows no evidence for such an effect, suggesting that there has been little evolution in the magnetic fields of galaxies since $z\\sim1$.}  \\FullConference{From Planets to Dark Energy: the Modern Radio Universe\\\\ October 1-5 2007\\\\ The University of Manchester, UK}  \\begin{document} ",
        "introduction": " ",
        "conclusions": "The large amount of multi-wavelength information available on regions such as the Spitzer extragalactic First Look Survey field mean that it is possible to create large samples of star-forming galaxies which contain a significant amount of spectral information.  This allows individual $k$-corrections to be performed, using the knowledge of the optical stellar emission, infrared thermal dust emission and radio synchrotron emission from each galaxy in the sample.  We find the relationship between radio luminosity and star formation rate has a slightly non-linear form with power law index of $1.13\\pm0.07$.  By comparing the radio luminosity and infrared-derived star formation rates of galaxies between $z \\sim 1$ and the present day, we find no evidence for significant evolution in the radio luminosity of star-forming galaxies.  Any systematic variation over time in the radio luminosity of galaxies with particular SFR would imply an evolution in one or more of the properties that controls synchrotron radiation, and lead to the tight infrared / radio correlation\\cite{Appleton04}.  In particular, an increase in the magnetic field strength of galaxies since redshift 1 ($\\sim$7.7~Gyr) would lead to greater synchrotron emission from the electron population around supernova remnants and a greater chance of electron `confinement' -- electrons radiating away all of their energy within a galaxy, rather than escaping.  This would lead to an increase in the radio luminosity for sources with given SFR at lower redshifts. No such effect is seen, suggesting that there has been little evolution in the magnetic fields of galaxies over this redshift range."
    },
    "0806/0806.3011_arXiv.txt": {
        "abstract": "In this paper we analyze the impact on the orbital motions of the outer planets of the solar system from Jupiter to Pluto of some velocity-dependent forces recently proposed to phenomenologically explain the Pioneer anomaly,  and compare their predictions (secular variations of the longitude of perihelion $\\varpi$ or of the semimajor axis $a$ and the eccentricity $e$) with the latest observational determinations by E.V. Pitjeva with the EPM2006 ephemerides. It turns out that while the predicted centennial shifts of $a$ are so huge that they would have been easily detected for all planets with the exception of Neptune, the predicted anomalous precessions of $\\varpi$  are too small, with the exception of Jupiter, so that they are still compatible with the estimated corrections to the standard Newton-Einstein perihelion precessions.  As a consequence, we incline to discard those extra-forces predicting secular variations of $a$ and $e$, also for some other reasons, and to give a chance, at least observationally, to those models predicting still undetectable perihelion precessions. Of course, adequate theoretical foundations for them should be found. ",
        "introduction": "The Pioneer anomaly\\cite{Nie06} (PA) consists of an unmodelled almost constant and uniform acceleration approximately directed towards the Sun of magnitude \\eqi\\apio = (8.74\\pm 1.33)\\times 10^{-10}\\ \\msm 2\\lb{pioacc}\\eqf detected\\cite{And98,And02,mark,Ol07} in the radiometric data from the Pioneer 10 (launched in March 1972) and Pioneer 11 (launched in April 1973) spacecraft after they passed the $\\approx$ 20 AU threshold moving with speed $v_{\\rm Pio} \\approx 1.2 \\times 10^4\\ \\msm 1$ along roughly antiparallel escape hyperbolic paths undertaken after their previous encounters with Jupiter ($\\approx 5$ AU) and Saturn ($\\approx 10$ AU), respectively. Concerning the possibility that PA started to manifest itself at shorter heliocentric distances\\cite{Nie05,Nie08}, efforts to retrieve and analyze early data from Pioneer 10/11 are currently being performed\\cite{Toth08,PioZARM08}.   The Pioneer spacecraft were particularly well suited for radioscience celestial mechanics experiments because they were spin-stabilized\\footnote{This was due to the fact that Pioneer 10/11 were equipped with Radioisotope Thermoelectric Generators (RTG) placed at the end of long booms to be away from the spacecraft and thereby avoid any radiation damage.}; in practice, they could be regarded as gyroscopes so that  only a few orientation maneuvers, easily  modeled, were needed every year to keep the antenna pointed towards the Earth.  On the contrary, 3-axis stabilized spacecraft like Voyager 1/2 undergo continuous, semi-autonomous, small gas jet thrusts to maintain the antenna facing the Earth; as a consequence, their navigation  is not as precise as that of the Pioneer 10/11.  The attempts performed so far to explain  PA in terms of known effects of gravitational\\cite{And02} and/or non-gravitational\\cite{Mur99,Kat99} origin were found to be not  satisfactory\\cite{And99a,And99b}, so that a vast number of  exotic explanations based on modified models of gravity were proposed (see, e.g., \\ctz{And02,Ber04,Dit05,Izz06}, and references therein). If PA is due to some modifications of the known laws of gravity, this must be due to a radial extra-force affecting the orbits of the planets as well, especially those moving in the region in which PA manifested itself in its presently known form. The impact of a Pioneer-like additional acceleration  on the motion of major and minor bodies in the outer regions of the solar system was  recently studied by numerous authors with different approaches\\cite{comets,Ior06,Ior07,Tan07,Wal07,Sta08}: it turned out that a constant and uniform extra-acceleration  with the magnitude of \\rfr{pioacc} would produce huge secular effects which are neatly absent in the planetary data.  It was recently suggested\\cite{Lamm07} that, from a purely phenomenological point of view, test bodies moving in the (outer) solar system could experience velocity-dependent extra-accelerations of the form \\eqi A_{v}=-|v_r| \\left(\\rp{\\apio}{v_{\\rm Pio}}\\right), A_{v}=-v_r \\left(\\rp{\\apio}{v_{\\rm Pio}}\\right),\\lb{extraa1}\\eqf and \\eqi\\ A_{v^2}=-v_r^2 \\left(\\rp{\\apio}{v^2_{\\rm Pio}}\\right),\\ A_{v^2}=-|v_r|v_r \\left(\\rp{\\apio}{v^2_{\\rm Pio}}\\right)\\lb{extraa}\\eqf where $v_r$ is the radial component of the test particle's velocity $\\bds v$; \\rfr{extraa1} and \\rfr{extraa} would reduce to \\rfr{pioacc} for the Pioneer 10/11 spacecraft whose  velocities  can be assumed entirely radial in the outer regions of the solar system in which PA was detected.  Standish in \\ctz{Sta08} put on the test such a hypothesis by fitting huge planetary data sets with the dynamical force models of the latest Jet Propulsion Laboratory (JPL) DE ephemerides modified $ad\\ hoc$ according to \\rfr{extraa1} and \\rfr{extraa}  and examining the results in terms, e.g, of the reliability of the estimated parameters. His conclusion was that the existence of extra-accelerations like those of \\rfr{extraa1} and \\rfr{extraa} at heliocentric distances $\\gtrsim 20$ AU cannot be ruled out by the present-day available data of the outer planets because \\rfr{extraa1} and, especially, \\rfr{extraa} would induce orbital effects on them too small to be detected. Their existence in the inner regions of solar system is, instead, ruled out.  In this paper we will follow a different approach by using the EPM2006 ephemerides produced by E.V. Pitjeva\\cite{Pit08} at the Institute of Applied Astronomy (IAA) of the Russian Academy of Sciences (RAS). First, we will  analytically work out the secular effects  of small perturbing accelerations like those of \\rfr{extraa1} and \\rfr{extraa} on the Keplerian orbital elements of a planet in order to gain as clear as possible insights about the modifications which the orbits would undergo if \\rfr{extraa1} and \\rfr{extraa} were real; should some implausible physical feature turn out, it would be more difficult to trust such proposed anomalous forces. Then, we will compare some of such predictions with the latest observational determinations for the outer planets estimated by Pitjeva with the EPM2006 ephemerides in a purely phemomenological way as corrections to the known effects due to usual Newton-Einstein laws, without modelling any additional force. In Table \\ref{Pittable} we quote some quantities we will use. They are the outcome of a global fit of more than 400,000 data points (1913-2006) performed by Pitjeva\\cite{Pit07,Pit08} with the EPM2006 ephemerides; about 230 parameters were estimated. \\begin{table} \\tbl{ Second column: formal standard deviations $\\delta a$, in m, of the semimajor axes of the outer planets from a fit of 400,000 data points spanning almost one century with the EPM2006 ephemerides (from Table 3 of Ref.~24). Third column: corrections to the standard Newton-Einstein secular precessions of the perihelia$^{25}$ in arcsec cy$^{-1}$. Fourth column:  their formal errors, in arcsec cy$^{-1}$, re-scaled by a factor 10.  For Neptune and Pluto no secular precessions have been estimated because the available modern data records for them do not yet cover an entire orbital revolution. } {\\begin{tabular}{@{}cccc@{}} \\toprule Planet & $\\delta a$ (m)  & $\\Delta\\dot\\varpi$ (arcsec cy$^{-1}$)& $\\delta\\Delta\\dot\\varpi$ (arcsec cy$^{-1}$) \\\\ \\colrule Jupiter &  615 & 0.0062 &  0.036\\\\ Saturn & 4,256 & -0.92 & 2.9 \\\\ Uranus & 40,294 & 0.57 & 13.0\\\\ Neptune & 463,307 & N.A. & N.A.\\\\ Pluto & 3,412,734 & N.A. & N.A.\\\\ \\botrule \\end{tabular}\\label{Pittable}} \\end{table} % It must be noted that the uncertainties $\\delta\\Delta\\dot\\varpi$ in the estimated corrections to the perihelion precessions are the formal ones re-scaled by a factor 10 in order to obtain realistic evaluations for them. ",
        "conclusions": "An   ingenious attempt recently proposed to explain the Pioneer anomaly as due to a modification of the usual Newton-Einstein laws of gravitation consists in postulating the existence of some velocity-dependent extra-forces linear or quadratic in the radial component $v_r$ of the velocity of a test body. We put on the test such empirical models in the outer regions of the solar system in which the Pioneer anomaly manifested itself in its presently known form  with the latest observational determinations of the planetary motions obtained by E.V. Pitjeva with the EPM2006 ephemerides. It turns out that the models yielding anomalous perihelion precessions cannot yet be ruled out, at least phenomenologically, for heliocentric distances larger than 5 AU. On the contrary, the models predicting secular variations of the semimajor axis $a$ and the eccentricity $e$ are much more difficult to be trusted not only because they would violate the conservation of energy but also because the centennial shifts for $a$ predicted by them are so large that they should have been  detected, given the present-day accuracy in determining such orbital element. However, it must be considered that sound theoretical justifications for such models must be given."
    },
    "0806/0806.1222_arXiv.txt": {
        "abstract": "A new method of determining galaxy star-formation histories (SFHs) is presented. Using the method, the feasibility of recovering SFHs with multi-band photometry is investigated.  The method divides a galaxy's history into discrete time intervals and reconstructs the average rate of star formation in each interval. This directly gives the total stellar mass.  A simple linear inversion solves the problem of finding the most likely discretised SFH for a given set of galaxy parameters. It is shown how formulating the method within a Bayesian framework lets the data simultaneously select the optimal regularisation strength and the most appropriate number of discrete time intervals for the reconstructed SFH. The method is demonstrated by applying it to mono-metallic synthetic photometric catalogues created with different input SFHs, assessing how the accuracy of the recovered SFHs and stellar masses depend on the photometric passband set, signal-to-noise and redshift. The results show that reconstruction of SFHs using multi-band photometry is possible, being able to distinguish an early burst of star formation from a late one, provided an appropriate passband set is used. Although the resolution of the recovered SFHs is on average inferior compared to what can be achieved with spectroscopic data, the multi-band approach can process a significantly larger number of galaxies per unit exposure time. ",
        "introduction": "\\label{sec_intro}  A significant step towards understanding how galaxies form and evolve can be made by measuring the variation in their star formation rate (SFR) with age.  Imprinted in every galaxy's integrated light is a record of its entire life from birth, through passive evolution, possible merging and recycling of material, up to the epoch at which it is observed. Star formation histories (SFHs) therefore play a crucial role in the quest for a complete and accurate model of the formation of stellar mass in the Universe and how distant systems relate to those locally.  Characterising galaxy SFHs has been a subject of much interest for several decades, with studies attempting to achieve this aim through a variety of different means. Approaches can be broadly divided into those using multi-band photometry and those using spectra. Recently the practice has seen a significant revival thanks to improvements in stellar synthesis modelling and the advent of large datasets such as the Sloan Digital Sky Survey \\citep[SDSS;][]{stoughton02}. Many new spectroscopic techniques have been developed \\citep[e.g.,][]{heavens00, vergely02,cid04,cid05,nolan06,ocvirk06,chil07,tojeiro07} and in their various forms, these have seen application to several sets of real data \\citep[e.g.,][] {reichardt01,panter03,heavens04,panter04,sheth06, cid07,nolan07,panter07,koleva08}. Similarly, there have been numerous recent studies conducted using multi-band photometry \\citep[e.g.,][]{borch06,schawinski07,salim07,noekse07,kaviraj07} including \\citet{kauffmann03} who combined multi-band photometry with measurements of the H$\\delta$ absorption line and 4000\\AA break strength.  In a similar vein to spectroscopic versus photometric redshift estimation, SFHs determined from spectra tend to have greater precision per galaxy, whereas those derived from multi-band photometry allow many more objects to be studied in the same amount of observing time but with a compromise in SFH resolution.  The method adopted by existing multi-band studies is to assume a parametric model for the SFH. The parameters are adjusted to find the set of model fluxes, computed from a spectral library of choice, that best matches the set of observed fluxes. This not only forces the SFH to adhere to a potentially unrepresentative prescribed form, it also necessitates a fully non-linear minimisation over all parameters.  In contrast, the majority of the recent spectroscopic methods divide up a galaxy's history into several independent time intervals and reconstruct the average SFR in each interval to give a discretised SFH. The advantage this brings, as shown in Section \\ref{sec_most_prob_SFH}, is that finding the best-fit SFR in every interval for a fixed set of galaxy parameters (such as redshift, extinction and metallicity) is a linear problem. The inefficient non-linear SFH minimisation with its risk of becoming trapped in local minima is therefore replaced with a simple matrix inversion guaranteeing that the global minimum for the fixed set of galaxy parameters is found.  The prescribed SFH models used by the multi-band methods are mainly driven by the small number of passbands used in many multi-band campaigns. With only a small number of passbands, the ability to constrain a galaxy's SFH is limited and a model SFH with only one or two parameters must be used. However, modern surveys are being carried out in many more passbands and over larger wavelength ranges than ever before \\citep[for example, the COMBO-17 survey of][]{wolf01}. Given these recent improvements, the possibility of recovering discretised SFHs from multi-band photometry alone is now worthy of investigation.  The purpose of this paper is twofold. Firstly, a new SFH reconstruction method that recovers discretised SFHs is presented. It is shown how the Bayesian evidence can be used to simultaneously establish the most appropriate number of discrete SFH time intervals and the optimal strength with which the solution should be regularised. The formalism is completely general and can be applied to spectra just as easily as multi-band photometry as well as a combination of both. The Bayesian evidence gives a more natural and simplified alternative to existing procedures for determining the optimal number of SFH intervals and for determining the correct level of regularisation.  Secondly, this paper presents results of an investigation into the feasibility of using the new method with multi-band photometry alone. By applying the method to synthetic galaxy catalogues created with different input SFHs and filtersets, the accuracy of the recovered discretised SFHs is demonstrated. This study focuses in particular on the dependence of the reconstruction on galaxy redshift, photometric signal-to-noise (S/N), the wavelength range spanned by the passbands, the number of passbands and the presence/absence of a new and/or old stellar population.  The layout of the paper is as follows. In Section \\ref{sec_method} the SFH reconstruction method is described. Section \\ref{sec_synthetic_cats} gives details of how the synthetic catalogues are generated. The method is applied to these catalogues in Section \\ref{sec_sims} to assess its performance.  Section \\ref{sec_summary} gives a summary of the findings of this paper to act as recommendations for applying the method to real data.  Throughout this paper, the following cosmological parameters are assumed; ${\\rm H}_0=100\\,{\\rm h}_0=70\\,{\\rm km\\,s}^{-1}\\,{\\rm Mpc}^{-1}$, $\\Omega_m=0.3$, $\\Omega_{\\Lambda}=0.7$. All magnitudes are expressed in the AB system. ",
        "conclusions": "\\label{sec_summary}  The primary aim of this study has been to assess reconstruction of discretised SFHs using a new method applied to multi-band photometric data. Although not tested in this paper, the method can also be applied to spectroscopic data as well as a mixture of both spectroscopic and multi-band data.  The method differs from existing methods by maximising the Bayesian evidence instead of minimising $\\chi^2$ (or maximising the posterior probability).  For regularised solutions, the evidence gives the unbiased relative probability of the fit between different model parameterisations. This is unlike the $\\chi^2$ statistic which suffers from an ambiguous number of degrees of freedom that changes between parameterisations when regularisation is applied.  This work has demonstrated that the evidence allows the data to correctly and simultaneously set the optimal regularisation strength and the appropriate number of blocks in the reconstructed SFH. Although negative SFRs can arise, the optimal level of regularisation ensures that the fraction of such cases is low. Negative SFRs are limited mainly to galaxies with low photometric S/N and inadequate filter sets (e.g., the optical set considered in this work).  Provided the filter set is adequate, negative SFRs are always consistent with a null SFR. This approach may be preferable to schemes that enforce positive SFRs. Enforcing positivity not only risks artificial ringing and biasing in the reconstructed SFH, it also hides problems that give rise to negative SFRs.  Application of the method to a range of synthetic galaxy catalogues generated with varying passband sets and SFHs demonstrates that use of multi-band data in constraining SFHs is feasible along with certain caveats. The scatter seen in the SFHs reconstructed in this work shows that occasional significant inaccuracies can occur even with a comprehensive filterset that extends up to near-IR and mid-IR wavelengths. Therefore, interpretation of SFHs recovered from solely multi-band photometry on a galaxy by galaxy basis should be conducted with some caution. The mean SFH of a sample of galaxies is therefore a more reliable quantity in order to average out uncertainties although this study indicates that averaging over only four galaxies readily allows a late burst to be distinguished from an early burst.  In comparison, studies using spectroscopic data show that reliable SFHs can be derived for individual galaxies. Nevertheless, multi-band photometry allows reconstruction of SFHs for many times more galaxies than spectroscopic methods for the same amount of observing time.  The most important factor governing the accuracy of the reconstructed SFHs is the wavelength range spanned by the filterset. The results show little difference between two filtersets that span approximately the same wavelength range (optical to mid-IR) despite one set having half the number of filters of the other. Conversely, SFHs based on only purely optical photometry are completely unreliable, it being impossible to distinguish any of the input SFHs investigated.  A filterset consisting of only near and mid IR filters ($Z'$ -- 4.5$\\mu$m) allows recovery of SFHs to within a comparable accuracy to that recovered when optical filters are also included, implying that the majority of the SFH constraints are provided by near and mid-IR data (for the SFHs tested here).  In terms of the ability of multi-band photometry to constrain different SFH types, the results show that apart from the case where only optical filters are used, early bursts of star formation can be differentiated from late bursts and both of these can be distinguished from dual bursts and constant SFRs.  However, early bouts of star formation activity are always artificially smeared to later times in the reconstructed SFH compared to the input SFH. These findings apply specifically to the SFHs considered in this work, where the early burst gives rise to bolometric luminosity that is one tenth that of the late burst. A quick test has revealed that a stronger early burst is more accurately recovered with less smearing to late times. In addition, although the dual burst SFH used here suggests that recovery of more than two bursts would be unfeasible with the filtersets tested, bursts with more similar bolometric luminosities can be more readily recovered.  This was demonstrated by \\citet{ocvirk06} who showed that CSP SEDs constructed from flux normalised bursts allow a higher SFH resolution on average than SEDs constructed from mass normalised bursts.  The results presented in this paper have been obtained using the \\citet{bruzual03} spectral libraries. Whilst the exact values of the numerical results quoted here will depend on the specific SED library of choice, there are no compelling reasons to suggest that the observed trends would not remain valid generally.  This study has considered a specific case where galaxy redshift and extinction are known prior to reconstructing the SFH. Also, mono-metallic stellar populations have been assumed where the metallicity does not evolve as the galaxy ages. Clearly, the more general problem necessitates maximising the evidence over extra parameters. The expected effect of this is that the maximum evidence would shift to lower SFH resolutions on average. Although generalising to a variable redshift and extinction is a relatively small expansion of the non-linear parameter space, incorporating a time-varying metallicity in addition results in a significantly larger and more complex non-linear parameter space. This increases the time required to locate the maximum evidence and increases the risk of becoming trapped at a local maximum.  However, there are two small reprieves.  The first is that the metallicity history can be regularised in a similar manner to the SFH, smoothing the evidence surface and therefore easing maximisation. The second exploits SED libraries with discrete metallicities. As shown by \\cite{tojeiro07}, finding the optimal metallicity within the range spanned by two tabulated values of metallicity is also a linear problem which can be directly combined with the linear inversion of the SFH. In this way, optimising the metallicity for each SFH block reduces to searching a smaller number of discrete values. A full investigation of the general case will be presented in forthcoming work.   \\appendix"
    },
    "0806/0806.4541_arXiv.txt": {
        "abstract": "We consider the examples of single-field inflation models predicting large amplitudes of the curvature perturbation power spectrum at relatively small scales. It is shown that in models with an inflationary potential of double-well type the peaks in the power spectrum, having, in maximum, the amplitude ${\\cal P}_{\\cal R} \\sim 0.1$, can exist (if parameters of the potential are chosen appropriately). It is shown also that the spectrum amplitude of the same magnitude (at large $k$ values) is predicted in the model with the running mass potential, if the positive running, $n'$, exists and is about $0.005$ at cosmological scales. Estimates of the quantum diffusion effects during inflation in models with the running mass potential are given. ",
        "introduction": "In last few months several papers appeared \\cite{Kohri:2007gq, Kohri:2007qn, Saito:2008em, Peiris:2008be}, in which single-field inflation models predicting (potentially) large amplitudes of the curvature perturbations on relatively small scales are discussed. It is shown in \\cite{Kohri:2007gq} that large class of such models exists, namely, the models with a potential of hill-top type (the idea of the hill-top inflation was proposed, to author's knowledge, in the earlier work \\cite{Boubekeur:2005zm}). In such models, the potential can be of concave-downward form at cosmological scales (in accordance with data) and be much flatter at the end of inflation when small scales leave horizon. Correspondingly, the amplitude of the perturbation power spectrum can be rather large. It is noticed in \\cite{Kohri:2007gq} that the running mass model, having the potential with the similar behavior, also can predict the large spectrum amplitude.  Authors of \\cite{Kohri:2007qn} discuss also more general scenarios of producing large amplitudes of perturbation spectrum. They show the limitedness of the standard procedure of the potential reconstruction which can easily miss the potentials leading to large spectrum amplitude and to noticeable primordial black hole (PBH) production.  In the recent paper \\cite{Saito:2008em} it was shown that PBH production is possible in single-field models of two-stage type (\"chaotic $+$ new\"). The idea was proposed ten years ago in \\cite{Yokoyama:1998pt}. Authors of \\cite{Saito:2008em} carried out the numerical calculation of the power spectrum using the Coleman-Weinberg (CW) potential.  In the present paper we continue a study of the problems discussed in the previous works \\cite{Kohri:2007gq, Kohri:2007qn, Saito:2008em}. We investigated thoroughly, as a particular example, the model of two-stage inflation with a potential of the double-well (DW) form, and showed that the characteristic features of the power spectrum in models of this type (such as an amplitude and a position of the peak, a degree of tuning of parameters of the potential) are very sensitive to an exact form of the potential. Further, we carried out the numerical calculation of the power spectrum in a running mass model and showed that the spectrum amplitude at small scales can be rather large. Our calculation differs from the previous one \\cite{Leach:2000ea} in several aspects: we express the results through the values of parameters $s$, $c$, which are used nowadays and prove to be very convenient for a comparison with data; we studied, in details, the difference in predictions of slow-roll and numerical approaches at high $k$-values; we exactly specified the value of the positive running, $n'$, which corresponds to our spectrum prediction. In the final part of the work we investigated the quantum diffusion effects in a model with the running mass potential.  A plan of the paper is as follows. In the second section we study predictions of two-stage inflation models with DW and CW potentials, with accent on a mechanism of the formation of peaks in the power spectrum. In the Sec. \\ref{sec-RM} all aspects connected with an obtaining of the predictions of running mass inflation models are discussed. In the Sec. \\ref{sec-Concl} we present our main conclusions.        \\begin{figure} \\includegraphics[width=0.9\\columnwidth]{DW-bkg.eps} % \\center \\caption{ \\label{DWbkg} The solution of the background equation for inflation with the double-well potential (\\ref{DW}). The parameters of the potential are: $v = 0.16286748 m_{Pl}, \\lambda=1.7\\times 10^{-13}$.} % \\end{figure} \\begin{figure}[!t] \\includegraphics[width=0.9\\columnwidth]{epst.eps} \\includegraphics[width=0.9\\columnwidth]{epsdelt.eps} \\center % \\caption{ \\label{DWed} The time dependence of the parameter $\\epsilon$ and the combination $1+\\epsilon-\\eta$ corresponding to the background field evolution shown in fig. \\ref{DWbkg} } % \\end{figure} ",
        "conclusions": "}  {\\bf 1.} It is shown, by numerical methods, that in the single-field inflationary model with a simple double-well potential the parameter values of this potential can be chosen in such a way that the power spectrum of curvature perturbations ${\\cal P}_{\\cal R} (k)$ has a huge peak (with amplitude $\\sim 0.1$) at large $k$ (and the right normalization and monotonic behavior at cosmological scales). The peak arises due to temporary interruption of the slow-roll near points of the minimum of the potential, $\\pm v$. The corresponding mass of PBHs produced in early universe in a case of the realization of such power spectrum is about $10^7$g.  The analogous behavior of the power spectrum was obtained by authors of \\cite{Saito:2008em} in a model with CW potential. There are some important differences in the results of \\cite{Saito:2008em} and ours, in the peak amplitude and PBH mass, connected, in particular, with a large flatness near the origin in a case of the CW potential.  {\\bf 2.} It is shown that the inflation model with running mass potential predicts a rather large amplitude of the power spectrum of curvature perturbations ($\\sim 0.1$) at $k$-values $\\sim 10^{16}$ Mpc$^{-1}$. For such a prediction, a very small positive spectral index running at cosmological scales is necessary, $n'~ \\sim 0.005$, as well as a small negative value for the slow-roll parameter $\\eta$ ($\\approx -0.02$). Both this numbers do not contradict with data. It is shown also that for an obtaining the correct quantitative results for the power spectrum at largest $k$-values an use of numerical methods is required because, in general, slow-roll formulas are not accurate enough at the end of inflation, where $\\eta \\approx 1$.  {\\bf 3.} Quantum diffusion effects in a model with the running mass potential are studied in details. It is shown that inflationary evolution of the universe in a model with a scalar field and the running mass potential can be described by the classic deterministic equations, and for a possibility of such a description the correct choice of the initial conditions is crucial. Concretely, an initial value of the inflaton field (at the beginning of the evolution) should not be too close to a point of the maximum of the potential. If this condition is satisfied, the quantum corrections to a total e-fold number and to a position of the mean value of the probability distribution function are small."
    },
    "0806/0806.1478_arXiv.txt": {
        "abstract": "We report the discovery of a 7.3 M$_{\\rm J}$ exoplanet WASP-14b, one of the most massive transiting exoplanets observed to date. The planet orbits the tenth-magnitude F5V star USNO-B1 11118-0262485 with a period of 2.243752 days and orbital eccentricity  $e = 0.09$. A simultaneous fit of the transit light curve and radial velocity measurements yields a planetary mass of 7.3$\\pm$0.5 M$_{\\rm J}$ and a radius of 1.28$\\pm$0.08 R$_{\\rm J}$. This leads to a mean density of about 4.6 g\\,cm$^{-3}$ making it densest transiting exoplanets yet found at an orbital period less than 3~days. We estimate this system to be at a distance of $160\\pm20$ pc. Spectral analysis of the host star reveals a temperature of $6475\\pm100$ K, log g = 4.07 cm\\,s$^{-2}$ and $v\\sin i = 4.9\\pm1.0$ km\\,s$^{-1}$, and also a high lithium abundance, $\\log N({\\rm Li})$ = 2.84$\\pm$0.05. The stellar density, effective temperature and rotation rate suggest an age for the system of about 0.5--1.0 Gyr. ",
        "introduction": "The giant exoplanets that transit across the disks of their host stars are of great interest due to their impact on our understanding of planetary structure.  Since the first discovery of a transiting exoplanet HD209458b \\citep{charbonneau, henry}, more than 40 transiting systems have been found around nearby stars. Transit light curves, along with their radial velocity motions, provide a wealth of information about the system including precise mass, radius and mean density of the planet. This in turn allows us to probe their internal structure by comparing their physical parameters with models of planetary structure and evolution \\citep{guillot, fortney}. Given the importance of these systems, several wide-field surveys are in progress to find transiting exoplanets, e.g. HAT \\citep{hat}, XO \\citep{xo}, TrES \\citep{tres}, and SuperWASP \\citep{sw}.  In this paper, we report the 14th exoplanet discovered in the SuperWASP survey which is orbiting around a 10th-mag F5-type star in the Northern Hemisphere. The paper is organised as follows. In \\textsection 2, we give details of the discovery light curve and outcome of our follow-up photometric and spectroscopic observations. Spectral analysis to determine system parameters is described in \\textsection 3. The stellar and planetary evolutionary status is discussed in \\textsection 4. Finally, our results are briefly summarized in \\textsection 5. ",
        "conclusions": "We have discovered a new massive exoplanet WASP-14b orbiting with a period of about 2.244 days around its host main-sequence star GSC 01482--00882. Spectral analysis of the star implies a spectral type of F5V with solar metallicity. High precision photometric and spectroscopic follow-up observations reveal that the planet has a mass of 7.3$\\pm$0.5 M$_{\\rm J}$, radius of 1.28$\\pm$0.08 R$_{\\rm J}$ suggesting a mean density of 3.5$\\pm$0.6 $\\rho_{\\rm J}$ ($\\sim$ 4.6 g\\,cm$^{-3}$), and an eccentricity of 0.091$\\pm$0.003. The mean density of WASP-14b is high in comparison with a typical Hot Jupiter density of 0.34-1.34 g\\,cm$^{-3}$ \\citep{loeillet}, and similar to that of rocky planets and makes it densest transiting exoplanets so far discovered with $<$ 3 d orbital period. The planet is too massive to fit the mass-period relation given by \\citet{torres}, but its radius is consistent with the theoretical radius expected from the \\citet{fortney} model. Spectral analysis reveals that WASP-14 has a high lithium abundance, $\\log N({\\rm Li})$ = 2.84. This is consistent with its stellar temperature ($T_{\\rm eff}$ = 6475$\\pm$100 K) and age of about 0.5--1.0 Gyr which is further supported by the fitting of stellar evolutionary models of the \\citet{girardi}.  High orbital eccentricities in close-in planets are often explained by perturbations from a third body (Jackson et al. 2008 and references therein) or internal structures that result in their tidal dissipation factor being significantly larger that that commonly assumed $Q_p \\sim 10^6$ (Matsumura et al. 2008, Hebb et al. 2008). Although we do not see any systematic variation in RV residuals during our two months of spectroscopic observations of WASP-14, we expect that long term radial velocity monitoring will help constrain the nature of the third body in this systems hence make it an important target both for future transit-timing variation studies and for longer-term RV monitoring to establish the mass and period of the putative outer planet.  WASP-14b is one of the most massive transiting planets known along with HAT-P-2b \\citep{hat2,winn1, loeillet} and XO-3b \\citep{christopher,winn2,hebrard} and its physical characteristics closely resemble with those of HAT-P-2 except the latter has a much longer orbital period and smaller radius. However, there is still no firm explanation about the formation of such highly massive planets (e.g., H\\'{e}brard et. al~2008). A quite interesting feature is that all these three massive exoplanets has unusually large eccentric orbit for their short orbital period. Matsumura et al. (2008) argues that the new class of eccentric, short period transiting planets are still in the process of circularization and speculates that $Q_p$ could be as large as 10$^9$ for these planets.  The success of theoretical models of planetary structure depends heavily on our precise knowledge of the basic physical parameters of the planetary systems. Transit surveys in recent times have produced a large sample of transiting planets that show a remarkable diversity in their mass, radius and internal structure. While most of the Jupiter-mass transiting planets are well explained by existing models, planets which show excessive mass for their small radius or relatively low mass for their large radius are yet to be explained satisfactorily by any available model.  WASP-14b is one of the few massive exoplanets, some of which are even in closer orbits than most other hot Jupiters. It poses a great challenge for theoretical models to explain their internal structure, atmospheric dynamics and heat distribution."
    },
    "0806/0806.0297_arXiv.txt": {
        "abstract": "Contemporary imaging air Cherenkov telescopes (IACT) for ground-based very high energy (VHE) gamma-ray astronomy have prime focus optical design. Typically these telescopes have a $(2-4)^{\\circ}$ wide field of view (FoV). They use F/0.7-F/1.2 optics and provide $(3-10)^{'}$ resolution in the FoV. Generally, a well designed telescope that includes more than one optical element will offer some advantages not available in prime focus designs, such as a wider FoV, a more compact size, a higher and more homogeneous resolution and a lower degree of isochronous distortion of light rays focused onto the focal plane. Also, they allow monitoring the gamma-ray activity in a sizeable portion of the sky in a single observation. This would allow one to perform a sensitive all-sky survey in a relative short time. We present an F/0.8 $15^\\circ$ wide FoV telescope design, which provides a high and near uniform resolution and low isochronous distortion across the entire FoV. ",
        "introduction": "The technique of employing Imaging Air Cherenkov Telescopes (IACTs) has been successfully used for ground based gamma-ray astronomy for about two decades, revealing more than 70 sources in the course of this period. The successful race started with the discovery of the first TeV gamma-ray signal from the Crab Nebula - now regarded as the standard candle in gamma-ray astronomy - by the Whipple Collaboration in 1989 \\cite{whipple}. Over 20 years following that event, the sensitivity of the IACTs has improved dramatically, leading to a large increase in the rate of scientific discovery made with these instruments. In the beginning a few tens of hours were needed to detect a significant signal from the Crab Nebula, whereas today{'}s installations operating in the same energy range require only a few minutes for the same signal strength. The main improvements are essentially due to a) a finer pixel size of photo sensors in the imaging camera, b) an improved trigger, c) a larger size of telescopes and improved optics providing stronger signals and revealing more structures in the images (helps to further suppress the backgrounds) and d) the use of multiple telescopes operating in coincidence mode (the so-called stereo mode of observations). However, the field of view (FoV) of IACTs has not undergone a similar evolution. The largest FoV telescope had a field of about $7^\\circ$ \\cite{sinitsyna}. Contemporaneous IACTs typically cover a $(2-4)^{\\circ}$ wide FoV. A wider FoV would enable sensitive all-sky surveys to be conducted within a relatively short time frame. In optical astronomy the trend to larger FoV is evident, just to mention the LSST \\cite{lsst} and the Pan-starrs projects \\cite{panstarrs}. An interesting example of the use of even moderately wide FoV telescopes has been demonstrated by the H.E.S.S. collaboration. While performing observations of scheduled astronomical targets, H.E.S.S. has discovered several new sources in the $\\sim 4^\\circ$ effective FoV of their instrument. Subsequently, when scanning the galactic plane, multiple sources were found in the $\\sim 300$ square degree band surveyed by H.E.S.S. - i.e. the region scanned was much larger than a single H.E.S.S. FoV. \\cite{hess}.  Along with the advantages of the wide FoV there are also a number of drawbacks, such as: i) compared to currently used simple prime focus constructions, they have a more complex optical and mechanical design, ii) the imaging camera will have a large transverse size and thus can vignet a significant fraction of the mirror and iii) the imaging camera will be composed of a very large number of light sensors and one will therefore need a large number of readout channels. These factors tend to make a wide FoV telescope expensive. In the following we shall present a concept for a wide FoV IACT, for which the complications due to i) and ii) are minimal. The challenge of building a camera with a large number of channels cannot be by-passed. ",
        "conclusions": "We have presented the principal design of a 7~m wide FoV IACT, which has excellent imaging characteristics over a $15^\\circ$ field diameter. The basic design is that of a Schmidt type telescope, F/0.8, i.e., comparatively fast. This design allows one to obtain an optical spot size of $1^{'}$ RMS everywhere in the field, with an isochronous distortion below 0.03~ns  in case a Fresnel lens is used as a corrector. It is straightforward to scale this design to larger apertures. The only aspect which changes by scaling is the isochronous distortion. For a 20~m diameter telescope, the isochronous distortion amounts to 0.06~ns. The limiting factor in the baseline design proposed here is chromatic aberration in the corrector plate. This can be overcome by implementing an achromatic corrector plate. The main challenge will however lie in filling the focal plane with detectors which would fully utilize the resolution provided by the telescope."
    },
    "0806/0806.0318_arXiv.txt": {
        "abstract": "With the goal of investigating the nature of \\os\\ absorbers at high redshifts, we study the effects of proximity to the background quasar. In a sample of sixteen quasars at \\zqso\\ between 2.14 and 2.87 observed at high signal-to-noise and 6.6~\\kms\\ resolution with VLT/UVES, we detect \\nsys\\ \\os\\ absorption-line systems (comprising over 100 individual \\os\\ components) lying within 8\\,000~\\kms\\ of \\zqso. We present component fits to the \\os\\ absorption and the accompanying \\hi, \\cf, and \\nf. The systems can be categorized into \\nstr\\ strong and \\nwea\\ weak \\os\\ absorbers. The strong (intrinsic) absorbers are defined by the presence of either broad, fully saturated % \\os\\ absorption or partial coverage of the continuum source, and in practice all have log\\,$N$(\\os)$\\ga$15.0; these systems are interpreted as representing either QSO-driven outflows or gas close the central engine of the AGN. The weak (also known as narrow) systems show no partial coverage or saturation, and are characterized by log\\,$N$(\\os)$<$14.5 and a median total velocity width of only 42~\\kms. % The incidence d$N$/d$z$ of weak \\os\\ systems within 2\\,000~\\kms\\ of the quasar down to a limiting equivalent width of 8~m\\AA\\ is % 42$\\pm$12. Between 2\\,000 and 8\\,000~\\kms, d$N$/d$z$ falls to 14$\\pm$4, % equal to the incidence of intervening \\os\\ absorbers measured in the same spectra. Whereas the accompanying \\hi\\ and \\cf\\ column densities are significantly lower (by a mean of $\\sim$1~dex) % in the weak \\os\\ absorbers within 2\\,000~\\kms\\ of \\zqso\\ than in those at larger velocities, % the \\os\\ column densities display no dependence on proximity. Furthermore, significant offsets between the \\hi\\ and \\os\\ centroids in $\\approx$50 per cent of the weak absorbers imply that (at least in these cases) the \\hi\\ and \\os\\ lines are not formed in the same phase of gas, preventing us from making reliable metallicity determinations, and ruling out single-phase photoionization-model solutions. In summary, we find no firm evidence that quasar radiation influences the weak \\os\\ absorbers, suggesting they are collisionally ionized rather than photoionized, possibly in the multi-phase halos of foreground galaxies. Non-equilibrium collisional ionization models are needed to explain the low temperatures in the absorbing gas, which are implied by narrow line widths ($b\\!<\\!14$~\\kms) in over half of the observed \\os\\ components. % ",
        "introduction": "Quasars are ideal backlights for absorption-line spectroscopy, thanks to their high luminosities and flat continua. Yet in addition to allowing us to detect foreground absorbers, quasars also create and influence them, through the effects of outflows and ionizing radiation. The enhanced level of ionizing radiation in the vicinity of quasars gives rise to the line-of-sight proximity effect, the observed decrease in the number density of \\lya\\ forest lines (equivalent to a decrease in the mean optical depth, and an increase in the level of hydrogen ionization) at velocities approaching the quasar \\citep{Ca82, Mu86, Ty87, Ba88, Lu91, Bh94, Ra98, Sc00, Gu07, FG08}.  To obtain an unbiased view of absorbing structures in the high-$z$ Universe, one needs to remove the effects of proximity from samples of metal-line systems. The broad absorption line (BAL) systems, easy to identify in QSO spectra because of their optically thick absorption extending over thousands of \\kms, are clearly connected to the quasar. BALs trace high-velocity QSO outflows at speeds of up to 0.1$c$ \\citep[e.g.][]{Tu84, Tu88}, and are seen in 10--15\\% of quasar spectra at $z\\!>\\!1.5$ \\citep{We91, Tr06}. A separate category of mini-BAL systems, with fully saturated absorption profiles and total velocity width $<$ 2\\,000~\\kms, has also been identified \\citep{Tu88, Ch99, Yu02, Mi07b}. As with BALs, mini-BALs are interpreted as tracing outflows ejected by QSOs. A more difficult task is to decide whether {\\it narrow} absorbers in quasar spectra are either created by or ionized by the QSO. The traditional way to do this employs the displacement velocity from the quasar, $\\delta v\\!\\equiv\\!v_{\\rm qso}\\!-\\!v_{\\rm abs}$. So-called `associated' or `$z_{\\rm abs}\\!\\approx\\! z_{\\rm qso}$' or `proximate' absorbers, typically defined as narrow absorbers at $\\delta v\\!<\\!5\\,000$~\\kms\\ \\citep{We79, Fo86, An87, HF99} are often removed from intervening samples because of the possibility they may be intrinsic\\footnote{We favour the use of the term `proximate' since it involves no assumption about the absorber origin.}. Arguments in favour of an intrinsic nature for many proximate absorbers include time-variability of absorption, partial coverage of the continuum source, detection of excited ionic states implying high gas densities, profiles that are smoother than intervening absorbers, and super-solar metallicities \\citep{Wa93, Mo94, Pe94, Sv94, Tr96, BS97, Ha97, Ha97a, Ha97c, Ha01, Ha00, PS99, SP00, Ga99, Ga01, Ga06, Na04}.  Using a velocity cutoff to differentiate intervening and intrinsic absorbers has two main problems: quasar-ejected intrinsic systems can appear at higher $\\delta v$ \\citep{Ha97b, Ri99, Ri01, Mi07a, Ne08}, and intervening systems can appear at lower $\\delta v$ \\citep{Mo98, Se04}. So are many genuine intervening systems lost when they are excluded from absorber samples just because of their proximity to the quasar? In this paper we investigate the transition between proximate and intervening absorbers. We focus on \\os\\ absorbers, which trace either warm-hot ($T\\!\\ga\\!10^5$~K) collisionally-ionized plasma, or photoionized gas subject to a hard ionizing spectrum extending to energies above 113.9~eV (the ionization potential to create O$^{+5}$). \\os\\ absorbers are of interest for many reasons, including the significant role they play in the baryon and metal budgets, the window they provide on intergalactic metal enrichment, and their ability to trace energetic galaxy/IGM interactions such as accretion and galactic winds. After the first detection of \\os\\ absorbers at $z$=2--3 by \\citet{LS93}, much work has been done in the era of 10m-class telescopes to characterize their properties \\citep*{Sy00, Ca02, Be02, Si02, Si04, Si06, Be05, Lo07, Ag08, Go08}.  Here we ask a simple question: is there a proximity effect in \\os? In other words, can we see the signature of quasar photoionization, in terms of correlations between the ionization properties of absorbers and their proximity to the quasar? To address this issue, we form a sample of proximate \\os\\ absorption systems at $z$=2--3 using a homogeneous set of high-resolution VLT/UVES quasar spectra, and we then compare their properties to a sample of intervening (i.e., non-proximate) \\os\\ absorbers observed in the same set of spectra; the intervening sample has been published by \\citet[][hereafter BH05]{Be05}, but it has since been enlarged from ten to twelve sight lines. % We also briefly compare our results on \\os\\ absorbers at $z$=2--3 to those obtained in the low-redshift Universe with space-based ultraviolet spectrographs. This paper is thus a study of both the proximity effect and the nature of \\os\\ absorbers in general, and is structured as follows. In \\S2 we describe the data acquisition and reduction, and our absorber identification and measurement processes. In \\S3 we discuss the observational properties of proximate \\os\\ absorbers. In \\S4 we discuss the implications of our results, and we then present a summary in \\S5. Throughout this paper we adopt a WMAP 3-year cosmology with $\\Omega_{\\rm M}$=0.27, $\\Omega_\\Lambda$=0.73, and $h_{\\rm 70}$=1 \\citep{Sp07}. ",
        "conclusions": "The BAL, mini-BAL, and intrinsic systems, which together we classify as strong systems based on their fully saturated \\os\\ absorption, strong accompanying \\os\\ and \\cf\\ absorption, and frequent evidence for partial coverage, are well-understood as being formed in either QSO-driven outflows or inflows in the immediate vicinity of the AGN central engine \\citep[see review  by][]{HF99}. Here we focus on the weak (narrow) proximate systems, which represent a separate population. Is there any evidence that these weak systems are directly photoionized by the quasar?  We report two results that % address this question, though their interpretation is not straightforward. First, in our sample of proximate \\os\\ absorbers at $z$=2--3, there is an enhancement by a factor of three % in the incidence d$N$/d$z$ of weak \\os\\ absorbers within 2\\,000~\\kms\\ of the quasar versus those in the interval 2\\,000 to 8\\,000~\\kms. A similar enhancement is seen near low-redshift quasars \\citep{Tr08}. While at face value this could be interpreted as a proximity effect, in which quasars preferentially ionize nearby clouds more often than they photoionize more distant clouds, this is not the only explanation. The enhancement in d$N$/d$z$ could also be explained by an over-density of galaxies near quasars, and where the \\os\\ absorbers are located in the gaseous halos of these galaxies \\citep{Yo82}. It is well-known that quasars are preferentially formed in cluster environments, particularly at high redshift \\citep[][]{Cr05, Sh07}. This idea is supported by the observation that the internal properties of the weak \\os\\ absorbers are not dependent on redshift or proximity to the quasar.  Second, there are statistically significant differences % between the \\hi/\\os\\ and \\cf/\\os\\ ratios measured in the weak \\os\\ populations above and below 2\\,000~\\kms, with the weak absorbers at $\\delta v\\!<\\!2\\,000$~\\kms\\ showing a median \\hi/\\os\\ ratio lower by a factor of $\\approx$10, and a median \\cf/\\os\\ ratio lower by a % factor of $\\approx$7 than those at $2\\,000\\!<\\!\\delta v\\!<\\!8\\,000$~\\kms. % Again, at face value, these results could be interpreted as proximity effects, in which gas closer to the quasar shows a higher ionization level than intervening gas. However, closer inspection finds that while $N$(\\hi) and $N$(\\cf) show a tendency to decrease as $\\delta v$ decreases, $N$(\\os) is uncorrelated with $\\delta v$, i.e. it is the behaviour of $N$(\\hi) and $N$(\\cf) alone that is driving the trends seen in the column density ratios, not the behaviour of $N$(\\os). Importantly, the \\hi\\ and \\cf\\ proximity effects imply that proximity in velocity does correlate with proximity in distance, which supports the idea that the velocities of the weak absorbers are dominated by the Hubble flow.  Because \\os-\\hi-\\cf\\ velocity centroid offsets are observed directly in $\\approx$50 per cent of the weak systems, single-phase photoionization models for the \\os, \\cf, and \\hi\\ are inadequate for at least half the absorbers in our sample. Indeed, the non-dependence of $N$(\\os) on proximity casts doubt on whether photoionization by the background QSO creates the \\os\\ at all. Photoionization by nearby stellar sources of radiation is also unlikely, since such sources do not emit sufficient fluxes of photons above 54\\,eV (the {\\mbox{He\\,{\\sc ii}}} ionization edge) to produce the observed quantities of \\os. We cannot rule out photoionization by the quasar in every individual case, for example the absorber at $z$=2.4183 toward HE~1122-1649, which shows strong \\os\\ but barely detectable \\cf, as do several proximate \\os\\ absorbers reported by \\citet{Go08}. Nonetheless, $N$(\\os) does not depend on proximity in the way that $N$(\\hi) and $N$(\\cf) do. Thus we infer that the sizes of spheres-of-influence around quasars, which are implied the shapes of Gunn-Petersen troughs \\citep{ZD95, Sm02}, and from studies of proximate absorption in lower ionization species, such as \\mgii\\ \\citep{Va08, Wi08} and \\cf\\ \\citep*{Fo86, Ve03, Ne08}, are dependent on photon energy, and the presence of a sphere of photons at energies above 113.9~eV (capable of ionizing O$^{+4}$ to O$^{+5}$) has yet to be demonstrated. %  Thus we turn to collisional ionization models. Since the line widths of a significant fraction ($\\approx$60\\%) of the components in the weak \\os\\ absorbers % are low enough ($b\\!<\\!14$~\\kms) to imply gas temperatures below 188\\,000~K, collisional ionization {\\it equilibrium} can be ruled out, because essentially no \\os\\ is produced in gas in CIE at these temperatures \\citep{SD93, GS07}. Indeed, the narrow line widths of many intergalactic \\os\\ absorbers at $z\\!\\approx\\!2$ have led various authors to conclude that photoionization is the origin mechanism \\citep{Ca02, Be02, Lv03, Be05, Re06, Lo07}. However, \\emph{non-equilibrium} collisional ionization models cannot be ruled out so easily. Indeed one expects that collisionally ionized gas at `coronal' temperatures of a few $\\times10^5$~K, where \\os\\ is formed through collisions, will be in a non-equilibrium state. This is because the peak of the interstellar cooling curve exists at these temperatures, and so the cooling timescales are short. When the cooling times are shorter than the recombination timescales, `frozen-in' ionization can result at temperatures well below those at which the ions exist in equilibrium \\citep{Ka73, SM76, EC86}, provided that there is a source of $\\sim10^6$~K gas in the first place.  There are at least two physical reasons why collisionally-ionized, million-degree regions of interstellar and intergalactic gas could arise in the high-$z$ Universe. The first is the (hot-mode) accretion and shock heating of gas falling into potential wells \\citep{BD03, Ke05, DB06}, a process which is incorporated into cosmological hydrodynamical simulations \\citep{CO99, Da01, FB03, Ka05, CF06}, and creates what is referred to as the Warm-Hot Intergalactic Medium. However, these models generically predict that the fraction of all baryons that exist in the WHIM rises from essentially zero at $z$=3 to 30--50\\% at $z$=0, so \\emph{little WHIM is expected at the redshifts under study here}. The second reason is the presence of galactic-scale outflows, which due to the energy input from supernovae are likely to contain (or even be dominated by) hot, highly ionized gas \\citep[see recent models by][]{OD06, Fa07, KR07, Sa08}. There is strong observational evidence for outflows at redshifts of $\\approx$2--3 \\citep{He02}, including blueshifted absorption in the spectra of Lyman break galaxies \\citep{Pe00, Pe02, Sh03}, the presence of metals in the low-density, photoionized IGM \\citep[the \\lya\\ forest; e.g.][]{Ar04, Ag05, Ag08}, though see \\citet{Sy07}, and the presence of super-escape velocity \\cf\\ components in the spectra of damped \\lya\\ (DLA) galaxies \\citep{Fo07b}. In addition, collisionally ionized gas in galactic halos at high redshift has been seen directly through detections of \\os\\ and \\nf\\ components in DLAs \\citep{Fo07a}, albeit with much broader system velocity widths than in the proximate absorbers discussed here. \\os\\ absorbers with lower \\hi\\ column density may probe the outer reaches of such halos or `feedback zones' \\citep[BH05;][]{Si02}.  We explore the ability of non-equilibrium collisional ionization models to reproduce the data in our proximate \\os\\ sample in Figure 10, which shows the \\cf/\\os\\ vs \\nf/\\os\\ ratio-ratio plane. We take the isobaric non-equilibrium model at log\\,$T$=5.00 (consistent with the observed \\os\\ component line widths) from \\citet{GS07}, computed using solar abundances, and then find the values of [N/O] and [C/H] that are required to reproduce the observations. The models for 0.1 and 0.01 solar absolute abundances, and for the isochoric case, predict similar values. The non-equilibrium models are not capable of reproducing the observed ratios when solar relative elemental abundances are used. However, if [N/O] takes values between $-$1.8 and 0.4, % and [C/O] between $-$1.9 and 0.6, it is possible to explain the observed column densities in each weak proximate system with a non-equilibrium collisional ionization model. Note that we have not corrected these model predictions for the effect of photoionization by the extragalactic background; the model represents the pure collisional ionization case. Hybrid collisional+photo-ionization models would help in making progress in this area \\citep[see discussion in][]{Tr08}.  \\begin{figure} \\includegraphics[width=9.2cm]{figures/f10.eps} \\caption{Comparison of predictions of non-equilibrium collisional ionization models with the observed column density ratios in our weak proximate sample. Non-detections of \\nf\\ (\\cf) lead to upper limits on the x- (y-) axis. We take the predictions from the isobaric non-equilibrium collisional ionization model of \\citet{GS07} for an absorber at $T$=$10^5$~K (higher temperatures are ruled out by the component line widths).  The model is fairly insensitive to changes in the overall abundance level, but is sensitive to non-solar {\\it relative} abundances. The dashed green lines show the range of values of [N/O] and [C/O] that are required to fit the proximate data with this collisional ionization model.} \\end{figure}"
    },
    "0806/0806.4227_arXiv.txt": {
        "abstract": "\\fontsize{10}{10.6}\\selectfont The R CrA region was observed in the 3.5 and 6.2 cm continuum with high angular resolutions (0.6--1.7 arcseconds) using the Very Large Array. Archival data sets were also analyzed for comparison, which provided angular resolutions up to 0.3 arcseconds. A cluster of young stellar objects was detected, and a rich array of star forming activities was revealed. IRS 7A showed an enhanced outflow activity recently. The main peak of IRS 7A positionally coincides with an X-ray source, which suggests that the X-ray emission is directly related to the central protostar. The Class 0 source SMA 2 is associated with a double radio source, B 9a and 9b, and seems to be driving two outflows. The B 9 complex is probably a multiple-protostar system. Both B 9a and 9b are nonthermal radio sources with negative spectral indices. IRS 7B is a compact radio source surrounded by an extended structure. The compact source corresponds to the Class 0/I source SMA 1, and it is also closely associated with an X-ray source, suggesting that magnetic activities start early in the protostellar stage of evolution. The extended structure of IRS 7B may be a bipolar outflow. IRS 5 was resolved into two sources with a separation of 0.9 arcseconds. Both IRS 5a and 5b display radio flares and X-ray emission, suggesting that energetic magnetic processes are active in both members. The month-scale active phase of IRS 5b implies that the flare activity must involve large-scale magnetic fields. During the strong flare event of IRS 5b in 1998, IRS 5a also showed an enhanced level of radio emission. This concurrent activity suggests that IRS 5 may be an interacting young binary system, but the interaction mechanism is unknown. Alternatively, what was seen in the radio images could be a circumbinary halo. The variable radio source B 5 was found to be a nonthermal source at times, and the size of the 6.2 cm source is about 1 arcsecond, suggesting that B 5 is a Galactic object. A radio outburst of IRS 6 was detected once, and the radio/X-ray source was identified as IRS 6a. The other member of the IRS 6 system, IRS 6b, was undetected in X-rays, suggesting that only IRS 6a has detectable magnetic activities. Properties of other radio sources, IRS 1, IRS 2, and R CrA, are discussed, and the radio detections of T CrA and WMB 55 are reported. The proper motion of R CrA was marginally detected. Also presented is the classification of infrared sources in the R CrA region based on an infrared color-color diagram. ",
        "introduction": "The Corona Australis molecular cloud is one of the nearest star-forming regions, and the R CrA region is the most active site of star formation in this cloud (Harju et al. 1993; Chini et al. 2003). The R CrA cloud core contains a cluster of young stellar objects known as the ``Coronet'' cluster (Taylor \\& Storey 1984; Wilking et al. 1992, 1997). Deeply embedded protostars and preprotostellar cores were detected in the (sub)millimeter continuum (Henning et al. 1994; Choi \\& Tatematsu 2004; Nutter et al. 2005; Groppi et al. 2004, 2007). Star forming activities, such as infall, accretion, molecular outflows, and Herbig-Haro objects, were revealed by millimeter spectroscopic observations and optical imaging (Strom et al. 1974; Hartigan \\& Graham 1987; Anderson et al. 1997; Groppi et al. 2004, 2007; Nisini et al. 2005; Wang et al. 2004). Centimeter continuum imaging with interferometers has been especially effective in probing the nature of the deeply embedded objects and determining their precise positions (Brown 1987; Suters et al. 1996; Feigelson et al. 1998; Forbrich et al. 2006; Miettinen et al. 2008). High-energy processes around some of the young stellar objects were discovered in X-ray observations (Koyama et al. 1996; Hamaguchi et al. 2005, 2008; Forbrich et al. 2006, 2007; Forbrich \\& Preibisch 2007). The distance to the CrA molecular cloud is $\\sim$170 pc (Knude \\& H{\\o}g 1998).  Among the near-IR sources discovered by Taylor \\& Storey (1984) in the R CrA cloud core, IRS 7 is the most deeply embedded and probably the youngest object. In the centimeter continuum imaging by Brown (1987), IRS 7 was resolved into three radio sources: IRS 7A, IRS 7B, and B 9. Single-dish observations of submillimeter continuum emission suggested that they are deeply embedded objects (Nutter et al. 2005). Several radio continuum sources were detected around IRS 7A (Choi \\& Tatematsu 2004; Forbrich et al. 2006), suggesting that this small region contains a multiple-protostar system. The IRS 7A region is interesting because it appears to contain objects with inhomogeneous spectral classes or in different stages of evolution (Choi \\& Tatematsu 2004). While IRS 7A appears to be a Class I protostar as it was detected in near-IR, the other members of the multiple system (CT 3 and B 9) were suggested to be probable Class 0 protostars. Recent interferometric observations in the millimeter continuum confirmed that B 9 (SMA 2) is a Class 0 source while IRS 7A and CT 3 were undetected (Groppi et al. 2007).  \\begin{deluxetable}{p{16mm}cccrccrc} \\tabletypesize{\\small} \\tablecaption{Parameters of the VLA Observing Runs}% \\tablewidth{0pt} \\tablehead{ &&& \\multicolumn{5}{c}{\\sc Synthesized Beam} \\\\ \\cline{4-8} && \\colhead{\\sc Frequency} & \\multicolumn{2}{c}{Natural weight} && \\multicolumn{2}{c}{Uniform weight} & \\colhead{\\sc Array}\\\\ \\cline{4-5} \\cline{7-8} \\colhead{\\sc Track} & \\colhead{\\sc Date} & \\colhead{(GHz)} & \\colhead{Size} & \\colhead{P.A.} && \\colhead{Size} & \\colhead{P.A.} & \\colhead{\\sc Configuration} }% \\startdata Tr 1\\dotfill  & 1996. 12. 29. & 4.9 & 1.21 $\\times$ 0.50 &  $-$6.5 && 0.86 $\\times$ 0.35 &  $-$7.1 & A \\\\ &               & 8.5 & 0.70 $\\times$ 0.31 &  $-$2.5 && 0.51 $\\times$ 0.20 &     0.1 & A \\\\ Tr 2\\dotfill  & 1997. 01. 19. & 8.5 & 0.82 $\\times$ 0.59 &    10.6 && 0.57 $\\times$ 0.20 &     0.8 & BnA \\\\ Tr 3\\dotfill  & 1997. 01. 20. & 8.5 & 1.04 $\\times$ 0.89 & $-$10.4 && 0.84 $\\times$ 0.22 &    10.6 & BnA \\\\ Tr 4\\dotfill  & 1998. 06. 27. & 8.5 & 1.13 $\\times$ 0.69 &    23.8 && 0.82 $\\times$ 0.52 &    41.0 & BnA \\\\ Tr 5\\dotfill  & 1998. 07. 19. & 8.5 & 2.61 $\\times$ 0.78 &  $-$1.9 && 2.38 $\\times$ 0.54 &  $-$0.8 & B \\\\ Tr 6\\dotfill  & 1998. 09. 07. & 8.5 & 2.61 $\\times$ 0.78 &     5.4 && 2.37 $\\times$ 0.54 &     4.7 & B \\\\ Tr 7\\dotfill  & 1998. 09. 19. & 8.5 & 2.54 $\\times$ 0.80 &     0.0 && 2.30 $\\times$ 0.53 &     0.4 & B \\\\ Tr 8\\dotfill  & 1998. 09. 27. & 8.5 & 2.71 $\\times$ 0.79 &    11.6 && 2.43 $\\times$ 0.52 &    11.8 & B \\\\ Tr 9\\dotfill  & 1998. 10. 02. & 8.5 & 2.58 $\\times$ 0.79 &  $-$5.4 && 2.25 $\\times$ 0.54 &  $-$4.6 & B \\\\ Tr 10\\dotfill & 1998. 10. 06. & 8.5 & 2.65 $\\times$ 0.79 &  $-$6.6 && 2.43 $\\times$ 0.54 &  $-$5.2 & B \\\\ Tr 11\\dotfill & 1998. 10. 10. & 8.5 & 2.54 $\\times$ 0.79 &     0.1 && 2.30 $\\times$ 0.53 &     0.4 & B \\\\ Tr 12\\dotfill & 1998. 10. 13. & 8.5 & 2.60 $\\times$ 0.80 &  $-$5.6 && 2.25 $\\times$ 0.54 &  $-$4.6 & B \\\\ Tr 13\\dotfill & 2005. 02. 03. & 8.5 & 1.01 $\\times$ 0.74 &    23.4 && 0.73 $\\times$ 0.54 &    33.9 & BnA \\\\ Tr 14\\dotfill & 2005. 02. 12. & 4.9 & 2.10 $\\times$ 1.35 &     9.9 && 1.45 $\\times$ 1.02 &    14.7 & BnA \\\\ \\enddata\\\\ \\tablecomments{Units of beam size and P.A. are arcseconds and degrees, respectively. Tracks Tr 2--3 correspond to the observing runs presented by Feigelson et al. 1998. Tracks Tr 4--12 correspond to epochs R1--9 in Table 1 of Forbrich et al. 2006.}% \\end{deluxetable}    IRS 7B attracted attention because it is a rare example of Class 0 protostars showing X-ray emission (Feigelson \\& Montmerle 1999; Hamaguchi et al. 2005). Recent interferometric observations in the millimeter continuum suggested that IRS 7B (SMA 1) may be a Class 0/I transitional object (Groppi et al. 2007). Hamaguchi et al. (2005) demonstrated that the X-ray emission from IRS 7B comes from two different components: one hot and variable, and the other cool and constant. The cool component was interpreted as originating from collisional shock of the protostellar jet with circumstellar gas, and such X-ray emission has been found around several Class 0/I protostars, for example, HH 2, L1551 IRS 5, and OMC 3 MMS 2 (Pravdo et al. 2001; Favata et al. 2002; Bally et al. 2003; Tsujimoto et al. 2004). The hot variable X-ray emission had not been detected from other Class 0 protostars. This component would be related to the reconnection of magnetic fields around the protostar during mass accretion (e.g., Montmerle et al. 2000). Spatially separating the jet collision region from the central protostar with a high resolution imaging is therefore important in understanding the structure of the protostellar system.  IRS 5 is an interesting young binary system. It was spatially resolved in the near-IR and X-ray images (Nisini et al. 2005; Hamaguchi et al. 2008). Both members of the system show quiescent X-ray emission, but flare activities of them were remarkably different: only IRS 5a frequently showed solar type X-ray flares (Hamaguchi et al. 2008). IRS 5 also displays radio outbursts and emits circularly polarized radio emission (Suters et al. 1996; Feigelson et al. 1998; Forbrich et al. 2006; Miettinen et al. 2008). However, the radio properties of each binary member need to be characterized separately to understand the radio activities.  The R CrA region also contains other radio sources displaying interesting star-formation activities of their own. IRS 1 and IRS 2 are Class I protostars and variable X-ray sources (Forbrich et al. 2006). IRS 1 is the driving source of the HH 100 outflow (Strom et al. 1974). WMB 55 is a Class I source without an X-ray detection (Wilking et al. 1997; Nutter et al. 2005). IRS 6 is probably an X-ray emitting T Tauri star driving a giant Herbig-Haro flow (Wang et al. 2004; Forbrich et al. 2006). R CrA and T CrA are Herbig Ae/Be stars, and each of them may be a binary system (Takami et al. 2003). B 5 is a radio source of unknown nature, and its flux varies rapidly (Brown 1987; Suters et al. 1996; Feigelson et al. 1998).  To obtain high-quality images of the young stellar objects despite the low declination, we observed the R CrA region with the Very Large Array (VLA) using a fast-switching phase calibration technique. In this paper, we present our centimeter continuum observations of the R CrA region. We describe our radio continuum observations in \\S~2 and archival data in \\S~3. In \\S~4 we briefly report the results of the continuum imaging. In \\S~5 we discuss the star forming activities of the radio sources in the R CrA region. A summary is given in \\S~6. ",
        "conclusions": ""
    },
    "0806/0806.3541_arXiv.txt": {
        "abstract": "A dust scattering model was recently proposed to explain the shallow X-ray decay (plateau) observed prevalently in Gamma-Ray Burst (GRB) early afterglows. In this model the plateau is the scattered prompt X-ray emission by the dust located close (about 10 to a few hundred pc) to the GRB site. In this paper we carefully investigate the model and find that the scattered emission undergoes strong spectral softening with time, due to the model's essential ingredient that harder X-ray photons have smaller scattering angle thus arrive earlier, while softer photons suffer larger angle scattering and arrive later. The model predicts a significant change, i.e., $\\D \\b \\sim 2 - 3$, in the X-ray spectral index from the beginning of the plateau toward the end of the plateau, while the observed data shows close to zero softening during the plateau and the plateau-to-normal transition phase. The scattering model predicts a big difference between the harder X-ray light curve and the softer X-ray light curve, i.e., the plateau in harder X-rays ends much earlier than in softer X-rays. This feature is not seen in the data. The large scattering optical depths of the dust required by the model imply strong extinction in optical, $A_V \\gtrsim $ 10, which contradicts current findings of $A_V= 0.1 - 0.7$ from optical and X-ray afterglow observations. We conclude that the dust scattering model can not explain the X-ray plateaus. ",
        "introduction": "{\\it Swift} has discovered a generic behaviour in X-ray afterglows of Gamma-Ray Bursts (GRB): the X-ray light curve (LC) first shows a steep decline during a few hundred seconds after the end of the $\\g$-rays, then it shows a shallow decay lasting $10^4 - 10^5$ s which is followed by a ``normal'' power-law decay (Nousek et al. 2006; O'Brien et al. 2006). The normal decay at late times is the canonical afterglow component due to the interaction of the decelerated GRB ejecta with the circumburst medium, i.e., the forward shock model. The steep decline is generally interpreted to have the same origin as the prompt $\\g$-ray emission (e.g., Kumar \\& Panaitescu 2000; Liang et al. 2006).  The intervening shallow decay, sometimes called the `plateau', is the most puzzling feature of the X-ray LC. The most straightforward interpretation is a late steady energy injection into the external shock, where the latter is produced by the decelerated early ejecta plunging into the medium. The late energy injection could be due to a new ejecta from the late activity of the central engine (e.g., Dai \\& Lu 1998a,b; Zhang \\& M\\'{e}sz\\'{a}ros 2001; Dai 2004; Yu \\& Dai 2007), or due to a slow trailing part of the outflow catching up with the already forward-shock-decelerated early part of the outflow when the outflow has a spread in its Lorentz factor distribution (e.g., Granot \\& Kumar 2006). If it is the first scenario, then this interpretation implies a steady, late activity of the central engine -- lasting as long as a day -- which poses a challenge to the models of the central engine. Moreover, according to the energy-injection interpretation, the plateau-to-normal transition in the LC corresponds to the cessation of the energy injection, thus the transition should be achromatic. But in about 1/3 of the X-ray plateau GRBs with optical afterglow observations, the optical LC does not show a simultaneous plateau-to-normal break, while in another smaller fraction of the plateau cases, the plateau-to-normal breaks in optical and X-ray are indeed simultaneous (Panaitescu 2007). In most cases the power-law decay following the plateau is consistent with the predictions (the closure relationships) of the forward shock model, which in turn is consistent with the energy injection interpretation. There is a long list of alternative models for the plateau phase, such as a slow energy transfer from the ejecta to the ambient medium (Kobayashi \\& Zhang 2007), a two-component jet model (e.g., Granot et al. 2006), a varying shock microphysical parameter model (e.g., Panaitescu et al. 2006), and a reverse shock dominated afterglow model (Uhm \\& Beloborodov 2007; Genet et al. 2007), etc. (see Zhang 2007 for a review), but none of them satisfy all the observational constraints.  An attractive possibility was suggested by Shao \\& Dai (2007) regarding the origin of the X-ray plateau. If the long-duration GRB progenitors are massive stars, it is very likely that dust exists in the vicinity of the GRB site since it is in a star forming region. The X-ray photons from the GRB and its afterglow can be scattered in small angles by the dust near the line-of-sight to the GRB, as analogous to the halo emissions of other X-ray sources (e.g., Smith \\& Dwek 1998). The GRB prompt emission scattered off the dust has been considered earlier by Esin \\& Blandford (2000) and M\\'{e}sz\\'{a}ros \\& Gruzinov (2000). Aside from the scattering by the dust local to the GRB site, Miralda-Escud\\'{e} (1999) considered the scattering of the X-rays from the GRB afterglows by the dust in the intervening galaxies along the line-of-sight to the GRB, but the flux turns out to be very low and difficult to detect for that case. Depending on the distance of the local dust region to the GRB site, a delayed emission component from the scattering can show up in the afterglows. Shao \\& Dai (2007) and Shao et al. (2008) recently used this scenario to interpret the plateau phase in the X-ray afterglow LC as to be the scattered prompt X-rays by the dust located at about ten to a few hundred pc from the GRB site. The scattering happens preferentially within a characteristic scattering angle $\\th_c$ which is dependent on the photon energy $E$ and the dust grain size. At larger angles the differential scattering cross section of the dust grains decays steeply. Therefore the scattering within $\\th_c$ gives rise to a plateau phase whose duration is determined by $\\th_c$ and the distance of the dust region to the GRB site. Larger angle scattering produces a $F(t)\\propto t^{-2}$ decay following the plateau. This model does not need to invoke a long steady central engine activity. In addition, since the scattering only works in the X-ray band, the lack of a simultaneous break in optical LC does not pose a problem for this model.  The purpose of this work is to carefully investigate the output of this dust scattering model - in terms of the spectral and temporal properties of the scattered emission - and to compare it with the data. The paper is structured as follows. We first calculate and quantify the softening expected from the dust scattering model in Section 2. Then, we search in the data for evidence in favour of the model including the spectral evolution in the plateau and post-plateau phases for a sample of GRBs in Section 3 and 4. An expected difference in hard X-ray and soft X-ray LCs is discussed in Section 5. We calculate and discuss the optical extinction for the dust in Section 6. Our conclusion and further discussion are presented in Section 7. Throughout the paper the spectral index $\\b$ and the time decay index $\\alpha$ of the emission flux are defined as in $f_{\\nu}(t) \\propto \\nu^{-\\b} t^{-\\alpha}$. ",
        "conclusions": "We have shown that in the dust scattering model the scattered X-ray emission must experience strong softening spectral evolution, with a significant change of the spectral index in 0.3 - 10 keV of $\\D \\b \\sim 2 - 3$ from the emerging of the plateau to its end. However, for a sample of GRBs with X-ray plateaus and with good quality data, no softening spectral evolution during the plateau phase is found, and in a few cases even traces of slight hardening are seen.  The change of $\\b$ according to the model does not depend on the spectral index of the source emission. The Rayleigh - Gans approximation is used in this paper to calculate the scattering cross section of the dust grain. It was claimed that this approximation tends to overestimate the scattering efficiency below 1 keV, typically by a factor of 4 at 0.5 keV and a factor of 2 at 1 keV, mainly due to the absorption of the soft X-ray photons by the K and L shell electrons in the dust grain (Smith \\& Dwek 1998), and that could change the spectral slope at the soft end ($< 1$ keV) and counteract against the softening (Shao et al. 2008). But we argue that this effect dose not alleviate the expected softening, because the discrepancy between the real scattering cross section and the Rayleigh - Gans approximation caused by this effect, which is mainly below 1 keV, must have been largely accounted for by the required neutral $H$ absorption in the routine power-law fit to the plateau spectra. The XRT spectral index is mainly determined by the photons with energy above 1 keV which is not affected by this effect. Moreover, this effect is time independent while the softening we consider is a strongly time dependent behaviour. Dust destruction by the GRB prompt emission is of very little relevance here because it happens within a distance smaller than the location of the dust considered in this work (e.g., Waxman \\& Draine 2000). Thus the Rayleigh - Gans approximation is sufficiently accurate for the effect considered in this work.  The dust scattering model also predicts very different temporal behaviours in the soft X-ray vs. hard X-ray LCs; the plateau lasts longer in soft X-rays. But this feature is not found in the data. Furthermore, the large scattering optical depth of the dust required by this model in order to explain the X-ray plateaus leads to extremely large extinction in optical - $A_V \\gtrsim 10$. This is inconsistent with the observed extinctions for GRBs.  We conclude that the dust scattering model, though very attractive, can not explain the X-ray plateaus seen in most GRB afterglows. Although it is very likely that dust exists near the site of GRBs, and will scatter some fraction of the prompt and afterglow X-rays, this scattered emission is not a dominant contributor to the observed X-ray plateau. For those cases where an achromatic break at the end of the plateau is seen, a late, steady energy injection to the external shock is a more likely mechanism for producing the observed X-ray plateau, though it may not be able to work well for the cases with chromatic breaks."
    },
    "0806/0806.3384_arXiv.txt": {
        "abstract": "% {In 2005, Scholz and collaborators (Scholz et al. 2005) discovered, in a proper motion survey, a young brown dwarf SSSPMJ1102-3431 (SSSPMJ1102) of spectral type M8.5, probable member of the TW Hydrae Association (TWA) and possible companion of the T Tauri star TW Hya. The physical characterization of SSSPMJ1102 was based on the hypothesis that it forms a binary system with TW Hya.  The recent discovery of a probable giant planet inside the TW Hya protoplanetary disk with a very short-period (Setiawan et al. 2008) and a disk around SSSPMJ1102 (Riaz \\& Gizis 2008) make it especially interesting and important to measure well the physical parameters of SSSPMJ1102.} {Trigonometric parallax and proper motion measurements of SSSPMJ1102 are necessary to test for TWA membership and, thus, to determine the mass and age of this young brown dwarf and the possibility that it forms a wide binary system with TW~Hya.} {Two years of regular observations at the ESO NTT/SUSI2 telescope, have enabled us to determine the trigonometric parallax and proper motion of SSSPMJ1102.} {With our accurate distance determination of $55.2^{+1.6}_{-1.4}$ pc and proper motions of ($-67.2,-14.0$)$\\pm0.6$ mas/yr, we could confirm SSSPMJ1102 as a very probable member of TWA.  Assuming the TW Hydrae association age of $5-10$~Myr, the evolutionary models compared to the photometry of this young brown dwarf indicate a mass of $\\rm{M}=25\\pm5~\\rm{M}_{\\rm{Jup}}$ and an effective temperature $T_{\\rm{eff}}=2550\\pm100$~K.} {Our parallax and proper motion determination allow us to precisely describe the physical properties of this low mass object and to confirm its TWA membership.  Our results are not incompatible with the hypothesis that SSSPMJ1102 is a binary companion of  the star TW Hya.} ",
        "introduction": "The TW Hydrae Association (TWA) is a young, nearby association consisting of about 25 known members. Due to its youth and proximity, this association has been intensively studied in the last decade revealing a great variety of systems: tight astrometric binaries good to calibrate PMS models, stars and brown dwarfs harbouring circumstellar disks, planetary and brown dwarf companions, and more recently a putative massive planet embedded in its own proto-planetary disks (TW Hya; \\cite{seti08}). Surprisingly, only five members have known trigonometric parallaxes. De la Reza et al. (2006) report a trace-back age of $8.3\\pm0.8$~Myr, independent of evolutionary models. Relying on astrometric and spectroscopic data, the Galactic space motions of TWA members are traced backward in time until they occupy a minimum volume in space. This age estimation would greatly benefit from parallax measurements of additional TWA members.  \\cite {scho05} discovered a new young sub-stellar object, SSSPMJ1102-3431 (SSSPMJ1102), a probable member of TWA. Its photometric and spectroscopic characteristics suggest a young brown dwarf of spectral type M8.5. Located $12\\arcmin$ from TW Hya and sharing similar proper motions, Scholz et al. (2005) suggested that SSSPMJ1102 could form a binary system with TW Hya. Assuming an age of 10 Myr (\\cite {webb99}) and the Hipparcos distance for TW Hya they derived for SSSPMJ1102 a mass of $\\approx 25 M_{Jup}$. Recently, a flat optically thick disk was discovered around SSSPMJ1102 (\\cite {riaz08}) based on a reconstructed mid-infrared spectral energy distribution using broad-band photometry (\\cite {ster04}; \\cite {riaz06}). Utilizing combined NASA IRTF and Spitzer spectroscopic observations, \\cite {morr08} argued in favor of high degrees of dust settling to the disk midplane as well as significant grain growth in the upper layers, suggesting rapid dust processing compared to disks around stars.  Characterization of SSSPMJ1102 itself and its disk properties and the question of binarity status with TW Hya make a distance determination of substantial interest. Since January 2006 we have conducted astrometric and photometric observations at the ESO NTT telescope to derive the trignometric parallax of SSSPMJ1102. Our observations are presented in Section~2. The data reduction and analysis and the result of this trigonometric parallax programme are given in Section~3.  Finally, membership in TWA, the physical properties of SSSPMJ1102 compared to other TWA substellar objects, and the binarity status with TW Hya are discused respectively in Sections~4, 5 and 6.  \\begin{table}[ht] \\caption{\\label{tpi}Astrometric parameters for SSSPMJ1102-3431 derived in this work. Proper motions and parallax are absolute quantities.} \\centerline{ \\begin{tabular}{cccc} \\hline\\hline\\noalign{\\smallskip}{\\smallskip} $\\mu^{*}_{\\alpha_{abs}}$(mas/yr)  & $\\mu_{\\delta_{abs}}$(mas/yr) & $\\pi_{abs}$(mas) & $d$(pc)\\\\ \\hline\\noalign{\\smallskip}{\\smallskip} -67.2$\\pm$0.6&-14.0$\\pm$0.6&18.1$\\pm$0.5&$55.2^{+1.6}_{-1.4}$\\\\ \\hline \\end{tabular} } \\end{table}  \\begin{table}[ht] \\caption{\\label{phot}Bessel (V, R and I derived in this work) and 2MASS (J, H and K, \\cite{cutr03}) apparent and absolute magnitudes for SSSPMJ1102-3431.} \\centerline{ \\begin{tabular}{rrr} \\hline\\hline\\noalign{\\smallskip}{\\smallskip} &\\multicolumn{1}{c}{m(mag)} & \\multicolumn{1}{c}{M(mag)}\\\\ \\hline\\noalign{\\smallskip}{\\smallskip} V & 21.46 $\\pm$ 0.04 & 17.75 $\\pm$ 0.19\\\\ R & 19.14 $\\pm$ 0.04 & 15.43 $\\pm$ 0.19\\\\ I & 17.90 $\\pm$ 0.03 & 14.19 $\\pm$ 0.18\\\\ J & 13.034 $\\pm$ 0.024 & 9.32 $\\pm$ 0.17\\\\ H & 12.356 $\\pm$ 0.022 & 8.65 $\\pm$ 0.17\\\\ K & 11.887 $\\pm$ 0.024 & 8.18 $\\pm$ 0.17\\\\ \\noalign{\\smallskip}\\hline \\end{tabular} } \\end{table} ",
        "conclusions": "Motivated by the need to have accurate distance determinations for members of the TW Hydrae Association and by specific interest in the brown dwarf SSSPMJ1102, we measured its trigonometric parallax with an error $<$3$\\%$, thus ensuring a precise distance determination and a refined physical characterization of this sub-stellar object. The good accordance of our precise distance and proper motions of SSSPMJ1102 with those characterizing TWA ensure that this object belongs to the TWA Association.  Our results show that SSSPJ1102 and TW Hya are in the same region of the space sharing the same spatial movement as suspected by \\cite {scho05} but this fact still does not allow us to conclude if these two objects belong to a wide binary system or not."
    },
    "0806/0806.4930_arXiv.txt": {
        "abstract": "We present number counts, luminosity functions (LFs) and luminosity densities of galaxies obtained using the Sloan Digital Sky Survey Sixth Data Release in all SDSS photometric bands. Thanks to the SDSS DR6, galaxy statistics have increased by a factor of $\\sim9$ in the $u$-band and by a factor of $\\sim4-5$ in the rest of the SDSS bands with respect to the previous work of \\cite{Blanton2003a}. In addition, we have achieved a high redshift completeness in our galaxy samples. Firstly, by making use of the survey masks, provided by the NYU-VAGC DR6, we have been able to define an area on the sky of high angular redshift completeness. Secondly, we guarantee that brightness-dependent redshift incompleteness is small within the magnitude ranges that define our galaxy samples. With these advances, we have estimated very accurate SDSS DR6 LFs in both the bright and the faint end. In the $^{0.1}r$-band, our SDSS DR6 luminosity function is well fitted by a Schechter LF with parameters $\\Phi_{*}=0.90 \\pm 0.07$, $M_{*}-5log_{10}h=-20.73 \\pm 0.04$ and $\\alpha=-1.23 \\pm 0.02$. As compared with previous results, we find some notable differences. In the bright end of the $^{0.1}u$-band luminosity function we find a remarkable excess, of $\\sim1.7$ dex at $M_{^{0.1}u}\\simeq-20.5$,  with respect to the best-fit Schechter LF. This excess weakens in the $^{0.1}g$-band, fading away towards the very red $^{0.1}z$-band. A preliminary analysis on the nature of this bright-end bump reveals that it is mostly comprised of active galaxies and QSOs. It seems, therefore, that an important fraction of this exceeding luminosity may come from nuclear activity. In the faint end of the SDSS DR6 luminosity functions, where we can reach $1-1.5$ magnitudes deeper than the previous SDSS LF estimation, we obtain a steeper slope, that increases from the $^{0.1}u$-band, with $\\alpha=-1.01 \\pm 0.03$, to the very red $^{0.1}z$-band, with $\\alpha= -1.26 \\pm 0.03$. These state-of-the-art results may be used to constrain a variety of aspects of star formation histories and/or feed-back processes in galaxy formation models. ",
        "introduction": "\\label{sec:intro}   From the pioneering work of \\cite{Humason1956} and \\cite{Sandage1978}, who measured redshifts of bright galaxies from the Shapley-Ames photometric catalog \\cite{Shapley1932}, much effort has been invested in mapping the luminous and matter contents of the Universe. The Center for Astrophysics survey (CfA, \\citealt{Huchra1983}) is considered as the first proper redshift survey, specifically designed for cosmology studies. More important, from the first slice of about $1,000$ galaxies the CfA Redshift Survey provided the community with observational evidence of an old theoretical, and at times controversial idea: the existence of a large scale structure of galaxies  in the Universe \\citep{Davis1985}. This first vision of cosmic complexity encouraged the development of new imaging and spectrometric technology and, consequently, gave rise to a number of other redshift surveys that followed different approaches and strategies. To name but a few, the Southern Sky Redshift Survey (SSRS, \\citealt{daCosta1988}); the Perseus-Pisces catalog \\citep{Giovanelli1991} or the catalogs based on data from the Infrared Astronomical Satellite (IRAS). In the last decade, the emergence of multi-fiber spectrographs set the scene for larger and deeper redshift surveys. Examples of these are the Las Campanas Redshift Survey (LCRS, \\citealt{Shectman1996}), consisting of 26,418 galaxies with an average redshift of $z\\sim0.1$; or the 2 degree Field Redshift Survey (2dF RS, \\citealt{Colless2001}), with about $222,000$ galaxies and  covering a sky area of $1500 ~deg^{2}$. Finally, the Sloan Digital Sky Survey (SDSS, \\citealt{York2000}) is the largest photometric and spectroscopic survey ever compiled, and represents the most accurate map of the nearby universe at $z\\lesssim0.3$. The SDSS Sixth Data Release \\citep{Adelman2008}, that we use in this paper, contains spectroscopic information for more than 1,000,000 galaxies and quasars which spread over $7425 ~deg^{2}$ on the sky. Only in recent years, with surveys like the DEEP2 Galaxy Redshift Survey \\citep{Davis2003} or the VIMOS-VLT Deep Survey (VVDS, \\citealt{LeFevre2003}), have we reached the stage where we can study the galaxy population in the distant ($z\\sim1$) Universe. Other high-z surveys are currently being completed.  The advances in the survey field also made it necessary to develop data reduction pipelines and analysis tools to process and understand increasingly larger data sets. These days, cosmologists use a number of statistics to characterize, for a particular survey, the distribution of galaxies in three-dimensional space. Number counts, selection functions, luminosity functions or correlation functions are just a few examples. In this work we focus on the number counts and the luminosity functions of galaxies, that we draw from the SDSS DR6. Number counts, which describe the distribution of fluxes of galaxies, have been calculated for a number of surveys. The general consensus is that, in the close-by universe, galaxy number counts look like what we expect from an euclidean, not-evolved Universe. \\cite{Yasuda2001} obtained number counts for the SDSS Commissioning Data in all ugriz bands. \\cite{Norberg2002} provided number counts for the 2dF survey in the $b_{j}$ band. \\cite{Feulner2007} not only estimated galaxy number counts for a set of catalogs based on the Munich Near-Infrared Cluster Survey (MUNICS, \\citealt{Drory2001}) but also presented a complete revision of this subject in the literature (see their Figure 8). In contrast to galaxy number counts, the luminosity function (LF), which is the number density of galaxies per unit absolute magnitude, has been historically a rather controversial issue. For example, \\cite{Marzke1994} - using the CfA RS -,  \\cite{Norberg2002} - 2dF RS - and \\cite{Blanton2003a} - SDSS DR2-, all obtained very different results.  Both the luminosity function and the luminosity density of galaxies are observational signs of the process of galaxy formation and evolution. A precise determination of these statistics is needed to constrain current theories. Consequently, new discoveries in observational cosmology could make a strong impact in our understanding of the physical processes that drive the birth and life of galaxies in the Universe. Nowadays, state-of-the-art models of galaxy formation invoke a number of galactic \"mechanisms\", which are connected through the so-call feed-back processes. Disentangling these relations is an ambitious but crucial task in modern Cosmology. In this sense, the semi-analytic models of galaxy formation (SAMs, e.g. \\citealt{Croton2006}), which are embedded in N-body simulations like the Millennium Run (see \\citealt{Springel2005}), are a very useful tool for cosmologists. These SAMs are a good ground for testing new theoretical ideas and understanding their observational implications.  The main purpose of this work is to take advantage of the large increase in the galaxy statistics thanks to the SDSS DR6 to obtain the number counts, LFs and luminosity densities of galaxies in the nearby universe. We intend to shed light both in the faint end of the LF, where most discrepancies come from, and in the bright-end, where statistics have always been poor and errors, consequently large. In section~\\ref{sec:dr6} we briefly describe the SDSS DR6, discuss our sample selection and comment on redshift completeness. In section~\\ref{sec:results} we present our results on the number counts, the LFs and the luminosity densities of galaxies in each one of the SDSS photometric bands. Finally, in Section~\\ref{sec:discussion} we discuss our results and in Section~\\ref{sec:sum} we present a summary of our work. Throughout this paper, unless otherwise stated, we assume a standard $\\Lambda$CDM concordance cosmology, with $\\Omega_m=0.3$, $\\Omega_\\Lambda=0.7$, $w=-1$, and $h=1$. In addition, we use AB magnitudes. ",
        "conclusions": "\\label{sec:discussion}  The main results presented in this work are the SDSS DR6 Luminosity Functions of galaxies in the nearby universe. A few years ago, \\cite{Blanton2003a} used an early version of the SDSS (DR2) to calculate the SDSS galaxy LFs. Now, with the SDSS DR6 available, galaxy statistics have improved by a huge factor of $\\sim9$ in the very blue $^{0.1}u$-band and by a factor of $\\sim4-5$ in the rest of the SDSS photometric bands. Moreover, we have ensured a high redshift completeness in our galaxy samples. Firstly, we have defined an area on the sky of high angular redshift completeness by making use of the survey masks, provided by the NYU-VAGC. Secondly, we guarantee that the effect of brightness-dependent redshift incompleteness is negligible within the magnitude ranges that define our galaxy samples. These advances make our SDSS DR6 LFs substantially more precise than those from \\cite{Blanton2003a} at both the bright and the faint end. This said, the LFs of \\cite{Blanton2003a} are compatible with our results. However, notable differences, which are surely physically significant, exist. In the bright end of the blue bands LFs (especially in the $^{0.1}u$-band) we find a remarkable excess, which was very noisy in \\cite{Blanton2003a} due to their lack of statistics. In the faint end, we obtain steeper slopes in all SDSS bands, especially in the $^{0.1}u$-band, where the DR6 statistics allow us to go about $1-1.5$ magnitudes deeper as compared to \\cite{Blanton2003a}. If the huge improvement in the statistics and/or the more accurate determination of the sample completeness were not behind this discrepancies, one possible explanation could come from the so-called evolution correction \\citep{Blanton2003a}, which we have not taken into account in the determination of our absolute magnitudes. Because we are dealing with relatively nearby objects, it seems unlikely that this correction could modify our results significantly. It is worth mentioning at this point that relatively small variations in the shape of the LF, which are probably not physically significant - given the uncertainties we are dealing with - , may translate into considerable changes in the values of the best-fit Schechter parameters. It is not convenient, therefore, to make comparisons between different LFs by just looking at these best-fit Schechter parameters.   \\begin{figure} \\begin{center} \\begin{tabular}{c} \\includegraphics[width=70mm,height=45mm]{./figures/cmd_bb_3.ps} \\end{tabular} \\end{center} \\caption{The $(^{0.1}u-^{0.1}r)$ vs. $M_{^{0.1}u}$ color-magnitude diagram for the three types of bright-end bump galaxies considered: Seyferts I/QSO's (left-hand plot), LINER's/SB's (middle plot) and Bulge-like galaxies (right-hand plot). The underlying CMD of the entire $^{0.1}u$-band sample from which BEBS galaxies are selected is shown with log-spaced contours.} \\label{fig:cmd_bb} \\end{figure}    \\begin{figure} \\plotone{./figures/redblue4.ps} \\caption{The $^{0.1}r$-band SDSS DR6 Luminosity Function for blue and red galaxies separately. The SWML LF estimates are shown in diamonds. The dashed lines represents the best-fit Schechter function. Best-fit values of Schechter parameters $\\alpha$, $M_{*}$ and $\\Phi_{*}$ for both blue and red galaxies are also shown in the figure. Shaded regions represent the 1$\\sigma$ uncertainty calculated using a bootstrapping technique.} \\label{fig:redblue} \\end{figure}   We have seen that the bright-end bump that we see in the blue bands (and marginally in the rest of the bands) is statistically very significant, according to our standard bootstrapping error analysis. Moreover, as we mentioned above, we also find clear - but noisy - evidence of its existence in \\cite{Blanton2003a}. We have also checked that this excess is not a consequence of any of the limits that we have imposed to define our samples. This is, therefore, a remarkable result that may have strong implications for galaxy evolution.  From a preliminary analysis on the nature of this bright-end bump in the $^{0.1}u$-band LF, we have seen that it is mostly populated by star-forming and active galaxies ($\\sim85\\%$). The spectra of these galaxies seem consistent with what we expect from QSOs/Seyferts I ($\\sim60\\%$) and LINERs/Starbursts ($\\sim25\\%$). It seems, therefore, that an important fraction of the light that we receive from the brightest galaxies in the $^{0.1}u$-band would come from nuclear activity. Only about $15\\%$ of galaxies in bright-end bump seem to be bulge-like galaxies. A more detailed study is needed, however, to fully understand the nature of this bright-end bump.  The implications of this new results could be investigated using semi-analytic models of galaxy formation and evolution (SAMs). With these models, we can, in principle, evaluate different processes and feed-back relations that could reproduce our results."
    },
    "0806/0806.4101_arXiv.txt": {
        "abstract": "A summary of starburst luminosities based on PAH features is given for 243 starburst galaxies with 0 $<$ z $<$ 2.5, observed with the $Spitzer$ Infrared Spectrograph. Luminosity $\\nu$L$_{\\nu}$(7.7$\\mu$m) for the peak luminosity of the 7.7$\\mu$m PAH emission feature is found to scale as log[$\\nu$L$_{\\nu}$ (7.7$\\mu$m)] = 44.63($\\pm$0.09) + 2.48($\\pm$0.28) log(1+z) for the most luminous starbursts observed.  Empirical calibrations of $\\nu$L$_{\\nu}$(7.7$\\mu$m) are used to determine bolometric luminosity $L_{ir}$ and the star formation rate (SFR) for these starbursts. The most luminous starbursts found in this sample have log $L_{ir}$ = 45.4($\\pm$0.3) + 2.5($\\pm$0.3) log(1+z), in ergs s$^{-1}$, and the maximum star formation rates for starbursts in units of \\mdot are log(SFR) = 2.1($\\pm$0.3) + 2.5($\\pm$0.3) log(1+z), up to z = 2.5.  The exponent for pure luminosity evolution agrees with optical and radio studies of starbursts but is flatter than previous results based in infrared source counts. The maximum star formation rates are similar to the maxima determined for submillimeter galaxies; the most luminous individual starburst included within the sample has log $L_{ir}$ = 46.9, which gives a SFR = 3.4 x 10$^{3}$ \\mdot. ",
        "introduction": "Understanding the evolution of star formation in the universe is a fundamental objective of observational cosmology.  Various efforts using data from ultraviolet through radio wavelengths have shown that the star formation rate (SFR) per unit volume of the universe increases rapidly with redshift \\citep[e.g. ][]{mad98,haa00,cal07,lef05,tak05,man07,mar08}. It is not yet well established, however, why this evolution occurs, at what redshift is the maximum, and whether the evolution in the star formation rate is primarily luminosity evolution (more star formation per galaxy at high redshift) or density evolution (more star-forming galaxies at high redshift).  Various observations with the Spitzer Space Telescope ($Spitzer$) are providing new insights into this problem, because the signatures of luminous, rapid star formation (starbursts) are easily seen in the spectra obtained with the The Infrared Spectrograph on $Spitzer$ (IRS; Houck et al. 2004).  These features are the strong emission from Polycyclic Aromatic Hydrocarbons (PAH) which dominate the spectra of starbursts in the mid-infrared spectral region from $\\sim$ 5\\,\\um to $\\sim$ 20\\,\\um \\citep[e.g. ][]{gen98,rig00,pee04,for04}. These features are remarkably similar in starbursts ranging from low-luminosity, nearby galaxies \\citep{bra06} to ultraluminous galaxies at z $\\sim$ 2 \\citep{pop08,far08}.  Because extinction corrections are much less significant in the mid-infrared compared to optical or ultraviolet, the PAH luminosities are more confident indicators of intrinsic starburst luminosity than optical or ultraviolet emission.  In addition, the broad wavelength range over which the same features can be consistently observed with the IRS makes it possible to trace PAH luminosity for 0 $<$ z $\\la$ 3 using the strong PAH features at 6.2\\,\\um and 7.7\\,\\um.  Observing these features in a wide variety of sources provides a unique tracer of star formation over most of the history of the universe.  In the present paper, we assemble a sample of 243 starbursts selected from a wide variety of IRS observing programs and tabulate their luminosities using a consistent measure of PAH luminosity.  The results trace the most luminous starbursts in the universe at redshifts 0 $<$ z $<$ 2.5 and determine the form of luminosity evolution for starbursts within these redshifts. ",
        "conclusions": "A sample of 243 starburst galaxies with infrared spectra obtained by The Infrared Spectrograph on $Spitzer$ has been assembled with measurements of PAH luminosities to determine the most luminous starbursts discovered.  The sample includes sources from a variety of $Spitzer$ observing programs and covers 0 $<$ z $<$ 2.5 (Figure 3).  Starburst luminosities are derived from the luminosity $\\nu$L$_{\\nu}$ (7.7$\\mu$m) as determined from the peak flux density of the 7.7$\\mu$m PAH feature. These luminosities for the most luminous starbursts scale with redshift as log[$\\nu$L$_{\\nu}$ (7.7$\\mu$m)] = 44.63($\\pm$0.09) + 2.48($\\pm$0.28) log(1+z). This result demonstrates that pure luminosity evolution for starbursts scales approximately with (1+z)$^{2.5}$, at least to z = 2.5.  This is less evolution than determined from previous infrared-derived source counts but agrees with the evolution determined from optical and radio samples of star-forming galaxies.  Transformations of $\\nu$L$_{\\nu}$ (7.7$\\mu$m) to bolometric luminosities $L_{ir}$ and to star formation rates are determined empirically from local starbursts and are shown to be the same as such transformations derived by others using a variety of star formation indicators and a variety of sources. Using the conversions that log $L_{ir}$ = log[$\\nu$L$_{\\nu}$ (7.7$\\mu$m)] + 0.78, and that log[SFR] = log[$\\nu$L$_{\\nu}$ (7.7$\\mu$m)] - 42.57, for luminosities in ergs s$^{-1}$ and SFR in \\mdot, we find that:  1. Bolometric luminosities of the most luminous starbursts in the universe scale with redshift as log $L_{ir}$ = 45.4($\\pm$0.3) + 2.5($\\pm$0.3) log(1+z).  2. The SFR of the most luminous starbursts in the universe scales with redshift as log(SFR) = 2.1($\\pm$0.3) + 2.5($\\pm$0.3) log(1+z), to z = 2.5.  The most luminous starbursts in the sample are similar in SFR to the most luminous starbursts previously found in submillimeter surveys; the maximum starburst in the sample has SFR = 3.4 x 10$^{3}$ \\mdot.  We also find that at the redshifts of IRAS ULIRGs, z $<$ 0.2, the most luminous ULIRG starbursts are similar in luminosity to the most luminous pure starbursts.  These results indicate that the starburst component in composite, heavily absorbed ULIRGs having both a starburst and  AGN component is not systematically different from the pure starbursts in other sources."
    },
    "0806/0806.0208_arXiv.txt": {
        "abstract": "N-body simulations of dark matter halos show that the density profiles of halos behave as $\\rho(r)\\propto r^{-\\alpha(r)}$, where the density logarithmic slope $\\alpha \\simeq 1\\sim1.5$ in the center and $\\alpha \\simeq 3\\sim 4$ in the outer parts of halos. However, some observations are not in agreement with simulations in the very central region of halos. The simulations also show that velocity dispersion anisotropy parameter $\\beta\\approx 0$ in the inner part of the halo and the so called \"pseudo phase-space density\" $\\rho/\\sigma^3$ behaves as a power-law in radius $r$. With these results in mind, we study the distribution function and the pseudo phase-space density $\\rho/\\sigma^3$ of the center of dark matter halos and find that they are closely-related. ",
        "introduction": "\\label{S1} The formation and evolution of the dark matter halo which can be treated as the self-gravitational collisionless stellar system have become a challenging issue in the study of dark matter.  Thanks to the improved computational power, N-body simulations of dark matter halos become more and more accurate and important with increasing resolution. N-body simulations such as the universal NFW profile\\cite{NFW95,NFW96} and others\\cite{M99,J2000,AR99} show that the density profiles of dark matter halos behave as $r^{-\\alpha(r)}$, where $\\alpha \\simeq 1\\sim1.5$ in the center and $\\alpha \\simeq 3\\sim 4$ in the outer parts of halos. However,the numerical inner behaviors of dark matter halos are not supported by observations.\\cite{R85,Cour97,PW2000,Blok01,Blok03,Sal01,Sw03,Cor03} Some work indicates that density profiles of dark matter halos might become shallower than $r\\simeq1$ in the center\\cite{Austin05}, and perhaps even tend to be a core with no cusp at all.\\cite{0608376,0603051}  N-body simulations not only provide us the density profiles of halos, but also give the relevant information of the velocity space of collisionless particles in the halos. Velocity dispersion and anisotropy profiles\\cite{Rasia04,Mc05,Merritt06,Cole96,Carl97,Col00,FM01} have been well described by simple analytical fits. Two interesting phenomena from simulations indicate that velocity dispersion and the density profiles of the haloes are not independent. First, Hansen and Moore\\cite{HM06,HS06} found that the density logarithmic slope $\\alpha(r)$ is correlated to the velocity anisotropy which is parameterized by the anisotropy parameter $\\beta(r)$ and they provided the empirical formula $\\beta\\approx1-1.15(1-\\alpha/6)$. Therefore $\\beta\\approx 0$ (isotropic velocity dispersion) in the inner part as $\\alpha\\approx 1$ and $\\beta\\approx 0.5$ in the outer part as $\\alpha\\approx 3$. Second, it has been argued\\cite{TN01} that the so called pseudo phase-space density follows a power law $\\rho(r)/\\sigma^3(r)\\propto r^{-\\gamma}$ with exponent $\\gamma=1.875$, where $\\rho(r)$ is the density profile and $\\sigma^2$ is the total velocity dispersion. Subsequent studies have confirmed that $\\rho(r)/\\sigma^3(r)$ is a power law in radius, but the best fitting values of the exponent $\\gamma$ diverse from each other\\cite{Rasia04,As04,Hof07,Mac06} and range from $\\gamma= 1.90\\pm0.05$ to $2.19\\pm0.03$. Because $\\rho/\\sigma^3$ has the same dimension as the phase-space density, $\\rho/\\sigma^3$ has been called pseudo phase-space density or \"poor-man's\" phase-space density\\cite{0510332}.  With these results in mind, much theoretical work has been done for the study of the relation between density profile behavior and pseudo phase-space density. Some authors examine this matter by solving the Jeans equation. Williams et al got a critical exponents\\cite{0412442} $\\gamma=35/18$ and Dehnen and McLaughlin calculated corresponding density profiles in both isotropic and anisotropic cases\\cite{DM05}. Some other work solved the Jeans equation to explore the relation between density profile and pseudo phase-space density \\cite{Austin05,0510332,0405371,0609784}. Recently, R. N. Henriksen\\cite{0709.0434} considered a series expansion for a dark matter distribution function in the spherically symmetric anisotropic limit to discuss pseudo phase-space density.  In this paper, we concentrate on the center of the dark matter halo where the velocity dispersion is almost isotropic and calculate the distribution function and pseudo phase-space density where $r\\rightarrow 0$ in the spherically symmetric case. In Section 2 we review the basic knowledge about distribution function which is needed for this paper. In Section 3, the distribution function, velocity dispersion and pseudo phase-space density in the center of the dark matter halo are calculated with the given density profile. We make the discussion and conclusion in Section 4. ",
        "conclusions": "N-body simulations show us some properties of the dark matter halo, such as density profile $\\rho(r)$, anisotropy parameter $\\beta(r)$, pseudo phase-space density, and so on. Although the simulation's resolution is limited, these information should be considered properly in the theoretic analysis in order to make theoretic work more realistic.  In this paper, we investigate the problem of the center of the halo based on the following facts: (1) the existence of universal density profiles such as NFW\\cite{NFW95,NFW96}, Moore99\\cite{M99}, and in particular the New generalized NFW profile used in our work; (2) anisotropy parameter $\\beta\\approx 0$ in the center of the halo; (3) the pseudo phase space density follows a power law in radius $r$; (4) limited resolution in the central region.  We should notice that in this paper all results are only valid in the very center of the halo. From a given density profile, we calculate the asymptotic approximation of the distribution function in the very center of the halo. Then the total velocity dispersion and the pseudo phase-space density are obtained from the distribution function. Eq. (\\ref{gamma2}) shows the relation between the two parameters $\\gamma$ and $a$. If we set $\\gamma=1.94$ as some simulations indicate, then $a$ should be 0.776 which is smaller than NFW's result. Otherwise, if we set $a\\simeq 1\\sim 1.5$ to adapt NFW profile and Moore's profile, then $\\gamma \\simeq 2.25 \\sim 2.5$ which is larger than the simulations' result. This apparent contradiction between theoretic result and simulation is not so weird since the resolution in the central region isn't high enough to describe the very central region of the dark matter halo.  Comparing the distribution function (real phase-space density) with the pseudo phase-space density, we find that they have the same asymptotic behavior if $2>a\\geq 1$. This result indicates that the distribution function and pseudo phase-space density are closely related, even though the assertion is rather premature  that the distribution function might have a power-law behavior just as the pseudo phase-space density in the whole range of radius $r$. This interesting suggestion may also give us more confidence in the study of pseudo phase-space density and some new clues to the construction of the distribution function. By the way, some authors\\cite{0709.0434} also study the pseudo phase-space density in from $\\rho/\\sigma^{3n}$. It is obvious that our conclusion in this paper is well-founded only if $n=1$. So, from this point of view, $\\rho/\\sigma^3$ may be the best choice of the pseudo phase-space density.  For the intensive study of the dark matter halo, the full range of radius $r$ should be involved. We need to consider how to extend the relationship between \"real\" and \"pseudo\" phase-space density from the very central region to other regions where the velocity dispersion is anisotropic. Furthermore, reasonable knowledge and new strategies should be introduced and developed for the construction of more realistic models. First, it is admitted that dark matter halos in the real world are not spherically symmetric. Second, the physics in the center of the halo may be very complicated. Third, the halo actually is a polycomponent system which contains the dark matter, baryonic matter and even a supermassive black hole in the center. It is not radical to say that new methods and ideas are still needed in the future research on this matter."
    },
    "0806/0806.0949_arXiv.txt": {
        "abstract": "We have investigated properties of the Quasi Periodic Oscillation (QPO) features in the accretion powered X-ray pulsar Cen X-3 over a period of about four years using observations carried out with the Proportional Counter Array (PCA) of the {\\it {Rossi X-ray Timing Explorer}}. The observations cover a wide range of X-ray intensity of the source in excess of the binary intensity modulation. We have detected QPOs in 11 out of a total 81 pointings with the PCA with rms intensity fluctuation upto 10\\%. The QPO peak frequency shows clustering around 40 and 90 mHz with the QPO frequency having no dependence on X-ray intensity. This indicates that either (a) the observed X-ray luminosity of the source is not related to the mass accretion rate or inner radius of the accretion disk or (b) that the QPO generation mechanism in Cen X-3 is different from the beat frequency model or Keplerian frequency model that is believed to be operational in most other transient and persistent X-ray pulsars. We have also found that, the rms variation in the 40 mHz QPO feature is not dependent on the X-ray energy, indicating that disk absorption related origin for the QPO is unlikely. ",
        "introduction": "The lightcurves of X-ray binary pulsars show periodic intensity variations with the spin of the neutron star and its orbital motion. But a few of the persistent X-ray binary pulsars also show a long term periodic intensity variation with time scales more than an order of magnitude greater than the orbital period of the binary. Periodic superorbital intensity variations are seen in Her X-1 (35 day: Still \\& Boyd 2004), LMC X-4 (30.5 day: Paul \\& Kitamoto 2002) and 2S 0114+650 (30.7 day: Farrell et al 2006). SMC X-1 shows quasi periodic superorbital intensity variations with a 50-60 day cycle (Clarkson et al. 2003). The intensity variations in these systems are understood to be due to obscuration of the central X-ray source by a warped precessing accretion disc. Spectral studies of Her X-1 and LMC X-4 show iron line intensity and equivalent width evolving during the superorbital periods. Also the absorption column density in the line of sight is found to be higher during the low intensity states indicating an excess of absorbing matter in the line of sight during these times (Naik \\& Paul 2003).  Cen X-3 is a high mass X-ray binary pulsar with very strong but aperiodic long term intensity variations (Figure 1). This is the first X-ray pulsar discovered (Giaconni et al. 1971) and is also the brightest persistent pulsar. It has a spin period of $\\sim$4.8 s and an overall spin-up trend with alternate spin-up and spin-down intervals which last from about 10 to 100 days (Finger, Wilson \\& Fishman 1994). It has an orbital period of 2.1 days and a companion star of about $20M_{\\odot}$ (Avni \\& Bahcall 1974). Though Cen X-3 is a persistent pulsar, its binary period averaged X-ray intensity varies by a factor of more than 40 (Paul, Raichur \\& Mukherjee 2005). As the long term intensity variation of Cen X-3 does not remotely appear to have any periodic or quasi-periodic nature (Paul et al. 2005), it is natural to assume that the X-ray flux variation is due to changing mass accretion rate. However, using a strong dependence of the orbital modulation and the pulsed fraction of Cen X-3 on its X-ray intensity state we have shown that the long term X-ray intensity variation in this source can be due to change in obscuration by an aperiodically precessing warped accretion disk (Raichur and Paul 2008). In this scenario, as the X-ray intensity decreases, reprocessed and unpulsed X-rays from a relatively large scattering medium progressively dominates the observed X-ray intensity. We further investigate this hypothesis using the Quasi Periodic Oscillations (QPO) in Cen X-3 with respect to its intensity state. In the accretion powered X-ray pulsars, the QPOs are understood to be due to inhomogenities in the inner accretion disk and therefore its frequency is expected to be related to the inner radius of the accretion disk. A correlation between the QPO frequency and X-ray luminosity (and hence mass accretion rate / inner disk radius) has been observed in several transient and persistent X-ray sources which show a large range of X-ray intensity (EXO 2030+375: Angelini, Stella \\& Parmar 1989, 3A 0535+262: Finger, Wilson \\& Harmon 1996, XTE J1858+034: Mukherjee et al. 2006, 4U 1626--67: Kaur et al. 2008).  QPOs are known to be present in all types of accreting X-ray pulsars. In most sources it is a transient phenonema and QPOs have been detected in about a dozen out of about a hundred known accreting X-ray pulsars. With a few exceptions (4U 1748-288: Zhang et al. 1996 and XTE J 0111.2-7317: Kaur et al. 2007) the QPO frequency is ususally in the range of 40-200 mHz, consistent with it being related to the inner radius of the accretion disk around a highly magnetised neutron star in its bright X-ray state. Previous studies of Cen X-3 power spectrum have shown Quasi Periodic oscillation (QPO) at $\\sim 40$ mHz (Takeshima et al. 1991).  In the present work, we have mainly studied the QPOs of Cen X-3 and its relation to the source intensity if any. For this we have analysed all the available archival data of the {\\it{Rossi X-ray Timing Explorer (RXTE)}} proportional courter array (PCA) from the year 1996 to 2000. ",
        "conclusions": "We can summarize the principal results of our analysis presented in the previous section as follows:  1. Cen X-3 shows intermittent QPOs in different frequency ranges namely 40 mHz, and 90 mHz.  2. The presence of the QPOs or the frequency of the QPOs are not related to the luminosity state of the source. The RMS fluctuations associated with the QPOs are not correlated with the luminosity of the source.  3. RMS fluctuation of the 40 mHz QPO is energy independent.  4. A weak coupling is measured between the low frequency aperiodic variabilities and the spin frequency.  In the discussion that follows we argue that the observed QPO properties of Cen X-3 is in agreement with the scenario in which the long term X-ray intensity variation is due to change in obscuration by an aperiodically precessing warped accretion disk (Raichur and Paul 2008).  The radius of the inner accretion disk around a magnetised neutron star with a mass of 1.4 $M_{\\odot}$ and a radius of 10 km can also be approximately expressed in terms of its magnetic moment and X-ray luminosity as (Frank, King and Raine  1992) \\begin{equation} r_M = 3 \\times 10^8 L_{37}^{-2/7}\\mu_{30}^{4/7} \\end{equation} where $L_{37}$ is the X-ray luminosity in the units of $10^{37}$ erg and $\\mu_{30}$ is the magnetic moment in units of $10^{30}$ G cm$^3 $. For disc accretion, often a scaling factor of 0.5 is used with the above expression of $r_M$. Then the radius of the inner accretion disk would be $R_M = 0.5 r_M$. Coburn et. al. (2002) have estimated the magnetic field of Cen X-3 neutron star to be $B \\simeq 3.4 \\times 10^{12}$ G (i.e, $\\mu_{30} =  3.4$) using the cyclotron absorption line in the X-ray spectrum of the source.  The lowest and the highest 3-30 keV X-ray flux at which the 40 mHz QPO feature is seen are $1.1 \\times 10^{-9}$ and $1.18 \\times 10^{-8}$ erg cm$^{-2}$s$^{-1}$ respectively. Assuming a distance of 8 kpc (Krzeminski 1974), these correspond to X-ray luminosity of $L_{low} = 2.64 \\times 10^{37}$ erg s$^{-1}$ and $L_{high} = 2.83 \\times 10^{38}$ erg s$^{-1}$. If the observed X-ray luminosity represents the true X-ray luminosity and a proportional mass accretion rate of Cen X-3, the inner accretion disk radius ($R_M$) will approximately vary between $3 \\times 10^8$ cm and $1.5 \\times 10^8$ cm. We note that the corotation radius of Cen X-3 ($P_{spin} \\sim 4.8$ s) is $4.7 \\times 10^8 cm$, larger than the inner disk radius for the lowest X-ray luminosity, and the QPO detections are outside a possible propeller regime.  The two largely used QPO models are the Beat Frequency model (BFM) and the Keplarian Frequency model (KFM). The BFM explains the QPO as the beat between the spin frequency $\\nu_{spin}$ and the keplarian frequency $\\nu_k$ of the inner accretion disk $\\nu_{QPO} = \\nu_{k}-\\nu_{spin}$. Thus the radius of the inner accretion disk according the BFM is given as follows. \\begin{equation} r_{M,BFM}=\\left(\\frac{GM_{NS}}{4\\pi^2(\\nu_{spin}+\\nu_{QPO})^2}\\right)^{1/3} \\end{equation} In KFM, the QPO occurs at the keplarian frequency of the inner accretion disk $\\nu_{QPO} = \\nu_{k}$. Then radius of the inner accretion disk due to KFM will be \\begin{equation} r_{M,KFM}=\\left(\\frac{GM_{NS}}{4\\pi^2\\nu_{QPO}^2}\\right)^{1/3} \\end{equation}  However, in the case of Cen X-3, as the $\\nu_{spin}$ is larger than the observed QPO frquencies, KFM is not applicable. This is because if the inner accretion disk rotates slower than the neutron star, propeller effect is expected to inhibit accretion of material from the accretion disk. Thus assuming an inner accretion disk origin of the QPOs and using equations 1 and 2, one can experss a relation between the QPO frequency and the X-ray luminosity. In the top panel of Figure 4, we have shown the expected QPO frequency against the measured X-ray for a source distance of 8 kpc. It is obvious from the figure that the QPO frequency of Cen X-3 does not have the flux dependence as expected in the beat frequency model.   From a study of the X-ray intensity dependence of the orbital modulation and pulsed fraction in Cen X-3, recently we have proposed that the different flux states of Cen X-3 are primarily due to varying degree of obscuration by an aperiodically precessing warped accretion disk (Raichur and Paul 2008). The nearly constant QPO frequency (ignoring the rare 90 mHz feature) reported here is indeed consistent with this hypothesis. We propose that the mass accretion rate and thus the inner accretion disk radius of Cen X-3 is not highly variable, thus the source produces a nearly constant QPO frequency. We note here that the frequencies predicted by the BFM are significantly larger than the measured ones. However, the expression used here for magnetospheric radius is only approximate and a different prescription for the magnetospheric radius (for example if it is considerably larger than $r_M$ given in equation 1) can explain the observed QPO frequencies at the highest observed X-ray flux state and also be consistent with the proposal that the QPO frequency is insensitive to the measured X-ray flux as the X-ray flux variation is primarily due to disk obscuration.  We also note the other possibility that the QPOs in Cen X-3 may not be due to any material inhomogenity in the inner accretion disk as is the case for a few other X-ray binary pulsars like A0535+262 (Finger et al. 1996), EXO 2030+375 (Angelini et al. 1989) XTE J1858+034 (Mukherjee et al. 2006) and 4U 1626--67 (Kaur et al. 2008). In these transient binary X-ray pulsars, the QPO frequency is well or somewhat correlated with the X-ray luminosity of the source and hence the QPO frequency variations are understood to be due to changes in the mass accretion rate and associated changes in the radius of the inner accretion disk.  One study which could give us more insight would be to study the emission lines from neutral, H-like and He-like iron atoms in the photoionised circumstellar material. Such spectral observations done during the eclipse ingress and egress of the source would give us knowledge of the distance at which the lines are produced. Measurements carried out at different intensity states of the source will tell us if the observed X-ray intensity is a true measure of the X-ray luminosity and the ionization parameter."
    },
    "0806/0806.3876_arXiv.txt": {
        "abstract": "The X-ray spectrum of NGC~7213 is known to present no evidence for Compton reflection, a unique result among bright Seyfert 1s. The observed neutral iron K$\\alpha$ line, therefore, cannot be associated with a Compton-thick material, like the disc or the torus, but is due to Compton-thin gas, with the Broad Line Region (BLR) as the most likely candidate. To check this hypothesis, a long \\textit{Chandra} HETG observation, together with a quasi-simultaneous optical spectroscopic observation at the \\textit{ESO} NTT EMMI were performed. We found that the iron line is resolved with a FWHM=$2\\,400^{+1\\,100}_{-600}$ km s$^{-1}$, in perfect agreement with the value measured for the broad component of the H$\\alpha$, $2640^{+110}_{-90}$ km s$^{-1}$. Therefore, NGC~7213 is the only Seyfert 1 galaxy whose iron K$\\alpha$ line is unambiguously produced in the BLR. We also confirmed the presence of two ionised iron lines and studied them in greater detail than before. The resonant line is the dominant component in the Fe \\textsc{xxv} triplet, therefore suggesting an origin in collisionally ionised gas. If this is the case, the blueshift of around $1\\,000$ km s$^{-1}$ of the two ionised iron lines could be the first measure of the velocity of a starburst wind from its X-ray emission. ",
        "introduction": "A narrow component of the iron K$\\alpha$ line is almost invariably present in \\textit{Chandra} high energy gratings and XMM-Newton CCD spectra of Active Galactic Nuclei \\citep[e.g.][and references therein]{bianchi07,yaq04}. The line, typically unresolved with upper limits of several thousands km s$^{-1}$ for the Full Width at Half Maximum (FWHM), must be produced far from the nucleus, either in the torus envisaged in the Unification Model \\citep{antonucci93} or in the Broad Line Region (BLR). In the former case, and if the torus is Compton-thick, a Compton reflection component should also be present. In the latter case, a much fainter Compton reflection component is expected, and the intrinsic width of the iron line should be the same as that of the optical broad lines. The torus hypothesis seems to be preferred for the majority of sources \\citep[e.g.][]{bianchi04,nan06}, but to unambiguously distinguish between these hypotheses high energy resolution observations are required, which at present only \\textit{Chandra} can provide.  The Seyfert 1/LINER NGC~7213 (z=0.005839) presents a negligible amount of Compton reflection \\citep[R=$\\Delta\\Omega/2\\pi<0.19$:][]{bianchi03b,bianchi04}. This result, confirmed by three other BeppoSAX observations, is unique among bright Seyfert 1s observed by BeppoSAX \\citep{per02,risa02,dad08}. Therefore, the observed neutral iron line, whose equivalent width is $\\simeq80$ eV, cannot be associated with a Compton-thick material, like the disc or the torus, but is due to Compton-thin gas, like the BLR. The line width is unresolved by XMM-Newton EPIC pn, giving only an upper limit of about $8\\,000$ km s$^{-1}$. In this Letter, we present a long \\textit{Chandra} grating observation of the iron line width, which turns out to be resolved and fully consistent with the FWHM of the broad component of the H$\\alpha$ line, as measured in a quasi-simultaneous optical observation. ",
        "conclusions": "Differently from all the other Seyfert 1s with broadband X-ray observations, NGC~7213 lacks a Compton reflection component, the clearest signature for the presence of a Compton-thick material. Therefore, the observed neutral iron line must originate in a Compton-thin material, be it the BLR or a Compton-thin torus.  The \\textit{Chandra} HEG spectrum of NGC~7213 allowed us to resolve, for the first time, the iron line width at $2\\,400^{+1\\,100}_{-600}$ km s$^{-1}$ (FWHM). The analysis of a quasi-simultaneous optical observation confirmed the presence of a broad component of the H$\\alpha$ line, for which we measured a FWHM=$2640^{+110}_{-90}$ km s$^{-1}$. The widths of the two lines are in very good agreement, which suggests that they are likely to be produced in the same material.  To test if this scenario is possible, we have to verify if the observed EW of the iron line is in agreement with an origin in the BLR. Following the detailed procedure described by \\citet{yaq01}, derived from \\citet{kk87} adopting updated atomic data, the expected EW for the BLR in NGC~7213 can be written as:  \\begin{equation} \\label{formulaew} \\mathrm{EW_{FeI}}\\simeq34\\,\\left(\\frac{f_c}{0.35}\\right)\\left(\\frac{N_\\mathrm{H}}{10^{23}\\,\\mathrm{cm}^{-2}}\\right)\\,\\mathrm{eV} \\end{equation}  This estimate assumes a spherically symmetric cloud distribution for the BLR, with a covering factor of $f_c$, a column density of $N_\\mathrm{H}$ for each cloud and an iron abundance relative to hydrogen of $A_\\mathrm{Fe}=4.68\\times10^{-5}$ \\citep{ag89}. The powerlaw index of $\\Gamma=1.69$ derived from the \\textit{Chandra} and the XMM-\\textit{Newton} observations was further adopted in the derivation of (\\ref{formulaew}). Assuming $f_c=0.35$, a column density around $3\\times10^{23}$ cm$^{-2}$ can reproduce an EW$\\simeq100$ eV, which is the order of magnitude found by \\textit{Chandra} and XMM-\\textit{Newton}. These values for $f_c$ and $N_\\mathrm{H}$ are within the ranges usually assumed in photoionization models of BLRs \\citep{netzer90}. In particular, $f_c=0.35$ is the value derived by \\citet{gk98} from observational constraints in the Seyfert 1 NGC~5548, even if more `canonical' values around 0.1 and 0.25 are generally found. The large covering factor needed to account for the iron line EW may be at odds with the regular intensity of the observed broad optical lines, but the weakness or absence of an UV bump in NGC~7213 \\citep[e.g.][]{starl05} may result in a deficit of UV photoionizing photons and compensate the geometrical factor. In any case, $f_c$ and $N_\\mathrm{H}$ can be lower, provided that $A_\\mathrm{Fe}$ is larger than solar and/or the X-ray illumination of the BLR is anisotropic \\citep[see e.g.][and references therein]{yaq01}. On the other hand, we cannot exclude a further contribution to the iron line EW from a Compton-thin torus, located on a pc-scale, but we stress here that we do not have any other piece of evidence to support its presence.  The presence of emission lines from highly ionised iron in the \\textit{Chandra} HEG spectrum of NGC~7213 confirms the XMM-\\textit{Newton} results by \\citet{bianchi03b}. As suggested by these authors, their origin may be in gas photoionised by the AGN, as found in many Seyfert 1s and 2s \\citep[e.g.][]{bm02,bianchi05}. However, the gratings data presented here revealed that the resonant line is the dominant component in the Fe \\textsc{xxv} triplet (unless its blueshift is significantly larger than the one measured for the Fe \\textsc{xxvi} line), which is suggestive of an origin in gas in collisional equilibrium. This is in agreement with the results found by \\citet{starl05} using diagnostic tools based on the O \\textsc{vii} triplet observed in the XMM-\\textit{Newton} RGS spectrum of NGC~7213. It is interesting to note that the Fe \\textsc{xxv}-\\textsc{xxvi} lines' blueshift of around $1\\,000$ km s$^{-1}$ \\citep[consistent with the centroid energies of the O \\textsc{vii} lines as measured by][]{starl05}, together with the evidence that the gas producing such lines is in collisional ionization equilibrium, suggest that the hot, line emitting, gas be in the form of a starburst driven wind \\citep[see e.g.][and references therein]{heck03}, rather than an intrinsic (typically photoionized) AGN wind.  The H$\\alpha$ morphology of NGC~7213 was studied in detail by \\citet{sb96} and \\citet{hameed01}. A circumnuclear ring of star formation is located at 20-30 arcsec from the nucleus (corresponding to a few kpc) and a giant H$\\alpha$ filament is present at 2.9 arcmin of distance (around 19 kpc). Neither of these structures is observed in the \\textit{Chandra} X-ray image, which only shows an unresolved (within few arcsec) nucleus. However, \\citet{sb96} reported the possible presence of a collimated outflow within 14 arcsec from the nucleus, with a velocity around 100 km s$^{-1}$. The gas producing the ionised iron lines reported in this Letter may be the inner, hotter and faster phase of this superwind. If this is the case, this could be the first measure of the velocity of a starburst wind from its X-ray emission."
    },
    "0806/0806.4384_arXiv.txt": {
        "abstract": "This paper discusses infrared and millimetre-wavelength observations of a Lyman-$\\alpha$ blob discovered by Smith \\& Jarvis, a candidate for ionization by the cold accretion scenario discussed in Fardal et al. and Dijkstra et al. We have observed the counterpart galaxy at infrared wavelengths in deep observations with the {\\it Spitzer Space Telescope} using the IRAC 3.6, 4.5, 5.8 \\& 8.0$\\mu$m and MIPS 24$\\mu$m bands, as well as using the Max-Planck Millimeter Bolometer Array at a wavelength of 1.2mm with the IRAM 30 metre telescope. These observations probe the $\\gtsim95$kpc Lyman-$\\alpha$ halo for the presence of obscured AGN components or the presence of a violent period of star formation invoked by other models of ionisation for these mysterious objects. 24$\\mu$m observations suggest that an obscured AGN would be insufficiently luminous to ionize the halo, and that the star formation rate within the halo may be as low as $<$140 \\Msolar yr$^{-1}$ depending on the model SED used. This is reinforced by our observations at 1.2mm using MAMBO-2, which yield an upper limit of SFR $<$550\\Msolar yr$^{-1}$ from our non-detection to a 3$\\sigma$ flux limit of 0.86 mJy beam$^{-1}$.  Finding no evidence for either AGN or extensive star formation, we conclude that this halo is ionised by a cold accretion process. We derive model SEDs for the host galaxy, and use the Bruzual \\& Charlot and Maraston libraries to show that the galaxy is well described by composite stellar populations of total mass 3.42 $\\pm\\ 0.13 \\times 10^{11}$\\Msolar\\ or 4.35 $\\pm\\ 0.16 \\times 10^{11}$\\Msolar\\ depending on the model SEDs used. ",
        "introduction": "Lyman-$\\alpha$ Blobs, discovered by Steidel \\etal\\ (2000), consist of amoebic structures emitting profusely at the rest-frame wavelength of the Lyman-$\\alpha$ emission line, 1216\\AA. Whilst these very large Lyman-$\\alpha$ emitting haloes are reminiscent of the extended emission line regions observed around powerful high redshift radio galaxies (e.g. Villar-Mart\\'in \\etal, 2007b), they typically have less than 1\\% of the associated radio flux, raising the question as to what ionises the neutral Hydrogen in order to enable the emission of Lyman-$\\alpha$ photons. Three of the most plausible explanations for this are:  \\begin{itemize} \\item LABs contain hidden QSOs (e.g. Haiman \\& Rees 2001; Weidinger, M\\o ller \\& Fynbo 2004; Weidinger et al. 2005; Barrio \\etal, 2008). The luminous nature of LABs with typically $L = 10^{44}$ erg s$^{-1}$ in the Lyman-$\\alpha$ emission line alone suggests that the hard spectra, and bolometric luminosity of QSOs are prime candidates to power such a galaxy's emission. \\item The Lyman-$\\alpha$ emission comes from a dust-enshrouded, extreme starburst galaxy with a large scale superwind due to large numbers of luminous and short-lived OB stars (e.g. Taniguchi \\& Shioya 2000; Ohyama \\etal ~2003; Wilman \\etal ~2005, Matsuda \\etal\\ 2007). Observational evidence from e.g. Chapman \\etal\\ (2001) and Geach \\etal\\ (2007) associating LABs with sub-millimetre galaxies suggests that some LABs contain enshrouded starbursts forming stars at rates of $\\sim 1000$\\Msolar yr$^{-1}$. \\item We are observing the cooling radiation of a collapsing proto-galaxy inside a dark matter halo's gravitational potential (the so-called ``cold accretion'' model - e.g. Haiman, Spaans \\& Quataert, 2000, Fardal \\etal, 2001, Matsuda \\etal, 2004, Dijkstra \\etal\\ 2006a,b, Nilsson \\etal, 2006, Dijkstra \\etal, 2007, Smith \\& Jarvis, 2007).  A cold accretion scenario invokes collisional excitation and re-radiation of accreting neutral gas to power the profuse Lyman-$\\alpha$ emission extended over the halo. \\end{itemize}  The sources of ionisation residing within these galaxies (which are sometimes extended over many tens of kiloparsecs), are currently thought to be diverse. This result is due to the discovery of a 200kpc LAB in Dey \\etal (2005) associated with a 24$\\mu$m detection attributed to an obscured AGN, to the discovery by Chapman \\etal\\ (2001) of a highly luminous sub-millimetre source (with $L_{\\rm bol} \\sim 10^{13}$\\Lsolar) indicative of a very high star formation rate ($\\sim 1000 \\Msolar$ yr$^{-1}$) residing within LAB1 from Steidel \\etal\\ (2000), and to the non-detection of any apparent source of ionization associated with a LAB in the GOODS-South field by Nilsson \\etal\\ (2006).  Smith \\& Jarvis (2007) discovered only the second known Lyman-$\\alpha$ Blob thought to be ionized by the process of cold accretion, residing at $z = 2.83$. In figure \\ref{blob_contour} we show the Sloan-g' band image of the LAB from our survey, overlaid with a contour map of the \\Heii\\ narrow-band data, sensitive to Lyman-$\\alpha$ emission at the redshift of the LAB. This LAB was found to be extended over at least $\\gtsim 95$kpc, with a Lyman-$\\alpha$ luminosity of $L_{Ly\\alpha} = 2.1 \\pm 0.5 \\times 10^{43}$ erg s$^{-1}$.  \\begin{figure} \\centering \\includegraphics[height=0.90\\columnwidth,angle=270]{blob_nboverlay.eps} \\caption{Narrow-band contour map overlaid on the Sloan-g' band data. This frame is $\\sim$33\\asec on a side and North is up while East is to the left. The object to the North and East of the core of the Lyman-$\\alpha$ emission is a low-redshift interloper, identified as being at $z=0.84$ due to the presence of [O{\\sc ii}]$_{3727}$, [Ne{\\sc iv}]$_{2424}$, Mg{\\sc ii}$_{2799}$, and [C{\\sc ii}]$_{2326}$ emission in our optical spectroscopy (for more details see Smith \\& Jarvis, 2007).} \\label{blob_contour} \\end{figure}  The LAB in question is located at 17$^h$09$^m$52.3$^s$ +59$^\\circ$13$^\\prime$ 21.72\\asec\\ (J2000), on the very edge of the Extragalactic component of the {\\it Spitzer First Look Survey} (Marleau \\etal, 2004, Fadda \\etal, 2006), and was covered with any S\\slash N at all only in the 4.5$\\mu$m band (IRAC channel 2 - Lacy et al., 2005). In any case, the FLS observations were not deep enough for the type of study proposed here (see below).  AGN themselves are expected to be particularly dusty due to the torus invoked by schemes of AGN unification (Antonucci, 1993) to explain the different observed species; this also makes them bright at mid-infrared wavelengths since the warm torus is thought to reprocess the X-ray and UV photons emitted by the central engine. Indeed, through their mid-infrared emission, powerful AGN (quasars) can be identified up to high redshifts (e.g. Lacy et al., 2004, Mart\\'inez-Sansigre \\etal, 2005), even when they are so heavily obscured that they are undetected at X-ray energies (e.g. Polletta \\etal, 2006, Lacy \\etal, 2007, Mart\\'inez-Sansigre \\etal, 2007).  In the event that there is a 24$\\mu$m detection residing within the LAB, through studying its IR SED we would be able to distinguish between an obscured AGN and the starburst SED that would be expected if there was extensive ongoing star formation in the LAB counterpart galaxy.  LABs ionized by a starburst are found to have very high star formation rates based on their fluxes at sub-millimetre wavelengths equivalent to $\\sim 1000$\\Msolar\\ yr$^{-1}$ (e.g. Chapman \\etal\\ 2001, Geach \\etal, 2007). The recent high-resolution Sub-Millimetre Array (SMA) observations of Matsuda \\etal\\ (2007) did not detect LAB1 from Steidel (2000), despite the bright sub-millimetre continuum measured by Chapman \\etal\\ (2001). Matsuda \\etal\\ (2007) argued that the sub-millimetre continuum was most likely resolved out by the high spatial resolution interferometric SMA observations, suggesting that the spatial extent of the sub-millimetre emitting region was $\\gtsim$30kpc, comparable to the extent of the Lyman-$\\alpha$ emission itself. This reinforces the super-wind model for this LAB, in which rapid star formation within the halo is widely distributed, and could be powering the Lyman-$\\alpha$ emission. Observing this new halo from Smith \\& Jarvis (2007) at millimetre wavelengths then provides an additional constraint on the properties of the ionizing source residing within.  Here we present the results of two independent tests for the presence of starbursting or AGN components enshrouded within the structure. This paper is organised as follows; in section \\ref{observations} we describe our {\\it Spitzer space telescope} and MAMBO-2 observations, while in section \\ref{sec:Results} we present our results, including our improved constraints on this galaxy's spectral energy distribution (SED), and section \\ref{conclusions} presents our conclusions. Throughout this paper the AB magnitude system is used (Oke \\& Gunn, 1983), and a standard cosmology is assumed in which $H_{0}$ = 71 km s$^{-1}$, $\\Omega_{M}$ = 0.27 and $\\Omega_{\\Lambda}$ = 0.73 (Dunkley \\etal, 2008). ",
        "conclusions": "\\label{conclusions}  With our new IRAC, MIPs and MAMBO-2 data, as well as our new deep U, B \\& V band observations, we have demonstrated that the counterpart galaxy associated with the newly-discoverd LAB announced in Smith \\& Jarvis (2007) is well-described by a composite stellar population. Using this new multi-wavelength data-set, we find no plausible evidence for the presence either of an obscured active galactic nucleus or of an energetic starburst that would be required to ionize the neutral Hydrogen gas in galaxy-wide super-wind schemes. The most likely source of ionization for this particular highly luminous Lyman-$\\alpha$ halo is a ``cold accretion'' scenario, which may be able to supply sufficient energy for this profuse emission without invoking either of the other power sources suggested in the literature (see Fardal \\etal\\ 2001, Dijkstra, \\etal, 2006a,b, 2007). We also place further constraints on the host galaxy, and find it to be well described by composite simple stellar populations with total masses of 3.42 $\\pm~0.13 \\times 10^{10}$ or 4.36 $\\pm~0.16 \\times 10^{11}$\\Msolar, depending on the models used."
    },
    "0806/0806.3271_arXiv.txt": {
        "abstract": "We present the first use of the Gemini North laser guide star adaptive optics (LGS AO) system and an integral field unit (IFU) to measure the stellar velocity dispersion of the host of a luminous quasar.  The quasar PG1426+015 ($z=0.086$) was observed with the Near-Infrared Integral Field Spectrometer (NIFS) on the 8m Gemini North telescope in the H-band as part of the Science Verification phase of the new ALTAIR LGS AO system.  The NIFS IFU and LGS AO are well suited for host studies of luminous quasars because one can achieve a large ratio of host to quasar light.  We have measured the stellar velocity dispersion of PG1426+015 from  $0.1''$ to $1''$ (0.16~kpc to 1.6~kpc) to be $217\\pm15 \\, {\\rm km \\, s^{-1}}$ based on high signal-to-noise ratio measurements of Si I, Mg I, and several CO bandheads. This new measurement is a factor of four more precise than a previous measurement obtained with long-slit spectroscopy and good, natural seeing, yet was obtained with a shorter net integration time. We find that PG1426+015 has a velocity dispersion that places it significantly above the \\msig\\ relation of quiescent galaxies and lower-luminosity active galactic nuclei with black hole masses estimated from reverberation mapping.  We discuss several possible explanations for this discrepancy that could be addressed with similar observations of a larger sample of luminous quasars. ",
        "introduction": "\\label{sec:intro} The link between central supermassive black holes and the properties of their host galaxies has been firmly established in both quiescent galaxies and active galactic nuclei \\citep[AGNs;][]{kormendy95, fm00, g00a, g00b, f01, graham01, marconi03, haring04}.  Arguably, the tightest correlation is the \\msig\\ relation, which relates the mass of the central black hole (\\mbh) and the stellar velocity dispersion of the host spheroid ($\\sigma_{\\ast}$).  In active galaxies, the stellar velocity dispersion is typically measured using the \\ion{Ca}{2} triplet stellar absorption features at rest wavelengths of 8498, 8542, and 8662\\AA. The most direct and broadly applicable method for measuring the central black hole mass in active galaxies is through reverberation mapping, where the time delay between continuum and emission line variations is used as a measure of the radius of the broad-line region ($R_{\\rm BLR}$) and the width of the H$\\beta$ emission line ($\\Delta V$) is used as a measure of the gas velocity within the BLR \\citep{blandford82, peterson93}.  Then one uses the virial equation to calculate the mass: \\begin{equation} \\label{eqn:virial} M_{\\rm BH}=f \\frac{R_{\\rm BLR} (\\Delta V)^{2}}{G}. \\end{equation} The scale factor $f$ accounts for the unknown geometry, kinematics, and inclination of the BLR.  \\citet[][hereafter O04]{onken04} estimated $\\langle f \\rangle$ statistically for the AGN population under the assumption that the zero-point of the \\msig\\ relation is the same for active and inactive galaxies. They derive a value of $f=5.5 \\pm 1.7$ from a sample of 14 active galactic nuclei (AGNs) with masses determined by reverberation mapping and stellar velocity dispersions determined using the \\ion{Ca}{2} triplet.  Using single-epoch spectroscopy, this scale factor is applied in black hole mass estimates for large samples of AGNs out to high redshift (see McGill et al. 2008 for a recent summary).  The calculation of $\\langle f \\rangle$ has unfortunately been limited to low-luminosity AGNs for two reasons: dilution of stellar features by the AGN continuum and the relative scarcity of higher-luminosity AGNs.  In particular, observations of higher-luminosity objects necessitate observations of higher-redshift objects, yet telluric absorption mostly limits the utility of the \\ion{Ca}{2} triplet to $z \\le 0.06$.  Both of these problems can be somewhat circumvented with stellar velocity dispersion observations using CO bandheads in the near infrared (NIR) H-band.  This wavelength range corresponds to the maximum ratio of the stellar emission from the host galaxy to the quasar continuum \\citep{wright94,elvis94}. Additionally, current adaptive optics (AO) systems function well in the NIR. This combination facilitates velocity dispersion measurements of luminous quasar hosts, which are essential to determine if these objects fall on the \\msig\\ relation and/or are characterized by the same $\\langle f \\rangle$ as lower-luminosity AGNs.  A suitable sample of higher-luminosity quasars for such a study is the reverberation-mapped sample discussed by \\citet{peterson04}.  The 16 PG quasars studied in this work are on average 40 times more luminous, with black holes 10 times more massive, than the AGNs in the O04 sample (the O04 and PG quasar samples have average luminosities of ${\\rm log} (\\lambda L_{\\lambda}(5100 {\\rm \\AA}) / {\\rm erg~s^{-1}}) = 43.6$ and $45.2$, respectively, and average black hole masses of ${\\rm log} (M_{\\rm BH}/M_{\\odot}) = 7.75$ and 8.65, respectively). Work has already begun to measure velocity dispersions in this luminosity regime. \\citet[][hereafter D07]{dasyra07} investigated the hypothesized evolutionary link between PG quasars and ultra-luminous infrared galaxies (ULIRGs) with H-band CO bandhead measurements of 12 PG quasars. These data were obtained using the ISAAC long-slit spectrometer \\citep{moorwood98} on the 8m Antu unit of the Very Large Telescope under good natural seeing conditions.  The H-band CO absorption features were detected in most of the quasars, including four reverberation-mapped quasars.  However, the observations were still of faint hosts with significant nuclear emission contamination. This resulted in rather noisy host-galaxy spectra for some objects and velocity dispersions with uncertainties ranging from $18$ to $67 \\, {\\rm km \\, s^{-1}}$ ($12 - 36 \\%$ error).  Over the next few years, integral field units (IFUs) combined with increasingly reliable AO systems should lead to substantial improvements in studies of the hosts of luminous quasars. The primary advantage of using an IFU is that more of the light in the region of interest (for example, we probe the central $3\\arcsec\\times3\\arcsec$) is dispersed rather than just the light within a normal single slit (e.g., $0.6\\arcsec\\times120\\arcsec$ in D07). Consequently, more host-galaxy light from near the galaxy's center is gathered in a single exposure. AO is a further aid because the contamination by the intrinsically point-like quasar can be confined to the central few pixels of the image, thus minimizing the quasar dilution of the stellar absorption features.  As a demonstration  of this approach, we present observations of PG1426+015 obtained with the Near-Infrared Integral Field Spectrometer (NIFS) during Science Verification Time for the new ALTAIR Laser Guide Star (LGS) AO system at Gemini North. We specifically chose PG1426+015 because it had a relatively uncertain velocity dispersion measurement in the earlier study of D07, yet one which placed it suggestively above the \\msig\\ relation (along with two of the three other luminous reverberation-mapped quasars).  In the next section we describe our observations, data reduction, and velocity dispersion measurement technique.  We discuss our results and their implications in Section~\\ref{sec:disc} and summarize our findings in Section~\\ref{sec:conc}.  Throughout this paper, we adopt $H_{0}=70\\, {\\rm km\\, s^{-1}}$, $\\Omega_{\\rm m}=0.3$, and $\\Omega_{\\rm total}=1$. ",
        "conclusions": "\\label{sec:disc}  Two copies of the host-galaxy spectrum of PG1426+015 are shown in the top and bottom panels of Figure~\\ref{fig:spectra}.  In the top section, the smooth curve shows the spectrum of the M5Ia velocity template, broadened by the best-fit LOSVD.  The bottom spectrum shows the host spectrum with the broadened K5III velocity template.  No systematic velocity offsets were present between the atomic and molecular features in the host spectrum, so we fit all six absorption features simultaneously.  The shaded areas designate spectral regions excluded from our pPXF fit.  We excluded the $1.658-1.682 \\mu {\\rm m}$ region because it did not show absorption features and had moderate telluric contamination, the $1.731-1.752 \\mu {\\rm m}$ region because it was affected by dense sky emission, and the region redward of $1.774 \\mu {\\rm m}$ because of increasingly strong telluric absorption features.  \\begin{figure} \\figurenum{1} \\plotone{f1.eps} \\label{fig:spectra} \\caption{Two copies of the observed frame host-galaxy spectrum of PG1426+015, normalized and offset for clarity.  The smooth curve shows either the M5Ia (top spectrum) or K5III (bottom spectrum) velocity template, broadened to fit the host-galaxy absorption features using the pPXF method.  The gray bands show regions of the spectrum that were excluded from the fit for various reasons (see text).} \\end{figure}  The host-galaxy spectrum has an average signal-to-noise ratio (SNR) of about 190 per pixel. This is larger than the SNR of the long-slit ISAAC spectrum, which was obtained with five hours on-source integration.  This increase in the SNR for a spectrum acquired in less time is evidence of the combined advantages of the NIFS IFU and the LGS AO system.  The velocity dispersion resulting from the M5Ia fit is $224 \\, {\\rm km \\, s^{-1}}$, with a reduced $\\chi^{2}$ of 0.91, while the velocity dispersion resulting from the K5III fit is $209 \\, {\\rm km \\, s^{-1}}$, with a reduced $\\chi^{2}$ of 0.92.  The CO(3-0) feature at $1.693 \\mu {\\rm m}$ is better fit by the M5Ia template while the \\ion{Mg}{1} feature at $1.632 \\mu {\\rm m}$ is better fit by the K5III template. The M1III fit resulted in a velocity dispersion of  $207 \\, {\\rm km \\, s^{-1}}$, with a reduced $\\chi^{2}$ of 0.98.  Because the M1III fit was somewhat worse than the M5Ia and K5III fits, we do not include it in Figure~\\ref{fig:spectra} and we use the average of only the K5III and M5Ia values as our final velocity dispersion.  The K0III was not a good match to the host spectrum.  \\citet{genzel01} found that K5-M1 giants and supergiants are the best templates for the young stellar populations in ULIRGs \\citep[also see][]{dasyra06a,dasyra06b} while D07 used K0-M0 giants and supergiants as templates for the presumably older stellar populations in quasars. The fact that later spectral-type templates provide better fits to the host spectrum of PG1426+015 may point to a young stellar population in the host galaxy. To determine the error in individual template fits, we varied the velocity dispersion from its best-fit value while holding the velocity and higher-order Gauss-Hermite coefficients constant.  In all three cases (using the K5III, M1III, or M5Ia), the location of $\\Delta \\chi^{2} = 1$ gave a $1\\sigma$ error bar of $\\sim 10~{\\rm km \\, s^{-1}}$.  As a measure of the error due to template mismatch, we use the standard deviation of the velocity dispersions determined from the three template fits, which is also $\\sim 10~{\\rm km \\, s^{-1}}$. Our final error bar is the quadrature sum of these two errors, resulting in a velocity dispersion for PG1426+015 of $217\\pm15 \\, {\\rm km \\, s^{-1}}$.  This measurement is within an effective physical aperture of 1.6~kpc. For comparison, the effective radius ($r_{\\rm e}$) for PG1426+015 is 4.86~kpc (based on a Galfit [Peng et al. 2002] analysis of Hubble Space Telescope Wide-Field Planetary Camera 2 data).  Our aperture is intermediate between $r_{\\rm e}$ and $r_{\\rm e}/8$, two commonly quoted aperture sizes in the literature.  We consequently choose to use our measured value for the remainder of the analysis but note that the velocity dispersion would be $208\\, {\\rm km \\, s^{-1}}$ if we corrected our value using the relation derived by \\citet{jorgensen95} for E and S0 galaxies.  This correction does not account for the contribution to the velocity dispersion that we miss by excluding the central 0.16~kpc; however, this correction is expected to be small. For comparison, D07 measured the velocity dispersion of PG1426+015 to be $185\\pm67\\, {\\rm km \\, s^{-1}}$.  This is about $0.5\\sigma$ lower than our measured value, although these measurements are consistent. Our higher precision is due to the higher SNR and broader wavelength coverage of the NIFS data.  Figure~\\ref{fig:m_sigma} shows that PG1426+015 is significantly above the \\msig\\ relation for AGNs with masses determined by reverberation mapping.  The filled squares represent \\ion{Ca}{2} triplet measurements from O04.  They have been updated with velocity dispersion data from \\citet{nelson04} and improved reverberation-based black hole masses from \\citet{bentz06,bentz07} and \\citet{denney06}. The open squares represent H-band CO bandhead measurements of the PG quasars studied in D07.  The solid line denotes the \\citet{tremaine02} fit to the quiescent galaxy \\msig\\ relation.  We have not used the \\citet{ferrarese05} fit to the quiescent galaxy \\msig\\ relation because $\\langle f \\rangle$ for the AGN population has not yet been computed relative to this fit.  The velocity dispersion presented in this work moves PG1426+015 to the position marked by the open star.  With our smaller error bar, PG1426+015 is now more significantly discrepant with the \\msig\\ relation.  Note that if we compare the location of PG1426+015 relative to the steeper \\citet{ferrarese05} \\msig\\ relation fit (and assume that $\\langle f \\rangle$ is not significantly different from 5.5), it is somewhat less significantly discrepant.  \\begin{figure} \\figurenum{2} \\plotone{f2_backup.eps} \\label{fig:m_sigma} \\caption{ The \\msig\\ relationship for reverberation-mapped AGNs. Filled squares represent AGNs with velocity dispersion measurements based on the \\ion{Ca}{2} triplet.  These are from O04, updated with additional velocity dispersion data and improved reverberation-based black hole masses. Open squares show the PG quasars from D07, with velocity dispersions measured from the CO bandheads. The open star represents our new measurement, while the open square to its immediate left is the D07 position of this object. The \\citet{tremaine02} fit to the quiescent galaxy \\msig\\ relationship is shown as a solid line.  We assume $\\langle f \\rangle = 5.5$. } \\end{figure}  D07 found that three of their four reverberation mapped PG quasars lie above the \\msig\\ relation and list three possible reasons for this.  First, the scale factor, $f$, from Equation~\\ref{eqn:virial} could be different for PG quasars and the lower-luminosity AGNs in the O04 sample.  In general, differences in scale factors could arise if different populations have different inclinations \\citep[e.g.,][]{wu01,jarvis06}, accretion rates \\citep[e.g.,][]{collin06}, radiation pressure \\citep{marconi08}, or some combination of factors. Secondly, measurement errors in either \\sigs\\ or $M_{\\rm BH}$ could spuriously drive these objects to lie above the \\msig\\ relation.  D07 note that underestimates of the velocity dispersion by 10\\% are expected because of quasar continuum dilution.  We found that continuum dilution does not significantly affect our measurement by varying our inner extraction aperture to include more or less quasar continuum.  The velocity dispersions resulting from these tests were not significantly different from the value derived using the optimal inner aperture.  In addition, preliminary results show that the black hole mass in PG2130+099---one of the D07 quasars that lies above the \\msig\\ relation---may be overestimated by a factor of several, based on a new measurement of the reverberation time lag (Grier et al., in preparation). Thirdly, D07 show that small number statistics can at least partly account for the PG quasars' location above the \\msig\\ relation.  Two final caveats should be considered.  First, this quasar is an interacting system, with a companion at a nuclear separation of $2.7\\arcsec$ (4.4~kpc).  We would therefore expect the velocity dispersion to be larger than predicted by the \\msig\\ relation, or opposite to its observed location.  Additionally, CO and \\ion{Ca}{2} triplet velocity dispersions could give inconsistent results (see Dasyra et al. 2006b for a discussion of this in the context of ULIRGs and Silge \\& Gebhardt 2003 for a discussion in the context of early-type quiescent galaxies).  If PG1426+015 has an atypically young stellar population in its nuclear region, as circumstantially suggested by the need for later-type spectral templates, these young stars may not trace the full velocity dispersion profile of the bulge and cause an underestimate of the true stellar velocity dispersion."
    },
    "0806/0806.0099_arXiv.txt": {
        "abstract": "The discovery of X-ray afterglows accompanying two short bursts from SGR\\,1900$+$14 is presented. The afterglow luminosities at the end of each observation are lower by 30-50\\% than their initial luminosities, and decay with power law indices $p \\sim $0.2-0.4. Their initial bolometric luminosities are $L \\sim 10^{34}$-$10^{35}$\\,erg\\,s$^{-1}$. We discuss analogies and differences between the X-ray afterglows of SGR short bursts and short gamma-ray bursts. ",
        "introduction": "Soft gamma repeaters (SGRs) are intriguing sources of very energetic (super-Eddington luminosity) high energy bursts. They exhibit repetitive, sporadic bursting activity with typical burst durations of $\\sim100$\\,ms \\citep[e.g.,][]{woods2006}. Four have been identified as definite SGRs, three are candidate SGRs, and the X-ray source AX\\,J1818.8$-$1559 is either a candidate SGR or an anomalous X-ray pulsar (AXP) \\citep{mereghetti2007, nakagawa2007b}. Quiescent X-ray emission has been observed from SGRs with a flux which exceeds the Eddington luminosity \\citep{nakagawa2007b}. The energy reservoir for the bursts and the steady emission is generally believed to be magnetic energy dissipation in the framework of the magnetar model \\citep[e.g.,][]{duncan1992}.  Among the SGR bursts, there is a minority population with a few seconds duration known as intermediate bursts \\citep[e.g.,][]{olive2004}. Some of the intermediate duration bursts from SGR\\,1900$+$14 have an X-ray tail with a duration of a few thousand seconds that is interpreted as a part of the burst itself \\citep{lenters2003}. More rarely, giant flares occur with an initial short, intense spike, several hundred milliseconds long, followed by a long pulsating tail lasting a few hundred seconds; these are the most exotic magnetar phenomena. A large flare from SGR\\,1900$+$14 on 2001 April 18 displayed an X-ray afterglow lasting about 11 days \\citep{feroci2003}. Radio afterglows were observed after the giant flares from SGR\\,1900$+$14 on 1998 August 27 and from SGR\\,1806$-$20 on 2004 December 27 \\citep{frail1999, cameron2005}. Indeed it is possible that some short-duration cosmic gamma-ray bursts (GRBs) could actually be extragalactic giant magnetar flares \\citep{hurley2005}. Radio, optical, and/or X-ray afterglows are a common phenomenon for the long-duration GRBs (typical durations of 2 seconds or more), but they have only been observed for 13 of 20 GRBs of duration less than 2 seconds (hereafter short GRBs). In addition, a flux increase and a slow decay of the quiescent emission ($F \\propto t^{-0.3}$ over 110 days, where $F$ is the flux and $t$ is the time since the burst) has been observed after a short burst from the AXP\\,CXOU\\,J164710.2-455216 \\citep[e.g.,][]{israel2007, nakagawa2007b}. A flux increase and a decay having two different components ($F \\propto t^{-4.8\\pm0.5}$ for $t < 0.5$\\,day and $F \\propto t^{-0.22\\pm0.01}$ for $t \\gtrsim 0.5$\\,day) has been observed in the quiescent emission after an outburst from AXP\\,1E\\,2259$+$586 in 2002 \\citep{woods2004}.  Many bursts from the known SGRs have been detected by {\\it Swift} \\citep{gehrels2004} because of the very wide field-of-view and excellent sensitivity of the Burst Alert Telescope \\citep[BAT;][]{barthelmy2005}. Thanks to the prompt, precise localization by the BAT, and the fast slew capability of {\\it Swift}, X-ray follow-up observations by the X-ray Telescope \\citep[XRT;][]{burrows2005} have taken place soon after short bursts from SGR\\,1900$+$14. In this paper, we report the {\\it Swift} discovery of X-ray afterglows from three of these short bursts. We present spectral analyses of the bursts and the X-ray afterglow using BAT and XRT data. We discuss analogies and differences between the short bursts from the SGRs and the short cosmic GRBs. Despite many satellite and ground-based telescope observations, the distance to SGR\\,1900$+$14 still remains very uncertain. In this paper, the distance is assumed to be 10\\,kpc \\citep{hurley1999, vrba2000}. ",
        "conclusions": "As shown in figure \\ref{ag_lc_summary}, the luminosities of the SGR short bursts are two or three orders of magnitude larger than the backwards-extrapolated values from the SGR X-ray afterglows. This implies that the decreasing SGR X-ray afterglow is not the tail of the short bursts. Therefore this result is different from the X-ray emission accompanying intermediate bursts from SGR\\,1900$+$14, which is interpreted as the tail of the burst itself \\citep{lenters2003}. Also, compared with the afterglow of the large flare from SGR\\,1900$+$14 \\citep{feroci2003}, or with the moderate flux decay of AXP\\,CXOU\\,J164710.2$-$455216 \\citep[e.g.,][]{nakagawa2007b}, the decay time scale of the afterglow accompanying the SGR short bursts (that is, the time to decrease by a factor of $\\sim2$) is shorter by a factor of $\\sim3$. Since the observing period of the SGR X-ray afterglow ($t < 0.5$\\,day) is different from that of the moderate flux decay of AXP\\,CXOU\\,J164710.2$-$455216 ($t \\gtrsim 0.5$\\,day), the decay indices and time scales are not exactly comparable. The initial ($t < 0.5$\\,day) decay index $p = 4.8\\pm0.5$ of AXP\\,1E\\,2259$+$586 is much steeper than that of the SGR X-ray afterglow, despite the fact that the time scale of the flux decay of AXP\\,1E\\,2259$+$586 \\citep{woods2004} is consistent with the decay time scale of the SGR X-ray afterglow (same definition as above). This suggests that the X-ray afterglows of the short bursts might be a different phenomenon. Note that the decay indeces at 0.5 day after the bursts for the two AXPs CXOU\\,J164710.2$-$455216 ($p = -0.3$) and 1E\\,2259$+$586 $(p = -0.22)$ are both consistent with the decay index of the SGR X-ray afterglow.  Here, we discuss the analogies and differences between the X-ray afterglows of short bursts from SGR\\,1900$+$14 and the X-ray afterglows of short GRBs. The decay indices of the SGR X-ray afterglows, $p = $0.2-0.4 (see $\\S$ \\ref{light_curves} and table \\ref{trig_summary}), are similar to those of the two short GRBs 050724 \\citep[$p = 0.6\\pm0.2$;][]{campana2006b} and 051221A \\citep[$p = 0.04_{-0.21}^{+0.27}$;][]{burrows2006}. On the other hand, the five short GRBs 050509B \\citep[$p = 1.10_{-0.53}^{+1.26}$;][]{gehrels2005}, 050709 \\citep[$p \\gtrsim 1$;][]{fox2005}, 051210 \\citep[$p = 2.58\\pm0.11$;][]{parola2007}, 060313 \\citep[$p = 1.46\\pm0.08$;][]{roming2006} and 061201 \\citep[$p = 1.90\\pm0.15$;][]{stratta2007} have much steeper temporal indices than those of the SGR X-ray afterglows. The X-ray afterglow luminosities of the SGR short bursts are $L \\sim 10^{34}$-$10^{35}$\\,erg\\,s$^{-1}$. Considering the cosmological distances to the short GRBs, the SGR X-ray afterglow luminosity is lower by a factor of $10^{6}$-$10^{12}$ than the luminosities of the four short GRBs 050509B, 050709, 050724 and 051221A \\citep[see the following literature for their redshifts;][]{berger2005, prochaska2005, bloom2006, covino2006}. For the other three GRBs, secure redshift measurements have not been reported yet. Thus the indices of the X-ray afterglows of GRBs 050724 and 051221A are similar to those of the SGR X-ray afterglows, while the luminosities are different.  Considering the super-Eddington luminosities of the short bursts ($L \\gtrsim 10L_{\\rm Edd}$ where $L_{\\rm Edd} = 1.8 \\times 10^{38}$\\,erg\\,s$^{-1}$ is the Eddington luminosity with $M = 1.4M_{\\sun}$ and $R = 10$\\,km), one possible explanation for the SGR X-ray afterglows might be a mechanism similar to an external shock in a GRB. That is, the short burst might be generated by a relativistic jet from a neutron star, and an interaction between the jet and an interstellar medium would generate the afterglow. The lower limit to the number density around SGR\\,1900$+$14 $n \\gtrsim 0.6(N_{\\rm H}/1.91 \\times 10^{22}$\\,cm$^{-2})(d/3 \\times 10^{22}$\\,cm$)^{-1}$\\,cm$^{-3}$ (where $d$ is the distance to SGR\\,1900$+$14) is reasonable to produce the SGR X-ray afterglow considering the number density for GRBs \\citep{sari1998}. We can estimate a bulk Lorentz factor assuming a total released energy $E \\sim T_{\\rm 90}L \\sim 7.6 \\times 10^{38}$\\,erg for Burst C and $n \\sim 1$\\,cm$^{-3}$ \\citep{fenimore1996}. The velocity of the material which is responsible for the external shock emission might be a weakly relativistic jet with $\\gamma \\sim 2$ at 1000\\,s after the short burst. Another possibility might be emission from the plasma remaining after the short burst, or an intrinsic quiescent X-ray flux increase. An alternative simple possibility might be the cooling of the surface heated by the plasma of the short burst \\citep{thompson2002}. Although some models have been proposed to explain the afterglows accompanying giant flares \\citep{yamazaki2005, lyutikov2006, cea2006}, it is not clear that these models can produce the afterglow emission from short SGR bursts. The presence of these afterglows may imply that short burst activity is in fact much longer than it appears to be. Indeed, X-ray afterglow emission accompanying short SGR bursts could be a key to understanding the origin and emission mechanism of these bursts. Further prompt multi-wavelength observations of these afterglows are needed to understand the radiation process."
    },
    "0806/0806.0116_arXiv.txt": {
        "abstract": "It has been suggested that Einstein's theory of General Relativity can be modified to accomodate mismatches between the gravitational field and luminous matter on a wide range of scales. Covariant theories of modified gravity generically predict the existence of extra degrees of freedom which may be interpreted as dark matter. We study a subclass of these theories where the overall energy density in these extra degrees of freedom is subdominant relative to the baryon density and show that they favour the presence of massive neutrinos. For some specific cases (such as a flat Universes with a cosmological constant) one finds a conservative lower bound on the neutrinos mass of $m_\\nu>0.31$ eV. ",
        "introduction": "\\noindent There is compelling evidence that the baryons in the Universe are unable to generate the gravitational potentials that we observe on a wide range of scales.   A simple paradigm can be used to explain this mismatch between light and gravity: the Universe is filled with an appreciable amount of matter which is cold (i.e. has non-relativistic velocities today) and does not interact with light. It has been shown that Cold Dark Matter (CDM) can explain a host of observation, from dynamics of clusters to the formation of the cosmic web \\cite{DM1}.  The CDM paradigm has been proposed within the context of Newtonian gravity and Einstein's theory of General Relativity. It has been argued that these theories may not be valid on all scales. Indeed, proposals for modifying gravity have been shown to fit much of the currently available data \\cite{MOND}.  A plethora of covariant theories have been studied in detail; TeVeS gravity, modified Einstein-Aether theories, conformal gravity, higher derivative actions, etc,  have been advocated as possible rival theories to the CDM scenario \\cite{Sanders,mannheim,Teves,Aether,Zlosnik,Zhao,gen_teves}. There has been considerable effort in studying the cosmological consequences of these theories \\cite{skordis,skordis_per,bourliot,Zlosnikcosmo}. Given the level of precision of current cosmological data, it is possible to find severe constraints on these alternative theories and compare their ability to describe nature with the CDM scenario.  There is an important, generic feature of covariant theories of modified gravity which is often overlooked: although they tamper with the gravitational sector of the equations of motion, they also {\\it inevitably} lead to the introduction of extra degrees of freedom which may be interpreted as an exotic form of dark matter. Let us exemplify. Theories which modify the Einstein-Hilbert action by, for example, replacing the Ricci scalar, $R$, by a function of different curvature invariants, $f(R,R_{\\alpha\\beta}R^{\\alpha\\beta},\\cdots)$, introduce higher derivative terms, and hence new modes. These new modes will contribute to the overall energy density. This is patently obvious in the case of theories where $f$ is simply a function of $R$; such theories can be mapped onto normal Einstein gravity with an additional scalar field. This also true of conformal gravity, where the action is now constructed from the Weyl tensor. A field must be added to fix the scale of gravity and the resulting low energy equations are fourth order \\cite{FofR}. More modern attempts at constructing theories of modified gravity have the same characteristics in a much more explicit way. In TeVeS \\cite{Teves}, a scalar field and a vector field is introduced which not only modify the gravitational field equations but also source the very same field through their stress energy tensor. In generalized Einstein-Aether theories, a time-like vector field is introduced \\cite{Zlosnik}.  Given what we have just said, there is an obvious question: aren't these extra degrees of freedom simply a contrived form of dark matter? It is conceivable that the extra degrees of freedom in modified theories of gravity may play such a role. If so, dark matter has been introduced through the back door. It turns out that the role of extra degrees of freedom in theories of modified gravity is more complicated than one might expect. In Skordis {\\it et al} \\cite{skordis}, it was shown that the extra degrees of freedom in TeVeS can make a negligible contribution to the background (or overall) energy density. Indeed, if TeVeS is to be consistent with big bang nucleosynthesis, the fractional energy density in these extra degrees of freedom, $\\Omega_X$, must be under a percent.  Yet even though $\\Omega_X\\ll 1$, fluctuations in the extra fields could have a significant impact on the growth of structure. In particular, due to the modified nature of gravity, they could source the growth of gravitational potentials and sustain them through Silk damping at recombination. These results were corroborated in Dodelson and Liguori \\cite{DL}, where the fluctuations in the vector field were found to play an important role.  Hence some theories of modified gravity can fit current observations of large scale structure, either from galaxy surveys or the cosmic microwave background, even though $\\Omega_X\\ll1$. We would like to point out that the latter property is not generic. In some incarnations $\\Omega_B\\ll\\Omega_X\\simeq 1$ where $\\Omega_B$ is the fractional energy density in baryons. These theories end up being a hybrid of the two paradigms, modified gravity and dark matter, and in principle should be harder to distinguish from dark matter theories (although there are some suggestions of specific tests) \\cite{Zlosnik,Zhao}.  In this paper we will try to expand on an important feature of TeVeS pointed out in Skordis {\\it et al} : if one assumes that the Universe is flat and the only form of non relativistic matter consists of $5\\%$ baryons (consistent with Big Bang Nucleosynthesis), the angular power spectrum of the Cosmic Microwave Background (CMB) will differ significantly from observations. The only way to resolve this discrepancy is to introduce some form of non-relativistic matter, and the only one allowed within the known menagerie of fundamental constituents of the Universe is a massive neutrino. To match observations of the CMB, neutrinos with a mass of approximately $2$ eV are needed. This result is clearly a hint of a more general statement that may be made about theories of modified gravity in which the extra degrees of freedom play a subdominant role: if these theoriess are to agree with measurements of the CMB then they require the presence of massive neutrinos. We wish to see if this implies a lower bound on the mass of the neutrino. ",
        "conclusions": "It has been claimed that modified theories of gravity inevitably require the presence of massive neutrinos and that these may be sufficiently massive to be measurable with up and coming neutrino experiments  such as KATRIN \\cite{KATRIN}. This claim has been triggered by two pieces of anecdotal evidence. Firstly that the simplest TeVeS model needs neutrinos to fit the angular power spectrum of the CMB as shown in \\cite{skordis}. And secondly, that attempts at reconciling observed and inferred masses of clusters requires the presence of a massive neutrino halo \\cite{sanders, famaey}. In this letter we have attempted to extend the remit of the first piece of evidence.  We have found that, although for a restricted set of models, we can place a lower bound on the mass of the neutrino, for more general ranges of parameters, it is possible to satisfy the subset of cosmological constraints without having to invoke massive neutrinos. This is not to say that specific models with, for example, a variable effective Newton's constant might not lead to a tight constraint on the neutrino mass. But it is clearly not possible to make a definitive statement on the mass of the neutrino for general theories of modified gravity. Theories must be studied case by case and we have shown how this can be done in an economical way.  It may be possible to come up with constraints on the neutrino masses from a different set of observables, related to the second piece of evidence. For example, in the simplest picture of a cluster in these theories, neutrinos seem to be inevitable to be able to make up dynamical mass measurements and weak lensing observations. This simple picture is incomplete and much of the work that has been done on clusters in the context of modified gravity has opted to ignore the extra degrees of freedom \\cite{Dai}. They can play a significant role and, in the same way as for large scale observations, may substantially weaken cluster constraints on the neutrino mass. A more detailed analysis of these systems must be undertaken before definitive conclusions can be inferred.  {\\it Acknowledgments}: We thank A. Cooray, A. Melchiorri, G. Starkman and T. Zlosnik for discussions. C. Zunckel is supported by a Domus A scholarship awarded by Merton College. Research at Perimeter Institute for Theoretical Physics is supported in part by the Goverment of Canada through NSERC and by the Province of Ontario through MRI.   \\vspace{-.3in}"
    },
    "0806/0806.0411.txt": {
        "abstract": "We give detailed predictions for the spectral signatures arising from photon-particle oscillations in astrophysical objects. The calculations include quantum electrodynamic effects as well as those due to active relativistic plasma. We show that, by studying the spectra of compact sources, it may be possible to {\\it directly} detect  (pseudo-)scalar particles, such as the axion, with much greater sensitivity, by roughly three orders of magnitude, than is currently achievable by other methods. In particular, if such particles exist with masses $m_a<10^{-2}$\\,eV and coupling constant to the electromagnetic field, $g>10^{-13}\\,{\\rm GeV}^{-1}$, then their oscillation signatures are likely to be lurking in the spectra of magnetars, pulsars, and quasars. ",
        "introduction": "It is well known that the combined product of the charge-conjugation (C) and parity (P) symmetry holds for electromagnetic interactions. Nevertheless, for some weak interactions it is known to be violated; so called CP-violation. This near (but incomplete) symmetry of nature in weak interactions is contrasted by the fact that the strong interaction, i.e., quantum chromodynamics (QCD), seems to be consistent with preserving CP symmetry to a very high precision since, otherwise, the electric dipole moment of the neutron would have been $\\sim 10$ orders of magnitude higher than current observed limits. In particular, the lack of broken symmetry requires fine tuning of QCD so that one of its parameters (treated as an effective angle in the Lagrangian) is very small instead of being of order unity, as may be a-priori expected. A way to overcome the fine tuning problem is via the Peccei-Quinn mechanism (Peccei \\& Quinn 1977) in which the angle becomes a dynamical field. This field, dubbed the axion (Wilczek 1978), gets a mass from QCD instanton effects, leading to a spontaneous relaxation of the angle  whereby the field energy is minimized (an insightful and entertaining review is given by Sikivie 1996). The axion is a (charge-less) pseudoscalar field that interacts very weakly with matter (e.g., Dine et al. 1981, Kim 1979, Sikivie 1983). Interestingly enough, pseudoscalar particles with axion-like properties and can be found in string theory (see among others, Banks \\& Dine 1997, Bar 1985, Choi \\& Kim 1985 and Green et al. 1987). Such particles also serve as cold dark matter candidates (e.g., Kolb \\& Turner 1990) forming as a Bose condensate  in the very early stages of our universe evolution, when Peccei-Quinn symmetry breaking takes place (e.g., Turner 1987 and references therein).  We note that besides pseudo-scalar particles, scalar particles may exist such that  their properties do not change sign under reflection. Neither pseudoscalar particles nor scalar particles have been detected to date  (a possible detection was reported recently by the PVLAS collaboration but the claim was later withdrawn; see below and also Asztalos et al. 2006 for a general review).   The properties of the pseudo-scalar particle (hereafter referred to sometimes as axions) are very loosely predicted by current theory. In particular, the axion mass cannot be predicted from first principles though various cosmological and astrophysical constraints, as well as experimental data, suggest it is probably $<1$\\,eV (Turner 1987). (For the axion to be a viable candidate for dark matter, its mass should be of order $10^{-5}$\\,eV.) In addition, the value of the coupling constant, $g$,  between axions and the electromagnetic field is unknown although quantum chromodynamics (QCD) models predict a relation between the (unknown) particle mass and the coupling amplitude such that more massive axions couple more strongly. Axions can, in principal, couple to photons, leptons, and baryons with the strength of the interaction depending on the Peccei-Quinn charges of the $u$ and $d$ quarks and the electron's. Different axion models predict different interaction strengths depending on the charges' values which are also not known from first principles.   As mentioned before, the axion mass and the coupling constant are essentially unknown and, in principal, may span a very large volume of the parameter space. Realistically, however, a part of the parameter space is less favored by astrophysical and cosmological considerations: axions whose mass is $<10^{-6}$\\,eV are likely to (but not necessarily) result in $\\Omega_M>0.3$ thereby closing the universe, in contrast to observations. Axions whose mass is in the eV range and which constitute a fair fraction of the mass of our galaxy (as any viable dark matter particle candidate would) are more likely to decay to photons - via their stronger coupling constants - and be observable as optical light glow over the sky, in contrast to observations (Turner 1987). These considerations suggest that the more relevant mass range for the axion is $10^{-6}<m_a<10^{-3}$\\,eV (though a more extended range cannot be firmly rejected). Additional constraints come from the coupling of  axions to photons in the interiors of stars which, for strong enough coupling, would enhance the cooling rate of those objects thereby shortening their lifetime. Measurements of old stellar populations in globular clusters [horizontal branch (HB) and red giant stars] suggest that the coupling constant is limited to $g<10^{-10}\\,{\\rm GeV}^{-1}$ (see Raffelt 1996). This agrees with current mass limits for standard axion models.  Apart from general theoretical arguments that indirectly constrain the physical parameters of axions, there is a considerable ongoing effort to directly detect these particles. We make no attempt to fully cover the numerous experimental methods  especially devised for this purpose but note a few general classes of them and that most utilize coupling of axions to photons. Microwave cavity experiments (Sikivie et al. 1983) are most relevant for detecting axions as dark matter candidates  and use the proposition that the number of axions through any given surface is large if they are to constitute dark matter hence, in the presence of a magnetic field, some of them would convert to photons. For axion mass of $\\sim 10^{-5}$\\,eV and energies typical of the virialized halos of galaxies, the photons with which such particles mix are the microwave band and a cavity is built with appropriate resonances so as to enhance their conversion. These experiments are sensitive to very low $g-$values yet scanning the entire mass range quickly and effectively is yet to be fulfilled (e.g., Duffy et al. 2006, Hagmann et al. 1990, Wuensch et al. 1989). That said, if axions exist but their density is low (and so do not constitute dark matter) then they cannot be easily detected by such experiments. The currently operating Cern Axion Solar Telescope (CAST; see Lazarus et al. 1992 and references therein for the general concept and Andriamonje et al. 2007 for recent findings) utilizes the fact that $\\sim$\\,keV photons in the sun's core convert to axions that reach the earth. Applying a $\\sim 10^5$\\,G magnetic field, this experiment attempts to reconvert solar axions to photons and observe them. This approach has yielded limits on the coupling constant of axions to photons which are comparable to indirect astrophysical arguments constraints suggesting that  $g<10^{-10}\\,{\\rm GeV}^{-1}$ (e.g., Raffelt 2007 and references therein). Photon regeneration experiments (\"light shining through walls\") are based on a similar concept and use a light source and a magnetic field to convert a fraction of the photons to axions (e.g.,  van Bibber et al. 1987, Cameron et al. 1993, Sikivie 1983, and also Adler et al. 2008). A mounted wall blocks the light ray from propagating but not the axion ray which is later reconverted to photons by a similar magnetic field (e.g., Rabad\\'an et al. 2006 and references therein). Using such a method, upper limits on the coupling constant, $g$, which are about 4 orders of magnitude higher than current astrophysical constraints are obtained (e.g., Robilliard et al. 2008; but see improvements recently suggested  by Sikivie et al. 2007). One such ongoing experiment is the PVLAS experiment (e.g., Zavattini et al. 2006). This experiment announced a tentative detection of a signal that may be interpreted as due to the elusive  axion. Nevertheless, a more thorough investigation  of the measurements suggested a problem with the experimental setup and the claim was later withdrawn by the same group (Zavattini et al. 2007).   We note that the coupling of the axion to other fields have also been investigated. In particular, a different type of experiment was initiated by Youdin et al. (1996) and utilizes a gravitational interaction potential between spin and matter (Moody \\& Wilczek 1984). Much heavier, axion-like particles that are predicted by some theories are probed by the DAMA experiment (e.g., Bernabei et al. 2008a,b and references therein). To date, the axion has not been detected and the current limits are about 10 orders of magnitude higher than most theoretical predictions.  A major limitation of most terrestrial experiments for the detection of photon-particle mixing has to do with  the fact that the expected axion-photon conversion probability $P_{\\gamma \\rightarrow a} \\propto g^2 B^2R^2$ is small (where $B$ is the magnetic field and $R$ the size of the system; see \\S4 for the accurate expression for the conversion probability in the general non-linear case, and also Raffelt \\& Stodolsky 1988 and Sikivie 1983). The small  probabilities require high signal-to-noise (S/N) data which render secure detection challenging. It is worthwhile to compare laboratory expected probabilities to those which may be expected from astrophysical objects under the assumption that the magnetic field is in equipartition with gravity i.e., $GM^2/R \\sim B^2R^3$; here, $G$ is Newton's constant and $M$ the object's mass (this seems to be a fair approximation for many astrophysical systems and seems to be supported by recent numerical simulations; e.g., Igumenshchev \\& Narayan 2002). In this case we obtain that, in the limit of small conversion probabilities and in the linear regime holds (see however \\S 4.1), the ratio between the conversion probabilities is \\begin{equation} \\frac{P_{\\gamma \\rightarrow a}^{\\rm laboratory~experiments}}{P_{\\gamma \\rightarrow a}^{\\rm celestial~objects}} \\simeq \\left ( \\frac{B_{\\rm lab}}{B} \\right)^2 \\left ( \\frac{R_{\\rm lab}}{R} \\right )^2 \\simeq  \\left ( \\frac{R_{\\rm obj}}{R_s} \\right )^{2} \\frac{B_{\\rm lab}^2 R_{\\rm lab}^2}{c^4/8G} \\ll 1, \\end{equation} where $R_s\\equiv 2GM/c^2$ is the Schwartzschild radius. For the particular case of a celestial object whose size is of the order of its Schwarzschild radius, and taking $B_{\\rm lab}=10^5$\\,G and $R_{\\rm lab}=10$\\,m, one obtains a ratio  $<10^{-30}$ (non-linear effects are treated in \\S 4).  This demonstrates the potentially greater sensitivity that may be achieved in the case of  astrophysical objects compared to laboratory experiments. We note that several other works have already taken an astrophysical approach for constraining the axion properties. An incomplete list includes: Deffayet et al. (2002) who considered photon-particle oscillations as an  explanation to supernovae dimming, M\\\"ortsell \\& Goobar (2003) who investigated the physical properties of very light axions (whose mass, $m_a\\sim 10^{-16}$\\,eV) using {\\it Sloan digital sky survey} quasar spectra,  Brockway et al. (1996) who deduced an upper limit on the coupling constant for very light axions of $\\sim 10^{-11}\\,{\\rm GeV}^{-1}$ from supernovae data, Rubbia \\& Sakharov (2008) who put more stringent constraints on heavy, $m_a>10^{-4}$\\,eV axions from polarization studies of the prompt emission in a $\\gamma$-ray burst, and Hochmuth \\& Sigl (2007) who investigated the observational implications of the (recently withdrawn) PVLAS experiment results. More relevant to our study is the recent work by Lai \\& Heyl (2006) who explored the possibility for axion detection in the case of magnetars.  In this work we wish to see whether, by studying at the spectra of various astrophysical objects, one can hope to observe the signatures of photon-particle conversion down to low values of the coupling constant and extend the physical parameter space accessible to us. This approach has been suggested in the past and was qualitatively treated in several works (e.g., Lai \\& Heyl 2006 and references therein). Nevertheless, the application of such methods is more complicated and requires that we have good understanding of the astrophysical object and can distinguish between photon-particle spectral oscillation features and other spectral imprints such as atomic lines, edges, and continuum features. In particular, detailed predictions for the spectral signatures of photon-particle mixing are crucial for correctly interpreting the observations. Unlike terrestrial experiments whose setup may be controlled and the results verified or refuted (e.g., PVLAS), an astrophysical experiment cannot be controlled and other corroborative means are required to assess the validity of the results. As we shall see, the observable properties of the photon-particle spectral  oscillation features depend rather sensitively on the physical properties of the object and can therefore be used to our advantage.  This paper is organized as follows: we start with a general layout of the formalism used in this work to calculate the conversion probabilities (\\S 2).  Various relevant physical processes leading to light refraction  are discussed in \\S 3 among which (cold and relativistic) plasma and quantum electro-dynamics (QED) effects. Section 4 focuses on the general properties of the predicted photon-particle oscillation features and study their dependence on the physical parameters of the problem (e.g., on medium stratification). Readers who are less interested in the technical aspect of the problem but wish to better understand the observational consequences for specific astrophysical systems (e.g., magnetars, pulsars, quasars, X-ray binaries, Ap stars, white dwarfs, and cataclysmic variables) may skip to section 5. The discussion follows in \\S 6 where a short \"user's manual\" is provided. Summary follows in \\S 7. ",
        "conclusions": "\\begin{figure*} \\plotone{f19.eps} \\caption{Spectral oscillation features for $m_a=4\\times 10^{-6}$\\,eV particles with $g=10^{-10}\\,{\\rm GeV}^{-1}$ for pulsars (left panel), magnetars (middle panel) and quasars (right panel). Over-plotted are the slope-normalized data for pulsars (Sollerman 2003) and for quasars (Mkn\\,279; Yaqoob \\& Padamabhan 2004). For magnetars we show the normalized and de-reddened broad band photometric data  for two objects (circles and diamonds) taken from Durant \\& van Kerkwijk (2005; error bars are small and cannot be seen and we have not included the marginal $B$-band data). Also shown in dashed line is a model for photon-particle conversion in pulsars/magnetars with $g=3\\times 10^{-11}\\,{\\rm GeV}^{-1}$. The dot-dashed line is a model for photon-particle conversion in quasars with $g=5\\times 10^{-12}\\,{\\rm GeV}^{-1}$. Note that different x-axis scale in each panel.} \\label{g10} \\end{figure*}   We have given detailed predictions for the spectral signatures of photon-particle oscillations in astrophysical objects. While these features appear as absorption features in the spectrum of an object, the physics behind them is {\\it very} different: photons are not a well determined state in highly compact and magnetized environments filled with dilute plasma and a more general definition is that of a photon-particle state where, under conditions prevailing in some objects, could manifest itself as a particle state showing an effective deficit of photons at particular energies across the electromagnetic spectrum.  We have demonstrated that spectroscopic observations of magnetars, pulsars, and quasars can be 2-4 orders of magnitude more sensitive to small values of the coupling constant between particles and the electromagnetic field leading to detectable photon-particle oscillations in cases where little or no signal is expected for current terrestrial experiments. This approach is also much more sensitive than indirect astrophysical constraints, such as cooling of old stars in globular clusters, providing limits on the coupling constant, $g$. In addition, the method presented here, probes particles in a cosmologically relevant mass range where they serve as viable dark matter candidates, and where other constraints (e.g., those related to optical diffuse light or the closure of the universe) are limited.  Unlike terrestrial experiments whose set up may be controlled and their results checked for consistency (for which the PVLAS experiment is an excellent example), astrophysical sources cannot be controlled and their parameters cannot be tweaked by an observer. Moreover, by using compact objects as effective laboratories, the experimental setup is only qualitatively understood. Nevertheless, several lines of reasoning suggest that despite these uncertainties and limitations, the approach explored here may prove very useful: first, the parameter space range probed by different astrophysical environments overlap providing further corroboration for the findings. In addition, the increased sensitivity is not by a factor of a few but rather by several orders of magnitude compared to current methods, hence less susceptible to the details of the system in question (clearly, if our understanding is not even in the order of magnitude range than considerable deviations from the above predictions may occur).    It is {\\it not} the purpose of this paper to undertake an exhaustive search of the parameter space in each class of objects by fitting a large range of spectral models to observations (while marginalizing over model uncertainties in each case and taking into account possible time-dependence of the spectral features). We do, however, wish to demonstrate the feasibility of this approach by showing spectral predictions for each class of objects using the current lowest limits on the coupling constant from CAST and indirect astrophysical arguments of $g=10^{-10}\\,{\\rm GeV}^{-1}$. To this end, we shall assume the aforementioned fiducial parameters for each class of systems considered above, and check the observational implications for the case of $m_a=4\\times 10^{-6}$\\,eV which remains  in the unexplored mass range of the microwave galactic haloscope experiments (whose limits are meaningful only if axions are numerous enough so that they constitute most of the dark matter in our galaxy). We emphasize that one should not draw any conclusions on the presence or absence of photon-particle oscillation features from the following demonstration of feasibility as only a very specific model is considered and we do not marginalize over model uncertainties.  Clearly, the spectral features are pronounced (see Fig. \\ref{g10}) and may be clearly detected even at low resolution and modest signal-to-noise conditions (under our set of assumptions concerning the physical properties of the objects).  A straightforward constraint on the axion properties can from pulsars where a very broad spectral feature is predicted yet is not seen in the data.  Magnetars can, in principal, provide similar constraints given if the densities in their magnetosphere is higher than the Goldreich-Julian value by several orders of magnitude. In this case, the broad features may extend to optical and UV energies (Fig. 19) were data for a few objects are available. Nevertheless, our current understanding of the various emission mechanisms contributing to the emission in these wavebands is at its infancy and different magnetars seem to have very different spectral behaviors (compare the two data sets in Fig. 10). These issues are likely to pose considerable difficulties when interpreting the spectra and attempting to draw robust conclusions of any kind.  At face value, the spectral energy distribution of both magnetars shown is inconsistent with the specific oscillation feature considered here.  For quasar, a broad X-ray feature is predicted yet is not seen in the data (Fig. 10). Interestingly, the oscillation feature, in this case, lies in the part of the spectrum close to the iron $K\\alpha$ line and a more detailed analysis including the effect of atomic features is in order. This is, however, beyond the scope of this paper. If the lack of discernible features in the spectrum is to be taken seriously, then, given the current quality of the data and given our restrictive set of model assumptions, a tentative limit (not marginalizing over model uncertainties!) on the coupling constant of $g<3\\times 10^{-11}\\,{\\rm GeV}^{-1}~(<5\\times 10^{-12}\\,{\\rm GeV}^{-1})$ may be obtained for pulsars and quasars, respectively.  We emphasize that these observations were not conducted to maximize the efficiency for the detection of photon-particle oscillations in those objects and that, in principal, much better data and analysis are required to reach meaningful limits. Specifically, high quality and high resolution X-ray data for quasars as well better understanding of the infrared to optical spectral energy distribution (via photometry and spectra) of magnetars may yield considerably better  limits in this case. We re-emphasize that the above limit on $g$ is given here {\\it only} as a proof-of-concept and applies {\\it only} within our restrictive set of assumptions concerning the physics of the relevant astrophysical objects.  Thus far we have considered pseudo-scalar particles such as the axion. The case of scalar particles is completely analogous to the one considered here with the interchange of ${\\bf e}_\\|$ and ${\\bf e}_\\bot$. By symmetry, all the predictions given here remain valid with the proper transformation. Naturally, the limits which can be obtained on such a class of particles are identical to the case of pseudo-scalar particles.  \\begin{figure*} \\plotone{f20.eps} \\caption{The particle parameter space (spanned by mass and coupling constant) which is probed by the spectroscopic constraints discussed in this work (hatched red surfaces whose difference is related to the assumptions concerning the density of the magnetosphere in magnetars; see \\S5.1.1) as compared to other currently used methods such as laser experiments, microwave experiments,  solar axion telescopes, and indirect astrophysical considerations. Also shown (hatched magenta region) is the range probed by compact objects under the assumption of uniform conditions (see text). Clearly, the method described here can {\\it directly} probe a considerably larger parameter range than is accessible by other methods. The proof-of-concept limits obtained for quasars and pulsars are also shown (for the case of cold plasma; see text).} \\label{total} \\end{figure*}    The higher sensitivity (assuming 5\\% detection threshold) of compact astrophysical objects for probing photon-particle oscillations over an interesting range of particle masses is summarized in figure \\ref{total} and is compared to the regions that can now be probed by other means (CAST, microwave resonance haloscopes, and laser experiments). Also shown is the sensitivity assuming uniform magnetic field and density conditions over a length scale $r_\\star$ across. Overall, significantly larger phase space may be probed by studying the spectra of compact objects which is unreachable by laboratory means. As such, the approach proposed here may allow us to {\\it directly} detect the long sought axion (and/or scalar particles) which provide perhaps the best solution to the strong-CP problem and may also solve the dark matter mystery. It might even be possible to probe the phase space relevant to some axion models.  We have shown that, given current limits on $g$, oscillation features in some objects may be rather broad and could, in principal, be detected by broad band photometry. Reaching down to low values of $g$ may require medium and high resolution spectroscopic observations of compact objects.   Apart from particle physics, there are also implications for the physics of compact objects: should light (pseudo-)scalar particles be detected via the route suggested here or by some other means, the spectral signature (or lack of) would teach us about the physics of inner engines of these objects. In particular, it may teach us about the large scale magnetic field configuration in those objects, the plasma properties in such extreme environments, as well as about their emission mechanism. It may also shed light on their time-dependent properties and could provide invaluable information on their evolution. This provides an additional link between particle physics and astrophysics.     \\subsection{Q\\&As}  Here we attempt to summarize a few of the topics related to the identification and interpretation of photon-particle oscillation features: \\begin{itemize} \\item {\\bf How can one tell photon-particle oscillation features from other spectral features?}  By their shape and variability. Calculations presented here show that the shapes of photon-particle oscillation features may be as narrow as atomic lines or broad as atomic edges or even continuum features. Nevertheless, the shape is, generally, very different from those produced by atomic physics. Nevertheless, secure identification requires good knowledge of the emission and absorption mechanism in the particular object probed. For systems with short variability timescales, distinct energy shifts and shape changes of the feature may be observed. Such variations would be set by the time-dependence of the system properties and are not expected to naturally occur for atomic features. \\item {\\bf What determines the properties of photon-particle spectral features?} Unlike atomic lines and edges, the shape of photon-particle conversion features depends rather weakly on the plasma temperatures for highly magnetized objects (but not for weakly magnetized systems such as quasars). The width of the feature is primarily determined by the stratification of the magnetic field and plasma density in the astrophysical object probed such that more stratified objects are likely to show broader spectral features. The depth of the feature (or its rest equivalent width, $W_0$) depends on the polarization of the emitted radiation with polarized sources potentially showing more extreme cases of conversion (from no conversion at all to conversion of almost 100\\% of the radiation leading to zero flux at certain frequencies). There could be dramatic temperature-induced variations in the properties of the feature  in objects where resonance conversion occurs above the electron cyclotron frequency and when the plasma is hot. An oscillatory pattern extending to higher photon energies can be observed in many cases when high resolution spectroscopy (and short time-frames in varying objects) is used. \\item {\\bf Which wavelengths should be searched?} Our calculations show that the feature may be observed over a very broad energy range. Generally, the relevant wavelength range is determined by the plasma density and magnetic field intensity. For more magnetized objects, the feature will be shifted to longer wavelengths where the refractive index is higher. We emphasize that several spectral features may be observed depending mainly on the composition and temperature of the plasma. In particular, while relatively broad features are expected off-cyclotron resonances, additional narrower features may overlap with electron and proton cyclotron lines. Generally, no features will be observed at photon energies below the plasma frequency where the object is optically thick. Care should be taken when attempting to identify oscillations in regions occupied by atomic features such as lines and edges.  Note that, unlike atomic features, the oscillation feature will shift in energy in objects whose magnetic field strength and density vary with time. This may prove crucial for correctly identifying and resolving oscillation features in compact objects. \\item {\\bf What kind of observations are needed?} This depends on how low a coupling constant, $g$, one wishes to be able to probe and how rapidly varying the astrophysical object probed is. We have shown that, for $g$-values of order the CAST limits, broad band photometry of magnetars and pulsars and high signal-to-noise non-grating X-ray spectra could suffice. Nevertheless, to be able to search for particles (axions) in a previously unexplored regions of the parameter space down to low $g$-values, medium to high resolution spectroscopy may be  required. This is particularly true for the case of magnetars in which a narrow feature is predicted, and for quasars, for which the feature is expected to fall in the (hard) X-ray band. As the X-ray ($<$10\\,keV) spectrum of low luminosity quasars shows various lines and edges (e.g., Chelouche \\& Netzer 2005), high resolution spectra are of paramount importance. As noted before, oscillation features may shift with time, especially in compact objects whose relevant timescales (e.g., dynamical ones) are short, and so multi-epoch short-exposure observations or time-tagging of individual photons may prove helpful in correctly identifying them and characterizing their shapes and positions. \\item {\\bf Can the particle parameters be directly inferred?} Unfortunately, this task is somewhat complicated since, unless one has a good understanding of the object in which efficient photon-particle conversion takes place, degeneracy between the particle parameters and the object parameters is a limitation. Some of the degeneracies may be alleviated if good S/N data are available in which case fitting of model templates to the data may reveal unique features (such as extended broad wings, oscillatory pattern).  It is particularly difficult to constrain the properties of  light particles  whose mass is considerably smaller than the plasma frequency. These very important issues, of extracting reliable particle properties (or limits to them) while marginalizing over model uncertainties, is beyond the scope of this introductory work and will be dealt with in a forthcoming publication. \\item {\\bf What can we learn about compact objects?} In the event that particles, such as axions, are detected by other means (e.g., using terrestrial experiments or in the spectra of several types of objects), then the presence or absence of an oscillation feature may teach us about the interiors of compact objects. As discussed here, the observed properties of photon-particle conversion features are sensitive to the value of the magnetic field, its spatial configuration (hence the effective size of the system), the photon propagation direction and polarization, and to the plasma properties. Therefore, studying those features in the spectra of celestial objects may shed light on the properties of astrophysical magnetospheres and enhance our understanding of such systems. \\end{itemize}"
    },
    "0806/0806.2325_arXiv.txt": {
        "abstract": "We investigate how ``extra'' or ``excess'' central light in the surface brightness profiles of cusp or power-law elliptical galaxies relates to the profiles of ellipticals with cores.  The envelopes of cusp ellipticals are established by violent relaxation in mergers acting on stars present in gas-rich progenitor disks, while their centers are structured by the relics of dissipational, compact starbursts.  Ellipticals with cores are formed by the subsequent merging of the now gas-poor cusp ellipticals, with the fossil starburst-components combining to preserve a dense, compact component in these galaxies as well (although mixing of stars smooths the transition from the outer to inner components in the profiles).  By comparing extensive hydrodynamical simulations to observed profiles spanning a broad mass range, we show how to observationally isolate and characterize the relic starburst component in core ellipticals. Our method recovers the younger starburst population, demonstrating that these dense concentrations survive spheroid-spheroid mergers and reflect the degree of dissipation in the initial mergers that formed the penultimate galaxy progenitors.  The degree of dissipation in the mergers that produced the cusp ellipticals is a strong function of stellar mass, roughly tracing the observed gas fractions of disks of the same mass over redshifts $z\\sim0-2$.  The strength of this component strongly correlates with effective radius at fixed mass: systems with more dissipation are more compact, sufficient to explain the discrepancy in the maximum phase-space densities of ellipticals and their progenitor spirals. The survival of this component together with scattering of stars into the envelope in re-mergers naturally explain the high-Sersic index profile shapes characteristic of very massive core ellipticals. This is also closely related to the kinematics and isophotal shapes: only systems with matched starburst components from their profile fits also reproduce the observed kinematics of boxy/core ellipticals. The final ``core-scouring'' phase of core formation occurs when a black-hole binary formed in the merger scatters stars out of the innermost regions of the extra-light component. It is therefore critical to adopt a physically-motivated profile-decomposition that accounts for the fossil starburst component when attempting to quantify scouring.  We show that estimates of the scoured mass that employ single-component forms fitted to the entire galaxy profile can be strongly biased. ",
        "introduction": "\\label{sec:intro}  Early work to incorporate gas physics and star formation in models of the formation of bulges and elliptical galaxies through mergers demonstrated \\citep{mihos:cusps,mihos:gradients,mihos:starbursts.94} that dissipation in disk-disk mergers leads to central starbursts like those observed in ULIRGs \\citep{joseph85}.  \\citet{mihos:cusps} predicted that this process should leave an observable signature by imprinting a two-component structure into the surface brightness profiles of merger remnants.  That is, the compact stellar remnant of the starburst should have a steeper density profile than the inward extrapolation of the outer \\citet{devaucouleurs} $r^{1/4}$-law shape of the main body of the galaxy.  The ``extra light'' in the starburst remnant is required \\citep{hernquist:phasespace} to explain the high mass and phase space densities of bulges and elliptical galaxies \\citep{ostriker80,carlberg:phase.space,gunn87,kormendy:dissipation} and their fundamental plane correlations \\citep{dressler87:fp,dd87:fp, kormendysanders92,rothberg.joseph:kinematics,hopkins:cusps.fp}.  When predicted by \\citet{mihos:cusps}, extra light was not known to exist in normal ellipticals. \\citet{sersic:profile} functions were replacing $r^{1/4}$ laws as standard machinery to interpret brightness profiles \\citep{caon:sersic.fits}. The steepest central profiles were called ``power~laws'' \\citep{ferrarese:type12,kormendy94:review, lauer:95,byun96:ell.profiles}. Both descriptions could be taken to imply, absent cores or nuclear star clusters,\\footnote{ It is important to distinguish nuclear clusters (stellar ``nuclei'') from extra light components -- we demonstrate the clear distinction in \\citet{hopkins:cusps.ell}. Nuclei are extremely small (typical radii $\\sim1-10\\,$pc, and mass fractions $\\sim10^{-3}\\,M_{\\rm gal}$) relative to starburst/extra light components (radii $\\sim0.2-2\\,$kpc and mass fractions $\\sim 0.1\\,M_{\\rm gal}$). The structural scalings and fundamental plane relations of stellar nuclei are similar to those of globular clusters, nearly perpendicular to and with an opposite physical sense from extra light components. Likewise their stellar population ages, metallicities, and kinematics are very different from extra light and elliptical galaxies \\citep{carollo99:nuclei.scalings, boker04:nuclei.scalings,walcher06:nuclei.ssp,cote:virgo,hopkins:cusps.ell}. } that the brightness profiles of elliptical galaxies are almost power laws with no breaks indicative of two-component structure. However, \\citet{kormendy99} discovered in a small sample of normal ellipticals -- and later \\citet{jk:profiles}, in all the known low-luminosity ellipticals in the Virgo cluster -- that the inner ``power-law'' or ``cusp'' indeed represented a central ``extra light'' component -- the inner profile of a dense, centrally concentrated rise above the inward extrapolation of an outer Sersic profile. \\citet{ferrarese:profiles} and \\citet{cote:smooth.transition} identified qualitatively similar features -- ``luminosity excess'' in the central regions of ellipticals \\citep[for a detailed comparison of these features, see the discussion in][]{hopkins:cusps.ell}. Extension of this to the profiles of recent gas-rich merger remnants observed by \\citet{rj:profiles} showed similar signatures. The ``power law'' profiles identified by previous authors form a part of this extra light; the conclusion that there are two components, however, did not become clear until observations from a variety of telescopes made it possible to create profiles with sufficiently large dynamic range. These authors suggested that the observed extra light is the signature of merger-induced starbursts, as predicted by \\cite{mihos:cusps}, and it is now becoming empirically established that essentially all ``power-law'' or ``cusp'' ellipticals show this behavior.  In \\citet{hopkins:cusps.mergers} (hereafter \\paperone) we showed that observed extra light can indeed be identified with the central density excess produced in simulations of gas-rich mergers.  We developed a formalism to fit surface brightness profiles with two-component models in a manner that accurately recovers the {\\em physical} decomposition between a central dissipational component (the remnant of the merger-induced starburst) and an outer dissipationless component (the product of violent relaxation acting on stars formed in the pre-merger disks before the coalescence of the merging galaxies).  We used this machinery in \\citet{hopkins:cusps.ell} (hereafter \\papertwo) to analyze observed ``cusp ellipticals''\\thinspace\\footnote{The term ``cusp elliptical'' is used in various ways in the literature.  Unless otherwise stated, we use this name to refer to ellipticals that do not have a central resolved core.  In published papers, these are called ``power-law ellipticals'', ``coreless ellipticals'', or ``extra light ellipticals.''}. We showed that this method could be used to recover the dissipational components in cusp ellipticals, and that such components were ubiquitous in the local cusp elliptical population (in samples of $\\sim100$s of such galaxies from \\citet{lauer:bimodal.profiles} and \\citet{jk:profiles}). In essentially every case, simulated dissipative merger remnants provided good matches to the observed profiles of cusp ellipticals. Mutually consistent decompositions of the observed galaxy profiles and the simulated merger remnants demonstrated that the structure of cusp ellipticals as a function of mass corresponds to that predicted for the remnants of gas-rich, spiral-spiral mergers with realistic properties.  That is, the degree of dissipation needed for mergers to explain the densities and scaling laws of ellipticals corresponds well with our empirically recovered starburst components and agrees with the gas fractions available in observed progenitor disk galaxies of appropriate masses.  However, it has been argued that the situation may be different for more massive ellipticals, which appear to exhibit central ``cores,'' and are known to display boxy isophotal shapes and slow rotation in contrast to the disky isophotal shapes and rapid rotation seen in less massive cusp ellipticals \\citep{kormendy:bulge.rotation, davies:faint.ell.kinematics,davis:85,jedrzejewski:87, bender:88.shapes,bender89,bender:ell.kinematics,peletier:profiles}. Cores were first seen in ground-based observations of nearby, high-luminosity ellipticals as central regions of nearly constant surface brightness \\citep{king78,young78,lauer85:cores,kormendy85:profiles,kormendy:cores.review}. They were cuspier than isothermal cores, but the functional form of the density profile as $r\\rightarrow0$ was unknown.  Later, {\\it HST} images showed that nearly all ellipticals have singular starlight distributions in the sense that the surface brightness diverges as $\\Sigma(r)\\sim r^{-\\gamma}$ \\citep{lauer91,lauer92:m32,lauer92,crane93,kormendy94:review,ferrarese:type12,lauer:95}. In low-luminosity, early-type galaxies, $\\gamma$ typically decreases only slowly as the center is approached, and a steep $\\gamma>0.5$ cusp continues in to the {\\it HST} resolution limit. \\citet{lauer:95} classified these as ``power-law'' galaxies.  In more luminous ellipticals, the steep outer density profile shows a robust break to a shallow, inner power-law cusp with $\\gamma<0.3$  \\citep{lauer:95}. The ``break radius'' $r_b$ corresponds to the ``core radius'' $r_c$ measured in ground-based observations \\citep{kormendy94:review}. \\citet{lauer:95} continued to call these ``core galaxies,'' even though the shallow cusps in projected brightness imply steep and singular cusps in luminosity density \\citep{lauer:95,gebhardt96}.  The division of central structure into two families was motivated by the observed bimodal distribution of cusp slopes $\\gamma$ \\citep{gebhardt96,lauer:bimodal.profiles}, but this aspect of central slopes participates in a larger, longer-recognized division of the elliptical population: the typical giant, core elliptical has different physical properties from the typical normal-luminosity, cusp elliptical \\citep[e.g.][]{davies:faint.ell.kinematics,davis:85, jedrzejewski:87,bender:88.shapes,bender89,bender:ell.kinematics, peletier:profiles,kormendybender96,faber:ell.centers, simien:kinematics.1,simien:kinematics.2,simien:kinematics.3, emsellem:sauron.rotation.data,emsellem:sauron.rotation, mcdermid:sauron.profiles,cappellari:anisotropy}. Massive giant ellipticals rotate slowly and have ``boxy'' (rectangularly-distorted) isophotal shapes, characteristic of systems supported by anisotropic velocity dispersions \\citep{schwarzschild:orbit.structure,dezeeuw:orbits,binneytremaine}; \\citet{faber:ell.centers} showed that these relate to the observed ``core'' population. In contrast, less massive ellipticals (and S0 galaxies), where the power-law population predominates, rotate more rapidly; they have more isotropic velocity dispersions and look like they have embedded disks (``disky'' isophotal shapes).  These differences thus led naturally to the idea, developed, in e.g.\\ \\citet[][and references therein]{faber:ell.centers,kormendy99, quillen:00,rest:01,ravindranath:01,laine:03,lauer:centers, lauer:bimodal.profiles,ferrarese:profiles,cote:smooth.transition,jk:profiles}, that disky, rapidly rotating cusp ellipticals are direct products of gas-rich (``wet'') mergers, whereas boxy, slowly rotating core ellipticals have been shaped by subsequent dissipationless (``dry'') re-mergers of two or more (initially cuspy gas-rich merger remnant) ellipticals. Several questions therefore arise.  How did core ellipticals form?  What were their progenitors? It has been shown that mergers of bulgeless disks fail to reproduce the shapes and kinematic properties of these galaxies \\citep[see e.g.][]{barnes:disk.halo.mergers, hernquist:bulgeless.mergers,hernquist:bulge.mergers,naab:gas,cox:kinematics}. Furthermore, if disks are their progenitors, then these systems would not be able to avoid dissipation, because spiral galaxies contain gas.  If lower-mass cusp ellipticals are the progenitors of core ellipticals, what happened to the extra light in those cusp ellipticals?  Nearly all numerical experiments find that light profile shapes are, to lowest order, preserved in dissipationless mergers \\citep[e.g.,][]{boylankolchin:mergers.fp}.  So, core ellipticals should be ``extra light'' ellipticals in the same sense as cusp ellipticals -- i.e. their profile reflects a combination of an outer, low phase-space density violently relaxed remnant of stellar disks, and an inner, compact component originally formed via dissipational processes (in whatever process formed the progenitors, that would allow gas to lose angular momentum and reach these densities in the first place). If their last merger was dissipationless, is the amount of dissipation in the ``original'' spheroid forming merger (i.e.\\ that which formed their progenitors, if these are re-merger remnants) important or relevant to the properties of the $z=0$ galaxy? Is any memory of the merger history preserved, and is there any way to observationally recover this information?  There have been suggestions of interesting behavior in the observations, but they have largely lacked an interpretive context. Models of core galaxies generally presume that the projected stellar brightness decreases monotonically outside the core with a monotonically increasing logarithmic slope -- the fact that many core ellipticals can be reasonably well-fitted by such descriptions has made the multi-component structure of such objects more ambiguous. A closer look at the existing data, however, shows evidence in some galaxies of characteristic features such as inflection points in the slope outside the core, similar to a ``smoothed'' version of the features often seen in cusp ellipticals where the profile transitions from being dominated by a dissipational (``extra light'') component superimposed on a background dissipationless outer envelope \\citep[see e.g.][]{ratcliff82,barbon84,lauer85:cores, lauer:bimodal.profiles,jk:profiles}. Moreover, even where profiles are smooth and monotonic, the inner portions of core galaxies typically rise well above an $r^{1/4}$ law fitted to the envelope. Instead, single profiles fitted to the galaxies are forced to higher Sersic indices that rise more steeply at small radii, reflecting and including the dense, high surface brightness central extra light component, and yielding the high Sersic indices ($n_{s}\\sim6-8$) characteristic of core ellipticals fitted in this manner \\citep[e.g.][]{prugniel:fp.non-homology,trujillo:sersic.fits, ferrarese:profiles,jk:profiles}. \\citet{cote:smooth.transition} point out that there appears to be a smooth transition from low-$n_{s}$ outer profiles with a steep central rise above the corresponding inward extrapolation, to high-$n_{s}$ outer profiles with a corresponding rise implicit in the Sersic profile, and hence not explicitly appearing above this threshold. Without a physical motivation for decomposing this inner rise and outer envelope, interpretation of this phenomenon has been restricted to the empirical notation of the best-fit Sersic indices. In the context of the present work, however, it highlights the potentially composite nature of core galaxies.  It is generally believed that the connection between the merger history of galaxies and their nuclear profile slope (``cusp'' or ``core,'' on scales much less than the effective radius or the scales of the extra light) arises because of ``scouring'' by a binary black hole \\citep[for a review, see][]{gualandrismerritt:scouring.review}. \\citet{begelman:scouring} first pointed out that binary black holes coalescing in a dissipationless galaxy merger stall (i.e.\\ are no longer efficiently transported to the center via dynamic friction) at radii $\\sim$\\,pc, larger than the radii at which gravitational radiation can efficiently dissipate energy and merge the binary -- the so-called ``last parsec problem.'' They noted that significant gas content can provide a continuous source of drag and friction and solve this problem in gas-rich mergers, but that in ``dry'' mergers, the binary will remain stalled for some time, and will harden by scattering stars in the nucleus in three-body interactions. This will continue, flattening the nuclear slope, until sufficient mass in stars ($\\sim M_{\\rm BH}$, by simple scaling arguments) is ejected to merge the binary. It is therefore of particular interest to estimate the stellar mass which must be scattered to explain the slopes of cores, as a test of scouring models and (in such models) a probe of the galaxy merger history. However, such estimates have been ambiguous \\citep[and often controversial; see e.g.][]{ferrarese:profiles,lauer:bimodal.profiles, lauer:massive.bhs,cote:smooth.transition,jk:profiles}, in large part because of the lack of an {\\em a priori} physical model for the profile shape. Understanding the global profile shapes and extra light in core ellipticals is therefore important to reliably estimating how their nuclear profiles have, in detail, been modified by scouring.  It is also important to recognize that there can be both continuity and bimodality in the cusp/core populations. Because the expected number of major mergers in the formation history of a typical elliptical is not large ($\\sim$ a couple), it should be a relatively Poisson process: a significant number of ellipticals (especially at low masses) will have experienced only the original, single major gas-rich merger that transformed them into ellipticals since $z\\sim2-3$ \\citep[see e.g.][]{maller:sph.merger.rates, hopkins:groups.qso,hopkins:groups.ell,somerville:new.sam,lin:mergers.by.type}; others (especially at high masses) will have experienced $\\sim$one or two subsequent major mergers, which will tend to be ``dry'' spheroid-spheroid mergers. Although there might be some intermediate cases, to the extent that properties (such as kinematics and isophotal shapes) are affected by the last major merger, there should be significant differences between those whose last merger was dissipational (with a disk that contains some mass in gas) or dissipationless (spheroid-spheroid).   However, although the last merger may be dissipational or dissipationless, the total amount of dissipation in the formation history -- the mass fraction formed dissipationally, from gas losing angular momentum in mergers and participating in nuclear starbursts -- should be continuous across either population (for e.g.\\ a given mass and original formation time). Although some objects may have had subsequent dry mergers, they still formed stars in a dense central concentration in the original, gas-rich merger that formed the progenitor spheroid; this process will be the same regardless of whether or not the system is destined for a future dry merger. As a function of mass, this should broadly reflect the gas fractions of ultimate progenitor disks at the spheroid formation times, and as such is a continuous function of mass, star formation history, and formation time. This continuity in dissipational content, to the extent that it effects the structure, fundamental plane correlations, and stellar populations of spheroids, is reflected in the continuity of e.g.\\ the fundamental plane \\citep[recently, see][]{ cappellari:fp,bolton:fp.update,hyde:stellar.mass.fp}, the stellar age and metallicity versus mass relations \\citep{trager:ages, nelan05:ages,thomas05:ages,gallazzi:ssps,gallazzi06:ages}, and the color-magnitude relation \\citep{strateva:color.bimodality, baldry:bimodality}. As there has been considerable observational debate regarding the degree of continuity or bimodality between cusp and core populations \\citep[see e.g.][]{ferrarese:profiles,lauer:bimodal.profiles, cote:smooth.transition,jk:profiles}, it is clearly of interest to identify properties that are or are not expected to be continuous across detailed merger and re-merger histories within the spheroid population.  In order to understand the structure and formation history of the core elliptical population, we therefore extend our study of merger remnants and cusp ellipticals in \\paperone\\ and \\papertwo\\ to the core population in this paper. We wish to test the hypotheses that these systems were, in fact, originally formed (i.e.\\ their progenitors were formed) in gas-rich mergers (albeit potentially modified in gas-poor re-mergers), and that the original degree of dissipation can be empirically recovered, and is the critical parameter that can explain their densities, scaling relations, profile shapes, and sizes.  In \\S~\\ref{sec:sims} and \\S~\\ref{sec:data} we describe our set of merger simulations and the observational data sets we consider, respectively.  In \\S~\\ref{sec:profile.evol} we study how light profiles of gas-rich mergers, which we studied in detail in \\paperone\\ and \\papertwo, evolve in subsequent re-mergers of such ellipticals, and demonstrate that our fitting procedures designed to recover the original dissipational/starburst component in gas-rich merger remnants can be applied to re-merger remnants. In \\S~\\ref{sec:gradients} we investigate how gradients in stellar populations, imprinted by the extra light and original gas-rich merger, are affected by re-mergers. Readers interested primarily in our comparison of the properties and scalings of dissipational components in observed core ellipticals may wish to skip to \\S~\\ref{sec:fitting}, where we compare our simulations with and apply our fitted galaxy decomposition to a wide range of observed systems. In \\S~\\ref{sec:scaling} we use these comparisons to study how structural parameters of the outer stellar light and inner extra light component scale with galaxy properties, and compare them with the extra light components in gas-rich merger remnants and cusp ellipticals, and examine how the existence and strength of the extra light component is related to galaxy structure and drives galaxies along the fundamental plane. In \\S~\\ref{sec:kinematics} we consider the global isophotal shapes and kinematic properties of re-merger remnants and how they depend on extra light content. In \\S~\\ref{sec:outer.sersic} we examine how re-mergers and issues of profile fitting relate to the outer Sersic profiles of core ellipticals and re-merger remnants, and compare results obtain with different choices of empirical fitting functions. In \\S~\\ref{sec:missing.light} we demonstrate that this can affect estimates of ``missing light'' on small scales, and explain how this effect arises and what it means for a proper physical understanding of nuclear light profiles. Finally, in \\S~\\ref{sec:discuss} we discuss our results and outline future explorations of these correlations.  Throughout, we adopt a $\\Omega_{\\rm M}=0.3$, $\\Omega_{\\Lambda}=0.7$, $H_{0}=70\\,{\\rm km\\,s^{-1}\\,Mpc^{-1}}$ cosmology, and appropriately normalize all observations and models shown, but note that this has little effect on our conclusions.  We also adopt a \\citet{chabrier:imf} initial mass function (IMF), and convert all stellar masses and mass-to-light ratios to this choice. The exact IMF systematically shifts the normalization of stellar masses herein, but does not substantially change our comparisons. All magnitudes are in the Vega system, unless otherwise specified.   \\breaker ",
        "conclusions": "\\label{sec:discuss}  In \\paperone\\ and \\papertwo, we demonstrated that ``extra light,'' in the sense of stellar populations formed from dissipation in e.g.\\ a gas-rich merger induced starburst, is ubiquitous in cusp ellipticals and gas-rich merger remnants, and studied how its properties are related to and (in some cases) drive those of the galaxy. Here, we show that this dissipational component is expected to survive subsequent re-mergers of those ellipticals, even major gas-poor spheroid-spheroid mergers, in the sense that it will continue to contribute substantially to the central light profile and can be empirically recovered. We apply this to a large observed sample of ellipticals with central cores (i.e.\\ flattening of their light profiles within the central $\\sim30-50\\,$pc), and show that they are consistent with a surviving dissipational ``relic extra light'' component which our adopted empirical fitting machinery and comparison with simulations allow us to recover and interpret.  \\subsection{Comparing Simulations and Observations: Empirical Decomposition of Light Profiles} \\label{sec:discuss:separating}  In \\papertwo, we argued that stars in cuspy ellipticals/gas-rich merger remnants should be separated into at least two distinct populations. First, stars which are formed in the disks (or otherwise in extended distributions in progenitor galaxies) before the final merger and coalescence of the progenitors. The final merger scatters these stellar orbits and they undergo violent relaxation. They dominate the light, even in highly gas-rich merger remnants, outside of $\\sim0.5-1\\,$kpc, and form a Sersic-law profile owing to their partial violent relaxation. Second, the starburst or dissipational population, formed in the central gas concentration in the final merger. This component is compact, and dominates the light inside a small radius $\\lesssim0.5-1\\,$kpc. These stars {\\em do not} undergo significant violent relaxation, but form in a nearly fixed background potential set by the dissipationless component of the galaxy. \\footnote{There is also a third component present in simulations and important for observed kinematics but not prominent in light profile fitting: gas moved to large radii temporarily either by feedback or tidal effects, which settles into the relaxed remnant and re-forms small rotationally supported components \\citep[embedded disks, kinematically decoupled cores, etc.; e.g.][]{hernquist:kinematic.subsystems,hopkins:disk.survival}.} We developed, tested, and studied in detail in \\paperone\\ a two-component fit decomposition that, in simulations, could reliably extract the properties of these two physically distinct components.  Here, we demonstrate that, to lowest order, the two primary components structuring the surface brightness profile are expected to survive in dissipationless spheroid-spheroid re-mergers, in the sense that the physical starburst component still forms a more compact distribution that dominates or contributes significantly to the profile at small radii. Re-mergers will generally expand or puff up the system by a factor $\\sim2$ in size, and smooth out the profile by mixing stellar populations and scattering stars by a typical factor $\\sim0.4\\,$dex in radius (see Figure~\\ref{fig:rf.of.ri}) -- this can make fitting the profile more sensitive to the prescription adopted, and may smooth an obvious break in the profile around the transition from dissipational to dissipationless components (Figures~\\ref{fig:highlownsdemo}-\\ref{fig:demo.ml.appearance}), but does not fundamentally remove the extra light in a physical sense.  We apply our same fitting procedure to simulated re-merger remnants, and find that, despite the re-merger, it is still able to statistically recover the physically meaningful components of the original, spheroid-forming gas-rich merger (i.e.\\ the original dissipational or starburst component and the original dissipationless or pre-starburst component; see Figures~\\ref{fig:demo.rem.appearance}-\\ref{fig:rem.recovery}). In other words, even after a re-merger, the surface brightness profiles of ellipticals retain information about the gas fractions and starburst mass fractions of their original gas-rich mergers. Our parametric fitting form is simple.  We consider the sum of two Sersic laws: an inner extra light component, for which a fixed $n_{s}=1$ works best in a mean sense when the data are not especially constraining in the central regions, but for which $n_{s}$ can be freed if the data are of sufficient quality; and an outer component with a free Sersic index. We explicitly demonstrate that this approach is successful even when (as is common in re-merger remnants) the profile is smoothed by re-mergers and obvious breaks might not be present in the profile (Figures~\\ref{fig:highlownsdemo}-\\ref{fig:dif.fit.qual.hists}).   We apply this to a large sample of ``core'' ellipticals, which (it is often argued) have been modified by re-mergers, and find that it is reliable and implies significant ``extra light'' (by which we mean the remains of the original gas-rich component) in almost all cases (Figures~\\ref{fig:shoulders}-\\ref{fig:lauerpp3}). We also match each of the observed profiles to our library of simulations (of both gas-rich mergers and re-mergers) -- i.e.\\ directly find the simulation mass profile which most closely resembles that observed, and consider non-parametric estimators of the mass in a central starburst component. We find that in all cases we have simulations which provide good matches to the observed systems, comparable to the typical point-to-point variance inherent in the simulation surface brightness profiles ($\\dmu\\lesssim0.1$). The physical starburst components in these best-fitting simulations are closely related to those that we fit directly to the observed profiles, lending further support to our attempt to physically decompose the profiles (Figure~\\ref{fig:extra.vs.sb}). Where available, stellar population models including extended star formation and a subsequent burst independently support our inferred starburst mass fractions (Figure~\\ref{fig:extra.vs.sb}). Likewise, metallicity, age, and abundance gradients, where available, support our decompositions, typically demonstrating a smooth transition to a younger, more metal rich population at the radii where the dissipational component begins to dominate the profile, as predicted in \\papertwo\\ (Figure~\\ref{fig:z.grad.demo}). A complete list of fit parameters and compiled galaxy properties is included in Table~\\ref{tbl:core.fits}.   \\subsection{Predictions and Observations} \\label{sec:discuss:predictions}   Fundamentally, we argue that {\\em all ellipticals -- including those with central cores -- are ``extra light'' ellipticals (dissipational systems)}, insofar as ``extra light'' refers to a component originally formed dissipationally on top of a background of dissipationlessly scattered envelope (a physical two-component nature).  Ellipticals with cores show just as much dissipation, at a given mass, as cuspy or power-law ellipticals. Over the core population, the mass contributed by a central starburst or dissipational population can vary substantially, but in a physical sense, it always exists in addition to the pre-starburst or violently relaxed stellar populations. In terms of their dissipational or starburst population properties, we demonstrate that core ellipticals appear to be, for the most part, a continuous extension of the cusp population. The correlations obeyed by the inferred dissipational component itself are continuous and agree where there is overlap in galaxy properties between the cusp and core populations, as expected in physical models where core ellipticals are formed by re-mergers of cusp elliptical progenitors.  Specifically, we find: \\\\  {\\bf (1)} {\\em The mass fraction in the dissipational or starburst component of both cusp and core ellipticals is a strongly decreasing function of mass} (Figures~\\ref{fig:fgas.needed}-\\ref{fig:mass.vs.fgas}). In detail, the mean starburst mass fractions can be approximated as Equation~(\\ref{eqn:fgas.m}): \\begin{eqnarray} \\nonumber & & \\langle f_{\\rm starburst} \\rangle \\sim {\\Bigl[}1+{\\Bigl(}\\frac{M_{\\ast}}{10^{9.15}\\,\\msun} {\\Bigr)}^{0.4}{\\Bigr]}^{-1}, \\end{eqnarray} with a factor $\\sim2$ scatter at each mass. The trend is similar for both cusp and core ellipticals, and the gas fractions needed span a range bracketed by the typical observed gas fractions of spiral galaxies at the same mass, at $z=0$ (bracketing the low end of the required gas fractions) and $z\\sim2-3$ (bracketing the high end). Core ellipticals have not only preserved their dissipational content, but its fraction reflects that of their ultimate progenitor disks. \\\\  {\\bf (2)} At each mass, the degree of dissipation strongly affects the sizes of the remnants. In both observations and simulations {\\em we demonstrate a tight correlation between effective (half-light) radius at a given stellar mass and the inferred dissipational/extra light fraction} (Figure~\\ref{fig:re.sigma.cusp}). This owes to the compact nature of the central dissipational component -- increasing the mass fraction in this component means that the half-light radius must be smaller.  The correlations obeyed between dissipational content and size at fixed mass are similar for cusp and core ellipticals.  Despite the fact that re-mergers puff up remnants by a factor $\\sim2$, their masses also double, so this only moves them by about $\\sim0.1\\,$dex off the mean size-stellar mass relation obeyed by gas-rich merger remnants (which has an intrinsic scatter $\\sim0.3\\,$dex). On the other hand, we demonstrate that changes in the original starburst or extra light fraction can alter the remnant size by nearly an order of magnitude at fixed mass. Thus, even if core ellipticals have undergone a moderate number of re-mergers, we expect that the degree of dissipation in their original formation should still be the most important factor setting their sizes today, and we demonstrate this in the observations. \\\\  {\\bf (3)} Re-mergers roughly preserve the size-mass relation of the dissipational/starburst component itself, approximately set by the radius at which the starburst component becomes self-gravitating ($G\\,M_{\\rm extra}/R_{\\rm extra} \\approx G\\,M_{\\ast}/R_{e}$); and we do find that the core ellipticals and cusp ellipticals obey nearly the same correlation (Figure~\\ref{fig:sizes}). If re-mergers preferentially puff out low binding energy material (i.e.\\ material at large initial radii), then we might expect that the quantity ($G\\,M_{\\rm extra}/R_{\\rm extra} \\approx G\\,M_{\\ast}/R_{e}$) should increase slightly in a re-merger (by a maximum factor $\\sim2$ in a 1:1 merger, if the inner component is not expanded at all, and the outer component expands by a factor of $2$). We see tentative evidence for such an offset, but it is small (a factor $\\sim1.4-1.5$; Figure~\\ref{fig:size.ratios}). \\\\  {\\bf (4)} In \\papertwo\\ we demonstrated that the extra light component gives rise to stellar population gradients in the remnant. We show here that these gradients are only weakly affected by re-mergers, in agreement with observations which find that cusp and core ellipticals do not show significant differences (at fixed galaxy properties) in their gradient strengths (Figure~\\ref{fig:grad.fx.demo}-\\ref{fig:grad.correlations}). To the extent that both radius and gradient strength scale with the contribution of the original dissipational component at fixed mass, we find that our prediction from \\papertwo\\ should also hold for core ellipticals (namely that at fixed mass, smaller ellipticals should, on average, exhibit stronger metallicity gradients), and early observational evidence appears to support this \\citep[e.g.][see also Figure~\\ref{fig:z.grad.demo}]{mehlert:ssp.gradients, sanchezblazquez:ssp.gradients,reda:ssp.gradients}. There are other lines of evidence for the survival of extra light: \\citet{lauer:centers} find that nuclear color gradients, ellipticities, and isophotal twists in the centers of cusp and core galaxies are continuous and trace similar distributions as a function of mass and luminosity. Analogous correlations should hold with integrated stellar populations, which we discuss in \\papertwo, provided that these are correlated with the gas fractions of the pre-merger disks. \\\\  {\\bf (5)} Given the appropriate dissipational mass fraction inferred from fits to the observed surface brightness profiles, simulated re-merger remnants reproduce the global kinematics and isophotal shape distributions of core ellipticals (specifically rotation $V/\\sigma$, ellipticity, and boxyness $a_{4}/a$), and their mutual correlations (Figures~\\ref{fig:kinematics.cuspcore}-\\ref{fig:kinematics.histograms}). These distributions are {\\em only} reproduced by simulations with matched dissipational mass fractions: systems with too little dissipation (lacking a high central density to make the potential more round and disrupt box orbits) are too elliptical, and systems with too much dissipation remain too disky and rapidly rotating even after several re-mergers (Figures~\\ref{fig:kinematics.before.after} \\&\\ \\ref{fig:kinematics.histograms}). More detailed comparison of simulations with full two-dimensional velocity fields, along with higher-order measures of the velocity field at a given point, are needed to robustly discern whether scenarios as simple as a single re-merger are viable formation mechanisms for slowly-rotating, boxy ellipticals, but similar dissipational content appears to be a basic prerequisite. \\\\  {\\bf (6)} We predict and find that the outer Sersic indices of ellipticals with cores (in our two-component decompositions) trace a roughly constant distribution with median $n_{s}\\sim3-4$ and scatter $\\Delta n_{s} \\sim1$, without a strong dependence on galaxy mass, effective radius, or other properties (Figure~\\ref{fig:ns.mass}). We found similar results in \\papertwo\\ for the cusp elliptical population, but with a lower median $n_{s}\\sim2-3$. The difference arises naturally in our re-merger simulations, as the scattering of stars in subsequent mergers will tend to broaden the outer regions of the light profile, leaving a less steep falloff and giving rise to a higher Sersic index. That there is an offset in the outer Sersic index distribution of these populations suggests that indeed some sizable fraction of core ellipticals have experienced a re-merger, but that the offset is relatively small suggests that there have not been a large number of such re-mergers.  We emphasize that these Sersic indices are not directly comparable to those in previous studies, which fit different functional forms to the light profile or fit e.g.\\ a Sersic profile to the entire light profile (including the extra light). If we fit our simulations to a single Sersic index or core-Sersic profile where it is the formal best fit, we recover a stronger dependence of Sersic index on galaxy mass, luminosity, or effective radius, similar to earlier claims \\citep{caon:sersic.fits,prugniel:fp.non-homology, graham:bulges,trujillo:sersic.fits,ferrarese:profiles}: our models are consistent with the results derived using these methods, however this does not strictly trace the physical dissipationless component of the galaxy. In fact, we show in \\S~\\ref{sec:outer.sersic} (Figures~\\ref{fig:highlownsdemo}-\\ref{fig:dif.fit.qual.hists}) that re-merger remnants are often formally well-fit by core-Sersic laws (Sersic profiles with a central deficit), in agreement with observations of core ellipticals. These cases obtain high Sersic indices, driven in part by the presence of a dissipational component at small radii, raising the central surface brightness (Figure~\\ref{fig:highlownsdemo}). These functional forms are a precise formal parameterization of the light profile, but the best-fit parameters determined in this manner reflect the combination of the dissipational and dissipationless outer components, and do not trivially translate to physical descriptions of the components of the galaxy.  Indeed, as demonstrated in \\papertwo, many of the results in studies using these fitting approaches are actually driven by a dependence of extra light on galaxy properties. The meaningful dependence of outer Sersic index as we quantify it (tracing the dissipationless component) on mass owes to the dependence of typical merger history on mass (dwarf spheroidals and pseudobulges will have lower typical $n_{s}$ than gas-rich merger remnants, which themselves have lower typical $n_{s}$ than re-merger remnants, and these populations, on average form an increasing sequence in mass). For a broadly similar merger history (in terms of e.g.\\ number of major mergers), the profile of the true dissipationless component is expected to be self-similar, as we recover (for both the cusp and core populations), since it is determined purely by gravity. \\\\  These and other correlations argue that core ellipticals should be properly thought of as containing ``extra'' or dissipational components in the same physical (although not necessarily observational) manner as cusp ellipticals; their central, high density regions formed via dissipational processes. We therefore refer to \\papertwo\\ for a number of other proposed observational tests of the models herein, which should hold for core ellipticals as well. This is a critical test of the merger hypothesis and supports the notion that even massive, slowly rotating, boxy ellipticals with cores were originally formed in gas-rich major mergers, albeit potentially modified by subsequent re-mergers.  Points {(1)} and {(2)} emphasize that, even if a system has experienced re-mergers, and even in observed massive core ellipticals, dissipation is the key driver of galaxies along the fundamental plane (in terms of galaxy stellar mass). The degree of dissipation in the original gas-rich merger is the most important factor determining the size of the object and the ratio of baryonic to dark matter within the effective radius of the stellar light, and the correlations obeyed by the observed systems agree well with our simulations.  We show that this degree of dissipation is a systematic function of mass, without a significant offset between cusp and core elliptical populations (as expected if it is set by the gas fractions of the progenitor disks, and there have not been a large number of re-mergers of individual objects). In other words, we are able to demonstrate that the amount of dissipation expected based on known disk gas fractions as a function of mass is precisely that needed to explain the extra light in the surface brightness profiles and to reconcile the densities and radii of disks and ellipticals as a function of mass. We investigate this further in \\citet{hopkins:cusps.fp}, and show how it gives rise to the fundamental plane scalings and ``tilt.'' Once on the fundamental plane, our simulations and other experiments \\citep{boylankolchin:mergers.fp,robertson:fp} find that re-mergers typically move more or less parallel to the fundamental plane, emphasizing that the tilt must arise in the initial, gas-rich spheroid-forming merger.  Although some of these effects will become increasingly scattered or smeared out by a large number of re-mergers, multiple lines of evidence above suggest that the typical core elliptical has undergone relatively limited dry re-merging. If re-mergers are indeed the mechanism by which cores are formed, our comparisons suggest a small number $\\sim1$\\,major re-mergers per object since its formation in a gas-rich merger.  This is in line with other observational indications from e.g.\\ the mass function evolution of ellipticals and red galaxies \\citep[e.g.][]{bundy:mfs,borch:mfs, pannella:mfs,franceschini:mfs,fontana:highz.mfs} and direct observational estimates of the dry merger rate \\citep{lin:merger.fraction,lin:mergers.by.type, vandokkum:dry.mergers,bell:dry.mergers}. It is also generally expected in cosmological models \\citep{delucia:ell.formation,zheng:hod.evolution,hopkins:groups.ell}, where only the most massive BCGs with $M_{\\ast}\\gg10^{12}\\,\\msun$ or so are expected to have a large number of major re-mergers. Even in these cases, it is not clear whether such mergers actually proceed efficiently, or whether the secondary is tidally destroyed before the merger and added as part of an extended envelope or intercluster light \\citep{gallagherostriker72,monaco:sat.destruction, conroy:sat.destruction,purcell:sat.destruction}. At these masses, the growth of the system may also be dominated by a large number of minor mergers, rather than a few major mergers \\citep[major mergers being the dominant ``dry'' growth mode at $\\lesssim$ a few $\\lstar$; see][]{maller:sph.merger.rates,masjedi:cross.correlations}.  Our constraints on this extreme of the population are, unfortunately, much weaker. Moreover, we do not intend our modeling to be extrapolated to this regime. Such systems formed, for the most part, at early times from potentially very different progenitors than those we model, with more complex subsequent merger histories our simulations do not capture \\citep[see, e.g.][]{li:z6.quasar}. More detailed analysis of individual objects, and more detailed models incorporating a fully cosmological context are called for to study the formation history of galaxies in this regime. However, such systems constitute only a small fraction ($\\lesssim10-20\\%$) of the mass density in even the core elliptical population (and just $\\sim3-5\\%$ of the total mass density in ellipticals), so the vast majority of the stellar populations of spheroids should be reasonably represented by our modeling.   \\subsection{Profile Fitting and Tests of Core Creation Models} \\label{sec:discuss:missing}  The existence of nuclear cores in these ellipticals is commonly attributed to ``scouring'' by a binary black hole in gas-poor re-mergers. This will eject stars in close encounters with the binary, flattening the initial steeply-rising power-law cusps into the observed shallow cores. To the extent that the stars in this initial cusp are ejected to larger orbits, leaving a core, these remnants can be thought of as having ``missing'' light in their centers. We emphasize that this is completely consistent with there also being ``extra light,'' as we define it, in these galaxies. The ``extra light'' we refer to is the remnant of stars formed in dissipational central star formation events, and dominates the profile within rather large radii $\\sim$kpc, blending smoothly onto the outer dissipationless profile. Scouring will flatten the nuclear peak or cusp of this dissipational remnant, but will not be likely to remove the $\\sim10\\%$ of the stellar mass that constitutes the dissipational starburst relic. Despite the terminology, ``extra'' and ``missing'' light are not mutually exclusive. In a strict physical sense, all ellipticals are dissipational/``extra light'' ellipticals -- they have all experienced some dissipational star formation. Of these dissipational ellipticals, there are ``un-cored'' or ``un-scoured'' (``power-law,'' ``cusp,'' or ``cuspy core'') ellipticals (presumably objects whose last major merger was gas-rich, allowing the black holes to merge quickly, and forming the central cusp via new star formation) and ``cored'' or ``scoured'' ellipticals (presumably the remnants of subsequent spheroid-spheroid ``dry'' re-mergers).  The ``missing light'' is a core -- a flattened nuclear profile as opposed to a continued steeply rising cusp -- within the very center of the dense dissipational component.  Our methodology is robust to re-mergers and scouring: our two-component profiles still recover the dissipational component in simulations allowing for any reasonable model of nuclear scouring. However, we demonstrate that the same profile can be well-fitted by different assumed functional forms, and the parameters obtained from these fits -- particularly in regards to the ``extra'' and ``missing'' light -- are not trivially related and require careful interpretation. In particular, observations find that core ellipticals can be reasonably well-fit in a formal statistical sense by cored Sersic laws (an outer Sersic profile with an inner ``flattening'' or deficit), generally with high Sersic indices $n_{s}\\gtrsim5$, and we obtain consistent results fitting our re-merger simulations with these methods. There is no explicit ``extra'' component in these fits: rather, in our simulations, we find that the high $n_{s}$ values -- which extrapolate to high central densities before the flattening or deficit -- implicitly reflect the extra light (that being the central, dense component that enables the high-$n_{s}$ fit; otherwise the system would have a high outer $n_{s}$ but a low central density, not resembling any Sersic profile). Because the re-merger has smeared out the light profile, removing characteristic features (or kinematic subsystems) that might have (in the original gas-rich merger remnant) been a more obvious indication of the transition to radii where stars formed dissipationally dominate the profile, the physical breakdown of these systems is not obvious from a strictly empirical standpoint. By calibrating different fitting methods with our simulations, we have attempting to provide an interpretive context and physical motivation for specific interpretations of these fits.  In a related manner, this raises some cautions regarding the purely empirical methods sometimes used to estimate how a profile has been modified by core scouring -- i.e.\\ how much mass is ``missing'' from the nuclear region (relative to the steep nuclear rise of the progenitor cusps). This is a subtle issue. There may of course be no light actually missing after scouring -- rather, stars are scattered from small radii to large, flattening the central profile (what is really desired is an estimate of how much stellar mass has been scattered -- i.e.\\ how much stellar mass must be moved from the nuclear region to larger radii to explain the difference in profiles). So any such estimate is sensitive to the model for the nuclear profiles of the progenitors -- equivalently, what the profile would be in the absence of scouring.  We explicitly consider one such (commonly adopted) model, where the implicitly assumed progenitor profile is based on the inwards extrapolation of an outer Sersic law (as in e.g.\\ core-Sersic fits). In these fits, the effects described above (effects on the profile at larger radii than the core itself) can lead to a very large outer Sersic index in the formal best fit; in the high-$n_{s}$ regime ($n_{s}\\gtrsim6$), the nature of the Sersic profile is such that these profiles then rise steeply at small radii (more steeply even than the power-law nuclear profiles of cusp ellipticals), and a small change in the outer profile (raising $n_{s}$) can substantially raise the extrapolated (and assumed ``pre-merger'') inner profile. This can bias the estimate of the ``missing light'' fraction towards very high values ($\\sim1-5\\%$) and literally interpreted can make it appear as if the ``core'' extends to radii $\\sim$kpc. In fact, we find that adopting this methodology in such cases, we would infer ``missing light'' at this level even in simulations where there is no physical scouring (i.e.\\ there is no actual ``missing light''). Moreover, if real, these values would imply scoured masses $\\sim10-50\\,M_{\\rm BH}$, in troubling disagreement with models of scouring which predict scoured mass deficits $\\sim0.1-1\\,M_{\\rm BH}$ per major re-merger \\citep[see e.g.][]{milosavljevic:core.mass, merritt:mass.deficit,sesana:binary.bh.mergers}. If stars from centrophilic orbits in triaxial potentials allow the binary to coalesce rapidly (providing stars from near $\\sim R_{e}$, where the loss of even $\\sim 10\\,M_{\\rm BH}$ makes no difference to the profile, to harden the binary), as suggested in some idealized calculations \\citep{berczik:triaxial.bh.mergers,holley:triaxial.loss.cone} and merger remnant simulations \\citep{hoffman:prep}, the difficulty for the models in explaining such large ``missing masses'' grows.  Given the steep dependence of the implicitly assumed central progenitor profile on Sersic index in this regime, and the appearance of ``missing light'' even in simulations without scouring given a literal interpretation of this fitting methodology, it is likely that the ``problems'' for the scouring models in these extreme cases reflect more the observational uncertainty regarding the appropriate ``un-scoured'' profile, rather than fundamental uncertainties in the physics of scouring. Future work, as observations and models improve, should attempt to carefully model the progenitor profiles and develop estimators that are less sensitive to the profile shape at radii much larger than the core itself. An important test for any such estimators should be that re-merged progenitors with initially cuspy profiles (down to the desired resolution limits), without scouring included, should typically yield little or no ``missing light.'' We briefly consider a couple of possibilities for such estimators of the scoured mass (less sensitive to the large scale radii), and obtain (re-analyzing the observations) typical scoured mass estimates of $\\sim0.5-3\\,M_{\\rm BH}$, in better agreement with scouring models.  Further study, in particular observation of the nuclear regions $\\sim R_{\\rm BH}$ is needed to test models of scouring and test the accuracy of different estimators for the scoured stellar mass. For example, scouring is expected to preferentially eliminate stars on radial orbits and leave a bias for tangential orbits within the radius affected \\citep[e.g.][]{quinlan:bh.binary.tang.orbit.bias}. \\citet{gebhardt:nuclear.anisotropies} see tentative evidence for this in a limited sample of ellipticals; the major-axis radii within which the effect appears are generally $\\sim0.5-3\\,R_{\\rm BH}$, as expected in scouring theories and our revised inferences from core profiles. A preliminary comparison supports our conclusions (and caveats) here, but the number of relevant ellipticals is small. If similar observations can be obtained for a sample of ellipticals with alternative estimates of the missing mass fraction, it can be determined whether some estimators are biased towards putting too much or too little of the profile into the scoured component.   \\subsection{Summary} \\label{sec:discuss:summary}  We have developed a paradigm to understand the structure of both cusp and core ellipticals, in which there are fundamentally two stellar components: a dissipational central starburst component and a more extended violently relaxed component. We have shown that the separation between these components can be inferred with observations of sufficient quality, and used to understand the formation history of ellipticals as a function of a wide range of properties. This allows us to demonstrate that dissipation is critical to understanding the properties of ellipticals, including (but not limited to) the structure of their surface brightness profiles, their sizes, ellipticities, isophotal shapes and rotation, age, color, and metallicity gradients (and their evolution), and the gas content and properties of their progenitors.  In particular, we argue here that this remains true for ellipticals with cores and re-merger remnants -- in other words, {\\em all ellipticals, including core ellipticals, are fundamentally ``extra light'' ellipticals}. The core ellipticals form a continuous family with the cusp ellipticals in terms of their dissipational content, as expected if these cusp ellipticals are their progenitors and if the number of re-mergers (if important) has not been large ($\\sim1-3$ major re-mergers) for the typical core elliptical. We demonstrate that, despite the possibility that these systems have expanded via re-mergers, the degree of dissipation in the original gas-rich merger (the memory of which is retained in the surface brightness profile) remains the most important factor determining the size, gradient strength, and other properties of the remnant.  We have studied the properties and identified robust trends of dissipational stellar remnants in the nuclei of elliptical galaxies with cores and remnants of gas-poor, spheroid-spheroid re-mergers, across a large set of simulations, in which we vary e.g.\\ the galaxy masses, initial gas fractions, concentrations, halo masses, presence or absence of bulges, presence or absence of black holes, feedback parameters from supernovae and stellar winds, orbital parameters and disk inclinations, and mass ratios of the merging galaxies. This range of parameters allows us to identify the most important physics.  As we found in \\paperone\\ and \\papertwo, the most important factor determining the structure of the remnant (insofar as the properties we consider are concerned) is how much mass is in the original (in the original or last spheroid forming, gas-rich merger) dissipationless (violently relaxed) component, versus the mass fraction in the dissipational (starburst) component. Orbital parameters and initial galaxy structure can, in principle, affect the remnant surface brightness profile significantly, but only indirectly, insofar as these help to set the amount of gas which will be available at the time of the final coalescence of the galaxy nuclei (i.e.\\ how much mass ends up in the starburst component, as opposed to being violently relaxed in this final merger).  Re-mergers will expand the original gas-rich merger remnants, and smooth the profile (scattering stars substantially about their mean final radii), but, as it is well established that dissipationless mergers conserve (in a mean sense) particle rank order in binding energy and therefore radius \\citep{barnes:disk.halo.mergers}, they will preserve these components. In other words, the product of a re-merger is, to lowest order, the sum of the two progenitor spheroid dissipationless components (constituting the nearly self-similar outer violently relaxed component) and their inner dissipational/starburst components (constituting the central dissipational/starburst remnant component of the final re-merger remnant, despite the fact that there might be little or no new dissipation in the re-merger).  We have demonstrated that this makes predictions for how fundamental plane scalings arise, which we study further in \\citet{hopkins:cusps.fp}.  We make a wide range of new predictions for the distributions of these properties and how they scale with the degree of dissipation, and how they should scale with each other and various other observational proxies for this degree of dissipation (which we define herein). We have predicted and shown (given these proxies) that dissipation is indeed more important (contributing a larger mass fraction) in low-mass ellipticals, in line with expectations based on how gas fractions are known to scale with disk mass.  Testing all of these with better observations should be possible in the near future, with well-defined samples of ellipticals and continued improvements in mapping e.g.\\ the surface brightness profiles, stellar populations and their gradients, and structural properties of ellipticals over a wide dynamic range.  Given our decompositions, we observationally confirm the long-standing prediction of the merger hypothesis, that sufficient dissipation should have occurred in the inner regions of ellipticals to explain the discrepancy between their central densities and those of their progenitor spirals, a confirmation that fits well in line with what is now well-established in gas-rich merger simulations and is also directly seen in progress in ongoing/recent mergers, which have (through clear recent central star formation) raised their phase space densities to be comparable to ellipticals \\citep{kormendysanders92,Doyon94, Genzel01,rothberg.joseph:kinematics, tacconi:ulirgs.sb.profiles,dasyra:pg.qso.dynamics}. We show that this is true for the core elliptical population just as we demonstrated it for merger remnants in \\paperone\\ and cusp ellipticals in \\papertwo. In other words, the {\\em same} mechanism can explain the central densities in both cusp and core ellipticals -- while core ellipticals may be modified by subsequent re-mergers, no alternative formation mechanism for them (in the sense of formation of their spheroid progenitors) is required.  We demonstrate some important caveats in observational studies of core ellipticals. These ellipticals are multi-component (dissipational plus dissipationless, in the physical sense above) galaxies, but they have shallow nuclear profiles (central cores) instead of steeply rising central cusps, and can therefore be thought of as ``missing light'' galaxies in the sense that some process (e.g.\\ scattering of stars from a black hole binary) has flattened the nuclear profile slope to create the core. This is fundamentally different from the processes involved in the evolution of the ``extra light'' (by which we mean the remnant of the starburst/dissipational star formation) and they should not be confused -- whatever process forms cores likely acts on small scales and involves a relatively small fraction of the galaxy mass, much less than the $\\sim$kpc scales and $\\sim10\\%$ of the galaxy mass that are characteristic of the dissipational component. Owing to the smoothing of the outer profile in a re-merger, re-merger remnants can often be better fit (in a purely formal sense) by e.g.\\ single Sersic profiles with central deficits or core-Sersic laws: where observed, we argue that this is an indication of such re-mergers, but we stress that it does not mean the galaxies are not, in fact, two-component objects in a physical sense (simply that the fitted function reflects some combination of the two components). For this reason, the two-component profiles adopted here lend themselves to more direct physical interpretation and comparison with galaxy formation models. Furthermore, we demonstrate that care is needed when using such core-Sersic fits as an estimator of the ``scoured'' nuclear mass -- when the outer Sersic index is large, the inferred ``missing mass'' can be quite sensitive to the details of the profile at large radii. This may explain estimates of scoured masses which are much larger than those predicted by models of core creation.  There are, of course, other changes to galaxy properties in re-mergers. As discussed in \\paperone\\ and \\papertwo, gas which survives the original spheroid-forming gas-rich merger will quiescently settle into the galaxy and form kinematic sub-components (in particular, embedded disks and kinematically decoupled nuclear components). We demonstrated that these do not contribute significantly to the surface brightness profile, and therefore they are not evident in our analysis. However, \\citet{cox:kinematics} and other numerical studies of gas-rich mergers \\citep{naab:gas,burkert:anisotropy} have demonstrated that these subsystems can contribute strongly to the isophotal shapes (in particular driving the diskyness of the remnant) and kinematics, yielding distributions of shape and kinematic properties (including ellipticity, isophotal shape $a_{4}/a$, rotation $(v/\\sigma)^{\\ast}$, anisotropy, triaxiality, and kinematic misalignments) in good agreement with observed cusp or disky, rapidly rotating $\\sim\\lstar$ ellipticals.  The evolution of these sub-systems in re-mergers is not readily apparent in the surface brightness profiles of the remnants -- their contribution to the profile is much less than the smoothing effects and variations introduced by the re-merger (shown in \\S~\\ref{sec:profile.evol}). However, their evolution in re-mergers is briefly discussed in \\S~\\ref{sec:kinematics} and will be studied in detail in \\citet{cox:remerger.kinematics}, who show that they are generically destroyed by major spheroid-spheroid re-mergers. This explains the significant effects on the global kinematic properties of re-merged ellipticals shown in \\S~\\ref{sec:kinematics} -- making them rounder, boxier, and more slowly rotating, in agreement with the observed properties of massive core or boxy, slowly rotating ellipticals. But effects on the velocity field in detail (as a function of radius, at higher order in the velocity moments, and in spatial distribution rather than just azimuthal average) will be more pronounced, and present opportunities not just to test whether or not re-mergers are good analogs to boxy/slowly-rotating ellipticals, but to distinguish the extent to which substructure and kinematic subcomponents contribute to these rotation/isophotal shape properties, as opposed to e.g.\\ global angular momentum content or features of the spheroid potential. Kinematic misalignments, kinematically decoupled subsystems, triaxiality, and trends of isotropy with radius are also more sensitive to orbital parameters \\citep[see e.g.][]{cox:kinematics,boylankolchin:dry.mergers, jesseit:kinematics} and may distinguish merger histories with preferences for certain orbital configurations (possible if e.g.\\ massive galaxies accrete minor companions preferentially along filaments). These more detailed, ``second-order'' structural parameters may therefore represent a more sensitive probe of the re-merger history (i.e.\\ the history of subsequent mergers after the original, spheroid forming gas-rich merger), while the ``first-order'' parameter (the light profile) retains a memory of the degree of dissipation imprinted in the {\\em original} gas-rich merger."
    },
    "0806/0806.1326_arXiv.txt": {
        "abstract": "{The magnetic topology of the solar photosphere in its quietest regions is hidden by the difficulties to disentangle magnetic flux through the resolution element from the field strength of  unresolved structures. The observation of spectral lines with strong coupling with hyperfine structure, like the observed Mn\\,{\\sc i} line at 553.7 {\\rm nm}, allows such differentiation.} {To analyse the distribution of field strengths in the network and intranetwork of the solar photosphere through inversion of the Mn\\,{\\sc i} line at 553.7 {\\rm nm}.} {An inversion code for the magnetic field using the Principal Component Analysis \\textit{(PCA)} has been developed. Statistical tests are run on the code to validate it. The code has to draw information from the small-amplitude spectral feature appearing in the core of the Stokes V profile of the observed line for field strengths below a certain threshold, coinciding with lower limit of the Paschen-Back effect in the fine structure of the involved atomic levels.} {The inversion of the observed profiles, using the circular polarization (V) and the intensity (I), shows the presence of magnetic fields strengths in a range from 0 to 2 {\\rm kG}, with predominant weak strength values. Mixed regions with mean strength field values of 1130 and  435 Gauss are found associated with the network and intranetwork respectively.} {The Mn\\,{\\sc i} line at 553 nm probes the field strength distribution in the quiet sun and shows the predominance of weak, hectoGauss fields in the intranetwork, and strong, kiloGauss fields in the network. It also shows that both network and intranetwork are to be understood at our present spatial resolutions as field distributions of which we hint the mean properties.} ",
        "introduction": "The absence of observables of the magnetism of the quiet sun is the main handicap in the determination of the magnetic topology of this region covering most of the solar photosphere. Over-interpretation of the scarce information contained in the usually observed spectral lines in magnetometry, and its comparison with numerical simulations of magnetoconvection,  has been the subject of long disputes about the nature of those fields. The addition of further observables, able to constrain the models and to unveil hidden ambiguities and biases in the diagnostic techniques was mandatory in such a situation. This work is a contribution in that direction, with a further step forward in the use of Mn lines with a strong hyperfine coupling.  In a short descriptive summary of the previous studies of the magnetism in quiet Sun regions,  we find from one side, the Fe lines  in the visible region whose circular polarization amplitudes are almost uniquely sensitive to magnetic flux, (e.g.  Keller et al. 1994; S\\'{a}nchez Almeida \\& Lites 2000; Lites 2002, Dom\\'{i}nguez Cerde\\~na et al. 2003; Orozco Su\\'{a}rez et al. 2007),   and on the other side, we find  the inversions of the Fe lines in the near-IR domain   whose profiles are almost universally split by Zeeman effect (Lin 1995; Lin \\& Rimmele 1999; Khomenko et al. 2003). Traditionally, studies based on the visible spectral lines concluded on the presence of fields with kG strengths and small surface coverage, while studies based on infrared lines favoured weaker hG fields and more spread in surface.  Recently, the picture has been made more realistic by considering a continuum distribution function for magnetic field strengths at each resolution element in the quiet sun \\cite{socas_SA_03}, instead of a single magnetic vector value. Additionally,  a  very interesting new approach has enriched the individual line analysis in quiet regions: the inversions of simultaneous and co-spatial observations of the Fe lines in the visible (630 nm) and in the near-IR (1.5 $\\mu$m).  Two such studies have already been carried out  founding  divergent results. In the first of these works, Dom\\'{i}nguez Cerde\\~na et al. (2006)  have employed a model with three magnetic components ({\\sc mismas}  model, S\\'anchez Almeida \\& Landi Degl'Innocenti 1996) to simultaneously invert the lines in both spectral ranges. The authors  retrieved field  distributions in the second and third magnetic components with peaks in the strong regime and they conclude that the kG fields contribution dominates the magnetic flux and energy transport.  In  the other of these studies, Mart\\'{i}nez Gonz\\'{a}lez et al. (2008) showed  that when  only the inversions of the Fe lines in the visible are considered, the magnetic distribution is dominated by strong strength fields (kG). However, when simultaneously near-IR and visible inversions were both considered, the kG contributions disappears  and the distributions  are mostly dominated by weak strength fields (hG). This last  conclusion is  in accord to the results obtained from near-IR inversions alone (see previous references). These results, suggesting a probable bias of the quiet sun Fe inversions when performed only in the visible at 630 nm had also been  found on numerical tests  \\cite{bellot_coll_03, marian_fe_06}.  In the counterpart, a possible bias in the inversions of the IR alone, has been also argued suggesting that this lines are slightly sensitive to the strong strength fields, such that when  mixed  and not resolved magnetic structures of hG and kG are both present, the inferred strengths of hG are preferred over those of kG (e.g. S\\'{a}nchez Almeida \\& Lites 2000; Dom\\' inguez Cerde\\~na et al. 2006).  Such discussions and differing results illustrate the difficulty and ambiguity of the measurements attempted, and have prompted the search for more observables.    Apart from those Zeeman-based techniques discussed above, the description should be completed with  the Hanle effect diagnostics,  pioneered by  Stenflo (1982). A paradigm of  not structured and  turbulent fields, e.g.  \\cite{manso04},  could result in  apparent contradiction with the tube-like structures used to describe Zeeman-based observations if one insists in the picture of a single vector magnetic field per resolution element.    Mn lines subject to strong hyperfine coupling appeared as an interesting observable to be added to the previous ones: their amplitudes are dependent on magnetic flux as the usual Fe lines in the visible regime, but under certain magnetic field regimes a new spectral feature appears in the Mn profile. Just the presence or absence of such a spectral feature allows the observer to determine the presence or absence of fields with strengths below a certain threshold thus adding a new observable on field strength to be compared to measurements with near-IR lines.  Quiet sun conditions limit the information on magnetic fields to circular polarization profiles, whose amplitudes stay too near to the usual signal-to-noise levels found in solar polarimetry. The amplitudes of the linear polarization  Stokes parameters use to be  below the noise level  at the typical spatial resolutions, and are thus useless for analysis purposes;  and intensity profiles are insensitive to the weak magnetic fluxes involved. Consequently one does not only need to add further spectral lines to the amount of observables, but also do it under observing conditions that guarantee very high signal-to-noise levels in polarimetry. We used the telescope THEMIS for that purpose in this work and we explored the capabilities of the Mn lines to provide further information on the quiet sun magnetic strength distribution. We simplified the formation of the Mn lines with a radiative transfer model using a Milne-Eddington atmosphere; the model  proves to be at the limit of present possibilities with yet too many of its parameters left undetermined by the available observables. We nevertheless prove through synthetic tests that we can ascertain the distribution of field strengths in strong and weak flux regions  (photospheric network and intranetwork respectively) with good accuracy, within the limits of the adopted Milne-Eddington modeling for the Sun.  With the results presented here, that show a field strength distribution with maximum number of occurrences at very weak fields (less than or around 100 G),  we propose in the conclusions a cartoon scenario that would be consistent with  both the Zeeman-based measurements (visible and near-IR) and the Hanle measurements.  In section 2, we review the signatures of the magnetic field strength in the line formation of the Mn\\,{\\sc i} profiles. In section 3, we describe  the inversion code used in this work to retrieve the strength and magnetic flux from the data, and  we test its capabilities with realistic noise conditions. The observations are described in section 4 and  the results presented in section 5. Conclusions and a final overview are included in the final section. ",
        "conclusions": "In this work  we contribute to  explore the diagnostic capabilities of Mn lines with strong hyperfine coupling for the magnetic topology of the quiet sun. We built an inversion code for the Mn line at 553.7 nm formed in a Milne-Eddington atmosphere, and tested it under the conditions typical of the quiet sun in terms of signal amplitudes and noise levels. Such conditions are highly constraining: in most of the cases the magnetic field information is contained exclusively in the Stokes V profile, the linear polarization signals being under noise levels and the intensity profile being dominated by the non-magnetic portion of the resolution element. Under such restrictive conditions the number of independent observables for the magnetic field is highly reduced \\cite{ARetal07} and it sets a limit to the number of free parameters of our model. Previous works have shown that the Mn lines increased the number of observables thanks to the spectral features appearing under certain magnetic regimes. The inversion code we tested in this work takes advantage of it and allows us to improve over the simplistic weak-field approximation of previous works into a more sophisticated Milne-Eddington atmosphere and yet to conserve the possibility of disentangling strong from weak fields.  Since most of the other parameters in the Milne-Eddington model, aside from field strength and flux, were not determined accurately under quiet-sun conditions, we doubt that any more sophisticated model can be used in the diagnostics at this point without the addition of further independent observables (e.g. with other spectral lines). However, more sophisticated models can shed light into implicit biases in such unrealistic picture of the solar photosphere as a Milne-Eddington atmosphere is, e.g. \\cite{SA_hfs_08}.  We have payed special attention throughout this work to the numerical tests made to ascertain whether, under present observational conditions of the quiet sun, the weak fields were disentangled from the strong fields for similar net fluxes. The answer is affirmative and therefore we proceeded to apply the code to real data taken with the THEMIS telescope under good seeing conditions. Even in a telescope like THEMIS, focused on performing the most sensitive polarimetry possible, the acquisition of the data at the required signal-to-noise levels implied a trade-off in spatial resolution (seldomly better than 1 arc sec) that limits our results. We therefore constrained ourselves to separate the photospheric network from the intranetwork in our study, and restrained of identifying smaller scales, like granules and intergranules.  To identify network from intranetwork we applied a rather poor rule fixing at $10^{-3}$ the smallest amplitude of the Stokes V profile respect to the continuum intensity for a point to belong to the photospheric network. Such a rule obviously classifies as network any high flux concentration, mostly found over the network, but not totally absent from the intranetwork. Conversely it calls intranetwork any low flux concentrations, mostly found all over the intranetwork but not absent from the network either. With such bias in mind we decided to keep the names of network and intranetwork for the two classes because a small change in such threshold value did not change the general shape of the histograms, but also because of the following reasons.  After inversion, the intranetwork points show a distribution of field strengths whose dominant feature can be easily fitted with an exponential tail. We use the noun \\textit{tail} because an exponential distribution would result in a non-zero probability for null fields, something unacceptable for a vector field. Such a distribution, showing that the quiet sun regions are mostly permeated of weak strength fields, would be in full agreement with previous results in both Hanle and Zeeman effects \\cite{manso04, LinRim99, khomenko03, MG_IRvsVis_08} and reveals a random or disorganised vector magnetic field at scales much smaller than present spatial resolution, if Hanle and Zeeman diagnostics are to be coherent with each other and with a field strength whose highest probability remains below 100G \\cite{SAEC03}.   The network points, on the other hand, show a bimodal distribution: the exponential tail found in the intranetwork distribution is still present, peaking at field strengths below 100G. But on top of that disorganised magnetic regime we observe the appearance of stronger fields, in the kiloGauss regime. This second distribution of fields would be in agreement with the presence of one or several concentration mechanisms, able to stabilise, organise and perhaps amplify the statistical fluctuations of the ubiquitous turbulent field in those places where photospheric dynamics or sheer accidents are able to maintain a temporal coherence bigger than the typical magnetic diffussivities  \\cite{parker_82}. One can heuristically think that the strong down drafting plumes in the vertex of the convection cells are such places \\cite{nordlund92,rast_03}, thus defining the magnetic photospheric network as a non-continuous succession of bright network points \\cite{MR_92} where magnetic fields are organised in vertical structures and made stable by the downdrafting plasma. Other smaller plumes found in mesogranular or granular scales \\cite{rast_03} may also survive long enough to allow the statistical concentration of high fluxes, to be found here and there throughout the intranetwork and that, in our analysis, would have been included in the network group.   The cartoon just drawn above is in agreement with the histograms measured in the present work over the quiet sun. In particular, we are prone to  the presence of an ubiquitous random field, as inferred from Hanle diagnostics, and whose statistical fluctuations can be seen as Zeeman signatures over deep magnetograms, but otherwise highly disorganized at present resolutions except when plasma movements create conditions, as around convection plumes, stable enough to concentrate and build up net magnetic flux.  In any case, with the inversions performed here on the Mn {\\sc i}  in the visible region of the spectra, at 553 nm, we have found that the distributions of the strength fields in the quiet Sun regions are dominated in occurrence by weak strength fields except in those areas related with the network."
    },
    "0806/0806.1921_arXiv.txt": {
        "abstract": "We present measurements of the low-temperature thermal conductivity of a number of polymeric and composite materials from 0.3 to 4~K.  The materials measured are Vespel SP-1, Vespel SP-22, unfilled PEEK, 30\\% carbon fiber-filled PEEK, 30\\% glass-filled PEEK, carbon fiber Graphlite composite rod, Torlon 4301, G-10/FR-4 fiberglass, pultruded fiberglass composite, Macor ceramic, and graphite rod.  These materials have moderate to high elastic moduli making them useful for thermally-isolating structural supports. ",
        "introduction": "Cryogenic instruments frequently require rigid mechanical support of significant mass while minimizing conductive heat flow. In particular, optical systems require very stiff structural mounts.  To increase the stiffness of a structure the designer generally has the choice of using more material or choosing materials with larger elastic moduli. However, materials with higher elastic moduli tend to have higher thermal conductivity.  This situation motivates the use of the ratio of elastic modulus to thermal conductivity as the figure of merit for thermally-isolating structural materials.  This work was motivated by the need to support the cryogenic microwave receiver of the \\spider balloon experiment\\citep{montroy2006,mactavish2008,crill2008}.  The \\spider instrument insert has massive components at 4.2, 1.4, 0.35 and 0.25~K.  Each of these temperature stages are thermally and structurally referenced to an outer liquid helium cryostat.  The cryostat will tip in elevation from zenith to horizon and the cryogenic support structures must minimize the gravitational deflection of the focal plane. Historically, we have used polyimide supports such as Dupont Vespel SP because of its low thermal conductivity \\citep{locatelli1976,olson1993}.  Advanced polymers like Vespel are generally quite expensive.  At the time of this writing, 1/2$^{\\prime\\prime}$ Vespel rods cost approximately half the price of gold by weight. Because the \\spider instrument will include six duplicate focal plane structures, the cost and availability of the material required to support them is a significant concern.  There are a number of materials properties that are potentially important when designing thermally-isolating structural supports.  As mentioned previously, low thermal conductivity and high modulus are desirable attributes.  The elastic modulus of materials increases as they cool (for example, see \\citet{flynn}).  This increase in stiffness is moderate in metals but can be significant for for some polymers.  In particular, the elastic modulus of Teflon increases by approximately a factor of 20 when cooled from room to cryogenic temperatures \\citep{corruccini1957}.  In addition to stiffness, the strength of materials is important for structures that may experience large forces or shocks.  The stability of the dimensions of the structure upon cooling may also be important.  In this case, materials with small coefficients of thermal contraction, such as carbon fiber composites \\citep{reed1997}, may be desirable. For larger structures the support members themselves may comprise a significant fraction of the mass, in which case the density may become important.  The enthalpy of the materials, particularly those with low thermal conductivity, may also be a significant concern since they will take longer to cool down. ",
        "conclusions": "We have measured the thermal conductivity of a number of polymeric and composite materials between 0.3 and 4~K and fit the results to a parametric model.  We then compare the ratio of the room--temperature elastic modulus and the low--temperature thermal conductivity to identify candidate materials for thermally--isolating structural support members.  By this metric, a number of the inexpensive materials performed better than much more expensive advanced engineering polymers across the relevant temperature range.  Of particular note is the pultruded carbon fiber Graphlite rod, which distinguishes itself at $T\\sim1-4$~K, as well as Macor ceramic at low temperatures."
    },
    "0806/0806.4390_arXiv.txt": {
        "abstract": "Spitzer IRAC mid-infrared photometry is presented for the globular cluster (GC) systems of the NGC~5128 (``Centaurus~A'') and NGC~4594 (``Sombrero'') galaxies.  Existing optical photometric and spectroscopic are combined with this new data in a comprehensive optical to mid-IR colour catalogue of 260 GCs.    Empirical colour-metallicity relationships are derived for all optical to mid-IR colour combinations.  These colours prove to be very effective quantities to test the photometric predictions of simple stellar population (SSP) models.  In general, four SSP models show larger discrepancies between each other and the data at bluer wavelengths, especially at high metallicities.  Such differences become very important when attempting to use colour-colour model predictions to constrain the ages of stellar populations.  Furthermore, the age-substructure determined from colour-colour diagrams and 91 NGC~5128 GCs with spectroscopic ages from \\citet{2008MNRAS.386.1443B} are inconsistent, suggesting any apparent GC system age-substructure implied by a colour-colour analysis must be verified independently.  Unlike blue wavebands, certain optical to mid-IR colours are insensitive to the flux from hot horizontal branch stars and thus provide an excellent metallicity proxy.  The NGC~5128 GC system shows strong bimodality in the optical R-band to mid-IR colour distributions, hence proving it is bimodal in metallicity.  In this new colour space, a colour-magnitude trend, a ``blue tilt'', is found in the NGC~5128 metal-poor GC data.  The NGC~5128 young GCs do not contribute to this trend.  In the NGC~4594 GC system, a population of abnormally massive GCs at {\\it intermediate metallicities} show bluer optical to optical colours for their optical to mid-IR colours, suggesting they contain extended horizontal branches and/or are younger than typical GCs.  Analysis of optical to mid-IR colours for a ultra-compact dwarf galaxy suggests its metallicity is just below solar. ",
        "introduction": "Massive star cluster formation takes place during major star bursts.  The age substructure of a particular star cluster system therefore provides a historical record of significant star formation events in its host galaxy.  Globular clusters (GCs) are a class of massive star clusters with a mean mass of a few $10^{5} M_{\\odot}$ and are notable for their compactness.  Observational evidences generally suggests the bulk of GC formation took place within roughly 2 gigayears of the Big Bang (see references in Brodie \\& Strader 2006).  Systems of GCs are known to host a blue and red GC subpopulations in terms of their photometric colour (e.g. Peng et al. 2006; Strader et al. 2006).  This has been interpreted as evidence for a bimodal metallicity distribution, hence most GC systems contain both a metal-poor (or blue) and metal-rich (red) GC subpopulation.  Given the close proximity of the GC formation epoch to the time when the universe was free of metals, the existence of two GC metallicity subpopulations likely implies that two distinct phases of GC formation took place at different epochs.  Obtaining absolute age dates of the two star formation events (and subsequent ones) can help constrain the detailed formation and assembly history of a galaxy.  This has encouraged efforts to age-date GCs and consequently a number of techniques have been employed to overcome the age-metallicity degeneracy \\citep[e.g. ][]{1994ApJS...95..107W}.  Analysis of absorption line indices from spectroscopy is one such technique and most of the few hundred or so GCs with good data are found to be as old as Galactic GCs \\citep[see summary of observations in ][]{2006ARA&A..44..193B}.  Unfortunately, the lengthy telescope time required for such analysis is costly, hence the number of GCs studied in this manner remains remarkably low.  A notable exception is the recent work of \\citet{2008MNRAS.386.1443B}, who have obtained spectroscopic ages and metallicities for $\\sim150$ GCs in the massive elliptical galaxy, NGC~5128.  Existing spectroscopic studies are also biased towards the brightest GCs, an important caveat that should be re-emphasised given the discovery of a GC colour-magnitude trend \\citep{2006ApJ...636...90H,2006AJ....132.2333S,2006AJ....132.1593S,2006ApJ...653..193M}.  On the other hand, the well-studied Milky Way GC system exhibits no such trend with GC luminosity.  GC colour data can also help overcome the age-metallicity degeneracy.  This involves a comparison between observational colour-colour data for individual GCs and theoretical stellar population model predictions.  Optical to infrared (IR) colours are essential for this work, as the degeneracy among optical colours is severe for ages $\\ga2$Gyr (e.g. Lee et al. 2007a).  From analysis of optical and near-IR photometry, GC system age substructure has been reported in NGC~1316 \\citep{2001MNRAS.328..237G}, NGC~4365 (\\citealt{2002A&A...391..453P,2005ApJ...634L..41K}; however c.f. \\citealt{2005A&A...443..413L,2005AJ....129.2643B}) and NGC~5846 \\citep{2003A&A...405..487H}.  In most instances, a subpopulation of ``intermediate-age'' GCs ($\\sim2-8$ Gyr) that make up a small fraction ($\\sim15\\%$) of the total GC system is claimed.  A notable exception is the GC system of NGC~4365, which might contain $\\sim40-80\\%$ of these younger GCs (Puzia et al. 2002; though cf. Larsen et al. 2005 \\& Brodie et al. 2005).  The systems for which no age substructure was reported using near-IR photometry include:  M87, NGC~4478 (Kissler-Patig, Brodie \\& Minniti 2002), NGC~3115 \\citep{2002A&A...391..453P}, NGC~4594, NGC~3585 and NGC~4472 (Hempel et al. 2007).  However, the data samples for each of these systems are typically small:   on order of a few tens of GCs per system.  Optical to IR colour data is also well-suited for characterising GC metallicity distributions, hence can help constrain the chemical properties of galaxies at early epochs.  Such colours span very large dynamical ranges and allow more precise metallicity measurements.  For example, 0.1 dex in metallicity corresponds to $\\sim0.03$ mag in B--R while 0.1 dex is $\\sim0.07$ mag in B--K (according to the preliminary models of Charlot \\& Bruzual 2008 priv. comm.).  The IR is also insensitive to hot horizontal branch stars.  This is a important point considering the cautionary work of Yoon, Yi \\& Lee (2006) who suggested GC colour bimodality might entirely result from a rapid transition between blue and red horizontal branches at intermediate GC metallicities.  The above issues are addressed here with an observational study of the Spitzer Space Telescope \\citep{2004ApJS..154....1W} IR Array Camera \\citep[IRAC; ][]{2004ApJS..154...10F} mid-IR properties of GC systems in the NGC~4594 and NGC~5128 galaxies.  NGC~4594, the ``Sombrero galaxy'' is classified as an Sa-type spiral, but has a galaxy stellar mass-normalised number of blue GCs \\citep[T$_{blue}$; ][]{1993MNRAS.264..611Z} comparable to massive, early-type galaxies (Spitler et al. 2008).  With an estimated $1900$ GCs \\citep{2004AJ....127..302R}, the NGC~4594 galaxy perhaps hosts the largest GC system within the local volume.  Recent Hubble Space Telescope (HST) Advanced Camera for Surveys (ACS) imaging revealed bi-modality in the optical colours of the GC system \\citep{2006AJ....132.1593S}, confirming previous work in the optical:  \\citet{1999AJ....118.1526G}, \\citet{2001MNRAS.327.1116L}, and Rhode \\& Zepf (2004).  Co-added spectra from 14 NGC~4594 GCs were found to be consistent with old ages by \\citet{2002AJ....124..828L}.  Comparing V--I and V--K photometric data of 26 NGC~4594 GCs to stellar population models, Hempel et al. (2007) concluded these GCs are predominantly old.  NGC~5128, ``Centaurus A'', is a peculiar S0 galaxy whose GC system shows optical colour bimodality \\citep{2001A&A...369..812R,2004ApJ...602..705P}.  Harris et al. (2006b) estimated the GC system contains $1550\\pm390$ GCs.  A intermediate-{\\it aged} subpopulation making up no more than 15\\% of the GC system total was recently reported by \\citet{2008MNRAS.386.1443B}.  In addition to providing spectroscopic evidence for the metal-poor and metal-rich GC subpopulations, \\citet{2008MNRAS.386.1443B} also found a small, intermediate-{\\it metallicity} GC subpopulation.  In the following, the distances to NGC~4594 and NGC~5128 galaxies are taken to be m$-$M$=29.77\\pm0.03$ or 9.0 Mpc \\citep{2006AJ....132.1593S} and m$-$M$=27.74\\pm0.14$ or 3.5 Mpc \\citep{2007ApJ...654..186F}, respectively. ",
        "conclusions": "In this work, a catalogue of 260 GC candidates with good optical and Spitzer [3.6] photometry is presented, enabling, for the first time, an examination of the optical to mid-IR properties of a large sample of GCs.  Following previous examples in the literature using K-band photometry, the age-substructure of these GC systems is examined with colour-colour diagrams including the [3.6]--band.  A spread of ages in the NGC~5128 metal-rich GC subpopulation might have been reported from photometry alone.  However, the youngest GCs found in the large spectroscopic study of \\citet{2008MNRAS.386.1443B} do not occupy a young region in colour-colour space and actually show photometric properties that are consistent with the old NGC~5128 GCs.  This observation likely demonstrates that the age-metallicity degeneracy is indeed difficult to overcome when using only typical photometric observations.  Spectroscopic age confirmation is required for any age-substructure to be interpreted from old stellar population in such colour-colour diagrams.  For the NGC~4594 GC system, for which no large sample of GCs with spectroscopic ages exists, the metal-rich GCs generally fall onto the the 5--7 Gyr SSP model predictions in \\br versus \\rf.  However, an empirical comparison between data and SSP models (see below) suggests the models likely under-predict the $B-$band flux relative to the data at these metallicities.  Until this discrepancy is understood, the apparent subpopulation of intermediate-aged NGC~4594 metal-rich GCs will remain unconfirmed.  Another notable feature in the NGC~4594 GC colour-colour diagram (Fig.~\\ref{figbrrch1}) is an apparent non-linearity traversal of the parameter space, which resemble the SSP model predictions of S. Yoon (2006, priv. comm.; Y06).  Driving this resemblance are a subpopulation of very luminous and extended GC candidates first noted by Spitler et al. (2006).  These objects have \\br colours similar to two Galactic GCs that host very extended HB morphologies (Rich et al. 1997; Busso et al. 2007).  These Galactic and the luminous NGC~4594 GCs are bluer in \\br than expected for their metallicities, perhaps suggesting they share similar HB morphologies (see Fig.~\\ref{figmw}).  Another possibility is that the massive NGC~4594 GCs are $\\sim5$ Gyr younger than the bulk of the GC system.  This would explain their high luminosities, but not the extended sizes they tend to show (Spitler et al. 2006).  A dwarf nuclei with a young stellar component that accreted onto NGC~4594 might match these characteristics, although testing this scenario with simulations is required.  Using the excellent mass proxy, \\f, and the largest colour dynamical range studied thus far for subpopulation colour-magnitude trends, the present work demonstrates the NGC~5128 blue GC subpopulation shows redder colours at higher luminosities.  The NGC~4594 sample does not span a large enough range of the GC luminosity function to constrain such trends.  The NGC~5128 blue GCs show a mass-metallicity proportionality of $Z\\propto M^{0.19}$; noticeably weaker compared to other galaxies.  This confirms initial suspicions of Spitler et al. (2006) that the strength of the blue tilt decreases with the host galaxy luminosity.  No such trend is found among the red GCs in NGC~5128.  The \\citet{2008MNRAS.386.1443B} NGC~5128 GC spectroscopic ages allow, for the first time, the blue tilt to be studied independent of any young to intermediate-aged GC.  It is concluded that age substructure in not likely playing a role in the blue tilt.  The scenario of Yoon et al. (2006) is tested with the \\rf GC data.  Such colours are insensitive to hot HB stars (see the Y06 models in Fig~\\ref{figspitmetal}), hence can determine whether GC colour bimodality implies metallicity bimodality without effects from a rapidly changing HB morphology.  While the NGC~4594 sample is too small to constrain this scenario, the old GCs in NGC~5128 provide strong evidence for \\rf colour bimodality, which can safely be interpreted as metallicity bimodality.  This confirms the spectroscopic analysis of NGC~5128 GCs by \\citet{2008MNRAS.386.1443B}.  Strader et al. (2007) and Kundu \\& Zepf (2007) provide further evidence for GC system metallicity bimodality in massive ellipticals, suggesting the Yoon et al. (2006) scenario does not apply to extragalactic (nor Galactic) GC systems.  To conclude, despite a low IRAC pixel resolution and relatively short exposure times (e.g. 240s -- NGC~4594, 72s -- NGC~5128), the present work demonstrates that interesting science in the field of extragalactic GC system astronomy is indeed possible with Spitzer IRAC imaging.  The largest prerequisite for such work is a good optical GC catalogue, because Spitzer photometry alone cannot be used for contamination removal for such faint objects."
    },
    "0806/0806.1260_arXiv.txt": {
        "abstract": "We report X-ray imaging of the powerful FR~II radio galaxy 3C~353 using the {\\it Chandra} X-ray Observatory. 3C~353's two $4$\\arcsec-wide and $2$\\arcmin-long jets allow us to study in detail the internal structure of the large-scale relativistic outflows at both radio and X-ray photon energies with the sub-arcsecond spatial resolution provided by the VLA and {\\it Chandra} instruments. In a $90$\\,ks {\\it Chandra} observation, we have detected X-ray emission from most radio structures in 3C~353, including the nucleus, the jet and the counterjet, the terminal jet regions (hotspots), and one radio lobe. We show that the detection of the X-ray emission associated with the radio knots and counterknots, which is most likely non-thermal in origin, puts several crucial constraints on the X-ray emission mechanisms in powerful large-scale jets of quasars and FR~II sources. In particular, we show that this detection is inconsistent with the inverse-Compton model proposed in the literature, and instead implies a synchrotron origin of the X-ray jet photons. We also find that the width of the X-ray counterjet is possibly narrower than that measured in radio bands, that the radio-to-X-ray flux ratio decreases systematically downstream along the jets, and that there are substantial (kpc-scale) offsets between the positions of the X-ray and radio intensity maxima within each knot, whose magnitudes increase away from the nucleus. We discuss all these findings in the wider context of the physics of extragalactic jets, proposing some particular though not definitive solutions or interpretations for each problem. In general, we find that the synchrotron X-ray emission of extragalactic large-scale jets is not only shaped by the global hydrodynamical configuration of the outflows, but is also likely to be very sensitive to the microscopic parameters of the jet plasma. A complete, self-consistent model for the X-ray emission of extragalactic jets still remains elusive. ",
        "introduction": "\\label{sec:intro}  Extragalactic jets constitute the longest collimated structures in the Universe. They transport huge amounts of energy from the nuclei of active galaxies out to kpc or Mpc distances, significantly affecting the properties of the surrounding intracluster/intergalactic medium. Jets have been extensively studied in the radio domain on different scales since the very beginning of the development of modern radio interferometers (Bridle \\& Perley 1984, Begelman et al. 1984). More recently, the excellent spatial resolution of the {\\it Chandra} X-ray Observatory (and, to a lesser extent, of other X-ray satellites like {\\it XMM-Newton}) has allowed us to image large-scale structures in powerful extragalactic radio sources at X-ray frequencies as well, and thus has opened a new era in studying the high energy emission of these objects. More than 100 radio-loud AGNs are now known to possess X-ray counterparts to their radio jets, hotspots or lobes on kpc-to-Mpc scales (e.g., Harris \\& Krawczynski 2002; 2006, Stawarz 2003, Hardcastle et al. 2004, Sambruna et al. 2004, Kataoka \\& Stawarz 2005, Marshall et al. 2005a, Croston et al. 2005, Tavecchio et al. 2005, Hardcastle 2006, and references therein). The X-ray emission observed from the extended lobes is well understood and modeled in terms of the inverse Comptonization of the cosmic microwave background photons by the low-energy electrons (IC/CMB; see the discussion in Kataoka \\& Stawarz 2005, and Croston et al. 2005), providing strong evidence for approximate energy equipartition between the radiating electrons and the lobe magnetic field (with the particle pressure dominating over the magnetic pressure by up to one order of magnitude). However, the origin of the $0.1-10$\\,keV radiation detected from the large-scale jets and hotspots is still widely debated and, to some extent, controversial.  The X-ray emission of the terminal regions (the hotspots) of \\emph{powerful} jets is consistent with the synchrotron self-Compton (SSC) emission model, in which radio-emitting electrons accelerated at the terminal shock inverse-Compton-scatter synchrotron radio photons to keV energies (e.g., Hardcastle et al. 2004, Kataoka \\& Stawarz 2005, and references therein). Modeling of X-rays from this emission process allows us to extract several crucial jet parameters. In particular, the observations imply rough equipartition of the energy densities stored in radiating electrons, $U_{\\rm e}$, and magnetic field, $U_{\\rm B} \\lesssim U_{\\rm e}$, and it has been argued that they indicate a dynamical role for non-relativistic protons in the outflow (see Stawarz et al. 2007 for the case of the radio galaxy Cygnus A). On the other hand, the X-ray emission of the hotspots in \\emph{low-power} jets is in disagreement with the predictions of the simple SSC model (see the discussion in Hardcastle et al. 2004). The complex X-ray morphology and spectral character of the broad-band hotspot emission in objects like 3C~227 or 3C~327 suggest instead that we are seeing synchrotron emission from high-energy electrons accelerated continuously in the extended and turbulent jet termination regions (Hardcastle et al. 2007a, Fan et al. 2008). The origin of the bright X-ray knots within the jets themselves (hereafter `jet knots') is also a matter of a debate. In nearby, low-power FR~I sources (like M~87 and Centaurus~A), the typical radio-to-X-ray spectra of the jet knots are consistent with a single smoothly broken power-law continuum, indicating synchrotron radiation from a single electron population extending up to the highest energies ($\\gamma \\equiv E_{\\rm e} / m_{\\rm e} c^2 \\sim 10^8$; e.g., Marshall et al. 2002, Hardcastle et al. 2003). Yet the particular acceleration processes involved, the multi-component character of the high-energy emission, and the main factors determining the spectral shape of the broad-band jet emission, are far from being understood (see the discussions in, e.g., Kataoka et al. 2006, Laing et al. 2006a, Honda \\& Honda 2007).  The most controversial issue, however, is the origin of the intense X-ray emission detected from knots in powerful quasar jets, such as PKS~0637$-$752 or 3C~273 (Schwartz et al. 2000, Marshall et al. 2001, respectively). Here the X-ray knot spectra are much brighter than expected from a simple extrapolation of the radio-to-optical synchrotron continua, indicating that an additional or separate spectral component dominates the jet's radiative output at high (X-ray) photon energies. Very often, this emission is modeled in terms of inverse-Comptonization of the CMB photon field by low-energy ($\\gamma \\lesssim 10^3$) electrons (Tavecchio et al. 2000, Celotti et al. 2001), which typically requires highly relativistic jet bulk velocities on kpc-Mpc scales in order to maintain rough equipartition between electrons and magnetic field and to fulfill the minimum power condition. The derived jet bulk Lorentz factors $\\Gamma_{\\rm jet} \\gtrsim 10$ are comparable to those inferred for the pc-scale jets; (Harris \\& Krawczynski 2002, Sambruna et al. 2004, Marshall et al. 2005a). Unfortunately, the global dynamics of large-scale jets in powerful sources are largely unknown, and different arguments in favor of and against highly relativistic bulk velocities on kpc-Mpc scales (obtained by means of the analysis of the radio or optical jet emission; e.g., Wardle \\& Aaron 1997, Hardcastle et al. 1999, Scarpa \\& Urry 2002), are not conclusive. Thus, the IC/CMB model may be considered to be strong support for the idea that these outflows do indeed propagate from sub-pc scales with little energy dissipation, efficiently transporting energy in the form of bulk kinetic motion, and depositing it far away from the active nuclei (Tavecchio et al. 2004; 2007, Sambruna et al. 2006; but see the discussion in Hardcastle 2006). This important conclusion is often claimed to be consistent with the observed one-sideness of the extragalactic X-ray jets detected by {\\it Chandra}, since jet one-sideness is a simple and natural consequence of relativistic beaming.   On the other hand, if there is significant beaming in powerful jets on large scales, then the detection of bright X-ray jet emission from FR~II radio galaxies, which are believed to be analogous systems to radio loud quasars but to be viewed with the jets at large angles to the line of sight, should be considered as unlikely. Such emission has, however, been detected in several objects (e.g., 3C~303, 3C~15, Pictor~A, or 3C~403; see Kataoka et al. 2003a,b, Hardcastle \\& Croston 2005, and Kraft et al. 2005, respectively). Obviously, detection of any X-ray counterjet would be of primary importance in this respect, since it would automatically exclude significant beaming, and thus impose very severe constraints on the jet emission models. In FR~I sources, where the emission mechanism is generally supposed to be synchrotron, possible or likely detections of X-ray counterjets have been reported in 3C~270 and Centaurus~A (Chiaberge et al. 2003, Hardcastle et al. 2007b, respectively). In the latter case the counterjet detection is almost certain, since \\emph{extended} X-ray emission coincident with the \\emph{extended} radio features in the receding jet rules out a claim that all of the observed X-ray counterparts to the radio counter-knots may be due to chance coincidence of the jet-related radio features with X-ray binaries of the Centaurus~A host galaxy. However, in more powerful sources, there has been no definitive counterjet detection to date. Possible counterjets have been reported in the intermediate FR~I/FR~II object 4C~29.30 (Sambruna et al. 2004), as well as in the broad-line FR~II radio galaxy Pictor~A (Hardcastle \\& Croston 2005). In these cases, it is unclear whether the X-ray emission is non-thermal in nature, because with the available very limited photon statistics any detailed spectral analysis is impossible, and in addition it is rather difficult to claim one-to-one morphological correspondence between the X-ray and radio structures (especially since the counterjet in the prototype FR~II source Pictor~A is not detected at radio frequencies). In X-ray maps of the classical double (narrow-line FR~II radio galaxy) Cygnus~A, linear features aligned --- but not co-spatial --- with the radio jet and the counterjet have been noted, but their direct connection with the radio jet plasma is extremely vague (Smith et al. 2002, Steenbrugge et al. 2008). It is then possible that in this, and some other possible X-ray counterjet sources, the aligned X-ray structures on the counterjet side are due to thermal radiation of shocked galactic/intergalactic gas interacting with the jet plasma, as observed in the distant ($z=2.2$) radio galaxy PKS~1138-262 (Carilli et al. 2002). We note in this context that Cygnus~A is located in the very center of a rich cluster environment, while the restarting 4C~29.30 radio galaxy is known to have clear signatures of the interaction between the radio jets and the ambient line-emitting gas (van Breugel et al. 1986, Jamrozy et al. 2007). A clear spatial association between the radio and X-ray jets and the capability to distinguish non-thermal and thermal emission are the key requirements for a convincing counterjet detection in a powerful source.  An alternative to the inverse-Compton model as an interpretation of the X-ray emission of powerful jets involves synchrotron emission from high-energy electrons, which, in order to be consistent with the constraints on optical emission, must often either be characterized by a `non-standard' (concave) energy distribution, or must constitute a separate population to the ones emitting the radio-to-optical continuum. It has been noted that such an electron population can arise due to the continuous and efficient stochastic acceleration processes expected to take place within the extended jet volumes (Stawarz et al. 2004), and will be seen preferentially within turbulent jet boundary layers with significant velocity shear (Stawarz \\& Ostrowski 2002). A very strong support for the synchrotron hypothesis was recently provided by detailed multiwavelength observations of the jet in the quasar 3C~273 (Jester et al. 2006, 2007, Uchiyama et al. 2006). The observations showed in particular that the X-ray spectra are significantly softer than the radio spectra in most regions of the 3C~273 outflow, and that they are compatible with extrapolating the X-ray power law down to the polarized (and therefore synchrotron in origin) UV/optical continuum. These findings are in strong disagreement with the predictions of the IC/CMB model. In addition, a detailed analysis of the broad-band emission of several other quasar jets supports the synchrotron hypothesis, pointing out an important role of the jet velocity structure (consisting of a fast spine and a slower boundary layer/outer sheath) in shaping the jet high-energy radiation (e.g., Hardcastle 2006, Jester et al. 2006, Siemiginowska et al. 2007). Such jet velocity structure is in fact always seen in numerical modeling of the evolution and propagation of extragalactic relativistic jets (e.g., Aloy et al. 1999, Leismann et al. 2005, Mizuno et al. 2007), since it results inevitably from the non-linear growth of Kelvin-Helmholtz instabilities on the jet surface (see the recent studies by Perucho et al. 2007, Hardee 2007, Meliani \\& Keppens 2007, and references therein). It is not yet established, however, if the jet boundary layers are indeed the places of the enhanced acceleration of high-energy particles. Nor is the exact velocity profile at the jet boundaries precisely known; in fact, there are reasons to believe it may be far from the widely expected monotonic one, and may even exhibit a sharp increase in the bulk velocity of the outflow at the jet edges (Aloy \\& Rezzolla 2006, Mizuno et al. 2008). The radiative signatures of such `anomalous' shear layers, in the context of the high-energy emission from extragalactic large-scale jets, were recently investigated by Aloy \\& Mimica (2008).   Observational studies of the velocity structure of extragalactic jets and its relationship to their broad-band emission are hampered by the difficulties in resolving the outflows transversely, especially at high energies. Detailed radio studies and sophisticated modeling of the polarization and total intensity in the radio of several nearby FR~I sources (like 3C~31, NGC~315, or 3C~296; Laing \\& Bridle 2004, Laing et al. 2006a,b, respectively) do reveal jet radial velocity structures, and indicate that the flatter-spectrum radio regions are indeed often associated with the jet boundary shear layers. Unfortunately, only a few kpc-scale low-power jets, e.g., the one hosted by the radio galaxy Centaurus~A, can be resolved at X-ray frequencies by the {\\it Chandra} instrument. The Cen A observations show interestingly limb-brightened X-ray morphology, but there is no obvious relation with the spectral shape of the synchrotron X-ray continuum (Hardcastle et al. 2003, 2007b, Kataoka et al. 2006, Worrall et al. 2008). In particular, the diffuse component of the Centaurus~A X-ray jet is characterized by a constant spectral index across the jet (Kataoka et al. 2006), while the X-ray spectra of isolated knots seem to be steeper within the jet sheath when compared to the X-ray spectra of the knots within the spine (Worrall et al. 2008). Recently, two `intermediate' FR~I/FR~II jets in the BL Lacertae objects 3C~371 and PKS~2201+44 were also resolved transversely in X-rays (Sambruna et al. 2007). The observations indicated a synchrotron origin of the detected keV photons from all the jet components, and suggested slightly steeper X-ray spectra of the jet edges when compared with the jet spines. Interestingly, the width of the X-ray jets in these two sources is not smaller than the widths of their optical and radio counterparts. Needless to say, no quasar or powerful FR~II jet has been resolved till now at X-ray frequencies. This is true even for the particularly bright and wide 3C~273 jet, which, although imaged in detail at optical and radio frequencies (Jester et al. 2005), cannot be resolved transversely at higher photon energies. For this source, however, the limits to the width of the X-ray jet are smaller than the sized measured in the optical and radio bands at the position of the bright knots (Jester et al. 2006), but not necessarily in the interknot regions (see Marshall et al. 2005b).  The FR II radio galaxy 3C~353 provides one of the best-known examples of FR II jets that can be resolved at both radio and X-ray wavelengths. 3C~353 ($z = 0.0304$) is the fourth strongest radio source in the 3C catalog, with a total flux density $S_{\\nu} \\approx 57$\\,Jy at $1.4$\\,GHz, and a projected size $\\sim 4.5\\arcmin$. It exhibits hotspots and a pair of large-scale FR~II-type jets, clearly visible within filamentary lobes (Swain et al. 1998). The jets in 3C~353, constituting about $1\\%$ of the entire source luminosity, are well collimated trains of knots with an average width $\\approx 4\\arcsec$ and jet-counterjet radio brightness asymmetry $\\approx 2$ (Swain 1996). Both total and polarized intensity profiles across the jets indicate that the bulk of the jet radio emission is produced at the jet edges, and so presumably within a boundary shear layer (Swain 1996, Swain et al. 1998). Since these jets are among only a very few FR~II jets wide enough to be resolved in X-rays, we planned and conducted a deep {\\it Chandra} observation of 3C~353, in order to investigate the multiwavelength structure of powerful FR~II outflows. We report on the results of these observations in the present paper.  \\begin{figure} \\begin{center} \\includegraphics[angle=0,scale=0.43]{f1a.eps} \\includegraphics[angle=0,scale=0.43]{f1b.eps} \\caption{Above: $1.4$\\,GHz VLA image of 3C~353 at $1.3''$ resolution (adapted from Swain et al. 1998). Below: a zoom in on the eastern jet ($0.44''$ resolution at $8.4$\\,GHz) shows that the jet is well resolved, and clearly extended compared to the unresolved nucleus. }\\label{fig:rad_img} \\end{center} \\end{figure}   \\begin{table*} \\small{ \\caption{X-ray sources detected with \\textsc{WAVEDETECT} along the 3C~353 jet.} \\label{tab:wavdetect} \\begin{center} \\begin{tabular}{lllcccc} \\tableline X-ray & associated & nearby & RA (J2000)$^a$ & DEC (J2000)$^a$ & Net counts$^b$ & Significance ($\\sigma$)$^d$\\\\ jet region & X-ray features & radio knot & & & & \\\\ \\tableline\\tableline EJ1 & E21/E23 & J1 & 17 20 29.7 & $-$00 58 41.6 & 46.9$\\pm$7.2 & 13.7 \\\\ EJ2 & E70/E73 & J4 & 17 20 32.8 & $-$00 58 26.6 & 12.0$\\pm$4.0 &  3.8 \\\\ EJ3 & E88 & H & 17 20 33.8 & $-$00 58 23.7 & 32.2$\\pm$6.3 &  8.2 \\\\ WJ1 & W47 & CJ2 & 17 20 25.0 & $-$00 58 55.1 & 32.0$\\pm$6.0 & 10.0 \\\\ WJ2 & W120a/b & ... & 17 20 20.3 & $-$00 59 09.9 & 12.3$\\pm$3.7 &  4.8 \\\\ \\tableline \\end{tabular} \\tablenotetext{}{Note: all the errors are $1\\sigma$.} \\tablenotetext{a}{Coordinate center of the detected X-ray sources.} \\tablenotetext{b}{Net photon counts given by the \\textsc{wavdetect} source detection algorithm.} \\tablenotetext{c}{Statistical significance of detected jet features.} \\end{center} } \\end{table*}     A complication in X-ray studies of 3C~353 is the X-ray emission from its large-scale environment. 3C~353 is situated at the edge of the cluster Zw~1718.1-0108, which is a complex and dynamic system captured at the moment of an on-going merger (Iwasawa et al. 2000). Based on a $40$\\,ks {\\it ASCA} observation, Iwasawa et al. reported that this cluster is characterized by a relatively low X-ray luminosity, $\\sim 5 \\times 10^{43}$\\,erg\\,s$^{-1}$, high temperature, $kT = 3-5$\\,keV, and disturbed morphology, extended over $\\sim 30\\arcmin$. The small field of view of {\\it Chandra}, together with the roll angle of {\\it Chandra} during our observations, prevent us from carrying out detailed studies of the interaction between the radio source and the cluster. Therefore, in this paper we focus solely on the analysis of the non-thermal emission from the nucleus, jets, hotspots and lobes. Analysis of the X-ray emission from the cluster and lobes based on recent {\\it XMM-Newton} observations is presented in Goodger et al. (2008: hereafter G08); they found slightly different temperatures for the southern and northern parts of the cluster, which supported the idea that these two components were originally separate features that are now undergoing a smooth merger, and also showed that the X-ray emission detected from the lobes of 3C~353 was non-thermal in origin, being consistent with the IC/CMB model prediction for the lobes if the magnetic field strength is slightly below the equipartition value, $U_{\\rm B} \\lesssim U_{\\rm e}$.   The present paper is organized as follows. In \\S\\ref{sec:observations}, we describe the {\\it Chandra} observations we performed and the data reduction process, and give a brief overview of the X-ray image we have obtained. A detailed analysis of the jet structures based on the X-ray/radio maps is presented in \\S\\ref{sec:spatial}, and spectral analyses of the nucleus, jets, and lobes are given in \\S\\ref{sec:spectral}. In \\S\\ref{sec:discussion}, we discuss our findings in the context of various jet emission/particle acceleration models. The final conclusions are presented in \\S\\ref{sec:conclusion}. Throughout this paper we adopt a modern cosmology with $\\Omega_{\\rm m} = 0.27$, $\\Omega_{\\rm \\Lambda} = 0.73$ and $H_0 = 71$\\,km\\,s$^{-1}$\\,Mpc$^{-1}$, leading to a luminosity distance of $d_{\\rm L}$ = 131.6\\,Mpc and a conversion scale of $0.60$\\,kpc$/''$ for the 3C~353 redshift, $z = 0.0304$. ",
        "conclusions": "\\label{sec:conclusion}  We have presented a detailed analysis of the data for the powerful FR~II radio galaxy 3C~353 obtained with the {\\it Chandra} X-ray Observatory. In a deep, $90$\\,ks {\\it Chandra} observation, we have detected the X-ray emission from most of the radio structures in this source, including the nucleus, isolated knots in the jet and the counterjet, the terminal jet regions (hotspots), and one radio lobe. Our major findings are as follows: \\renewcommand\\theenumi{(\\roman{enumi})} \\begin{enumerate} \\item Non-thermal X-ray emission associated with a knot in the counterjet of a powerful FR~II source has been detected for the first time. This detection, in agreement with the established large inclination of 3C~353, strongly disagrees with the inverse-Compton model proposed in the literature, and points to a synchrotron origin for the X-ray jet photons. \\item We find that the width of the X-ray knots is narrower ($\\gtrsim 4\\sigma$ effect) than that measured at radio wavelengths, suggesting that the production of the X-ray emission is associated with the central jet spine rather than with the jet boundaries. This conclusion, however, must be confirmed with X-ray data with a much better signal-to-noise ratio. \\item The radio-to-X-ray flux ratio decreases systematically downstream along the outflows, as is often observed in other extragalactic sources. We argue that this is due to the fact that the particle acceleration conditions associated with microphysical plasma parameters changes within the central part of the jet, rather than being due to a gradual decrease of the jet bulk velocity resulting from a smooth jet deceleration. \\item Substantial (kpc-scale) offsets between positions of the X-ray and radio intensity maxima within each knot are found, with the magnitude of the offsets increasing away from the nucleus. We have speculated that these offsets can be explained if radio knots are moving portions of jet material which are produced by the central engine during an epoch of enhanced activity. Depending on the velocity and density ratio between these knots and the outflow produced at times of quiescence, a complex double-shock structure may form as a result of the interaction between the two phases of the ejecta, with the reverse shock (propagating within the faster portion of the jet) being associated with the peak of the X-ray emission. \\end{enumerate} Although we cannot provide definitive solutions or interpretations for each problem, we argue that the 3C~353 data strongly suggest that the synchrotron X-ray emission of extragalactic large-scale jets is not simply shaped by the global hydrodynamical configuration of the outflows, but is also very sensitive to the microphysical parameters of the jet plasma."
    },
    "0806/0806.1056_arXiv.txt": {
        "abstract": "We study the relation between the internal structure of early-type galaxies and their environment using 70 strong gravitational lenses from the Sloan ACS Lens Survey (SLACS). The Sloan Digital Sky Survey (SDSS) database is used to determine two measures of overdensity of galaxies around each lens: the projected number density of galaxies inside the tenth nearest neighbor ($\\Sigma_{10}$) and within a cone of radius one $h^{-1}$ Mpc ($D_1$). Our main results are: 1) The average overdensity is somewhat larger than unity, consistent with lenses preferring overdense environments as expected for massive early-type galaxies (12/70 lenses are in known groups/clusters). 2) The distribution of overdensities is indistinguishable from that of ``twin'' non-lens galaxies selected from SDSS to have the same redshift and stellar velocity dispersion $\\sigma_*$. Thus, within our errors, lens galaxies are an unbiased population, and the SLACS results can be generalized to the overall population of early-type galaxies. 3) Typical contributions from external mass distribution are no more than a few per cent in local mass density, reaching 10-20\\% ($\\sim0.05-0.10$ external convergence) only in the most extreme overdensities. 4) No significant correlation between overdensity and slope of the mass density profile of the lens galaxies is found. 5) Satellite galaxies (those with a more luminous companion) have marginally steeper mass density profiles (as quantified by $f_{\\rm SIE}=\\sigma_*/\\sigma_{\\rm SIE}=1.12\\pm0.05$ vs $1.01\\pm0.01$) and smaller dynamically normalized mass enclosed within the Einstein radius ($\\Delta \\log M_{\\rm Ein}/M_{\\rm dim}$ differs by $-0.09\\pm0.03$ dex) than central galaxies (those without). This result suggests that tidal stripping may affect the mass structure of early-type galaxies down to kpc scales probed by strong lensing, when they fall into larger structures. ",
        "introduction": "The observed properties of galaxies correlate with their environment. For example, the mix of morphological types depend on the local number density of galaxies \\citep{Dre80a,Dre80b,P+G84,Dre++97,Tre++03,Pos++05,Smi++05,Cap++07,Cas++07} and dark matter \\citep{Lan++07}, with elliptical galaxies dominating in high density regions. The star-formation rate and colors of galaxies also scale with their local number density, or distance from the center of clusters \\citep[e.g.,][]{Bal++04,Got++03,Hog++04, Coo++08}.  The physical origin of these environmental trends has been studied for over thirty years. A number of mechanisms have been proposed to shut off or trigger star-formation as well as to modify the mass-dynamical structure of galaxies. They include interactions with other galaxies (mergers or harassment), with the dark-matter potential of clusters and groups (tidal stripping; tidal compression, harassment) and with the intracluster medium when present (starvation or strangulation, ram pressure stripping). A review of the main mechanisms, their effects on galaxies, and the range of environments over which they operate is given by \\cite{Tre++03}.  Although the interaction between galaxies and their environment is not yet fully understood, comprehensive analyses of large imaging and spectroscopic datasets over a range of environments and cosmic times have shown that a variety of mechanisms are at work, starting from the very outskirts of clusters. Physical processes such as starvation and harassment start to be effective at the group stage or when galaxies infall onto clusters as groups, beyond the cluster virial radius. At higher densities other mechanisms such as ram pressure stripping become effective, further modifying the properties of galaxies.  To complement the progress based on traditional luminous tracers of star-formation and morphology, it is clear that one can gain additional insights by following the modifications to the mass-dynamical structure of galaxies induced by the environment. Empirical scaling laws connecting stellar populations with kinematics, such as the fundamental plane and the Tully Fisher relation, have suggested mild trends of the star-formation history of early-type galaxies with environment, and that spiral galaxies are dynamically perturbed as they enter massive clusters \\citep[e.g.][and references therein]{Mor++07c}.  Gravitational lensing provides an additional tool to address the connection between galaxies and their environment. By measuring total mass directly --- rather than through optical tracers --- gravitational lensing gives a direct handle on the transformations of the mass dynamical structure. For example, weak galaxy-galaxy lensing studies in clusters have shown that dark-matter halos of infalling galaxies are tidally truncated \\citep{NKS02,Gav++04,Nat++07}.  In this paper we exploit the large and homogeneous sample of strong gravitational lenses discovered by the Sloan Lenses ACS Survey \\citep[][hereafter collectively SLACS, or SLACS Papers I through VII]{Bol++06, Tre++06, Koo++06, Gav++07, Bol++08a, Gav++08, Bol++08b} to study the connection between environment and mass-dynamical structure at the typical scale of galaxy strong lensing (i.e. $\\sim10$kpc). At these scales the mass distribution is more directly connected with morphology and stellar populations than at the larger scales probed by weak lensing.  Two main effects are expected. On the one hand we expect that tidal truncation by an external potential could steepen the local mass density slope, as suggested for the case of PG1115+080 \\citep{Imp++98,T+K02b}, by numerical simulations \\citep{DKF07}, and by an analysis of the first 15 lenses in the SLACS sample \\citep{Aug08}. If this is the case, then we would expect empirical scaling relations ---- such as the correlation between velocity dispersion of the best fit lens model and stellar velocity dispersion and the mass plane \\citep[MP;][]{Bol++07}--- to depend on the environment. On the other hand, since lensing only measures projected mass, a high density environment with a relatively smooth and shallow embedding dark-matter halo could mimic a shallower local slope and therefore skew measurements in the opposite direction.  One of the goals of this paper is to clarify these issues in order to improve our understanding of the internal structure of lens galaxies, and therefore of early-type galaxies in general if lenses are an unbiased subset.  In addition to providing a diagnostic of the interaction between galaxies and their environment -- and of the internal structure of early-type galaxies -- this study also has repercussions for a number of applications of gravitational lensing.  For example, the degeneracy between local mass density slope and external mass density is the dominant source of error in measuring the Hubble Constant from gravitational time delays \\citep[e.g.,][]{Koc02,Koo++03,Mou++07,Ogu07a,Suy++08}. Discovering trends with the local environment may help reduce systematic errors in these measurements. Similarly, the effects of the local environment need to be taken into account to interpret weak galaxy-galaxy lensing results, as well as to do precision cosmography based on lensing statistics.  In previous Papers (I,II,IV,V) we showed that the SLACS lenses are statistically indistinguishable within the current level of measurement errors from control samples in terms of properties such as size, luminosity, surface brightness, location on the Fundamental Plane, and weak lensing signal, and thus our results could be generalized to the overall population of early-type galaxies. In this paper, we address the question of whether the environment of SLACS lens galaxies differs from that of non-lens galaxies with the same properties, using samples of ``twin galaxies'' (or simply ``twins'') selected to have the same stellar velocity dispersion and redshift.  The paper is organized as follows. In Section~\\ref{sec:data} we briefly review the SLACS sample, describe the selection of a sample of twin galaxies from the Sloan Digital Sky Survey (SDSS), and present our new measurements of local environment.  In Section~\\ref{sec:comparison} we discuss whether the environments of lens galaxies are special compared to those of non-lens galaxies with similar properties. In Section~\\ref{sec:intstuc} we explore the dependence of the internal structure of early-type galaxies on the environment.  In Section~\\ref{sec:disc} we discuss our findings, and in Section~\\ref{sec:sum} we provide a summary.  Throughout this paper magnitudes are given in the AB scale. We assume a concordance cosmology with matter and dark energy density $\\Omega_m=0.3$, $\\Omega_{\\Lambda}=0.7$, and Hubble constant H$_0$=100$h$kms$^{-1}$Mpc$^{-1}$, with $h=0.7$ when necessary. ",
        "conclusions": "\\label{sec:sum}  We have used the SDSS database to measure the environment around seventy strong gravitational lens galaxies selected from the SLACS Survey.  We adopt two standard estimators: the surface density of galaxies within the tenth nearest neighbor ($\\Sigma_{10}$) and the density of galaxies within a cone of radius one $h^{-1}$Mpc ($D_1$). Both are normalized in terms of the corresponding quantities for random fields, and are referred to as ``local'' and ``global'' overdensities. For comparison purposes, we also selected from the SDSS database a sample of 100 ``twins'' for each lens galaxy, i.e.\\ galaxies with virtually the same redshift and velocity dispersion. The new observables are combined with measurements of internal properties of the lens galaxies from SLACS papers to investigate the relationship between early-type galaxy structure and environment. The main results of this study can be summarized as follows:  \\begin{enumerate} \\item SLACS lens galaxies appear to live generally in somewhat overdense environments. Twelve of the seventy lenses ($17\\pm5$\\%) are associated with known clusters/groups at the same redshifts. This is consistent with the notion that lens galaxies are massive and therefore clustered.  \\item The distribution of overdensities for SLACS lens galaxies is consistent within the errors with the corresponding distribution for ``twin'' non-lens galaxies. This is consistent with lens galaxies being an unbiased population of massive early-type galaxies with respect to their environment.  \\item The misalignment of mass and light is found to correlate with the local overdensity of galaxies. The misalignment is negligible for most lens galaxies except for those in the most overdense regions. Randomly oriented external shear of order of 0.05-0.06 is required to reproduce the observed misalignment in the most overdense environments.  \\item The small departures from the average relation between Einstein mass, scaled by dimensional mass, and Einstein radius, scaled by the effective radius, support the previous conclusions. The contribution of the environment to the local potential of the main lens is estimated to be below the current detection threshold for most lenses except for those residing in the densest environments where it can reach at most 10-20\\% in terms of surface mass density at the Einstein radius (0.05-0.1 external convergence).  \\item No significant correlation is found between local and global overdensity and measures of internal structure, such as the slope of the total mass density profile--- quantified in terms of $f_{\\rm SIE}$--- and the difference between the observed Einstein mass and that predicted based on dimensional mass, effective radius, and Einstein radius. Thus -- within the current level of precision -- the internal structure of early-type galaxies does not appear to be biased by projection effects.  \\item The properties of ``satellite'' galaxies (i.e., those with a nearby companion with $i'<i'_{\\rm lens}-1$) are found to be different than those of ``central'' galaxies at the 95-99\\% CL. The average slope of the total mass density profile is $f_{\\rm SIE}=1.12\\pm0.05$ for the satellites and $1.01\\pm0.01$ for the central galaxies. Similarly, the ratio between Einstein Mass and M$_{\\rm dim}$ --- as a function of the Einstein radius in units of the effective radius ratio --- is $0.09\\pm0.03$ dex lower for the satellites with respect to the central galaxies. This suggests that the outer parts of satellite galaxies are perturbed by the environment down to the kpc scales probed by strong lensing, consistent with tidal stripping. \\end{enumerate}"
    },
    "0806/0806.2986.txt": {
        "abstract": "A model of inhomogeneous baryogenesis based on the Affleck and Dine mechanism is described. A simple coupling of the scalar baryon field to the inflaton allows for formation of astronomically significant bubbles with a large baryon (or antibaryon) asymmetry. During the farther evolution these domains form compact stellar-like objects, or lower density clouds, or primordial black holes of different size. According to the scenario, such high baryonic number objects occupy relatively small fraction of space but despite that they may significantly contribute to the cosmological mass density. For some values of parameters the model allows the possibility the whole dark matter in the universe to be baryonic. Furthermore, the model allows the existence of the antibaryonic B-bubbles, i.e. a significant fraction of the mass density in the universe can be in the form of the compact antimatter objects (e.g. anti-stars). ",
        "introduction": "}  Though the universe in our neighbourhood looks as 100\\% baryonic, the observations do not exclude subdominant but astronomically large domains and objects of antimatter. For a recent analysis and literature see e.g. ref.~\\cite{Casadei:2004pi,Bambi:2007cc}. The standard mechanism of baryogenesis~\\cite{Sakharov:1967dj} predicts baryon asymmetry expressed as the ratio of baryon-to-photon number densities: \\be \\beta = n_B/n_\\gamma \\label{beta} \\ee to be constant both in space and, after baryogenesis is over, also in time, however, more complicated models can lead to not so dull pattern. The outcome of the baryogenesis scenario is closely related to the mechanism of CP-violation realized in cosmology. For a review of the latter see ref.~\\cite{Dolgov:2005wf}. If CP is broken spontaneously~\\cite{Lee:1974jb}, the universe would be charge symmetric with equal weight of baryonic and antibaryonic domains so that the average value of $\\beta$ over whole universe would be zero~\\cite{Zeldovich:1974uw,Brown:1979dq}. However, in the standard cosmology the domain size would be very small by astronomical standards and some exponential (inflationary) expansion is necessary to avoid immediate problems~\\cite{Sato:1980yn,Sato:1981ds}. Another cosmological problem of domain walls with huge energy density inside our horizon~\\cite{Zeldovich:1974uw} can be solved either by non-restoration of CP invariance at high temperature~\\cite{Mohapatra:1979vr,Mohapatra:1979qt}, but in this case antimatter domains would not be created, or by a possible mechanism of wall destruction~\\cite{Kuzmin:1981ip,Kuzmin:1981xu}. However, even if these problems are resolved, it can be shown that baryo-symmetric cosmology demands the domain size above (a few) Gpc. Otherwise too intensive gamma ray background would be created by the annihilation on the domain boundaries ~\\cite{Cohen:1997ac}.  If in addition to spontaneous CP-violation there exists an explicit one, the universe would be baryo-asymmetric and, if the magnitude of spontaneous CP breaking happened to be larger than the explicit one, there could exist subdominant antimatter domains ~\\cite{Kuzmin:1981ip,Kuzmin:1981xu,Chechetkin:1982xy}.  Here we explore another mechanism of inhomogeneous baryogenesis suggested in ref.~\\cite{Dolgov:1992pu}. This mechanism is based on the Affleck and Dine~\\cite{Affleck:1984fy} idea that in supersymmetric theories a complex scalar field $\\chi$ with non-zero baryonic number may be displaced far from the origin, $\\chi=0$, along flat directions of the potential $U(\\chi)$. Such flat directions are generic features of supersymmetric models. Inflationary stage is especially favorable for such a travel along the flat directions due to the infrared instability of massless or light, $m^2_\\chi \\ll H^2$, fields in De Sitter space-time. According to the results of ref.~\\cite{Linde:1982uu, Vilenkin:1982wt, Starobinsky:1982ee} (see however, ref.~\\cite{Dolgov:2005se}), vacuum expectation value of massless scalar field rises as: \\be \\langle \\chi^2 \\rangle = H^3 t /4 \\pi^2\\pkt \\label{chi-2-vac} \\ee Inflation not only drives $\\chi$ away from the origin but also allows for formation of the exponentially big domains with a large value of $\\chi$.  After the Hubble parameter, $H$, dropped below $ m_\\chi$, the scalar baryon (AD-field) $\\chi$ would evolve down to $\\chi = 0$ transforming the accumulated baryonic charge into quarks by $B$-conserving decays. Typically such scenario leads to a large baryon asymmetry, $\\beta \\sim 1$ and theoretical efforts should be taken to diminish it down to the observed value $\\beta \\approx 5\\cdot 10^{-10} $.  The modification of this mechanism, according to ref.~\\cite{Dolgov:1992pu}, was an introduction of the quadratic coupling of $\\chi$ to the inflaton field $\\Phi$ in the following renormalizable form, $\\lambda_1 |\\chi|^2 (\\Phi^2 +\\tilde{\\mu} \\Phi )$, where $\\tilde{\\mu}$ is a constant parameter with dimension of mass. This term can be rewritten as: \\be U_{int} (\\chi,\\Phi) = \\lambda_1 |\\chi|^2 \\left( \\Phi - \\Phi_1\\right)^2 \\label{U-int} \\ee by redefinition of the mass term of $\\chi$.  When the inflaton field evolves down to zero, it passes through $\\Phi=\\Phi_1$ and at that moment the window to flat direction in $\\chi$ potential would be open but only for a rather short period of time. Correspondingly, the probability for $\\chi$ to evolve away from the origin would be small and baryogenesis resulting in high $\\beta$ could proceed only in a minor fraction of space. So we should expect that in the bulk of space the usual homogeneous baryogenesis took place, and compact B-bubbles, with very high baryonic number density could be created over this background. Despite their small size, B-bubbles might contribute significant amount into the total cosmological mass density and even make all dark matter of the universe. They would either form compact stellar-like objects, or high density clouds, or primordial black holes (BH). In the simplest version of the model there should be an equal number of baryonic and antibaryonic B-bubbles. So cosmology would be almost baryo-symmetric but the bounds derived in ref.~\\cite{Cohen:1997ac} are not applicable here because of compactness of the bubbles. An interesting feature of this scenario is that dark matter (DM) can be made from the usual particles, baryons and antibaryons in the form of evolved stellar type objects or primordial black holes.  The paper is organized as follows. In the Sec. \\ref{sec-model} we give a detailed description of the model and it's parameter space. The B-bubble evolution and the generation of the baryon asymmetry are discussed in Sec. \\ref{s-b-evol}. The next section, Sec. \\ref{s-size}, gives a detailed description of the development of quantum fluctuations of the Affleck and Dine field and the distribution of B-bubbles over the length and mass. In Sec. \\ref{s-perturb} the fate of B-bubbles during different cosmological epochs is considered. In Sec.\\ref{s-cosm-astro} the observational bounds on the dark matter in the form of primordial black holes (PBH) and the amount af antimatter in the compact stellar objects are presented. Our conclusions are hosted in the Sec. \\ref{s-conclusions}. ",
        "conclusions": "We reconsidered and further developed the scenario of ref.~\\cite{Dolgov:1992pu} and argued that a simple coupling of the Affleck-Dine scalar baryon to the inflaton may lead to an efficient creation of PBHs either with relatively small masses, $M_{BH}=10^{14}-10^{20}g$, or much larger, $M_{BH}\\geq M_\\odot$. The predicted log-normal mass spectrum of PBH (\\ref{dn-dM}) very much differ from the power law one which could be created by other mechanisms considered in the existing literature. This difference is crucial for relaxing the observational bounds on PBH abundance. The suggested mechanism may explain formation of all observed heavy black holes from millions to thousands or tens solar masses. There is no unique and compelling explanation in the existing literature. There is an interesting possibility that all cosmological DM is made of PBH or compact stellar mass objects (early formed and now dead stars) with log-normal mass spectrum. In particular, the observed MACHOS may be solar mass PBHs, or dead (anti)stars.  In the simplest version the mechanism considered here leads to significant creation of astronomical objects made of antimatter. Such objects should have anomalay abundances of (primordially formed) light elements. The search for cosmic antimatter \\cite{Picozza}, if successed would be crucial for the conformation of the scenarion."
    },
    "0806/0806.2675_arXiv.txt": {
        "abstract": "{The possibility to constrain fractal space dimensionality from Astrophysics and other areas is briefly reviewed. Assuming such dimensionality to be $3 + \\epsilon$, a limit to $\\epsilon$ can be inferred from COBE satellite data. The available data for the cosmic microwave background radiation spectrum are fitted by a Planck's radiation distribution generalized to non integer space dimensionality. Our analysis shows that the shape of the CMBR spectrum, which does not depend on the absolute normalization, is correctly described from this distribution provided the absolute temperature is equal to 2.726 $\\pm$ $0.003\\times 10^{-2}$~K and $\\epsilon = - (0.957 \\pm 0.006) \\times 10^{-5}$. This value for $\\epsilon$ is shown to be consistent with what is found on a very different spatial scale based on a quantum field phenomenon. The  $|\\epsilon|$ is interpreted as an upper limit for how much space dimensionality could have deviated from three. In other words, this is the maximum fluctuation space dimensionality should have experienced in a spatial and temporal scale compared to that of the decoupling era.}  \\noindent {\\it Subject headings:} cosmic microwave background  \\vspace{1.5cm} ",
        "introduction": "The general idea that Astrophysics has something to say about space dimensionality can be traced back to Kant's conjecture (Handyside 1929) that three-dimensionality of space can be related to the structure of Newtonian gravitational law. As a second step, one can quote the work of William Paley (1802), which can be considered the first attempt to shed light on the space dimensionality problem clearly from Anthropic arguments. In his work, Paley analyzes the consequences of changes in the form of Newton's gravitational law and of the stability of the solar system on human existence. His speculations take into account a number of mathematical arguments for an anthropocentric design of the World, which rest all upon the stability of the planetary orbits in our solar system and on a Newtonian mechanical {\\it Weltanschauung}, as should be expected at that time.  A scientific approach to this issue was introduced by Ehrenfest (1917, 1920) who tried to understand the threefold nature of space by formulating the question: ``{\\it How does it become manifest in the fundamental laws of Physics that space has 3 dimensions?}\". Following this approach, the first investigations concerning General Relativity and a heuristic model of a pulsating Universe were made by Tangherlini (1963, 1986), who tried to impose constraints to integer space dimensionality by searching for bound stable states of the Universe.  The above question can obviously be reversed and we can try to answer the following: ``{\\it How do the fundamental laws of Physics entail space dimensionality?}\" (Caruso \\& Moreira, 1997). In this alternative approach, space dimensionality may be taken as an unknown quantity or can be admitted to have a non integer value $d = 3 + \\epsilon$, $\\epsilon$ being a parameter to be experimentally determined. Since the introduction of the concept of fractal dimension by Mandelbrot (1977) this became an interesting possibility to be explored. ",
        "conclusions": "\\label{disc}  Strictly speaking, since absolute measurements for the fluxes are not yet available, the physical meaning of the parameter $\\epsilon$ should be reinterpreted. Indeed, since the experimental device compares the background radiation to that of a local black-body with controlled temperature, at least in principle, one can figure out that the value of $\\epsilon$ could be locally (in this case, where the reference black-body is placed) different from the fractal dimensions of space on the horizon scales. In this case, $\\epsilon$ should be interpreted as a difference between these two fractal dimensions at two far away spatial scales, as did Grassi, Sironi \\& Strini (1986). Alternatively, one can suppose that locally space is just three-dimensional and, in this case, we are putting a limit on how much space dimensionality could be different from three in the farthest corner of the universe one can look into.  In this paper, we established a bound for $\\epsilon$ by fitting data on CMBR taking the absolute temperature $T$, $\\epsilon$, and the normalization as free parameters in eq.~(\\ref{fit_function}). The overestimated value we get for $\\epsilon$, equation~(\\ref{epsilon}), is two order of magnitude less than the more accurate upper limit value estimated by Torres \\& Herrej\\'on (1989) with small errors. Another important difference found here, in respect to this last paper and to that of Grassi, Sironi \\& Strini (1986), is that we are able to fix also the {\\it sign} of the $\\epsilon$ parameter, showing that it is {\\it negative}. To the best of our knowledge, this is the first time that the sign of $\\epsilon$ can be determined by analyzing experimental results at astrophysical or cosmological scale. Up to now, the strongest constraint comes from microphysics. Indeed, Zeilinger \\& Svozil (1985) were able to determine not only the value of $|\\epsilon| = |d-3|$ but the sign of this parameter too, from quantum field theory, getting $ \\epsilon = - (5.3 \\pm 2.5) \\times 10^{-7}$, which is approximately one order of magnitude less than our upper limit for $|\\epsilon|$. It is important to remember that they pointed out that a value of space-time dimension less than 4 (assuming just one time dimension) would render all the logarithmic divergences in relativistic quantum field theory finite. Their result seems to resolve the discrepancies between the theoretical and experimental values of the anomalous magnetic moment of the electron ({\\it cf.} also Svozil \\& Zeilinger, 1986). Another evidence in favor of arbitrarily small non-zero $|\\epsilon|$ comes from the study of Ising gauge theories in non-integer dimensions (Bhanot \\& Salvador, 1986). This was a first evidence, based on a numerical analysis, that for abelian gauge theories, the space-time dimension $d+1 =4$ may, in some sence, be identified as an upper critical dimension.    In the concluding remarks of their 1985 paper, Zeilinger \\& Svozil wrote: ``{\\it It is certainly a challenge for future research to investigate whether or not the deviation of the dimension of space-time from four can be made more statistically significant than the present work suggests. Furthermore, the question of possible evidence for such a small deviation in other areas of Physics deserves attention}\". If we assume time to be one-dimensional, as usually done, the main result of the present paper can be seen as an answer to both challenges.  Finally, since the small deviation from three-dimensionality of space obtained here is extract directly from the CMBR data, it actually suggests that space dimensionality did not vary significantly in a huge temporal scale, once this background radiation is expected to be related to the Big Bang. This time scale can be safely put on the later epoch where the universe was about $3\\times 10^{5}$ yr old ($z \\simeq 10^{3}$). In addition, due to its isotropy, this is the {\\it only} experimental situation that, being confronted to {\\it any} other local experiment on Earth, aimed to measure the number of space dimensions, drive to the conclusion that dimensionality did not vary also in a very large spatial scale (a cosmological scale)."
    },
    "0806/0806.2505_arXiv.txt": {
        "abstract": "We develop a technique to study relativistic perturbations in the generalised brane cosmological scenario, which is a generalisation of the multi-fluid cosmological perturbations to brane cosmology. The novelty of the technique lies in the inclusion of a radiative bulk which is responsible for bulk-brane energy exchange, and in turn, modifies the standard perturbative analysis to a great extent. The analysis involves a geometric fluid -- called the Weyl fluid -- whose nature and role have been studied extensively both for the empty bulk and the radiative bulk scenario. Subsequently, we find that this Weyl fluid can be a possible geometric candidate for dark matter in this generalised brane cosmological framework. ",
        "introduction": "As mentioned, we shall concentrate on the most general bulk scenario, for which the bulk geometry is given by a Vaidya-anti de Sitter metric \\begin{equation} d S_5^2 = - f(r, ~v) ~dv^2 + 2 dr ~dv + r^2 d \\Sigma_3^2 \\label{eqd1} \\end{equation} where $\\Sigma_3$ is the 3-space having flat, spherical or hyperboloidal symmetry, $f(r, ~v) =  k - \\frac{\\Lambda_5}{6}r^2 - \\frac{m(v)}{r^2}$, and $m(v)$ is the resultant of the variable mass of the Vaidya black hole and radiation field. This type of bulk  can exchange energy with the brane as a null flow along the radial direction \\cite{maartbulk, langrev, lang1, chamb, lang2, vads5gen, asymgen, nonradial, radbh5, sustr, sucol}. Consequently, the brane matter conservation equation is modified to \\be \\dot\\rho + 3 \\frac{\\dot a}{a} (\\rho + p) = -2 \\psi \\label{rho} \\label{vaidden} \\ee where $\\psi$ is the null flow characterising the VAdS$_{5}$ bulk by the  radiation field of a null dust $T_{AB}^{\\rm bulk} = \\psi q_{A} q_{B}$, which leads to the above equation by using \\be \\nabla^\\mu T_{\\mu \\nu} = -2 ~T_{AB}^{\\rm{bulk}} ~n^A g_\\nu^B \\ee (where $n^A$ are the normals to the surface), which gives \\be \\nabla^\\mu T_{\\mu \\nu} = -2 \\psi u_{\\mu} \\ee ($u_{\\mu}$ are the unit velocity vectors), and readily leads to Eq (\\ref{rho}). This modified conservation equation, with the help of  the Bianchi identity on the brane $\\nabla^\\mu G_{\\mu \\nu} =0$, leads to another constraint equation \\be \\nabla^\\mu {\\cal E}_{\\mu \\nu} = \\frac{6 \\kappa^2}{\\lambda} \\nabla^\\mu {\\cal S}_{\\mu \\nu} + \\frac{2}{3} \\left[ \\kappa_5^2 \\left( \\dot \\psi + 3 \\frac{\\dot a}{a} \\psi \\right) - 3 \\kappa^2 \\psi \\right] u_{\\mu} +\\frac{2}{3} \\kappa_5^2 \\overrightarrow \\nabla_\\mu \\psi \\ee where $\\lambda$ is the brane-tension and ${\\cal E}_{\\mu \\nu}$ and ${\\cal S}_{\\mu \\nu}$ are, respectively, the projected bulk Weyl tensor and the quadratic contribution from the brane energy-momentum tensor to the Einstein equation on the brane \\cite{sustr}. The above equation governs the evolution of the Weyl fluid $\\rho^*$ (so named since it is a fluid-like contribution from the bulk Weyl tensor to the brane). For FRW geometry on the brane, this is given by \\be \\dot\\rho^* + 4 \\frac{\\dot a}{a} \\rho^* = 2 \\psi - \\frac{2}{3} \\left ( \\frac{\\kappa_{5}}{\\kappa} \\right )^{2} \\left( \\dot \\psi + 3 \\frac{\\dot a}{a} \\psi \\right) \\label{vaidweyl1} \\ee so that this quantity evolves as  \\cite{lang2, sustr} \\be \\rho^* = \\frac{C(\\tau)}{a^4} \\propto \\frac{1}{a^{(4 - \\alpha)}} \\label{vaidweyl2} \\ee which gives a general, physically relevant behaviour for the Weyl fluid. Here, $\\tau$ is the proper time on the brane. Obviously, contrary to the Sch-AdS$_5$ bulk, here $C(\\tau)$ is evolving, and consequently, the Weyl fluid no longer behaves like radiation. To a brane-based observer, the cosmological dynamics is now governed by an effective perfect fluid, the  components of which  are given by \\cite{chamb, sustr} \\bea \\rho^{\\text{eff}} &=& \\rho +\\frac{\\rho^2}{2\\lambda} + \\frac{C(\\tau)}{a^4} \\label{eqb15} \\\\ p^{\\text{eff}} &=& p  + \\frac{\\rho}{2\\lambda} (\\rho +2p)+\\frac{C(\\tau)}{3 a^4} \\label{eqb16} \\eea The anisotropic components of the Weyl fluid, {\\em viz.}, $q^*_\\mu$ and $\\pi^*_{\\mu\\nu}$ vanish, in order that the VAd$S_5$ bulk be compatible to FRW geometry on the brane. The Friedmann equation and the covariant Raychaudhuri equation, expressed in terms of these effective quantities, are  respectively \\cite{sucol} \\bea H^2 &=& \\frac{\\kappa_4^2}{3} \\rho^{\\rm eff} + \\frac{\\Lambda}{3} - \\frac{k}{a^2} \\label{fr} \\\\ \\dot H &=& -\\frac{\\kappa_4^2}{2} \\left(\\rho^{\\rm eff} + p^{\\rm eff} \\right) + \\frac{k}{a^2} - \\frac{\\kappa_5^2}{3} \\psi \\label{rc} \\eea    In the brane-based Newtonian analysis of perturbations by considering small fluctuations of the effective density $\\rho^{\\text{eff}}(\\overrightarrow x, ~\\tau) = \\bar\\rho^{\\text{eff}}(\\tau) (1 + \\delta^{\\text{eff}}(\\overrightarrow x, ~\\tau)) $ and the effective gravitational potential $\\Phi^{\\text{eff}}(\\overrightarrow x, ~\\tau) = \\Phi_0^{\\text{eff}} + \\phi^{\\text{eff}}$ on the hydrodynamic equations for  this effective perfect fluid, one obtains for a barotropic fluid a single second order equation in terms of Fourier mode \\be \\frac{d^2 \\delta_k^{\\text{eff}}}{d \\tau^2} + 2\\frac{\\dot a}{a} \\frac{d \\delta_k^{\\text{eff}}}{d \\tau} - \\left[4 \\pi G \\bar\\rho^{\\text{eff}} - \\left(\\frac{c_s^{2{\\text{eff}}} k}{a} \\right)^2\\right]\\delta_k^{\\text{eff}} = 0 \\label{eqc11} \\ee where $c_s^{2{\\text{eff}}}$ is  the square of the effective sound speed \\cite{maartrev, sustr}. The above perturbation equation of the effective fluid can account for the required amount of gravitational instability  if the Weyl density redshifts more slowly than baryonic matter density, so that  it can dominate over  matter at late times, which is realised when $1 < \\alpha < 4$ in Eq (\\ref{vaidweyl2}). Now, for late time behaviour, we can drop the quadratic terms in equations (\\ref{eqb15}) and (\\ref{eqb16}) so that the effective density is given by $\\rho^{\\text{eff}} = \\rho^{(b)} + \\rho^*$ which is now constituted of the usual matter (baryonic) density $\\rho^{(b)}$ and an additional density contribution from the Weyl fluid. This Weyl density, being geometric, is essentially non-baryonic. Consequently, we can decompose Eq (\\ref{eqc11}) to get the individual evolution equations for the perturbation for each of the fluids \\begin{eqnarray} \\frac{d^2 \\delta^{(b)}}{d \\tau^2} + 2\\frac{\\dot a}{a} \\frac{d \\delta^{(b)}}{d \\tau} = 4 \\pi G \\bar\\rho^{(b)} \\delta^{(b)} + 4 \\pi G \\bar\\rho^* \\delta^* \\label{eqc18} \\\\ \\frac{d^2 \\delta^*}{d \\tau^2} + 2\\frac{\\dot a}{a} \\frac{d \\delta^*}{d \\tau} = 4 \\pi G \\bar\\rho^* \\delta^* + 4 \\pi G \\bar\\rho^{(b)}  \\delta^{(b)} \\label{eqc19} \\end{eqnarray} where $\\delta^{(b)}$ and $\\delta^*$ are the fluctuations of baryonic matter and Weyl fluid respectively. With $\\Omega^{(b)}  \\ll \\Omega^*$, the relevant growing mode solutions are given by \\cite{sustr} \\bea \\delta^*(z) &=& \\delta^*(0) (1 + z)^{-1} \\label{eqc20} \\\\ \\delta^{(b)}(z) &=& \\delta^*(z) \\left(1 - \\frac{1+z}{1+z_N}\\right) \\label{eqc23} \\eea with the input that the late time behaviour of the expansion of the universe in RS II is the same as the standard cosmological solution for the scale factor \\cite{maartbulk, maart6} where  the scale factor is related to the redshift function by $a \\propto (1+z)^{-1}$.   The solutions reveal that at a redshift close to $ z_N$, the baryonic fluctuation $\\delta^{(b)}$ almost vanishes but the Weyl fluctuation $\\delta^*$ still remains finite. So, even if the baryonic fluctuation is very small at a redshift of $z_N \\approx 1000$, as confirmed by CMB data \\cite{cmb}, the fluctuations of the Weyl fluid still had a finite amplitude during that time, whereas at a redshift much less than  $z_N$ the baryonic matter fluctuations are of equal amplitude as the Weyl fluid fluctuations. This is precisely what is required to explain the formation of structures we see today. Thus, the Newtonian analysis of perturbations on the brane is capable of explaining structure formation (within its limit) by Weyl fluid, devoid of any material existence of dark matter. Hence, the Weyl fluid acts as a possible geometric candidate for dark matter. ",
        "conclusions": "In this article, we have developed a technique for relativistic  perturbations valid for a general brane cosmological scenario. The essential distinction of our analysis from the studies on brane cosmological perturbations available in the literature is that, here the geometrical effect of the bulk on the brane -- the so-called Weyl fluid  -- plays a very crucial role in determining the nature of the evolution of density perturbations. This is materialised from the realisation that in the general brane cosmological scenario obtained from Vaidya-anti de Sitter bulk, the Weyl fluid plays a significant role in controlling the dynamics on the brane, contrary to the earlier results based on dark radiation. Our results are, in a sense, a generalisation of the multi-fluid covariant perturbation formalism in brane cosmological framework. Further, we have solved the perturbation equations and found that the perturbation of the Weyl fluid grows at late time, and thus, this component of braneworld gravity plays a significant role in late time cosmology to act as a possible geometric candidate for dark matter. We have discussed some of the implications of fluctuations involving it and have   mentioned some observable sides of this model as well.   An important issue is to fit this theoretical model with current observational data. Recently there has been some progress in this direction \\cite{obs}. An extensive study on confronting this braneworld model with observations in a more rigorous method can provide us with necessary information on the merits and demerits of the formalism. To this end, a thorough study of different parameters related to cosmological perturbation is to be performed. As mentioned in the previous section, the different cosmological parameters need to be reformulated in this framework. The next step is to estimate them and confront them with observations. For example, it is to be seen if the power spectrum, redefined in this paradigm with the Weyl fluid acting as a dark matter candidate, fits with the highly accurate observational data.   Further, analysis of different types of metric-based perturbations, namely, scalar, vector and tensor as well as related issues like CMB anisotropy, Sachs-Wolfe effect etc has to be be studied in details in this brane cosmological framework with a significant  Weyl fluid. An extensive study in this direction is essential, which we hope to address in near future. Also, to apply this formalism in the framework of branworld models of dark energy  \\cite{de} remains as  another interesting issue."
    },
    "0806/0806.1730_arXiv.txt": {
        "abstract": "We characterize the problem of artificial polarization for the Submillimeter High Angular Resolution Polarimeter (SHARP) through the use of simulated data and observations made at the Caltech Submillimeter Observatory (CSO). These erroneous, artificial polarization signals are introduced into the data through misalignments in the bolometer sub-arrays plus pointing drifts present during the data-taking procedure. An algorithm is outlined here to address this problem and correct for it, provided that one can measure the degree of the sub-array misalignments and telescope pointing drifts. Tests involving simulated sources of Gaussian intensity profile indicate that the level of introduced artificial polarization is highly dependent upon the angular size of the source. Despite this, the correction algorithm is effective at removing up to 60\\% of the artificial polarization during these tests. The analysis of Jupiter data taken in January 2006 and February 2007 indicates a mean polarization of $1.44\\%\\pm0.04\\%$ and $0.95\\%\\pm0.09\\%$, respectively. The application of the correction algorithm yields mean reductions in the polarization of approximately 0.15\\% and 0.03\\% for the 2006 and 2007 data sets, respectively. ",
        "introduction": "Submillimeter polarimetry provides a means to investigate the morphology of interstellar magnetic fields that are highly embedded in dusty clouds. Such an investigative tool is extremely useful for the study of astrophysical phenomena in which magnetic fields are suspected to play a significant role. Such areas of interest include: star formation \\citep{Shu 1987,Hildebrand 1984}, circumstellar disks and jets \\citep{Davis 2000}, filamentary structure in molecular clouds \\citep{Fiege and Pudritz 2000}, and galactic-scale field morphology \\citep{Greaves and Holland 2002}. In the particular case of low mass star formation, the current leading model places great emphasis upon the presence of embedded magnetic fields to regulate the entire process \\citep{Mouschovias 2001}. Hence any further understanding of these magnetic fields may yield a clearer understanding of the origins of \\textquotedblleft{}solar-like\\textquotedblright{} stellar/planetary systems.  Current work in this field is being carried out at the Caltech Submillimeter Observatory (CSO) using the Submillimeter High Angular Resolution Polarimeter (SHARP). SHARP is a fore-optics module designed to be used in conjunction with the SHARC-II camera to form a highly sensitive, dual-wavelength (350-$\\mu\\mathrm{m}$ and 450-$\\mu\\mathrm{m}$), polarimeter \\citep{Novak 2004,Li 2006}. SHARC-II employs a 12 $\\times$ 32 pixel bolometer array that is optically \\textquotedblleft{}split\\textquotedblright{} by SHARP into three zones: two 12 $\\times$ 12 pixel regions that record orthogonal states of linear polarization (which are labeled \\textquotedblleft{}H\\textquotedblright{} and \\textquotedblleft{}V\\textquotedblright{} for horizontal and vertical, respectively), and a 12 $\\times$ 8 pixel central zone that is not used with SHARP. The horizontal and vertical components are combined during data reduction to yield the $I$, $Q$, and $U$ Stokes parameters.  The simultaneous measurement of the H and V polarization components allows for the effective removal of the sky background signal \\citep{Hildebrand 2000}. However, it does not negate the possibility of erroneous polarization signal generation. The combination of misalignments between the two sub-arrays (i.e., H and V) and pointing drifts during the observation cycle, can result in the generation of artificial polarization. The generation of these erroneous signals may place limitations on the sensitivity of SHARP and thus could reduce data-gathering efficiency. This would hurt efforts to rapidly survey large extended objects, such as giant molecular clouds (GMC\\textquoteright{}s), where many observations would be required to properly survey the source and thus a high data-taking efficiency is required.  A correction algorithm has been designed in an attempt to model and correct for this problem in the SHARP data reduction pipeline. This paper will go over in detail the problem of artificial polarization in dual-array polarimeters and the algorithm by which a correction is attempted, with simulated and planetary data being used to test the proposed method. Section \\ref{sec:2} will describe the means by which artificial polarization is generated in a dual-array polarimeter. Section \\ref{sec:3} will describe the algorithm employed to treat this problem. Section \\ref{sec:4} of this paper will discuss the magnitude of the problem and cover the results obtained thus far from the testing of simulated and planetary data. Section \\ref{sec:5} covers the concluding remarks. ",
        "conclusions": ""
    },
    "0806/0806.1506_arXiv.txt": {
        "abstract": "We present source code for the computer algebra system Mathematica that analyzes the distribution of nearby Galaxies using SDSS data. Download instructions are given, thus within 10 minutes, the reader can verify that galaxies are distributed in an essentially non-homogeneous manner and  cluster on 2-dimensional structures. The short code uses a simple method inspired by Minkowski functionals: the distances to the next neighbors are calculated and compared to random distributions in three and two dimensions. The observed distance distribution corresponds clearly to the latter case. The paper may also be helpful for nonexpert scientists to get started with SDSS data analysis. ",
        "introduction": "Still 50 years after Hubble's discovery of the expanding universe the distribution of galaxies was assumed to homogeneous - a seemingly obvious consequence of the cosmological principle. Pioneering investigations \\cite{Gel:86, Gel:89} however showed that distribution of galaxies is all but homogeneous. Rather there seems to be a hierarchy of galaxy groups, clusters and superclusters that concentrate on twodimensional structures, while there are large voids in between (`sponge structure', \\cite{Got:86}). There is an ongoing discussion whether the universe becomes homogeneous for scales larger than 100 $Mpc$, % or if it has the properties of a fractal with $D=2$ even on larger scales \\cite{Syl:08}.  Here we do not present any new results that help to decide that question and our approach cannot compete with the detailedness of the expert's analysis (\\cite{Pee:01,Syl:08, Got:08} and references herein). From a point of view of general scientific methodology, we find it however desirable that important results of fundamental physics that require extensive numerical treatment can be repeated by a broad public of non-expert scientists\\fo{See, e.g. \\cite{Unz:07a} for a similar approach.}. In particular, the unique quality of the free accessible SDSS data supports such an approach we would like to ease further. Two-point statistics are frequently used to extract information on the dimensionality. Here we use just next neighbor statistics. Imagine spheres with growing radius $r$ around each galaxy. When $r$ reaches a critical radius $r_c$, the spheres will overlap to a connected manifold of the size of the whole sample (see fig.~\\ref{minkowski}).  \\begin{figure}[h] \\includegraphics[width=8cm]{minkowski.jpg} \\caption{Example of Minkowski functionals: at a critical radius, the manifold becomes connected. Picture taken from \\cite{Ost:07}.} \\label{minkowski} \\end{figure} Obviously, this transition must occur much earlier (at a small $r_c$) when the distribution is not homogeneous in three dimensions. We do not fully implement this method of Minkowski functionals \\cite{Mec:94}, but next neighbor distances obviously do yield significant information.  The code of about 100 lines given below reproduces the results given in section~\\ref{results}. It can easily be run with different data sets and future data releases of SDSS. We plan to add some refinements for a second version, but also the reader should be able to do slight modifications or extensions of the code. A quick description for getting started is found in section~\\ref{sbs}. Though we cannot give a detailed description of the program, some clarifying comments are included in the quite self-explaining code (see ~\\ref{sbs2}). ",
        "conclusions": "The distribution of galaxies in the universe is still a riddle and theoretically not fully understood \\cite{Pee:01}. The discovery of the two-dimensional structure of the galaxy distribution reminds us from something deep and mysterious, such as from Dirac's observation of a twodimensional density in the universe.\\footnote{From Dirac's large number hypothesis follows that the surface of all protons is of the same order as the surface of the horizon. Today this is usually considered as pure coincidence.} Particulary $\\Lambda$CDM simulations have problems to account for the observed structure. Observational progress in this field is therefore very much triggered by new high-quality data like SDSS. The ongoing discussion can thus benefit from a broad accessibility of those data and a transparent processing which is not limited to a few groups. If one day the data allow a definite answer to decide whether the large scale structure is homogeneous of even fractal, this must become evident also for the non-expert scientist. We hope this is a little step towards repeatability and transparency for the efforts to answer that important question.  \\paragraph{Acknowledgement.} Though we are grateful for any comments, please understand that we cannot guarantee functionality or give further support for getting this program to run on your computer."
    },
    "0806/0806.3359_arXiv.txt": {
        "abstract": "R-mode instabilities lead to specific signatures in the evolution of rotating pulsars, and may provide a unique signature which can identify the phase of matter that exists in the interiors of such objects.  The contributions of quarks' bulk and shear viscosities to the dissipation of r-mode instabilities are studied. It is shown that (contrary to earlier works) the quark contribution is \\emph{not} fully suppressed in the presence of fully gapped color-superconducting phase due to the fact that the bulk viscosity is a resonance effect between the density oscillation on one side, and the interaction rates that try to restore $\\beta$-equilibrium during these oscillations on the other. This will have an effect on the structure of the r-mode instability window and will affect the angular momentum evolution of pulsars. ",
        "introduction": "Compact (neutron) stars provide a natural laboratory of matter under extreme conditions. In the central regions of such stars the baryon density of matter could reach values up to 10 times the nuclear saturation density (i.e., $10 \\rho_0$ where $\\rho_0\\simeq 0.15~\\mbox{fm}^{-3}$). At such high densities matter is likely to be in a deconfined state in which quarks, rather than hadrons, are the natural dynamical degrees of freedom \\cite{Ivanenko1965,Ivanenko1969, Itoh1970,Iachello1974,Collins1975}. At sufficiently low temperatures, the ground state of deconfined quark matter is a color superconductor. (For reviews on color superconductivity, see Refs.~\\cite{RajWil, Alfordreview,Reddy2002,Rischke2003,Buballa2003,Huangreview,Shovkovy2004, Alford:2006fw}). Many phases of color superconductivity are known that could possibly be realized in dense matter. It remains unclear, however, which of these describe the ground state of matter under the specific conditions in stars. This is because of theoretical uncertainties in treating the strongly coupled, non-perturbative dynamics in QCD at the baryon densities of relevance.  There are many instabilities associated with a star's rotation that can be operative in a neutron star. The most familiar of those is that if the star rotates with frequencies above a critical value, known as the Kepler frequency, mass shedding sets in from the star's equator. This frequency is $\\nu_K \\sim 1200 \\mbox{Hz} $ in a typical neutron star.  Of concern to us are the r-mode instabilities. (For reviews on the r-mode instabilities see REfs.~\\cite{Anderssonreview, Lind-lect}). R-modes are fluid waves in rotating stars with the Coriolis force acting as the restoring force. These waves couple to gravitational radiation emission causing the star to lose angular momentum. Only when dissipative phenomena dampen these r-modes can the star rotate without losing any angular momentum. This sets a limit on the maximum rotation frequency of the star at any given temperature. This also affects the rotation frequency evolution of the star, see, for instance Ref.~\\cite{Zheng:2006gm}.  There can be many dissipative phenomena that will damp the r-modes, these phenomena depend strongly on the phase of matter expected and on the structure of the star. In this paper we study the quark contribution to the r-mode instability window of quark stars for different phases of deconfined quark matter.  At first, we consider the r-mode instability window for unpaired quark matter in comparison with the r-mode instability window of proton-neutron-electron matter (\\emph{npe} matter), the results are very similar to the earlier work in Ref.~\\cite{Madsenprl2}.  It was suggested \\cite{Madsenprl2} that the gapped quark contribution in the color-flavor locked (CFL) phase is strongly suppressed even if the energy gap is as small as $\\Delta_{CFL} \\sim 1 \\mbox{MeV}$. This will be reconsidered more carefully in this article. Moreover, the r-mode instabilities for the two-flavor color superconductor (2SC) phase will be studied, the results are different from those of Ref.~\\cite{Madsenprl2} due to the fact that the bulk viscosity may increase with the reduction of weak interaction rates as was seen in Ref.~\\cite{SSR1}. The quark contribution to the r-mode instability window in the color-spin-locked (CSL) phase will also be studied.  This paper is organized as follows: A quick review of the r-mode instabilities is presented in Sec. \\ref{rm}. Quarks contribute to the r-mode instability window mainly through two dissipative phenomena, the bulk viscosity and the shear viscosity, these are reviewed in Sec. \\ref{dp}. In order to relate the dissipative factors ``the bulk and shear viscosities\" to the r-mode instabilities, one has to calculate the dissipation timescales related to these factors, these are discussed in Sec. \\ref{dts}. The results for the r-mode instability window are shown in Sec.~\\ref{results}. The discussion of the results is given in Sec.~\\ref{d}. ",
        "conclusions": "\\label{d} In this article, the contribution of quarks to the r-mode instabilities was studied. The main conclusion is that one \\emph{cannot} neglect the quark contribution to r-mode instabilities even if all quarks were gapped.  It is shown in Fig.~(\\ref{nu_c_normal}) that the r-mode instability window is very different for quark matter than for neutron star matter. This might provide us with an opportunity to distinguish between quark stars and neutron stars. Further investigation is clearly needed.  Each different color-superconducting phase has a different structure for the r-mode instability window. Quarks in the CFL phase generate a narrow window of stability at temperatures which are, incidentally, the temperatures at which a young pulsar is formed, see Fig.~(\\ref{fig:BV_nu_c_cfl}). This implies that a newly born quark star in the CFL phase is generated in a stable region, and has to cool down first to reach the instability region which starts at $T \\simeq 10 \\mbox{MeV}$. One cannot draw more conclusions because the CFL phase has several other contributions to be considered.  The r-mode instability window for the CSL phase provides a wider stability region which ends at a temperature $T\\sim 0.01 \\mbox{MeV}$, see Fig~(\\ref{fig:BV_nu_c_csl}). The situation is not fully resolved in the case of the CSL phase due to the uncertainties in the values of the energy gap and the critical temperature, and the fact that there are other dissipative phenomena to be considered.  The 2SC phase has a structure very similar to unpaired quark matter, see Fig.~(\\ref{fig:BV_nu_c_2sc}), and the conclusions for the unpaired phase can be similarly drawn for the 2SC phase. This is because in the 2SC phase there is always an ungapped quark color. This means that $\\frac 1 3$ of the quarks behave exactly as the unpaired quark matter phase. Moreover, the strange quark always remains gapless. This leads to a suppression for the non-leptonic weak interaction close to $\\frac 2 9$, for the strange quark \\emph{Urca} processes it is $\\frac 2 3$, and for the down quark \\emph{Urca} processes it is $\\frac 1 3$. Ref.~\\cite{Madsenprl2} suggests that the 2SC phase is marginally inconsistent with LMXB's data, we have shown that the 2SC is completely consistent with the  LMXB's observations, see Fig.~(\\ref{fig:BV_nu_c_2sc}).  To better understand the reasone for the enhancement of the bulk viscosity (and of the r-mode stability window) due to color superconductivity, we need to look deeper into the behavior of the bulk viscosity with respect to the oscillation frequency $\\omega$ and the difference in interaction rates $\\lambda_i$. It is hard to see that from Eq.~(\\ref{zeta1-3}). Let us consider the case of a single-interaction. If only ruled by the non-leptonic weak interaction, the bulk viscosity takes the form \\cite{SSR2}  \\begin{equation} \\zeta_{\\rm non} \\simeq \\frac{\\lambda_{1}C_{1}^2}{\\omega^2+\\left(\\lambda_{1}A_{1}/n\\right)^2}. \\end{equation}  As can be seen from this formula, there are two regions with two different dependences on the interaction rate: 1) The region where $\\omega^2  \\gg \\left(\\lambda_{1}A_{1}/n\\right)^2$, in this region the bulk viscosity can be approximated by $ \\zeta_{1}\\simeq \\frac{\\lambda_{1}}{\\omega^2}C_{1}^2$ and any reduction to $\\lambda_1$ leads to a similar reduction of the bulk viscosity ($e^{-\\frac{2 \\Delta_{CFL}}{T}}$ in the CFL phase). 2) The region where $\\omega^2  \\ll \\left(\\lambda_{1}A_{1}/n\\right)^2$, in this region the bulk viscosity is approximated by $\\zeta_{1} \\simeq \\frac{n C_1^2}{A_1^2 \\lambda_1}$ and any reduction to the interaction rate ($e^{-\\frac{2 \\Delta_{CFL}}{T}}$ in the CFL phase) leads to an enhancement of the bulk viscosity ($e^{\\frac{2 \\Delta_{CFL}}{T}}$).  Another way of describing the argument above is to say the bulk viscosity is a resonance effect between the density oscillation on one side, and the interaction rates that try to restore $\\beta$-equilibrium during these oscillations on the other \\cite{Alford:2006gy}.  The thorough reader might notice that our results for the normal quark matter are a bit different -albeit close- to those presented in Ref.~\\cite{Madsenprl2}, there are several reasons for this discrepancy, unlike Ref.~\\cite{Madsenprl2}, which considers the bulk viscosity only due to the non-leptonic weak interaction, the bulk viscosity coefficient considered here takes into account the contribution of both the non-leptonic and the \\emph{Urca} processes. Another source of discrepancy was due to the assumptions used to calculate the chemical potentials of quarks, in here we determine the chemical potential of each quark species from the charge neutrality and $\\beta$-equilibrium conditions in a fashion similar to Ref.~\\cite{SSR2}.  This manuscript represents a first step towards a more complete treatment of the r-mode instabilities of color-superconducting quark matter. As a follow up to this work, one could attempt to calculate other dissipative phenomena in deconfined quark matter, such as the shear viscosity in superfluid phases like the CFL \\cite{cristina_shear} and the CSL phase, the bulk viscosity due to superfluidity in the CFL \\cite{cristina_bulk} and the CSL phase, see also Ref.~\\cite{Mannarelli:2008jq}, or dissipation due to electron-electron scattering in the system, or due to surface friction between a rotating core and a nearly static crust in the case of hybrid stars. Supplementing these results with the ones in this article should provide us with more accurate figures for the r-mode instability window in the superfluid phases, which would then allow us to draw more firm astrophysical conclusions.  In this work a constant density approximation of the star was assumed, however, one can, in principle, generalize the calculations to include different models that might result in a density profile of the star, which will lead to a radis-dependent viscosities and dissipative timescales.  While preparing this manuscript, an article which addresses the dissipation of r-mode instabilities in the CFL phase due to the viscosities calculated in Refs.~\\cite{cristina_shear, mark_bulk} was published, Ref.~\\cite{Jaikumar:2008kh}. Some discrepancies between their treatment and the one presented here need to be pointed out, first, in order to calculate the dissipative timescale associated with bulk viscosity they used the approximation that was used in Ref.~\\cite{lindblom_owen_morsink,kokkotas_stergioulas} in which the Lagrangian perturbation of the fluid is approximated by an Eulerian perturbation. This approximation greatly overestimates the dissipative timescale of the bulk viscosity, as was shown in Refs.~\\cite{Anderssonreview, Lind2nd}. Another point worth mentioning was that the shear viscosity formula that was used is different, it has a temperature dependence of $T^{-2}$ which was derived in \\cite{Haensel:1989ja}, however, it was shown in \\cite{Heiselberg:1993cr} that a more proper treatment leads to the formula we have used in here, \\emph{i.e.} Eq.(\\ref{sqm-sv})."
    },
    "0806/0806.4053_arXiv.txt": {
        "abstract": "New high-resolution spectra of FU Ori, obtained with the HIRES spectrograph at the Keck I telescope in 2003--2006, make it possible to compare the optical line profiles with those predicted by the self-luminous accretion disk model. A dependence of line width on excitation potential and on wavelength, expected for a Keplerian disk, is  definitely not present in the optical region, nor is the line duplicity due to velocity splitting. The absorption lines observed in the optical region of FU Ori must originate in or near the central object, and here their profiles are shown to be those expected of a rigidly rotating object. They can be fitted by a rapidly rotating ($v$ sin $i$ = 70 km s$^{-1}$) high-luminosity G-type star having a large dark polar spot, with axis inclined toward the line of sight.  Over these years, the radial velocity of FU Ori has remained constant to within $\\pm$0.3 km s$^{-1}$, so there is no indication that the star is a spectroscopic binary.  These results apply to the optical region ($\\lambda< 8800$ \\AA); more distant, cooler regions of the disk contribute in the infrared. ",
        "introduction": "FU Ori is the prototype of a small group of pre-main-sequence stars, the so-called FUors.  As explained in the Introduction to \\citet{herb03}, the first members were optically detectable and so could be studied in some detail.  They were distinguished by flareups of several magnitudes followed either by near-constancy or slow decline; by broad-lined absorption spectra resembling F- to K-type supergiants in the optical region; by powerful outflowing winds shown by P Cyg structure at H$\\alpha$ and the \\ion{Na}{1} D lines; and by high \\ion{Li}{1} abundances (an indication of youth). All the known FUors are located in star-forming regions.  These characteristics defined the FU Ori class; their observed properties have been reviewed in detail by \\citet{herb77} and by \\citet{hart96}.  In recent years, some authors have chosen to take a broader and more inclusive view of  the FUor phenomenon. For example, the list of \\citet{abr04} contains some stars that are certainly not FUors in the original sense.  We see no advantage in lumping such incompletely characterized  objects in with those about which we know a great deal, such as FU Ori and V1057 Cyg. So we prefer to call them ``FU Ori-like\" (following \\citealt{gre08}) until their status can be clarified.  This paper deals with FU Ori.  One of best-documented FUors is V1057 Cyg, which before its flare-up by about 6 mag in 1971 was a faint variable star with an emission-line spectrum like that of a T Tauri star. After the flare-up, it appeared as a high-luminosity early-type star with only H$\\alpha$ in emission. It has since been in slow decline, and is now (2008) only about 1 mag (in B) above its pre-outburst brightness.  FU Ori itself increased in brightness from $m_{\\rm pg} = 16$ to 10 in 1937--1939, and since that time has remained near maximum, except for a slow decrease in optical brightness of about 0.15 mag per 10 years \\citep{cl05}.  Before the flare-up, FU Ori was an irregular variable of unknown spectral type. Optically, it is now  classified as G0, about luminosity class Ib. The spectral energy distribution shows an IR excess, typical for a pre--main-sequence star with circumstellar disk and/or dust shell. There is also excess flux in the UV region ($\\lambda$ $<$ 2600 \\AA) indicating the presence of a hot continuum of $T_{\\rm eff} \\approx  9000$ K \\citep{krav07}.  FU Ori has a faint red companion, probably a pre--main-sequence star, at a projected separation of 0\\farcs5 = 250 AU \\citep{wang04}. It is unknown whether the presence of this star has a connection to the outbursts.  A widely known concept is that FUors are solar mass pre--main-sequence stars that have increased dramatically in brightness due to a sudden episode of enhanced accretion triggered by some kind of disk instability. The rate of accretion has been estimated as $2 \\times 10^{-4} M_\\odot$ yr$^{-1}$ by \\citet[hereafter PKHN96]{pop96} by fitting an accretion disk model to the observed spectral energy distribution.  That is three orders of magnitude larger than in a classical T Tauri star. The optical spectra of FUors contain broad metallic lines that, according to \\citet[hereafter KHH88]{ken88}, originate in the atmosphere of the inner accretion disk at radii of 8--80 $R_\\odot$.  There is general agreement that the IR excess seen in the SED of FU Ori originates in an extended envelope---presumably a disk---whose outline is seen in the $K$-band interferometric image of \\cite{mal05}. In the model of Hartmann \\& Kenyon, this structure extends inward to the central star and overwhelms it: ``during outburst, the disk outshines the central star by factors of 100--1000\" \\citep{hart96}.  Thus the optical spectrum of FU Ori is thought to be formed in an atmosphere atop the surface of this optically thick disk, and the contribution of the central star is small or negligible.  Because of the plausibility of this hypothesis, and the many subsequent papers that built upon it,  it largely escaped the critical scrutiny that usually follows upon such an interesting new proposal. A simple yet critical test of the self-luminous accretion disk model can be provided by the higher-resolution optical spectra that have now become available.  These make it possible to approach the key question: whether the {\\it optical} line profiles are in fact those of a Keplerian disk or not.  One of the most explicit predictions of that FUor disk model is that line width (or $v$ sin $i$) should depend on wavelength and excitation potential, because the warm central regions of the disk rotate faster than the cooler regions farther out. Observational evidence for such relationships in the optical spectra of FU Ori and V1057 Cyg were reported by \\citet{welt92}, but not confirmed later by spectroscopic data of higher quality and resolution (\\citealt{herb03}, hereafter HPD03).  This paper deals with that issue in more detail.  Another prediction of the model was that FUor absorption lines should appear double because of Keplerian splitting.  No such splitting is observed in the high-resolution spectra of FU Ori described in this paper.  In V1057 Cyg, some (but not all) absorption lines appeared to be double, but as the star faded, it became clear that the duplicity  was caused by the emergence of emission cores in the line centers that are produced in a hot region (which was called a ``chromosphere\" by HPD03 for lack of a better word). ",
        "conclusions": "There is no doubt that FU Ori possesses a disk: radiation from the disk and envelope dominates the IR, the outline of the disk was resolved at AU scale by IR interferometry \\citep{mal05}, and the distribution of energy in the infrared is compatible with a disk model. So our question was: what is it that is seen in the optical region, a central star or an inner accretion disk?  We found that (1) all weak photospheric lines have the same line width and profile, as expected for a rigidly rotating body, but in definite conflict with prediction for the self-luminous Keplerian disk model; (2) those profiles can be explained if the central object is a rapidly rotating high-luminosity star with a dark polar spot; and (3) there is no sign of the line doubling, or the dependence of line width on wavelength (in the optical region) that is expected for the disk model.  In any disk, the temperature and rotational velocity change with radius, and this defines the width of spectral lines in integrated light. It would be strange if the disk temperature and velocity profiles were designed by nature in such a way that all atomic lines formed in different parts of the disk were of equal width.  It is more natural to assume that the optical spectrum of FU Ori is produced at the central object, while the IR spectrum is formed in the circumstellar disk (as in classical T Tauri stars), and the 2.3 $\\mu$m CO lines---a characteristic feature of FUors---are formed in a cool disk or a distant shell. Certainly there is no evidence of TiO bands or other M-type features  in the optical region: see Figure 9.  Do these results undermine the self-luminous accretion disk hypothesis in a significant way? No, but they do demonstrate that some modification is in order.  It is remarkable that the hypothesis still stands even though two of the strongest observational arguments that were originally urged in its favor are now seen, in the case of the prototype, to be invalid."
    },
    "0806/0806.0071_arXiv.txt": {
        "abstract": " ",
        "introduction": "The ion cyclotron resonance  and the ion cyclotron  mode have  been discussed   in the context of  problems related to the heating of the solar corona, e.g.,  Marsch et al. (1892), Cranmer et al. (1999), Cranmer (2000),  Tu $\\&$ Marsch (2001), Markovskii (2001), Isenberg (2001),  Hollweg $\\&$ Isenberg (2002).   This is due to various reasons, such as  the evidence obtained from in situ measurements in the solar wind and coronal holes of resonant ion cyclotron heating  (Hollweg $\\&$ Isenberg 2002), and a preferential heating of coronal ions (with respect to electrons) which is most dominant in the direction perpendicular to the magnetic field lines. The damping of such IC waves is believed to be a good candidate for the consequent coronal heating and solar wind acceleration Markovskii (2001). However,  as in many wave-heating scenarios of the corona proposed in the past, there is usually the problem of the source for the required massive generation of such IC waves. In the literature, various effects have been proposed as sources for the IC mode, like currents (Forslund 1970, Toichi 1971), global resonant MHD modes Markovskii (2001), etc. Yet, we note that these effects themselves need some source, so the problem is not solved but merely shifted to another problem. The multi-fluid description of the  beam-driven  IC  wave  excitation is presented in Mecheri $\\&$ Marsch (2007), where a tiny ion-beam  population is assumed to originate from small scale reconnections.  A rather large growth rate (of the order of $0.1 \\Omega_p$, where $\\Omega_p$ is the proton gyro-frequency) is obtained for relatively large values of the wave number, and for angles of propagation below $60^{\\circ}$.    In one of the recent studies dealing with the generation of  ion cyclotron waves in the solar atmosphere Markovskii (2001) it is pointed out that the solar atmosphere contains density inhomogeneities of various scales with density gradients in the direction perpendicular to the magnetic field. Indeed, detailed studies performed in the past, which include direct observations, reveal the existence of ray-like structures that span from very large perpendicular cross sections, like in the case of polar plumes and streamers, down to very fine filamentary structures with a cross section of the order of a kilometer. More details on that issue are given in   Woo $\\&$  Habbal (1997), Woo (1996), November $\\&$  Koutchmy (1996), Karovska $\\&$ Habbal (1991), and in references cited therein.  Fine filaments are visible even from ground-based observations like those during the eclipse in 1991 (November  and  Koutchmy  1997), showing a slow radial enlargement of the structures (i.e., in the direction out of the Sun) which is consistent with a low-beta plasma. In situ measurements by Voyager~1 and Voyager~2  (Woo $\\&$  Habbal 1997), show that  the finest structures in the slow solar wind at around $9\\;R_{\\odot}$ are about 3 times finer than those in the fast wind. Assuming a radial expansion, they conclude that the transverse sizes of these highly elongated structures at the Sun are below 1$\\;$km. In   Karovska $\\&$  Habbal (1991) an image restoration  method is used to study the structures of the quiet Sun, with a spatial resolution of $5\"$ and a temporal resolution of 5.5 minutes. The contour maps presented in the work reveal the existence of numerous structures of various sizes.   The presence of density gradients in the direction perpendicular to the magnetic field lines implies the possibility of  the existence of a  drift wave, a mode with the unique feature of being unstable both in the kinetic and the fluid domain (hence the term {\\em universally unstable mode}). Within  kinetic theory, the drift mode is unstable even for a Maxwellian electron distribution. This is due to the inverse electron Landau damping effect. Within fluid theory, the mode is also unstable, due to the common effect of the electron collisions and the ion inertia, in the presence of a background density gradient.  A comparison between the kinetic and the fluid instability (Goldston $\\&$ Rutherford 1995) reveals that the resistive one is dominant provided that the electron parallel mean-free path is smaller than the parallel wavelength. This is essentially the reason for the interest in the drift modes with very large parallel wavelengths and relatively short perpendicular wavelengths.  One particularly interesting feature of the drift wave is related to its intrinsic nonlinearity (Hasegawa $\\&$ Sato 1989). The nonlinear terms, which follow from the convective derivatives in the momentum equations, are comparable to the linear ones provided that $k_\\bot^2\\rho_s^2 e\\phi/(\\kappa T_e)$ is of the same order as $\\rho_s/L_n$. Here, $\\phi$ is the electrostatic wave potential, $\\rho_s=c_s/\\Omega_i$, $\\Omega_i$ is the ion gyro-frequency,   $c_s^2=\\kappa T_e/m_i$, $L_n$ is the scale length of the density gradient,  and $\\rho_s/L_n$ is usually a small quantity.  Hence, the mode becomes nonlinear even for very small perturbations provided that $k_\\bot \\rho_s$ is not much larger than unity. Nevertheless, in the literature dealing with the waves and instabilities in solar plasma, the effects related to the density gradient and the consequent drift waves are usually disregarded.  In our recent publications, some  aspects of the drift wave instability  in the solar plasma have been discussed.  The collisional coupling between the drift and kinetic Alfv\\'en waves in the upper solar atmosphere (Vranjes $\\&$ Poedts 2006)  reveals a strongly  growing drift mode in the chromospheric plasma. The similar collisional instability of the mode, though with a much  smaller increment, has been demonstrated also for the coronal plasma. In both cases, the kinetic Alfv\\'en part of the mode is shown to be collisionally damped. In the presence of a plasma flow along the magnetic field lines the drift mode is subject to a reactive-type instability, provided that this background velocity has a gradient in the perpendicular direction. Such a problem has been discussed in our recent work (Saleem, Vranjes $\\&$ Poedts 2007) dealing with solar spicules. The drift mode has been shown to be unstable for typical spicule characteristic lengths of the density and the shear flow gradients, i.e.\\ in the range of a few hundred meters up to a few kilometers, yielding wave frequencies of the order of a few Herz.  One of the basic properties  of the drift mode  is the low frequency $|\\partial/\\partial t|\\ll \\Omega_i$.  In the case when this  low-frequency limit is not well satisfied, there appears a coupling between the drift and ion-cyclotron modes. In the limit when  $|\\partial/\\partial t|\\sim \\Omega_i$,  this  coupling is very effective and well known (Ichimaru 1973). In the simplest case, it yields a coupled unstable drift-cyclotron  mode with the instability driven by the plasma density gradient. The first demonstration of this instability was  performed by Mikhailovskii $\\&$  Timofeev (1963).  To the best of our knowledge, the   perpendicular density gradients, and the coupling of the drift and IC modes in the regime $\\omega \\sim \\Omega_i$ together with the consequent wave instability, have never been discussed in the context of the heating of the solar corona and the acceleration of the solar wind. In the present work we  describe the basic instability of the mode  to  focuss attention   on  the existence of such an instability that may be widespread and may contribute substantially to the problems discussed above. The term 'fine structures',  used in the text above for the structures observed in the solar corona still  denotes spatial scales that are much larger than those used in our  present study, in Sections 2 and 3. In other words, there is no  observational support for the existence of such tiny density inhomogeneities, and it is not expected in the near future. Yet, in view of the  variety of density structures at larger and observable spatial scales, and having no obvious physical effect that would prevent inhomogeneities at  scales that are even shorter than those observed so far,  such micro scale plasma inhomogeneities cannot be excluded. ",
        "conclusions": "The instability discussed here implies very short scale lengths for the inhomogeneity of the equilibrium density and/or a  very weak magnetic field. Only in such circumstances can the frequencies of the drift and IC modes  become close to each other so that the two modes can effectively couple. Short density scale lengths are presently not directly observable in the solar atmosphere. However, we have learned   that  improvement of the resolution in observations of phenomena in the solar atmosphere tends to result in an increased variety and complexity of the density and/or magnetic field structures at short scales, showing very tiny filaments at scales below 1~km. Therefore, one may expect a larger variety at even shorter scales. The present analysis clearly demonstrates the instability of perpendicularly propagating modes at frequencies in the range of the ion cyclotron frequency and at wavelengths of the order of the ion gyro radius. The tiny filaments are very elongated, they may extend to many solar radii, and the growth of the mode and the consequent dissipation and heating of ions  may take place over  large distances. The numbers used here are for the solar corona however, the large radial (from  the Sun) length of the structures and the consequent decreasing of the magnetic field intensity implies larger density scale lengths at which the instability takes place. This can be easily shown by reducing the magnetic field to $10^{-4}\\;$T and the number density by one order of magnitude. As a result, the necessary density scale length $L_n$ for the unstable modes becomes of the order of 60 meters.  Therefore the development of unstable growing modes may take place at large distances along the same density filaments that pervade the corona and spread within the solar wind. The  presence of hotter ions in and around such filamentary structures should be interpreted as an indication and a signature of the instability.  There have  been  intensive searches for possible mechanisms for the excitation of such modes. Our analysis  is based on one  such mechanism that is well known, but  not used in the context of the solar plasma, and it requires a step  beyond the widely used MHD model. At higher  densities or temperatures the plasma beta becomes higher so that electromagnetic effects should be included. However, it is known that eventual bending of the magnetic field lines implies an additional obstacle for electron motion along the magnetic lines, making the drift-type modes even more unstable (Vranjes et al. 2007b). The parameter values that we use here are realistic and the instability that has been demonstrated for the given cases is thus physically very likely.    Acknowledgements:  These results are  obtained in the framework of the projects G.0304.07 (FWO-Vlaanderen), C~90205 (Prodex), GOA/2004/01 (K.U.Leuven),  and the Interuniversity Attraction Poles Programme - Belgian State - Belgian Science Policy."
    },
    "0806/0806.4008_arXiv.txt": {
        "abstract": "We report the discovery of a planet transiting a moderately faint ($V=12.3$\\,mag) late F star, with an orbital period of \\hatcurP\\,days. From the transit light curve and radial velocity measurements we determine that the radius of the planet is $\\rpl = \\hatcurr\\,\\rjuplong$ and that the mass is $\\mpl = \\hatcurm\\,\\mjuplong$. The density of the new planet, $\\rhopl = \\hatcurrho\\,\\gcmc$, fits to the low-density tail of the currently known transiting planets. We find that the center of transit is at $T_{\\mathrm{c}} = \\hatcurT$ (HJD), and the total transit duration is \\hatcurdur\\,days. The host star has $\\mstar=\\hatcurms\\,\\msun$ and $\\rstar=\\hatcurrs\\,\\rsun$. ",
        "introduction": "\\label{sec:intro}  Transiting extra-solar planets are important astrophysical objects as they allow to test planetary structure and evolution theory \\citep[e.g.,][]{fortney08a, burrows08, baraffe08} especially because they yield a measurement of the planetary mass and radius. Over the last few years, the sample of known transiting planets has grown substantially, leading to an improved theoretical understanding of their physical nature \\citep[e.g.,][]{Guillot06, burrows07, fortney08b, chabrier07}. In addition, the increasing sample of transiting planets enabled the discovery of some interesting correlations, such as the mass-period relation \\citep{mazeh05, gaudi05}. The astrophysics behind these correlations is not fully understood, and a clear way towards progress is the discovery of many more transiting planets.  We report here the discovery of another transiting planet detected by the HATNet project\\footnote{http://www.hatnet.hu} \\citep{bakos02,bakos04}, labeled \\hatcurb, and our determination of its parameters, including mass, radius, density and surface gravity. In \\secr{anal} we describe our photometric and spectroscopic observations and in \\secr{stelpar} we derive the stellar parameters. The orbital solution is performed in \\secr{rv} and a discussion of possible blend scenarios is brought in \\secr{blend}. The determination of the \\lc\\ parameters and the planet's physical parameters is described in \\secr{planpar} and we bring a discussion in \\secr{disc}. ",
        "conclusions": "\\label{sec:disc}  \\ifthenelse{\\boolean{emulateapj}}{\\begin{figure}[t]}{\\begin{figure}[t]} \\ifpdf \\plotone{f5.pdf} \\else \\plotone{f5.eps} \\fi \\caption{ {\\it Top panel:} The radius-mass diagram for the known transiting planets, in log-log scale, with \\hatcurb\\ marked by an open circle. Constant density lines are over-plotted. The diagram visually shows that the new planet has a relatively low density. The Neptune-mass planet GJ~436b, orbiting an M star, is positioned beyond the lower-left corner of this diagram.\\\\ {\\it Bottom panel:} The mass-period diagram for the known transiting planets, in linear scale, with \\hatcurb\\ marked by an open circle. Three planets are positioned outside the boundaries of this diagram: The long-period planet \\hd{17156b}, and the massive planets HAT-P-2b (a.k.a.~\\hd{147506b}) and XO-3b. The planet with the lowest mass in this diagram is GJ~436b, orbiting an M star, and the one with the largest mass is CoRoT-Exo-2b.\\\\ This figure is based on data taken from http://www.inscience.ch/transits/ on 2008 June 1st. \\label{fig:rmp}} \\ifthenelse{\\boolean{emulateapj}}{\\end{figure}}{\\end{figure}}  We present here the discovery of a new transiting extra-solar planet, \\hatcurb, with a mass of \\hatcurm\\,\\mjup, radius of \\hatcurr\\,\\rjup\\ and orbital period of \\hatcurP\\,days. The $V=12.3$\\,mag host star is among the faintest planet-host stars identified with small-aperture, wide-field, ground-based campaigns. In fact, together with WASP-5 \\citep{anderson08} they are currently the faintest planet-host stars discovered with such campaigns, where the transit depth is below 2\\%.  To show the main characteristics of the new planet relative to the currently known transiting planets we plot in \\figr{rmp} the position of the latter on the radius-mass diagram, on the top panel, and the mass-period diagram, on the bottom panel. In both panels \\hatcurb\\ position is marked by an open circle.  The radius-mass diagram, presented here in log-log scale, visually shows that the new planet has a low density, of \\rhopl = \\hatcurrho\\,\\gcmc, similar to that of \\hd{209458b} \\citep[e.g.,][]{knutson07, mazeh00}, WASP-1b \\citep{cameron07, shporer07, charbonneau07}, HAT-P-1b \\citep{bakos07, winn07b} and CoRoT-Exo-1b \\citep{barge08}.  Comparing the measured radius to the modeled radius of \\citet{burrows07}, for the same stellar luminosity, semi-major axis and (core-less) planet mass as the \\hatcur\\ system, shows the measured value is larger than the modeled one by 1-2\\,$\\sigma$. In this comparison we assumed solar opacity for the planetary atmosphere. As shown by \\citet{burrows07}, the transit radius effect together with enhanced planetary atmospheric opacities, relative to solar opacity, can be responsible for inflating the planetary radius. Since atmospheric opacities are related to the host star metallicity, and \\hatcur\\ metallicity seems to be super-solar, it is likely that these two effects account for the radius difference, as is the case for WASP-1b \\citep[][see their Fig.~7]{burrows07}. Other mechanisms that may explain the increased radius involve downward transport and dissipation of kinetic energy \\citep{showman02} and layered convection \\citep{chabrier07}.  According to its \\teq\\ and Safronov number, listed on \\tabr{orb}, \\hatcurb\\ is likely a Class~II planet, as defined by \\citet{hansen07}. This classification is in accordance with its low density, since members of this class usually have smaller masses and larger radii than those of Class I.  The stellar incident flux at the planet is \\fpl\\ = \\hatcurfp\\,10$^9$\\,erg cm$^{-2}$ s$^{-1}$, which is at the lower end of the \\fpl\\ range for pM planets, according to the pL/pM planet classification of \\citet{fortney08b}. Since it is positioned near the transition region in \\fpl\\ between these two classes of planets it may be used to further characterize them.  In the mass-period diagram the new planet is positioned near XO-1b \\citep{mccullough06, holman06}, OGLE-TR-182b \\citep{pont08}, OGLE-TR-111b \\citep{pont04, winn07a} and HAT-P-6b \\citep{noyes08}. The position of \\hatcurb, along with that of neighbouring planets, suggests that the mass-period relation \\citep{mazeh05, gaudi05}, i.e., the decrease of planetary mass with increasing orbital period, levels-off at periods $\\gtrsim 3$\\,days.  The ephemeris obtained here is based on the four follow-up \\lcs. We compared it to the mid-transit time of the HATNet \\lc, obtained some 1000\\,days earlier, by propagating the error on the period. The result shows a difference of about 1.5\\,$\\sigma$ between the mid-transit times of the follow-up \\lcs\\ and that of the HATNet \\lc. Future follow-up \\lcs\\ will be able to study more thoroughly this possible discrepancy.  Finally, we note that due to the line-of-sight stellar rotation velocity of \\hatcurvsini\\,\\kms\\ and the non-zero impact parameter \\hatcurimp, this planet is a good candidate for the measurement of the Rossiter-Mclaughlin effect \\citep[e.g.,][]{winn05}. Eq.~6 of \\citet{gaudi07} gives an expected amplitude for the effect of more than 100\\,\\ms, which is larger than the orbital amplitude. However, since this is a relatively faint star, with $V=12.3$\\,mag, such a measurement will be challenging."
    },
    "0806/0806.4474.txt": {
        "abstract": "We present echelle long-slit optical spectra of a sample of objects evolving off the Asymptotic Giant Branch (AGB), most of them in the pre-planetary nebula (pPN) phase, obtained with the ESI and MIKE spectrographs at the 10\\,m Keck\\,II and 6.5\\,m Magellan-I telescopes, respectively. The total wavelength range covered with ESI (MIKE) is $\\sim$3900 to 10900\\AA\\ ($\\sim$3600 to 7200\\AA). In this paper, we focus our analysis mainly on the \\hal\\ profiles. Prominent \\hal\\ emission is detected in half of the objects, most of which show broad H$\\alpha$ wings (with total widths of up to $\\sim$4000\\,\\kms). In the majority of the \\hal-emission sources, fast, post-AGB winds are revealed by P-Cygni profiles. %For these objects the %jet-sculpting of the AGB envelope is, therefore, most likely currently %active. In $\\sim$37\\% of the objects \\hal\\ is observed in absorption. In almost all cases, the absorption profile is partially filled with emission, leading to complex, structured profiles that are interpreted as an indication of incipient post-AGB mass-loss. The rest of the objects ($\\sim$13\\%) are \\hal\\ non-detections. We investigate correlations between the \\hal\\ profile and different stellar and envelope parameters. All sources in which \\hal\\ is seen mainly in absorption have F-G type central stars, whereas sources with intense \\hal\\ emission span a larger range of spectral types from O to G, with a relative maximum around B, and also including very late C types.  Shocks %resulted from the interaction between the post-AGB %wind and the AGB envelope may be an important excitation/ionization agent of the close stellar surroundings for objects with late type central stars. Sources with pure emission or P Cygni \\hal\\ profiles have larger $J-K$ color excess than objects with \\hal\\ mainly in absorption, which suggests the presence of warm dust near the star in the former. The two classes of profile sources also segregate in the IRAS color-color diagram in a way that intense \\hal-emitters have dust grains with a larger range of temperatures. Spectral classification of the central stars in our sample is presented. For a subsample (13 objects), the stellar luminosity has been derived from the analysis of the \\ion{O}{1}\\,7771-5\\AA\\ infrared triplet. The location in the HR diagram of most of these targets, which represent $\\sim$30\\% of the whole sample, is consistent with relatively high final (and, presumably, initial) masses in the range $M_{\\rm f}$$\\sim$0.6-0.9\\msun\\ ($M_{\\rm i}$$\\sim$3-8\\msun). %We report and discuss the possible raid %post-AGB evolution of some pPNs. %objects in the last 15 years. ",
        "introduction": "\\label{intro}  Intermediate mass stars ($\\sim$1-8\\msun) evolve from the Asymptotic Giant Branch (AGB) to the Planetary Nebula (PN) phase through a short-lived ($\\sim$\\,10$^3$\\,yr) and fascinating evolutionary stage designated as the post-AGB (pAGB) or pre-planetary nebula (pPN) phase.  At some point in the late-AGB or early pAGB stage, a process (or processes) becomes operative that accelerates and imposes severe asymmetries upon the slow, spherical AGB winds: the spherical, slowly expanding ($V_{\\rm exp}$\\,$\\sim$\\,15\\,\\kms) AGB circumstellar envelope (CSE) becomes a PN with clear departures from sphericity and fast ($\\ga$\\,100\\,\\kms) outflows directed along one or more axis.  Although there is no consensus yet for what causes this spectacular metamorphosis, fast jet-like winds have been hypothesized to play an important role \\cite[see e.g.\\,the review paper on PNs shaping by][]{bal02}. These outflows carve out an imprint within the AGB CSE producing and shaping the fast, bipolar lobes observed in most pPNs and PNs \\citep{sah98}. The mechanism that powers and collimates pAGB jets is a fundamental issue on stellar evolution that remains a mystery.  %The mechanism that drives and %collimates the pAGB winds remains still unknown. %What produces bipolarity in these objects and at what stage does bipolarity %manifest itself are key questions that remain yet poorly understood.  Optical spectroscopic observations of pPNs and PNs have allowed considerable advances in our understanding of pAGB evolution and, in particular, are very useful for probing the interaction between pAGB winds and the CSE formed in the previous AGB phase.  Our current (very limited) knowledge of pAGB evolution and, more particularly, of pAGB winds is derived mainly in two ways. The first is {\\sl indirect}, that is, based on the effects of the pAGB winds on the AGB CSEs.  Many pPNs/PNs show extended lobes with, often, bow-shaped features at their tips that are visible through optical recombination and forbidden emission lines.  Spectroscopic observations have been crucial for understanding the origin of these regions, which are found to be excited by the passage of fast ($\\ga$100\\,\\kms) shocks \\citep[e.g.][]{san00,vaz00}. From these results we infer the existence of fast, pAGB winds that interact hydrodynamically with the AGB CSE leading to the formation of shocks and, ultimately, to the acceleration and shaping of the nebular material.  {\\sl Direct} detection of pAGB winds is limited to a very few objects. In some cases, there is Balmer and forbidden line emission arising in a set of compact, shock-excited regions located along the nebular axis \\citep[e.g.][]{buj98,rie06}. These `knots', which usually move away from the star at high velocity, are thought to result from the propagation of shocks in the pAGB wind itself, suggesting that pAGB winds are collimated and directed along the nebular axis. In some pPNs/PNs, fast, pAGB winds are also revealed by P-Cygni profiles close to the central star. One of the most exciting results recently provided by long-slit spectra with $HST$/STIS is the first direct observation of the spatio-kinematic structure of a fast collimated pAGB wind (jet) in a `pristine' stage, i.e$.$, not strongly altered by the interaction with the AGB shell \\citep{san01}. The characterization of the pristine jet in Hen\\,3-1475, % with the largest velocity ever measured in a pPN which is collimated and expands at more than 2000\\kms\\,at only $\\approx$10$^{16}$\\,cm from the central star, has had a direct impact on the different theories of pAGB wind collimation, ruling out purely hydrodynamical processes as the collimation agent in this object.  Recognizing the importance of optical spectroscopy of pAGB objects we have recently obtained optical long-slit spectra a large sample of objects evolving off the AGB, most of them pPNs and young PNs, in order to characterize the pAGB mass-loss process and the jet-sculpting of AGB winds. This forms part of our extensive, multi-wavelength survey program of imaging and spectroscopic observations of OH/IR stars (evolved mass-losing stars with OH maser emission), which may be the youngest pPNs \\citep[e.g.][]{sahs07,san04}.  In this paper, we present the spectroscopic database resulting from our observations.  Our work complements recent optical spectroscopic surveys of pPNs/PNs \\cite[e.g.][]{sua06,per07} by providing deep, high spectral resolution spectra covering a wide wavelength range.  Here, we focus on the characterization and analysis of the different types of H$\\alpha$ profiles and their correlation with several stellar and envelope parameters. %, namely, the spectral type and luminosity of the %central star, the NIR and IRAS colors, the chemistry, the nebular %morphology, and the galactic distribution. The shape of the line profile is interpreted in terms of present-day mass-loss in these objects, which enable obtaining information on their on-going pAGB winds. We have also derived spectral types based on the absorption line spectrum of our targets. Whenever available, the \\ion{O}{1}\\,7771-5\\AA\\ infrared triplet have been used as a luminosity indicator.    %{\\bf VLSR are obtained and used to derive kinematical %distances -- UNDER CONSIDERATION}. second paper - need to do it on %a large number of stellar lines. *** WARNING *** if the stellar continuum is %scattered then the absorption  lines will be REDSHIFTED. %% *** WARNING *** %The analysis of the nebular emission lines to derive the physical %properties of the envelope and relative contribution to shocks and UV %stellar radiation to the gas excitation/ionization as well as modeling %of the spatio-kinematic structure of some objects with spatially resolved %nebulosities is deferred to a second paper. %Also, we intend to perform detailed studies of the stellar parameters and %chemical composition through model-atmosphere methods of individual %targets.  %The remainder of this paper is organized as follows: %Section\\,\\ref{obs} describes the observations, data reduction, and the %object sample; In Section\\,\\ref{res}, we show the extracted one-dimensional %spectra and describe our %observational results, including the different types of \\hal\\ profiles %and other non-\\ion{H}{1} lines present in the spectra of our %targets. In Sections\\,\\ref{anal} and \\ref{discuss} we describe the %data analysis, which includes \\hal\\ line parameterization, spectral %typing, and luminosity estimates, and discuss our results, %respectively. A summary of our main results from this work are provided in %Section\\,\\ref{sum}.  %% discusses... %% gives ",
        "conclusions": "\\label{discuss}   \\subsection{Fast pAGB winds probed by P Cygni like profiles} \\label{dis_pcyg}  P Cygni profiles are known to be indicative of on-going mass-loss in the form of a stellar wind. %can be studied to derive information on the velocity law and mass-loss %rate of the wind. The formation of typical P Cygni profiles can be understood qualitatively by a simple model of a spherically symmetric wind in which the velocity increases outwards up to a terminal velocity of the outflow $v_\\infty$.  The wind material in front of the star is moving towards the observer and, therefore, absorbs the stellar continuum producing a blue-shifted absorption feature with a Doppler shift between $-$$v_\\infty$ and 0, relative to the source systemic velocity (\\vsys). Except for the part that is occulted by the star, the rest of the wind, which forms a halo around the star, will produce broad line emission centered around \\vsys\\ with wings extending up to $\\pm$$v_\\infty$.  The sum of the emission and the absorption at each Doppler velocity range leads to a so called P Cygni profile. % with a blue-shifted absorption component extending up %to $-$$v_\\infty$ plus an emission component with its red emission wing %extending up $+$$v_\\infty$. These P Cygni profiles are often seen in UV resonance lines of the central stars of PNs \\citep[e.g.][]{peri93,kud97}.   The \\hal\\ profiles observed in the majority of our targets differ from the typical P Cygni profiles described above. The emission component is not centered at \\vsys\\ but, rather is red-shifted in all cases, except for IRAS\\,21282$+$5050 (Table \\ref{t_pcyg}).  Moreover, the terminal velocity of the outflow, as measured from the blue-shifted edge of the absorption feature, is smaller than that obtained from the edge of the red emission wing, and in some cases, e.g.\\ the yPN M\\,1-92 (see below), the blue absorption is well detached from the line emission core. %% ---------------------------------------------------------------------- %% v$_\\infty$ is given by the edge of the feature... %% --- %,  typically $\\ga$1000\\kms.... Unsaturated P-cyg profiles are good for mass-loss estimate %% but saturated profiles are best for deriving accurate velocity laws. %% --- These peculiar H$\\alpha$ line-shapes can be explained if the emission and absorption form in two distinct nebular components, with different velocity fields and, more generally, affected by different broadening mechanism. In particular, we propose the following model: the broad line emission, characterized by a given $intrinsic$ (not necessarily P Cygni) profile, arises from a compact central source; the line emission (and stellar continuum) is scattered by dust in the walls of the nebular lobes, which should produce an overall red-shift of the line profile as a whole relative to \\vsys\\ as observed. The blue-shifted absorption is due to neutral or partially ionized outflowing gas inside the lobes absorbing the scattered photons along the line of sight. This results in a P Cygni like profile where the blue-shifted absorption is produced against the intrinsic emission profile. Such a model has been successful in producing a detailed fit to the spatio-kinematic distribution of the blue-shifted absorption and the red-shifted scattered line core in the H$\\alpha$ line profiles observed with STIS/$HST$ in the yPN Hen\\,3-1475 (S\\'anchez Contreras \\& Sahai 2001; see their Figure\\,1d for a sketch of our geometrical model). A similar scenario is proposed by \\cite{arr05} to explain the P Cygni like profile of \\hal\\ and other recombination lines in the yPN M\\,1-92, which shows a very broad, blue absorption feature well detached from the emission component that cannot result from the single stellar wind scenario used to explain typical P Cygni profiles. % that extends few$\\times$10$^2$ to $\\sim$$\\times$10$^3$\\,\\kms. Our interpretation is also consistent with the fact that most pPNs exhibit intrinsic polarization of the continuum as well as some emission lines, which are most likely due to scattering off dust grains \\citep{tram94}. In particular, Balmer lines seem to come from an \\ion{H}{2} region close to the star, which is blocked from direct view by, e.g.,\u00a0the thick equatorial dust waists commonly present in these objects, and are seen in reflection off the nebular dust \\cite[see also, e.g.,][]{san02}.  %{\\bf it should be relative to \\vsys}, which is expected to represent the projected %velocity of the bulk of the outflow ($v$). %The terminal velocity of the pAGB outflow $v_\\infty$ is given by the edge of the absorption %feature relative to the systemic velocity, \\vsys. %The values obtained for $v_\\infty$ imply the presence of stellar winds with velocities of %the order of $\\approx$100\\kms\\ in all pcyg sources.   The nature of the outflow producing the absorption in pcyg sources is unknown, however, as in the case of Hen\\,3-1475, we argue that it is most likely a ``pristine\" pAGB wind from (or near) the central star that has not been strongly altered by its interaction with the progenitor AGB wind. The projected speed of the bulk of the material in the pAGB outflow, derived from the centroid of the blue-shifted absorption feature relative to the systemic velocity in our sample, is approximately $\\sim$100\\kms\\ (Table\\,\\ref{t_pcyg}); however the absorption features are quite wide and reach larger velocities at the edge indicating the presence of smaller amounts of material expanding as fast as, e.g., $v_\\infty$\\,$\\sim$800\\kms\\ in the case of IRAS\\,19520$+$2759. %% These velocities are much larger than the typical expansion speed %% of AGB winds. As for Hen\\,3-1475, it is plausible that a radial and or latitudinal velocity gradient exists in the outflows and/or that distinct wind components are present; the latter is suggested by the two different absorption features observed in the \\hal\\ profiles of a few objects (\\S\\,\\ref{res_hi}). Unfortunately, ground-based observations do not allow us to conclude on the isotropic or collimated nature of the pAGB winds probed by the observed P Cygni profiles, however, it is plausible that these are bipolar/multipolar ejections, as in Hen\\,3-1475 and, probably, M\\,1-92.  %The velocities of pAGB winds as derived from the \\hal\\ P Cygni %profile are uncertain (most likely lower limits) because these are %affected by the broad emission wings. %% High angular resolution spectroscopy is needed to see whether or not the line emission follows the same %% distribution as the scattered continuum and confirm the second scenario. High angular resolution long-slit %% spectroscopy is also needed to check whether these pAGB outflows are collimated.  \\subsection{Broad \\hal\\ emission wings} \\label{dis_wings}   The emission component observed in the \\hal\\ profile in our pcyg sources consists of an intense core plus weaker broad wings that reach widths (measured as the line FWZI) of up to $\\sim$4000\\kms. The yet unclear origin of the very extended \\hal\\ wings, which are also observed in a number of PNs, has been discussed by \\cite{arr03}. These authors consider several possible mechanisms for line broadening and conclude that Raman scattering of photons with wavelength close to Ly$\\beta$ by neutral hydrogen is the most probable one in 12 of the 13 objects in their sample.  In this scenario, Ly$\\beta$ photons with a given velocity width are converted to optical photons that fill the \\hal\\ wing region. %The width of the scattered \\hal\\ is proportional to the initial width %of the Ly$\\beta$ line, the Raman-scattering broadening factor being of %$\\lambda_{\\rm H_\\alpha}$/$\\lambda_{\\rm Ly_\\beta}$=6.4. Therefore, to produce %$\\sim$1000\\kms-wide \\hal\\ wings, the Ly$\\beta$-emitting region must be characterized by a typical velocity of the %order of 150\\kms. The conclusion of \\cite{arr03} is mainly supported by a) the fit of the \\hal-wing profile to a $\\propto \\lambda^{-2}$ law; b) the presence of other Raman features produced in emission; and c) the presence of a significant neutral hydrogen component.  Criteria a) to c) are to be satisfied for the line wings to be attributable to Raman scattering.   We have checked whether or not the same conclusion can be drawn from our sample. As we have shown in Section \\ref{proffit}, the \\hal\\ wings observed follow a $\\propto \\lambda^{-2}$ law in all cases, except for IRAS\\,20462$+$3416 and IRAS\\,19306$+$1407. In these two objects, which happen to display \\hal\\ wings with similar widths and shapes, Raman scattering can be ruled out as the main line broadening mechanism. \\cite{arr03} found the same disagreement between the observed and predicted Raman profile for IRAS\\,20462$+$3416, which was also in their sample, concluding that a strong stellar wind is the most probable mechanism for line broadening in this case. We have searched for other emission features produced by Raman scattering in our spectra. In particular, we looked at the 6830 and 7088\\AA\\ bands that correspond to the Raman-scattered \\ion{O}{6} $\\lambda\\lambda$1032,1038 doublet. Only four objects, IRAS\\,08005$-$2356, M\\,1-92, Hen\\,$-$1475, and IRAS\\,22036$+$5306, show emission around 6830\\AA\\ with FWZI$\\sim$150\\kms, and only in two of them, Hen\\,$-$1475 and IRAS\\,22036$+$5306, a weak emission feature is tentatively detected around 7088\\AA\\ with FWZI$\\la$100\\kms\\ (Fig.\\,\\ref{f_fullspec}). In contrast, \\cite{arr03} find these and other Raman features with large FWZIs of a few$\\times$10$^2$-10$^3$\\,\\kms, i.e$.$\\ comparable to the width of the \\hal\\ line wings, in most of their objects.  Since both the $\\lambda^{-2}$ profile of the wings and the presence of a large column density of neutral hydrogen are necessary but not sufficient conditions for attributing the observed \\hal-wings to Raman scattering, we cannot conclude that in our sample this is the main mechanism for line broadening except maybe for the four objects listed above showing other Raman-scattered emission features. Note, however, that the FWZIs of such features are significantly smaller than those measured in Arrieta's sample and smaller than the \\hal\\ wings themselves. %% In two of the four objects, there are other indications that stellar %% winds are maybe a more plausible mechanism for line broadening than %% Raman scattering.  As suggested by \\cite{arr03}, further investigation of the effect of Raman scattering producing the line wings could be obtained by measuring the widths of UV lines produced in the same region as the Ly$\\beta$ photons. This is because the width of the scattered \\hal\\ is proportional to the initial width of the Ly$\\beta$ line, the Raman-scattering broadening factor being of $\\lambda_{\\rm H\\alpha}$/$\\lambda_{\\rm Ly\\beta}$=6.4. We searched the IUE archive for UV spectra of our targets around the line \\ion{Si}{3}]$\\lambda$1892\\AA, which presumably arises in the same region as Ly$\\beta$. Unfortunately, such data only exist for four objects and only in one case, Hen\\,3-1475, the Si III] line is clearly detected. The line, which was observed with the low spectral resolution ($\\Delta\\lambda$$\\sim$6\\AA) SWP, is found to be spectrally unresolved. Nevertheless, for Raman scattering being able to produce the observed $\\sim$4000\\kms-wide \\hal\\ wings in Hen\\,3-1475, the Ly$\\beta$-emitting region must be characterized by a typical velocity of $\\sim$600\\kms, i.e$.$\\ the Si III] line should be spectrally resolved with a FWHM of $\\sim$7\\AA\\ as observed with SWP, in constrast to what is found. Therefore it seems that Raman scattering is not the main mechanism affecting the profile of the \\hal\\ wings in this case.  One additional requirement for Raman scattering to be efficient is that a relatively strong incident Ly$\\beta$ flux has to be produced by the central source. This may be satisfied by Arrieta's sample, formed by objects with central stars with \\teff$\\geq$20,000\\,K, except maybe for the hypergiant IRC$+$10420 with a stellar spectral type of A5 (\\teff$\\sim$8500\\,K) at the time Arrieta's data were obtained.\\footnote{The spectral type of the central star of IRC$+$10420 changed from F8 to A5 in 4 years, from 1992 to 1996 \\citep{klo97}.} An estimate of the Ly$\\beta$ luminosity expected around the central star of the PN IC\\,4997 using the photoinization code CLOUDY has been done by Lee \\& Hyung (2000). In this case, for a very hot (\\teff=60,000\\,K) and quite small ($R_*$=\\rsun) central star, a compact \\ion{H}{2} region of high density ($\\approx$10$^9$-10$^{10}$\\cm3) together with a column density of neutral hydrogen of $N_{\\rm HI}$=few$\\times$10$^{20}$\\,\\cm2 in the scattering region are needed in order to have a sufficient number of Ly$\\beta$ photons and efficient Raman scattering. In our sample, a significant number of objects showing broad \\hal\\ wings have relatively cool central stars, with spectral types F-G (\\teff$\\sim$5000-7000\\,K) and, therefore, are unlikely to produce enough UV photons around Ly$\\beta$. This is in fact supported by their featureless nebular spectra, which are remarkably different from those observed in more evolved PNs (\\S\\,\\ref{res_nohi}). For the objects in our sample with the hottest central stars (with late-O and early-B spectral types, \\teff$\\sim$20,000-30,000\\,K) estimates of the Ly$\\beta$ flux will be presented in future individual papers. We will also investigate whether or not shocks resulting from the wind interaction process are able to provide a strong Ly$\\beta$ flux.   The presence of fast stellar winds is another plausible mechanism to explain the broad \\hal\\ wings observed. As already pointed out by \\cite{arr03}, this is indeed the most likely mechanism for IRAS\\,20462$+$3416, for which the \\hal\\ wing profile was found not to follow a $\\propto \\lambda^{-2}$ law. For similar reasons, a fast ($\\sim$1400\\kms) stellar wind is also very likely being expelled by the central star of IRAS\\,19306$+$1407 leading to its \\hal\\ wings. In more than one third of the objects in our sample, the presence of stellar winds is, in fact, confirmed by P Cygni profiles in \\hal\\ and other lines. The question is whether or not these winds reach much higher velocities (of 1000-2000\\kms) in certain regions that can explain the broad wings observed. In the case of M\\,1-92 and Hen\\,3-1475 there is direct evidence for fast winds with expansion velocities of up to $\\sim$800 and 2300\\kms, respectively \\citep{arr05,san01}, so it is plausible that these winds are also responsible (at least partially) for their \\hal\\ wings. In the rest of the objects in our sample, the terminal speed of the fast wind measured from the blue-shifted absorption feature of the \\hal\\ P Cygni profile ($v_\\infty \\approx V_{\\rm max}$$\\approx$10$^2$\\,\\kms, Table\\,\\ref{t_pcyg}) is never as large as that implied by the broad wings if these were to be attributed to fast outflows. However, we find that $V_{\\rm max}$ is normally larger than the FWHM of the \\hal\\ core indicating that indeed there are winds moving faster than the region emitting the bulk of the \\hal\\ line (with an average velocity of $\\sim$120\\,\\kms; Table\\,\\ref{t_pcyg}).  This clearly suggests the presence of velocity gradients in the stellar winds of these objects. Moreover, we have found a weak correlation between $V_{\\rm max}$ and the FWZI of the broad \\hal\\ wings: the sources with the largest values of $V_{\\rm max}$, namely, IRAS\\,19520$+$2759, IRAS\\,20462$+$3416, M\\,1-92, and Hen\\,3-1475, exhibit the broadest wings, reaching widths of $\\sim$2600, $\\sim$2800, $\\sim$3000, and $\\sim$4000\\kms, respectively.  (For IRAS\\,20462$+$3416 and Hen3-1475, the \\hal\\ absorption feature is partially masked out by emission, therefore, values of $V_{\\rm max}$=$-$300\\kms\\ are measured from their H$\\beta$ P Cygni profiles; Fig.\\,\\ref{f-hb}.)  This could be an indication that broad \\hal\\ wings are partially due to fast stellar winds.  %% Our measurements of the FWHM of the core emission component in pcyg sources (Table\\,\\ref{t_pcyg}) implies %% an average velocity of $\\sim$120\\,\\kms\\ %% in the the line core emitting region.  %% Alternative mechanisms for line broadening: %We believe %that a fast ($\\sim$1400\\kms) stellar wind is also very likely being %expelled by the central star of IRAS\\,19306$+$1407 leading to the %\\hal\\ wings in this object.  %Accordingly, we believe that in our sample kinematical broadening of the %line wings due to the presence of pAGB winds with large velocity %gradients is very probable... % %For Raman scattering to be efficient, a dense, compact HI shell around %the star is therefore needed since atomic hydrogen is the main %scatterer.  \\subsection{Incipient mass-loss probed by efA profiles} \\label{dis_efA}  Structured \\hal\\ absorption profiles partially filled in with emission are common in F- and G-type supergiants and are considered as a possible evidence for mass-loss \\citep[e.g.][]{tam93,are85,sow90}. Although most efA sources in our sample are surrounded by extended nebulosities (see Table\\,\\ref{t_parms} and references therein), the observed structured \\hal\\ profile, in particular, the weak emission that fills in the stellar absorption, arises in the bright nebular nucleus, i.e$.$\\ in a compact region very close to the central star, and not in the extended nebulae. This compact nuclear region is directly seen but is spatially unresolved in our ESI spectra for IRAS\\,04296$+$3424, IRAS\\,17440$-$3310, IRAS\\,18167$-$1209, IRAS\\,19114$+$0002, IRAS\\,19475$+$3119, and IRAS\\,20136$+$1309. %The limited spatial resolution of the %present observations does not allow us to derive any conclusions on %the structure of the compact %\\hal-emitting region resulted (most likely) from on-going mass-loss, %i.e$.$\\ whether it is a shell, torus-like %structure, a collimated wind, etc. Among these, IRAS\\,17440$-$3310 is the only efA source for which our ESI spectrum also shows weak \\hal\\ emission arising in the extended lobes (we note, however, that the spectrum of this object is yet dominated by the bright nucleus). In this case, therefore, part of the emission filling in the absorption profile in Fig.\\,\\ref{f-gaussf3c} may be locally produced in the lobes. Some contribution by nebular \\hal\\ emission to the efA profile cannot be ruled out either for IRAS\\,04296$+$3429, which is known to have a $\\sim$1\\farcs5 extended nebula \\citep{uet00} that is unresolved in our long-slit spectrum. For objects with optically thick equatorial regions, the nuclear region is blocked from direct view. %% and its light is scattered by nebular dust. This is the case of IRAS\\,17317$-$2743, IRAS\\,17441$-$2411 and IRAS\\,20028$+$3910. Our 2D spectra of IRAS\\,17441$-$2411 and IRAS\\,17317$-$2743 show, in fact, continuum emission from two extended lobes separated by a dark equatorial region but there is no \\hal\\ emission locally produced in the lobes. The similar spatial distribution along the lobes of the stellar continuum and the structured efA \\hal\\ profile indicate that both the continuum and the \\hal\\ line originate at the nebular core and are scattered by dust in the lobes. This is most likely the case of IRAS\\,20028$+$3910, for which our ESI spectrum shows unresolved continuum emission arising in the brightest of its two reflection lobes.  %% Halpha profiles in pAGB objects are usually: shell-like, p-cygni %% or inverse p-cygni. %% Shell profiles have several possible explanations: %% a disc around a star, infall of material, or pulsation.  % Except maybe for IRAS 19114, supergiants in our sample show IR excesses % and other properties that are most % consistent with being pAGB objects, i.e$.$ evolved low- or intermediate-mass stars rather than % young massive stars just evolved off the main sequence. % In some cases, the emission component has a P Cygni or inverse % P Cygni profile or a shell-like profile suggesting the presence of % outflows, inflows, pulsation, or discs around the central stars of % these objects. % Asymmetries in the abs. profile are frequent   As shown in Section\\,\\ref{proffit}, the emission filling in the absorption in efA sources is well reproduced by a Gaussian function with an average width of FWHM$\\sim$65\\kms\\ (in deriving this value we have ignored objects for which the width of the emission component is uncertain; see Table\\,\\ref{t_gaussfit}).  After deconvolution with the spectral resolution in our data (37\\kms, \\S\\,\\ref{obs}) and thermal line broadening ($\\sim$15\\kms\\ for a kinetic temperature of 10$^4$\\,K), this value implies moderate speed ($\\sim$50\\kms) motions in the \\hal-emitting region. %The upper limit to the velocity %is because the FWHM of \\hal\\ has not been corrected for line thermal %broadening (the thermal velocity for a kinetic temperature of %10$^4$\\,K is $\\sim$15\\kms). This value is larger than the typical expansion velocity of the slow (5-15\\kms) AGB wind but smaller than the velocity of the fast pAGB outflows observed in pcyg sources (typically $\\ga$100\\kms; see \\S\\,\\ref{dis_pcyg}).  For the majority of the objects with a reliable estimate of the Doppler shift ($\\Delta\\lambda$) for the emission feature (\\S\\,\\ref{proffit}), the emission is displaced bluewards from the narrow absorption core: the blue-shift found ranges from $-$10 to $-$43\\kms\\ (Table\\,\\ref{t_gaussfit}). Only in one case, IRAS\\,20136$+$1309, the emission and the narrow absorption components are centered around the same wavelength within the expected errors in our measurements. %% The shift and W of 20136 are consistent with a larger shell. Among the three objects with the weakest \\hal\\ emission features and, therefore, most uncertain emission line parameters, two show blue-shifted emission and one, IRAS\\,04296$+$3429, red-shifted (by $+$7\\kms) emission. We note, however, that for the latter, the parameters derived from the Gaussian analysis of the \\hal\\ line are somewhat uncertain due to the weakness of the emission component and the fact that some contamination of the efA profile by nebular \\hal\\ emission cannot be ruled out.  %(Weak emission features are also %observed for IRAS\\,17441 and IRAS\\,18167, but in these cases the %blueshift of the emission is quite evident already from the observed %profile.)  We interpret the weak emission partially filling the efA \\hal\\ profile, which arises in the stellar vicinity, as an indication of present-day (i.e$.$\\ pAGB) mass-loss. (The material in the extended nebulosities around most objects resulted predominantly from the large-scale mass-loss process during the AGB -- \\S\\,\\ref{intro}.)  For some of our efA sources, the presence of nebular material near the star is confirmed by detection of a small number of weak non-Hydrogen emission features, e.g.\\ IRAS\\,04296$+$3429, IRAS\\,19114$+$0002, and IRAS\\,20136$+$1309 (and, tentatively, IRAS\\,19475$+$3119; \\S\\,\\ref{res_nohi}). %The strength of %the \\hal\\ emission in efA is much smaller than in pcyg sources, which %probably indicates lower mass-loss rates in the former (but see next %Section). Since the stellar spectral types of the objects with efA \\hal-profiles are F and G (i.e$.$\\ their central stars are relatively cool, with temperatures ranging between 5000-7000\\,K), the emission component is likely to be formed in the de-excitation region behind a shock wave, presumably produced by wind interaction, rather than in a photoionized \\ion{H}{2} region. The fact that the emission is systematically doppler shifted bluewards with respect to the narrow absorption core can be explained if there is significant occultation of the receding part of the outflowing stellar wind.  Note that for stellar occultation effects to be significant the size of the \\hal-emitting region has to be comparable to the radius of the star. %(region could be a spherical shell, %a disk or a collimated wind...)  Alternatively, the weak \\hal\\ emission observed in efA sources could arise, in principle, in a long-lived reservoir of gas near the star, such as a rotating disk, rather than in a stellar wind. We believe, however, that the systematic blue-shift of the emission would be more difficult to explain under this scenario since a long-lived circumstellar structure would necessarily have to be rotating (to persist throughout a substantial amount of time) and, in this case, there is no obvious reason why the emission from the receding edge of such structure would be preferentially depressed/reduced.   Finally, although, both the narrow and broad absorption components of the efA and pA \\hal\\ profiles have most likely a photospheric origin, they are expected to arise in different layers of the stellar atmosphere, the broad feature probably arising in deeper, denser regions. %, whereas the narrow core could arise %in a shell at the highest part of the %atmosphere. Lebre and Gillet 1992. The observed Doppler shifts between both absorption components, therefore, could result from the complex kinematics across the stellar photosphere. Complex, large amplitude atmospheric motions, involving e.g.\\ the presence of shock waves propagating across the atmosphere, are derived from asymmetrical absorption lines and radial velocity measurements in supergiants, including those that are pPN candidates \\citep[e.g.][]{leb91,leb92}. %% (e.g$.$ the RV Tau star R Sct; Lebre \\& Gillet 1991, 1992). Such complex motions are also evidenced by variable Doppler shifts of the Stark photospheric absorption relative to \\vsys\\ (see references above). % In particular, the Stark photospheric absorption is largely blue-shifted % w.r.t. the narrow core for W Vir (Lebre...) %% Halpha is produced over a large scale-height. %(broad=Stark, %narrow=low surface-gravity) ==>  indicate  %Unfortunately, in the absence of spectroscopic observations at %different epochs, it is not possible to reach definitive conclusions %on the origin of the emission and absorption line shifts for our efA %sources. For example, time variable line shifts between the emission %and absorption components are expected if there are periodic shock %waves in the envelopes of these objects. Finally, modeling of the %observed composite profiles with non-LTE radiative-transfer models are %needed for studying the spatio-kinematic structure of the extended %atmosphere and inner circumstellar layer relevant to the formation of %\\hal\\ absorption and emission features in these complex objects.    \\label{sum}   We present echelle long-slit optical spectra of a sample of evolved intermediate-mass stars in different evolutionary stages: 5 AGB stars, 17 pPNs, and 6 yPNs. Our sample also includes the object IRAS\\,19114$+$0002, which has a controversial classification as a pPN or a yellow hypergiant, %but that, in any case, follows a %mass-loss driven evolutionary path similar to that of pAGB stars, and one YSO, IRAS\\,05506$+$2414, which was serendipitously discovered in our multi-wavelength survey of pPNs. (The spectrum of the latter is presented here for completeness but it is not discussed except for spectral typing of its central star.)  We have analyzed extracted 1D spectra of our targets with special focus on the characteristics of the \\hal\\ line profile arising in the vicinity of the central source, i.e.\\,the nebular nucleus. In this section, we summarize the main results obtained from this work:  \\begin{itemize}  \\item[-] Fifteen objects in our sample show relatively intense \\hal\\ emission, whereas eleven targets show \\hal\\ mainly in absorption. We have also found three sources (the AGB stars IRC$+$10011, V656 Cas, and IK\\,Tau) %IRAS\\,01037$+$1219, IRAS\\,02316$+$6415, and IRAS\\,03507$+$1115) with neither \\hal\\ emission or absorption, and one object (the pPN IRAS\\,19292$+$1806) in which the low S/N of the spectrum prevents us from determining whether \\hal\\ emission or absorption is present.  \\item[-] Based on the shape of the \\hal\\ line, we have defined four main types of sources. Among the \\hal\\ emitters, those with a symmetric \\hal\\ profile are referred to as pure emission sources (pE), whereas objects with asymmetric P Cygni like profiles are denoted as pcyg. Objects with a pure absorption \\hal\\ profile are named pure absorption targets (pA). Objects with an absorption profile partially filled with weak emission are referred to as emission filled absorption sources (efA).  %Among the objects with \\hal\\ absorption, two show a %pure absorption \\hal\\ profile (pA) and the rest show an absorption %profile partially filled with weak emission -- the latter are referred %to as emission filled absorption sources (efA).  \\item[-] The presence of \\hal\\ emission from the compact nebular core displaying either a pE, pcyg, or an efA profile is interpreted as an indication of on-going (i.e.\\,pAGB) mass-loss most likely in the form of a stellar wind. The observed \\hal\\ profiles have been parameterized and analyzed to derive information on the current stellar wind and other processes that may be affecting the observed line shape. Among the objects that have already left the AGB (i.e.\\,pPNs and yPNs) there are only two pA sources. Except for these two objects, the rest show evidence of current stellar winds, which supports the idea that pAGB winds are generally present in pPNs and yPNs and are most likely responsible for the nebular shaping.  \\item[-] We interpret the peculiar P Cygni like profiles observed in some of our objects in a similar manner as we did for He3-1475 \\citep{san01}. In this scenario the emission and absorption form in two distinct nebular components. The broad line emission, characterized by a given $intrinsic$ profile, arises from a compact central source. This line emission (and stellar continuum) is scattered by dust in the walls of the nebular lobes.  The blue-shifted absorption is due to neutral or partially ionized outflowing gas inside the lobes absorbing the scattered photons along the line of sight. This results in a P Cygni like profile where the blue-shifted absorption is produced against the intrinsic emission profile.  \\item[-] For pE and pcyg sources the FWHM of the intense \\hal\\ core emission indicates gas motions with velocities in the range [50:200]\\kms. The mean velocity of the bulk of the material producing the blue-shifted absorption in pcyg targets varies between $V$$\\sim$50 and 500\\kms, however, larger outflow terminal velocities of up to $v_\\infty \\approx$\\,1000\\kms\\ are observed. It is possible that a radial and or latitudinal velocity gradient exists in the pAGB outflows and/or that distinct wind components are present.  \\item[-] Broad \\hal\\ emission wings, with widths of up to $\\pm$2000\\kms, are observed in most pE and pcyg sources. The yet unclear origin of the very extended wings is investigated following a similar analysis to that performed by \\cite{arr03}.  Unlike these authors, we cannot conclude that in our sample Raman scattering is the main mechanism for line broadening. The presence of fast stellar winds, which is confirmed by P Cygni profiles in more than one third of our targets, is another plausible mechanism that could contribute to the broad \\hal\\ wings observed. These winds, however, would have to reach velocities larger than those derived from the P Cygni absorption features ($v_\\infty$) to explain the broad \\hal\\ wings.  %This is indeed %the most likely mechanism for IRAS\\,20462$+$3416 and %IRAS\\,19306$+$1407 for which the \\hal\\ wing profile do not follow %the $\\propto \\lambda^{-2}$ law typical of Raman-scattered features.  %In fact, the shape of the wings in the \\hal\\ profile %of IRAS\\,20462$+$3416 and IRAS\\,19306$+$1407 cannot be reproduced by a %$\\lambda^{-2}$ law, expected for Raman scattering. Although the wing %profile can be reproduced by a $\\lambda^{-2}$ law in the rest of the %cases, the presence of other Raman-scattered features, which is needed %to confirm this broadening mechanism, are identified only in four %objects: IRAS\\,08005$-$2356, M1$-$92, He3$-$1475 and %IRAS\\,22036$+$5306. The FWZIs of such features ($\\ls$150\\kms) are, %however, significantly smaller than those measured in Arrieta's sample %and smaller than the \\hal\\ wings themselves. In one of these four objects, %He3$-$1475, the width of the \\ion{Si}{3}]$\\lambda$1892\\AA\\ observed with the IUE %and expected to arise in the same region as the Ly$\\beta$ photons %is too small to  \\item[-] The \\hal\\ emission filling in the absorption profile of efA indicates moderate speed ($\\sim$50\\kms) motions in the nuclear \\hal-emitting region. For most efA sources, the emission ``hump'' is displaced bluewards from the narrow stellar absorption feature which can be explained by occultation of the receding part of an outflowing stellar wind by the central star.  \\item[-] We have found differences in the \\hal\\ profiles of some of our targets with respect to earlier observations, namely, IRAS\\,19114$+$0002, IRAS\\,19475$+$3119, IRAS\\,20462$+$3416, and the AGB stars IRC+10216, CIT\\,6, and IK\\,Tau. These variations most likely represent real changes on the physical properties of the nuclear \\hal-emitting region (e.g$.$ density, excitation, ionization fraction, geometry and size) presumably induced by the evolution of the central star and/or its current stellar wind. %% Need this??  \\item[-] We briefly discuss other Hydrogen and non-Hydrogen lines observed towards our targets.  We note several prominent C$_2$ Swan bands in the AGB stars IRC$+$10216 and CIT\\,6 and the pPN IRAS\\,04296$+$3429. Some of these bands have been reported for the first time in this work. It is also worth mentioning the detection of emission by the [\\ion{Ca}{2}] $\\lambda$7291,7324\\AA\\ doublet in the absorption-line dominated spectrum of the YHG IRAS\\,19114$+$0002. Observation of these lines, also present in the spectra of IRAS\\,17516$-$2525, M\\,1-92, Hen\\,3-1475, IRAS\\,22036$+$5306, IRAS\\,08005$-$2356, and (tentatively) IRAS\\,19520$+$2759, is consistent with dust grain destruction by moderate-velocity shocks in the stellar wind. %% Need this??   \\item[-] We have estimated the spectral type of the central stars of the objects in our sample by comparing their normalized spectra with those of template stars from published stellar libraries to investigate the correlation of this fundamental stellar parameter with the type of \\hal\\ profile. In a number of objects with F- and G-type central stars, we found discrepancies (typically of one type) with respect to previous spectral type assignations obtained from low spectral resolution studies. Such differences may reflect the limitations of low resolution spectroscopy for accurate spectral typing, although a real time evolution of the stellar effective temperature cannot be ruled out.  %% The latter could be induced by intense, episodic mass-loss %% events during the pAGB phase %(Appendix\\,\\ref{starevol}). %%   \\item[-] The stellar luminosity has been estimated from the luminosity-denpendent \\ion{O}{1}\\,7773\\AA\\ infrared triplet, which is observed in absorption in the spectrum of some of our targets (13 out of 29). The obtained values, $L_{\\rm bol}$$\\approx$5$\\times$10$^3$-10$^5$\\ls, are consistent with pAGB stars with initial masses in the range 1-8\\msun, with a significant fraction (10 out of 13) of objects with masses $>$3\\,\\msun. Such massive pAGB objects represent $\\sim$30\\% of our whole sample. The rest of the sources, which is also the majority ($\\sim$70\\%), may well have low-mass ($\\la$3\\msun) progenitors -- this may be part of the reason why the \\ion{O}{1} triplet is not observed in absorption. For the highest masses/luminosities derived in our sample, the theory of pAGB stellar evolution predicts rapid changes of the effective stellar temperature that should be observable in a time-scale of $\\approx$10-100\\,years.  \\item[-] We investigated correlations between the type of \\hal\\ profile and some stellar and envelope parameters. The shape of the \\hal\\ line emission is correlated with the stellar spectral type as well as the the NIR and IRAS colors. %, which are indicative of differences in the AGB and %  most recent pAGB mass loss history. No correlation has been found with the chemistry, galactic latitude, or nebular morphology.   \\item[-] All sources in which \\hal\\ is seen mainly in absorption (i.e., pA and efA) have F-G type central stars, whereas sources with intense \\hal\\ emission (i.e., pE and pcyg) span a larger range of spectral types from O to G, with a relative maximum around B, and also including very late C types. The measured equivalent widths of the \\hal\\ emission in objects with O- and B-type stars are consistent with UV stellar radiation being the main ionizing agent. \\hal\\ emitters with cooler central stars lack enough ionizing radiation, therefore, the emission component in these cases is likely to be formed in the de-excitation region behind a shock wave, presumably produced by wind interaction. %shocks produced at the close stellar %  surroundings are likely the main ionization mechanism.   \\item[-] Pcyg and pE sources are found to exhibit a larger $J-K$ color excess than pA and efA objects. Moreover, while the NIR colors of efA+pA sources are consistent with the major contribution to the emission in the NIR bands being the reddened stellar photosphere, the position of pcyg+pE sources in the ($J-H$)-($H-K$) color-color diagram highlights the presence of warm dust. This component of warm dust probably results from substantial present-day mass-loss evidenced by the strong \\hal\\ emission of pE and pcyg targets.  \\item[-] Intense \\hal\\ emitters (i.e.\\,pcyg+pE) and objects with \\hal\\ mainly in absorption (i.e.\\, pA+efA) also segregate in the IRAS color-color diagram in a way that the former have dust grains with a larger range of temperatures. Such a different temperature range may result from: 1) a different distribution of the grain density in the AGB envelope; 2) a different spatial distribution of grain sizes; 3) a different grain composition; or any combination of these situations.   \\item[-] The differences found between pE+pcyg and pA+efA sources described above indicate dissimilarities in their pAGB and AGB mass-loss histories. The intense \\hal\\ emission and NIR color excess of pE and pcyg sources may indicate larger pAGB mass-loss rates compared with those in efA and pA targets.  The lack of a mass/luminosity segregation of the two profile sources suggests that pAGB mass-loss is not very dependent on intrinsic stellar properties but may be dictated by extrinsic factors like, for example, the presence of a binary companion. % The velocity of the stellar winds is found to be larger % in \\hal\\ emitters than in objects with \\hal\\ absorption. The faster winds and earlier spectral types of the central stars of pcyg and pE sources suggest that these are older and/or have a larger speed traversing the pAGB evolutionary path than pA and efA.  \\item[-] We have not found any indications of the extended optical nebulosities (observed in most objects) being older for pcyg+pE sources than for efA+pA targets. %since there are no obvious differences in % the nebular morphology. Such nebulosities show many signs that the jet-sculpting process of the AGB CSE is currently active or has been active in the past in all cases. For the only two pA targets in our sample, IRAS\\,19477$+$2401 and IRAS\\,17150$-$3224, the winds that shaped their AGB CSEs to its current aspherical morphology are not seen at present. The on-going pAGB winds probed by the nuclear \\hal\\ emission in pE+pcyg and efA objects, which represent the vast majority of our sample, may not be the same that carved and accelerated the much more extended (probably older) optical lobes, however, shaping of the innermost layers of the AGB envelope (by interaction with the present-day pAGB winds) is probably still at work.  \\item[-] In principle, pAGB winds may have been ejected in a continuous or episodic manner since the shaping of the AGB envelope began. The fact that pcyg and efA objects with central stars of similar late types (F and G) exhibit quite different \\hal\\ emission profiles may suggest episodic pAGB ejections: pcyg would represent targets that are caught at times of intensive jet activity leading to energetic shocks and intense \\hal\\ emission, whereas efA could be currently in a status of ``mild'' wind interaction (resulted from smaller velocities and/or densities in the pAGB wind?). In pA targets, the pAGB winds that presumably shape their aspherical AGB envelopes may have temporarily or permanently ended.   \\end{itemize}  %% The \\notetoeditor{TEXT} command allows the author to communicate %% information to the copy editor.  This information will appear as a %% footnote on the printed copy for the manuscript style file.  Nothing will %% appear on the printed copy if the preprint or %% preprint2 style files are used.   %% Included in this acknowledgments section are examples of the %% AASTeX hypertext markup commands. Use \\url without the optional [HREF] %% argument when you want to print the url directly in the text. Otherwise, %% use either \\url or \\anchor, with the HREF as the first argument and the %% text to be printed in the second."
    },
    "0806/0806.4187_arXiv.txt": {
        "abstract": "In many theoretical scenarios it is expected that intermediate-mass black holes (IMBHs, with masses $M\\sim 10^{2-4}~M_\\odot$) reside at the centers of some globular clusters.  However, observational evidence for their existence is limited.  Several previous numerical investigations have focused on the impact of an IMBH on the cluster dynamics or brightness profile.  Here we instead present results from a large set of direct N-body simulations including single and binary stars. These show that there is a potentially more detectable IMBH signature, namely on the variation of the average stellar mass between the center and the half-light radius.  We find that the existence of an IMBH quenches mass segregation and causes the average mass to exhibit only modest radial variation in collisionally relaxed star clusters. This differs from when there is no IMBH. To measure this observationally requires high resolution imaging at the level of that already available from the Hubble Space Telescope (HST) for the cores of a large sample of galactic globular clusters. With a modest additional investment of HST time to acquire fields around the half-light radius, it will be possible to identify the best candidate clusters to harbor an IMBH. This test can be applied only to globulars with a half-light relaxation time $\\lesssim 1$~Gyr, which is required to guarantee efficient energy equipartition due to two-body relaxation. ",
        "introduction": "Theoretical work has suggested that some globular clusters may harbor intermediate-mass black holes (IMBHs; $M\\sim 10^{2-4}~M_\\odot$) in their centers (e.g., \\citealt{por02}). If this is indeed the case, there are significant consequences for ultra-luminous X-ray sources,  gravitational wave emission from dense star clusters, and the dynamics of globular clusters (GCs) in general (see \\citealt{vdm04,mc04} for an overview).  Definitive evidence for IMBHs has, however, been elusive.  For example, \\citet{grh02, grh05} argued for an IMBH in G1 based on the analysis of HST line-of-sight velocity data and Keck spectra, but an alternative analysis by \\citet{bau03a} points out that acceptable dynamic models without a large central object also fit the observations. \\citet{ger03a,ger03b} argued that the kinematics of M15 seem to slightly favor the presence of an IMBH, but for this cluster alternative interpretations exist \\citep{bau03b,dul03}.  More recently, the observed line-of-sight kinematics of Omega Cen have also been used to argue for the presence of an IMBH \\citep{ngb08}.  A more secure identification of an IMBH in a GC can, in principle, be provided by also measuring the proper motion of central stars in order to reconstruct their orbits and thus firmly establish if a central massive point object is present. Several HST-GO programs based on this idea have been approved in past cycles (e.g. GO10474, PI Drukier; GO10401 \\& GO10841 PI Chandar; GTO/ACS10335 PI Ford), but to date they have not yielded any indisputable detections. The limitation for such studies is the need to carry out multi-year observations, thus progress is slow. To maximize the chances of success it is thus of primary importance to focus the observations on the candidates most likely to harbor an IMBH.  Candidate selection is possible if one focuses on the indirect influence of the IMBH on the dynamics of its host.  Direct N-body simulations by \\citet{bau04} and \\citet{trentiea07b} found that the presence of an IMBH acts as a central energy source that is able to prevent gravothermal collapse and thus maintain a sizable core to half-mass radius ratio throughout the entire life of the GC. The existence of such a large ($\\gtrsim 0.1$) core to half-mass radius ratio in a collisionally relaxed cluster might be due to the presence of an IMBH (see also \\citealt{heg07}). However, the picture becomes more complicated when this signature is transferred from the ideal world of N-body simulations, where a complete knowledge of the system is available, to real observations, where essentially only main sequence and red giant branch stars define the light profile of the system. In fact, an analysis by \\citet{hur07} cautioned that the difference between mass and light distributions can lead to a large observed core to half-light radius ratio for GCs with single stars and binaries only.  Here, we continue the search for indirect IMBH fingerprints by focusing on the consequences of the presence of an IMBH on mass segregation.  Through direct N-body simulations we show that the presence of a large (of order 1\\% of the total mass), central mass significantly inhibits the process of mass segregation, even among only visible main sequence stars and giants.  To the best of our knowledge this effect was first briefly mentioned in \\citet{bau04}, but left without further quantitative analysis. Quenching of mass segregation is present in all of our simulations with an IMBH, independent of the initial conditions of the cluster, including variations in initial mass function, density profile, strength of the galactic tidal interaction, number of particles and initial binary fraction.  We find that a differential measurement of the average mass between the center and the half-light radius is effective in separating star clusters with and without an IMBH, provided that the stellar system is at least 5 initial half-mass relaxation times old.  This measure is observationally feasible with current data (for example, see \\citealt{dem07} and references therein), and can lead to the selection of a promising set of IMBH host candidates.  A direct observational application of this approach is left to a companion paper. Here, we focus instead on building the theoretical framework for such analysis. In \\S~2 we describe our numerical simulations, in \\S~3 we discuss our results, and in \\S~4 we present our conclusions. ",
        "conclusions": "\\label{sec:conclusion}  We have carried out a large set of direct N-body simulations of star clusters with and without an IMBH including a realistic mass spectrum and primordial binaries.  While previous research has focused its attention mainly on the effects of an IMBH on the surface brightness and velocity dispersion profiles of the clusters --- signatures that are difficult to observe --- we searched instead for a different fingerprint of the presence of an IMBH. The existence of a massive, central object quenches mass segregation and this effect manifests itself in collisionally relaxed clusters through decreased radial variation in the average mass of main sequence stars. This effect does not depend on the mass of the black hole as long as it is dominant over the typical mass of a star, nor on the details of the initial configuration of the system such as initial mass function, density profile and tidal field strength. The amount of mass segregation is only weakly dependent on the binary fraction of the cluster. This result allows us to use the amount of mass segregation to separate collisionally relaxed clusters with and without an IMBH without the need of additional modeling assumptions.  A critical requirement for the proposed signature is that the system be well-relaxed, so that it has already attained equilibrium with respect to mass segregation. From our simulations it turns out that this takes about $5 t_{rh}(0)$. However we can only observe the current half-mass relaxation time and this might be shorter than its initial value if the system has lost a large fraction of its original mass. To compare our simulations to observations, we must thus conservatively restrict ourselves to GCs that: \\begin{enumerate}  \\item Are not too influenced by the galactic tidal field (that is, with a tidal to half-light radius $r_t/r_{hl} \\gtrsim 10$, which corresponds to tidal fields weaker than the weakest field in our simulations).  \\item Have half-mass (3D) relaxation times below $\\approx 1.5$~Gyr, i.e. an age above $8 t_{rh}$. This leaves room for a mass loss of about $50\\%$ of the initial mass while still giving an integrated age of about $5 t_{rh}$. In terms of observable quantities this translates into a half-light relaxation time below $\\approx 1$~Gyr.  \\end{enumerate} Based on the \\citet{har96} catalog, 31 galactic GCs satisfy these stringent requirements in terms of relaxation time and $r_t/r_{hl}$. The proposed diagnostic could probably be applied to more clusters after properly evaluating a dynamical model for their configuration and eventually accepting some uncertainty in the selection of likely candidates to harbor an IMBH.  Thanks to the HST treasury survey of galactic GCs, data exist for the cores of many clusters that explore deep enough to see main sequence stars down to around $0.2 M_{\\odot}$.  Along the same lines, \\citet{dem07}, among others, have also acquired images of clusters around the half-light radius, in order to calculate the global mass function of the system. The existing data from \\citet{dem07} are sufficient to apply this diagnostic to a few actual clusters and the results from such a comparison will be presented in a companion paper.  In closing, we stress again that while the amount of mass segregation has been proven here to be a viable indicator for the presence of an IMBH in simulated star clusters, we cannot use this method alone to claim the detection of an IMBH. However, by combining the measure of mass segregation with all other constraints from the velocity dispersion and surface brightness profiles, we can select the clusters that seem most likely to harbor an IMBH, while at the same time excluding some others from further scrutiny. Once we have identified those clusters that are most promising, future observations, such as proper motion studies, can focus their efforts to secure a robust detection."
    },
    "0806/0806.3072_arXiv.txt": {
        "abstract": "We report the discovery of a $\\sim$500 kpc HI extension southwest of the Virgo Cluster HI-rich pair NGC 4532/DDO 137, detected as part of the Arecibo Legacy Fast ALFA (ALFALFA) Survey.  The feature is the longest and most massive HI tail structure so far found in the Virgo Cluster and, at 1.8 Mpc from M87, the most distant from the main concentration of the intracluster medium.  The structure is spatially and spectrally separated into two ridges and is defined by diffuse emission and discrete clumps of mass 2.5 - 6.8 x 10$^7$ \\msol. All emission is blue-shifted with respect to the NGC 4532/DDO 137 pair emission.  Including diffuse emission, the structure has a total mass of up to 7 x 10$^8$ \\msol, equivalent to $\\sim$10\\% of the system's HI mass. Optical $R$-band imaging finds no counterparts to a level of 26.5 mag arcsec$^{-2}$.  The characteristics of the structure appear most consistent with a tidal origin. ",
        "introduction": "Galaxies in clusters experience a variety of environmental interactions that affect their evolution (see Boselli \\& Gavazzi 2006 for a review). Intracluster medium (ICM) interactions such as ram pressure stripping (Gunn \\& Gott 1972) and starvation (Larson, Tinsley, \\& Caldwell 1980) prematurely remove gaseous reservoirs. Tidal interactions, including nearby, slower encounters in cluster and group substructures (Toomre \\& Toomre 1972), nearby high velocity encounters (Duc \\& Bournaud 2008), and galaxy harassment (Moore \\etal~ 1996, 1998; Bekki \\etal~2005), rearrange stellar and gaseous contents. These interactions potentially explain the observed gas deficiencies (Giovanelli \\& Haynes 1983, Cayatte \\etal~ 1990), reduced star formation (Kennicutt 1983; Koopmann \\& Kenney 2004), and morphology-density relation (Dressler 1980) in clusters.  Environmental interactions can produce low surface brightness stellar and gas tails. Chung \\etal~(2007) find one-sided HI tails in seven Virgo spirals, attributing them to the influence of the ICM.  Oosterloo \\& van Gorkom (2005) report a 110 x 25 kpc plume of HI gas extending away from the HI-deficient Virgo spiral NGC 4388, also attributing the feature to ICM interaction.  Mihos \\etal~(2005) find stellar streams associated with several presumably tidal events near the Virgo Cluster core. Similar features have been found in Coma and Centaurus (Gregg \\& West 2004). The characteristics of tail features are related to the details of the environmental interaction that produced them, helping to determine the relative importance of different interactions.  Tidal tails may allow the formation of tidal dwarf galaxies, contributing to the dwarf galaxy population.  The Arecibo Legacy Fast ALFA (ALFALFA) Survey, a sensitive blind survey of the Arecibo sky (Giovanelli \\etal~2005), is providing a complete and unbiased view of HI content and structures in the entire Virgo cluster region. The survey has revealed the presence of several HI clouds without optical components (Kent \\etal~2007) and a 250 kpc extended tidal arc emerging from the Sc galaxy NGC 4254. This structure, encompassing the Virgo HI21 cloud (Minchin \\etal~ 2007), is likely due to a high velocity close galaxy encounter (Haynes \\etal~2007; Duc \\& Bournaud 2008).  In this Letter, we report the detection by ALFALFA of an even larger tidal feature associated with the Virgo Cluster pair NGC 4532/DDO 137. This pair of SmIII/SmIV (Binggeli, Sandage, \\& Tammann 1985) galaxies is located in the Virgo B Cloud (Binggeli, Popescu, \\& Tammann 1993), 6$^{\\circ}$ south of the Virgo center.  NGC 4532 is the brightest Sm cataloged in the Virgo Cluster Catalog (Binggeli, Sandage \\& Tammann 1985; hereafter VCC) and has a high star formation rate (Koopmann \\& Kenney 2004) and an asymmetric stellar morphology. The galaxies share a common HI envelope extended over 150 kpc (Hoffman \\etal~1992, 1993).  Hoffman \\etal~(1999) found that the HI envelope contains three additional discrete HI clumps that have no optical counterparts as well as a significant diffuse HI component, some of which appeared as a tail-like extension to the southwest. We show that there is indeed an extended HI structure, stretching $\\sim$ 500 kpc beyond the pair.  Section~\\ref{obs} describes ALFALFA and optical followup observations of the extended HI structure and Section~\\ref{disc} addresses possible formation mechanisms. We assume a Virgo Cluster distance of 16.7 Mpc (e.g., Mei \\etal~ 2007) throughout. ",
        "conclusions": "\\label{disc}  We have discovered an extremely long (500 kpc) HI stream of low $N_{\\rm HI}$, dotted by higher density clumps appearing as isolated HI clouds, most with no optical counterpart, apparently associated with the galaxy pair NGC 4532/DDO 137. The characteristics of the system are reminiscent of those reported by Kent et al~(2007), Haynes et al~(2007), and Tripp (2007), also found mainly in the periphery of the Virgo cluster.  The feature is to our knowledge the most extreme HI tail structure found in a cluster, both in terms of its length and its position in the cluster.  It is located at $\\ge$1.6 times the distance from the Virgo Cluster center as other galaxies with tail features. The projected extent is a factor $\\ge$ 14 times as large as the one-sided HI tails discovered in the VLA Imaging of Virgo Galaxies Survey (Chung \\etal~2007). It is several times larger than the stellar tails found by Mihos \\etal~(2005) and Gregg \\& West (2004) and the HI tail described by Osterloo \\& Van Gorkom (2005), and twice as large as the feature near NGC 4254 (Haynes \\etal~2007). The total HI mass is a factor of 1.5-16 times larger than that reported for other HI tails, although the fractional mass of $\\sim$10\\% of the presumed host system is similar.  The NGC 4532/DDO 137 pair is located at a projected distance of 1.8 Mpc south of M87 and 0.5 Mpc southwest of M49 (NGC 4472).  Binggeli \\etal~(1993) identify the galaxies as members of the Virgo B cloud, which is centered near M49 and lies at about the same distance as the more massive Virgo A Cloud centered near M87  (Binggeli \\etal~1993; Mei \\etal~2007).  The subclump contains about $\\sim$1\\% of the total ICM mass in the cluster (Schindler \\etal~1999) and has a spiral-rich population with a mean line-of-sight velocity of $\\sim$1040 km s$^{-1}$ and a velocity dispersion of $\\sim$500 km s$^{-1}$ (Binggeli \\etal~1993).  We estimate the ram pressure force due to the ICM at the position of NGC 4532/DDO 137 to be 2 - 25 times smaller than the gravitational restoring force on their ISM (following a similar approach as Chung \\etal~2007 with dynamical properties of the galaxies given by Hoffman\\etal~1999).  The extended HI envelope is presumably less tightly bound and would be more susceptible to stripping. A challenge for an ICM interpretation is the length of the feature: it is an order of magnitude longer than other observed and simulated (e.g., Vollmer \\etal~2001; Roediger \\& Br\\\"uggen 2008) features.  In addition, it extends south of the pair, implying a trajectory that did not take the pair through the densest and hottest part of the ICM, as traced by ROSAT (B\\\"{o}hringer \\etal ~1994) and ASCA (Shibata \\etal~2001).  Tidal interactions naturally produce long, gas-rich tails (e.g., Toomre \\& Toomre 1972).  NGC 4532 and DDO 137 appear to be a bound pair and could be interacting.  NGC 4532 shows other symptoms of tidal interaction: it is optically asymmetric and has a high star formation rate (Koopmann \\etal~2004) and disturbed velocity field (Rubin \\etal 1999; Chemin \\etal~2005; Hoffman \\etal~1999).  The HI extension described here displays a highly ordered velocity field. However low velocity tidal interactions between galaxies tend to produce symmetric tails of gas and stars (e.g., Toomre \\& Toomre 1972; Hibbard \\etal~2001). In this case no stellar tail has yet been found and the HI extension and excess HI envelope gas not identified with the galaxies (Hoffman \\etal~1999) display strong kinematic and spatial asymmetries.  These peculiarities could be consistent with a higher velocity encounter with another massive galaxy.  Models of high velocity ($\\sim$ 1000 km s$^{-1}$) close galaxy encounters (Duc \\& Bournaud 2008) share some similarities to low velocity encounters, e.g., the length of the tail and the formation of dense clumps along tails, but produce lower mass, asymmetric, gas-dominated tails.  Duc \\& Bournaud (2008) are able to reproduce the 250 kpc long HI tail extending northward from NGC 4254 (Haynes \\etal~2007) via an encounter 750 Myr ago at a speed of 1100 km s$^{-1}$ with a galaxy 50\\% more massive. There are $\\sim$10 massive (M$_B <$ -18.1) galaxies within 1.5$^{\\circ}$ (440 kpc) of NGC 4532/DDO 137 and the HI extension, including M49 (NGC 4472) and 5 other galaxies identified with Virgo B (Binggeli \\etal~1993).  NGC 4532 and DDO 137 have line-of-sight velocities of $\\sim$ 2000 km s$^{-1}$ (2$\\sigma \\sim$ 1000 km s$^{-1}$ greater than Virgo B mean) so that a high speed encounter with a B member is possible.  As argued by Duc \\& Bournaud (2008), the perturber may be further away; a galaxy moving at 1000 km s$^{-1}$ can travel a projected distance of $\\sim$ 1 Mpc in 1 Gyr.  We note that ALFALFA observations, to date complete to +4$^{\\circ}$00$'$, show no other extended HI features associated with other galaxies in the vicinity.  Based upon the available models, we suggest that the structure associated with NGC 4532/DDO 137 is most consistent with a tidal interaction, possibly a high velocity encounter.  Determining the exact nature of these very long HI tails and the extended HI envelope will require detailed simulation of the system in the entire Virgo Cluster environment, an exercise outside the scope of this paper."
    },
    "0806/0806.2148_arXiv.txt": {
        "abstract": "The electron mean free path in the intracluster medium (ICM) of galaxy clusters is much larger than the gyroradius; thus, heat is transported anisotropically along magnetic field lines.  We show that the intracluster medium is unstable to the magnetothermal instability (MTI) using MHD simulations with anisotropic thermal conduction.  As a result of the MTI, we find that the temperature profile of the ICM can be substantially modified on timescales of several billion years while the magnetic field is amplified by dynamo action to more than fifty times the original energy.  We also show that the instability drives field lines to become preferentially radial leading to conduction that is a highly efficient fraction of the Spitzer conductivity.  As such, we present the first self-consistent calculation of the effective thermal conductivity in the ICM. ",
        "introduction": "\\label{sec:intro} Clusters of galaxies are the most massive gravitationally-bound objects in the universe, yet a mere 3\\% of their mass is in the form of stars.  The majority of their mass (84\\%) is in dark matter with the remaining 13\\% consisting of a hot, low density, magnetized plasma called the intracluster medium (ICM). Due to their enormous masses, surveys of clusters provide key tests of cosmology.  With the Chandra X-ray Observatory, we are capable of measuring the spatially resolved X-ray luminosity ($10^{43}$--$10^{46}$ erg/s) emitted by the ICM to determine its density as a function of radius.  In order to fully understand the fundamental astrophysics of the largest bound objects in the universe, we must develop an accurate quantitative picture of the heating and cooling mechanisms as well as the dynamical processes that govern the ICM.   Thus, we are motivated by the use of clusters as a cosmological tool to study the physics of the ICM much more rigorously.  The intracluster medium is heated by gravitational infall such that typical temperatures are 1--15 keV and typical densities are in the range of $10^{-3}$--$10^{-2} \\;\\mbox{cm}^{-3}$ \\citep{pf06}.  Since the ICM is fully ionized and radiative forces are negligible, we can treat radiation with optically thin cooling and neglect electrical resistivity.  Estimates for the magnetic field strength in clusters range from roughly 1--10~$\\mu$G at the center and 0.1--1.0~$\\mu$G at a radius of 1 Mpc, values that correspond to a plasma beta, $\\beta = 8\\pi P/B^2\\approx 200$--$2000$, a dynamically weak magnetic field \\citep{ct02}.  In this dilute plasma, however, the mean free path for electron collisions can be twelve orders of magnitude larger than the gyroradius \\citep{nm01}.  In this circumstance, the equations of MHD must be supplemented with anisotropic terms that include the near free-streaming motions of particles along magnetic field lines \\citep{brag65}.  The parallel thermal conductivity of the electrons is larger than that of the ions by a factor proportional to $(m_i/m_e)^{1/2}$, whereas the parallel viscosity of the ions is larger than that of the electrons by the same factor.  \\citet{bal00} has predicted that gravitationally bound plasmas in this regime are susceptible to the magnetothermal instability (MTI).  This convective instability is driven by anisotropic thermal conduction along field lines with major dynamical consequences \\citep{ps05,ps07a,ps07b,cd06}.  Thus, we are no longer justified in considering the weak magnetic field in the ICM as a purely passive component in a fluid governed solely by hydrodynamics.  There are two long-standing puzzles about the ICM.  First, the radial temperature profiles show a slight decline in temperature with radius \\citep{pratt07} that is smaller than that expected from structure formation calculations, e.g. \\citep{lok02}.  Second, the magnetic field has been significantly amplified in the cluster from its primordial value of $< 1$ nG.  The ICM is unstable to the magnetothermal instability on scales of tens of kpc and larger, as calculated from the current observed magnetic field strengths. In three-dimensional simulations, we have shown that the MTI is capable of generating convective motions and a magnetic dynamo while efficiently transporting heat.  In addition, the MTI can make the temperature profile more isothermal as it exhausts its source of free energy.  These similarities provide a strong motivation for studying the ICM with this type of simulation.  Hydrostatic equilibrium dictates that \\begin{equation} M(r) = -\\frac{k_B r^2}{G\\mu m_p}\\left(T_e(r)\\frac{\\dif n_e(r)}{\\dif r} + n_e(r)\\frac{\\dif T_e(r)}{\\dif r}\\right), \\label{eqn:clust:m-cluster} \\end{equation} where $\\mu$ is the mean molecular weight \\citep{laroque06}. When an X-ray observation is made, the density profile (first term of eqn. [\\ref{eqn:clust:m-cluster}]) can be constrained fairly well due to the property that $L_x \\propto n_e^2$ for free-free emission, the dominant cooling mechanism. Often, the gas density profile is fit with an isothermal-$\\beta$ model \\citep{cff76}. On the other hand, the temperature profile (second term of eqn. [\\ref{eqn:clust:m-cluster}]) can only be constrained by spectra, which are increasingly difficult to measure at larger redshifts.    For high redshift clusters, the virial masses are often constrained simply by assuming isothermality, in contradiction to local observations. Why do we care about determing mass with such accuracy?  A flux-limited X-ray survey can be used to constrain $\\Omega_0$, the matter parameter, e.g. \\citet{borg99}.  Locally, however, there is a degeneracy between the parameters $\\Omega_0$, the matter fraction of the universe, and $\\sigma_8$, the variance of the power spectrum on $8 h^{-1}$ Mpc scales.  Thus, one must examine the evolution of the cluster mass function with redshift to break this degeneracy.  In \\S\\ref{sec:MTIphys}, we review the physics of the MTI then procede to construct a model of the ICM for numerical simulation in \\S\\ref{sec:model}.  In \\S\\ref{sec:method} we outline our numerical methods.  We then show in \\S\\ref{sec:sim} that the MTI is capable of rearranging the temperature profiles of clusters on cosmologically important timescales, \\textit{i.e.} 0.5--10 billion years.  We also show that a magnetic dynamo operates as a result of the turbulent convection driven by the MTI.  A comparison to runs with purely isotropic conduction at a specified fraction of the Spitzer conductivity show that much less turbulence and no dynamo result.  Finally, we point out that the magnetic field is driven to a highly radial geometry and speculate on possible observational implications before concluding. ",
        "conclusions": "\\label{sec:discuss} We have shown in the previous section that properly accounting for anisotropic thermal conduction in the intracluster medium drives unique physics with potential observational consequences.  The resultant physics is significantly different from the pure MHD case without conduction or the isotropic conduction case.  Of course, without heating or cooling mechanisms included, the simulations presented here are insufficient to make predictions for any particular cluster.  Instead, we briefly outline here two potential observational consequences that will be a focus of future work. \\subsection{Turbulent Velocities}\\label{subsec:discuss:velocities} In this work, we predict rms turbulent velocities of order 3--4\\% the speed of sound and peak velocities of 30--40\\% of the sound speed.  In a relaxed cluster that has not undergone a major merger, this may be the dominant velocity component in the ICM; although, certainly wakes from passing galaxies may also be a significant contribution.  The typical ICM velocities for the MTI are predicted to be approximately an order of magnitude larger than the velocities that would result from pure isotropic conduction alone.  To determine the measurability of these hydrodynamic motions, we must compare the Doppler broadening due to hydrodynamic motions to the simple thermal Doppler broadening.  The most success will come in measuring hydrogen- and helium-like iron lines since the thermal broadening is much smaller for the heavy iron nuclei.  We can express the ratio of the broadening mechanisms as \\begin{equation} \\frac{\\Delta \\nu_{\\textrm{turbulent}}}{\\Delta \\nu_{\\textrm{thermal}}} = 0.46 \\left(\\frac{v_{rms}}{60 \\,\\textrm{km}\\textrm{s}^{-1}}\\right) \\left(\\frac{k_B T}{5 \\,\\textrm{keV}}\\right)^{-1/2} \\left(\\frac{m_i}{56 \\,m_H}\\right)^{1/2}. \\label{eqn:broadening} \\end{equation} These two contributions are roughly equal for an rms turbulent velocity of 60 km s$^{-1}$ and a temperature of 1 keV.  For the peak velocity fluctuations, the hydrodynamic motions dominate; however, at the rms level, the velocity fluctuations will be harder to see.  Clever modeling of the actual line shapes, as in \\citet{sun03}, may improve the detectability threshold of hydrodynamic turbulence.  As of yet, there have been no successful measurements of ICM velocities.  Had the X-ray calorimeter on \\textit{Astro-E2} (\\textit{Suzaku}) not failed, it would have been able to image the velocity fluctuations of the ICM at the level of the thermal or hydrodynamic broadening.  The next chance for confirming any of these predictions will occur with \\textit{Constellation-X}.  Its spectral resolving power is energy-dependent, but based on current performance requirements, should be able to resolve the predicted level of velocity fluctuations in iron lines with energies around 6 keV. \\subsection{Magnetic Field Geometry}\\label{subsec:discuss:fields} Several authors have noted a discrepancy in the measurements of cluster magnetic fields, finding the fields derived from rotation measure (RM) observations are typically a factor of four to ten higher than fields derived from inverse Compton (IC) measurements \\citep{pet01, ct02}.  The efforts at explaining this discrepancy are not particularly satisfactory.  First, consider fields derived from a rotation measure.  The RM can be written as \\begin{equation} \\mathrm{RM} = 812 \\int_{0}^{L} n_e \\boldsymbol{B}\\cdot\\boldsymbol{\\dif l}\\;\\mbox{radians}\\,\\mbox{m}^{-2}, \\label{eqn:clust:RM} \\end{equation} where $n_e$ is given in cm$^{-3}$, $\\boldsymbol{B}$ is measured in $\\mu$G, and $\\boldsymbol{\\dif l}$ is measured in kpc \\citep{ct02}.  The rotation measure is proportional to the component of the magnetic field along the line of sight times the electron density.  Typically, the field geometry is assumed to be isotropic in direction.  An enhancement in the radial field along the line of sight would represent an enhancement in rotation measure, and an overestimate of the magnetic field strength.  Alternatively, the magnetic field may be derived from inverse Compton emission from the CMB off relativistic electrons.  While the inverse Compton emission has no direct dependence on the magnetic field, the magnetic field strength may nonetheless be derived by comparing the inverse Compton emission to the synchrotron emission, provided they are both produced by the same population of relativistic electrons.  More illumination is provided by examining the formulae for the luminosity of both phenomena: \\begin{eqnarray} L_{sync}& = &\\frac{4}{3}\\beta^2\\gamma^2 c \\sigma_T U_B, \\\\ L_{IC}& = &\\frac{4}{3}\\beta^2\\gamma^2 c \\sigma_T U_{\\gamma}, \\label{eqn:clust:IC-lum} \\end{eqnarray} where $\\beta=v/c$, $\\gamma=(1-\\beta^2)^{-1/2}$, $\\sigma_T$ is the Thompson scattering cross section, $U_B=B^2/8\\pi$ is the magnetic energy density, and $U_{\\gamma}$ is the energy density of the photon field.  The similarity of these two equations is quite evident, giving \\begin{equation} \\frac{L_{IC}}{L_{sync}} \\propto \\frac{U_{\\gamma}}{U_{B}}. \\label{eqn:clust:IC-rat} \\end{equation} The inverse Compton emission peaks in the X-ray around 20 keV, corresponding to electrons with a $\\gamma\\sim 5000$.  The corresponding synchrotron emission peaks in the radio around 100 MHz \\citep{bag98}.  The IC emission is entirely independent of the magnetic field geometry; however, the synchrotron emission is a weak function of magnetic field geometry. Clearly, more detailed analysis will be required in the future; however, the basic idea is intriguing.  By simply positing a magnetic field that is preferentially radial to the line of sight, the rotation measure--derived field is overestimated while the IC/synchrotron--derived magnetic field is not significantly changed.  Future work will examine whether the action of the MTI can explain this observational discrepancy. \\subsection{Temperature Profiles} \\label{subsec:temp} It may seem as if the temperature profile is the most obvious observerable, however we have a ``chicken and egg\" problem.  Our initial conditions posited in \\S\\ref{subsec:method:IC} reflect our observations of clusters \\textit{today}.  Thus what we are observing are relaxed clusters that could be in the saturated state of any instability that has taken place.  In fact, the initial conditions for a relaxed cluster today should probably be taken from a typical disturbed cluster (e.g. Perseus) that has recently undergone a major merger.  Additionally, we have at this point neglected heating sources due to minor mergers, AGN activity, and infall that can be important in heating a cluster.  Nonetheless, we find that significant evolution of the temperature profile can occur in the relevant time between cluster major merger events.  In future work, we will work to more systematically connect our predicted temperature profiles with observations. \\subsection{Thermal Evaporation and Missing Baryons} \\label{subsec:discuss:missing-baryons} There have been several recent suggestions that Sunyaev-Zel'dovich (SZ) and X-ray measurements of galaxy clusters are incompatible.  As mentioned before, the X-ray measurements are sensitive to the line integral of the square of the density; whereas, the SZ signal is directly related to the line integral of the pressure (or thermal energy) of the baryons.  Recently, \\citet{afshordi07} have used composite WMAP data to find that $35 \\pm 8$\\% of the cluster baryons are missing from the ICM. Namely, the authors argue that the gas fraction of clusters is consistently poor relative to the universe as a whole.  The validity of this measurement will likely be confirmed or refuted in the near future as much more accurate SZ surveys, including the Atacama Cosmology Telescope (ACT) and the South Pole Telescope (SPT) as well as numerous others, begin to make maps and measurements of clusters.  This piece of observational evidence has motivated proposals for thermal evaporation of galaxy clusters, such as the work by \\citet{loeb07}.  This work supposes that the outer portions of galaxy clusters could thermally evaporate, losing up to 10\\% of their mass.  The primary proposed mechanism for this is leakage of suprathermal particles with large Coulomb mean free paths.  While this work is perhaps not theoretically robust, it does require magnetic fields that are preferentially radial in direction in order for this large mean free path argument to translate into suprathermal particles escaping.  Radial magnetic fields are of course one of the predictions of the nonlinear MTI simulations to date, thus satisfying one key requirement for this process.  Reconsidering this problem with a methodology similar to Parker's treatment of the solar wind could place this process on a more firm theoretical footing.  Of course, observations and theory are not yet well-developed on this topic, but it is worth keeping these facts in mind until the observations become more certain."
    },
    "0806/0806.1527_arXiv.txt": {
        "abstract": "NGC604 is the largest \\hii-region in M33, second only within the Local Group to 30\\,Dor, and is important as a laboratory for understanding how massive young stellar clusters interact with the surrounding interstellar medium. Here, we present deep (300ks) X-ray imagery of NGC604 obtained as part of the \\cxo\\ ACIS Survey of M33 (ChASeM33), which show highly structured X-ray emission covering $\\sim$70\\% of the full \\ha\\ extent of NGC604. The main bubbles and cavities in NGC604 are filled with hot ($kT=0.5$\\,keV) X-ray emitting gas and X-ray spectra extracted from these regions indicate that the gas is thermal. For the western part of NGC604 we derive an X-ray gas mass of $\\sim$4300\\,$M_{\\odot}$ and an unabsorbed (0.35\\,--\\,2.5\\,keV) X-ray luminosity of $L_{\\rm X}=9.3\\times 10^{35}$\\,erg\\,s$^{-1}$. These values are consistent with a stellar mass loss bubble entirely powered by about 200 OB-stars. This result is remarkable because the standard bubble model tends to underpredict the luminosity of X-ray bright bubbles and usually requires additional heating from SNRs. Given a cluster age of $\\sim$\\,3\\,Myr it is likely that the massive stars have not yet evolved into SNe. We detect two discrete spots of enhanced and harder X-ray emission, which we consider to be fingerprints from a reverse shock produced by a supersonic wind after it collided with the shell wall. In the eastern part of NGC604 the X-ray gas mass amounts to $\\sim$1750\\,$M_{\\odot}$. However, mass loss from young stars cannot account for the unabsorbed X-ray luminosity of $L_{\\rm X}=4.8\\times 10^{35}$\\,erg\\,s$^{-1}$. Off-center SNRs could produce the additional luminosity. The bubbles in the east seem to be much older and were most likely formed and powered by young stars and SNe in the past. A similar dichotomy between east and west is seen in the optical, implying that a massive wall of neutral and ionized gas shields the dynamically quiescent east from the actively star forming west. ",
        "introduction": "NGC604 is the second largest giant \\hii-region in the Local Group and has been studied throughout all wavelengths, from radio \\citep{church99,tosa07}, infrared \\citep{hig03}, optical \\citep{rosa82, teno00}, ultraviolet \\citep{rosa80,keel04}, to the X-ray regime \\citep{apel04}. The main focus of these studies was the investigation of the effects that massive young stars impose on the surrounding multi-phase interstellar medium.  It is commonly believed that NGC604 is ionized by the radiation field and winds from a young and very massive stellar population. A detailed analysis of the stellar content indicates the existence of about 200 O and WR-stars with a WR/O-ratio of about 0.075 \\citep{dris93,hunt96,gonzo00,bruh03}. Although the stellar content, the gas mass, and the average gas density of NGC604 are different compared to other giant \\hii-regions in, e.g., the LMC, II\\,Zw\\,40, or NGC4214 \\citep{teno06}, the blister-like structures and overall \\ha\\ morphology resulting from stellar winds and/or SNe activity is very similar among all these regions. In this regard the similarity with 30\\,Dor in the LMC is the most striking \\citep[see, e.g.,][]{wang99,town06}.  Besides the warm ionized phase of the ISM in NGC604, the hot ionized medium (HIM) is another crucial, yet poorly investigated, phase of the ISM which tells us much about the internal energetics, the wind-ISM interaction, the evolution of the bubble, and the sources which power it. A difficulty in the past originated from the fact that several X-ray bright superbubbles ($\\log L_{\\rm X}>35$) in the LMC show an excess X-ray luminosity in comparison to predictions from the standard bubble model \\citep{cas75,weaver77}. There is general agreement that this excess emission can be attributed to off-center SNRs interacting with the shell walls \\citep{chu90,wang91}. On the other hand, X-ray dim bubbles are found to be consistent with the standard model and do not require additional heating from SNRs \\citep{chu95}. Another subdivision of stellar bubbles was introduced by \\citet{oey96} who distinguished between high velocity bubbles, i.e., those whose expansion velocity is underestimated by the standard model and those which are consistent with it. The emerging picture seems to be that X-ray bright, high velocity bubbles are additionally heated by SNRs and that X-ray dim, slowly expanding bubbles do not require additional heating.  With the current study, which is part of the \\cxo\\ ACIS Survey of M33 \\citep[ChASeM33,][]{plu08}, we present the first detailed X-ray analysis of NGC604 in order to explore the origin of the X-ray emission and to test the wind-blown bubble scenario. Although NGC604 \\citep[$D$\\,=\\,817\\,kpc,][]{freed01} has been detected in previous major X-ray surveys \\citep{long96,misa06,plu08}, the only significant observation that has been published is a \\cxo\\ contour map of the central part \\citep{apel04},  which revealed that the bubbles seen in the optical are filled with X-ray emitting gas. However, the limited sensitivity from this 90\\,ks observation (ObsID \\dataset[ADS/Sa.CXO#obs/02023]{2023}) inhibited further analysis.  The paper is structured as follows: after an outline of the data reduction steps for imaging (Sect. 2.1) and spectroscopy (Sect. 2.2), we present the main results derived from both techniques in Sect. 3. In Sect. 4.1 we discuss the origin of the X-ray gas. By means of spectral modeling, we determine electron densities and filling factors of the HIM and derive X-ray gas masses for individual regions (Sect. 4.1.1). In Sects 4.1.2 and 4.1.3, we cross-check whether the stellar O/WR population in NGC604 can account for the observed X-ray gas masses via continuous stellar mass loss and if the observed X-ray luminosity is consistent with the one predicted from the ``standard'' model of a stellar wind-blown bubble \\citep{weaver77}. Sect. 4 is concluded by evaluating whether SNe produced a significant fraction of the observed X-ray gas mass (Sect. 4.2) and by presenting a consistent picture of NGC604 which emerges from the multi-wavelength data analysis (Sect. 4.3). A summary is provided in Sect. 5. ",
        "conclusions": "We presented the first X-ray images of NGC604, that are deep enough to reveal the complex interplay among different constituents of the ISM. A strong negative morphological correlation was found between \\ha\\ and X-ray emitting gas. All bubbles and cavities are filled with hot coronal gas ($3.3-6.4\\times 10^6$\\,K), suggesting high volume filling factors of about $f_{\\rm X}=0.8$ and densities $n_{\\rm e}<0.6$\\,cm$^{-3}$.  A total, absorption-corrected, diffuse X-ray luminosity in the soft band (0.35\\,--\\,2.5\\,keV) of $L_{\\rm X}=1.43\\times 10^{36}$\\,erg\\,s$^{-1}$ was found for NGC604, which is about 14\\% of the corresponding value for 30\\,Dor \\citep{town06}. The X-ray gas mass in NGC604 amounts to about 6000\\,$M_{\\odot}$, which is only 15\\% of the corresponding mass in 30\\,Dor \\citep{wang99}, and emphasizes the differences between the two largest \\hii-regions of the Local Group even more.  The sum of the detected X-ray gas mass in the western part of NGC604, containing regions B1, B2, C1, and C3, amounts to $\\sim$$4300\\pm1100$\\,M$_{\\odot}$ and the unabsorbed X-ray luminosity of $L_{\\rm X}=9.25^{+0.84}_{-1.14}\\times 10^{35}$\\,erg\\,s$^{-1}$ places this region into the regime of X-ray bright bubbles. Both, the X-ray gas mass and the luminosity agree well with expectations from the standard bubble model and are consistent with mass loss from about 200 O/WR-stars. In view of an estimated cluster age of 3\\,Myr, the consistency between observed and expected luminosities, and the predictions from the starburst models from \\citet{leit92}, the current stellar generation seems to be too young for SNe to have made a significant contribution to heating. The western hemisphere contradicts the general view in which the luminosity of X-ray bright bubbles is underestimated by the standard model and additional contributions from SNRs are required. The X-ray gas likely originates from the stellar O-type population in clusters A and B. We find that mass load from evaporation due to thermal conduction contributes about 95\\% to the X-ray gas mass.  In the eastern hemisphere the estimated mass of X-ray emitting gas amounts to $\\sim$$1750\\pm550$\\,M$_{\\odot}$. The sum of the predicted X-ray luminosity from regions B3 and C2 which one would expect from a shocked stellar wind during the lifetime of the donor stars, largely underestimates the observed unabsorbed luminosity of $L_{\\rm X}=4.84^{+0.62}_{-0.75}\\times 10^{35}$\\,erg\\,s$^{-1}$ by about a factor of 30. This result supports the hypothesis that the luminosity of X-ray bright bubbles is inconsistent with the one predicted by the standard bubble model and argues for additional contributions from SNRs. We rule out a shocked stellar wind produced by the current stellar generation as the main source of the X-ray emission. We cannot rule out that the gas was created during an era when the stellar population was powerful enough to drive a continuous wind and/or SNe heated the gas. The number of expected SNe in C2 and B3, respectively, remains unconstrained due to the lack of spectrophotometric data of the stellar population.  The dichotomy between east and west was also found for the warm ionized gas. \\citet{teno00} noted that the dynamics and ionization of the gas in the eastern part of NGC604 is completely different from the western part. As in the case of the HIM, the \\ha\\ ridge appears to play an important role by separating both hemispheres. The eastern region seems to be an old part of NGC604 in which the warm gas appears to be in a dynamically relaxed condition. In contrast, the western region is dominated by young stars which create a highly dynamical ISM via strong stellar winds and contribute essentially $100\\%$ to the X-ray gas mass and luminosity. The freely-streaming winds in NGC604 should produce characteristic emission features when they interact with the shell walls. Therefore, we interpret the medium X-ray emission along the \\ha\\ ridge and at the \\ha-'cap' as direct fingerprints of such shocks. Provided this is correct, these spots of enhanced emission would confirm the expected supersonic motions of the gas."
    },
    "0806/0806.3714_arXiv.txt": {
        "abstract": "Recent models of spectral formation in magnetars called renewed attention on electron-photon scattering in the presence of ultra-strong magnetic fields. Investigations presented so far mainly focussed on mildly relativistic particles and magnetic scattering was treated in the non-relativistic (Thomson) limit. This allows for consistent spectral calculations up to a few tens of keVs, but becomes inadequate in modelling the hard tails ($\\la 200$ keV) detected by {\\em INTEGRAL} from magnetar sources. In this paper, the second in a series devoted to model the X-/soft $\\gamma$-ray persistent spectrum of magnetar candidates, we present explicit, relatively simple expressions for the magnetic Compton cross-section at resonance which account for Landau-Raman scattering up to the second Landau level. No assumption is made on the magnetic field strength.  We find that sensible departures from the Thomson regime can be already present at $B\\sim 5\\times 10^{12}$~G. The form of the magnetic cross section we derived can be easily implemented in Monte Carlo transfer codes and a direct application to magnetar spectral calculations will be presented in a forthcoming study. ",
        "introduction": "The recent discovery with the {\\em INTEGRAL} satellite of hard X-ray tails in the (persistent) spectra of the magnetar candidates (the anomalous X-ray pulsars, AXPs, and the soft $\\gamma$-repeaters, SGRs; e.g. \\citealt{sandrorev}) provides evidence that a sizeable fraction (up to $\\sim 50\\%$) of the energy output of these sources is emitted above $\\sim 20$ keV. Up to now, high-energy emission has been detected in two SGRs, 1806-20 and 1900+14 \\citep{sandro05, diego06}, and three AXPs, 4U 0142+614, 1RXS J170849-4009 and 1E 1841-045 \\cite[][see also for an updated summary of {\\em INTEGRAL} observations \\citealt{diego08}]{kui04, kui06} \\footnote{Very recently \\cite{Leyder08} reported the {\\em INTEGRAL} detection of the AXP 1E~1048.1-5937, but no spectral information are available yet.}. This result come somehow unexpected, since the spectra of SGRs/AXPs in the soft X-ray range ($\\sim 0.1$--10 keV) are well described by a two component model, a blackbody at $kT\\sim 0.2$--0.6 keV, and a rather steep power-law with photon index $\\Gamma_{soft}\\sim 1.5$--4. {\\em INTEGRAL} observations have shown that in SGRs the power-law tail extends unbroken (or steepens) in the $\\sim 20$--200 keV range, $\\Gamma_{hard}\\ga\\Gamma_{soft}\\sim 1.5$. In AXPs, which have steeper spectra in the soft X-ray band, a spectral upturn appears, i.e. the high-energy power-law is harder than the soft one, $\\Gamma_{hard}\\sim 1$ while $\\Gamma_{soft}\\sim 3$--4 \\cite[see e.g.][for a joint spectral analysis of {\\em XMM-Newton} and {\\em INTEGRAL} data]{nanda08}.  Within the magnetar scenario, the persistent 0.1--10 keV emission of SGRs and AXPs has been successfully interpreted in terms of the twisted magnetosphere model \\citep{tlk02}. In an ultra-magnetized neutron star the huge toroidal field stored in the interior produces a deformation of the crust. The displacement of the footprints of the external (initially dipolar) field drives the growth of a toroidal component which, in turn, requires supporting currents. Charges flowing in the magnetosphere provide a large optical depth to resonant cyclotron scattering (RCS) and repeated scatterings of thermal photons emitted by the star surface may then lead to the formation of a power-law tail. The original picture by \\cite{tlk02} has been further explored by \\cite{lg06}, \\cite{ft07} and \\cite{papI}. Recently \\cite{nanda08} presented a systematic application of the 1D, analytical RCS model of \\cite{lg06} to a large sample of magnetars X-ray spectra finding a good agreement with data in the 0.1--10 keV range. As shown in paper I, more sophisticated 3D Monte Carlo calculations of RCS spectra successfully reproduce soft X-ray observations.  Despite the twisted magnetosphere scenario appears quite promising in explaining the magnetars soft X-ray emission, if and how it can account also for the hard tails detected with {\\em INTEGRAL} has not been unambiguously shown as yet. \\cite{tb05} suggested that the hard X-rays may be produced either by thermal bremsstrahlung in the surface layers heated by returning currents, or by synchrotron emission from pairs created higher up ($\\approx 100$ km) in the magnetosphere. \\cite{bh07,bh08} have recently proposed a further possibility, according to which the soft $\\gamma$-rays may originate from resonant up-scattering of seed photons on a population of highly relativistic electrons. Previous investigations (\\citealt{lg06}; \\citealt{ft07}; paper I) mainly focussed on mildly relativistic particles and magnetic scattering was treated in the non-relativistic (Thomson) limit. This is perfectly adequate in assessing the spectral properties up to energies $\\ll m_ec^2/\\gamma$ (here $\\gamma$ is the typical electron Lorentz factor) since i) the energy of primary photons is low enough ($\\approx 1$ keV) to make resonant scattering onto electrons possible only where the magnetic field has dropped well below the QED critical field, ii) up-scattering onto electrons with $\\gamma\\approx 1$ hardly boosts the photon energy in the hundred keV range, so electron recoil is not important. However, some photons do actually gain quite a large energy (because they scatter many times on the most energetic electrons) and fill a tail at energies $\\ga 50$ keV. We caveat that, despite in previous works spectra have been computed up to the MeV range, they become untrustworthy above some tens of keV and can not be used to assess the spectral features that can arise due to electron recoil effects (i.e. a high-energy spectral break). Proper investigation of the latter demands a complete QED treatment of magnetic Compton scattering. This is mandatory if highly relativistic particles are considered because a photon can be boosted to quite large energies in a single scattering and, if it propagates towards the star, it may scatter again where the field is above the QED limit.  Monte Carlo numerical codes for radiation transport in a magnetized, scattering medium, as the one we presented in paper I, make an excellent tool to investigate in detail the properties of RCS in the case in which energetic electrons are present in addition to the mildly relativistic particles which are responsible for the formation of the soft X-ray spectrum. The Compton cross-section for electron scattering in the presence of a magnetic field was first studied in the non relativistic limit by \\cite{clr71}, and the QED expression was derived long ago by many authors \\citep{her79,dh86,bam86,hd91}. However, its form is so complicated to be often of little practical use in numerical calculations. Moreover, because of their inherent complexity, many of the published expressions are affected by misprints and the comparison between different derivations is often problematic. On the other hand, the use of the full expression of the cross section is especially needed when a detailed model of cyclotron line formation has to be evaluated, including expectations for the line profile \\citep[see e.g.][]{dv78, ah99}. In the situation we are considering, it is reasonable to assume that nearly all photons scatter at resonance. Non-resonant scattering contributions have negligible effects on shaping the overall spectrum, except possibly in the very neighborhood of a cyclotron line peak.  Motivated by this, we present here explicit, relatively simple expressions for the magnetic Compton cross-section at resonance that can be then included in Monte Carlo calculations such as that described in Paper I. In doing so, we investigate the behaviour of the different terms and assess their relative importance. The final expressions have been cross-checked by comparing different published formulations, when specialized at resonance. A direct application of the results discussed here to spectral calculations will be presented in a forthcoming study (Nobili, Turolla \\& Zane, in preparation). The paper is organized as follows. In \\S\\ref{first} we formulate the problem and summarize the main ingredients needed for a Monte Carlo simulation. In \\S\\ref{Master} we compute the relevant cross sections, specified at the resonance, while the transition rates (which are related to the natural width of the excited resonant levels) are given in \\S\\ref{TranRate}. The creation of photon via spawning effects is discussed in \\S\\ref{Spawning}, while \\S\\ref{abs} contains a brief comparison between scattering and absorption. In \\S\\ref{mean} we compute the optical depth, which yields the probability of scattering. Conclusions follow in \\S\\ref{conc}. ",
        "conclusions": "\\label{conc}  Recent models of spectral formation in magnetars called renewed attention on electron-photon scattering in the presence of ultra-strong magnetic fields. The complete expression for the QED cross section is known since long \\cite[e.g.][]{hd91} but its practical application in the relativistic regime is numerically challenging. In many astrophysical problems, including the one which motivated us, it is reasonable to assume that scattering occurs only at resonance, i.e. when the incident photon frequency equals the cyclotron frequency (or one of its harmonics). Restricting to resonant scattering introduces a major simplification, and here we presented explicit expressions for the magnetic Compton cross section in this particular case. Our main goal has been to provide a complete, workable set of formulae which can be used in  Monte Carlo simulations of photon scattering in strongly magnetized media.  Our results are fairly general and can be applied to resonant photon scattering under a variety of conditions. In particular, no assumption is made on the field strength. Having in mind applications to spectral modelling in the $\\,\\sim 0.1$--200 keV range, resonant scattering necessarily occurs where $\\, {\\cal B}/{\\cal B}_{cr}\\la 1$. Under these conditions, and although the expressions we derived are still valid for $\\, {\\cal B}/{\\cal B}_{cr}>1$, we restricted to the case in which the electron is excited at most up to the second Landau level. We find that deviations from the non-relativistic limit in both the first and second resonant contributions to the cross section become significant for $\\, {\\cal B}/{\\cal B}_{cr} \\ga 0.1$. The probability that scattering occurs at the second resonance, which is negligible below $\\, {\\cal B}/{\\cal B}_{cr} \\la 0.01$,  becomes sizeable at higher $B$ and, depending on the scattering angle, it can be up to $\\,\\sim 30\\%\\, $ for $\\, {\\cal B} \\sim {\\cal B}_{cr}$. In case the second Landau level is excited, its is more likely that the recoiled electron is left in the first than in the ground level with the ensuing emission of a new photon (spawning). Using our results for the cross section together with known expressions for the transition rates, we checked under which conditions resonant Compton scattering can be treated as the combination of two first-order processes, photon absorption followed by emission. While the scattering and absorption cross sections differ by at most $\\,\\sim 20\\%\\, $  for $\\, {\\cal B} \\la {\\cal B}_{cr}$, the angular distribution of the scattered/emitted photons shows deviations already at $\\,{\\cal B} \\sim 0.1{\\cal B}_{cr}$. Finally, having in mind the implementation in a Monte Carlo code  (Nobili, Turolla \\& Zane, in preparation), we presented an explicit derivation of the scattering optical depth along the photon path."
    },
    "0806/0806.4302_arXiv.txt": {
        "abstract": "The energy spectrum of cosmic rays above~\\unit{$2.5\\times10^{18}$}{\\eV}, derived from 20,000 events recorded at the Pierre Auger Observatory, is described. The spectral index $\\gamma$ of the flux, $J\\propto E^{-\\gamma}$, at energies between \\unit{$4\\times10^{18}$}{\\eV} and \\unit{$4\\times10^{19}$}{\\eV} is $2.69\\pm0.02\\usk\\text{(stat)}\\usk\\pm0.06\\usk\\text{(syst)}$, steepening to $4.2\\pm0.4\\usk\\text{(stat)}\\usk\\pm0.06\\usk\\text{(syst)}$ at higher energies, consistent with the prediction by Greisen and by Zatsepin and Kuz'min. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0806/0806.1657_arXiv.txt": {
        "abstract": "We report extensive photometry of the frequently outbursting dwarf nova RZ Leo Minoris. During two seasons of observations we detected 12 superoutbursts and 7 normal outbursts. The $V$ magnitude of the star varied in range from 16.5 to 13.9 mag. The superoutbursts occur quite regularly flashing every 19.07(4) days and lasting slightly over 10 days. The average interval between two successive normal outbursts is 4.027(3) days. The mean superhump period observed during the superoutbursts is $P_{\\rm sh}=0.059396(4)$ days ($85.530\\pm0.006$ min). The period of the superhumps was constant except for one superoutburst when it increased with a rate of $\\dot P/P_{\\rm sh} = 7.6(1.9)\\cdot 10^{-5}$. Our observations indicate that RZ LMi goes into long intervals of showing permanent superhumps which are observed both in superoutbursts and quiescence. This may indicate that decoupling of thermal and tidal instabilities play important role in ER UMa systems. No periodic light variations which can be connected with orbital period of the binary were seen, thus the mass ratio and evolutionary status of RZ LMi are still unknown.   \\noindent {\\bf Key words:} Stars: individual: RZ LMi -- binaries: close -- novae, cataclysmic variables ",
        "introduction": "Dwarf novae are believed to be unmagnetized close binary systems containing white dwarf primary and low mass main sequence secondary. The secondary fills its Roche lobe and looses the material through the inner Lagrangian point. This matter forms an accretion disc around the white dwarf.  One of the most intriguing classes of dwarf novae are SU UMa stars which have short orbital periods (less than 2.5 hours) and show two types of outbursts: normal outbursts and superoutbursts. Superoutbursts are typically about one magnitude brighter than normal outbursts, occur about ten times less frequently and display characteristic tooth-shape light modulations i.e. so called superhumps (see Warner 1995 for review).  The behavior of SU UMa stars in now quite well understood within the frame of the thermal-tidal instability model (see Osaki 1996 for review). Superhumps occur at a period slightly longer than the orbital period of the binary system. They are most probably the result of accretion disk precession caused by gravitational perturbations from the secondary. These perturbations are most effective when disk particles moving in eccentric orbits enter the 3:1 resonance. Then the superhump period is simply the beat period between orbital and precession rate periods. Although in the last decades significant progress has been made in explaining the behaviour of dwarf novae light curves, some physical processes ongoing in these systems are still not fully understood (see for example Smak 2000, Schreiber and Lasota 2007).  In the beginning of 90ties of XX century, SU UMa stars were believed to be quite uniform group of variables with common properties. These objects went into superoutburst every year or so and between two successive superoutbursts showed $\\sim$10 ordinary outbursts. However, there were some exceptions like WZ Sge, which show infrequent and large amplitude superoutburst followed by the period of quiescence with no single eruption lasting even 30 years.  In 1995 astronomical community was alerted about the presence of stars characterized by complete opposite behavior. First, Misslet and Shafter (1995) reported observations of PG 0943+521 (later called ER UMa), which allowed to detect superhumps with period of 0.0656 days and include this object into the SU UMa group of variables. The most intriguing feature of the long term light curve of ER UMa was very short interval between two successive superoutbursts (so called supercycle) reaching only 44 days. This value was about three times shorter than shortest previously known supercycles. This work was quickly followed by paper of Robertson et al. (1995), who confirmed all findings of Misslet and Shafter (1995) and precisely determined the value of supercycle of ER UMa to be equal to 42.95 days. Moreover, they found two more objects with similar properies - V1159 Ori with supercycle of 44.5 days and RZ LMi with supercycle as short as 18.87 days! In the same year Nogami et al. (1995) published paper which confirmed extremely short supercycle of RZ LMi and showing that it belongs to SU UMa variables exhibiting clear superhumps with period of 0.05946 days.  One year later the number of these unusual variables increased to four objects. Kato et al. (1996) reported the discovery that DI UMa has a supercycle of 25 days and shows clear superhumps with period of 0.0555 days.  The fifth ER UMa-type variable - IX Dra - was discovered by Ishioka et al. (2001). Their observations revealed a supercycle length of 53 days and an interval between normal outbursts of 3-4 days. Olech et al. (2004) determined precisely both superhump and orbital periods of the binary and estimated the supercycle length to 54 days.  The basic properies of five known up-today members of ER UMa group are summarized in Table 1.  \\begin{table}[!h] \\caption{\\sc Basic properties of ER UMa variables. $P_{\\rm orb}$ and $P_{\\rm sh}$ denote orbital and superhump periods, $\\epsilon$ is a period excess, $T_s$ and $T_n$ are supercycle and cycle periods, $T_{\\rm sup}$ is duration of the superoutburst, $A_{\\rm sup}$ and $A_{\\rm n}$ are amplitudes of superoutburst and normal outburst.} \\vspace{0.1cm} \\begin{center} \\begin{tabular}{|l|c|c|c|c|c|c|c|c|c|} \\hline \\hline Star & $P_{\\rm orb}$ & $P_{\\rm sh}$ & $\\epsilon$ & $T_s$ & $T_n$ & $T_{\\rm sup}$ & $A_{\\rm sup}$ & $A_{\\rm n}$ & Ref\\\\ & [days] & [days] & [\\%] & [days] & [days] & [days] & [mag] & [mag] & \\\\ \\hline \\hline RZ LMi    & ?        & 0.05946  &  ?   & 18.9      & 3.8 & 11 & 2.5 & 2.0 & (1,2)\\\\ DI UMa    & 0.054564 & 0.0555   & 1.72 & 25.0      & 5.0 & 12 & 2.9 & 2.1 & (3,4)\\\\ ER UMa    & 0.06366  & 0.065552 & 2.97 & 43.0      & 4.4 & 23 & 2.6 & 2.2 & (2,5,6)\\\\ V1159 Ori & 0.062178 & 0.064284 & 2.11 & 44.6-53.3 & 4.0 & 16 & 2.2 & 1.4 & (2,6,7,8)\\\\ IX Dra    & 0.06646  & 0.066968 & 0.76 & 54.0      & 3.1 & 16 & 2.2 & 1.7 & (9,10)\\\\ \\hline \\hline \\multicolumn{10}{l}{\\small 1. Nogami et al. (1995), ~2. Robertson et al. (1995), ~3. Kato et al. (1996), ~4. Thorstensen et al. (2002)}\\\\ \\multicolumn{10}{l}{\\small 5. Kato et al. (2003), ~6. Thorstensen et al. (1995), ~7. Kato (2001), ~8. Patterson et al. (1995)}\\\\ \\multicolumn{10}{l}{\\small 9. Ishioka et al. (2001), ~10. Olech et al. (2004)}\\\\ \\end{tabular} \\end{center} \\end{table} \\bigskip  It is clear that ER UMa stars consist a group of variables with common properties such as extremely short supercycles, small amplitudes of eruptions and relatively long superoutbursts lasting even longer than half of the supercycle. However, the period excess $\\epsilon$ defined as $P_{\\rm sh}/P_{\\rm orb}-1$, which is connected with mass ratio by the following relation:  \\begin{equation} \\epsilon\\approx\\frac{0.23q}{1+0.27q} \\end{equation}  \\noindent suggests different evolutionary status of the particular members of ER UMa stars. For example, ER UMa and V1159 Ori seem to have normal secondaries and evolve towards the shorter orbital periods. On the other hand, DI UMa and IX Dra are much more evolved objects with sub-stellar secondaries (possibly degenerate brown dwarfs) and evolve towards the longer orbital periods (Patterson 2001, Olech et al. 2004). The question why DI UMa and IX Dra are so active, while WZ Sge stars having similar period excesses have longest supercycles, is still open. ",
        "conclusions": "\\subsection{Evolutionary status of RZ LMi}  From our Table 1 summarizing main properies of ER UMa stars, it is clear that these objects have many common properties but may be divided into two subgroups probably with different evolutionary status. Fig. 15, repeated after Patterson (1998, 2001) and Olech et al. (2004), shows correlation between period excess (i.e. mass ratio) and orbital period of the system. The solid line shows the evolutionary track of a dwarf nova with a white dwarf of mass 0.75 $\\cal M_\\odot$ and secondary component with effective radius 6\\% larger than that of single main sequence star. The nova evolves towards the shorter periods first due to the magnetic braking, next due to the emission of gravitational waves. After reaching the period minimum, the secondary becomes degenerate brown dwarf and system starts to increase its orbital period.  \\vspace{10.2cm}  \\special{psfile=fig15.ps hoffset=20 voffset=-208 vscale=80 hscale=80 angle=0}  \\begin{figure}[h] \\caption {\\sf The relation between the period excess and orbital period of the system. The solid line corresponds to the evolutionary track of a binary with a white dwarf of $0.75 {\\cal M}_\\odot$ and a secondary with effective radius 6\\% larger than in the case of an ordinary main sequence star. Calculations were made under the assumption that below the orbital period of two hours the angular momentum loss in only due to gravitational radiation. Triangles denote the positions of ER UMa and V1159 Ori.} \\end{figure}   It seems that DI UMa and IX Dra (both belonging to ER UMa stars) are such evolved period bouncers, which in fact should be similar to old and inactive WZ Sge stars (WZ Sge, AL Com and EG Cnc showed in the plot). On the other hand, ER UMa and V1159 Ori, shown as filled triangles, seem to be much younger objects still evolving towards shorter periods.  Where is the place of RZ LMi? It is difficult to answer this question without knowledge about the orbital period of the system. Our photometric data showed no other short term modulations than these corresponding to the ordinary superhumps. It would be very tempting to make the spectroscopic observations of the star in quiescence. With minimum brightness of 16.5 mag it can be done with 2-3-meter class telescope.  \\subsection{Stability of the supercycle}  The comprehensive analysis of the global light curve of RZ LMi made by Robertson et al. (1995) and  based on almost three years observing period showed that supercycle of RZ LMi is not stable. Their $O-C$ diagram for supermaxima was characterized by clear decreasing trend with $\\dot P = -1.7 \\cdot 10^{-3}$. However graph shows also occasional jumps where particular superoutburst occur even 5 days before or after the predicted moment. If this decreasing trend would continue to the epoch of our observations the supercycle should be then around 18.5 days, which is in disagreement with determined value of 19.07 days.  Our global light curve spans only two seasons and has no enough data to construct reliable $O-C$ diagram for supermaxima. However, quick look at Fig. 1, could draw some valuable conclusions. In 2004 the 19-day periodicity is preserved through all superoutbursts except eruption number IV. In this case, we, in fact, are not certain whether we deal with superoutburst which occurred slightly before predicted moment or exceptionaly bright normal outburst lasting longer than usual. Vicinity of eruption number IV is also the time when disk could loose its eccentricity, expel the matter via this long outburst and rebuilt eccentricity again in superoutburst no. V.  Data from 2005 seem to confirm stability of 19-day supercycle. The superoutburst no. XXIV, which has the best observational coverage, occurs at right time according to 19-day ephemeris. The problem is  with superoutburst no. XXIII, where instead of supermaximum we noted two ordinary outbursts. Our light curve, however, does not exclude possibility that supermaximum occurred a few days earlier according to the ephemeris.  Mass transfer from the secondary to the disk, building the eccentricity, ignition of the outbursts and superoutbursts due to the thermal and tidal instabilities are stochastic processes, which are far for regularity. The question is why RZ LMi is so regular? Even if we observe some shifts in time of the start of particular supermaximum, the clock returns to stability without shift of the phase of whole pattern. This is hard to explain from the point of view of standard thermal-tidal instability model and might need some help from, for example, external force. The present number of known SU UMa systems reached the level for which the statistics tells us that some of these close binaries might be orbited by a third body. Is this in case of RZ LMi? We do not know. But the hypothesis that 19-day period is the orbital period of the third body (or some kind resonant value) and cause of both the stability of supercycle and high activity of the star, which without this body would be quiet WZ Sge object, is tempting.  \\subsection{Permanent superhumper?}  The standard thermal-tidal instability model is unable to produce supercycles shorter that 40 days. Activity of the ordinary SU UMa variable can be increased by increasing a mass transfer rate. But when it reaches $\\dot M \\approx 3\\cdot 10^{16}$ g/s the supercycle starts to lenghten again due to the fact that superoutburst lasts longer. Further increasing of mass transfer causes transition of the star to the group of permanent superhumpers which are in permanent state of supermaximum and show infinite value of supercycle.  Osaki (1995) tried to explain properties of RZ LMi by artificial ending the superoutburst at the moment, when the disk had shrunk from $0.46a$ to only $0.42a$, whereas a typical value used for ordinary SU UMa stars is $0.35a$.  Hellier (2001) suggested that the source of the premature end of superoutburst might be a cooling wave propagating from the region outside $0.46a$ causing transition of the disk to the cold state when still eccentric. This decoupling of tidal and thermal stability brings the star to the minimum light with still precessing and elliptical disk. This hypothesis is confirmed by observations of ordinary superhumps both in quiescence and normal outbursts of V1159 Ori and ER UMa (Patterson et al. 1995, Gao et al. 1999, Zhao et al. 2006).  Our observations shows that RZ LMi also shows superhumps in minimum light. Additionally, for the first time, we demonstrated that in interval covering at least 60 days (including superoutbursts numbers I, II and III) the star was showing superhumps with constant period which can be described by common ephemeris and phased without any phase shift. It indicates that decoupling could have place in this case and the disk of RZ LMi was eccentric and precessing in the entire 60-day period."
    },
    "0806/0806.3470_arXiv.txt": {
        "abstract": "The Pierre Auger Observatory reports that 20 of the 27 highest energy cosmic rays have arrival directions within 3.2$^\\circ$ of a nearby galaxy in the Veron-Cetty \\& Veron Catalog of Quasars and Active Galactic Nuclei (12$^{\\rm th}$ Ed.), with $\\sim 5$ of the correlations expected by chance.  In this paper we examine the correlated galaxies to gain insight into the possible UHECR sources.  We find that 14 of the 21 correlated VCV galaxies are AGNs and we determine their bolometric luminosities.  The remaining 7 are primarily star-forming galaxies.  The bolometric luminosities of the correlated AGNs are all greater than $ 5 \\times 10^{42} \\,{\\rm erg \\,s^{-1}}$, which may explain the absence of UHECRs from the Virgo region in spite of the large number of VCV galaxies in Virgo, since most of the VCV galaxies in the Virgo region are low luminosity AGNs.  Interestingly, the bolometric luminosities of most of the AGNs are significantly lower than required to satisfy the minimum condition for UHECR acceleration in a continuous jet.  If a UHECR-AGN correlation is substantiated with further statistics, our results lend support to the recently proposed ``giant AGN flare\" mechanism for UHECR acceleration. ",
        "introduction": "The origin of ultra-high energy cosmic rays (UHECRs) has been an important question in astrophysics, for many decades.  Early observations \\citep{AGASA1998} suggested a violation of the Grizsen-Zatsepin-Kuzmin (GZK) prediction that the UHECR energy spectrum must drop off at energies above $\\sim 60 $ EeV. If confirmed, this would have been a major challenge for theory, but recent high-statistics observations by the HiRes and Pierre Auger collaborations find a downturn in the spectrum consistent with the GZK prediction \\citep{HRspec08,augerSpec08}.  Nonetheless, the puzzle of which astrophysical sites are capable of accelerating UHECRs remains open, and the answer will be of fundamental importance for our understanding of gamma ray bursts (GRBs), active galactic nuclei (AGNs) and quasars, and other extreme systems.  Efforts to find angular correlations between UHECR arrival directions and candidate astrophysical sources have been bedeviled until recently by a combination of inadequate statistics and the fact that UHECRs, being charged particles, are deflected by magnetic fields en route from their sources.  The Pierre Auger Observatory's discovery of a significant correlation between the highest energy cosmic rays and nearby galaxies in the \\citet{VCV} (VCV) Catalog of Quasars and Active Galactic Nuclei (12$^{\\rm th}$ Ed.) \\citep{augerScience07, AugerLongAGN} (Auger07a,b below) is an important step toward identifying the sources of UHECRs.  Of the 27 cosmic rays above 57 EeV recorded prior to Aug. 31, 2007, twenty are within 3.2$^\\circ$ of a VCV galaxy with $z \\leq 0.018$ (about 75 Mpc).  Restricting to $|b|>10^\\circ$, where the VCV catalog is more complete, there are 22 UHECRs of which 19 are correlated. About 5 of these correlations would be expected by chance if the arrival directions of UHECRs were isotropic.  The distribution of magnetic deflections is not known {\\em a priori} and it is not yet known with certainty whether the UHECRs are protons or nuclei. Consequently the expected correlation between UHECR energy and maximum source distance is not certain.  The strategy used by Auger was to search for correlations by scanning over UHECR energy threshold, maximum angular separation, and maximum source redshift, to find the values of these parameters that maximize the significance of the correlations.  The VCV catalog and the parameters given above were identified using data through May 31, 2006 and a ``prescription\" formulated for an independent test.  The correlation was confirmed with independent subsequent data, taken from June 1, 2006 to Aug. 31, 2007, with a probability of less than 1\\% of occurring by chance.  This Auger result is of fundamental importance to particle astrophysics, because the correlation with nearby extragalactic structure clearly demonstrates that UHECRs are of extragalactic origin and that the highest energy cosmic rays have a horizon, consistent with the GZK effect. (The downturn in the spectrum might merely be due to a maximum energy of the accelerators.)  However as stressed by the Auger collaboration, the observed correlation may not mean that the correlated UHECRs are produced by galaxies with which they are correlated:  the VCV galaxies may just be tracers of the true sources.  Our purpose here is to examine the VCV galaxies that correlate with the 20 Auger UHECRs, as a step toward elucidating properties of the sources of UHECRs.  We emphasize that the correlation observed by Auger is only statistical: $\\sim$1/4 of the correlations are expected simply by chance.  We cannot be confident that any given one of the correlated VCV galaxies is a source.  However the degree of correlation observed by Auger is so high that galaxy clustering alone cannot account for it (GRF, A. Berlind and IZ, in preparation); of order half or more of the correlated galaxies are most likely the sources of their associated UHECR, enabling us to obtain statistically useful information on source properties even though we do not have a pure sample of sources.  The underlying principle is that the chance correlations must be representative of the ensemble of VCV galaxies, and therefore their presence does not distort the conclusions when properly used.  The VCV catalog is a list of the active galactic nucleus, quasar and BL Lac candidates reported in the literature, based on heterogeneous selection criteria.  While it is the largest available collection of known AGNs, especially for the southern hemisphere where most of Auger's exposure lies, it has several deficiencies compared to an ideal catalog. The VCV catalog is known to be incomplete and non-uniform. Furthermore, it is not pure.  Here we assemble the correct classifications of the VCV galaxies correlated with the Auger UHECRs.  In addition, we determine the bolometric luminosity of each AGN, which is an important diagnostic of the maximum UHECR energy an AGN can produce in conventional jet acceleration.  Finally, we look for evidence of a bolometric luminosity threshold for AGNs which may be responsible for UHECR production.  We examine the AGNs in the Virgo region, from which no UHECRs were detected, and find that most are much lower in their bolometric luminosity than the AGNs which are correlated with UHECRs.  Thus the lack of UHECRs from the direction of Virgo may simply reflect the existence of a minimum bolometric luminosity of AGNs responsible for accelerating UHECRs.  We also estimate the fraction of low-luminosity AGNs (LLAGNs, taken here to have $L_{\\rm bol} \\leq 5 \\times 10^{42} {\\rm erg\\, s^{-1}}$) in the VCV catalog in the HiRes exposure region, finding that about half the AGNs are below the luminosity threshold of AGNs correlated with Auger UHECRs.  The fraction of known LLAGNs in VCV is higher in the HiRes field of view than in Auger's, due to the sensitive northern hemisphere survey by \\citet{Ho1995}; this may contribute to HiRes' not observing  a significant correlation between VCV galaxies and their UHECR data \\citep{HRAGN08}. ",
        "conclusions": "We have examined the 21 galaxies from the VCV catalog with $z \\leq 0.018$ which are within 3.2$^\\circ$ of the 27 highest-energy cosmic rays, and 8 additional AGNs from NED introduced by \\citet{MSPC}.  We find that 14 of the 21 VCV galaxies are AGNs and the other seven are either star forming or quiescent galaxies.  X-ray or radio observations are needed to find out if any of these seven have obscured nuclear activity at an interesting level.  Two of the 8 additional NED galaxies can be optically established as AGNs (one is a Seyfert 2, one a LINER, and 6 do not have adequate optical spectra to make a determination).  The new optically identified AGN, WKK 4374, increases by one the number of UHECR-AGN correlations, because it is within 2.8$^\\circ$ of a previously uncorrelated UHECR.  We have determined the spectral type and bolometric luminosity of the 14 VCV AGNs and the newly identified AGN.  Their luminosities range from $5 \\times 10^{42}$ erg s$^{-1}$ to 1 $\\times$ $10^{46}$ erg s$^{-1}$. Five are broad line, nine are narrow-line and one is a BL Lac.  \\citet{haoAGN05} finds that the numbers of broad line and narrow-line AGNs are comparable at low and moderate luminosity, while at high luminosity, broad line AGNs far outnumber narrow-line ones.  Thus the AGNs possibly producing the observed UHECRs are representative of the population of moderate luminosity AGNs both in luminosity and type.  We have also identified and determined the bolometric luminosities of the Virgo AGNs and LINERs. Most have luminosities more than an order of magnitude below the lowest luminosity of any AGN correlated with a UHECR.  (A large number of low luminosity AGNs and LINERs are known in Virgo, due to the \\citet{Ho1995} high-sensitivity spectroscopic survey of the near-by galactic nuclei in the northern hemisphere.)  A minimum bolometric luminosity requirement for UHECR-accelerating systems may therefore explain the absence of UHECRs correlated with Virgo AGNs.  If there is such a threshold, the significance of correlations between UHECRs observed by northern hemisphere observatories and VCV galaxies would be reduced in comparison to that observed by Auger South, due to a greater dilution of the VCV catalog in the northern hemisphere by large numbers of known very low luminosity AGNs.  A threshold in the bolometric luminosity is consistent with the Giant AGN Flare scenario \\citep{FarrarandGruzinov}, but is not predicted by it, since the properties of the remnant have not yet been modeled.  Altogether, 21 of the 27 Auger UHECRs are within $3.2^\\circ$ of a Veron-Cetty Veron galaxy or the newly identified AGN.  Of these, 17 of the 21 correlated UHECRs can be attributed to an identified AGN within $3.2^\\circ$ and $z \\leq 0.018$.  However few of the correlated AGNs satisfy the condition $L_{\\rm bol} \\geq 10^{45} E_{20}^2~{\\rm erg \\, s^{-1}}$, required to confine a cosmic ray as it is accelerated \\citep{FarrarandGruzinov}.  If correlated AGNs with inadequate bolometric luminosities are in fact the sources of most of these UHECRs, their luminosities present a puzzle to the conventional picture of AGN acceleration in a continuous jet.  The recent analysis of \\citet{MSPC} finds that few of the correlated AGNs have powerful radio jets or lobes, corroborating this conclusion.   If the AGN-UHECR correlation is a real one, these results provide evidence in favor of the new mechanism of UHECR acceleration by giant AGN flares proposed by \\citet{FarrarandGruzinov}, in which a modest AGN has an intense flare producing a jet of the required luminosity (initiated for instance by a stellar tidal disruption rapidly heating the accretion disk) and then reverts to a mild-mannered existence."
    },
    "0806/0806.1828_arXiv.txt": {
        "abstract": "We reexamine the production of gravitational waves by bubble collisions during a first-order phase transition. The spectrum of the gravitational radiation is determined by numerical simulations using the \"envelope approximation\".  We find that the spectrum rises as $f^{3.0}$ for small frequencies and decreases as $f^{-1.0}$ for high frequencies. Thus, the fall-off at high frequencies is significantly slower than previously stated in the literature. This result has direct impact on detection prospects for gravity waves originating from a strong first-order electroweak phase transition at space-based interferometers, such as LISA or BBO. In addition, we observe a slight dependence of the peak frequency on the bubble wall velocity. ",
        "introduction": "Colliding bubbles in a first-order phase transition constitute one possible source of stochastic gravitational wave (GW) radiation~\\cite{Witten:1984rs, Hogan}. If the electroweak phase transition is strongly first-order, for instance, the kinetic energy stored in the Higgs field and the bulk motion of the plasma is partially released into gravity waves. This happens mostly at the end of the phase transition, when collisions break the spherical symmetry of the individual Higgs field bubbles.  This possibility was systematically analyzed in a series of papers~\\cite{Kosowsky:1991ua, Kosowsky:1992rz, Kosowsky:1992vn, Kamionkowski:1993fg}. The first simulation~\\cite{Kosowsky:1991ua} consisted hereby of the full scalar field dynamics of two bubbles in vacuum, where the essential observation was made that the emitted radiation depends only on the gross features of the problem, namely the kinetic energy stored in the uncollided bubble regions. This observation is the basis of the so-called envelope approximation that opened up the possibility of simulating phase transitions with a large number of bubbles. This was subsequently exploited in refs.~\\cite{Kosowsky:1992rz, Kosowsky:1992vn} and further refined for a thermal environment in ref.~\\cite{Kamionkowski:1993fg}.  In the case of only two colliding bubbles, the spectrum decreases as $f^{-1.8}$ for high frequencies. But this result might be special to the case of two bubbles, where the collision never finishes, which makes the introduction of a time cutoff function mandatory~\\cite{Kosowsky:1991ua}. The situation is different if a realistic phase transition with a large number of bubbles is simulated. There were hints in refs.~\\cite{Kosowsky:1992rz, Kosowsky:1992vn, Kamionkowski:1993fg} that the spectrum of multi-bubble simulations might be more flat than in the two bubble case, but the numerical accuracy prohibited a conclusive statement. As a result, the frequency fall-off of the two bubble case is still being used in the present day literature~\\cite{Grojean:2006bp, KonHub, Kahniashvili:2008pf}.  The aim of the present work is to reexamine the generated spectrum of GWs by simulating a phase transition with a large number of bubbles, making use of the aforementioned envelope approximation. Compared to ref.~\\cite{Kamionkowski:1993fg}, the numerical accuracy will be considerably improved, and a larger portion of the spectrum will be determined to allow for a careful analysis of the high frequency behavior. ",
        "conclusions": "We reexamined the spectrum of gravitational wave radiation generated by bubble collisions during a first-order phase transition in the envelope approximation.  Using refined numerical simulations, our main finding is that the spectrum falls off only as $f^{-1.0}$ at high frequencies, considerably slower than appreciated in the literature. This behavior is most probably related to the many small bubbles nucleated at a later stage of the phase transition~\\footnote{The effect has been mentioned in refs.~\\cite{Kosowsky:1992rz, Kosowsky:1992vn, Kamionkowski:1993fg}, but could never be clarified because of numerical uncertainties.}. This result is especially interesting in the light of recent investigations~\\cite{Grojean:2006bp, KonHub} that indicate that in the case of a first-order electroweak phase transition (obtained by a singlet sector~\\cite{Huber:2000mg, Espinosa:2007qk} or higher dimensional operators~\\cite{Grojean:2004xa, Bodeker:2004ws, Delaunay:2007wb}) the peak frequency of the produced radiation is below the best sensitivity range of planed satellite experiments, such as LISA and BBO~\\cite{Danzmann:2003tv, Corbin:2005ny}. This effect is shown in Fig.~\\ref{fig_spectrum} for several typical parameter sets for the phase transition in the nMSSM~\\cite{KonHub}. Notice that the discussion in ref.~\\cite{KonHub} suggests that stronger phase transitions in general lead to smaller peak frequencies due to a decrease in the parameters $\\beta/H$ and $T_*$. This amplifies the importance of the high frequency part of the gravitational wave spectrum. Notice also that a flatter spectrum simplifies the distinction from other sources of stochastic gravitational waves, such as turbulence~\\cite{Kosowsky:2001xp, Dolgov:2002ra, Caprini:2006jb, Gogoberidze:2007an} or preheating after inflation~\\cite{Khlebnikov:1997di, Easther:2006gt, Easther:2006vd, GarciaBellido:2007dg, GarciaBellido:2007af, Dufaux:2007pt}. Besides, we found that the peak frequency slightly depends on the expansion velocity of the bubbles and decreases for higher wall velocities. Our quantitative results are summarized by eqs.~(\\ref{res1})-(\\ref{res3}).  Finally, we would like to comment on the recent paper \\cite{Caprini:2007xq}, where an analytic approach to the GW production by collisions based on stochastic considerations was presented. In this approach, assumptions have to be made about the time-dependence of unequal time correlations of the velocity field. In their favored model, the authors obtain a scaling as $\\omega^{-2}$ for the high frequency part of the spectrum.  We suspect that this disparity is due to conceptual differences.  First, notice that the treatment presented here is based on two main ingredients: The thin wall and the envelope approximations. Even though the stochastic approach in ref.~\\cite{Caprini:2007xq} does not require the thin wall approximation, the results are also valid in this limit, such that this approximation cannot be responsible for the different spectra. However, in the approach of ref.~\\cite{Caprini:2007xq} the collided and uncollided regions of the bubbles are treated equally but in a stochastic manner. Breaking of the spherical symmetry, necessary for GW production, is encoded in assumptions on the velocity correlation functions.  This is in contrast to our analysis in which the well tested \\cite{Kosowsky:1992rz} envelope approximation breaks the spherical symmetry in a realistic way.  Furthermore, only the time dependent mean bubble radius enters into the stochastic calculation, while in the analysis at hand bubbles with a realistic size distribution are simulated. The occurrence of many small bubbles probably enhances the high frequency part of the spectrum, as already argued in refs.~\\cite{Kosowsky:1992rz, Kosowsky:1992vn, Kamionkowski:1993fg}.  Finally, it is interesting to see that even in the stochastic treatment the high frequency part of the spectrum can scale as $\\omega^{-1}$ (in agreement with our results) if one assumes that the source is fully uncorrelated at unequal times. Especially at late stages of the phase transition, many small bubbles are generated, start to collide and are absorbed by neighboring bubbles at a large rate such that in this regime the assumption of non-correlation might indeed be plausible."
    },
    "0806/0806.4911_arXiv.txt": {
        "abstract": "Until recently, it was considered by many that ground-based photometry could not reach the high cadence sub-mmag regime because of the presence of the atmosphere. Indeed, high frequency atmospheric noises (mainly scintillation) limit the precision that high SNR photometry can reach within small time bins. If one is ready to damage the sampling of his photometric time-series, binning the data (or using longer exposures) allows to get better errors, but the obtained precision will be finally limited by low frequency noises. To observe several times the same planetary eclipse and to fold the photometry with the orbital period is thus generally considered as the only option to get very well sampled and precise eclipse light curve from the ground. Nevertheless, we show here that reaching the sub-mmag sub-min regime for one eclipse is possible with a ground-based instrument. This has important implications  for transiting planets characterization, secondary eclipses measurement and small planets detection from the ground. ",
        "introduction": "Last year, the first transit of a `hot Neptune' was detected \\cite[(Gillon et al. 2007c)]{Gillon07c}. This detection was not obtained with an expensive space instrument but with a commercial CCD camera mounted on a 60cm telescope located in Swiss mountains and mostly devoted to outreach activities. Most of the data were obtained in non-optimal transparency conditions (see Fig.\\,\\ref{fig1}) and are thus far to represent the best photometric quality that can be obtained with commercial equipment. Indeed, some amateur astronomers have demonstrated that they can obtain mmag transit photometry\\footnote{see Bruce Gary's Amateur Exoplanet Archive {\\tt http://brucegary.net/AXA/x.htm}} and they play an important role in the detection and characterization of transiting planets in the context of the {\\tt TransitSearch.org} network \\cite[(Barbieri et al. 2007)]{Barbieri07}  and the XO transit survey \\cite[(Mc Cullough et al. 2006)]{xo2006}.  \\begin{figure}[t] \\begin{center} \\includegraphics[width=4.4in]{stluc.ps} \\caption{Phase-folded GJ~436 OFXB photometry ($black$) from  \\cite[Gillon et al. (2007c)]{Gillon07c}. } \\label{fig1} \\end{center} \\end{figure}  Detecting an eclipse shallower than 1\\% is thus now possible with commercial equipment, but many efforts are undertaken to allow {\\it the detection of eclipses shallower than 1 mmag} with professional ground-based instruments. It is indeed highly desirable to push the precision limit of ground-based photometry towards the sub-mmag regime. While we are presently able to detect and characterize from the ground gazeous giant planets transiting solar-type stars and Neptune-size planets transiting red dwarfs, what we should find below the mmag limit looks very exciting: transits of hot Neptunes around solar-type stars and of Super-Earths around M-dwarfs, secondary eclipse measurements in the visible and near-IR that would nicely complement the $Spitzer$ measurements, very accurate timing measurements allowing to detect Earth-mass planets via the Transit Timing Variations (TTV) method, and much more.  With a depth of 3 mmag, the shallowest eclipse detected so far from the ground is the transit of the core-dominated Saturn-mass planet HD\\,149026b \\cite[(Sato et al. 2005)]{Sato05}. The aim of this contribution can be summarized by a simple question: can we detect eclipses ten times smaller from the ground? ",
        "conclusions": "While the `near-IR' and `big telescope' approaches allow to get  high-precision highly-sampled eclipse photometry for one event observed from the ground but are limited to specific cases, the `composite light curve' approach has a much broader applicability. We are entering a new era of ground-based eclipse photometry, and we bravely predict that the first detection of a sub-mmag eclipse from the ground will be announced in the next future."
    },
    "0806/0806.2530_arXiv.txt": {
        "abstract": "We have conducted optical and X-ray simultaneous observations of \\swift\\ with RXTE and ULTRACAM, while the system persisted in its relatively bright low/hard state. In the cross-correlation function (CCF), we find that the optical leads the X-rays by a few seconds with a broad negative peak, and has a smaller positive peak at positive lags. This is markedly different from what was seen for the similarly interesting system \\xte, and the first time such a correlation function has been so clearly measured. We suggest a physical scenario for its origin. ",
        "introduction": "X-ray and optical emission from astrophysical objects are produced by very different means, with different energetics and time-scales. X-ray binaries are prolific sources of both, with measurable rapid variability. The emission modes cannot be wholly independent, so by comparing their inter-connection, we can learn about the physical conditions from the main emission region: the inner disc around a neutron star or black hole.  Whilst high time resolution ($\\leq1$\\,ms) for X-ray observations has been achieved from the very earliest observations, the photon rates for most sources prevented statistically significant timing work. In the optical domain, exposure times in the ms domain have only become possible relatively recently, with low enough noise and dead-time and high enough quantum efficiency (i.e., CCDs as opposed to photometers) to achieve good signal-to-noise ratios per exposure. The final problem has been simply to schedule simultaneous X-ray and optical observations.  The black hole X-ray Transient (XRT)  \\swift\\ is a system which has been of great interest recently following its outburst episode and detailed observations with the {\\em SWIFT} satellite. First discovered by the {\\em SWIFT}/BAT (Burst Alert Telescope; Palmer et al., 2005) in 2005, pointed $\\gamma$-ray, X-ray, UV, optical and radio observations all detected a new bright source at this location (Morris et al., 2005; Still et al., 2005; Halpern  2005; Fender et al., 2005). Following the early report of a 0.6\\,Hz quasi-periodic oscillation (QPO; Morgan et al. 2005; Ramadevi \\& Seetha, 2005), persistent for some time after the bursting episode, we applied to observe the system simultaneously in X-rays and optical as it faded\\footnote{ESO and RXTE observation IDs 079.D-0535 and 93119-02-02-00, respectively.}. The source was still relatively bright however, especially in the optical, at the time of our project. Following the burst, Cadolle-Bel et al. (2007) measured the spectrum from radio up to 600\\,keV, and Miller et al. (2006) showed spectroscopically that a disc reaching down to small radii was likely.  Further details of our observational campaign, including spectroscopy, longer-term multi-band optical photometry and detailed periodogram analysis are to be published separately in Durant et al. (2008), and an analysis of the long-term R-band variability and orbit-like modulation are presented in Zurita et al (2008). Here, we intend to make the minimum number of processing steps and assumptions to obtain the X-ray/optical CCFs a unique phenomenological hint of black hole accretion physics. ",
        "conclusions": "Very few X-ray/optical CCFs are recorded in the literature, mostly of systems in low states, where the optical lags the X-rays, and is assumed to be the reprocessing signature, possibly from the large inner radius of a truncated accretion disc. This can be used for tomography of the disc and companion by tracking the lag evolution with spectral range and orbital phase (see particularly Hynes, 2005). The dearth of cross-correlations is in good part due to the logistics of arranging simultaneous observations with the few instruments capable, within the short window following an outburst. In this respect, \\swift\\ has been unique, by persisting in its low/hard state (relatively bright) for years after outburst (Zurita et al. 2008). We may find that similar relations exist in other systems, which have gone undetected for technical reasons. Notably, Hynes et al. (2006) report the optical lagging the X-rays in this system by $\\delta t<10$\\,s during outburst, so clearly it was in some different mode during our observations.  Most typically, one expects X-ray flaring to occur in the innermost, hottest regions, and power optical emission by reprocessing of the X-ray flux further out in the optically thick accretion disc, or on the surface of the companion. This picture is a natural consequence of the typical ADAF model geometry, where the disc truncates at a large inner radius, inside of which material is hot and low-density (the ADAF itself), where X-rays are produced. One would expect from this a CCF positive response with a steep rise and slower fall. The positive part of our CCFs do not show this, so simple reprocessing does not dominate over the positive lag region.  Alternatively, X-ray emission may be from a magnetically driven corona around the dense accretion disc (ADC), where reconnection produces energetic particles, and the energy release is dominated by hard X-rays, from the up-scattering of photons by these highly energetic coronal particles. The different energy outputs can thus be coupled and interrelated in a complex manner, e.g., the jet-disc coupling model of Malzac et al. (2004). The corona can form a self-limiting feedback system (see e.g., Uzdensky \\& Goodman 2007, 2008) wherein particles can evaporate from the disc to the corona or condense back to the disc, and the reconnection rate is determined by the density of the corona and magnetic loop movement. Reconnection occurs preferentially in a {\\em marginally collision-less} coronal medium where free magnetic and particle kinetic energies are comparable, which results in a stable equilibrium for a given parameter set (energy transfer rate, sheering etc.).  If the emission is composed of localized micro-flares ({\\em flickering}), each flare might proceed thus: between flares, a patch of the accretion disc cools, and particles from the corona are able to condense; as the coronal density drops, it becomes less collisional and magnetic energy dominated towards the marginally collisionless state, where particle density cannot inhibit magnetic reconnection. At a critical point, cool X-rays flash at the moment of reconnection, decreasing the stored magnetic energy; the disc is heated again, and the coronal particle population replenished by evaporation. Thus, the X-ray emission of the micro-flare is timed after a period of condensation, when the optical emission would be decreasing as more particles are shielded behind cyclotron absorption (indeed, free particles in the corona may also radiate by cyclotron/curvature); immediately afterwards, a population of corona particles and possibly the surface area of the disc patch is larger, so the optical emission is enhanced again. Only a fraction of the released X-rays and accelerated particles heat the disc, most escape or are emitted at higher energies. The self-regulating aspect,, in which the X-ray luminosity acts in the opposite sense and reverses a dip in the optical; magnetic energy driving and disc evaporation/condensation do seem consistent with our CCF. How magnetic field is transmitted from the disc to the corona is not known.  This agrees roughly with the model by Fabian et al. (1982), where emission is controlled by the growing and contraction of optically thick, cyclotron-emitting plasma clouds.  This model fitted well with the early measurements of \\gx\\ (see below). Since the energy emitted by cyclotron depends on the electron number and magnetic energy density, as field is expelled from dense regions, the optical emission would decrease while the energy available for reconnection increases. If the cyclotron emission is predominantly optically thick, then the emitted luminosity in the optical depends on the surface area of these clouds - as they contract and expel magnetic field to the more tenuous, hot medium, the optical emission decreases as the X-ray emission increases, they are naturally anti-correlated.  If the optical emission is not, then, predominantly jet-like, then there is no longer a reason to expect radio emission to correlate with luminosity: Cadolle-Bel et al. (2007) noted (confirmed by Soleri et al. 2007) that the radio emission was unusually low, and that its synchrotron-like spectrum fell below the optical emission. Furthermore, reprocessing of the emitted X-rays will be affected by the acceleration of plasma in magnetic reconnection events. If there is mildly relativistic bulk motion away from the denser matter, then the reprocessing would be weakened and show a different time-response (Beloborodov, 1999).  In the recent work by Liu et al (2007), they show that \\swift\\ in particular, and also \\gx\\ can be modelled as a cool inner disc where thermal conduction and Compton cooling are important in this disc's interaction (condensation and evaporation) with the surrounding low-density corona. They do not specifically consider timescales and driving in their model, but their model is at least consistent with the results here and similar to the picture presented above, and to the dynamic picture of Fabian et al. It will be interesting to see further development of their model.  In our source, the high-energy  emission dominates the total luminosity (up to INTEGRAL energies, see e.g., Cadolle-Bel et al. 2007). The picture is, therefore, of emerging magnetic flux from the disc, releasing its energy in the corona in a self-regulating way. Optical emission is from the dense, hot, magnetically active disc (by cyclotron/synchrotron) and by particles in the corona (by cyclotron and curvature). We believe this scenario accounts qualitatively for what is seen, but is short of a proof.  \\subsection{Comparison with \\xte\\ and \\gx} Two objects that have been extensively studied, including simultaneous X-ray/optical projects, are \\xte\\ and \\gx.  \\xte\\ seems initially very similar in many characteristics to \\swift: persistent low/hard state, high galactic latitude, X-ray spectral and timing characteristics. Kanbach et al. (2001) found that the optical emission lags the X-rays by a small amount ($\\sim$0.5\\,s), but there is an interesting 'precognition dip' in the CCF which is difficult to explain. These features may be qualitiatively similar to our CCF, but in our case, the main feature is a strong anti-correlation, optical before X-ray. The alternative view would be that we see the same precognition dip and response signal, but with very different intensities. Interestingly, they find that the dip is stronger for longer wavelength optical data, which would also fit our data (with the caveat on the quality of the g' band observations above). It was to describe this system that Malzac et al. (2004) developed their jet-disc coupling model.  In the earliest such measurement made that we are aware of, Motch et al (1983) derived the X-ray/optical CCF for \\gx. From a very short simultaneous observation window (96\\,s), they suggested a optical-leading anti-correlation, but only at energies $E<13$\\,keV. It was in this context that the model of Fabian et al. (1982, above) was fairly successful. The CCF was not independently confirmed, but appears similar to our work. Later, Gandhi et al. (2008 in prep.) repeated these measurements over a longer time base-line for the source presumably in quiescence, and found that the CCF similar to \\xte, with the strongest feature a weak positive peak showing a slight lag of optical behind X-rays (by $\\sim$0.2\\,s), but the peak has a markedly different shape with a shallow rise and steep fall, followed by negative correlation in the 1--3\\,s lag range.  These comparisons are suggestive that the CCF we find is symptomatic of the accretion mode in our object, at the time of observation.  \\subsection{Conclusions} Notwithstanding the technical difficulties of constructing X-ray/optical cross-correlations, of the few capable instruments and simultaneous scheduling, this work presents the functions for \\swift\\ which challenge our understanding of the physical processes in the immediate vicinity of a black hole. We find a strong anti-correlation, with the optical preceding the X-rays on time-scales of 1-10\\,s. This demonstrates that there exists a causal link between the optical and X-rays, aside from simple reprocessing, and detailed dynamical modelling will be required to describe the system more fully.  \\medskip\\noindent{\\bf Acknowledgements:} MD and TS are funded by the Spanish Ministry of Science. PG is a Fellow of the Japan Society for the Promotion of Science (JSPS). ULTRACAM was designed and built with funding from PPARC (now STFC), and used as a visiting instrument at ESO Paranal, and RXTE is operated by NASA. Partially funded by the Spanish MEC under the Consolider-Ingenio 2010  Program grant  CSD2006-00070: ``First Science with the GTC''  ({\\tt http://www.iac.es/consolider-ingenio-gtc/})."
    },
    "0806/0806.0703_arXiv.txt": {
        "abstract": "We review the theories and observations of high-mass star formation emphasizing the differences with those of low-mass star formation. We hereafter describe the progress expected to be achieved with \\emph{Herschel}, thanks notably to  Key Programmes dedicated to the earliest phases of high-mass star formation. ",
        "introduction": "High-mass stars, also called OB stars, have luminosities larger than $10^3~\\lsun$, spectral types B3 or earlier, and stellar masses roughly spanning the range $10-100~\\msun$. From their birth to their death, high-mass stars are known to play a major role in the energy budget of galaxies via their radiation, their wind, and the supernovae. Despite that, the formation of high-mass stars remains an enigmatic process, far less understood than that of their low-mass (solar-type) counterparts.  Theoretically, the copious UV flux emitted by a stellar embryo of more than $8~\\msun$ heats and ionizes its parental molecular cloud, leading to the formation and development of a hot core and an \\hii region. These physical and chemical feedback processes are difficult to treat but must be added in  analytical and numerical simulations classically used for low-mass star formation. These further difficulties, summing up with those associated with the star formation process itself (see chapter on low-mass star formation), have long delayed any effort in this direction.  Observationally, the main difficulty arises from the fact that high-mass stars are fewer in number than low-mass stars (see the shape of the IMF in, e.g., Fig.~2 of the chapter on low-mass star formation). Therefore, molecular clouds able to form high-mass stars are statistically more remote (typically at $d_{\\mbox{\\scriptsize Sun}}> 1$~kpc) than those of well-studied low-mass star-forming regions. Current observational studies of high-mass star formation therefore suffers from a lack of spatial resolution and from our bare knowledge of remote star-forming regions.  The past ten years have seen an increasing interest in approaching the issue of the formation of high-mass stars, from both the theoretical and observational sides. Here we review the recent progress made in this domain, and especially in the context of \\emph{Herschel} programmes. ",
        "conclusions": "With the advent of the \\emph{Herschel} satellite and soon after the ALMA interferometer we are entering a very promissing era for the studies of the earliest phases of star formation. Large-scale imaging surveys complemented by high-resolution studies will, without any doubt, revolutionize our knowledge of high-mass star formation."
    },
    "0806/0806.1313_arXiv.txt": {
        "abstract": "In a long-term observing project we have imaged a complete sample of FRII quasars and radio galaxies with $z < 1.0$ at high resolution and high sensitivity with the VLA and MERLIN. This sample of 98 sources includes 15 quasars, 11 broad line radio galaxies and 57 narrow line radio galaxies, allowing unification to be considered in terms of source morphological properties. Radio maps of all the targets have been presented in earlier papers. Here we carry out a systematic analysis of the properties of the jets, cores, lobes and hotspots of objects in the sample. The majority of the tests that we perform show that the data are consistent with a model in which quasars and broad-line radio galaxies are unified with narrow-line objects. Relativistic beaming is the main effect that determines the properties of kiloparsec-scale jets, and it may also have some effect on hotspots. However, some properties of the sample are difficult to account for in simple unified models. ",
        "introduction": "\\label{sec:intro} Fanaroff \\& Riley (1974) type II radio sources (hereafter FRIIs) are powerful sources associated with bipolar outflows that extend great distances from the central engine, remaining highly collimated as they do so. They can be divided into different classes based on features of their optical spectra: radio-loud quasars (Qs), broad line radio galaxies (BLRGs), narrow line radio galaxies (NLRGs) and low excitation radio galaxies (LERGs) can all be FRIIs. A principal defining characteristic is the presence, or absence, of broad line emission, with the Qs and BLRGs having both broad and narrow line emission lines, the NLRGs having narrow line emission only and the LERGs lacking strong high-excitation lines of either type (Hine \\& Longair 1979; Laing et al. 1994).  The current standard unification scheme proposes that the Qs, BLRGs and (at least some of) the NLRGs are intrinsically part of the same population (Scheuer 1987; Barthel 1987, 1989). In this model, the broad emission line region lies close to the very compact central engine and is surrounded by a dusty torus, whereas the narrow line emission region lies further out. Sources that are viewed along or close to the axis of the torus show both broad and narrow line emission -- these are the Qs and BLRGs, which we refer to collectively as broad-line objects -- but the broad line emission region is obscured for those that are oriented closer to the plane of the sky, the NLRGs. Thus differences in the orientation of the source axis to the observer's line of sight are the origin of the three spectral classes. The LERGs lie outside of this scheme; it has been suggested (e.g. Barthel 1994) that LERGs form part of the parent population of BL Lac objects rather than core-dominated Qs and should not show broad line emission at any angle to the line-of-sight, a model consistent with their nuclear properties at other wavebands (Chiaberge, Capetti \\& Celotti 2002; Hardcastle, Evans \\& Croston 2006).  An important detail of the model arises from the fact that the observed luminosity distributions of Qs and BLRGs are not the same. Qs are more powerful and found at higher redshifts (or, equivalently in a flux-limited sample, higher radio luminosities) than the BLRGs; for example, in the 3CR sample (Bennett 1962) Qs are found only with $z \\gtrsim 0.3$, while BLRGs are found with $z \\lesssim 0.3$. It has been suggested that BLRGs may be the low-luminosity equivalents of Qs, or that they lie near the critical angle dividing the quasars and radio galaxies (Barthel 1989; Hardcastle et al. 1998, hereafter H98). While some high-luminosity BLRGs may indeed be intermediate objects, it is clear that at low luminosity, where there are no Qs, BLRGs are the only candidate for the aligned counterpart of the population of low-luminosity NLRGs.  Often FRIIs exhibit a bright linear feature called a jet that extends at least some of the distance between the central feature, the core, and the bright hotspot at the end of the lobe. The jets in most FRIIs are one-sided: either no counterjet is seen or it is much fainter than the feature that is identified as the jet. Relativistic beaming of the jet emission is invoked to explain this asymmetry, as the large scale lobe morphology appears otherwise roughly symmetric. The jet detection rate is higher for Qs and BLRGs than for NLRGs; this can be explained in unified models, since for the broad-line objects the beamed jet is aligned closer to the line of sight and appears brighter. The jet detection rate for LERGs is the highest of all classes (e.g. Mullin, Hardcastle \\& Riley 2006) which may be related to systematic environmental differences between some of the LERGs and the other emission-line types (Hardcastle 2004).  Further evidence in support of relativistic beaming in jets is provided by the Laing-Garrington effect (Laing 1988; Garrington et al. 1988), which is the association of the jet side with the less depolarized lobe. High-resolution multi-frequency observations indicate that the depolarization occurs in an external Faraday screen, so that the less depolarized lobe is expected to be the lobe closer to us; any tendency for the (brighter) jet to be associated with this lobe then implies that beaming is an important factor in jet detection (Scheuer 1987).  While various aspects of the unification and beaming model have been tested and discussed in the literature, there has been little work using complete samples of radio sources free from orientation bias that include sufficient numbers of objects of all spectral classes to give statistically significant results. Good quality observations of such a sample, with both high resolution and sensitivity, are therefore vital, and this has been the rationale behind a long-term observational project in which we have mapped a complete sample of the brightest FRII radio sources with $z<1.0$. The sample, which is defined in section \\ref{sec:data}, includes 98 sources. Maps of these have been presented in a series of papers: Black et al. (1992), Leahy et al. (1997), Hardcastle et al. (1997), Gilbert et al. (2004) and Mullin, Hardcastle \\& Riley (2006). These maps are available online\\footnote{See http://zl1.extragalactic.info/} along with a database of all measurements analysed and discussed in this paper. The sample includes 15 Qs, 11 BLRGs and 57 NLRGs, thus enabling aspects of unification to be tested along with trends in source properties over the wide redshift and luminosity range spanned by the data. In section \\ref{sec:data} we also define a number of morphological and flux parameters corresponding to the observed source properties and describe our measurement methods.  We examine the properties of the lobes, cores and jets and hotspots in sections \\ref{sec:lobes}, \\ref{sec:cores_jets} and \\ref{sec:hotspots} respectively. For each feature, observational effects are considered as well as trends across the power, redshift and source size range of the sample and we quantify these where appropriate with statistical tests. The significance of linear correlation is tested for using Spearman's rank correlation coefficient. For the core and jet prominence data, however, this is not possible as only upper limits on these parameters are available for some sources: in statistical terminology, the data are censored. Instead, a modified Kendall's $\\tau$ rank correlation coefficient as implemented in the survival-analysis package {\\sc asurv} (LaValley, Isobe \\& Feigelson 1992) is used for these data. The Kolmogorov-Smirnov (hereafter K-S) test determines if it is the case that the culmulative distribution function of two samples differ and is used to address the question of whether some property of two subsamples of the data (that is, subsamples defined by power, redshift and size cutoffs or by spectral class) differ significantly. It is sensitive to differences in both location and shape of the functions. No modification of the K-S test to take account of censoring is available to us, and so we do not use it in situations where censoring is important. The Wilcoxon-Mann-Whitney (hereafter W-M-W) test is also used to determine if two defined subsamples differ, but in this case the null hypothesis tested is that the probability of an observation of one population exceeding an observation from the second is 0.5. Thus the W-M-W test is used to determine whether there is a significant difference in the magnitude of the quantity of interest between the two subsamples, that is, if one dataset has significantly smaller or larger values than the other. In order to treat censored data correctly when testing for such differences, a generalized W-M-W test is used, the Peto-Prentice test, which is implemented in {\\sc asurv}. Finally, the binomial test is used to determine the statistical significance of any correlation with jet or longer lobe side for a number of properties. The significance of all test results is discussed in the text and the results are tabulated. We take a result to be significant enough to be discussed if the null hypothesis is rejected at better than the 95 per cent confidence level.  The interpretation of the observed properties of our sample sources, and the evidence for and against unified models, is discussed in section \\ref{sec:discussion}. The quantitative implications of our results for beaming in the cores and jets of powerful radio galaxies will be discussed in a separate paper.  The spectral index, $\\alpha$, is defined throughout the paper in the sense that $S = v^{-\\alpha}$ (where $S$ is the flux and $v$ denotes the frequency) and  we assume that $H_{0} = 70\\ {\\rm kms^{-1}Mpc^{-1}}$, $\\Omega_{\\rm m}=0.3$ and $\\Omega_{\\Lambda}=0.7$. ",
        "conclusions": ""
    },
    "0806/0806.4536_arXiv.txt": {
        "abstract": "We present the results of numerical simulations of continuum-driven winds of stars that exceed the Eddington limit and compare these against predictions from earlier analytical solutions. Our models are based on the assumption that the stellar atmosphere consists of clumped matter, where the individual clumps have a much larger optical thickness than the matter between the clumps. This `porosity' of the stellar atmosphere reduces the coupling between radiation and matter, since photons tend to escape through the more tenuous gas between the clumps. This allows a star that formally exceeds the Eddington limit to remain stable, yet produce a steady outflow from the region where the clumps become optically thin. We have made a parameter study of wind models for a variety of input conditions in order to explore the properties of continuum-driven winds.  The results show that the numerical simulations reproduce quite closely the analytical scalings. The mass loss rates produced in our models are much larger than can be achieved by line driving. This makes continuum driving a good mechanism to explain the large mass loss and flow speeds of giant outbursts, as observed in {$\\eta$~Carinae} and other luminous blue variable (LBV) stars. Continuum driving may also be important in population III stars, since line driving becomes ineffective at low metalicities. We also explore the effect of photon tiring and the limits it places on the wind parameters. ",
        "introduction": "Massive, hot stars continuously lose mass through radiation driving. The most commonly explored mechanism is line driving, wherein the scattering of photons by ions in the stellar atmosphere transfers momentum from the radiation field to the gas. This mechanism results in a quiescent mass loss with mass loss rates ranging up to about $10^{-4}~\\mso$~yr$^{-1}$ \\citep{so06}, which can explain the winds of most massive, hot stars. In fact, for most stars, the observed mass loss rates are considerably less \\citep{vk05}.  However, some stars, most notably luminous blue variable stars (LBVs) such as {$\\eta$~Carinae}, experience  outbursts with mass loss rates several orders of magnitude higher than can be explained through line driving \\citep{dh97, s02, om08}. In the case of $\\eta$~Carinae, the 1840's outburst is inferred to resulted in the ejection of ca. 10-20~$M_{\\odot}$ over a time lasting several years, even up to a decade \\citep{dh97}. While short compared to evolutionary timescale of millions of years, this is much longer than the typical dynamical timescale of hours, characterized by either free-fall time  or interior sound travel time across a stellar radius. Thus in contrast to SN ``explosions'' that are effectively driven by the overpressure of superheated gas in the deep interior, explaining LBV outbursts requires a more sustained mechanism that can drive a quasi-steady mass loss, a stellar wind, from near the stellar surface. In the case of $\\eta$~Carinae the outburst was accompanied by a strong increase in radiative luminosity, very likely making it well above the Eddington limit for which {\\em continuum} driving by just electron scattering would exceed the stellar gravity. The reason for this extended increase in luminosity is not yet understood, and likely involves interior processes beyond the scope of this paper. Instead, the focus here is on the way such continuum driving can result in a sustained mass loss that greatly exceeds what is possible through line opacity.  A key feature of continuum driving is that, unlike line driving, it does not become saturated from self-absorption effects in a very dense, optically thick region. Indeed, since both continuum acceleration and gravity scale with the inverse square of the radius, a star that exceeds the Eddington limit formally becomes gravitationally unbound not only at the surface but throughout. Clearly, this is in contradiction to the steady surface wind mass loss observed for these stars. N.B. Contrary to what is sometimes claimed, this would not automatically destroy the entire star. Although radiative acceleretion might overcome gravity locally, the total energy in the radiation field would not  suffice to drive the entire envelope of the star to infinity (this is known as `photon tiring'). Instead, the outward motion of the gas would quickly stagnate and matter would start to fall back. Nevertheless, the net result would not resemble a steady wind.  This problem can be resolved by assuming that the stellar material is clumped rather than homogeneous, with the individual clumps being optically thick - and therefore self-shielding from the radiation - whereas the medium in between the clumps is relatively tranparent to radiation. This so called `porosity effect' can lead to a reduced coupling between matter and radiation \\citep{s98,s00}. The photons tend to escape through the optically thin material between the clumps without interacting with the matter inside the clumps. This implies that a star that formally exceeds the Eddington limit can remain gravitationally bound and would only exceed the effective Eddington limit at the radius where the individual clumps themselves become optically thin.  The structure of such a star should therefore look as follows \\citep{s01}. Deep inside the super-Eddington star, convection is necessarily excited \\citep{jso73} such that the radiation field remains sub-Eddington through most of the stellar interior. At low enough densities where maximally efficient convection cannot sufficiently reduce the radiative flux, the near Eddington luminosity necessarily excites at least one of several possible instabilities \\citep[e.g.,]{a92,s01a} which give rise to a reduced opacity. This `porous'  layer has a reduced effective opacity and an increased effective Eddington luminosity. Thus, the layer remains gravitationally bound to the star. At lower densities still, the dense clumps become optically thin and the effective opacity approaches its microscopic value. From this radius outwards, the matter is gravitationally unbound, and is part of a continuum-driven wind.  A detailed analytical study of this paradigm was carried out by {Owocki, Gayley \\& Shaviv (2004)}, hereafter \\citet{ogs04}. This predicted that continuum-driven winds can produce high mass loss rates ($\\geq10^{-3}~\\msoy$) at intermediate wind velocities ($10^2-10^3~\\kms$). Here we test these analytical predictions with numerical simulations of winds from super-Eddington stars.  In addition to LBVs, which are the specific objects we study here, other types of astronomical objects can exceed the Eddington limit and therefore experience similar continuum-driven winds. These include for example classical novae \\citep{s01} or high accretion rate accretion disks around black holes \\citep{b06}. In fact, classical nova eruptions clearly exhibit steady continuum-driven winds, indicating that the mass loss rate is somehow being regulated and likely to be described by the porosity model and the same continuum-driven wind analyzed here.   The layout of this paper is as follows. In  \\S\\ref{sec-subEdd} we show the effect of continuum scattering on a sub-Eddington, line-driven wind. In  \\S\\ref{sec-analytic} we summarize the analytic results obtained by \\citet{ogs04}. In \\S\\ref{sec-numeric} we describe the numerical methods that we have used for our simulations. \\S\\ref{sec-result} shows the results of our simulations and the comparison with the analytical predictions. In \\S\\ref{sec-phtir} we discuss the effect of photon tiring and show how it influences the results of our calculations. Finally, in \\S\\ref{sec-disc} we end with a summary and a discussion. \\begin{figure*} \\centering \\resizebox{\\hsize}{!}{\\includegraphics[width=\\textwidth,angle=-90]{f01.eps}} \\caption{ Left: The effect of continuum driving on a sub-Eddington, line-driven wind for $\\Gamrad=0.1$. The solid lines depict the wind parameters for the CAK force only, the dashed lines show the joint effect of both CAK line driving and continuum scattering. Clearly, the influence of the continuum term is negligeable. Right: The same variables , but for $\\Gamrad=0.5$. Here the difference caused by continuum driving is significantly more pronounced.} \\label{fig:line1} \\end{figure*} ",
        "conclusions": "\\label{sec-disc} A series of numerical simulations was carried out in order to test the analytical approximations for the porosity length formalism of continuum driving, published by \\cite{ogs04}. The numerical results coincide well with the analytical results and demonstrate that this mechanism allows for powerful radiation-driven winds. This effect can explain the mass loss rates and wind velocities observed in Luminous Blue Variables such as {$\\eta$~Carinae}.  The simulations also confirm that the photon tiring effect plays an important role in continuum driving as it places an upper limit on the mass loss rate. The effects of actually crossing the photon tiring limit have not yet been explored. This situation is much more complicated, since the simulations will no longer be able to reach a steady state solution. Ideally, such simulations should be done in two, or even three dimensions, to investigate the effect of interactions between different layers of the stellar wind as they move back and forth. Note that at this point we cannot predict how the star itself would react to such a situation. All our simulations have been done under the assumption that stellar parameters do not change significantly over time. It is possible that conditions in the outer layers of the star would change to reduce the driving force.  Porosity reduced continuum driving can also be important for the winds of other super-Eddington objects such as Novae \\citep{s01}, accretion disks \\citep{b06} and transients like {M85OT2006-1} \\citep{ketal07}.  In a companion paper, \\citep{mos08b}, we explore the situation where the star exceeds the photon tiring limit. We also intend eventually to generalize the simulations to multiple dimensions and explore the influence of stellar rotation on continuum-driven winds."
    },
    "0806/0806.2126_arXiv.txt": {
        "abstract": "The detection of Earth-skimming tau neutrinos has turned into a very promising strategy for the observation of ultrahigh-energy cosmic neutrinos. The sensitivity of this channel crucially depends on the parameters of the propagation of the tau neutrinos through the terrestrial crust, which governs the flux of emerging tau leptons that can be detected. One of the characteristics of this propagation is the possibility of regeneration through multiple $\\nu_\\tau \\leftrightarrow \\tau$ conversions, which are often neglected in the standard picture. In this paper, we solve the transport equations governing the $\\nu_\\tau$ propagation and compare the flux of emerging tau leptons obtained allowing regeneration or not. We discuss the validity of the approximation of neglecting the $\\nu_\\tau$ regeneration using different scenarios for the neutrino-nucleon cross sections and the tau energy losses. ",
        "introduction": "\\label{sec:intro}  With the advent of a new generation of large-scale detectors of cosmic radiation, the observation of high-energy cosmic neutrinos produced in distant astrophysical sites or possibly by other, more exotic, mechanisms has become one of the major challenges of astroparticle physics. In both astrophysical and exotic models, substantial fluxes of electron and muon neutrinos are expected from the disintegration of charged pions (and kaons) produced in the interaction of accelerated particles with ambient matter and radiation, either at the source location or between the source and the observer. Given the large distances traveled by these cosmic neutrinos, approximately equal fluxes in $\\nu_e , \\nu_\\mu$ and $\\nu_\\tau$ are expected on Earth as a result of flavor mixing and oscillations~\\cite{nuoscil,Learned:1994wg}. Important efforts are ongoing to build dedicated neutrino telescopes both in the Southern~\\cite{Andres:1999hm,ice3} and Northern~\\cite{antares,nemo,nestor,Aynutdinov:2005sc} hemisphere, opening a window on the neutrino sky in the energy range $10^{-6}\\ \\mathrm{EeV} \\leq E_\\nu \\leq 10^{-1}\\ \\mathrm{EeV}$. At even higher energies, other promising experiments are developing the detection of coherent radio emission produced by neutrino-induced showers in matter~\\cite{anita,rice,forte,glue,hankins,james,swarup}.  Interestingly, recent studies have shown that the new generation of ultrahigh-energy (UHE) cosmic ray detectors such as the Pierre Auger Observatory~\\cite{pao} and the HiRes Fly's Eye detector~\\cite{hires} have a comparable detection potential for UHE neutrinos in the range of energies $10^{-1}\\,\\, \\mathrm{EeV} \\leq E_\\nu \\leq 10^{2}\\,\\, \\mathrm{EeV}$, where neutrinos are expected to be produced in the interaction of UHE cosmic rays with the cosmic microwave background~\\cite{cosmogenic}. It is long known that downward-going showers induced by neutrinos that penetrate deep in the atmosphere can in principle be identified at large zenith angles ($\\theta > 75^\\circ$) where no background is expected from hadronic primaries~\\cite{beresmirnov,capelle}. It was also pointed out more recently that the presence of $\\nu_\\tau$ in the cosmic neutrino flux provides another promising channel of detection for air shower detectors~\\cite{antoine,fargion}. Upward-going UHE tau neutrinos that graze the Earth just below the horizon (often referred to as ``Earth-skimming neutrinos'') are indeed likely to interact in the crust and produce a tau lepton which may emerge and initiate an observable air shower, provided it does not decay too far from the detector.  The sensitivity to such UHE Earth-skimming neutrinos crucially depends on the conditions of the $\\nu_\\tau$ propagation through the terrestrial crust, and on the correct estimation of the flux of $\\tau$ leptons that emerge from the Earth. This propagation problem has been widely discussed in different contexts and with different approximations~\\cite{DRSS,bertou,feng,bottai,tseng,montaruli,aramo}. An exhaustive treatment should account for $\\tau$ and $\\nu_\\tau$ neutral-current (NC) and charged-current (CC) interactions with nucleons, $\\tau$ decay and energy losses. However the full coupled transport equations admit no analytical solution, and even in the case of Monte Carlo calculations, simplifications are usually made such as dropping the $\\tau$ weak interactions and neglecting multiple regenerations of the $\\nu_\\tau$.  Such approximations are defendable in the standard case where the characteristic lengths for the $\\tau$ CC interaction ($\\sim 600$ km, at 1 EeV) and for the $\\tau$ decay ($\\sim 50$ km) are larger than that for energy losses ($\\sim 6$ km). However, they might be challenged in other, more exotic scenarios. The knowledge of the neutrino cross section and the tau energy losses in the EeV energy range is indeed limited~\\cite{sarkar, parente}, and these could be significantly affected at center-of-mass energies beyond the TeV by the onset of new physics beyond the standard model. Several studies have even suggested that the comparison of the flux rates between down-going and Earth-skimming neutrinos could actually help constraining the neutrino properties at ultrahigh energies, where no direct measurements exist~\\cite{irimia,anchordoqui}.  All these considerations pinpoint the necessity of an accurate determination of the flux of emerging $\\tau$ leptons. In this context, it is important to correctly describe all the contributions to the $\\tau$ flux and to assess the impact of simplifications in the description of the propagation problem.  The present paper therefore focuses on understanding the effects of the $\\nu_\\tau$ and $\\tau$ regeneration while skimming the Earth on the flux of emerging $\\tau$ leptons in different scenarios for the neutrino-nucleon cross section and for the $\\tau$ energy losses. In Sec. \\ref{sec:calcul}, we present the transport equations for the $\\nu$ and $\\tau$ propagation, and the scenarios to be studied. In Sec. \\ref{sec:results}, we compare the flux of emerging taus for different conditions of propagation. Finally, we present our conclusions in Sec. \\ref{sec:conclusion}. ",
        "conclusions": "\\label{sec:conclusion}  In this paper, we have studied the mechanism of regeneration of the $\\nu_\\tau$ flux while crossing the Earth, to investigate its effect on the flux of emerging $\\tau$ leptons. Assuming a detector with an energy threshold of 0.1 EeV and sensitive only to taus with an emerging angle below 3$^\\circ$, the effect is negligible for a flux of incident neutrinos $dN/dE \\propto E_\\nu^{-2}$ with the standard values of cross sections and energy losses. But we have shown that this is not valid for other assumptions on the detector performance, neither for less standard values of the weak cross section or tau energy losses. On the one hand, neglecting the regeneration leads to about 30$\\%$ underestimation on the integrated flux for a detector sensitive up to 15$^\\circ$. On the other hand, the underestimation of the $\\tau$'s emerging from the Earth with an energy greater than 0.1 EeV and an emerging angle below 3$^\\circ$ may reach 70$\\%$ for the different values of $\\sigma_{\\nu_\\tau}$ and $\\beta_\\tau$ studied in this paper. Moreover, the error made in the computation of the flux of emerging taus can be even larger for a harder flux of incident neutrinos since the contribution of the regeneration effect increases with the energy of the incident neutrino.  The simplification of neglecting the regeneration is thus only safe for particular values of the physical properties playing a role on the propagation and specific detectors. It may lead to a significant underestimation of the flux of emerging $\\tau$'s when looking at nonstandard values of the weak cross section or tau energy losses. Therefore, it should be carefully treated and accounted for when studying the systematics due to the uncertainties on those properties or while using the Earth-skimming technique to test for instance higher weak cross-sections. Similarly, one should carefully check the effect for the characteristics of the actual detector before neglecting the regeneration."
    },
    "0806/0806.2060_arXiv.txt": {
        "abstract": "{The main observational results from radio continuum and polarization observations about the magnetic field strength and large-scale pattern for face-on and edge-on spiral galaxies are summarized and compared within our sample of galaxies of different morphological types, inclinations, and star formation rates (SFR). We found that galaxies with low SFR have higher thermal fractions/ smaller synchrotron fractions than those with normal or high SFR. Adopting an equipartition model, we conclude that the nonthermal radio emission and the \\emph{total magnetic field} strength grow nonlinearly with SFR, while the \\emph{regular} magnetic field \\emph{strength} does not seem to depend on SFR.\\  We also studied the magnetic field structure and disk thicknesses in highly inclined (edge-on) galaxies. We found in four galaxies that -despite their different radio appearance- the vertical scale heights for both, the thin and thick disk/halo, are about equal (0.3/1.8~kpc at 4.75~GHz), independently of their different SFR. This implies that all these galaxies host a galactic wind, in which the bulk velocity of the cosmic rays (CR) is determined by the total field strength within the galactic disk. The galaxies in our sample also show a similar large-scale magnetic field configuration, parallel to the midplane and X-shaped further away from the disk plane, independent of Hubble type and SFR in the disk. Hence we conclude that also the \\emph{large-scale} magnetic field \\emph{pattern} does not depend on the amount of SFR.}  \\resumen{}    \\addkeyword{galaxies: magnetic fields} \\addkeyword{galaxies: spiral} \\addkeyword{galaxies: individual (NGC~891)} \\addkeyword{galaxies: individual (NGC~4631)} \\addkeyword{galaxies: halos} \\addkeyword{polarization}  \\begin{document} ",
        "introduction": "\\label{sec:intro} Radio observations of the continuum emission in the cm-wavelength regime are the best way to study magnetic fields in galaxies. The total intensity of the synchrotron emission gives the strength of the total magnetic field. The linearly polarized intensity reveals the strength and the structure of the resolved regular field in the sky plane (i.e. perpendicular to the line of sight).  However, the observed polarization vectors suffer Faraday rotation and depolarization on their way from the radiation's origin to us. Correction for Faraday rotation is possible with observations at two -- or better more -- wavelengths by determining the rotation measure RM (being proportional to $\\int n_{\\rm e} B_{\\parallel} dl$). $B_{\\parallel}$ is the coherent magnetic field parallel to the line of sight, and its sign gives the direction of this magnetic field component. Both field components, parallel and perpendicular to the line of sight, enable us in principle to deduce a 3-dimensional picture of the large-scale magnetic field.\\ Note, however, that the polarized intensity is only sensitive to the field orientation, i.e. it does not distinguish between parallel and antiparallel field directions in the plane of the sky, whereas the RM is large for parallel fields along the line of sight, but zero for parallel and antiparallel fields (of equal strength).\\ Magnetic fields consist of regular and turbulent components. The total magnetic field strength in a galaxy can be estimated from the nonthermal radio emission under the assumption of equipartition between the energies of the magnetic field and the relativistic particles (the so-called {\\em energy equipartition}) as described in Beck \\& Krause (2005).\\ ",
        "conclusions": ""
    },
    "0806/0806.1006_arXiv.txt": {
        "abstract": "{VO-Neural is the natural evolution of the Astroneural project which was started in 1994 with the aim to implement a suite of neural tools for data mining in astronomical massive data sets. At a difference with its ancestor, which was implemented under Matlab, VO-Neural is written in C++, object oriented, and it is specifically tailored to work in distributed computing architectures. We discuss the current status of implementation of VO-Neural, present an application to the classification of Active Galactic Nuclei, and outline the ongoing work to improve the functionalities of the package. ",
        "introduction": "One of the main goals of the International Virtual Observatory (VOb) is the federation under common standards of all astronomical archives available worldwide \\cite{IVOA}. Once this meta-archive will be completed, its exploitation will allow a new type of multi-wavelenght, multi-epoch science which can only be barely imagined \\cite{george_1}, but will also pose unprecedented computing problems. From a mathematical point of view, in fact, most of the operations performed by the astronomers during their every-day life can be reconduced (either consciously or unconsciously) to standard data mining tasks such as, for instance, clustering, classification, pattern recognition and trend analysis. All these tasks scale very badly when either the number $N$ of records to be processed or the number $D$ of features characterizing each record, increase: \\begin{itemize} \\item clustering scales as $\\sim N \\times \\log N \\times N^2$, and as $\\sim D^2$; \\item search for correlations scales as $\\sim N \\times \\log N \\times N^2$, and as $\\sim D^k$ with  $k \\geq 1$; \\item bayesian or likelihood  algorithms scale as $\\sim N^m$ with $m\\geq 3$ and as$\\sim D^k$ with  $k \\geq 1$. \\end{itemize} To get an idea of the computational demands posed by the VOb we shall just notice that a modern digital survey can easily produce datasets having $N\\sim 10^{9}$ and $D\\gg 10^2$ and leave to the reader to imagine what could be the demands of a multiwavelenght, multi-epoch survey. It is apparent that the extraction of knowledge from such data sets cannot be performed with traditional SW\\cite{astroweka} \\& HW, and requires some form of high performance computing (HPC). The traditional HPC approach based on parallel multi-CPU software running on dedicated clusters, is however against the very same phylosophy of the VOb which aims at opening the exploitation of its data archives also to scientists who do not have access to large HPC centers. In this respect, the GRID seems to offer the most natural and democratic answer since, at least in theory, it allows any user possessing a personal certificate to access the distributed computing resources. The VOb, however, for the same fact of being open to use by the community at large, does not match the security requirements of the GRID and this limitation strongly undermines its effectiveness.  In \\cite{deniskina2008} we discuss the first version of $GRID-Launcher$, a tool which interfaces the UK-ASTROGRID \\cite{astrogrid} with the GRID-SCOPE \\cite{scope}. In this contribution we discuss instead the structure of the data mining package VO-Neural \\cite{voneural} which is specifically designed to perform complex data mining (DM) tasks on astronomical (but not only) massive data sets (MDS). As an exemplification, in Sect.\\ref{applications} we also show how the methods so far implemented can be used to address the challenging task of obtaining an objective classification of Active Galactic Nuclei (AGN). Finally, in the last Section we shortly outline some ongoing and planned developments. ",
        "conclusions": ""
    },
    "0806/0806.1189_arXiv.txt": {
        "abstract": "{Magnetic fields are proposed to play an important role in the formation and support of self-gravitating clouds and the formation and evolution of protostars in such clouds.} {We attempt to understand more precisely how the Pipe nebula is affected by the magnetic field.} {We use $R$-band linear polarimetry collected for about 12\\,000 stars in 46 fields with lines of sight toward the Pipe nebula to investigate the properties of the polarization across this dark cloud complex.} {Mean polarization vectors show that the magnetic field is locally perpendicular to the large filamentary structure of the Pipe nebula (the `stem'), indicating that the global collapse may have been driven by ambipolar diffusion. The polarization properties clearly change along the Pipe nebula. The northwestern end of the nebula (B59 region) is found to have a low degree of polarization and high dispersion in polarization position angle, while at the other extreme of the cloud (the `bowl') we found mean degrees of polarization as high as $\\approx$15\\% and a low dispersion in polarization position angle. The plane of the sky magnetic field strength was estimated to vary from about 17\\,$\\mu$G in the B59 region to about 65\\,$\\mu$G in the bowl.} {We propose that three distinct regions exist, which may be related to different evolutionary stages of the cloud; this idea is supported by both the polarization properties across the Pipe and the estimated mass-to-flux ratio that varies between approximately super-critical toward the B59 region and sub-critical inside the bowl. The three regions that we identify are: the B59 region, which is currently forming stars; the stem, which appears to be at an earlier stage of star formation where material has been through a collapsing phase but not yet given birth to stars; and the bowl, which represents the earliest stage of the cloud in which the collapsing phase and cloud fragmentation has already started.} ",
        "introduction": "\\label{int}  Understanding the role that magnetic fields play in the evolution of interstellar molecular clouds is one of the outstanding challenges of modern astrophysics. One problem related to star formation concerns the competition between magnetic and turbulent forces. The prevailing scenario of how stars form is quasi-static evolution of a strongly magnetized core into a protostar following influence between gravitational and magnetic forces. By ambipolar diffusion, i.e., the drift of neutral matter with respect to plasma and magnetic field, gravity finds a way to overcome magnetic pressure and eventually win the battle \\citep[e.g., ][]{MS56, Na79, MP81, LS89}.  However, doubts about the validity of this theory were expressed because of the apparent inconsistency between the expected and inferred lifetimes of molecular clouds. This inconsistency inspired some researchers to propose a new theory in which star formation is driven by turbulent supersonic flows in the interstellar medium. Magnetic fields may be present in this theory, but they are too weak to be energetically important \\citep[e.g. ][]{ES04,MK04}. It must be noted, however, that some results \\citep{TM04,MTK06} demonstrate that the ambipolar--diffusion--controlled star formation theory is not in contradiction with molecular cloud lifetimes and star formation timescales.  \\begin{figure*}[ht] \\sidecaption \\resizebox{12cm}{!}{\\includegraphics{10091fig1.eps}} \\caption[]{Mean polarization vectors, for each of the observed 46 fields, overplotted on the dust extinction map of the Pipe nebula obtained by \\citet{LAL06}. The lengths of these vectors are proportional to the scale indicated in the top left-hand corner. Only stars showing ${\\rm P}/\\sigma_P \\ge 10$ were used in the calculus of the mean polarization and position angle. The dashed-lines indicate the celestial meridians defined by $17^{\\rm h} 14^{\\rm m} 30\\fs0$ and $17^{\\rm h} 27^{\\rm m} 40\\fs0$ (see text and Fig.\\,\\ref{fig2}).} \\label{fig1} \\end{figure*}  Previous optical polarimetric observations toward well-known forming molecular clouds have enabled the large-scale magnetic field associated with these regions to be studied \\citep[e.g.][]{GBM90}. In this work, we introduce the general results of a polarimetric survey conducted for the Pipe nebula, a nearby \\citep[130--160\\,pc,][]{LAL06, AF07} and massive ($10^4\\,{\\rm M}_{\\sun}$) dark cloud complex that appears to provide a suitable laboratory for investigating magneto-turbulent phenomena. The Pipe nebula exhibits little evidence of star formation activity despite having an appropriate mass. Until now, the only confirmed star-forming region in this nebula was B59 \\citep{BHB07}, an irregularly-shaped dark cloud located at the northwestern end of the large filamentary structure that extends from ($l,b$) $\\approx$ ($0\\degr, 4\\degr$) to ($l,b$) $\\approx$ ($357\\degr, 7\\degr$). This apparently low efficiency in forming stars may be an indication of youth. \\citet{ALL07} identified, in this cloud, 159 cores of effective diameters between 0.1 and 0.4\\,pc, and estimated masses ranging from 0.5 to 28 M$_{\\sun}$, supposedly in a very early stage of development. A further investigation of these cores \\citep{LMRAL} discovered that most of them appeared to be pressure confined and in equilibrium with the surrounding environment, and that the most massive ($\\ga2$~M$_{\\sun}$) cores were gravitationally bound. They suggested that the measured dispersion in internal core pressure of about a factor of 2--3 could be caused by either local variations in the external pressure, or the presence of internal static magnetic fields with strengths of less than 16 \\,$\\mu$G, or a combination of both. The results derived from our optical polarimetric observations indicate that the magnetic field probably plays a far more important role in the Pipe nebula. ",
        "conclusions": "We have described the global polarimetric properties of the Pipe nebula as an increasing polarization degree along the filamentary structure from B59 towards the bowl, while the dispersion in polarization angles decreases along this way. Our results appears to indicate that there exist three regions in the Pipe nebula of distinct evolutionary stages: since the mean orientation angle of the mean polarization vectors is perpendicular to the longer axis of the cloud, this implies that the cloud collapse is taking place along the magnetic field lines. We can subdivide the Pipe nebula into the following components:  \\begin{itemize} \\item B59, the only active star-forming site in the cloud. For the observed fields, we measure a large dispersion in polarization angle and low polarization degree. \\item The stem, which collapsed by means of ambipolar diffusion but has not yet given birth to stars. It appears to represent a transient evolutionary state between B59 and the bowl. \\item The bowl, which contains the fields of the highest values of mean polarization and the lowest values of dispersion in polarization angle. These values imply that the dust grains in the bowl are highly aligned by a rather strong magnetic field. For this reason, the bowl may represent the start of the contraction phase during a very early evolutionary stage. \\end{itemize}"
    },
    "0806/0806.1140_arXiv.txt": {
        "abstract": "s{ The Tully-Fisher relation is a correlation between the luminosity and the HI 21cm line width in spiral galaxies (LLW relation). It is used to derive galaxy distances in the interval 7 to 100 $Mpc$. Closer, the Cepheids, TRGB and Surface Brightness Fluctuation methods give a better accuracy. Further, the SNIa are luminous objects still available for distance measurement purposes, though with a dramatically lower density grid of measurements on the sky. Galaxies in clusters are all at the same distance from the observer. Thus the distance of the cluster derived from a large number of galaxies (N)  has an error  reduced according to $\\sqrt{N}$. However, not all galaxies in a cluster are suitable for the LLW measurement. The selection criteria we use are explained hereafter; the important point being to avoid Malmquist bias and to not introduce any systematics in the distance measurement.} ",
        "introduction": " ",
        "conclusions": ""
    },
    "0806/0806.3974_arXiv.txt": {
        "abstract": "We develop and implement an observational test of the theoretical notion that dissipation in major mergers of gas-rich galaxies produces the fundamental plane (FP) and related correlations obeyed by ellipticals.  Observations have shown that the ``tilt'' of the FP involves more than a simple non-homology or stellar population effect: lower-mass ellipticals have a higher ratio of stellar to dark matter within their stellar effective radii.  Theoretical models have attempted to explain this via dissipation: if ellipticals are formed in major mergers of disks, then mergers between disks having a larger gas content (typically observed to be lower-mass disks) will yield remnants with a larger mass fraction formed in a central, compact starburst, giving a smaller stellar $R_{e}$ and lower $M_{\\rm tot}/ \\mstar$ within that $R_{e}$.  Such starbursts leave a characteristic imprint in the surface brightness profiles of ellipticals, in the form of a central excess above the outer profile established by the dissipationless, violent relaxation of disk stars.  In previous work, we implemented a purely empirical method to use such features in the observed profiles of ellipticals to robustly estimate the amount of dissipation involved in the original spheroid-forming merger. Applying this to a large sample of ellipticals with detailed kinematic and photometric observations, we demonstrate that the location of ellipticals on the FP and its tilt are in fact driven by dissipation. We show that at fixed mass, ellipticals formed in more dissipational events, as indicated by their observed profiles, are smaller and have a lower ratio $M_{\\rm tot}/\\mstar$. {\\em At the same (fixed) degree of dissipation, there is no tilt in the FP} -- i.e.\\ ellipticals formed with a similar level of dissipation have the same ratio of enclosed stellar to total mass within $R_{e}$.  We further demonstrate that observations and these models obey the ``homology assumption,'' i.e.\\ that the true enclosed mass $\\mtrue(R_{e})\\propto \\sigma^{2}\\,R_{e}$.  Measured at the radii of disks of the same mass, we show that ellipticals have the same total enclosed masses as those disks -- i.e.\\ that the FP tilt can be effectively removed. Therefore, the fundamental plane tilt cannot primarily owe to non-homology or to changes in the dark matter distribution: it {\\em must} arise as a result of a contraction of the baryonic component relative to the dark matter in the process that transforms disks to ellipticals, as predicted by dissipational mergers. If we allow for the observed cosmological dependence of disk gas fraction on mass, the observed FP, size-mass, and velocity dispersion-mass correlations are reproduced by our models, as are the observed homology constraints and profile shapes. {\\em Dissipation is both necessary and sufficient to explain the observed FP correlations of ellipticals.}  These observations all favor theories in which ellipticals are formed in major mergers of disks with gas fractions, sizes, and dark matter content similar to that observed as a function of mass in low-redshift disks: unusually compact disks are {\\em not} required to make $\\sim0.01-10\\,L_{\\ast}$ ellipticals.  We present a number of associated predictions that can be used to further test these assertions. ",
        "introduction": "\\label{sec:intro}  Understanding the scaling relations between the photometric and kinematic properties of galaxy spheroids -- their masses, sizes, velocity dispersions, and luminosities -- is fundamental to explaining the origin of early-type galaxies.  \\citet{fj76} demonstrated that ellipticals obey a relatively tight correlation between optical luminosity and central velocity dispersion, and \\citet{kormendy77:correlations} found an analogous relationship between their effective surface brightness and radii. With improved observations and the advent of stellar population modeling, these observed trends can be translated into robust correlations between physical parameters: a velocity dispersion-stellar mass ($\\sigma-\\mstar$) and a size-stellar mass ($R_{e}-\\mstar$) relation.  \\citet{dd87:fp} and \\citet{dressler87:fp} demonstrated that the scatter in either the \\citet{fj76} or \\citet{kormendy77:correlations} relation could be reduced by adopting a three parameter correlation of the form $\\log{(R_{e})}= a\\,\\log{(\\sigma)} - 0.4\\,b\\,\\mu_{e} + c$ (equivalently $R_{e}\\propto\\sigma^{a}\\,I_{e}^{b}$), with best fit scalings $a\\sim1.3-1.4$, $b\\sim-0.8$ to $-0.9$.  This defines the ``fundamental plane'' (FP) of elliptical galaxies: a correlation relating stellar mass or luminosity (implicit in the surface brightness), effective radius, and velocity dispersion (effectively the dynamical mass of the system). With a small observed scatter $\\sim0.1\\,$dex, the FP has presented as a long-standing, and still unresolved challenge to observations and theoretical models of spheroid formation.  In developing a physical understanding of the FP and associated elliptical scaling laws, \\citet{djorgovski:fp.tilt, jorgensen:fp.scatter} and others demonstrated that the FP could be represented as a ``tilted'' virial plane. If ellipticals were perfectly homologous systems with constant stellar mass-to-light ratios $\\mstar/L$, then a virial correlation $L \\propto \\mstar = k\\,\\sigma^{2}\\,R_{e}/G \\equiv \\mdyn$, with constant integral factor $k$, would be expected.  Since $I_{e} \\propto L/R_{e}^{2}$, this translates to an expected ``virial FP'' $R_{e}\\propto\\sigma^{2}\\,I_{e}^{-1}$.  The observed FP is similar to this, but not exactly so; it is equivalent to and can be represented as a ``tilted'' version of this correlation, namely \\begin{equation} \\mdyn \\propto \\mstar^{1+\\tilt} \\label{eqn:tilt} \\end{equation} with some small, but non-zero $\\alpha$. Equivalently, the difference between the best-fit observed FP (in any projection) and the virial FP can be expressed as a mass dependent mass-to-light (or, for our purposes, total mass-to-stellar mass) ratio \\begin{equation} \\frac{\\mdyn}{\\mstar}\\propto \\mstar^{\\tilt} \\end{equation} where the quantity $\\tilt$ quantifies the tilt, or deviation of the FP from the virial relation.  Various independent measurements find similar values of $\\tilt\\approx0.2$ \\citep[e.g.][]{pahre:nir.fp, gerhard:giant.ell.dynamics, borriello03, padmanabhan:mdyn.mstar.tilt,gallazzi06:ages}.  Although this is not strictly identical to the best-fit relation $R_{e}\\propto \\sigma^{a}\\,I_{e}^{b}$ if both $a$ and $b$ are fit as free parameters, multiple observations have shown that it is statistically an equivalent representation (i.e.\\ has the same scatter in {\\em physical} quantities), and that there is no additional information in the best-fit FP beyond this tilt (i.e.\\ once this tilt is accounted for, there is no additional systematic scaling in $R_{e}$ or $\\sigma$ that can reduce the scatter in predicting the other quantities).  It is now well-established that part of the observed tilt in optical bands is a consequence of stellar population effects: lower-mass ellipticals tend to be younger, yielding lower stellar mass-to-light ratios \\citep[see e.g.][]{trager:ages}.  However, various constraints imply that only a small fraction % of the optical tilt owes to these effects \\citep[see e.g.][] {pahre:nir.fp,gerhard:giant.ell.dynamics,bertin:weak.homology, borriello03,padmanabhan:mdyn.mstar.tilt,trujillo:non-homology, gallazzi06:ages,vonderlinden:bcg.scaling.relations}.  For example, in the $K$-band, the tilt is still substantial: observations indicate $\\mdyn\\propto L_{K}^{1.25\\pm0.05}$ \\citep{pahre:nir.fp}, whereas the systematic dependence of $\\mstar/L_{K}$ is quite weak \\citep[most estimates suggest $\\mstar/L_{K}\\propto L_{K}^{0.03}$;][]{bell:mfs}. It is now possible to combine high-resolution spectra and stellar population synthesis models, allowing reliable stellar mass estimates, and almost all such studies yield a similar relation: \\begin{equation} \\mdyn \\propto \\mstar^{1.2}, \\end{equation} i.e.\\ $\\alpha\\approx0.2$ as described above. It has been demonstrated that this result is robust to the details of the stellar population model, spectral coverage, or even simplifying assumptions such as the use of a single color to derive a mean $\\mstar/L$.  There are only two reasonable explanations for this finding (some combination of the two is also possible).  First, the true mass enclosed within the stellar effective radius ($\\mtrue$) could in fact be proportional to the stellar mass $\\mstar$, but owing to e.g.\\ changes in the profile shape or kinematics of galaxies with mass (traditional non-homology), the relation between actual mass and the dynamical mass estimator $\\mdyn\\propto\\sigma^{2}\\,R_{e}/G$ is a changing function of mass.  In other words, $\\mdyn \\propto \\mtrue^{1.2}$.  However, observations appear to rule out this possibility, at least as the origin of most of the tilt.  Integral modeling of the mass distribution from two-dimensional kinematic maps \\citep[which should recover any systematic difference between $\\mdyn$ and $\\mtrue$ without reference to any homology assumptions, e.g.][]{cappellari:fp}, as well as mass distributions estimated from gravitational lensing \\citep{bolton:fp,bolton:fp.update,nipoti:homology.from.mp}, independently give $\\mdyn \\propto \\mtrue^{1.00\\pm0.03}$.  That is, the allowed contribution of non-homology to the FP tilt is small.  The only remaining explanation is that the FP tilt reflects a meaningful physical change, namely that the ratio of total enclosed mass within $R_{e}$ ($\\mtrue$, represented reasonably well to within a normalization constant by $\\mdyn$) to the stellar mass is an increasing function of mass ($\\mtrue\\propto\\mstar^{1.2}$).  In other words, low-mass ellipticals are more baryon-dominated within their stellar $R_{e}$, and high-mass ellipticals have higher dark matter fractions.  We emphasize that these constraints apply {\\em within the stellar effective radii}. The change in dark matter fraction is not required to be global: if one were to contract the stellar mass distribution but keep the dark matter halo relatively fixed, for example, it would significantly decrease the dark matter fraction (and correspondingly $\\mdyn/\\mstar$) within $R_{e}$.  This trend is {\\em contrary} to that followed by disks.  For disk galaxies, an opposite (negative) tilt $\\mdyn\\propto\\mstar^{0.7-0.8}$ ($\\tilt \\approx-0.2 $ to $-0.3$) is observed \\citep[see e.g.][]{persic96,belldejong:tf,shen:size.mass,courteau:disk.scalings} -- low mass disks (and dwarf spheroidals) are the most dark-matter dominated systems \\citep[][]{persic88,persic90,persic96:data,persic96,borriello01}. Such a scaling is expected if the properties of disks track those of their dark matter halos: lower-mass halos are more compact \\citep[][and references therein]{neto:concentrations}, and it is also well-established that lower-mass disks experience less efficient star formation \\citep{belldejong:disk.sfh,gallazzi:ssps}.  Consequently, disks and ellipticals have similar ratios $\\mdyn/\\mstar$ at high masses ($\\sim$a few $L_{\\ast}$), but disks are more dark-matter dominated (have much higher $\\mdyn(R_{e})$) at low (stellar) masses. This difference in scaling laws also relates to their stellar size-mass relations: disks obey a shallow relation $R_{e}\\propto\\mstar^{0.25-0.35}$, roughly consistent with the scaling of halo effective radii as a function of mass, whereas spheroids obey a much steeper relation $R_{e}\\propto\\mstar^{0.6}$ \\citep{shen:size.mass}.  Again, at $\\sim$a few $L_{\\ast}$, disks and ellipticals have similar sizes and densities, but at low masses ($\\ll L_{\\ast}$), ellipticals are smaller (in their stellar/baryonic mass distributions) and more dense.  This difference has, for $\\sim30$ years, represented a major challenge for theory -- especially models which posit that ellipticals are formed through the merger of disk galaxies \\citep[the ``merger hypothesis'';][]{toomre72,toomre77}.  In particular, it has been noted that purely dissipationless mergers of stellar disks cannot raise the mass and phase-space densities of ellipticals, and so cannot change the scaling laws of ellipticals to something different from disks \\citep{ostriker80,carlberg:phase.space,gunn87}.  However, these arguments do not pertain if ellipticals are formed from mergers of {\\em gas-rich} disks.  Particularly at low masses, where ellipticals are more compact than spirals, disks have a large fraction of their mass in gas, so {\\em mergers must account for dissipation}.  In a merger of two disks containing both gas and stars, the stars are dissipationless and they cannot increase their phase space density and so violently relax to a distribution with an $R_{e}$ similar to the progenitor disks.  Gas, on the other hand, can radiate, and tidal torques excited during a merger can remove its angular momentum \\citep{hernquist.89,barnes.hernquist.91,barneshernquist96}. The resulting inflows produce a dissipational, merger-induced starburst which is compact (typical size scales $\\sim0.5-1$\\,kpc).  If a significant fraction of the final stellar mass is formed in this manner, the scale length of the stellar component will be much smaller than that of its progenitor.  \\citet{onorbe:diss.fp.model,robertson:fp} and \\citet{dekelcox:fp} argued that, because low mass disks are more gas rich (on average) than high mass disks, dissipation will be more important in low-mass systems.  That is, lower-mass ellipticals (the merger products of low-mass spirals) should have smaller effective radii relative to their progenitor disks, and, since the halo mass distribution is not strongly affected by this process, a correspondingly smaller dark matter fraction (lower $\\mdyn/\\mstar$) within the stellar $R_{e}$. On the other hand, high-mass disks are observed to be gas-poor: at $\\gg L_{\\ast}$ gas fractions become negligible ($\\ll 10\\%$) -- this is precisely where ellipticals are {\\em not} more compact than disks.  Together, \\citet{robertson:fp} and \\citet{dekelcox:fp} argued that the dependence of dissipational fraction on mass is sufficient, in principle, to explain the tilt of the FP, the size-mass and velocity-dispersion mass correlations of ellipticals.  The importance of gas dynamics and triggered star formation in mergers is reinforced by observations of ultraluminous infrared galaxies (ULIRGs) \\citep[e.g.][]{soifer84a,soifer84b}, which are invariably associated with mergers in the local Universe \\citep{joseph85,sanders96:ulirgs.mergers}.  The infrared emission from ULIRGs is thought to be powered by intense starbursts in their nuclei, originating in central concentrations of gas \\citep[e.g.][]{scoville86, sargent87,sargent89}, which will leave dense stellar remnants \\citep{kormendysanders92,hibbard.yun:excess.light,rj:profiles}, as predicted theoretically \\citep{mihos:cusps}.  Moreover, observations of merging systems and gas-rich merger remnants \\citep[e.g.,][]{LakeDressler86,Doyon94,ShierFischer98,James99}, as well as post-starburst (E+A/K+A) galaxies \\citep{goto:e+a.merger.connection}, have shown that their kinematic and photometric properties are consistent with them eventually evolving into typical $\\sim L_{\\ast}$ elliptical galaxies. The correlations obeyed by these mergers and remnants \\citep[e.g.,][and references above]{Genzel01,rothberg.joseph:kinematics, rothberg.joseph:rotation} are similar to e.g.\\ the observed fundamental plane and \\citet{kormendy77:correlations} relations for relaxed ellipticals, and consistent with evolution onto these relations as their stellar populations age, as well as the clustering and mass density of ellipticals \\citep{hopkins:clustering}.  Unfortunately, the consequences of the models are less clear; as such, there has been some ambiguity regarding whether or not recent observations of the FP support or disagree with theory. In particular, while merger remnants may fall on the FP, it is not obvious that a differential role of dissipation is in fact responsible for the FP tilt in the manner predicted by \\citet{robertson:fp}, or that this applies to all ellipticals.  However, there is hope: \\citet{mihos:cusps} predicted that these dissipational starbursts should leave an observable signature in the surface brightness profiles of remnants, in the form of a steep departure from the outer \\citet{devaucouleurs} $r^{1/4}$-law distribution in the inner regions: i.e. a central ``extra light'' component above the inwards extrapolation of the outer profile. Observations have now uncovered distinctive evidence for this two-component structure in local ellipticals \\citep{kormendy99,jk:profiles,ferrarese:profiles}, classical bulges \\citep{balcells:bulge.xl}, and recent merger remnants \\citep{hibbard.yun:excess.light,rj:profiles}. With the combination of HST and ground-based photometry, it now appears that such components are ubiquitous \\citep{jk:profiles}, with mass ranges and spatial extents comparable to those expected from observations of ongoing merger-induced starbursts and numerical simulations.  In a series of papers, \\citet{hopkins:cusps.mergers,hopkins:cusps.ell, hopkins:cores} (hereafter \\paperone, \\papertwo\\ and \\paperthree, respectively), we used these simulations and data to develop and test a method to empirically determine the degree of dissipation involved in the formation of a particular elliptical galaxy -- i.e.\\ the mass fraction in the stellar remnant of a central, compact nuclear (dissipational) starburst. In \\paperone, we demonstrated that observed merger remnants can be robustly decomposed into two components: an outer, dissipationless (violently relaxed) component with a Sersic law-like profile, comprising the pre-merger stars, and an inner, compact starburst remnant, produced in a starburst.  Combining large ensembles of observations with a library of simulations that enabled us to calibrate various empirical methods, we developed a purely empirical technique to separate the inner ``excess'' owing to the true physical starburst component in the observed surface brightness profile from the outer profile.  In \\papertwo, we showed that this method -- given photometry of sufficient quality and covering a large dynamic range (from $\\lesssim100\\,$pc to $\\gtrsim20-50$\\,kpc) -- could be extended to observed ``cusp'' ellipticals (i.e.\\ ellipticals with steep nuclear profiles), commonly believed to be the direct remnants of gas-rich mergers \\citep{faber:ell.centers}.  Separating the observed surface brightness profile in this manner, we demonstrated that simulations and independent observations (e.g.\\ distinctions in stellar populations evident in kinematics, colors, stellar ages, metallicities, or abundances) confirm that the component of the elliptical formed via dissipation (in a nuclear starburst) could be reliably (statistically) determined.  In \\paperthree, we showed that the same methods can robustly recover the dissipational starburst remnants in ``core'' ellipticals (ellipticals with shallow nuclear profiles). In general, even if other processes such as e.g.\\ scattering of stars by a binary black hole create the core, their impact on the overall starburst component is negligible (by both mass and radius, the scales of the core, typically $\\lesssim30-50\\,$pc, are much smaller than the mass and size of the starburst). We also showed that even if core or other ellipticals have subsequently been modified by spheroid-spheroid ``dry'' re-mergers, profile shape is preserved to a sufficient degree that the original nuclear excess (i.e.\\ the indicator of the degree of dissipation in the original, spheroid-forming merger) remains.  By applying our methodology to observations of ellipticals over a wide range in mass and size, we can, for the first time, empirically compare the degree of dissipation (starburst or dissipational mass fraction) in ellipticals to their global properties and locations on the FP.  If a differential effect of dissipation as a function of mass is the explanation for the FP and elliptical scaling relations, as the models suggest, then we should be able to see and quantify the signatures of this directly in their {\\it observed} profiles.  Therefore, in this paper we present a critical examination of the relationship between spheroid properties and FP correlations and ``extra light'' components, in both simulations and observed galaxies.  The combination of a large number of observations, together with an ensemble of hydrodynamic gas-rich merger simulations sampling the entire observed range in e.g.\\ mass, gas content, and other properties, enables us to develop and apply new, detailed, empirical tests of these models for the origin of the FP correlations.  In \\S~\\ref{sec:sims} we summarize our library of merger simulations, and in \\S~\\ref{sec:data} we describe the compilation of observations used to test the models.  In \\S~\\ref{sec:data:proxies} we review different approaches to fit the surface density profile and recover the physically distinct (dissipational versus dissipationless) components in merger remnants.  We use a set of simulations to infer how galaxy properties are predicted to scale with dissipation in \\S~\\ref{sec:diss.fx}. We then compare with observed systems: examining how observed sizes and masses scale with gas content in \\S~\\ref{sec:obs}, and the scaling relations obeyed at fixed dissipational content in \\S~\\ref{sec:obs.tests}.  We combine the observed dependences on dissipation and gas content in \\S~\\ref{sec:obs.tests.2} to determine whether this is sufficient to explain the tilt and scatter of the FP and its projected correlations. In \\S~\\ref{sec:obs.tests.mtot} we compare dynamical and total masses and consider possible non-homology effects.  In \\S~\\ref{sec:remergers} we examine the impact of subsequent re-mergers on the FP correlations. Finally, in \\S~\\ref{sec:discuss} we discuss our results and outline future explorations of these correlations.  Throughout, we adopt a $\\Omega_{\\rm M}=0.3$, $\\Omega_{\\Lambda}=0.7$, $H_{0}=70\\,{\\rm km\\,s^{-1}\\,Mpc^{-1}}$ cosmology, and normalize all observations and models accordingly.  We note that this has little affect on our conclusions, however.  We also adopt a \\citet{chabrier:imf} stellar initial mass function (IMF), and convert all stellar masses and mass-to-light ratios to this choice. The exact form of the IMF systematically shifts the normalization of stellar masses herein, but does not substantially influence our comparisons. All magnitudes are in the Vega system, unless otherwise specified.   \\breaker ",
        "conclusions": " \\citet{cappellari:fp} \\citep[see also][]{vandermarel:ml.models, kronawitter:ml.models,haringrix} estimate true enclosed masses within $R_{e}$ based on two and three-integral modeling, from two-dimensional velocity maps of observed ellipticals. Alternatively, \\citet{bolton:fp,bolton:fp.update} measure strong lensing gravitational masses.  In both cases, the authors find $\\mtrue\\propto\\mdyn$, without any significant dependence on mass -- in other words, the FP is unchanged regardless of whether the true total mass enclosed in $R_{e}$ is used, or whether the dynamical mass proxy $\\mdyn$ is used. It appears that at most, at the $\\sim3\\,\\sigma$ limits on the observed $\\mtrue-\\mdyn$ relation (and based on the fitted $\\mtrue-\\mstar$ relations in Figure~\\ref{fig:fp.pred.trueM}), these traditional forms of non-homology may contribute $\\sim1/4$ the observed tilt. \\\\  {\\em Genuine Change in $M/L$:} The observations therefore imply that the FP must reflect a genuine physical difference as a function of mass: namely, that low-mass ellipticals have a higher ratio of stellar or baryonic mass to total (baryonic plus dark matter) mass within the stellar $R_{e}$, relative to high-mass ellipticals.  All means of achieving this end can be classified into one of three categories:  {\\em (a) Varying Global Baryon Fractions:} One could imagine that the stellar and halo mass distributions of different ellipticals are separately self-similar, but that the {\\em total} ratio of stellar to halo mass changes as a function of galaxy mass. In other words, there are no structural changes implied: one decreases the halo mass relative to galaxy mass in lower-mass systems.  While this may seem plausible at the highest galaxy masses (it is well established that in bright clusters, the total stellar to dark matter mass ratio decreases with mass), this is relevant over much larger scales than the stellar $R_{e}$ of the central galaxy.  Moreover, over most of the mass range of interest, the known trends in global stellar to dark matter fractions are {\\em opposite} to that needed to explain the FP tilt.  At masses below $\\sim\\mstar$, where the FP is observed to be continuous (recall, the observed FP and its tilt extend to systems with masses $< 10^{-2}\\,\\mstar$), all $\\Lambda$CDM models \\citep[e.g.][]{conroy:monotonic.hod, zentner:substructure.sam.hod,zheng:hod, valeostriker:monotonic.hod,shankar06,vandenbosch:concordance.hod} and observational constraints \\citep[][]{eke:groups,yang:obs.clf,mandelbaum:mhalo, weinmann:obs.hod,wang:sdss.hod} {\\em require} that the ratio of {\\em total} dark matter halo mass to stellar mass is {\\em higher} in lower-mass systems -- i.e.\\ that star formation is less efficient in low-mass systems.  This is contrary to the effect desired here, and demonstrates that the FP tilt cannot owe to a simple global change in baryon fraction.  Indirectly, however, the cosmological trend of stellar to dark matter mass is in fact important -- this lower star formation efficiency in low-mass systems means that they have larger gas fractions when they undergo mergers, which we show does give rise to the FP tilt.  If the global baryonic mass ratio is unchanging at small radii (or changes in the opposite sense needed to tilt the FP), then the only other possibility is that the size/shape of the halo and stellar distributions vary relative to one another: i.e.\\ one of the two components is made more or less compact, relative to the other, altering the ratio of stellar to total mass within the stellar $R_{e}$.  Recall, the radii of interest are small relative to the halo virial radii, and contain only a small fraction of the total halo dark matter mass, so regardless of the total halo to stellar mass ratio, changing the compactness of one component relative to the other can make a significant difference.  There two possibilities for this relative contraction/expansion:  {\\em (b) Baryons are Fixed, Halos Contract:} In this scenario, the galaxy stellar mass distributions are ``scale free'' (i.e.\\ do not change owing to external factors), but the dark matter halos are less compact in low-mass ellipticals, so that within the stellar $R_{e}$ (i.e.\\ the central regions of the halo) the total dark matter mass fraction enclosed is lower. This is quickly ruled out: all cosmological models \\citep[e.g.][]{nfw:profile,bullock:concentrations,wechsler:concentration, dolag:concentrations,kuhlen:concentrations,neto:concentrations} and observations from e.g.\\ weak lensing and X-ray mass measurements \\citep{buote:obs.concentrations,schmidt:obs.concentrations, comerford:obs.concentrations} find that halo concentration is a weak, {\\em decreasing} function of galaxy mass, the opposite of the effect desired.  Furthermore, if this were the case, one would expect to see it in progenitor disks, as well (the halos are insensitive to the morphological transformation of their central galaxies) -- however, the baryonic Tully-Fisher relation and constraints on the dark matter halos of disks \\citep[e.g.][and references therein]{persic90,persic96,belldejong:tf,mcgaugh:tf} reflect the expected cosmological trends (namely, baryon fractions and halo sizes scaling as predicted in $\\Lambda$CDM, in a weaker and opposite sense from that necessary if the FP tilt were to be explained in this manner).  In short, effects {\\em (a)} and {\\em (b)} would predict that disks should follow the same FP scalings as ellipticals (modulo possible normalization offsets) -- while in fact, they obey different scalings with, in many cases, an {\\em opposite} qualitative sense \\citep[e.g.\\ their size-mass, velocity-mass, surface brightness-mass, and FP scalings;][]{fj76,kormendy:dissipation,shen:size.mass}.   \\begin{figure} \\centering \\scaleup \\plotter{f24.ps} \\caption{{\\em Top:} Mean stellar size-mass relation of ellipticals, from Figure~\\ref{fig:fp.projections}, compared to that of disks (from \\citet{shen:size.mass}). {\\em Middle:} FP of spheroids and disks -- i.e.\\ ratio of $\\mdyn$ (evaluated at $R_{e}$) to $\\mstar$. We show the observed ellipticals from Figure~\\ref{fig:fp.pred}, and the best-fit power law $\\mdyn\\propto\\mstar^{1+\\tilt}$ to both relations (with corresponding uncertainties). Disk correlation is from the data in \\citet{belldejong:tf}. At a given stellar mass, (low mass) spheroids have more compact stellar mass distributions, and less total mass ($\\mdyn$) enclosed within their stellar $\\re$. {\\em Bottom:} Same, but measuring $\\mdyn$ for the same observed ellipticals (from {\\em middle}) at the mean expected radius of an equivalent (similar stellar mass) disk (from {\\em top}). The best-fit relation to these data is shown; it is indistinguishable from the relation for disks. At the radii of their equivalent disks, ellipticals and disks have the {\\em same} enclosed total (dark matter plus stellar) masses. The distribution of dark matter is {\\em not} significantly different -- the smaller sizes and $\\mdyn$ of ellipticals must reflect a contraction of the baryonic material relative to the dark matter, as predicted to occur in dissipative mergers. \\label{fig:equiv.disks}} \\end{figure}   Figure~\\ref{fig:equiv.disks} demonstrates this explicitly. We show the stellar size-mass relation of ellipticals (from Figure~\\ref{fig:fp.projections}), compared to that of disks (from \\citet{shen:size.mass}, but for our purposes here different sources agree well). Modulo a small normalization offset (owing to the profile shape), ellipticals would necessarily obey the same correlation if they were formed in purely dissipationless mergers \\citep[see also][]{ciotti.van.albada:msig.constraints.on.gal.form}.  Obviously, ellipticals at low mass (where we empirically estimate and theoretically expect dissipation to be important) are much more compact in their {\\em stellar} distributions than disks.  We also compare the FP of both types of objects; i.e.\\ dynamical mass $\\mdyn$ (measured within $R_{e}$ of the stellar light) versus stellar mass. We estimate the relation for disks based on the data and best-fit relations in \\citet{belldejong:tf} \\citep[see also][who construct a similar correlation]{persic90,persic96}. Note that, for our purposes here, it makes little difference whether we plot the baryonic or stellar mass (there is almost no difference for ellipticals, owing to their small gas content, and for disks, the relations fall within the quoted uncertainties in either case).  It also makes no significant difference whether we use the same virial constant $k$ to estimate $\\mdyn$ for both disks and ellipticals, or attempt to make some correction for either profile shapes or the use of a velocity dispersion as opposed to a circular velocity (the difference is small, a factor $\\lesssim2$, and we are ultimately interested in the qualitative scalings).  It is clear (in agreement with previous work) that the scaling of $\\mdyn/\\mstar$ in disks is opposite that of ellipticals: there is either no tilt or {\\em inverse} tilt in the FP (i.e.\\ $\\mdyn/\\mstar$ is the same or {\\em higher} in low-mass disks).  For ellipticals where we have kinematic data as a function of radius, we can test whether the difference in the disk and elliptical scalings owes to effects {\\em (a)} or {\\em (b)}; we do so by evaluating their enclosed mass $\\mdyn$ not at the observed stellar effective radius of the spheroid, but at the radius of an equivalent disk (i.e.\\ at the mean $R_{e}$ of a disk of the same stellar mass). The resulting trend of $\\mdyn$ versus $\\mstar$ is indistinguishable from that of observed disks -- i.e.\\ by considering elliptical properties at their equivalent disk radii (radii they would have in the absence of dissipation), we effectively remove the tilt of the FP, and recover the observed correlations of disks.  At the radius of a disk of similar mass (equivalently, at the radius the elliptical would have, if it were the product of a purely dissipationless merger of stellar disks), ellipticals have the same enclosed total mass as equivalent disks. In other words, at the same (equivalent dissipationless) radius, disks and ellipticals of the same mass have the same dark matter mass content and distribution. In general, it is observationally well-established that the FP correlations of ellipticals become less distinct from those of disks as their properties are measured at larger radii. These correlations clearly rule out scenarios {\\em (a)} and {\\em (b)} above: if case {\\em (b)} were true, disks and ellipticals should obey a similar stellar size-mass relation, and at the $R_{e}$ of an equivalent disk, ellipticals should still have much lower $\\mdyn$ than spirals. If case {\\em (a)} were true (i.e.\\ both stellar and dark matter were more compact in ellipticals, but with lower total dark matter to stellar mass ratios), then we would again expect much lower $\\mdyn$ at the equivalent disk radius in ellipticals.   {\\em (c) Halos are Fixed, Baryons Contract:} The only remaining possibility is our conclusion in this paper: namely, that the FP tilt arises because lower-mass spheroids have more compact {\\em stellar} mass distributions, relative to their halos (equivalently, relative to what their stellar mass distribution would be in the absence of dissipation).  We have demonstrated that this outcome is a natural consequence of dissipation in mergers -- regardless of the initial scalings obeyed by progenitor disks, the fact that low mass disks are more gas-rich implies that, on average, there will be more dissipation in their mergers, yielding more compact baryonic remnants (while having little effect on the halo), and therefore increasing the ratio of stellar to dark matter mass inside the stellar $R_{e}$ in lower-mass systems. We further demonstrate that this reproduces precisely the observed tilt and scalings of elliptical properties with mass, while being consistent with all other observational constraints on the FP, including the ``homology constraint,'' that $\\mtrue\\propto\\mdyn$.  We have also shown that dissipation is {\\em necessary} in achieving this. It has been known for some time that dissipationless mergers cannot alter the phase-space density of ellipticals relative to their disk progenitors; consequently, ellipticals produced through dissipationless mergers obey the same scaling relations as their spiral progenitors (modulo normalization offsets), in stark contrast to the observed FP scalings. We demonstrate this explicitly: systems with the same dissipational fraction have the same ratio of total to stellar mass within $R_{e}$; i.e.\\ do not internally exhibit any FP ``tilt.''  Furthermore, we show that dissipation is the dominant factor determining the effective radii of ellipticals at fixed mass, even allowing for differences in formation and merger history -- therefore if the explanation for the FP invokes any systematic change in elliptical sizes, it {\\em must} involve dissipation.  Only when the observed dependence of dissipation on mass is included is the observed tilt recovered.     \\subsection{Additional Predictions} \\label{sec:discuss:pred}  We have extensively considered the role of dissipation in setting the FP tilt, effectively changing the ratio of $\\mdyn/\\mstar$.  To the extent that other properties also trace the degree of dissipation, we predict that these should similarly correlate with $\\mdyn/\\mstar$.  In \\paperone\\ and \\papertwo\\ we develop an extensive set of predictions for elliptical properties that relate to the degree of dissipation in the spheroid-forming merger, and show how these relate to, e.g., the observed extra light (i.e.\\ the tracer of the degree of dissipation in the spheroid forming merger-induced starburst).  We refer to those papers for details and a large number of observational proxies of the degree of dissipation which can be used to further test the ideas herein.  To the extent that the degree of dissipation in the original spheroid-forming merger reflects the gas fractions of the progenitor disks, it must also reflect the progenitor star formation history. Broadly speaking, disks with more extended star formation histories would be expected to have larger gas fractions at the time of their merger, with younger stellar population ages and lower $\\alpha$-enhancement. If we ignore the effect of the merger on these stellar populations (a reasonable assumption if the system is observed at times significantly later than the merger, and if the mass fraction formed in the merger-induced starburst is not large, which are true for most of the observed systems of interest here), then these should be reflected in the stellar populations of the elliptical remnant. Preliminary observational comparisons from \\citet{graves:prep} appear to support these predictions, and can provide powerful independent tests and constraints for models of dissipation in spheroid formation: we therefore outline some relevant quantitative predictions.  Consider the following highly simplified toy model: identical progenitor disks with initial gas fraction $\\fgas=1$ follow an exponential, closed-box star formation history with time scale $\\tau$, i.e.\\ $\\dot{M}_{\\ast}\\propto\\exp{(-t/\\tau)}$, and merge at time $t_{m}$, when the remaining gas is rapidly consumed in a central starburst. The gas fraction at the time of the merger (and correspondingly, the dissipational fraction in the starburst) will be $\\fsb=\\exp{(-t_{m}/\\tau)}$ (giving $\\tau = t_{m}/\\ln{(1/\\fsb)}$), and the mass-weighted mean formation time of the stars will be $t_{\\rm form} = \\tau\\,[1-\\exp{(-t_{m}/\\tau)}] = t_{m}\\, (1-\\fsb)/\\ln{(1/\\fsb)}$.  For systems with a similar redshift range of their last major merger (similar $t_{m}$, expected for systems of similar mass), their stellar formation times should therefore correlate with the starburst fraction $\\fsb$ (both depending implicitly on the pre-merger star formation timescale $\\tau$).  We have shown in \\S~\\ref{sec:obs} how $\\mdyn/\\mstar$ is predicted to scale with the dissipational fraction $\\fsb$ (this can be roughly approximated as $\\mdyn/\\mstar\\sim(\\fsb/0.2)^{-1/2}$ over the range of interest); combining the two, this yields an expected correlation between formation time and $x\\equiv\\mdyn/\\mstar$ of the form $t_{\\rm form} = 0.5\\,t_{m}\\, (1-0.2\\,x^{-2})/\\ln{(2.24\\,x)}$.  Observed at $z=0$ (i.e.\\ with age $t_{\\rm H}-t_{\\rm form}$), ellipticals with larger $x=\\mdyn/\\mstar$ at fixed mass (similar $t_{m}$) should be older -- for a typical $t_{m}\\sim10\\,{\\rm Gyr}$, systems with $\\mdyn/\\mstar\\approx 2$ are predicted to be $\\sim 2\\,$Gyr older than systems with $\\mdyn/\\mstar\\approx1$.  A shorter star formation history also implies higher $\\alpha$-enrichment in the progenitors. If we adopt the correlation for simple star formation models in \\citet{thomas05:ages}, $[\\alpha/{\\rm Fe}]\\approx 1/5-1/6\\,\\log{\\Delta\\,t}$ (where $\\Delta\\,t$ is the star formation timescale for a Gaussian burst, but for our purposes here can be replaced by $\\tau$ modulo a conversion constant), and the scalings above, then we obtain the result (independent of the merger time $t_{m}$) that ellipticals with higher $\\mdyn/\\mstar$ should be more $\\alpha$-enriched. Specifically, objects with $\\mdyn/\\mstar\\approx2$ are predicted to have $[\\alpha/{\\rm Fe}]$ values $\\approx 0.04-0.05$ higher than systems with $\\mdyn/\\mstar \\approx 1$.  Predictions for the absolute metallicities are more ambiguous (but see \\papertwo). In general, the trend of total/mean metallicity with mass will be dominated by the metallicity of pre-merger disks, which observations show (excluding the most gas-rich disks, where self-enrichment in the merger will dominate the final total metallicity) trace a similar mass-metallicity correlation to ellipticals \\citep{gallazzi:ssps}. In the absence of outflows or recycling, the metallicity would be the same for any systems with the same total accreted gas content and stellar mass (with no dependence on the ``starburst'' content at fixed mass). However, to the extent that outflows in low-mass systems are responsible for the mass-metallicity relation (as is generally believed), then the detailed interplay between these outflows and merger-induced starbursts will be important. In broad terms, if outflow strengths and velocities are similar, then we expect dissipational star formation at the center of the galaxy to retain a higher metal content in comparison to the same star formation in a more extended disk (since escape velocities from the galactic center and densities leading to radiative losses in the outflows are higher). In this case, at fixed mass, more dissipational (lower $\\mdyn/\\mstar$, higher surface brightness) systems should have slightly higher metallicities than their less dissipational counterparts (and, as demonstrated in \\paperone, this should be correlated with the effects above -- at fixed mass, the dependence of both quantities on dissipation should give rise to an inverse correlation between metallicity and stellar population age or $\\alpha$-enrichment). Experimenting with e.g.\\ different degrees of dissipation, outflow strengths, and initial disk metalliticies in our simulations, we estimate this to be a relatively small ($\\sim0.1$\\,dex) effect -- not sufficient to dramatically effect the global mass-metallicity relation, but potentially visible in detailed studies.  There is one important caveat here: systems of the same mass might also have had more gas-rich progenitors because they formed from mergers at very early times (i.e.\\ had a different merger time $t_{m}$, in the toy model illustrated above), making them older and more $\\alpha$-enriched. However, cosmological estimates \\citep[e.g.][]{hopkins:groups.ell} suggest this process is not dominant at a given stellar mass -- i.e.\\ systems of comparable stellar mass (and correspondingly similar total halo masses) tend to have similar merger histories. Specifically, the relatively large scatter in star formation history and disk gas fractions at fixed mass and redshift (a factor $\\sim2$ in $\\fgas$) is larger than the scatter introduced by the combination of a scatter in formation times and the systematic cosmological evolution of disk gas fractions with time. Furthermore, systems with such early mergers will usually have multiple subsequent mergers at later times, so they will grow significantly in mass and have their effective radii substantially modified by these additional processes (such that they should and will be compared to different systems at $z=0$).  Considering higher order effects, we demonstrate in \\papertwo\\ that the strength of stellar population gradients is correlated with the degree of dissipation in the original spheroid-forming gas-rich merger, and show in \\paperthree\\ that this holds even for remnants of subsequent gas-poor ``dry'' re-mergers. The most useful gradients in this sense are metallicity gradients -- stellar age gradients and (especially) color gradients evolve strongly with time even in a fixed, passively evolving elliptical (owing to the change in relative $M/L$ for young and old stellar populations) and as such are more ambiguous, and gradients in $\\alpha$-enhancement are more sensitive to the gradients and overall star formation histories in the pre-merger disks. Metallicity gradients are, on the other hand, generally dominated by the degree of dissipation and imprinted in the gas-rich merger, and are not sensitive to the star formation history of the pre-merger disks, making them a more robust diagnostic for our purposes. At fixed mass, stronger gradients indicate more dissipation, and so we predict that, at fixed mass, ellipticals with higher $\\mdyn/\\mstar$ should have weaker metallicity gradients (see \\papertwo).   \\subsection{Summary} \\label{sec:discuss:summary}  We have demonstrated from observations that the tilt of the FP owes to differential degrees of dissipation as a function of mass.  Lower mass disks are more gas-rich, so their mergers are more dissipational: a larger fraction of the remnant mass is formed in a dissipational, merger-induced compact central starburst in the final stages of a major merger. This yields a remnant with a more compact stellar mass distribution, i.e.\\ smaller $R_{e}$ relative to their progenitor disks, in lower-mass ellipticals.  The dark matter distribution is only weakly affected -- implying that the stellar distribution in low-mass ellipticals is more compact, relative to the dark matter, than in equivalent disks or higher-mass ellipticals. Consequently, relatively less dark matter mass is enclosed within the stellar effective radius in low-mass ellipticals, so the ratio of enclosed mass $\\mdyn/\\mstar$ is an increasing function of mass ($\\mdyn/\\mstar\\propto\\mstar^{\\tilt}$). This is the ``tilt'' in the FP. Given the observed, quantitative dependence of gas fractions on mass, the tilt is predicted to be exactly that observed, $\\tilt\\approx0.2$.  Using a new empirical method, we robustly estimate the amount of dissipation involved in the formation a given elliptical. Specifically, with data of sufficient quality, we separate the observed surface brightness profile into an outer, violently relaxed component, which was established in a dissipationless manner, and an inner ``starburst remnant'' or ``extra light'' component.  We demonstrate in \\paperone, \\papertwo\\ and \\paperthree\\ that observations of both evolved ellipticals and recent merger remnants support the proposal that this compact nuclear mass component a good proxy for the true mass formed in a dissipational starburst. Using this proxy, we show that the observed sizes of ellipticals, at fixed mass, depend strongly on the degree of dissipation involved in their formation (more so than even e.g.\\ the number of mergers in their formation history). Correspondingly, we show that the ratio of total to stellar mass within the stellar effective radius, $\\mdyn/\\mstar$, is a function of dissipation, both globally and at fixed mass (in the sense that elliptical sizes and $\\mdyn/\\mstar$ are decreasing functions of the amount of dissipation).  These observed dependences are highly significant ($P_{\\rm null}\\lesssim10^{-7}$).  Motivated by this, we show that by removing the mean systematic dependence of dissipation on mass, we can empirically remove the tilt of the FP.  Considering ellipticals with the same dissipational extra light fractions, we show that they obey a relation $\\mdyn\\propto\\mstar$ (i.e.\\ ellipticals with the same extra light content have the same ratio of dynamical to stellar mass within $R_{e}$, independent of mass).  In the proposed dissipational models of the FP \\citep{bekki:fp.origin.tsf,onorbe:diss.fp.model,robertson:fp,dekelcox:fp}, the tilt of the FP, and its projected correlations (e.g.\\ the steepness of the stellar size-mass correlation of ellipticals relative to that of disks), arise because low-mass disks are more gas-rich, and therefore low-mass mergers and ellipticals will have (on average) systematically higher degrees of dissipation, and therefore smaller (relative) $R_{e}$ and $\\mdyn/\\mstar$. If we consider e.g.\\ simulations that obey the observed correlation between disk gas content and mass (as opposed to being dissipationless, or having all the same gas fractions independent of mass -- neither of which is consistent with observations), then the FP predicted has exactly the observed tilt. Equivalently, the observed mean dissipational fractions of ellipticals, as a function of mass, agree well with the observed gas fractions of progenitor disks of the same masses, over the redshift range $z\\sim0-2$. In other words, {\\em dissipation is both necessary and sufficient to explain the FP tilt and differences between disk and elliptical scaling relations}.  To our knowledge, this is the first {\\em explicit} observational test of these theoretical models.  We further demonstrate that other mechanisms cannot be responsible for the majority of the FP tilt. For example, observations have demonstrated that the ``homology assumption,'' namely that $\\mtrue\\propto\\sigma^{2}\\,R_{e}$, is valid, and we show that simulations predict this -- the homology breaking introduced by dissipation is negligible.  In other words, the FP tilt reflects the ratio of stellar to true mass enclosed within $R_{e}$: $\\mtrue/\\mstar\\propto\\mstar^{\\alpha}$ (this ratio is {\\em not} constant), rather than an ``apparent'' effect.  We also show that, if we measure elliptical properties at the radius of an equivalent disk, the tilt of the FP is removed: within the radii of disks of the same mass, ellipticals have {\\em the same} total and dynamical masses. That is, the dependence of $\\mdyn/\\mstar$ cannot be driven primarily by changes in the dark matter distribution at fixed baryonic properties, nor by changes in the total dark matter to stellar mass ratio (integrated over the entire halo).  The variation in $\\mdyn/\\mstar$ within $R_{e}$ {\\em must} predominantly reflect the change in size of the baryonic component: low-mass ellipticals have much more compact stellar distributions than similar-mass disks, and therefore have less enclosed dark matter within that stellar $R_{e}$, as predicted in dissipational theories.   Together, these observational tests represent an important vindication of the ``merger hypothesis,'' that ellipticals are formed by the gas-rich mergers of disk galaxies, and models for the origin of the FP in dissipational major mergers \\citep{robertson:fp,dekelcox:fp}. We explicitly demonstrate that, regardless of subsequent gas-poor (spheroid-spheroid or ``dry'') re-mergers, the location of systems with respect to the FP and e.g.\\ elliptical size-mass and velocity dispersion-mass relations is primarily determined by the amount of dissipation involved in their formation: i.e.\\ the gas content involved in the original, spheroid-forming merger. Gas rich mergers cannot be ignored in the formation of ellipticals.  Not only have we demonstrated that the FP is consistent with the merger hypothesis, but that (given the systematic dependence of disk gas fractions on mass), a FP tilted in the manner observed is a necessary prediction of the theory.  We have also shown that elliptical sizes, inferred dissipational fractions, and the FP are completely consistent with the formation of ellipticals in mergers of disks with similar properties (sizes, gas fractions, dark matter halo masses and sizes) to those observed in {\\em low-redshift} ($z\\sim0-1$) disks. In other words, the sizes of ellipticals and their FP correlations {\\em do not require elliptical progenitors to be more compact than observed, low-redshift disks}. Dissipation is sufficient to explain the differences in their densities and sizes. The fact that, within the radius of an equivalent disk, ellipticals obey the same correlation between total and stellar mass actually implies that their dark matter halos (and presumably those of their progenitors) are {\\em not} significantly more compact than those of low-redshift disks.  This is important for the viability of the merger hypothesis, given the observations indicating that disk and halo sizes do not evolve strongly with redshift \\citep{ravindranath:disk.size.evol,barden:disk.size.evol, trujillo:size.evolution,zirm:drg.sizes}.  This is not to say that elliptical sizes might not evolve with redshift \\citep[which is easily possible if e.g.\\ disk gas fractions systematically evolve; see][]{khochfar:size.evolution.model,hopkins:bhfp.theory}, nor that ellipticals all formed at low redshift (indeed, if the disk size evolution is weak, then ellipticals form could rapidly at relatively early times and still resemble the products of low-redshift disks). However, it does imply that exotic progenitors -- progenitors not found in the local universe -- are not required to explain the observed correlations, surface brightness profiles (see \\papertwo\\ and \\paperthree), or kinematics \\citep[see][]{cox:kinematics} of typical local ellipticals."
    },
    "0806/0806.0370_arXiv.txt": {
        "abstract": "Observational limits on the high-energy neutrino background have been used to place general constraints on dark matter that annihilates only into standard model particles. Dark matter particles that annihilate into neutrinos will also inevitably branch into electromagnetic final states through higher-order tree and loop diagrams that give rise to charged leptons, and these charged particles can transfer their energy into photons via synchrotron radiation or inverse Compton scattering. In the context of effective field theory, we calculate the loop-induced branching ratio to charged leptons and show that it is generally quite large, typically $\\agt 1$\\%, when the scale of the dark matter mass exceeds the electroweak scale, $M_W$. For a branching fraction $\\agt 3$\\%, the synchrotron radiation bounds on dark matter annihilation are currently stronger than the corresponding neutrino bounds in the interesting mass range from 100 GeV to 1 TeV.  For dark matter masses below $M_W$, our work provides a plausible framework for the construction of a model for ``neutrinos only\" dark matter annihilations. ",
        "introduction": "\\label{AppI}  Here we present the amplitude for the box diagram of Figure (d), and a calculation of the divergent part of this amplitude. Then we construct the ratio of rates, \\beq {\\mathcal R}=\\left[ \\frac{\\langle v\\,\\sigma(\\chi\\bar\\chi\\rightarrow l^+l^-)\\rangle} {\\langle v\\,\\sigma(\\chi\\bar\\chi\\rightarrow \\nu\\ \\bar\\nu\\ )\\rangle} \\right]_{DivPart}\\,. \\eeq The subscripted qualifier ``DivPart'' is a reminder that only the leading (divergent) contribution to the electromagnetic rate will be included here.  In unitary gauge, the $W$-propagator is \\beq\\label{Wpropagator1} \\frac{-i\\,(g_{\\mu\\nu}-k_{\\mu}k_{\\nu}/M_{W}^2)}{k^2-M_{W}^2+iM_{W}\\Gamma_{W}}\\,, \\eeq and the matrix element for the amplitude of the box (d) is given by \\begin{eqnarray}\\nonumber i \\mathcal{M}_{\\rm box} &=&\\int\\frac{d^4 k}{(2\\pi)^4}\\left(\\frac{i}{(k-p+q)^2-M_{B}^2 + iM_B\\Gamma_B}\\right)\\left(\\frac{-i\\,(g_{\\mu\\nu}-k_{\\mu}k_{\\nu}/M_{W}^2)}{k^2-M_{W}^2+iM_{W}\\Gamma_{W}}\\right)\\\\ &&\\hspace{1.0cm} \\times\\,\\left(\\frac{i}{(k+q)^2 -M_{\\nu}^2 + i\\epsilon}\\right)\\left(\\frac{i}{(\\bar{q}-k)^2 -M_{\\bar{\\nu}}^2 + i\\epsilon}\\right)\\\\\\nonumber && \\\\\\nonumber &&\\times[\\bar{u}(q)\\Gamma_{W}^{\\mu}(\\slashed{k} + \\slashed{q} + M_{\\nu})\\Gamma_B u(p)]\\, [\\bar{v}(\\bar{p})\\Gamma_B(\\slashed{\\bar{q}} - \\slashed{k} + M_{\\bar{\\nu}})\\Gamma_{W}^{\\nu}v(\\bar{q})]\\,, \\end{eqnarray} where the momenta assignments to the particles are $\\chi(p)$, $\\overline{\\chi}(\\bar{p})$, $l^{-}(q)$, $l^{+}(\\bar{q})$, $W(k)$, which in turn determine the further assignments $B(k-p+q)$, $\\nu(k+q)$, and $\\bar{\\nu}(\\bar{q}-k)$. Here $\\Gamma^{\\mu}_W=(\\frac{g}{\\sqrt{2}})\\gamma^\\mu \\frac{(1-\\gamma_5)}{2}$ is the usual charged-current electroweak vertex with the $SU(2)$~coupling constant $g=e/\\sin\\theta_w$, $p$ is the four-momentum of the dark matter, $q$ is the four-momentum of the electron, and the internal loop four-momentum of the virtual $W$-boson is given by $k$. In addition, $\\Gamma_B$ in the numerator is the coupling times Lorentz structure assigned to a $B\\chi\\nu$ vertex, not to be identified with the $B$ width in the denominator. The value of the coupling we do not require, as the coupling cancels out when divided by the amplitude of the tree diagram; accordingly, we let simplicity be our guide and choose a scalar Lorentz structure for the $B$~field. We will neglect the neutrino and charged lepton masses, and define $M_\\chi$ to be the dark matter mass.  We introduce four Feynman parameters $\\xi_i$ (one per internal line) and find for the denominator $D$: \\begin{eqnarray} \\frac{1}{D} &=& (4-1)!\\int_{0}^{1}d\\xi_1\\int_{0}^{1}d\\xi_2\\int_{0}^{1}d\\xi_3\\int_{0}^{1}d\\xi_4\\, \\delta(1-[\\xi_1+\\xi_2+\\xi_3+\\xi_4])\\\\\\nonumber &&\\times\\,\\left[\\xi_1((k-p+q)^2-M_{B}^2)+\\xi_2(k^2-M_{W}^2) + \\xi_3(k+q)^2 + \\xi_4(k-\\bar{q})^2\\right]^{-4} \\end{eqnarray} Here we have dropped the $B$ and $W$ widths from the denominator, as they play no important role in the present calculation. Upon rotating $k$ to Euclidean space, we find that the divergent part of the amplitude, coming from the longitudinal mode of the $W$~propagator, is \\beq \\mathcal{M}_{\\rm box}=\\frac{-g^2}{8 M_{W}^2}\\int \\frac{d^4k_E}{(2\\pi)^4}\\int d\\xi\\frac{k_{E}^4}{(k_{E}^2 + \\Delta_B^2)^4} \\times[\\bar{u}(q)(1 + \\gamma_5)\\Gamma_Bu(p)]\\,[\\bar{v}(\\bar{p})\\Gamma_B(1-\\gamma_5)v(\\bar{q})]\\,. \\eeq Here we have defined the functional \\begin{eqnarray} \\int d\\xi &\\equiv& 3!\\int_{0}^{1}d\\xi_1\\int_{0}^{1}d\\xi_2\\int_{0}^{1}d\\xi_3\\int_{0}^{1}d\\xi_4\\, \\delta\\left(1-[\\xi_1+\\xi_2+\\xi_3+\\xi_4]\\right)\\\\\\nonumber &=& 6\\left[ \\int_{0}^{1}d\\xi_1\\int_{0}^{1-\\xi_1}d\\xi_2\\int_{0}^{1-\\xi_1-\\xi_2}d\\xi_3 \\right]_{\\xi_4=1-\\xi_1-\\xi_2-\\xi_3} \\end{eqnarray} and the parametrized mass-squared  \\beq\\label{DB} \\Delta_{B}^2 \\equiv \\xi_1\\,M_{B}^2 + \\xi_2\\,M_{W}^2 -(\\xi_1-\\xi_{1}^2 - 2\\xi_1\\xi_2 +4\\xi_{3}\\xi_4)\\,M^2_\\chi\\,. \\eeq In (\\ref{DB}) we have set the invariant $t\\equiv (q-p)^2$ equal to the value appropriate for a non-relativistic $\\chi\\bar\\chi$~annihilation, namely, $t\\approx -M_\\chi^2$.  Performing the integral over $k_E$ will result in a logarithmic divergence, and therefore we introduce a Pauli-Villars regulator for the $B$ propagator \\beq \\frac{i}{(k-p+q)^2-M_{B}^2} \\rightarrow \\frac{i}{(k-p+q)^2-M_{B}^2} - \\frac{i}{(k-p+q)^2-\\Lambda^2} \\eeq As is well-known, the Pauli-Villars regularization preserves local and global symmetries of the interaction. We then perform the integral over $k_E$, and find that the divergent amplitude after regularization becomes \\begin{eqnarray}\\label{Mbox} \\mathcal{M}_{\\rm box}=\\frac{-g^2}{2^7\\pi^2 M_{W}^2}[\\bar{u}(q)(1 + \\gamma_5)\\Gamma_B u(p)]\\, [\\bar{v}(\\bar{p})\\Gamma_B (1-\\gamma_5)v(\\bar{q})]\\int d\\xi\\,\\ln\\left|\\frac{\\Delta_{\\Lambda}^2}{\\Delta_{B}^2}\\right| \\,, \\end{eqnarray} where \\beq \\Delta_{\\Lambda}^2 = \\Delta_B^2 (M_B^2\\rightarrow\\Lambda^2) = \\xi_1\\,\\Lambda^2 + \\xi_2\\,M_{W}^2 -(\\xi_1-\\xi_{1}^2 - 2\\xi_1\\xi_2 +4\\xi_{3}\\xi_4)\\,M^2_\\chi\\,. \\eeq Due to the assumed scalar nature of the $B$ particle, the $\\gamma_5$'s in Eq.~(\\ref{Mbox}) may be omitted.  The tree-level amplitude for $\\chi\\bar\\chi\\rightarrow\\nu\\bar\\nu$ is \\beq\\label{Mtree} i {\\mathcal M}_{\\rm tree}= [\\bar{u}(q)\\Gamma_B u(p)]\\left(\\frac{i}{t-M_B^2}\\right)[\\bar{v}(\\bar{p})\\Gamma_B v(\\bar{q})]\\,. \\eeq Dividing the box amplitude by the tree-level amplitude (with $t\\approx -M_\\chi^2$, again) and squaring, we arrive at the result \\beq\\label{rateratio2} {\\mathcal R} = \\left(\\frac{M_{B}^2+M_\\chi^2}{128\\cdot 8\\pi\\cdot M_W^2}\\right)^2 \\left[\\ \\int d\\xi \\ln\\left|\\frac{\\Delta_{\\Lambda}^2}{\\Delta_{B}^2}\\right|\\ \\right]^2\\,. \\eeq For the running $SU(2)$ coupling we use the value applicable at the weak scale: \\beq g^2=\\frac{e^2}{\\sin^2\\theta_w}\\approx 4\\,e^2 =16\\pi\\,\\alpha \\approx \\frac{16\\pi}{128}\\,. \\eeq Notice that the spin-parity assignment of the $B$-meson appears only via the Dirac structure $\\Gamma_B$, which cancels out of the ratio ${\\cal R}$. Thus, our results to follow are independent of the $B$-meson's spin and parity.   There are some interesting, exactly calculable limiting cases. For $M_B^2\\gg M_\\chi^2,\\ M_W^2$, one has $\\Delta_\\Lambda^2/\\Delta_B^2\\approx \\Lambda^2/M_B^2$, independent of $\\int d\\xi$, and so Eq.~(\\ref{rateratio2}) becomes to lowest non-vanishing order, simply \\beq\\label{rateratio3} {\\mathcal R} \\longrightarrow 0.97\\times 10^{-7}\\, \\left(\\frac{M_B}{M_W}\\right)^4 \\left[\\ \\ln\\left(\\frac{\\Lambda^2}{M_{B}^2}\\right)\\ \\right]^2\\,, \\quad {\\rm for\\ }M_B^2\\gg M_\\chi^2,\\ M_W^2\\,. \\eeq We note that Eq.~(\\ref{massratiobnd}) allows this mass ordering, but does not require it.  Another interesting calculable limit is $M_\\chi^2\\ll M_B^2, \\Lambda^2, M_W^2$. Neglecting $M_\\chi^2$ in $\\Delta^2_B$ and $\\Delta^2_\\Lambda$, one finds that to lowest non-vanishing order, \\beq\\label{rateratio4} {\\mathcal R} \\longrightarrow  \\left(\\frac{M_{B}^2}{128\\cdot 8\\pi\\cdot M_W^2}\\right)^2 \\left[\\, \\ln\\left(\\frac{\\Lambda^2}{M_B^2}\\right) +\\frac{\\ln\\left(\\frac{M_W^2}{\\Lambda^2}\\right)}{\\left(1-\\frac{\\Lambda^2}{M_W^2}\\right)} -\\frac{\\ln\\left(\\frac{M_W^2}{M_B^2}\\right)}{\\left(1-\\frac{M_B^2}{M_W^2}\\right)} \\,\\right]^2\\,, \\ {\\rm for\\ }M_\\chi^2\\ll M_B^2, \\Lambda^2, M_W^2\\,. \\eeq If in addition to $M_\\chi^2\\ll M_B^2,\\Lambda^2, M_W^2$, one includes $M_B^2,\\Lambda^2 \\ll M_W^2$, i.e., a low-mass model with a low-mass cutoff, then there results \\beq\\label{rateratio5} {\\mathcal R} \\rightarrow  0.97\\times 10^{-7} \\left(\\frac{M_B}{M_W}\\right)^4 \\left[\\, \\left(\\frac{\\Lambda^2}{M_W^2}\\right)\\ln\\left(\\frac{M_W^2}{\\Lambda^2}\\right) -\\left(\\frac{M_B^2}{M_W^2}\\right)\\ln\\left(\\frac{M_W^2}{M_B^2}\\right) \\,\\right]^2, \\ {\\rm for\\ }M_\\chi^2\\ll M_B^2,\\Lambda^2\\ll M_W^2\\,. \\eeq Another limit, for a low-mass model with a high-mass cutoff, is \\beq\\label{rateratio6} {\\mathcal R} \\rightarrow  0.97\\times 10^{-7} \\left(\\frac{M_\\chi^2+M_B^2}{M_W^2}\\right)^2 \\left[\\, \\frac{ \\ln\\left(\\frac{\\Lambda^2}{M_W^2}\\right) }{ \\left(1-\\frac{M_W^2}{\\Lambda^2}\\right)} \\,\\right]^2, \\quad\\quad\\quad\\ {\\rm valid\\ for\\ }M_\\chi^2, M_B^2\\ll \\Lambda^2, M_W^2\\,. \\eeq The bracketed quantity in (\\ref{rateratio6}) is a monotonically increasing function of $\\Lambda/M_W$, equal to zero at $\\Lambda=0$ (as it must at $\\Lambda=M_B$), to 1 at $\\Lambda=M_W$, and growing as $\\ln(\\Lambda^2/M_W^2)$ at $\\Lambda^2\\gg M_W^2$. It is clear that with the ordering  $M_B^2\\ll M_W^2$, i.e., when the tree graph has a light-mass ($M_B$) propagator while the box graph has a heavy-mass ($M_W$) propagator, then the resulting electromagnetic branching fraction is negligible for any value of the effective field thoery cutoff $\\Lambda$, even up to the Planck mass. ",
        "conclusions": ""
    },
    "0806/0806.2834_arXiv.txt": {
        "abstract": "\\noindent Using the \\citet{NP} spin coefficient (NP) formalism, we examine the full Bianchi identities of general relativity in the context of gravitational lensing, where the matter and space-time curvature are projected into a lens plane perpendicular to the line of sight. From one component of the Bianchi identity, we provide a rigorous, new derivation of a Poisson equation for the projected matter density where the source term involves second derivatives of the observed weak gravitational lensing shear. We also show that the other components of the Bianchi identity reveal no new results. Numerical integration of the Poisson equation in test cases shows an accurate mass map can be constructed from the combination of a ground-based, wide-field image and a Hubble Space Telescope image of the same system. ",
        "introduction": "\\label{intro:sec}  In the current age of precision cosmology, weak gravitational lensing is an important tool in understanding cluster mass morphology.  Because all gravitating mass (baryonic or dark) influences the path of light rays on the same footing, gravitational lensing studies are the most direct measure of mass distribution within galaxies or galaxy clusters.  Current weak gravitational lensing studies relate the projected matter distribution to the observed gravitational shearing of images through an integral relation derived in \\citet{miralda95} and elsewhere. Previous work by \\citet{seitz} and others introduces a PDE approach where the gravitational shearing is related first to the gravitational potential, from which the underlying matter distribution can be found.  A recent paper by \\citet{kk} introduced a new PDE approach to weak gravitational lensing that directly relates the gravitational shear to the projected matter distribution.  By examining the weak gravitational perturbations on a flat background space-time in the NP spin coefficient formalism for a tetrad of constant null vectors, \\citet{kk} showed that one component of the Bianchi identity provides a first-order, complex partial differential equation relating the gravitational shearing of images to the projected matter density. Using a Green's function, the authors were able to show that their PDE was equivalent to the integral equation derived by \\citet{miralda95}.  This paper provides a more detailed derivation of the relevant PDE to first order in gravitational potential and using a non-constant NP null tetrad. To this order, we show that the other components of the Bianchi identity (that were not examined in \\citet{kk}) yield no new information related to gravitational lensing.  We show that a real Poisson equation naturally arises from this complex, first-order PDE that relates the gravitational shear to the projected matter density.  The Poisson equation derived here is related to the Poisson equation of \\citet{seitz}, but is derived from first principles of general relativity.  We numerically integrate the Poisson equation using relaxation methods for a matter distribution representing a truncated, isothermal sphere with a core radius. These numerical integrations are used to study the feasibility of the method for application to observational scenarios. ",
        "conclusions": "This paper gives a rigorous derivation of a Poisson equation, Eq.~\\ref{poisson3}, that directly relates the mass density to derivatives of the weak lensing shears.  We show that no new information is obtained by examining the remaining components of the Bianchi identity, so that the full theory of weak gravitational lensing and image distortion is most completely described by the optical scalar equations discussed in \\citet{fkn2}.  Using the simplest possible integration scheme, we show that the Poisson equation can be used to detect weak lensing signals, but fails to give accurate mass maps for wide-field, ground based images.  However, simple relaxation methods were shown to be successful when a ground and HST image were both present.  This method of determining the matter distribution could prove very useful in developing highly accurate mass maps.  Further examination of the Poisson equation is warranted. Specifically, it would be interesting to know whether maximum likelihood methods could yield accurate mass maps for ground based images alone.  Also, the application of the two grid method to mock background images that include appropriate levels of noise would solidify the usefulness of the method."
    },
    "0806/0806.3415_arXiv.txt": {
        "abstract": "We analyse $f(R)$ modifications of Einstein's gravity as  dark energy models in the light of their connection  with chameleon theories. Formulated as scalar-tensor theories, the $f(R)$ theories imply the existence of a strong coupling of the scalar field to matter. This would violate all experimental gravitational tests on deviations from Newton's law. Fortunately, the existence of a matter dependent mass and a thin shell effect allows one to alleviate  these constraints. The thin shell condition also implies strong restrictions on the cosmological  dynamics of the $f(R)$ theories. As a consequence, we find that the equation of state of dark energy  is constrained to be extremely close to $-1$ in the recent past. We also examine the potential effects of $f(R)$ theories in the context of the E\\\"ot-wash experiments. We show that the requirement of a thin shell for the test bodies is not enough to guarantee a null result on deviations from Newton's law. As long as dark energy accounts for a sizeable fraction of the total energy density of the Universe, the constraints which we deduce also forbid any measurable deviation of the dark energy equation of state from -1. All in all, we find that both cosmological and laboratory tests imply that $f(R)$ models are almost coincident with a $\\Lambda$CDM model at the background level. \\pacs{04.50.Kd, 95.36.+x, 12.20.Fv} ",
        "introduction": "The acceleration of the Universe expansion was discovered ten years ago and is still a deep mystery (see e.g. \\cite{de1} for recent results on observations of dark energy and e.g. \\cite{durrer,de2} for theoretical overviews). Two types of approaches have been considered. One can either introduce a new kind of matter whose role is to trigger acceleration or modify the behaviour of gravity at cosmological distances. In the first approach, dark energy is a new energy form, with all its well-known puzzles such as the cosmological constant problem, the coincidence problem and the value of the equation of state. This approach is subject of intense experimental investigation and any deviation from -1 would be a smoking gun for new physics beyond the standard models of particle physics and cosmology. On the other hand, in the second approach, various attempts to modify gravity have been presented (see e.g. \\cite{Carroll}-\\cite{Amendola}; the literature is vast, see \\cite{fara} for a recent overview and further references). Up to now, they are all plagued with various theoretical problems such as the existence of ghosts or instabilities. In this paper, we will consider a modification of Einstein's gravity, the so--called $f(R)$ theories, which do not seem to introduce any new type of matter and can lead to late time acceleration. In fact, these theories can be reformulated in terms of scalar-tensor theories with a fixed coupling of the extra scalar degree of freedom to matter. As theories of dark energy, they suffer from the usual problems and are also potentially ruled out by gravitational tests of Newton's law.  The only way-out for these models is to behave as chameleon theories \\cite{chamKA}, i.e. develop an environment dependent mass \\cite{navarro,brookfield,faulkner,li}. When the density of the ambient matter in which the scalar field/chameleon propagates is large enough, its mass becomes large and the smallness of the generated fifth force range is below the detectability level of gravitational experiments. On the other hand, for planetary orbits or any other situations in which gravity is at play in a sparse environment, one must impose the existence of the so--called thin shell effect. In this case, the fifth force is attenuated as the chameleon is trapped inside very massive bodies (the Sun for instance). It has been argued that the existence of thin shells is usually enough to salvage $f(R)$ models \\cite{navarro,faulkner}. We show that thin shells do not always guarantee null results in experimental tests of Newton's law. We exemplify this fact using the E\\\"ot-wash setting and obtain strong constraints on the models which translate into stringent bounds on the present dark energy equation of state, preventing any detection of a deviation from -1 in the foreseeable future  $\\vert 1+w\\vert \\le 10^{-4}$, where $w$ is the equation of state of dark energy in the recent past. This corroborates a similar bound obtained from the existence of  thin shell for objects embedded in a super-cluster. It should be noted that this result holds at the background level. For higher red-shifts where the effective dark energy density fraction, $\\Omega_{\\rm de}$, may become small (or even vanish), larger deviations can be present as exemplified in the models in \\cite{tsu1,tsu2}  where the equation of state can deviate from -1 for red-shifts of order $z=2-3$.  In all these models however $\\vert 1 + w\\vert \\Omega_{\\rm de} \\ll 1$, and so even if $w$ deviates significantly from  $-1$, deviations of the homogeneous cosmology from $\\Lambda$CDM are still very small. Detectable deviations from $\\Lambda$CDM are envisageable  at the perturbative level as the growth factor is anomalous at small scales (see e.g. \\cite{green} for a discussion of this point for the original chameleon model). Some consequences of this fact on the matter power spectrum and the CMB spectrum of $f(R)$ models have been presented in Ref. \\cite{Hu1,spergel,Hu2}.  The paper is organized as follows: In the subsequent section, we review $f(R)$ models and chameleon theories. In section III we derive the cosmological thin shell bound on the equation of state. In section IV, we consider tests of the inverse square law. Finally, we apply these considerations to particular models in section V. The appendices contain some technical details. ",
        "conclusions": "In recent years, modifications of General Relativity have been suggested as a possible explanation for the observed accelerated expansion of the universe. A popular class of models are the so--called $f(R)$ theories. While cosmologically viable theories can be found, local constraints on such theories have to be worked out, since the gravitational sector is modified, which could result in unacceptable deviations from Newton's law of gravity.  In this paper we have constrained $f(R)$ theories, using the well known equivalence between these and scalar-tensor theories. For an $f(R)$ theory to be consistent with both cosmology and local gravity experiments, the equivalent scalar-tensor theory must be a chameleon field theory. We have shown that the requirement of the thin-shell mechanism at work in E\\\"ot-Wash experiments results in an equation of state for dark energy very near to that of a cosmological constant. Thus, viable $f(R)$ models (those which are compatible with local experiments) behave on cosmological scales similarly to the standard $\\Lambda$CDM model and deviations are expected only on very small (sub-galactic) scales. The expected deviations from the cosmological constant equation of state $w=-1$ now in viable $f(R)$ theories are unmeasurably small (at least with current technologies). As examples, we have studied $f(R)$ theories with logarithmic potentials (based on \\cite{kaloper} for a fixed coupling $\\beta=1/\\sqrt 6$) as well as power-law potentials (such as those presented in \\cite{Hu2,Staro}). The former are ruled out by local gravitational tests, while there is still room for the latter models.  To conclude, while on cosmological scales viable $f(R)$ theories behave like $\\Lambda$CDM, deviations are expected on scales which could be large enough to be within the reach of next generation galaxy surveys \\cite{spergel}. Hopefully, future measurements of the dark matter distribution on those scales can be used to find such deviations from the standard $\\Lambda$CDM model. For this, a detailed understanding of galaxy formation is necessary, including an understanding of both the dynamics of baryons as well as that of dark matter in $\\Lambda$CDM and $f(R)$/chameleon theories.  \\vspace{0.5cm}  \\noindent{\\bf Acknowledgements:} CvdB, ACD and DJS are supported by STFC. We are grateful to Lisa Hall for a careful reading of the manuscript and to Nemanja Kaloper for useful comments.  DJS would like to thank  Kotub Uddin for helpful discussions, and D. J. Kapner for discussions and providing further details about the E\\\"{o}t-Wash experiment.  \\appendix"
    },
    "0806/0806.4832_arXiv.txt": {
        "abstract": "The effects of the in-flight behaviour of the bolometer arrays of the Herschel/PACS instrument under impacts of Galactic cosmic rays are explored. This instrument is part of the ESA-Herschel payload, which will be launched at the end of 2008 and will operate at the Lagrangian L2 point of the Sun-Earth system. We find that the components external to the detectors (the spacecraft, the cryostat, the PACS box, collectively referred to as the `shield') are the major source of secondary events affecting the detector behaviour. The impacts deposit energy on the bolometer chips and influence the behaviour of nearby pixels. 25\\% of hits affect the adjacent pixels. The energy deposited raises the bolometer temperature by a factor ranging from 1 to 6 percent of the nominal value. We discuss the effects on the observations and compare simulations with laboratory tests. ",
        "introduction": "The European Space Agency Herschel satellite will be launched at the end of 2008  with an Ariane-5 rocket and will operate at the Lagrangian L2 point of the Sun-Earth system (see ESA web page: {\\it www.rssd.esa.int/SA-general/Projects/Herschel}). Herschel is the ESA fourth cornerstone mission and will perform imaging photometry and spectroscopy in the far infrared and submillimetre part of the spectrum. The Herschel payload consists of two cameras/medium resolution spectrometers (PACS and SPIRE) and a very  high resolution heterodyne spectrometer (HIFI).  The PACS (Photo-conductor Array Camera and Spectrometer) instrument will perform efficient imaging and photometry in three wavelength bands in the range 60--210$\\mu$m, and spectroscopy and spectroscopic mapping with spectral resolution between 1000 and 2000 over the same wavelength range or short segments.  PACS is made of four sets of detectors: two Ge:Ga photoconductor arrays for the spectrometer part and two Si-bolometers for the photometer part. On each instrument side each detector covers roughly half of the PACS bandwidth\\footnote{ More about PACS can be found in the PACS Web pages: {\\it pacs.mpe-garching.mpg.de} and {\\it pacs.ster.kuleuven.ac.be}}.  It is well known that cosmic rays may influence strongly the detector behaviour in space. The performances of Infrared Space Observatory (ISO) detectors were largely affected for time periods long enough to corrupt a large amount of data \\cite{Heras}. Our goal is to exploit our present knowledge about the detectors and the cosmic environment to understand how the detector behaviour changes and how we may retrieve the lost information and/or extract the astronomical signal from the sources through an appropriate data analysis tool.  In this paper we focus our attention on the Si-bolometer arrays of the PACS instrument. These detectors differ from those on board of ISO as they are not of photoconducting type and very likely less affected by cosmic ray hitting. In a previous paper we dealt with the Ge:Ga photoconductors of the same instrument (\\cite{Bon06}).  The paper is organized as follows: in Sections \\ref{simulations} through \\ref{PL}, we briefly describe the used simulation toolkit, the input detector design, the Galactic cosmic ray spectra and the physics list. In section \\ref{checks} we discuss very simple tests to check the reliability of the present simulations. The results are reported in section \\ref{results} and discussed in section \\ref{discussion}, while in section \\ref{IPN} we compare simulations with laboratory tests. Conclusions about the effect on observations are to be found in \\ref{Conclusion}. ",
        "conclusions": "\\begin{itemize} \\item Red bolometers are less affected by GCR glitches than the blue ones. The pixel grid of the red bolometer is a factor of 1.569 larger so that its heat capacity (volumetric heat capacity times the volume) is correspondingly higher for the red pixels. The number of impacts is proportional to the detector surface. Both the red and the blue bolometer are affected by the same impinging flux but they have different collecting areas: 2048 grids for the blue and 512 grids for the red and the expected ratio of hit rate is: $\\frac{A_{\\rm blue} \\cdot 2048}{A_{\\rm red} \\cdot 512} = \\frac{173599.7 \\mu ^2 \\cdot 2048}{272392.1 \\mu ^2 \\cdot 512} \\sim 2.55$, where $A_{\\rm blue}$ and $A_{\\rm red}$ are the array areas of the blue and red bolometers respectively.  \\item We have quantified the contribution of secondary events produced by the 'shield'. This latter acts as the major source of secondary events, while the bolometer itself produces approximately 1/3 of them.  \\item The sensitivity of the bolometer cameras may be limited by glitch impacts for very faint sources; we expect a glitch contamination $\\le 5$--6\\%  of the background. The contamination is lower for the red camera.  \\item A test with the instrument simulator shows that in case of very faint sources a suitable observing technique and glitch removal algorithm must be put in place. In particular this latter should work on each chopping time interval in order to avoid false detections due to the sudden change of the flux level. \\end{itemize}"
    },
    "0806/0806.4690_arXiv.txt": {
        "abstract": "{} {We present a comparative study of several molecular lines and of the dust contiunuum at 1.2mm in a pre-stellar core that is embedded in the Galactic cirrus cloud MCLD123.5+24.9.  Previous studies found that the core is gravitationally stable and shows signs of inward motion. } {Using the Owens Valley (OVRO) and Plateau de Bure (PdB) interferometers we obtained high-angular resolution maps of the core in the carbon monosulfide CS $(2\\to1)$ and the cyanoacetylene HC$_3$N $(10\\to9)$ transitions. Together with CS $(5\\to4)$, C$^{34}$S $(3\\to2)$, and bolometer data obtained with the IRAM 30\\,m telescope, we analyse the excitation conditions and the structural properties of the cloud.} {On the one hand, the new CS $(J=2\\to1)$ observations reveal significant substructure on a scale of about $7''$, i.e., the beam size, corresponding to about 1050 AU at an adopted distance of 150\\,pc. On the other hand, the interferometric observations in the HC$_3$N $(J=10\\to9)$ transition shows just one single well resolved clump in the inner part of the core.  This core is well described by an intensity profile following from a centrally peaked volume density distribution. We find no evidence for depletion of CS onto dust grains. The inward motion seen in the CS $(2\\to1)$ occurs one-sided from the middle of the filamentary cloud towards the HC$_3$N core.} {} ",
        "introduction": "Dense cores in molecular clouds are the birthplaces of stars. One of the questions of current interest focusses on the evolution from a core to a protostar. Different classifications have been proposed, which are based on different observational diagnostics. Based on the detection or non-detection of IRAS point sources, cores have been separated into stellar or starless cores (Beichman et al. \\cite{beichman86}, Benson \\& Myers \\cite{benson89}). Because this classification is limited by the sensitivity and angular resolution of the IRAS satellite, in recent years some of the cores originally classified as ``starless'' showed the presence of embedded point sources (e.g., Young et al. \\cite{young04}).  A different classification scheme was proposed by Ward-Thompson et al. (\\cite{wardthompson94}, \\cite{wardthompson99}), who classified cores on the basis of the non-existence or existence of submm point sources as pre-protostellar (later on named 'pre-stellar' for brievity, Ward-Thompson et al. \\cite{wardthompson99}) or protostellar.  Unlike the classification of a core as being protostellar or stellar, prestellar cores cover a much wider range in their evolutionary states. Not all are necessarily gravitationally bound, while some of them could be in an early stage of contraction. To better quantify their state it is therefore necessary to study more pre-stellar cores in different environments.   \\begin{figure*} \\includegraphics[angle=-90,width=17.5cm]{8557fig1.eps} \\caption {Comparison of integrated intensity maps in different CS lines with maps in the dust continuum at 1.2mm and in the HC$_3$N $(10\\to9)$ line. Beamsizes are indicated in the lower right corner of each map. The open square marks the centre position of HC$_3$N-B, the open circle that of CS-C. Contours are every 0.4\\,K~km~s$^{-1}$ starting at 0.8\\,K~km~s$^{-1}$ for the CS\\,$(2\\to1)$ line (top left); every 0.1\\,K~km~s$^{-1}$ starting at 0.2\\,K~km~s$^{-1}$ for the HC$_3$N\\,$(10\\to9)$ line (top middle); and every 2\\,mJy/beam starting at 2\\,mJy/beam for the dust continuum map (top right); every 0.14\\,K~km~s$^{-1}$ starting at 0.14\\,K~km~s$^{-1}$ for the CS\\,$(5\\to4)$ line (bottom left); every 0.25\\,K~km~s$^{-1}$ starting at 0.25\\,K~km~s$^{-1}$ for the CS\\,$(3\\to2)$ line (bottom middle); and every 0.06\\,K~km~s$^{-1}$ starting at 0.06\\,K~km~s$^{-1}$ for the C$^{34}$S $(2\\to1)$ line (bottom right). The white line in the bolometer map indicates the major axis of the filamentary cloud.} \\label{integratedmaps} \\end{figure*} \\begin{figure} \\includegraphics[angle=0,width=9cm]{8557fig2.eps} \\caption{Channel maps of HC$_3$N-B in the HC$_3$N\\,(10$\\to$9) line as obtained with the Plateau de Bure interferometer.  The beam is indicated in the top left map. The centre velocity of each channel ($v_{\\rm LSR}$) is indicated in the lower left corner of each map. Contours are every 0.3\\,K starting at 0.3\\,K.} \\label{hc3nmaps} \\end{figure}  We continue our study of dense cores in Galactic cirrus clouds (e.g., B\\\"ottner \\cite{boettner2005}), and concentrate on MCLD123.5+24.9. Based on the above discussed definitions the core is both starless and pre-stellar. Nevertheless, CS observations towards part of the core obtained at high spectral resolution with the IRAM 30m telescope (Heithausen \\cite{heithausen99}, hereafter paper 1) revealed double-peaked CS profiles, which are interpreted as signature for infall motion (Myers et al. \\cite{myers:etal96}; Choi et al. \\cite{choi:etal95}).  This suggested a scenario of a possibly collapsing fragment embedded in a cloud that is gravitational unbound on larger scales. Further support for this idea was given by dust continuum and HC$_3$N spectral line observations, which showed that the cloud fragment was indeed gravitionally bound (Heithausen et al. \\cite{heithausen02}, hereafter paper 2).  The core itself is chemically evolved, showing a large number of molecules with abundances similar to other dense cores with on-going star-formation (Gro{\\ss}mann et al. \\cite{grossmann90}; Gro{\\ss}mann \\& Heithausen \\cite{grossmann92}; Heithausen et al. \\cite{heithausen95}; Gerin et al. \\cite{gerin97}).  Based on NH$_3$, HCN, HNC, and HCO$^+$ observations, a low kinetic temperature of less than 15\\,K was derived and its chemical abundances are well described by low-temperature chemistry with no indication of shock chemistry (Gro{\\ss}mann \\& Heithausen \\cite{grossmann92}).  The low temperature was confirmed with observations of the submm- and far-infrared continuum radiation, which revealed a dust temperature of only 13\\,K (Bernard et al. \\cite{bernard99}).  Subsequent observations of SO, CS, CCH, and C$_3$H$_2$ confirmed high column densities ($N({\\rm H_2})\\approx 10^{22}$cm$^{-2}$) and volume densities ($n({\\rm H}_2)\\approx10^5$cm$^{-3}$) of the core (Heithausen et al.  \\cite{heithausen95}; Gerin et al. \\cite{gerin97}).  At an adopted distance of 150pc (Heithausen \\& Thaddeus \\cite{heithausen90}) the core as seen in the dust continuum at 250GHz or in C$^{18}$O emission has a size of 0.18pc$\\times$0.03pc (paper 2). On that scale the core shows significant variation of its chemical abundances (Gerin et al. \\cite{gerin97}), e.g., cyanoacetylene, HC$_3$N, was found only at the ends of the filamentary structure, which was interpreted as being caused by different chemical ages of the different regions (paper 2).  All these results so far were obtained with moderate angular resolution of single dish telescopes. With such observations it is difficult to get detailed information on the density structure of the core or to study the connection between the inward motion and the densest region of the core.  Therefore, we present the first high-angular resolution data of this core in several spectral lines obtained with the Plateau de Bure (PdB) and the Owens Valley (OVRO) radio interferometers. At the distance of the core we reach an linear resolution of about 1000\\,AU. Details on the observations are presented in Sect.  \\ref{observations}. The results of the observations are discussed in Sect. \\ref{results}. Implications on the density structure and chemistry of the core are described in Sect. \\ref{discussion}. The paper is ended with some conclusions. ",
        "conclusions": "We presented a high-angular resolution study of the pre-stellar core in MCLD123.5+24.9 in several CS transitions and in the HC$_3$N $(10\\to 9)$ line as well as in the dust continuum. Our observations sharpens  the picture of gravitational collapse in that core. The major results of our study are the following: \\begin{itemize} \\item the CS emission shows structure down to the smallest scale oberved, \\item the HC$_3$N emission comes from a single well resolved core, \\item based on an analysis of the intensity profiles in the dust continuum and in the HC$_3$N line we find that the core is centrally condensed, \\item towards the upper half of the HC$_3$N core we have clear signs of inward motion, \\item towards the lower half the CS profiles could be interpreted as outward motion. \\end{itemize} The interpretation of the spectral line profiles follows the standard line discussed in the literature for pre-stellar cores. Other possibilities for the interpretation may exist. Clearly, a radiative transfer modelling is demanded to support or discard this interpretation."
    },
    "0806/0806.3079_arXiv.txt": {
        "abstract": "We report results of {\\em Chandra} X-ray and VLA radio observations of the Galactic accreting black hole \\vq\\ (GS~2023$+$338) in its quiescent state. \\vq\\ is detected at its faintest level of radio and X-ray emission with a 0.5--10 keV unabsorbed luminosity of 8.3 $\\times$ 10$^{32}$ (d/3.5 kpc)$^2$ \\ergs .  The X-ray spectrum fit with an absorbed power-law model yields a photon index of 2.17 $\\pm$ 0.13. Contrary to previous findings, this clearly indicates that \\vq\\ undergoes -- like most black holes in quiescence -- a softening of its X-ray spectrum at very low luminosity compared to the standard hard state.  The quiescent radio emission is consistent with the presence of self-absorbed compact jets.  We have also reanalyzed archival data from the decay of the 1989 outburst of \\vq\\ in order to quantify more precisely the correlation between radio and X-ray emission in the hard state of \\vq.  We show that this correlation extends over five decades in X-ray flux and holds down to the quiescent state of \\vq .  The index of this correlation ($\\sim$ 0.5) may suggest that synchrotron self-Compton emission is the dominant physical process at high energy in \\vq .  However, this index is also  consistent with scale invariant jet models coupled to an inefficiently radiating accretion disc. We discuss the properties of the quiescent state of black holes and highlight the fact that some of their properties are different from the standard hard state. ",
        "introduction": "Accreting black holes in X-ray binaries are known to undergo transitions between various ``X-ray'' spectral states (see \\citet{McClintock06}  for a review), mainly (but not only, see \\citealt{Homan01}) due to variation of the accretion rate within the accretion disk.  The quiescent state is the lowest luminosity state and is a factor of $\\sim$10$^6$  or more fainter than the brightest outburst state. In addition, a hard state is usually observed in the initial and final phases of an outburst with typical luminosity in the range 10$^{-3}$--10$^{-1}$ of the Eddington luminosity.   The quiescent and hard states share similar properties  (e.g. \\citealt{Tomsick04}).  Indeed, the quiescent state is often viewed as a lower luminosity version of the hard state. The compact jet observed in the hard state \\citep{Corbel00} appears also in quiescence as inferred from the characteristics of the radio spectrum \\citep{Gallo06}. The strong correlation between radio and X-ray emissions in the hard state \\citep{Corbel03} seems to be maintained down to quiescence \\citep{Corbel03,Gallo03,Gallo06}.    However, current X-ray satellites (especially  {\\em Chandra} and {\\em XMM-Newton}) have revealed new details of the spectrum of quiescent black holes. It appears that a fraction of them display a softer X-ray spectrum compare to the standard hard state \\citep{Corbel06}. In addition, deviations to the standard radio/X-ray flux correlation have been observed in the black hole \\gx\\ at very low luminosity (Corbel et al. in prep.). These peculiarities might imply that the quiescent state has to be considered as distinct from the standard hard state.   The universal radio/X-ray flux correlation presented by  \\citet{Gallo03} is dominated by two sources (\\gx\\ from \\citet{Corbel03} and \\vq ) plus additional points from other sources.  \\a0\\ in quiescence is, remarkably, consistent with an extrapolation of the \\vq\\ and \\gx\\ correlations down to quiescence \\citep{Gallo06}, but the exact track of \\a0\\ in outburst is unknown as no radio observations of a hard state were conducted in that time.  The correlation observed in the hard state of accreting galatic black holes has been extended to active galactic nuclei by  including  an additional correction for taking into account the mass of the black hole \\citep{Merloni03,Falcke04,Kording06}.  This fundamental plane of black hole activity relies strongly on the correlation observed in  \\vq\\ \\citep{Gallo03} and \\gx\\ \\citep{Corbel03}.  Thus, it is important to assess the reliability of the correlation for Galactic systems.   \\vq\\ has the longest orbital period of any black hole system detected in quiescence to date. \\citet{Corbel06} reported that long orbital period systems have quiescent spectra consistent with the hard state, contrary to short orbital period systems that have softer spectra.  However, the long orbital period group was statically dominated by the spectrum of \\vq\\ and recent {\\em XMM-Newton} observations of \\vq\\  show a soft quiescent spectrum \\citep{Bradley07}.  Therefore, we used a new {\\em Chandra} observation to reconsider the X-ray spectrum of \\vq\\ and the properties of black holes in quiescence.   In this paper, we describe the results of radio and X-radio observations of \\vq\\ in quiescence  and re-examine archival observations of \\vq\\ during the decay of its 1989 outburst (\\S 2). The observations in quiescence provide a detailed measurement of  the X-ray spectrum of \\vq\\  as well as a radio detection at its  faintest level of emission (\\S 3).  These first simultaneous radio and X-ray observations of  \\vq\\ in quiescence led us to revisit the radio/X-ray flux correlation in \\vq\\  (\\S 4). For that purpose, we reconsidered all X-ray observations of \\vq\\  during its 1989 outburst, allowing us to study this correlation in much finer detail. Our conclusions are summarized in section \\S 5. ",
        "conclusions": "We publish the results of a {\\em Chandra} and VLA observation of \\vq\\ in quiescence and we reevaluate the correlation between radio and X-ray emission.  The main conclusions of our work can be summarized as follows:  \\begin{itemize} \\item The {\\em Chandra} and VLA observations have allowed a detection of \\vq\\ at its faintest level of emission yet observed.  The characteristics of the radio emission remain consistent with the presence of self-absorbed compact jets in the quiescent state.  \\item The {\\em Chandra} observation of \\vq\\ confirms the softening of its X-ray spectrum in quiescence.  This implies that all black holes in quiescence have a softer X-ray spectrum in quiescence than in the standard hard state.  \\item We revisit and improve the correlation between radio and X-ray emission in \\vq\\ and found an index of the correlation that may suggest that the X-ray emission in the hard state of \\vq\\ is due to synchrotron self-Compton emission.  Compared to other sources, this would imply that the X-ray emission in black holes could be the results of interplay of various emission mechanisms (synchrotron, SSC or external Comptonisation) with different mechanisms favoured under conditions that still need to be understood. \\end{itemize}"
    },
    "0806/0806.2887_arXiv.txt": {
        "abstract": "The first billion years of the history of the Universe, commonly referred to as the Cosmic Dark Ages and the Epoch of Reionization, constitute a crucial missing link in our understanding of the evolution of the intergalactic medium and the formation and evolution of galaxies.  As gravity gradually amplified the small density contrasts in this dark-matter-dominated universe, structure arose hierarchically, with small objects condensing out first, then merging with each other to make ever larger objects in a clustered pattern in space known as the Cosmic Web.  Less than a hundred million years after the Big Bang, the stars forming in the first dwarf galaxies began to release ionizing radiation, which leaked out of these galaxies and caused ionization fronts to sweep outward through the surrounding primordial atomic gas, gradually transforming the cold, neutral intergalactic medium into a hot, ionized one. This epoch of reionization ended before the universe was a billion years old, at a time from which the light now reaching us has its wavelength redshifted by about a factor of 7.  Due to the complex nature of this global process it is best studied through large-scale numerical simulations. A number of large observational efforts trying to detect this process directly are currently under way, and their success depends critically upon correctly modelling its observable signatures.  Such simulations present considerable challenges, however, related to the large dynamic range required and the necessity to perform fast and accurate radiative transfer calculations. The hierarchical nature of cosmological structure formation in the Cold Dark Matter paradigm means that, at these early times, the dominant contributors of ionizing radiation were dwarf galaxies of around one billion Solar masses. These tiny galaxies must be resolved in very large cosmological volumes in order to derive their clustering properties and the corresponding observational signatures correctly, which makes this one of the most challenging problems of numerical cosmology.  We have recently performed the largest and most detailed simulations of the formation of early cosmological large-scale structures and their radiative feedback leading to cosmic reionization. This was achieved by running extremely large, $1728^3$- to $3072^3$-particle (5.2 to 29 billion) N-body simulations of the formation of the Cosmic Web, with enough particles and sufficient force resolution to resolve all the galactic halos with total masses larger than 100 million Solar masses in computational volumes of $(64/h\\,\\rm Mpc)^3$ to $(114/h\\,\\rm Mpc)^3$, respectively. These results were then post-processed by propagating the ionizing radiation from all (up to millions of) sources by using fast and accurate ray-tracing radiative transfer method. For these simulations, we utilized a P$^3$M N-body code called CubeP$^3$M and a radiative transfer code called C$^2$-Ray, both developed by us. Both codes are parallelized using a combination of MPI and OpenMP and to this date have been run efficiently on up to 2048 cores (CubeP$^3$M) and up to 10000 cores (C$^2$-Ray) on the newly-deployed Sun Constellation Linux Cluster (Ranger) at the Texas Advanced Computing Center.  In this paper we briefly present our codes, methods and parallelization strategies, discuss our recent simulations and  results and outline our future plans. ",
        "introduction": "Reionization is generally believed to be the outcome of the release of ionizing radiation by galaxies undergoing star formation (see for example \\citep{2006Sci...313..931B,2005SSRv..116..625C} for recent reviews). Current theory suggests that the galaxies responsible for most of this radiation are dwarf galaxies more massive than about $10^8M_\\odot$. In the Cold Dark Matter (CDM) universe nonlinear structures form from initially small-amplitude, Gaussian-random-noise density fluctuations, by a continuous hierarchical sequence of mergers and infall with the smallest galaxies forming first and merging to yield larger ones, forming the Cosmic Web of structures (Fig.~\\ref{web}).  While dark matter dominates the gravitational forces which cause this structure formation, ordinary atomic matter must join the dark matter in making galaxies for star formation to be possible.  Once the atomic gas in the intergalactic medium (IGM) in some region is heated by reionization, however, gas pressure opposes gravitational collapse, and, thereafter, the smallest galaxies form without atomic matter and cannot make stars. The minimum mass of star forming galaxies in such regions is about one billion solar masses.  Due to the complex nature of the reionization process it is best studied through large-scale numerical simulations \\citep{2006MNRAS.369.1625I, 2007ApJ...654...12Z, 2007ApJ...657...15K, 2007ApJ...671....1T}. Such simulations present considerable challenges related to the large dynamic range required and the necessity to perform fast and accurate radiative transfer calculations. The tiny galaxies which are the dominant contributors of ionizing radiation must be resolved in very large cosmological volumes, large enough to contain billions of times more total mass than one dwarf galaxy and up to tens of millions of such galaxies, in order to correctly derive their numbers and clustering properties. The correct number densities and clustering of these ionizing sources impact strongly the corresponding observational signatures. The ionization fronts expanding from all these millions of galaxies into the surrounding neutral IGM should then be tracked with a 3D radiative transfer method and the full non-equilibrium chemistry should solved in order to derive the resulting ionization state of the IGM. The combination of all these requirements makes this problem a formidable computational task.   \\begin{figure*} \\begin{center} \\includegraphics[width=3.0in]{6.0xy_red_enh3_small.ps} \\includegraphics[width=3.0in]{10x10Mpc_114Mpc_box_small.ps} \\end{center} \\caption{(left) Slice of the Cosmic Web of Dark Matter at redshift $z=6$ from our CubeP$^3$M Code simulation with $3072^3$ particles (29 billion) on a $6144^3$ fine grid in a comoving volume of 163~Mpc on a side. (right) zoom-in of 10~Mpc $\\times$ 10~Mpc subregion. Slices are 20 Mpc thick. \\label{web}} \\end{figure*}  We have recently performed the largest and most detailed simulations of the formation of early cosmological large-scale structures and their radiative feedback leading to cosmic reionization. This was achieved by running extremely large, $1728^3$- to $3072^3$-particle (5.2 to 29 billion) N-body simulations of the formation of the Cosmic Web, with enough particles and sufficient force resolution to resolve all the galactic halos with total masses larger than 100 million Solar masses in computational volumes of $(64/h\\,\\rm Mpc)^3$ to $(114/h\\,\\rm Mpc)^3$, respectively. These results were then post-processed by propagating the ionizing radiation from all (up to millions of) sources by using fast and accurate ray-tracing radiative transfer method. For these simulations, we utilized a P$^3$M N-body code called CubeP$^3$M and a radiative transfer code called C$^2$-Ray, both developed by us.  Our original simulations of the same process \\citep{2006MNRAS.369.1625I} resolved the formation of all galaxies more massive than about one billion solar masses, those which are expected to form stars even after reionization has heated their environment.  With this new generation of simulations we explore the role of less massive dwarf galaxies. These are important sources of ionizing radiation if they form before their neighborhood is reionized but are prevented from being sources if they form after it is reionized. Our preliminary results show that the inclusion of these sources and of their suppression changes the outcome of reionization substantially. ",
        "conclusions": ""
    },
    "0806/0806.1871.txt": {
        "abstract": "We generalize  to non-flat geometries the formalism of Simon et al. (2005) to reconstruct the dark energy potential. This formalism  makes use of quantities similar to the Horizon-flow parameters in inflation, can, in principle,  be made non-parametric and is general enough to be applied outside the simple, single scalar field quintessence. Since presently available and forthcoming data do not allow a non-parametric and exact reconstruction of the potential, we consider a general parametric description in term of Chebyshev polynomials. We then consider present and future measurements of $H(z)$, Baryon Acoustic Oscillations surveys and Supernovae type 1A surveys, and investigate their constraints on the dark energy potential. We find that, relaxing the flatness assumption increases the errors on the reconstructed dark energy evolution but does not open up significant degeneracies, provided that a modest prior on geometry is imposed.   Direct  measurements of $H(z)$, such as those provided by BAO surveys,  are crucially important to constrain the evolution of the dark energy potential and the dark energy equation of state, especially for non-trivial deviations from the standard $\\Lambda$CDM model. ",
        "introduction": "\\label{sec:intro} Recent observations e.g., \\cite{SpergelWMAP03,SpergelWMAP06,wmap5dunkley,wmap5komatsu,  Wood-VaseySN07, TegmarkLRGDR4,PercivalLRG} indicate  that the  present-day energy density of the universe is dominated  by a ``dark energy\" component, responsible for the current accelerated expansion.  The leading  dark energy candidates are a cosmological constant or a slowly varying rolling scalar field  e.g.,\\cite{Caldwell98, Quint99,Wang00,RatraPeebles88a,RatraPeebles88b,RatraPeebles95,Wetterich95} and \\cite{Steinhardt03} for a review, although explanations in terms of modifications of the Friedman equation (\\cite{DGP,DDG02,Carroll03,Carroll05,Capozziello05,NojiriOdintsov04,Dolgov03,HuSaw07, Acquaviva05,AV07} and references therein, under the name of ``modifications of gravity\" and e.g.,\\cite{Enqvist08, rasanen06,KKNNY07,KolbMatarreseRiotto06}  under the name of inhomogeneous models) are also being widely investigated.  An extended observational effort is being carried out  (e.g., SNLS, SDSS,  PanStarrs etc.)  and ambitious plans for the future are being proposed or planned (e.g., DES, PAU, BOSS, WFMOS, SNAP,  ADEPT,  DUNE, SPACE, SKA, LSST) with the goal of shedding light on the nature of dark energy. With few exceptions, current constraints on the nature of dark energy  measure an integrated value over time of the Hubble parameter, $H(z)$, which in turn is an integral of  its equation of state parameter ($w=p/\\rho$, with $p$ denoting pressure  and $\\rho$ density).  While these constraints are very tightly centered around the cosmological constant value, with a ~15\\% error, the finding that the average value of $w$ is consistent with $-1$ does not exclude the possibility that $w$ varied in time e.g., \\cite{Maor02}. An emerging technique in dark energy studies uses observations of the so-called baryon acoustic oscillations (BAO) \\cite{EH98, Eisenstein05, 2df05}.  The BAO  yield a measurement of the sound horizon at recombination,  a standard ruler visible at  different epochs in the lifetime of the Universe: at the last scattering surface through cosmic microwave background observations and  at lower redshifts  through galaxy clustering. In galaxy surveys,  the BAO scale  can be measured both along and perpendicular to the line sight. In particular the line-of-sight measurement offers the unique opportunity to measure directly $H(z)$,  rather than  its integral.  To improve our understanding of dark energy, it is important   not only to ask whether this dark energy component is dynamical or constant, but also,  to constrain possible shapes of the dark energy potential. As different theoretical models are characterized by different potentials, a reconstruction of the dark energy potential  from cosmological observations could help discriminating among different, physically motivated, models.  Different approaches to the reconstruction of the dark energy equation of state or potential have been proposed in the literature e.g., \\cite {S1,HutererTurner01,HS03, Simonetal05,S2,HP07}. In this paper we build upon and generalize the reconstruction technique proposed in \\cite{Simonetal05}  to non-flat universes. In fact \\cite{Polarski:2005jr,ClarksonCortesBassett07} showed that there can be a degeneracy between geometry and dark energy properties, thus  analyses to constrain dark energy parameters  should be carried out  varying  jointly the geometry of the Universe.  We then apply  this reconstruction  to  existing determinations of $H(z)$  from ages of passively evolving galaxies \\cite{Simonetal05}, to new  Supernovae data \\cite{Wood-VaseySN07}, and  we forecast the constraints on the dark energy potential achievable with  the next generation of BAO  and Supernova  surveys. We find that relaxing the assumption of flatness increases the error in the reconstructed dark energy evolution but does not open up degeneracies, provided that a modest prior on geometry is imposed: a Gaussian prior on  $\\Omega_k$ with  {\\it r.m.s.} $\\sigma_k=0.03$. Measurements of $H(z)$ such as those provided by BAO surveys, are crucial to constrain the dark energy evolution and to break degeneracies among dark energy parameters, for non-trivial deviations from the simplest $\\Lambda$CDM model.  The rest of the paper is organized as follows: in \\S 2  we present our reconstruction of the dark energy potential.  In \\S 3 we describe the priors  used in addition to the  present and future  data sets we consider. Our results on the dark energy potential reconstruction  are reported in \\S 4.  For completeness, we present a similar reconstruction  applied to the dark energy equation of state parameter as a function of redshift (in \\S 5). We conclude in \\S 6. ",
        "conclusions": "\\label{sec:conclusions}  We have generalized to non-flat geometries the formalism of \\cite{Simonetal05} to reconstruct the dark energy potential. This approach makes use of quantities similar to the horizon flow parameters used to reconstruct the inflation potential \\cite{schwarz01, Leach02}. The method can, in principle, be made non-parametric, but present and forthcoming data do not  allow a fully non-parametric reconstruction. We have therefore considered a parametric description in term of Chebyshev  polynomials which, for all our applications, we have truncated to  second order. For completeness we have also considered a reconstruction of the dark energy equation of state redshift dependence in terms of Chebyshev  polynomials, also generalizing to non-flat geometries the results of \\cite{Simonetal05}.   We have considered present  measurements of $H(z)$ from ages of passively evolving galaxies \\cite{ages03,Simonetal05},  future Baryon Acoustic Oscillation (BAO) surveys and present and future type IA supernova surveys and investigated their constraints on dark energy properties.  We present present and forecast constraints both on $V(z)$ (Fig.~\\ref{fig:potential}) and, more interestingly, on $V(\\phi)$ (Fig.~\\ref{fig:potentialphi}), in sec. 4. Model building  for dark energy which  rely on simple single-field models and  provide physically motivated potentials  should satisfy the constraints shown in the  left top and middle panels of Fig.~\\ref{fig:potentialphi}. In the future, the expected constraints can be as tight as those shown in the two bottom panels of Fig~\\ref{fig:potentialphi}. More complicated models (multi fields etc.) should  produce a redshift evolution of the effective dark energy potential   which  satisfies the  constraints in the left top and middle panel of  Fig.~\\ref{fig:potential}. The expected future constrains can be as tight as shown in the bottom panels of Fig.~\\ref{fig:potential}.   We find that  relaxing the flatness assumption increases slightly the errors on the reconstructed dark energy evolution, but does not generate significant degeneracies, provided that a modest prior on geometry is imposed $\\sigma_k=0.03$ (e.g., Fig.~\\ref{fig:omegak}).  Dark energy properties are best constrained at $z\\lap 0.3$: this is the result of the late-time dominance of dark energy. Under the assumptions made here, the most crucial being the assumption of a smooth $V(z)$ or $w(z)$, we find that high redshift ($z<2$) measurements of both $H(z)$ and $d_A$ are more powerful than low $z$ measurements.  When constraining the redshift evolution of both the dark energy potential $V(z)$  or the dark energy  equation of state parameter $w(z)$  with measurements of integrated quantities such as $d_A$ or $d_L$, there are large degeneracies among the parameters. These degeneracies are greatly reduced or  removed with measurements of $H(z)$, such as those provided e.g., by future BAO surveys. This is illustrated in Figs.~\\ref{fig:Hvsda_ground} and \\ref{fig:Hvsda_space} for the potential reconstruction.   While the $H(z)$ constraint is generally weaker than the $d_A(z)$ constraints, $H(z)$ is more directly related to the dark energy properties and thus offers more powerful dark energy constraints. We have illustrated this with an example of a model which lies on the ``$d_A$-degeneracy\" for the reconstruction of $w(z)$. This model  produces $d_A(z)$ and $d_L(z)$ virtually indistinguishable from that of the $\\Lambda$CDM, however the $H(z)$ are different and easily distinguishable from BAO measurements with $H(z)$ information.  This highlights an important open issue in dark energy studies: constraints on dark energy parameters coming from measurement of integrated quantities such as $d_A$ or $d_L$ depend crucially on the choice of  the dark energy parameterization [e.g., \\cite{Maoretal01}]: in the absence of a theoretical motivation for a parameterization of dark energy properties, forecasts and constraints become crucially model-dependent. The dependence of the constraints on the assumed dark energy parameterization becomes evident only when considering non-trivial deviations from a $\\Lambda$CDM model (e.g., deviations from a constant $w$  or generic shape of the potential). This issue is greatly alleviated  by measurements that carry information on $H(z)$."
    },
    "0806/0806.0285_arXiv.txt": {
        "abstract": "Color-flavor locked (CFL) quark matter at high densities is a color superconductor, which spontaneously breaks baryon number and chiral symmetry.  Its low-energy thermodynamic and transport properties are therefore dominated by the $H$ (superfluid) boson, and the octet of pseudoscalar pseudo-Goldstone bosons of which the neutral kaon is the lightest. We study the CFL-$K^0$ phase, in which the stress induced by the strange quark mass causes the kaons to condense, and there is an additional ultra-light ``$K^0$'' Goldstone boson arising from the spontaneous breaking of isospin.  We compute the bulk viscosity of matter in the CFL-$K^0$ phase, which arises from the beta-equilibration processes $K^0 \\leftrightarrow H+H$ and $K^0 +H \\leftrightarrow H$. We find that the bulk viscosity varies as $T^7$, unlike the CFL phase where it is exponentially Boltzmann-suppressed by the kaon's energy gap. However, in the temperature range of relevance for $r$-mode damping in compact stars, the bulk viscosity in the CFL-$K^0$ phase turns out to be even smaller than in the uncondensed CFL phase, which already has a bulk viscosity much smaller than all other known color-superconducting quark phases. ",
        "introduction": "\\label{intro}  The color-flavor locked (CFL) phase of quark matter is the densest predicted state of matter \\cite{Alford:2007xm}; it may occur in nature, in the core of compact stars, which are expected to reach several times nuclear saturation density. Quark matter is described by the theory of Quantum Chromodynamics (QCD), which, because of  asymptotic freedom \\cite{Gross:1973id,Politzer:1973fx}, becomes weakly coupled and hence perturbatively tractable at asymptotically high densities. In that regime, the CFL phase can be shown to be the ground state. However, at compact-star densities the coupling is strong, so first-principles calculations are not possible. We therefore follow a more phenomenological approach. We start with the hypothesis that color-flavor-locked quark matter occurs in compact stars, and calculate phenomenologically relevant properties which could allow us to find astrophysical signatures of its presence. We will calculate the bulk viscosity, which is relevant for the damping of pulsations and (indirectly) the spin down of the star.  If the interior of the star is a perfect (dissipationless) fluid, then a certain class of oscillating modes, $r$-modes, are unstable with respect to the emission of gravitational waves \\cite{Andersson:1997xt,Lindblom:2000jw,Owen:1998xg}. This emission acts as a brake on the rotation of the star. Since we know from observations that there are compact stars with very large rotation frequencies, we conclude that the instability must be damped. One damping mechanism is a nonzero viscosity of the fluid. Both shear and bulk viscosity can affect the $r$-modes and typically act in different temperature regimes. Therefore, it is of physical interest to compute shear and bulk viscosity of various candidate phases in a compact star as a function of temperature. Several calculations exist in the literature, for nuclear \\cite{Haensel:2000vz,Haensel:2001mw,Andersson:2004aa,Chatterjee:2007qs,Gusakov:2007px,Chatterjee:2007iw,Chatterjee:2007ka} and hyperonic \\cite{Lindblom:2001hd,Haensel:2001em} as well as for unpaired quark matter \\cite{Madsen:1999ci,Madsen:1992sx,Sa'd:2007ud} and various color-superconducting phases \\cite{Manuel:2004iv,Alford:2006gy,Sa'd:2006qv,Manuel:2007pz}, for a review see \\cite{Dong:2007ax}. Bulk viscosity due to thermal kaons in the CFL phase has been computed in Ref.\\ \\cite{Alford:2007rw}. Here, we extend this calculation to the case of condensed kaons.  At asymptotically large quark chemical potentials $\\mu_q$, three-flavor quark matter is in the CFL phase \\cite{Alford:1998mk}. This phase breaks the local color gauge group $SU(3)_c$, giving rise to Meissner masses for all eight gluons. It also breaks the global chiral symmetry group $SU(3)_L\\times SU(3)_R$, leaving a residual global $SU(3)_{c+L+R}$ which contains simultaneous rotations in color and flavor space, hence ``locking'' color with flavor. Moreover, the global baryon number conservation symmetry $U(1)_B$ is broken, rendering the CFL phase a superfluid. For massless quarks there are thus 8+1 massless Goldstone bosons (and an additional Goldstone boson for the $U(1)_A$ symmetry, which, however, is expected to be explicitly broken at moderate chemical potentials). At moderate densities the mass of the strange quark $m_s$ cannot be neglected and the bosons of the octet associated with chiral symmetry breaking acquire masses. These masses are small compared to the energy gaps of the fermions, and therefore the low-energy properties of the CFL phase can be described within an effective theory for the Goldstone bosons \\cite{Son:1999cm,BedaqueSchaefer,Kaplan:2001qk}. We shall make use of this theory in this paper. If $m_s^2/\\mu_q$ is large enough, it is expected that the lightest pseudo-Goldstone bosons, namely the neutral kaons, condense. The CFL phase with $K^0$ condensation is called the ``CFL-$K^0$'' phase. Goldstone bosons in CFL and their condensation have also been studied in a different approach, using a Nambu-Jona-Lasinio model \\cite{Forbes:2004ww,Buballa:2004sx,Warringa:2006dk,Ruggieri:2007pi,Ebert:2007bp,Kleinhaus:2007ve}.   Besides giving rise to masses for the meson octet, a nonzero strange mass induces a mismatch in the Fermi momenta of the quarks that form Cooper pairs in the CFL phase. In fact, the strange mass induces a mismatch in any possible spin-zero color-superconducting phase \\cite{Rajagopal:2005dg}. This means that at lower densities, the particularly symmetric CFL phase may be replaced by a less symmetric pairing pattern. It is currently not known whether, going down in density, the CFL phase is superseded by nuclear matter or by a different, more exotic, color-superconducting phase. Candidate color-superconducting phases have Cooper pairs with nonzero angular momentum \\cite{Schafer:2000tw,Alford:2002rz,Schmitt:2002sc,Schmitt:2004et} or nonzero momentum \\cite{Alford:2000ze,Rajagopal:2006dp,Mannarelli:2006fy}. In this paper, we shall only consider the CFL and CFL-$K^0$ phases.  The paper is organized as follows. In Sec.\\ \\ref{low} we give a brief overview over the properties of the CFL-$K^0$ phase, in particular we present the low-energy excitations at finite temperature. As an application, we discuss the resulting specific heat of the system in Sec.\\ \\ref{specific}. The calculation of the bulk viscosity is presented in Sec.\\ \\ref{bulk}. We define the bulk viscosity in Sec.\\ \\ref{definition}. In Secs.\\ \\ref{density} and \\ref{rates} we collect the ingredients needed for the bulk viscosity, namely the kaon density and susceptibility and the rate of the processes $K^0 \\leftrightarrow H+H$ and $K^0 +H \\leftrightarrow H$. We put these ingredients together in Sec.\\ \\ref{results} to obtain the result for the bulk viscosity and give our conclusions in Sec.\\ \\ref{conclusions}. ",
        "conclusions": "\\label{conclusions}  We have computed the bulk viscosity of kaon-condensed CFL quark matter (CFL-$K^0$ phase). Kaon condensation affects the low-energy properties significantly and therefore has a significant effect on thermodynamic and transport properties of color-flavor locked quark matter. In particular, the CFL-$K^0$ phase has a massless bosonic excitation associated with kaon condensation which is absent in the pure CFL phase. In both CFL-$K^0$ and CFL phases, there is also a massless Goldstone mode $H$ associated with superfluidity. For most of the thermodynamic properties, the effect of the additional Goldstone mode is important, but rather easy to predict. We have shown this for the specific heat, which acquires a contribution of the kaon mode which, as expected, has the same temperature dependence as the contribution of the superfluid mode. The prefactor of the former, however, is typically larger than that of the latter such that the kaon contribution in fact dominates the specific heat at low temperatures.  The effect of the additional Goldstone mode on the bulk viscosity is more complicated. We have used the results of our earlier work \\cite{Alford:2007qa} which provides a self-consistent description of the CFL-$K^0$ phase for arbitrary temperatures. Using the resulting thermal kaon mass and excitation energy, we have computed the density and the susceptibility of kaons, both of which are needed for the bulk viscosity. Moreover, we have computed the rate of the processes $K^0\\leftrightarrow H+H$ and $K^0+H\\leftrightarrow H$, where we denote by $K^0$ both the neutral kaon in the CFL phase and the massless Goldstone mode arising upon kaon condensation in the CFL-$K^0$ phase. These weak processes serve to re-establish chemical (flavor) equilibrium in response to an external volume oscillation, hence giving rise to a nonvanishing bulk viscosity.  At very high temperatures, $T\\gtrsim 10\\,{\\rm MeV}$, the difference in the kaon excitations in the CFL and CFL-$K^0$ phase is negligible. Consequently, in this case the kaon production (and annihilation) rate is almost identical for the two phases. At smaller (but not too small) temperatures, $10\\,{\\rm keV}\\lesssim T\\lesssim 10\\,{\\rm MeV}$, the masslessness of the Goldstone mode in the CFL-$K^0$ phase suppresses this rate because of a smaller available phase space for the weak process. Since the timescale of the rates in both phases is smaller than the typical oscillation (and rotation) frequency in a compact star, this effect decreases the bulk viscosity of the CFL-$K^0$ phase compared to the CFL phase. Another effect is given through the different susceptibilities. In the condensed system, the susceptibility at low temperatures is much larger than that of the uncondensed system. This effect works in the same direction, further decreasing the bulk viscosity compared to the uncondensed system. For even smaller temperatures, $T\\lesssim 10\\,{\\rm keV}$ the phase space actually is {\\em larger} in the CFL-$K^0$ phase and consequently the bulk viscosity is larger too. It is interesting to note that for the neutrino emissivity the effect of kaon condensation is quite different: neutrino emissivity in the CFL-$K^0$ phase is larger than in the CFL phase for all temperatures $T\\lesssim 10\\,{\\rm MeV}$ \\cite{Reddy:2003ap}.  We now have a fairly complete understanding of bulk viscosity in color-flavor-locked phases of quark matter. The suppression of the bulk viscosity due to the absence of ungapped fermionic excitations was predicted in Ref.\\ \\cite{Madsen:1999ci}. Subsequent more careful calculations took into account the contribution of the superfluid \\cite{Manuel:2007pz} and kaonic \\cite{Alford:2007rw} Goldstone modes. With the result of the present paper we have shown that the conclusion already drawn in Ref.\\ \\cite{Madsen:1999ci} is, for temperatures $T\\lesssim 1\\,{\\rm MeV}$, not changed by the contribution of the Goldstone modes: color-flavor locked quark matter, even in the presence of kaon condensation, has a much lower bulk viscosity than all other known phases of dense quark matter and than nuclear matter. Only at large temperatures, and thus in very young neutron stars, can the contribution of the Goldstone modes render the bulk viscosity comparable to that of unpaired quark matter.  We finally mention that besides the bulk (and shear) viscosity, other properties of color-superconducting quark matter also deserve attention. Its equation of state may be used to put constraints on the mass-radius relation of hybrid stars with a quark matter core and a hadronic mantle. These calculations rely on simple models whose parameters are poorly known in the strong-coupling region of interest. While NJL model calculations, mainly due to their relatively large predicted strange quark mass, tend to find no stable hybrid star with a CFL core \\cite{Baldo:2002ju,Klahn:2006iw}, other parametrizations of the equation of state allow for hybrid stars with masses compatible with the observations \\cite{Alford:2004pf}. Other observables that may distinguish between certain phases of color-superconducting quark matter or between quark matter and nuclear matter are for instance the cooling curve of the star or glitches (sudden spin-ups). The corresponding transport properties of color superconductors have already been computed in the literature, see for instance Refs.\\ \\cite{Jaikumar:2002vg,Schmitt:2005wg,Jaikumar:2005hy,Anglani:2006br} and Ref.\\ \\cite{Mannarelli:2007bs} for neutrino emissivity and shear modulus, respectively. It is an interesting and promising task for the future to extend these calculations and compare them with more and better astrophysical data in order to understand matter inside a compact star, and, ultimately, map out the phase diagram of cold and dense quark matter."
    },
    "0806/0806.2280_arXiv.txt": {
        "abstract": "{}{We  present   the  abundance  measurements  of several elements (Fe,  Ca, Na, Ni, Ti, Al,  Cr, Si) for 20 solar--type   stars  belonging   to  four   Galactic  open clusters: NGC~3680,  IC~4651, Praesepe, and  M~67.  Oxygen abundances  were in  addition measured  for most  stars in each cluster  apart from IC~4651.   For NGC~3680, accurate abundance  determinations  using high--resolution  spectra covering  a large  spectral  domain are  computed for  the first  time.}   {  We  used  UVES  high--resolution,  high signal--to--noise  (S/N)  ratio  spectra and  performed  a differential  analysis   with  respect  to   the  sun,  by measuring equivalent widths  and assuming LTE.}  {The most surprising   result  is   a  measurement   of  significant supersolar         metallicity         for        Praesepe ([Fe/H]=0.27$\\pm$0.10).   As for  the  other clusters,  we confirm    a    supersolar    metallicity   for    IC~4651 ([Fe/H]=0.12$\\pm$0.05),  a   solar  metallicity  for  M~67 ([Fe/H]=0.03$\\pm$0.04)  and a slight  subsolar metallicity for  NGC~3680 ([Fe/H]=-0.04$\\pm$0.03).   We find  that the abundance ratios  of almost  all elements are  solar, with the notable exception of  oxygen in NGC~3680 and Praesepe, supersolar in the former cluster ([O/Fe]=0.2$\\pm$0.05) and as low as [O/Fe]=-0.4$\\pm$0.1 in the latter.  Observations of  several  objects per  cluster  is  required to  obtain robust  results,  especially  for  those elements  with  a limited number of suitable lines. }{} ",
        "introduction": "Understanding the  chemical evolution of the Galaxy  requires both the development of theoretical models  and comparison of their predictions with observational results. Several different models have been created of  increasingly more  complex and  realistic scenarios,  for instance including two infall episodes that formed the halo-thick disk and thin disk   \\citep{cmg97,cmg01,pat04}.    Observationally,   new-generation telescopes such as VLT and  Keck and the availability of multi--object spectroscopy  have  enabled  a  significant  amount  of  high  quality spectroscopic  A long-standing question concerns  the existence and evolution  of the chemical  abundance gradient   in  the Galactic  disk.   Galactic open clusters are probably the best tool for  understanding whether and how the gradient  slope changes with time because  they have formed at all epochs.  Their distances can be measured  more accurately and are less affected by observational biases than other classes of objects.  \\cite{dl05}   identified the   birthplace of  612   open  clusters and determined the spiral pattern  rotation  speed of the Galactic   disk. The possibility of tracing open clusters back  to where they formed is interesting because, coupled with  metallicity measurements,  it allow us to reconstruct  the chemical distribution of  the Galactic disk  in the past.  \\cite{dl05} based their work on large photometric databases in the visible.  They  achieve important conclusions on the  structure and   dynamics of  the spiral   arms,  but to   estimate precisely the deviation of open cluster  motions  from circular orbits, and  improve the scale of  Galactocentric distances at  the place of birth, a  more complete  and precise database   of age, distance,  proper  motion and radial velocity    determinations   is required.         Unprecedented opportunities    are provided by   deep   photometric  surveys in  the infrared, such as UKIDSS \\citep{ukidss} and VISTA \\citep{vista}. These surveys   can be used to  complete  deep, precise and homogeneous open cluster  databases of location in the   3-D space, age, and reddening. The calibration  of photometric metallicity  indicators could  be used for clusters not observable with high--resolution spectroscopy.  Open clusters have the added advantage of providing a sample of coeval stars that formed from  the same material,  which means, in particular for main--sequence stars,  that they should  have the same atmospheric chemical   composition.  As a result,   the chemical composition of an open cluster can   be studied by  several  stellar spectra.  With  all these  advantages,  open  clusters are   ideal   objects to probe  the chemical evolution of the Galactic disk.  A large dataset, spanning as wide a range of Galactocentric distances and ages,  as possible, is of course required.  Finally, clusters are  at the basis  of  our understanding of  stellar evolution,  and  their colour--magnitude diagrams, and well-determined abundances  can be   used to test  stellar   evolution models directly \\citep[see e.g.][]{naa97}.   To  date, accurate chemical composition data  have been determined for only a few open clusters using high  resolution spectra of a number of primarily giant stars  \\citep{gratton00}.   The situation is  evolving rapidly due  to coordinated  efforts   using facilities available   at 8--M--class  telescopes   \\citep[][]{largeprog,    BOCCE}.  This  work combines with   other  efforts  attempting  to   establish  a   robust open-cluster  metallicity  scale.  Random  and  systematic errors have given rise  to  -- sometimes dramatic  --  discrepancies  in abundance determinations  of a  given open  cluster  as  computed from different groups.  We confirm  that this is no longer   the case when  comparing independent      metallicity   measurements    from  high   resolution spectroscopy. ",
        "conclusions": "This   work is one of   several  aimed  at collecting  high--precision abundance measurements in open clusters, based on high--resolution and high--signal--to--noise  spectra.  We present  data for four clusters, and confirm all of the abundance trends present  in the literature for IC~4651  and M~67.  For these   clusters, all determinations agree to within  a few hundredths  of dex, irrespective of whether dwarfs or giants have been observed.  We report slightly different abundances for NGC~3680  and mostly Praesepe, which is  shown to be substantially metal rich.  The  modern  iron determinations  agree  very  well  in most of cases, though there  is still no precise agreement  about the extent to which Praesepe and Hyades are supersolar.  Metallicity measurements for open clusters  are in general    consistent, irrespective of   the analysis method and research group.  As for abundance  ratios of elements other than iron,  significant  differences  between   the results  of    the different groups still arise.  The dispersion within each cluster is limited to a few hundredths of a dex, and differences of up to 0.1 dex or more are measured between the individual stars.  These differences  have to be interpreted as  being due to limitations in the observations and in the analysis rather than true, intrinsic  chemical  variations   in  the  cluster  composition. Caution should therefore be taken when adopting analyses that consider only one or a few stars per cluster.  We     find  a   significant   supersolar   metallicity    for IC~4651 ([Fe/H]=0.12$\\pm$0.05)  and   a    solar    metallicity    for    M~67 ([Fe/H]=0.03$\\pm$0.04)    and   a slightly   subsolar  metallicity for NGC~3680 ([Fe/H]=-0.04$\\pm$0.03).  The  surprising  result is that  of Praesepe,  which  is  as metal  rich   as  [Fe/H]=0.27$\\pm$0.10.   The photometric analysis  of  \\cite{an07} also indicates that  Praesepe is metal rich. The reasons for the discrepancy with other results need to be further investigated.  For other elements, the composition is solar within the errors for all elements and for  all the clusters,  except for  aluminum (subsolar in NGC~3680  and   IC~4651),  nickel  (also   subsolar in NGC~3680),  and oxygen. We   find [O/Fe] to  be  $\\approx 0.2$  and  $\\approx -0.4$ in NGC~3680 and Praesepe respectively, and  just $\\approx 1 \\sigma$ below the solar value in M~67.  We confirm the extraordinarily similar overall chemical composition of the Sun and the star cluster M~67."
    },
    "0806/0806.2777_arXiv.txt": {
        "abstract": "{ {\\sl Aims.} We have conducted a study in radio wavelengths and in X-rays of the pulsar wind nebula (PWN) in the  supernova remnant (SNR) G0.9+0.1 with the goal of investigating in detail its morphology and to accurately determine its characteristic parameters.  {\\sl Method.} To carry out this research we have observed the PWN at $\\lambda$ 3.6 and 6 cm using the Australia Telescope Compact Array (ATCA) and combined these data with existing multiconfiguration VLA data and single dish observations in order to recover information at all spatial scales. We have also reprocessed VLA archival data at $\\lambda$ 20 cm. From all these observational data we have produced high-fidelity images at the three radio frequencies with angular resolution better than 3\\s. The radio data were compared to X-ray data obtained with {\\it Chandra} and in two different observing runs with {\\it XMM-Newton}.  {\\sl Results.}  The new observations revealed that the morphology  and symmetry suggested by {\\it Chandra} observations (torus and jet-like features) are basically preserved in the radio range in spite of the rich structure observed in the radio emission of this PWN, including several arcs, bright knots, extensions and filaments. The reprocessed X-ray images show for the first time that the X-ray plasma fills almost the same volume as the radio PWN. Notably the X-ray maximum does not coincide with the radio maximum and the neutron star candidate \\cxou~ lies within a small depression in the radio emission. From the new radio data we have refined the flux density estimates, obtaining S$_{\\rm PWN}\\sim 1.57$ Jy, almost constant between $\\lambda$ 3.6 and $\\lambda$ 20 cm. For the whole SNR (compact core and shell), a flux density S$_{\\rm 20 cm}$= 11.5 Jy was estimated. Based on  the new and the existing $\\lambda$ 90 cm flux density estimates, we derived a spectral index $\\alpha_{\\rm PWN}=-0.18\\pm 0.04$ and $\\alpha_{\\rm shell}=-0.68\\pm 0.07$. From the combination of the radio data with X-ray data, a spectral break is found near $\\nu\\sim 2.4 \\times 10^{12}$ Hz. The total radio PWN luminosity is L$_{\\rm radio}=1.2\\times10^{35}$ erg s$^{-1}$ when a distance of 8.5 kpc is adopted. By assuming equipartition between particle and magnetic energies, we estimate a nebular magnetic field B$ = 56~ \\mu$G. The associated particle energy turns out to be U$_{\\rm part}=5 \\times 10^{47}$ erg and the magnetic energy U$_{\\rm mag}=2 \\times 10^{47}$ erg.  The high ratio between magnetic and particles flux energy density suggests that the pulsar wind just started to become particle dominated. Based on an empirical relation between X-ray luminosity and pulsar energy loss rate, and the comparison with the calculated total energy, a lower limit of 1100 yr is derived for the age of this PWN. ",
        "introduction": "Radio composite supernova remnants (SNRs)  consist of a shell and a spectrally distinct inner nebula, presumably a pulsar wind nebula (PWN), powered by the wind of relativistic electron/positron pairs from a central pulsar. Only in a few cases, however, has the central pulsar been detected \\citep[see][~for a review]{kaspihelfand02}.  Recent observations of several composite SNRs carried out with {\\it Chandra} X-ray Observatory have resolved out on arcsec scales complex structures in the interior of several PWNe. These structures include toroids, axial bipolar jets, wisps, etc. \\citep[][~etc.]{weiss00,helfand01,roberts03}. Images in the different spectral domains are essential to understanding the physics of PWNe. Particularly, the radio emission depends on the history of the nebula and represents the combination of the efficiency of the pulsar in providing accelerated particles and magnetic fields, and the expansion history. The expansion history, in turn, depends on the density and geometry of the medium that confines the relativistic particles and fields (i.e. the interior of the SNR, that includes stellar ejecta and the presence of forward and reverse shocks).  The detailed analysis of the geometry and structure of the PWN and the parent SNR, can shed light on the coupling mechanisms between the neutron star, the relativistic wind nebula and the surrounding SNR plasma.  G0.9+0.1 (RA= 17\\hh 47\\mm 21\\ss, dec= -28\\d 09\\m,  J2000) is a composite SNR located in the direction of the Galactic center and at about the same distance (assumed through this paper to be   8.5 kpc).  It is characterized by a bright, centrally condensed synchrotron nebula, approximately 2$^\\prime$ in size, and a weak surrounding radio shell, about 8$^\\prime$ in size, for which radio spectral indices $\\alpha_{\\rm core}\\sim-0.12$ and $\\alpha_{\\rm shell}\\sim-0.6$ (where S $\\propto \\nu^\\alpha$), have been proposed for the core and shell respectively \\citep{helfbec87,larosa00}.  In the X-rays domain, the first detection was reported by \\citet{helfbec87} based on IPC-{\\it Einstein} observations, who concluded that the observed flux could come either from the compact core or from a combination of core and part of the bright limb of the shell of G0.9+0.1. The  core X-ray emission  was detected  by \\citet{mere98} using {\\it BeppoSAX} satellite. \\citet{sidoli00} later confirmed these results on the basis of better quality data. These early detections are indicative of the presence of a young neutron star powering the nebula, although no coherent pulsations are found.  \\citet{gaensler01} presented the results of 35 ksec ACIS {\\it Chandra} observations of the PWN, between 0.5 and 8.0 keV. From these images, the authors identify a faint semicircular arc and a jet-like feature that define a symmetry axis, which they interpret  as evidence of a torus of emission in the pulsar's equatorial plane and a jet directed along the pulsar spin axis. No X-ray emission is detected in correspondence with the radio-shell nor its interior. Based on these observations the authors propose that the hard point-like X-ray source CXOU J174722.8-280915 detected at energies above 3 keV, is the best candidate for a central pulsar that would be powering the inner nebula. This point source has a rather low ratio of magnetospheric pulsar emission to surrounding nebular emission, with a  luminosity that amounts only 0.5 \\% of the total PWN luminosity in the energy range 2-10 keV.  \\citet{porquet03} carried out an X-ray study of the PWN within G0.9+0.1 using the {\\it XMM-Newton} EPIC-MOS and EPIC-PN cameras. The images obtained in the energy band 1.5-12.0 keV show an amorphous nebula with a bright maximum towards the east surrounded by extended diffuse emission. At the spatial resolution of {\\it XMM-Newton} (8\\arcsec), the arc and jet-like features noticed by \\citet{gaensler01} are not obvious.  The X-ray spectrum within the PWN softens from the core to the outskirts, consistent with synchrotron radiation losses of high energy electrons as they diffuse through the nebula. The {\\it XMM-Newton} study also reveals spectral variations across the ``arc-like feature'' identified by \\citet{gaensler01}, with the eastern part of the arc having clear indications of a very hard photon index ($\\Gamma \\sim 1.0$), opposite to the western part with  a very soft spectrum ($\\Gamma \\sim 3.2$).  \\citet{aharonian05} reported the detection, for the first time, of gamma-ray emission in the direction of G0.9+0.1 at energies greater than 100 GeV at a level of significance of 13 $\\sigma$. The very high energy gamma-rays, discovered using the H.E.S.S instrument, appear to originate in the pulsar wind nebula. The photon spectrum is compatible with a power law with photon index $\\Gamma$ = 2.4.  In radio wavelengths, G0.9+0.1 is prominent at 57.5 MHz and 80 MHz \\citep{larosa85}. It has also been observed at 843 MHz \\citep{gray94}, at 330 MHz as a part of the high-resolution imaging of the Galactic Center region  \\citep{nord04} (Fig.~\\ref{fig:radio90cm}), and at 1.5 GHz and 4.8 GHz  \\citep{helfbec87}.  \\begin{figure} \\centering \\includegraphics[width=8 cm]{figure1.epsi} \\caption{The  SNR G0.9+0.1 at 330 MHz as taken from \\citet{nord04}.} \\label{fig:radio90cm} \\end{figure}   This paper attempts to analyze the morphology and spectral properties of the PWN in G0.9+0.1 in the radio range, based on new high-resolution radio images obtained from observations at 4.8 GHz ($\\lambda$ 6 cm) and at 8.4 GHz ($\\lambda$ 3.6 cm) carried out with the Australia Telescope Compact Array\\footnote{The Australia Telescope Compact Array (/ Parkes telescope / Mopra telescope / Long Baseline Array) is part of the Australia Telescope which is funded by the Commonwealth of Australia for operation as a National Facility managed by CSIRO.} and from reprocessed archival VLA\\footnote{The VLA of the National Radio Astronomy Observatory is a facility of the NSF operated under cooperative agreement by the Associated Universities Inc.}  data at 1.4 GHz ($\\lambda$  20 cm). Also the X-ray emission associated with this nebula has been re-analyzed including new {\\it XMM-Newton} data and reprocessed {\\it Chandra} observations. ",
        "conclusions": "This paper presents new high-resolution and high-sensitivity images of the PWN in the SNR G0.9+0.1 obtained at different radio frequencies. The study is complemented with reprocessed X-ray images based on {\\it XMM-Newton} and {\\it Chandra} data. The new radio images have revealed interesting structures in the nebula, like bright knots, rings and elongated filaments which might be showing instability regions at the sites where the expanding nebula interacts with the surrounding ejecta.  From the comparison of the radio images with the reprocessed X-ray images it is found that the X-ray emitting electrons largely fills the volume delineated by the radio PWN. Also, the new  detailed radio images have confirmed  the symmetry suggested by the {\\it Chandra} X-ray observations, with a bright central band aligned with the X-ray ``toroidal'' feature and a narrow elongated north-south structure that appear as the counterpart of the ``jet-like'' X-ray feature. These good radio/X-ray correspondences are, however, accompanied by notable disagreements, the most important of which is the separation observed between the radio and the X-ray maxima.  Based on the new radio images, with contributions from all spatial scales adequately recovered, we estimated the multispectral flux densities and performed a spectral study. In Table~ \\ref{table:summary} we summarize these results together with other observed and derived characteristic parameters of the PWN in G0.9+0.1. Our study revealed a quite uniform distribution of radio spectral index across the nebula, with only small fluctuations around the mean value of $\\alpha_{\\rm r}= -0.18$.  \\begin{table}[htdp] \\caption{Characteristic parameters of the PWN in G0.9+0.1}  \\begin{center} \\begin{tabular}{r l} S$_{\\rm 3.6 cm}=$ & 1.35$\\pm$0.50 Jy\\\\ S$_{\\rm 6 cm}=$  & 1.45$\\pm$0.23 Jy\\\\ S$_{\\rm 20 cm}=$  & 1.72$\\pm$0.30 Jy\\\\ $\\alpha_{\\rm radio}=$& $-0.18\\pm0.04$\\\\ L$_{\\rm radio}\\simeq$ & $ 1.2 \\times 10^{35}$ erg s$^{-1}$\\\\ $\\nu_{\\rm b}\\simeq$ & $2.4\\times10^{12}$ Hz\\\\ B ~$\\simeq$ & 56 $\\mu$G\\\\ U$_{\\rm particles}\\simeq$ & 5 $\\times 10^{47}$ erg\\\\ U$_{\\rm magnetic}\\simeq$ &  2 $\\times 10^{47}$ erg\\\\ age $\\geq$ &1100 yr\\\\ \\end{tabular} \\end{center} \\label{table:summary} \\end{table}  From the combination of observations in the radio regime with X-ray data we  traced a broadband spectrum which suggests a spectral break at $\\nu_{\\rm b} \\approx 2.4\\times10^{12}$ Hz. On the basis of this information, together with the observed luminosities and the assumption of equipartition between particles energy and Poynting vector energy, we investigated the energetics and the magnetic field  in the nebula."
    },
    "0806/0806.4367_arXiv.txt": {
        "abstract": "We examine a variety of observations that shed light on the orientation of the semi-major axis of the $\\eta$ Carinae massive binary system. Under several assumptions we study the following observations: The Doppler shifts of some He~I P-Cygni lines that is attributed to the secondary's wind, of one Fe~II line that is attributed to the primary's wind, and of the Paschen emission lines which are attributed to the shocked primary's wind, are computed in our model and compared with observations. We compute the hydrogen column density toward the binary system in our model, and find a good agreement with that deduced from X-ray observations. We calculate the ionization of surrounding gas blobs by the radiation of the hotter secondary star, and compare with observations of a highly excited [Ar III] narrow line. We find that all of these support an orientation where for most of the time the secondary$-$the hotter less massive star$-$is behind the primary star. The secondary comes closer to the observer only for a short time near periastron passage, in its highly eccentric ($e\\simeq 0.9$) orbit. Further supporting arguments are also listed, followed by discussion of some open and complicated issues. ",
        "introduction": "\\label{sec:intro}  The $P=5.54 \\yr$ ($P =2022.7 \\pm 1.3~$d; Damineli et al. 2008a) periodicity of the massive binary system \\astrobj{$\\eta$ Car} is observed in the radio (Duncan \\& White 2003), IR (Whitelock et al. 2004), visible (e.g., van Genderen et al. 2006), X-ray (Corcoran 2005), and in many emission and absorption lines (e.g., Damineli et al. 2008a, b). According to most models the periodicity, e.g., of the spectroscopic event and of the X-ray minimum, follows the 5.54~years periodic change in the orbital separation in this highly eccentric, $e \\simeq 0.9$, binary system (e.g., Hillier et al. 2006). The spectroscopic event is defined by the fading, or even disappearance, of high-ionization emission lines (e.g., Damineli 1996; Damineli et al. 1998, 2000, 2008a,b; Zanella et al. 1984). The rapid changes in the continuum, lines, and in the X-ray properties (e.g., Martin et al. 2006,a,b; Davidson et al. 2005; Nielsen et al. 2007; van Genderen et al. 2006; Damineli et al. 2008b; Corcoran 2005) are assumed to occur near periastron passages, although less rapid variations occur along the entire orbit.  It is generally agreed that the orbital plane lies in the equatorial plane of the bipolar structure$-$the Homunculus (Davidson et al. 2001). The inclination angle (the angle between a line perpendicular to the orbital plane and the line of sight) is $i \\simeq 45 ^\\circ$, with $i=41^\\circ-43^\\circ$ being a popular value (Davidson et al. 2001; Smith 2002). However, there is a disagreement about the orientation of the semimajor axis in the orbital plane$-$the periastron longitude. We will use the commonly used periastron longitude angle $\\omega$: $\\omega=0^\\circ$ for a case when the secondary is toward the observer at an orbital angle of $90^\\circ$ after periastron, $\\omega=90^\\circ$ for a case when the secondary is toward the observer at periastron, $\\omega=180^\\circ$ for a case when the secondary is toward the observer at an orbital angle $90^\\circ$ before periastron, and $\\omega=270^\\circ$ for a case when the secondary is toward the observer at apastron, and so on.  While some groups argue that the secondary (less massive) star is away from us during periastron passages, $\\omega=270^\\circ$ (e.g., Nielsen et al. 2007; Damineli et al. 2008b), others argue that the secondary is toward us during periastron passages, $\\omega=90^\\circ$ (Falceta-Gon\\c{c}alves et al. 2005; Abraham et al. 2005; Kashi \\& Soker 2007b [hereafter KS07], who use the angle $\\gamma=90^\\circ-\\omega$). Other semimajor axis orientations have also been proposed (Davidson 1997; Smith et al. 2004; Dorland 2007; Henley et al. 2008; Okazaki et al. 2008a). Abraham \\& Falceta-Gon\\c{c}alves (2007) have obtained the orientation angle in the range $\\omega=60- 90^\\circ$, independent of the orbital inclination. They did not assume that the binary and the Homunculus orbital plane coincide. This contradicts the binary interacting model we support. In addition, Falceta-Gon\\c{c}alves et al. (2005) suggested a model for the X-ray emission with the usage of a problematic expression for the X-ray emission (see Akashi et al. 2006). We therefore will not use the arguments of these two papers to support our claim for $\\omega \\simeq 90^\\circ$.  In the present paper we consider all observations that can shed light on the periastron longitude. As will be shown, all of them support a value of $\\omega \\simeq 90^\\circ$. Namely, the secondary is behind the primary most of the time, and passes in front of the primary for a short time near periastron passage. In section \\ref{sec:doppler} we discuss the Doppler shifts of several lines, in section \\ref{sec:N_H} we discuss the hydrogen column density as deduced from X-ray observations, and section \\ref{sec:narrow} contains a discussion of the variation in the intensity of high excitation narrow lines. Our discussion and predictions are in section \\ref{sec:diss}.  Our assumptions concerning the location where some lines are produced is different from those assumed by other research groups. We therefore present in Figure \\ref{fig:map} a map of the origins of the different lines according to our assumptions. This figure should be consulted when reading the sections to follow. \\begin{figure}[!t]% \\resizebox{0.49\\textwidth}{!}{\\includegraphics{orientation_f1.eps}} \\caption{\\footnotesize A schematic map of the $\\eta$ Car system at apastron. The origins of the different lines are marked on the map. According to our suggested model ($\\omega=90^\\circ$), the observer is on the left side of the system.} \\label{fig:map} \\end{figure} ",
        "conclusions": "\\label{sec:diss}  \\subsection{Main results}  Our goal was to learn about the orientation of the semimajor axis of the $\\eta$ Car binary system from several different observations. We found that an orientation where the hotter secondary star is closer to us for a short time at periastron, i.e., a periastron longitude of $\\omega=90^\\circ$, fits these observations. During most of the time the primary is closer to as, as depicted in Figure \\ref{fig:map}.  In section \\ref{sec:doppler} we attributed the high excitation He~I lines to the secondary wind, and by that could reasonably fit the variation of the Doppler shift with the orbital motion (for more detail see KS07). The Doppler shift of the low excitation Fe~II~$\\lambda6455$ line could be fitted by attributing its origin to the primary stellar wind. The orbital motion explanation for the Doppler shift accounts also for the absence of the Doppler shift variation toward the polar directions. The doppler shift of the Paschen lines could be accounted for if they are assumed to be formed in the shocked primary's wind near the stagnation point (see Fig. \\ref{fig:map}). We also note that the variation in the Doppler shift of the X-ray lines might also be caused by the orbital motion (Behar et al. 2007)  Trying to explain the lines' behavior, we used two sets of primary and secondary masses which we consider possible (see KS08) and obtain nice fits without over-playing with other parameters. We emphasize that claims for exaggerated values for those masses are no problem for the qualitative model we present. Using lower masses we were still able to fit all the lines in the paper by slightly adjusting other parameters (eccentricity, inclination, magnetic-compression, cone opening angle, mass-loss rates, etc.) well within their acceptable range.  In section \\ref{sec:N_H} we examined the hydrogen column density toward the hard X-ray emitting gas, $N_H(>5 \\kev)$. The column density at several times along the orbit is given by Hamaguchi et al. (2007). The column density is sensitive to the mass loss rate and velocity of the primary's wind, and to the nature of the wind interaction process, e.g., where exactly the gas emitting the hard X-ray resides. We cannot reproduce the exact variation of $N_H$ with orbital phase, but could reproduce the approximate value at each phase (Fig. \\ref{fig:NH}).  In an opposite binary orientation, where the secondary is toward us during most of the time (near apastron), the expected value of $N_H$ near apastron is much smaller than the observed value. It is also expected to rise toward periastron by a much larger factor than what is observed. Over all, although our fit is not perfect and requires further work, it has less severe problems than what a model based on an opposite orbital orientation would have.  We did not deal with the value of $N_H(1 \\keV)$ toward the gas emitting the soft X-ray ($\\sim 1 \\keV$), as we are not sure where this gas is located. However, during most of the time $N_H(1 \\keV) \\simeq 0.3 N_H(>5 \\kev)$. This suggests that if it was not for the dense primary's wind around the secondary, the hot X-ray emitting gas would have had lower $N_H$ than the observed value. It is not easy to account for the high $N_H$ value in a model where during most of the time we observed the shocked secondary's wind through the tenuous secondary wind ($\\omega \\sim 270^\\circ$).  In section \\ref{sec:narrow} we calculated the hard ionizing radiation that is emitted by the secondary and reaches the Weigelt blobs (WBs). We found that when the absorption by the undisturbed and shocked primary's wind is considered, the qualitative behavior of the high excitation [Ar III]~$\\lambda 7135$ line that is assumed to be emitted mainly by the WBs (Damineli et al. 2008b) can be reproduced.  The absorption by the shocked primary's wind$-$ in the conical shell$-$depends on the compression of the post-shock gas (Fig. \\ref{fig:WB_D_dependance}). In our model (Kashi \\& Soker 2007a) the compression is constrained by the post-shock magnetic pressure. The post-shock magnetic pressure is determined by the ratio of the pre-shock magnetic pressure to the wind's ram pressure $\\eta_B$. To fit the observed behavior (Fig. \\ref {fig:WB}) we had to assume that the magnetic field in the primary's wind evolves according to equation (\\ref{eq:etaB}). In general, it is expected that the magnetic field in stellar wind will change over time scales of years to tens of years, as it is very well known for our sun.  To summarize, using the model that has been proposed in our study of the radio emission from $\\eta$ Car (Soker \\& Kashi 2007a), we reproduced the basic properties of the ionizing radiation that are required to form the high excitation narrow lines.  \\subsection{Further considerations}  We list four observational results that our proposed $\\omega=90^\\circ$ periastron longitude can account for.  \\subsubsection{The evolution of the radio emission}  The evolution of the radio emission also supports the orientation proposed here. The radio image contains some bright knots, with fainter emission in areas between and around these knots (Duncan \\& White 2003). In particular we note a bright radio knot to the same direction relative to the center as the WBs are. If the orientation was opposite to what we claim here, then the ionizing radiation of the secondary toward this bright region would be constant until $\\la 3$~months before periastron. The reason is that the radiation from the secondary to the knot would propagate through the tenuous secondary's wind. Namely, in the model where the secondary is toward us near apastron the bright radio region would stay more or less at the same maximum brightness until phase $\\sim -0.05$, while region to the sides would decline slowly, as the secondary dives into the denser part of the primary's wind. However, a careful examination of the radio movie (White et al. 2005) shows that the radio knot fades after apastron passage toward the 1998 and 2003 minima, as the rest of the nebula does. This is expected in the periastron longitude $\\omega=90^\\circ$ proposed by us, when the knot is being irradiated through the primary's wind, which continuously becomes denser as the system approaches periastron. This fading of the radio emission as a result of the secondary `diving' into the dense primary's wind has been discussed before by Duncan \\& White (2003) and Abraham \\& Falceta-Gon\\c{c}alves (2007).   \\subsubsection{The purple haze} Smith et al. (2004) presented the ultraviolet images of the Homunculus at 6 epochs around the 2003 minimum. They found that just before periastron the bright UV region extended to the south-east, along the symmetry axis of the Homunculus. Smith et al. (2004) assumed that the UV emitting-reflecting regions are in the equatorial plane, and from that they deduced that the semi-major axis is perpendicular to our line of sight, namely, a periastron longitude of $\\omega=0$.  Abraham \\& Falceta-Gon\\c{c}alves (2007) also attributed the excess UV radiation to the secondary star. Close to periastron this radiation, according to {Abraham \\& Falceta-Gon\\c{c}alves (2007), is confined by dust absorption to the interior region of the colliding-winds cone; the cone points eastward before, and westward after, the event.  We suggest a different interpretation of that behavior. We note that the south-east lobe is toward us, and attribute the UV bright region in the south-east direction to the polar direction, rather than to the equatorial region. For example, a transient polar outflow (Behar et al. 2007) that opens a cone for the ionization and UV illumination might be the cause of this illumination, rather than the orbital orientation at this phase.  We also note the following interesting result: The region of the Weigelt blobs became UV-bright for a short time about a month before periastron passage (Smith et al. 2004). This is at the same time as the peak in the narrow He~I line intensity; the He~I peak is seen a month before the 2003.5 minimum (event 11), but not in the previous two minima (Damineli et al. 2008b). According to the orientation proposed in our model, the secondary is toward this side (closer to us) near that time, and it is possible that for a short time the ionizing radiation from the secondary toward the blobs actually increased, e.g., due to clumpy primary's wind or instabilities an opening in the dense conical shell was formed.  \\subsubsection{The He~I $\\lambda10830$ line}  The He~I $\\lambda10830$ high excitation line has a P Cygni profile with absorption changing from $-640 \\km \\s^{-1}$ to $-450 \\km \\s^{-1}$ (Damineli et al. 2008b). Just before periastron a wide wing in absorption appears, reaching $\\sim - 1000 \\km \\s^{-1}$ a month before periastron, and $\\sim - 1400 \\km \\s^{-1}$ at periastron. This shows that there is a flow of a high velocity gas toward us at periastron. The primary star is not expected to blow such a fast wind toward us, and if it does, why it cannot be seen in other phases as well? We attribute this behavior to the shocked secondary's wind or polar outflow. In our proposed orientation the secondary is closer to us at periastron. The shocked secondary's wind will flow toward us near periastron passage. This is why the wide wing of the He~I $\\lambda10830$ line is seen only close to periastron passage. In any case, as suggested by the recent study of Teodoro et al. (2008), the interpretation of the He~I $\\lambda10830$ requires careful attention, and will be the subject of a forthcoming paper.  \\subsubsection{The location of the Weigelt blobs}  Chesneau et al. (2005) observed the sub-arcsec butterfly-shaped dusty environment surrounding $\\eta$ Car with the VLTI, using the mid-IR instrument MIDI, and the adaptive Optics system NACO. As shown in their figure 7, the lower density `SE filament' (named so by Smith et al. 2003a) is apparently aligned in the same direction as the NW Weigelt blobs complex.  The exitance and the large mass of the Weigelt blobs might be related to their location close to the periastron of the system orbit, where the secondary's wind is strongly disturbed by the primary's dense wind. The counterpart (the SE direction) of this high density region is a very low density region, but has some filaments inside it. In the NE direction has even a lower density. The higher density in the NW might be the consequence of the eccentric orbital motion of the secondary. We predict that 3D numerical simulations using our proposed orientation will reproduced this density asymmetry.  \\subsection{Unclosed issues}  The proposed orbital orientation and our study might have some weak points:  (1)\\emph{Coincidence.} The coincidence that the secondary's wind velocity where the lines are formed, $v_{\\rm zone} \\simeq -430 \\km \\s^{-1}$, is practically the same as the primary wind terminal velocity. Although this is a somewhat weak point in our model, our answer to that coincident (KS07) is that the same can be said about the model where the He~I lines originate in the primary's wind: How come the changes in the velocity of the regions where the lines are formed in the primary's wind exactly mimic the secondary velocity around the center of mass? Another coincident is that the best orientation is exactly for the secondary to be toward us at periastron, and not even several degrees from that direction. However, we found (KS07) that several degrees deviation from $\\omega=90^\\circ$ is possible. Some other weak points regarding the orbital motion interpretation for the Doppler shifts are listed by Davidson (1997). However, like Damineli (1997) he attributed the He~I lines to the primary, while we attribute them to the secondary.  (2) \\emph{Magnetic field.} The absorption of ionizing radiation depends strongly on the magnetic field in the primary's wind. Our need for magnetic field adds a parameter to the model. As it stands now it is a weak point. However, if magnetic fields are detected in the primary wind, then this becomes a strong point. In particular, the magnetic field becomes strong in the shocked primary's wind. We encourage a search for magnetic fields in the shocked primary region, e.g., by looking for its influence on some lines. The strength of the magnetic field near the stagnation point is expected to be $B_{\\rm shock} \\sim 100 (r/1 \\AU)^{-1}$~G, if we assume that the post-shock magnetic pressure is about equal to the wind ram pressure. This field can be detected, e.g., in the Paschen lines very close to periastron passage.  We note the following. Our need for magnetic field comes from the following consideration. If the post-shock primary gas has no magnetic field, it is compressed to a very high density in the conical shell, such that it absorbs too much of the ionizing radiation. Instabilities in the conical shell can reduce this absorption even without magnetic fields, because most of the gas might be concentrated in dense clumps, and the ionizing radiation escapes between this clumps. the study of this process requires 3D numerical simulations.  (3) \\emph{Helium lines from the secondary's wind.} The secondary luminosity is only $\\sim 20 \\%$ of the total luminosity. This might cause some problems in attributing the He~I lines to its wind. However, we note that most of the radiation from $\\eta$ Car is in the IR anyhow. A detail study of the line formation in the secondary is required, similar to the one Hillier et al. (2001, 2006) conducted for the primary star. For the time being we note that many stars similar to the secondary are known to have P-Cygni He~I lines with velocity much smaller than their terminal velocity (see discussion in KS07)."
    },
    "0806/0806.0708_arXiv.txt": {
        "abstract": "{} {We present a method that allows us to estimate the distance from the continuum source located in the center of AGN to the highly ionized gas called warm absorber.} {We compute a set of constant total pressure photoionization models compatible with the warm absorber conditions, where a metal-rich gas is irradiated by a continuum  in the form of a double power-law. The first power-law is hard up to 100 keV and represents a radiation from an X-ray source, while the second power-law extends downwards from several eVs and illustrates a radiation from an accretion disk. } {When the ionized continuum is dominated by the soft component, the warm absorber is heated by free-free absorption, instead of Comptonization, and the transmitted spectra show different absorption line characteristics for different values of the hydrogen number density at the cloud illuminated surface. } {This fact results in the possibility of deriving the number density at the cloud illuminated side from observations and hence the distance to the warm absorber.} ",
        "introduction": "Many Active Galactic Nuclei (AGN) exhibit numerous absorption features of highly ionized material in their UV/X-ray spectra. Such gas, called Warm Absorber (hereafter WA), is located on the line-of-sight towards the observer and it is illuminated by radiation originating from the active nucleus.  Most of the observed absorption lines are blushifted, suggesting that the metal-rich gas is outflowing. We still don't know how such a wind is powered and where it is lunched, nevertheless, due to high resolution spectroscopic observations of AGN from {\\it FUSE}, {\\it Chandra}, {\\it XMM} and other satellites, we are able to make some diagnostics on the physical conditions of the WA \\citep[for a review see:][]{kriss2004,blustin2005,rozanska2007}.  The observed absorption lines velocity shifts are of the order of $10^4$ km s$^{-1}$ in case of UV absorbers \\citep{gabel2003} and  of the order of $10^2$ to $10^3$ km s$^{-1}$ in the case of X-rays \\citep{kaspi2001,kaastra2002}. The column density of the WA is generally estimated to be about $10^{21-23}$ cm$^{-2}$. The absorbing gas comprises different ionization phases, corresponding to temperatures from about $10^5$ K, when iron is partially ionized, up to a few $10^7$ K, when iron is almost completely ionized \\citep{netzer2003,krongold2003,steenbrugge2005}. From the above hints, photoionization modeling of the WA is made, trying to answer to the most important and unsolved question: how far from the continuum source is the absorbing gas located?  However, there is one  difficulty which prevents us from answering this question. It is well known that photoionized models of a cloud illuminated by a single X-ray power-law, typically ranging from 0.01 keV up to 100 keV (hereafter we will call this a hard X-ray illumination)  are degenerate \\citep{rozanska2007}. We cannot distinguish between clouds with the same ionization parameter, but different hydrogen number density $n_0$, at the cloud illuminated surface, and at different locations. The transmission spectrum from a rarefied cloud ($n_0 \\sim 10^5$ cm$^{-3}$) located farther away i.e. at $\\sim 0.1$ pc, is identical to the  spectrum of a dense cloud ($n_0 \\sim 10^{10}$ cm$^{-3}$), situated at $\\sim 0.0001$ pc  from a continuum source. The first case is consistent with the WA being co-spatial with the NLR (narrow line region) or dusty torus, while the second case corresponds to the closest neighbourhood of an accretion disk and BLR (broad line region).  Some estimations on the distance to the WA were done using variability studies  \\citep{netzer2003,krongold2005}. However, using the same 900 ksec {\\it Chandra} data of NGC 3783, \\citet{netzer2003} found upper limits for the location of three ionization phases to be 3.2, 0.63, and 0,18 pc respectively, while \\citet{krongold2005} estimated the location of the high ionization phase at 0.0029 pc and low ionization phase at 0.0004-0008 pc.  In this paper we show that the degeneracy of photoionized models breaks down when the WA is illuminated by an AGN continuum including the thermal disk radiation, represented here as a second power-law component. The disk contribution is particularly important in the case of quasars, since their broad band spectra are clearly dominated by the soft disk emission. In this article we explain why double power-law models are not degenerate and claim that using such models provides a way to determine the distance to the WA in some objects.  The structure of the paper is as follows: In Sec.~\\ref{sec:odl} we explain how to estimate the distance to the WA from photoionized models. Section~\\ref{sec:mod} describes the photoionization models used in this work. Results for different spectral shapes are presented in Sec.~\\ref{sec:res}. Discussion of the models is presented in Sec.~\\ref{sec:dis}, and main conclusions of our work are drown in Section~\\ref{sec:con}. ",
        "conclusions": "\\label{sec:con}  Photoionized models of a WA usually assume that the illuminating continuum has the shape of a  single power-law with the spectral index derived from  X-ray observations of a particular AGN. In this paper, we have shown that it is important to include the broad-band continuum from the active nucleus to achieve a proper modeling of the transmission spectrum through the ionized WA gas.  We have computed a set of models assuming illumination of the metal rich material by a double power-law incident continuum, where one component mimics the emission from an accretion disk, and the second component represents the hard X-ray emission coming most probably from a hot corona, or magnetic flare. This allows to break the degeneracy generally observed in photoionization models using a single hard X-ray illumination. By breaking this degeneracy we observe that models display different ionization and temperature structures for clouds with different number densities at the illuminated surface. Therefore, transmitted spectra look different and in principle it should be possible to identify the photoionized model which better fits a given observation.  We have shown that this degeneracy breaks down due to a switch in the radiative mechanism heating the ionized cloud when $F_X/F_{soft}$ decreases. For a single power-law irradiation, hard photons mostly participate to the Compton heating, and this process becomes responsible for the hot equilibrium temperature of the illuminated layers. When $F_X/F_{soft}$ decreases, i.e. when the soft component starts to dominate, Compton heating is not efficient anymore, since the cloud begins to be more efficiently cooled by Comptonization, and a hot temperature equilibrium is established due to bremsstrahlung absorption.  In conclusion, by taking into account broad band illumination from the active nucleus it is possible to determine the  number density at the illuminated side of WA and thus its distance from the irradiating continuum source. This method is complementary to the method based on variability presented by \\citet{netzer2003} and \\citet{krongold2005}, and can be applied to bright quasars with a strong disk component."
    },
    "0806/0806.0364_arXiv.txt": {
        "abstract": "We have modeled the spatial distribution of luminous X-ray binaries (XRBs) in Milky Way like spiral galaxies with an evolutionary population synthesis code developed by \\citet{Hurley00,Hurley02}. In agreement with previous theoretical expectations and observations, we find that both high- and low-mass X-ray binaries show clear concentrations towards the galactic plane and bulge. We also compare XRB distributions under the galactic potential with dark matter halo and the Modified Newtonian Dynamics (MOND) potential, and suggest that the difference may serve as potential evidence to discriminate these two types of models. ",
        "introduction": "X-ray binaries (XRBs) contain a neutron star (NS) or a black hole (BH) accreting from a normal companion star. They are conventionally divided into low-mass X-ray binaries (LMXBs) and high-mass X-ray binaries (HMXBs) according to the masses of the optical companions \\citep[e.g.][]{verbunt94}. In HMXBs, the evolved (super)giant companions, generally $M_{\\rm optical}\\ga 10M_{\\odot}$, have strong stellar wind mass-loss to power a  bright X-ray source for $\\sim 10^{5}-10^{6}$ yr; while LMXBs, in which $M_{\\rm optical}\\la 1.5 M_{\\odot}$, experience mass transfer via Roche-lobe overflow (RLOF) at a rate of $\\sim 10^{-10}-10^{-8} M_{\\odot}$ yr$^{-1}$. Between them are intermediate-mass X-ray binaries (IMXBs), in which companion stars' masses are in the range $\\sim 2 - 10 M_{\\odot}$ \\citep{heuvel75}. Mass transfer in these binaries often occurs on a (sub)thermal timescale of $\\sim 10^4-10^5$ yr through RLOF.   Using distance estimates and angular distribution of LMXBs, \\citet{paradijs95} and \\citet{white96} investigated the spatial distribution of NS and BH LMXBs in our Galaxy, and suggested that the compact objects had received a kick during the supernova (SN) explosions. More recent work by \\citet{grimm02} using the {\\it RXTE\\/} data showed that HMXBs were concentrated towards the Galactic plane with a vertical scale height of $\\sim 150$ pc while the vertical distributions of LMXBs was significantly broader with a scale height of $\\sim 410$ pc, and the radial distribution of LMXBs peaked strongly at the Galactic bulge. But this sample suffers from some incompleteness of the optical identifications/distance measurements at the large distances from the Sun \\citep{jonker04}. Fortunately, today's sensitive, high-resolution X-ray observations allow the study of luminous XRBs in galaxies even beyond the Local Group, and make it possible to examine XRB populations in a wide range of galactic environments with different star formation histories. For example, {\\it XMM-Newton\\/} and {\\it Chandra\\/} observations of NGC 891, a nearby edge-on spiral galaxy which is very similar in many respects to our own Galaxy, present a straightforward look from outside \\citep{temple05}. The spatial distribution of luminous discrete point sources in this galaxy also shows clear concentrations towards the galactic plane and bulge. From the locations of 154 discrete non-nuclear ultraluminous X-ray sources (ULXs) identified in 82 galaxies observed with {\\it Chandra\\/}, \\citet{swartz04} found that the ULXs in their host galaxies were strongly peaked toward their galaxy centers. Statistical analysis of the X-ray point sources from the {\\it ROSAT HRI} survey of nearby galaxies by \\citet{liujf06} showed that there is a significant concentration of ULXs towards galactic center in late-type galaxies. They also suggested that regular ULXs are likely to be a high-luminosity extension of the ordinary HMXB/LMXB population in late-type galaxies through luminosity function (LF) study.  The spatial distribution of XRBs in a galaxy is determined by the initial kick velocity due to any asymmetry in SN explosion at the birth of a NS/BH, the galactic gravitational potential, and the mass transfer process in a binary. In the present work, we investigated the dynamical consequences of XRBs in spiral galaxies like the Milky Way in an theoretical view. We employed an evolutionary population synthesis (EPS) code to calculate the expected number and luminosity distributions of XRBs in the galaxies. Then, following the approach of \\citet{paczynski90}, we calculated the spatial distribution of XRBs with luminosities $>10^{37}$ ergs$^{-1}$. For the galactic gravitational potential, we adopted both the standard cold dark matter (CDM) model and the Modified Newtonian Dynamics (MOND) model. The objective of this study is to present an integrated picture of XRB distribution in spiral galaxies under the two kinds of galactic potential models, and to explore the difference in the predicted spatial distributions, which could be testified by comparison with future high-resolution observations of XRB distribution in nearby galaxies. A recent related work is to use the detection of LMXBs in the Sculptor dwarf spheroidal galaxy to probe the dark matter halo \\citep{dehnen06}.  This paper is organized as follows. In \\S 2 we describe the population synthesis method and the input physics for XRBs in our model. The calculated results are presented in \\S 3. Our discussion and conclusions are in \\S 4. ",
        "conclusions": "This study shows that, with current understanding of binary evolutions and galactic structure, it is possible to investigate both the luminosity function and spatial distribution of luminous XRBs in nearby galaxies, although the results are subject to many uncertainties and simplified treatments. For example, in our calculations, only primordial binaries were considered while in dense environments like the galactic bulge, dynamical formation channels such as tidal capture, exchange encounters, and direct collisions, may play an important role in binary formation and change the distribution of the XRBs \\citep{voss07}. Additionally, we adopted a simplified radial distribution of newborn binaries in the disk and  in the bulge. The actual initial distribution is related to the structure of the spiral arms, the distribution and evolution of the giant HII regions and warm CO clouds, which are still poorly known. Finally, recent dynamical encounter of galaxies may also lead to the prevalence of ULXs population and make a great influence in their spatial distributions \\citep[e.g.,][]{fabbiano01,wolter03,belczynski04,colbert05,fabbiano06}. Although detail comparison between observations and theoretical predications is not available at present, rough agreement can be obtained. In particular, our calculations show that XRBs in CDM and MOND potentials may have distinct radial distribution around the galactic bulge, suggesting a new way to constrain the nature of DM and test the law of gravity. Our work motivates further efforts to explore the origin of the spatial distributions of luminous XRBs around the galactic center regions."
    },
    "0806/0806.2157_arXiv.txt": {
        "abstract": "We study the mid-infrared properties of \\ngal~ spectroscopically confirmed members in eight massive ($M_{vir}\\gtrsim5\\times10^{14}$\\msun) galaxy clusters covering the redshift range from 0.02 to 0.83. The selected clusters all have deep {\\it Spitzer} MIPS $24\\mu$m observations, {\\it Hubble} and ground-based photometry, and extensive redshift catalogs.  We observe for the first time an increase in the fraction of cluster galaxies with mid-infrared star formation rates higher than 5 \\myr~ from 3\\% at $z=0.02$ to 13\\% at $z=0.83$ ($R_P\\leqslant 1$Mpc).  This increase is reproduced even when considering only the most massive members ($M_{\\ast} \\geqslant 4 \\times 10^{10}$\\msun). The $24\\mu$m observations reveal stronger evolution in the fraction of blue/star-forming cluster galaxies than color-selected samples: the number of dusty, strongly star-forming cluster galaxies increases with redshift, and combining these with the optically-defined Butcher-Oemler members [$\\Delta(B-V)<-0.2$] doubles the total fraction of blue/star-forming galaxies in the inner Mpc of the clusters to $\\sim23$\\% at $z=0.83$.  These results, the first of our {\\it Spitzer}/MIPS Infra-Red Cluster Survey (SMIRCS), support earlier studies indicating the increase in star-forming members is driven by cluster assembly and galaxy infall, as is expected in the framework of hierarchical formation. ",
        "introduction": "\\citet{bo78,bo84} observed that galaxy clusters at intermediate redshift have a higher fraction of members with blue optical colors than clusters in the local universe, thus providing a key piece of evidence supporting galaxy evolution.  This increase in blue members with redshift, named the Butcher-Oemler (BO) effect, was intensely debated for two decades \\citep[e.g.][]{mathieu81,dressler82}. However, multiple optical studies based on spectroscopic observations have since confirmed the increase in blue, star-forming galaxies in higher redshift clusters \\citep[e.g.][]{couch87,caldwell97,fisher98,ellingson01}, and found that BO galaxies reveal signs of recent and ongoing star formation.  The paramount question now is have we seen only the tip of the iceberg?  Most studies of star-forming galaxies in clusters rely on rest-frame ultraviolet or optical tracers \\citep[e.g.][]{balogh98,poggianti06}, but UV/optical tracers can suffer severely from dust obscuration, especially when star formation is concentrated in the nuclear regions \\citep{kennicutt98}.  For example, ultraluminous infrared galaxies have SF rates of $\\gtrsim1000$\\msun, yet many ULIRGs fail to even be detected at UV and optical wavelengths \\citep[e.g.][]{houck05}. Although corrections for dust attenuation are possible, reliable estimates of SF rates cannot be achieved solely using rest-frame UV/optical observations \\citep{bell02,cardiel03}.  A substantially more robust method of determining total SF rates is with mid-infrared (MIR) imaging.  The first MIR imaging of galaxy clusters at intermediate redshifts was taken with ISO's ISOCAM camera, and \\citet{duc02} found that at least 90\\% of the star formation was hidden at optical wavelengths.  The first handful of galaxy clusters observed with the MIPS camera on the {\\it Spitzer Space Telescope} (SST) have also revealed strong dust-obscured star formation \\citep{geach06,marcillac07,bai07}.  It remains unclear as to what causes the increase in star-forming galaxy cluster members.  Detailed morphological studies of blue galaxies [defined as having $\\Delta(B-V)<-0.2$]\\footnote{$\\Delta(B-V)$ is the color offset from the red sequence fit to the cluster ellipticals.} with the {\\it Hubble Space Telescope} (HST) find that most are disk systems similar to those in local clusters \\citep[e.g.][]{dressler94,couch94}; past studies also find that many show signs of interactions or mergers \\citep{lavery88,lavery92,couch94,oemler97}.  More recently, studies indicate that galaxy infall is a viable explanation for the significant numbers of blue galaxies and their disturbed morphologies in intermediate redshift clusters \\citep[e.g.][]{vandokkum98,ellingson01,tran05}, a scenario supported by hierarchical clustering models \\citep{kauffmann95}.  In this case, galaxy clusters that are accreting a significant number of new members should have a higher fraction of star-forming galaxies, especially at higher redshifts when the amount of activity was enhanced also in the field.  Here we present the first comprehensive study of SST/MIPS $24\\mu$m imaging of galaxies that are spectroscopically confirmed members of eight massive ($M_{vir}\\gtrsim5\\times10^{14}$\\msun) X-ray luminous clusters spanning a wide redshift range ($0.02<z<0.83$). After presenting the data in \\S2, we focus our analysis in \\S3 and \\S4 on the evolution of star-forming members with redshift.  A cosmology with $(H_0,\\Omega_{M},\\Omega_{\\Lambda})=(70$ \\kms$, 0.3,0.7)$ is assumed throughout the paper; at $z=0.83$, the look-back time is $\\sim7$ Gyr.\\\\ ",
        "conclusions": "Having established an increase in the fraction of MIR-detected galaxies from $z\\sim0$ to $z\\sim0.8$, we stress that optical studies are likely underestimating the increase in star-forming cluster galaxies with redshift.  As seen in Fig.~\\ref{fig1}, an increasing number of strong dust-obscured star-forming members appear on or near the red sequence at higher redshifts; these are not included in traditional color-selected BO studies.  The late-type morphologies of these members supports our intepretation of dusty star formation and red colors due to extinction \\citep[see also A901/902;][]{wolf05}.  Using the standard BO definition of $\\Delta(B-V)<-0.2$, the fraction of blue galaxies with $M_B\\leq-19.5$ and $R_P<1$~Mpc at $z\\sim0.8$ is $\\sim11$\\%; however, including the red, massive, star-forming members raises the total fraction of blue/star-forming members to $\\sim23$\\%.  We note that for the five main clusters, the increase in the blue/star-forming fraction due to these red, star-forming members is \\{1.1, 1.2, 1.3, 1.7\\} at $z=$\\{0.02, 0.33, 0.59, 0.83\\}, $i.e.$ the relative importance of including red, dusty star-forming members increases with redshift.  Is this increase linked to galaxy infall?  In Figure~\\ref{fig2}, both CL0024 ($z\\sim0.4$) and MS2053 ($z\\sim0.6$) are above the general trend established by the other six clusters.  Both CL0024 and MS2053 have enhanced star formation compared to other clusters at similar redshift, and both have bimodal redshift distributions.  CL0024 is made of two colliding subclusters \\citep{czoske02}, and has an unusually large number of luminous infrared galaxies \\citep{coia05}.  Similarly, \\citet{tran05} conclude from that MS2053 has a significant number ($>25$\\%) of infalling galaxies; these members tend to be blue and star-forming. Both CL0024 and MS2053 are accreting a large number of new members and have high fractions of dusty star-forming galaxies.  We speculate that the increase in star-forming members reflects the recent accretion of new members, $i.e.$ galaxy infall, and that such events are more frequent at higher redshift due to the process of cluster assembly \\citep{ellingson01,tran05,loh08}. As further evidence of this, 80\\% of the MIPS-detected galaxies in the $z\\sim0.8$ clusters are more than 700 kpc from the cluster cores in projected distance and thus the MIR Butcher-Oemler effect is significantly altered by only considering the inner 500 kpc of the clusters (open symbols in Fig.\\ref{fig2})."
    },
    "0806/0806.1648_arXiv.txt": {
        "abstract": "{Strong observational evidence indicates that all extragalactic jets associated with AGNs move at relativistic speed up to $100$ pc - 1 kpc scales from the nucleus. At larger distances, reflecting the Fanaroff-Riley radio source classification, we observe an abrupt deceleration in FR-I jets while relativistic motions persist up to Mpc scale in FR-II. Moreover, VLBI observations of some object like B2 1144+35, Mrk501 and M87 show limb brightening of the jet radio emission at the parsec scale. This effect is interpreted kinematically as due to the presence of a deboosted central spine at high Lorentz factor and of a weakly relativistic external layer. } {In this paper we investigate whether these effects can be interpreted by a breaking of the collimated flow by external medium entrainment favored by shear instabilities, namely Kelvin-Helmholtz instabilities. We examine in details the physical conditions under which significant deceleration of a relativistic flow is produced.} {We investigate the phenomenon by means of high-resolution three-dimensional relativistic hydrodynamic simulations using the PLUTO code for computational astrophysics.} {We find that the parameter of utmost importance in determining the instability evolution and the entrainment properties is the ambient/jet density contrast. We show that lighter jets suffer stronger slowing down in the external layer than in the central part and conserve a central spine at high Lorentz factor.} {Our model is verified by constructing synthetic emission maps from the numerical simulations that compare reasonably well with VLBI observations of the inner part of FR-I sources.} ",
        "introduction": "Extragalactic radio sources are traditionally divided in two morphological classes according to their intrinsic power \\citep{FR74}: low luminosity sources (Fanaroff-Riley type I, FR-I) are brighter close to the nucleus of the parent galaxy and their jets become dimmer with distance, while high power sources (Fanaroff-Riley type II, FR-II) show the maximum brightness in the hot spots at the jet termination. The different morphology is generally accepted to reflect a difference in how the jet energy is dissipated during propagation in the extragalactic medium and produces the observed radiation. For FR-I sources, it was quickly accepted that entrainment and deceleration of the jet must play an important role in shaping their morphology \\citep{Bicknell84, Bicknell86, DeYoung96, Komissarov94}, while in FR-II sources energy and momentum are transported without losses to the front working surface.  More recently a large body of evidence has accumulated showing that jets are relativistic at their base, not only in FR-II radio sources but also in FR-I. Superluminal motions are observed on milliarcsecond scales in several FR-I jets \\citep{Giovannini01} and on arcsecond scales in M87 and Cen A \\citep{Biretta95, Hardcastle03} and finally, on small scales, there are also observations of one-sidedness and brightness asymmetry between jet and counter-jet, due most likely to Doppler boosting effects \\citep{Laing99}. FR-I sources are also thought to be the parent population of BL Lac objects, for which the presence of relativistic velocities on parsec scales is well established \\citep{UP95}.  The Lorentz factors of the jet bulk motion at sub-pc scales cannot be deduced directly from the observations, but are inferred from assumption on the physical emission mechanism. \\citet{Harris06} in their review indicate that FR-II radiogalaxies have jets with bulk Lorentz factors between 5 and 40 \\citep[see also][]{UP95} and that FR-I jets have less constrained values somewhat lower than those of the FR-II jets. \\citet{Giovannini01} using data on proper motions and  brightness ratio between jet and counter jet in a sample of radiogalaxies  do not  find any  systematic difference between low and high power radio sources, and the values they derive are between 3 and 10. \\citet{CG08} derive  jet physical properties modelling the spectral energy distribution in a sample of blazars and  find no difference between BLLac's (associated to low power radio sources in the unified model) and radio-quasars (associated to high power radio sources), with an average value of about 15.   On the other hand, relativistic motions from the inner regions all the way to larger scales, with Lorentz factors of about ten, appear  to be present in powerful FR-II jets, as indicated by Chandra discovery of bright X-ray emission at kpc scales \\citep{Tavecchio04, Harris06}; instead sub-relativistic velocities are found at kpc scales \\citep{Bicknell94} in low power radio-sources, and the decrease of brightness asymmetry along the jet suggests that a deceleration must occur \\citep{Laing99, Laing02, Laing05}.  Both these morphological and kinematical data are indicative of the interaction between collimated outflows and the surrounding medium. In particular jet deceleration can be obtained by redistributing the bulk momentum through some form of mass entrainment  \\citep{Bicknell94, Bicknell95}. Velocity shear instabilities are the most likely triggering mechanisms of entrainment \\citep{DeYoung96, DeYoung05}. \\footnote{An alternative and most likely complementary view for the origin of the entrainment has been presented by \\citet{Komissarov94} and \\citet{Komissarov96}, who consider the possibility of entrainment by injection of mass lost by stars within the jet volume. } In fact the nonlinear development of velocity shear or Kelvin-Helmholtz instabilities leads to an exchange of mass, momentum and energy at contact discontinuities between fluids in relative motion \\citep{khcyl94, slab, jet3d}; understanding the details of this process is essential for modeling the jet dynamics and for giving clues on the determination of the jet physical parameters. One of the key features of this form of interaction is the formation of a mixing layer at the interface between jet and surrounding medium, where the external material is entrained and accelerated at the expenses of the jet momentum. This process leads to the formation of a transverse velocity profile where the internal layers feel less the effects of the interaction with the ambient medium and keep an higher velocity, while the external layers are more decelerated. The presence of such transverse velocity structure has been already suggested for explaining some of the observational properties of radio sources such as their magnetic field configuration \\citep{Komissarov90, Laing93}, limb brightening effects \\citep{Giroletti04} and to overcome problems in unifying radiogalaxies with BL Lac objects \\citep{Chiaberge00}.  For studying in detail the instability evolution and the subsequent entrainment process one has to resort to three-dimensional numerical simulations, since the mechanisms at the base of the entrainment are inherently three-dimensional, as shown in preliminary analyses by \\cite{BodoNAR03} and \\cite{Rossi04}.  In this paper we present results of high-resolution hydrodynamic simulations in which we follow the evolution of a perturbed relativistic jet as it propagates in a homogeneus stationary ambient medium. We assume that the jet is in pressure equilibrium with the outside medium, although this choice is not critical for the final results. Perturbations grow as a consequence of the velocity shear instabilities and lead to entrainment of external medium and to jet deceleration. The main questions we address are the dependence of the entrainment and deceleration processes on the jet physical parameters and the kind of structure that the jet acquires as the result of these processes.  The plan of the paper is the following: in \\S\\ref{sec:numsetup} we present the numerical setup adopted for calculations and the parameter space covered by the simulations; in \\S\\ref{sec:results} we discuss the results of our simulations focusing our attention on the dependence of the efficiency of deceleration on the physical parameters; in \\S\\ref{sec:entr} we analyze in more detail the entrainment properties for the case that appears more successful in decelerating the jet. Preliminary comparisons of our simulations with FR-I source of different morphologies are reported in \\S\\ref{sec:implic} and conclusions of our study are drawn in \\S\\ref{sec:disc}. ",
        "conclusions": "\\label{sec:disc}  In this paper we have presented the 3D nonlinear dynamical evolution of relativistic light jets, as a result of a perturbation introduced at the jet inlet. The perturbation grows because of KHI and gives rise to a strong interaction of the jet with the external medium with a consequent mixing and deceleration. The two main parameters controlling the jet dynamics are the Mach number $M$ and the density ratio $\\eta$ between the ambient medium and the jet. The Lorentz factor has been set equal to $10$ in all computations.  We have explored the parameter plane ($M, \\eta$), finding that the deceleration becomes more efficient increasing $\\eta$. A preliminary analysis of the parameter space suggests that only jets with a large density ratio ($\\eta > 10^{2}$) can undergo appreciable deceleration, while the Mach number does not seem to play a fundamental role in this respect.  We have focused our attention on three extreme cases in the parameter plane, namely: case A with $M=3$ and $\\eta=10^2$, B with $M = 3$ and $\\eta = 10^4$ and E with $M = 30$ and $\\eta = 10^2$. The comparison of these simulations show in fact that case B (and D) undergoes the strongest deceleration. Cases A  and E retain a high Lorentz factor spine with propagation velocities essentially unchanged from the injection value, while some deceleration has been observed only in the outer layers with the formation of a wide transverse velocity structure. In case B  we observe the formation of a similar pattern although on a much shorter distance and with a significant stronger decrease of the maximum Lorentz factor.   The fact that larger values of $\\eta$ (i.e. lower jet densities) lead to prominent deceleration may have direct astrophysical implications. Observational data seems to indicate that the jet kinetic power associated with FR-I radio sources is, on the average, $\\sim 10^3$ times lower than in FR-II radio sources \\citep{Celotti03}. On the other hand, there seems to be no difference in the value of the initial Lorentz factor in the two classes \\citep{Giovannini01, CG08}. Since lighter jet beams imply reduced jet kinetic powers, our model leaves the density contrast as the most likely candidate to account for the discrepancies in the deceleration process efficiency.  Using some astrophysically relevant units we can rescale our models and come up with a rough estimate for the critical value of the jet power $P^{*}_{j}$ that separates FRI from FRII radiosources: \\begin{equation} P^{*}_{j}  \\sim 1.3 \\times  10^{44} \\left( \\frac{r_{j}}{1 pc} \\right)^{2} \\left( \\frac{\\gamma_{b}}{10} \\right)^{2} \\left( \\frac{n}{1 cm^{-3}} \\right) \\left( \\frac{\\eta^{*}}{10^{3}} \\right)^{-1} {\\rm erg s^{-1}} \\end{equation} where we assumed a jet radius of $1 pc$, an external density of $1 cm^{-3}$ which is typical at distances below $100pc$ \\citep{BBC08}, and a critical density ratio separating the two behaviors of order $10^{3}$. The value of $P_{j}^{*}$ is affected by many uncertainities but is in agreement, for example, with the estimates given by \\citet{CG01} based on the results by \\citet{W99}, who gave for $P_{j}^{*}$ the value \\begin{equation} P_{j}^{*} \\sim 10^{44} \\left( \\frac{M_{BH}}{10^{8} M_{\\sun}} \\right) {\\rm erg s^{-1} } \\end{equation} where $M_{BH}$ is the mass of the central black hole.    Considering the case with $\\eta = 10^{4}$  we obtained a spine-layer structure similar to that deduced from observations. Looking at the synthetic maps produced from the simulations, it is evident that many of the salient features are fairly well reproduced. The main difference lies in the terminal Lorentz factor of the slow layer, typically smaller than the observational estimates. In view of this first promising result we intend to further pursue this investigation, trying to better constrain jet parameters and introducing another essential ingredient: the magnetic field."
    },
    "0806/0806.0565_arXiv.txt": {
        "abstract": "{} {This paper is a continuation of an ongoing study of the evolutionary processes affecting cluster galaxies.} {Both CCD R band and H$\\alpha$ narrow-band imaging was used to determine photometric parameters ($m_{r}$, $r_{24}$, H$\\alpha$ flux, and equivalent width) and derive star formation rates for 227 CGCG galaxies in 8 low-redshift clusters. The galaxy sample is a subset of CGCG galaxies in an objective prism survey (OPS) of cluster galaxies for H$\\alpha$ emission.} {It is found that detection of emission-line galaxies in the OPS is 85\\%, 70\\%, and 50\\% complete at the mean surface brightness values of $1.25\\times 10^{-19}$, $5.19\\times 10^{-20}$, and $1.76\\times 10^{-20}$ W m$^{-2}$ arcsec$^{-2}$, respectively, measured within the R band isophote of 24 mag ${\\rm arcsec}^{-2}$ for the galaxy.} {The CCD data, together with matched data from a recent H$\\alpha$ galaxy survey of UGC galaxies within $v \\le 3000$ km ${\\rm s}^{-1}$, will be used for a comparative study of R band and H$\\alpha$ surface photometry between cluster and field spirals.} ",
        "introduction": "\\label{intro}  While the transformation of cluster disc galaxies from predominantly spiral to mainly lenticular galaxies over the past $\\sim$ 5 Gyr is well established (e.g. Butcher \\& Oemler 1978, 1984; Dressler et al. 1997; Fasano et al. 2000), the mechanism or mechanisms that have affected this transformation are not so clear. However, a comparative study of the rate, distribution, and morphological dependence of star formation between cluster and field spirals appears to be a promising enquiry that can help to disentangle some of the suggested transformation processes. For example, ram-pressure stripping of the cold interstellar gas of spirals by the hot ionised intracluster medium (e.g. Gunn \\& Gott 1972; Quilis, Moore \\& Bower 2000) should be most effective in the centres of rich clusters, and may lead to a rapid truncation of the star forming disc, but provides no obvious mechanism to promote circumnuclear star formation. On the other hand, strangulation, i.e. the stripping of an hypothesised hot halo gas of spirals (e.g. Larson et al. 1980; Bower \\& Balogh 2004) should be a more gradual process; simulations by Bekki et al. (2002) have shown that this would lead to anemic spirals rather than truncation. Tidal interactions with the cluster potential can induce star formation across both bulge and disc (e.g. Byrd \\& Valtonen 1990), whereas low-velocity interactions between galaxies can be efficient at triggering star formation in central regions (e.g. Kennicutt et al. 1987; Mihos et al. 1992; Iono et al. 2004). In contrast, galaxy harassment, i.e. frequent galaxy high-speed encounters within a cluster, are expected to trigger modest disc-wide response of star formation for giant spirals (see Moore et al. 1999; Mihos 2004).  Systematic comparative studies of the massive star formation properties of cluster and field galaxies have already produced interesting results.  Moss \\& Whittle (2000, 2005) undertook an objective prism survey (OPS) of a complete magnitude-limited sample of 727 CGCG galaxies (Zwicky et al. 1960--68) in 8 low-redshift clusters. These authors show an enhancement of circumnuclear starburst emission in cluster spirals associated with a disturbed morphology that is attributed to slow galaxy--galaxy encounters and major and minor mergers (see also Moss 2006).  Koopmann \\& Kenney (2004a,b) have pioneered a comparative study of the massive star formation properties of Virgo cluster and isolated bright ($M_{\\rm B} < -18$) S0--Scd galaxies via analyses of R and H$\\alpha$ surface photometry.  They show that the median total normalised massive star formation rate is reduced by a factor of 2.5 for cluster galaxies as compared to the field. Few of the cluster or isolated galaxies are anemic, suggesting that strangulation is not a major contributory factor in the reduced star formation rates of Virgo spirals; rather, this reduction is caused by spatial truncation of the star forming discs. In addition, several of the truncated galaxies show evidence of recent tidal interaction or minor mergers, such as enhanced central star formation rates and disturbed stellar discs.  It is intended to extend the analyses of Koopmann \\& Kenney to the clusters studied by Moss \\& Whittle using R band and narrow-band H$\\alpha$ imaging obtained for a 227 subset of CGCG galaxies of mainly types Sa + later in the OPS. Clusters in the OPS include those of greater central galaxy density (most especially Abell 1367 and the Coma cluster) where on-going environmental effects on galaxy morphology and transformation may be expected to be even more pronounced than for the Virgo cluster. In the present paper, we discuss observational data and data reduction for imaging data, and present global photometric properties and derived star formation rates. Completeness limits for the OPS are also determined.  A second paper (Bretherton et al., in preparation) will give results of a comparative study of R band and H$\\alpha$ surface photometry between the sample of 227 CGCG cluster galaxies, and sets of galaxies, matched according to morphology and absolute magnitude to the cluster sample, taken from the recent H$\\alpha$ galaxy survey of UGC galaxies within $v \\le 3000$ km ${\\rm s}^{-1}$ (H$\\alpha$GS, Shane 2002; James et al. 2004).  In the present paper, sample selection and observations of the cluster data are discussed in section \\ref{sselect}. The data reduction procedures and photometry are outlined, and global parameters derived for all sample galaxies, in sections \\ref{datred} and \\ref{glp} respectively. Section \\ref{complet} uses a complete sample of Sa--Sc galaxies within the cluster data to investigate the completeness of the objective prism survey (OPS, see Moss \\& Whittle 2000, 2005) on which present sample selection is based. Conclusions of this paper are given in section \\ref{concl}. ",
        "conclusions": "\\label{concl}  H$\\alpha$ and R band continuum CCD observations have been completed for a sample of 227 CGCG galaxies associated with 8 low-redshift Abell clusters, which were the subject of an objective prism survey (OPS) by Moss and collaborators (Moss et al. 1998; Moss \\& Whittle 2000, 2005). The sample galaxies were generally restricted to those with velocities within 3$\\sigma$ of the cluster mean, and known AGN have been excluded. R band magnitudes, H$\\alpha$ fluxes and EWs, and star formation rates for the sample are listed in Table \\ref{gparams}.  The dominant constraint on the detection efficiency of emission-line galaxies (ELGs) by the OPS is shown to be H$\\alpha$ surface brightness. Detection of ELGs is 85\\%, 70\\%, and 50\\% complete at the mean surface brightness values of $1.25\\times 10^{-19}$, $5.19\\times 10^{-20}$, and $1.76\\times 10^{-20}$ W m$^{-2}$ arcsec$^{-2}$ respectively, where the mean H$\\alpha$ surface brightness was measured within the R band isophote of 24 mag ${\\rm arcsec}^{-2}$ for the galaxy.  The present data, together with matched sets of data from a recent H$\\alpha$ galaxy survey of UGC galaxies within $v \\le 3000$ km ${\\rm s}^{-1}$ (H$\\alpha$GS, Shane 2002; James et al. 2004) will be used for a forthcoming comparative study of R band and H$\\alpha$ surface photometry between cluster and field spirals (Bretherton et al., in preparation)."
    },
    "0806/0806.2560_arXiv.txt": {
        "abstract": "We discuss the differences and analogies of gravitational clustering in finite and infinite systems. The process of collective, or violent, relaxation leading to the formation of quasi-stationary states is one of the distinguished features in the dynamics of self-gravitating systems. This occurs, in different conditions, both in a finite than in an infinite system, the latter embedded in a static or in an expanding background. We then discuss, by considering some simple and paradigmatic examples, the problems related to the definition of a mean-field approach to gravitational clustering, focusing on role of discrete fluctuations. The effect of these fluctuations is a basic issue to be clarified to establish the range of scales and times in which a collision-less approximation may describe the evolution of a self-gravitating system and for the theoretical modeling of the non-linear phase. ",
        "introduction": "As discussed in various papers in this volume (see e.g.\\cite{intro,campa}) equilibrium properties of long-range interacting systems require a non-trivial analysis as standard thermodynamics techniques do not simply apply when dealing with pair-interactions decaying with sufficiently small exponents.  Many interesting and unsolved problems lie in the out-of-equilibrium dynamics of systems with long-range interaction about which very little is known from a theoretical point of view (see also \\cite{chavanis} in this volume). The understanding of the thermodynamics and dynamics of systems of particles interacting only through their mutual Newtonian self-gravity is of fundamental importance in cosmology and astrophysics. It encompasses the range of physical scales relevant to the formation of the largest structures in the Universe, down to those relevant to stellar dynamics. The statistical mechanics of systems dominated  by gravity has been studied and applied in many different contexts in astrophysics and cosmology (see e.g. \\cite{lyndebell,pad_physrep,chavanis,chandra,chandra_revmodphy, pad_book,pad_dtslri,saslaw1,saslaw2,binney,peebles,thierry}): for example in the studies of globular clusters, galaxies and the clustering in the expanding universe.   While systems with short range interactions can be usually studied through laboratory experiments, gravitational systems can only be observed in astrophysical contexts. Alternatively one may set up numerical experiments which then represent the unique instrument to study the dynamics of gravitational clustering. In this respect the astrophysicist's perspective is usually to model some intricate realistic systems, such as stellar or galaxy systems, having the aim of understanding a specific set of observations.  For example in the cosmological context one uses very complicated initial conditions (described by a large number of parameters) and needs a certain number of important assumptions, from the way the universe expands to the amount and type of dark matter which dominates the dynamics on the relevant scales. This is so because, by studying gravitational clustering, one would like to understand the relations between some important observations of the cosmos. For example the studies of the cosmic microwave background radiations provide with the information about the initial conditions of the matter density field. The large scale geometrical properties of the universe are deduced, for example, through the measurements of the supernovae magnitude-redshift relation. Galaxy redshift surveys map the present-day matter distribution. The estimations of the mass-to-light ratio of astrophysical objects is ultimately related to the abundance of dark matter. The task of the model of cosmological structure formation is thus to build a unified and coherent picture to explain these (and other) observations of the universe at the largest scales \\cite{peacock}.   In statistical physics the problem of the evolution of self-gravitating classical bodies has been relatively neglected, primarily because of the intrinsic difficulties associated with the attractive long-range nature of gravity and its singular behavior at vanishing separation.  When approaching the problem of gravitational clustering in the context of statistical mechanics it is natural to start by reducing as much as possible the complexity of the analogous cosmological or astrophysical problem. For example, in order to focus on the essential aspects of the problem one may study gravitational clustering without the expansion of the universe, and starting from particularly simple initial conditions. With respect to the motivation from cosmology/astrophysics, there is of course a risk: in simplifying we may loose some essential elements which change the nature of gravitational clustering. Even it were, it seems unlikely that we will not learn something about the more complex cosmological/astrophysical situations  in addressing slightly different and simplified problems.    A fundamental distinction has to be made between finite and infinite systems.  They have in common that the gravitational force on a arbitrary point has contributions coming from all scales in the system. However they differ for the fact that in the case of the finite system there is a mean field force generated by the system as a whole, which is related to its internal symmetries (e.g., spherical symmetry) and which eventually will give rise to a global collapse of the entire system. In an infinite space, in which the initial fluctuations are non-zero and finite at all scales, the collapse of larger and larger scales will continue ad infinitum: being no geometric center there will not be a global collapse of the entire system. The mean field dynamics is driven by system's fluctuations which determine the gravitational force at different spatial scales. The collapse occurring on larger and larger scales will clearly happen at different times but it is characterized by the unique time-scale in the system that is \\be \\label{tau} \\tau \\sim \\sqrt{G\\rho_0}^{-1} \\,. \\ee This is in general the typical characteristic time scale of any (finite or infinite) gravitational system with average mass density $\\rho_0$. For instance, as we discuss below, this is the time scale predicted by the self-gravitating fluid approximation.  The infinite system can therefore never reach a time independent state, and in particular it will never reach a thermodynamic equilibrium. Although so different, the finite and the infinite systems share some subtle and important analogies which we briefly discuss in what follows.  We firstly discuss the general difficulties related to the long-range character of gravity and the usual way to make a mean-field approximation. Then we consider two basic examples of a finite and of an infinite system: the simplest example of a finite system is represented by an initially isolated spherical distribution of randomly distributed points (i.e. a Poisson distribution) --- this will be also discussed in the contribution by Morikawa in this book \\cite{morikawa}.  Analogously the simplest example of an infinite system is a Poisson distribution in an infinite space. For this second case one may consider that the space background is static (as Joyce in his contribution in this book \\cite{joyce}) or it is expanding (as Saslaw in his contribution \\cite{saslaw}). We will discuss these two cases briefly, outlining the analogies and the differences. In the conclusions we try to point out which are the main problems of the field. ",
        "conclusions": "\\label{conclusions}  The dynamics of infinite self-gravitating systems is a fascinating theoretical problem of out of equilibrium statistical mechanics, directly relevant both in the context of cosmology/astrophysics and, more generally, in the physics of systems with long-range interactions.  We discussed some of the many problems encountered in the study of the gravitational clustering in both finite and infinite systems.  We would like to stress two important open problems.  The first concerns the extent to which such numerical simulations of a finite number of particles, reproduce the mean-field/Vlasov limit which is usually used to describe the evolution from a theoretical point of view. That is, the theoretical question that arises is about the validity of this collisionless limit.  Another major question is that of the understanding of halo structure observed in simulations of infinite systems: while these show strongly universal characteristics, their dynamical origin is not yet understood from a theoretical point of view.    \\begin{theacknowledgments} I wish to thank Bruno Marcos, Michael Joyce and Andrea Gabrielli for useful comments and fruitful collaborations on the subject. I also thank Bill Saslaw for the very many discussions over the years we had together on this topic. \\end{theacknowledgments}"
    },
    "0806/0806.1755_arXiv.txt": {
        "abstract": "We present a preliminary calibration and flight performance of the Long-Slit Imaging Dual Order Spectrograph (LIDOS), a rocket-borne instrument with a large dynamic range in the 900 - 1700 \\AA\\ bandpass. The instrument observes UV-bright objects with a CCD channel and fainter nebulosity with an MCP detector. The image quality and the detector quantum efficiencies were determined using the calibration and test equipment at the Johns Hopkins University, and further monitored using an on-board electron-impact calibration lamp. We review results from each of the three flights of the instrument. ",
        "introduction": "\\label{sec:intro}  %  One of the ongoing problems in astrophysics is to constrain the interactions between hot young stars and the nebular material in which they form. Characterization of the properties of gas and dust in such environments using spectroscopy is enhanced by knowledge of the spectral energy distribution of the illuminating star(s) in the far-ultraviolet (far-UV) bandpass (900 -- 1650 \\AA). Such knowledge allows us to quantify the relationship between  extinction of the far-UV radiation field by dust and molecular hydrogen (H$_{2}$) fluorescence by simultaneously measuring the spectra of the exciting stars and the scattered light from the surrounding nebular material. The Long-slit Imaging Dual Order Spectrograph (LIDOS) is uniquely equipped for this task, employing a charge-coupled-device (CCD) channel to observe UV-bright objects and an microchannel plate (MCP) detector for the fainter emission. The long-slit configuration combined with the resolving power of the telescope, allows for the extraction of information as a function of the angular offset from the point source. The instrument is designed to be flown on a sounding rocket in order to avoid UV attenuation by the atmosphere. LIDOS has flown three times and has obtained data on the $\\gamma$~Cas/IC~63 system (36.208UG, December 2003), the Trifid Nebula (36.220UG, August 2007) and the Orion Nebula (36.243UG, January 2008). The targets were chosen to provide a proving ground for the instrument, with stars that would have exceeded the brightness limits of the $Far-Ultraviolet$ $Spectroscopic$ $Explorer$ ($FUSE$). All flights were successful, with the MCP collecting data on all three objects, and the CCD acquiring the spectrum of $\\gamma$~Cas\\cite{France:05} and $\\theta^1$~Ori~C, the brightest star in the central region of the Orion nebula. All flights were preceded and followed by refurbishing and end-to-end characterization of the instrument in the facilities developed at the Johns Hopkins University. The status of the instrument was monitored during all phases of the integration through the use of an on-board e$^{-}$-impact lamp\\cite{McCandliss:03} mounted on the side of the spectrograph and a similar lamp on the shutter door of the telescope.  The scope of this paper is to review the characteristics of the LIDOS instrument as measured during flight and during extensive laboratory calibrations, and to show the preliminary results obtained from the flight data. It serves as a performance study of the LIDOS design and calibration techniques for future UV instruments, and as a reference for the data reduction. A brief description of the instrument is given in Section 2, followed by the description of the calibration procedures and the derived instrument characteristics in Section 3. The flight performance is presented in Section 4 and the flight results follow in Section 5. ",
        "conclusions": "The field of far-UV spectroscopy is continuously seeking innovative designs to maximize the throughput and dynamic range for astrophysical observations. As proven by the success of the $FUSE$ and the $Galaxy$ $Evolution$ $Explorer$, many research areas benefit form the understanding of molecular and atomic signatures in this bandpass, including Solar System objects, star formation, galaxy evolution, and cosmic reionization. A good handle on the processing of UV radiation in the local universe translates into better predictions and observing plans for objects at high redshift. Obtaining high quality data in this spectral range is limited not only by the necessity of using space-based instruments, as to avoid atmospheric absorption, but also by the low efficiencies of the detectors and coatings available. By incorporating new design concepts, the LIDOS instrument has given us the possibility to test and characterize recent developments in technology within the limited resources available for a rocket flight.  The comprehensive analysis of the performance and characteristics of the LIDOS instrument has shown a good agreement with the design expectations. The total effective area and image quality was limited by the performance of the telescope. The pre-flight calibrations have been essential in determining the observing plan and setting the working parameters for the detectors, while customized post-flight studies provided better constraints for data reduction. The rocket experiments have proven that the dual-order design is a viable option for future far-UV instruments, and the combined use of the two detectors is efficient in providing redundancy and a significantly increased dynamic range. The CCD detector has an excellent uniform response across the chip, very good image fidelity, and a very high saturation limit. As such, a CCD spectrograph would benefit from the increase in throughput in a two-bounce design, or from a grating blaze transferring power from the symmetric order. On the down side, its ability to detect faint sources is limited by read noise and dark current. The CCD performance in uniformly illuminated fields is degraded somewhat by a red leak. Also, in our launch environment, the necessity to operate at a low temperature puts constraints on launch preparations and observing time. The MCP is a robust detector, with a very low background level, but becomes quickly non-linear and manifests poor imaging performance compared to the CCD. The pixel mapping suffers from distortions in both spatial and spectral directions, and the measured line spread functions show a scatter in pixel distribution much higher than expected. However, the photon counting ability of the MCP with its instantaneous response provides an unrivaled advantage in detecting the faintest emissions and making real-time observing decisions. An ideal detector would combine the imaging performance and linearity of the CCD with the low background and photon counting ability of the MCP.  Overall, in spite of the identified shortcomings pertaining to cooling issues, focus consistency, and contamination control, the design of the mission and instrument have proven to be fault tolerant. An in-depth analysis of the instrument capabilities and associated error margins has led to the acquisition of valuable scientific information. Preliminary analysis of the flight data has been performed, but the final data products are not yet obtained. We will seek further clarification on the absolute flux scale of the stellar spectra, and reliable estimates for the background and slit orientation for the extraction of diffuse nebular emission. We believe that our observations will contribute to constructing better models for the interaction between UV radiation and dust particles in the interstellar medium."
    },
    "0806/0806.4346_arXiv.txt": {
        "abstract": "We revise a magnetic buoyancy model that has recently been proposed as a mechanism for extra mixing in the radiative zones of low-mass red giants. The most important revision is our accounting of the heat exchange between rising magnetic flux rings and their surrounding medium. This increases the buoyant rising time by five orders of magnitude, therefore the number of magnetic flux rings participating in the mixing has to be increased correspondingly. On the other hand, our revised model takes advantage of the fact that the mean molecular weight of the rings formed in the vicinity of the hydrogen burning shell has been reduced by $^3$He burning. This increases their thermohaline buoyancy (hence, decreases the total ring number) considerably, making it equivalent to the pure magnetic buoyancy produced by a frozen-in toroidal field with $B_\\varphi\\approx 10$\\,MG. We emphasize that some toroidal field is still needed for the rings to remain cohesive while rising. Besides, this field prevents the horizontal turbulent diffusion from eroding the $\\mu$ contrast between the rings and their surrounding medium. We propose that the necessary toroidal magnetic field is generated by differential rotation of the radiative zone, that stretches a pre-existing poloidal field around the rotation axis, and that magnetic flux rings are formed as a result of its buoyancy-related instability. ",
        "introduction": "During their first ascent on the red giant branch (RGB), a majority of low-mass stars (those with $M\\la 2\\,M_\\odot$) experience extra mixing in their radiative zones separating the H burning shell from the bottom of convective envelope (\\citealt{sm79,chdn98,dv03}). Despite 30 years of effort, however, the underlying physical mechanism is still not understood. Observationally, the RGB extra mixing manifests itself through changes of the surface abundances of Li, C, N, and of the isotopic ratio $^{12}$C/$^{13}$C correlating with an increasing luminosity (\\citealt{grea00,sm03}). These changes are produced by the joint operation of thermonuclear reactions that take place in the vicinity of the H shell and a nonconvective mixing process that transports reaction products through the radiative zone to the convective envelope. Observations support the idea that this mixing process starts (or, gets much more efficient) when an RGB star reaches a luminosity at which the differential luminosity function for a population of stars having the same age and chemical composition shows a prominent bump (a local pile-up of stars). The luminosity bump results from a temporary slowing down of the star's evolution caused by its structural readjustment. This happens when the H shell crosses and erases a discontinuity in the H-abundance profile left behind by the bottom of convective envelope at the end of the first dredge-up. During the first dredge-up, that occurs on the subgiant branch and on the lower RGB, the convective envelope grows in mass, which causes its bottom to penetrate the layers whose chemical composition had been altered yet on the main sequence (MS). This produces changes of the surface abundances of Li, C, N, and of the $^{12}$C/$^{13}$C ratio similar to but by far less substantial than those incurred from the subsequent operation of the RGB extra mixing.  Until recently, it has been thought that the only reason why the RGB extra mixing does not manifest itself below the bump luminosity is a strong gradient of the mean molecular weight $\\mu$ caused by the onset of a deep convective envelope (e.g., \\citealt{chea98}). Any mixing mechanism has to overcome the stable thermal stratification of the radiative zone, and in the presence of a positive $\\nabla_\\mu$ such mixing is correspondingly more difficult. In an RGB star above the bump luminosity, the H shell has already crossed the H-profile discontinuity, therefore the radiative zone is now chemically uniform everywhere except in a very close neighborhood of the H shell. This circumstance was repeatedly emphasized in the past. In particular, it has been used to model the RGB extra mixing with rotation-driven meridional circulation and turbulent diffusion. It was not until recently that it has become clear that rotational mixing fails to explain the chemical element transport in the radiative zones of upper RGB stars (\\citealt{chea05,pea06}). In short, this failure is due to the following main causes: firstly, rotation period measurements for young cluster stars and helioseismic data indicate that MS stars with $M\\la 1\\,M_\\odot$ loose a great amount of their initial angular momentum via magnetized stellar winds and that they most likely become slow and nearly solid-body rotators before leaving the MS; secondly, the chemical element transport by meridional circulation is strongly hindered by rotation-induced horizontal turbulence in stellar radiative zones (\\citealt{chz92}); thirdly, the vertical turbulent diffusion powered by differential rotation in the radiative zones of RGB stars operates at a low level too because it also redistributes the angular momentum, thus reducing the degree of differential rotation in a self-regulating way.  A new class of RGB extra mixing models has emerged since \\cite{eea06} noticed that a tiny $\\mu$-gradient inversion ($\\nabla_\\mu \\approx -10^{-4}$) becomes visible at the outer tail of the H burning shell precisely at the moment when the H shell erases the H-profile discontinuity. This inversion is produced by the reaction $^3$He($^3$He,\\,2p)$^4$He that locally reduces the mean molecular weight by $\\Delta\\mu\\approx\\mu^2\\Delta X_3/6$, where $X_3$ is the $^3$He mass fraction. The mechanism can be effective in the low-mass RGB stars because their MS progenitors synthesize large amounts of $^3$He in their outer radiative cores through non-equilibrium pp burning. Even though this $^3$He-rich material gets diluted in the convective envelope during the first dredge-up, the radiative zone of a low-mass RGB star above the bump luminosity can still have $X_3$ increased up to a value of $2\\times 10^{-3}$ (the solar initial $^3$He abundance is $3\\times 10^{-5}$). For this mass fraction, the $^3$He burning leads to $\\Delta\\mu\\approx -10^{-4}$, assuming that $\\mu\\approx 0.6$ and $\\Delta X_3\\approx -X_3$. Below the bump luminosity, the $\\mu$-gradient inversion is overridden by the strong positive $\\mu$-gradient built up on the MS. It shows up and may come into play only when the $^3$He burning shell, advancing in mass in front of the major H shell, finds itself in the chemically homogeneous part of the radiative zone. This happens at the bump luminosity. \\cite{eea06} found that a rapid mixing process occurred in their 3D simulations above this point, although the underlying cause was not identified (see \\citealt{dp08b}).  Inspired by this work, \\cite{chz07a} have proposed that the $\\mu$-gradient inversion maintained by the $^3$He burning drives thermohaline convection in the radiative zones of low-mass RGB stars above the bump luminosity and that this is the long-sought physical mechanism for the RGB extra mixing. Thermohaline convection is a mixing process triggered by a double diffusive instability (e.g., \\citealt{v04}). Consider a stratified ideal gas with a stable temperature gradient ($\\nabla\\equiv d\\ln T/d\\ln P < \\nabla_{\\rm ad}$, where ``ad'' stands for adiabatic changes) but with an unstable composition gradient ($\\nabla_\\mu < 0$). If we isolate a gas blob and shift it up in the vertical direction then its further motion will depend on how fast the blob exchanges heat and composition with its surrounding medium horizontally. Indeed, the relative difference in density between the surrounding medium and the blob is $\\Delta\\rho/\\rho \\approx \\Delta\\mu/\\mu - \\Delta T/T$, assuming that $\\Delta P = 0$. For the blob to continue rising, we need $\\Delta\\rho > 0$. Our assumptions about the gradients mean that $\\Delta\\mu > 0$ and $\\Delta T > 0$ in the absence of both heat and molecular diffusion. Because these differences grow when the blob rises, $\\Delta\\rho$ may stay positive or it may ultimately become negative depending on the ratio $r_\\mu = |\\nabla_\\mu|/(\\nabla_{\\rm ad}-\\nabla)$. In our particular case, $r_\\mu\\ll 1$. Therefore, our idealized impermeable and adiabatic blob will rise a short distance and then stop, when the accumulated difference in $T$ compensates that in $\\mu$. In reality, the heat exchange, whose rate is specified by the radiative diffusivity $K$, constantly works to reduce the difference in $T$. On the other hand, molecular diffusion $\\nu_{\\rm mol}$ tries to smooth out the difference in $\\mu$. The double diffusive instability may therefore develop only if $K\\gg \\nu_{\\rm mol}$. In this case, the blob's rising speed can be estimated as $v\\sim l/\\tau_{\\rm th}$, where $l$ is the mean path that the blob travels before it gets dissolved, while $\\tau_{\\rm th}\\sim d^2/K$ is the characteristic thermal time scale for a spherical blob of the diameter $d$. An approximate expression for the thermohaline diffusion coefficient can be obtained as $D_{\\rm thc}\\sim lvr_\\mu \\sim Kr_\\mu(l/d)^2$. \\cite{chz07a} and \\cite{dp08b} have demonstrated that the observed RGB mixing patterns can be explained by stellar evolutionary models with the $^3$He-driven thermohaline convection only if $l/d\\ga 10$\\,--\\,$30$. A similarly large parameter ratio for thermohaline convection in stellar radiative zones was postulated by \\cite{u72}, as opposed to a ratio $l/d\\sim 1$ advocated by \\cite{kea80}. Besides, the double-diffusive instability has been shown to result in formation of elongated (large $l$ to $d$ ratios) structures known as ``salt fingers'' in laboratory experiments with the saltier and warmer water overlying the fresher and colder water (\\citealt{s60}).  However, there appears to exist observational and theoretical arguments challenging this model. First of all, a large number of old metal-poor MS stars with $M\\la 0.9\\,M_\\odot$, both in globular clusters and in the halo field, that had accreted He- and C-rich high-$\\mu$ material from their evolved cluster or binary companions do not seem to have been thoroughly mixed by thermohaline convection (\\citealt{nt07,dp08a,aea08}), as it would be expected even if the less efficient prescription by \\cite{kea80} were used for $D_{\\rm thc}$. Second, thermohaline convection is expected to be suppressed by the rotation-induced horizontal turbulence that works together with the molecular diffusion to reduce the $\\mu$ contrast between the rising gas blob and its surroundings (\\citealt{dp08b}). Third, strong differential rotation is predicted to hinder thermohaline convection as well, because the ``salt fingers'' may be tilted by the rotational shear so rapidly that they will get damped before they produce significant mixing (\\citealt{c99}, and references therein). We anticipate that a similar effect is also produced by the Coriolis force in a uniformly rotating radiative zone. Contrary to these expectations, a much larger fraction of Li-rich objects has been found among rapidly rotating ($v\\sin\\,i\\geq 8$\\,km\\,s$^{-1}$) K giants than among their more common slowly rotating ($v\\sin\\,i\\la 1$\\,km\\,s$^{-1}$) counterparts (\\citealt{dea02}). The Li-rich K giants are low-mass stars located above the bump luminosity (\\citealt{chb00}) in which large amounts of Li are thought to be synthesized via the $^7$Be-transport mechanism (\\citealt{cf71}). To be efficient, this mechanism needs a 10 to 100 times faster mixing than that required to reproduce the abundance patterns in the majority of upper RGB stars (\\citealt{dh04}). It is not clear how thermohaline convection can explain the phenomenon of Li-rich K giants given that its efficiency should be lower in the more rapidly rotating stars. These arguments have motivated our search for an alternative RGB mixing mechanism.  In this paper, we use a simple model of toroidal magnetic field generation in a differentially rotating radiative zone of a bump-luminosity RGB star to obtain order-of-magnitude estimates demonstrating that the buoyant rise of magnetic flux rings being formed close to the local minimum in $\\mu$ may be a good alternative to the $^3$He-driven thermohaline convection. A similar model has recently been proposed by \\cite{bea07} (hereafter, referred to as BWNC). However, they assumed that a rising ring always stays in thermal equilibrium with its surrounding medium, and we argue that this leads to a substantial overestimate of the ring's radial velocity. We account for the impact of $\\mu$ gradients and discuss the origin of the magnetic rings.  We show that, as a mechanism for the RGB extra mixing, magnetic buoyancy has some advantages over thermohaline convection and, therefore, it is worth further investigating by means of multidimensional MHD simulations. ",
        "conclusions": "In this work, we have presented a simple model of the formation and buoyant rise of magnetic flux rings in the radiative zone of the bump luminosity RGB star. Our model is based on ideas and equations published by \\cite{svb82}, \\cite{sr83}, \\cite{mw87}, \\cite{chg87}, \\cite{chmg93}, \\cite{s99}, \\cite{eea06}, BWNC, and \\cite{dp08b}. It qualitatively describes a possible mechanism for the RGB extra mixing, which we call the magneto-thermohaline mixing, as an alternative to the pure $^3$He-driven thermohaline convection that has recently been proposed by \\cite{chz07a}. For our mechanism to work, the radiative zone has to possess a strong differential rotation and a poloidal magnetic field $B_{\\rm p}\\ga 1$\\,--\\,10 G. We assume that the differential rotation stretches the poloidal field around the rotation axis, thus creating a strong toroidal magnetic field $B_\\varphi\\approx 0.1$\\,--\\,1 MG. When the latter exceeds a critical value, the buoyancy-related undular instability comes into play to form magnetic flux rings. These rings turn out to be buoyant, therefore they rise toward the bottom of convective envelope.  We have shown that, when the radiative heat exchange between the ring and its surrounding medium is taken into account, the ring's buoyant rising time increases by about five orders of magnitude compared to the case considered by BWNC, when the ring and its surrounding medium are assumed to be in thermal equilibrium all the time. However, given that our model still neglects possible internal heating of the ring's material by residual nuclear reactions and anisotropic thermal exchanges in the presence of strong oriented magnetic field, while it uses the aerodynamic drag coefficient that is 20 times as large as the one employed by BWNC, it is fair to say that BWNC might have fixed a safe upper limit while our paper fixes a conservative lower limit for the magnetic ring's rising velocity.  We have found that the number of rings needed to be present in the radiative zone at the same time to produce the observationally constrained rate of the RGB extra mixing is unrealistically large unless these rings originate from the region of the $\\mu$ inversion maintained by the $^3$He burning. Such rings have a deficit of the mean molecular weight compared to the bulk of the radiative zone through which they move. Their buoyancy is mainly caused by the difference in $\\mu$ rather than by a deficit in density due to the excess magnetic pressure. The frozen-in toroidal magnetic field is still needed for the rings to remain cohesive while rising. That is why we have coined the term ``magneto-thermohaline'' mixing. Our model has some advantages over the pure thermohaline mixing model, the most important of which being the robustness of the magnetic rings against the eroding effect produced by the horizontal turbulent diffusion. Leaving aside the problem of the \"parasitic\" rings that are formed at $r\\ga r_{\\rm mix} + 0.05\\,R_\\odot$, our model looks promising. However, because it is based on a number of assumptions whose legality is impossible to confirm in the framework of our 1D computations we call for its future verification by 3D MHD simulations."
    },
    "0806/0806.4493_arXiv.txt": {
        "abstract": "Sky temperature map of the cosmic microwave background (CMB) is one of the premier probes of cosmology. To minimize instrumentally induced systematic errors, CMB anisotropy experiments measure temperature differences across the sky using paires of horn antennas with a fixed separation angle, temperature maps are recovered from temperature differences obtained in sky survey through a map-making procedure. The instrument noise, inhomogeneities of the sky coverage and sky temperature inevitably produce statistical and systematical errors in recovered temperature maps. We show in this paper that observation-dependent noise and systematic temperature distortion contained in released Wilkinson Microwave Anisotropy Probe (WMAP) CMB maps are remarkable. These errors can contribute to large-scale anomalies detected in WMAP maps and distort the angular power spectrum as well. It is needed to remake temperature maps from original WMAP differential data with modified map-making procedure to avoid observation-dependent noise and systematic distortion in recovered maps. ",
        "introduction": "The COBE and WMAP missions measure temperature differences between sky points using differential radiometers consisting of plus-horn and minus-horn~\\cite{smo90,ben03a}. Let denote $t_i$ the temperature anisotropy at a sky pixel $i$. The raw data in a certain band is a set of temperature differences {\\bf d} between pixels in the sky. From $N$ observations we have the following observation equations \\begin{equation} \\label{dt} \\begin{array}{c@{\\:-\\:}c@{\\;=\\;}c} t_{1^+} & t_{1^-} & d_1 \\\\ t_{2^+} & t_{2^-} & d_2 \\\\ \\multicolumn{3}{c}{\\dotfill}\\\\ t_{N^+} & t_{N^-} &~~d_N~. \\end{array} \\end{equation} The above equation system can be expressed by matrix notation \\begin{equation} \\label{dt1} \\mathbf{At=d}~. \\end{equation} Where the scan matrix of the experiment {\\bf A}$ =(a(k,i)),~k=1,\\cdots,N$ and $i=1,\\cdots,L$ with $L$ being the total number of sky map pixels. The most of elements $a(k,i)=0$ except for $a(k,i=k^+)=1$ and $a(k,i=k^-)=-1$, where $k^+$ denotes the pixel observed by the plus-horn and $k^-$ the pixel observed by the minus-horn at an observation $k$.  The normal equation of Eq.~\\ref{dt} or Eq.~\\ref{dt1} is \\begin{equation} \\label{ne} \\mathbf{Mt=A^Td} \\end{equation} with $\\mathbf{M=A^TA}$. The least-squares estimate of the sky map results from solving Eq.~\\ref{ne} \\[ \\mathbf{\\hat{t}=M^{-1}A^Td}~. \\] The WMAP team \\cite{hin03} uses the following approximate formula to compute the iterative solution \\begin{equation} \\label{mm-w} \\mathbf{t^{(n+1)}=\\tilde{M}^{-1}(A^Td-A^TAt^{(n)})+t^{(n)}}~, \\end{equation} where $\\mathbf{\\tilde{M}^{-1}}=$ diag$(\\frac{1}{N_1},\\frac{1}{N_2},\\cdots)$ is an approximate inverse of {\\bf M} with $N_i$ being the total number of observations for pixel $i$.  The use of approximate inverse matrix  $\\mathbf{M^{-1}}$ is not necessary. Here we derive an iterative formula directly from the normal equation. The Eq.~\\ref{ne} can be expressed as \\begin{eqnarray*} N_i^+t_i-\\sum_{k^+=i}t_{k^-}-\\sum_{k^-=i}t_{k^+}+N_i^-t_i =\\sum_{k^+=i}d_k-\\sum_{k^-=i}d_k~,  \\\\ (i=1,2,\\cdots,L)~. \\end{eqnarray*} Where $\\sum_{k^+=i}$ means summing over $N_i^+$ observations while the pixel $i$ is observed by the plus-horn and $\\sum_{k^-=i}$ means summing over $N_i^-$ observations while the pixel $i$ is observed by the minus-horn, and the total number of observations for the pixel $i$ is $N_i=N_i^++N_i^-$. From the above equations we can derive the following iterative formula \\begin{eqnarray} \\label{mm-l} t_i^{(n+1)}=\\frac{1}{N_i}(\\sum_{k^+=i}(d_k+t_{k^-}^{(n)})-\\sum_{k^-=i} (d_k-t_{k^+}^{(n)}))~, \\nonumber\\\\ (i=1,2,\\cdots,L)~. \\end{eqnarray}  With Eq.~\\ref{mm-w} or Eq.~\\ref{mm-l} when the number $n$ of iteration is large enough, we get the final solution $\\hat{t}_i=t_i^{(n)}$ for each pixel $i$. The Eg.~\\ref{mm-w} used by the WMAP team is an approximate formula and Eq.~\\ref{mm-l} is an exact one, but both has good performance for the differential data of a noiseless instrument. With Eq.~\\ref{mm-l} we can easily study the statistical and systematical errors induced by instrument noise, inhomogeneity of sky coverage, inhomogeneity of sky temperature, and unbalance between two sky side measurements. ",
        "conclusions": "\\subsection{Can the exposure induced anisotropy be corrected ?} It has to be pointed out that Eq.~\\ref{rms} can be used to modify the effect of exposure dependent noise, like what we do for the latitude distribution of rms variation shown in Fig.~2, only for the case that the map rms fluctuation $\\<{\\hat{t}}^2\\>$ itself is the directly analyzed quantity. However, almost all analysis works are based on CMB maps of temperature. We know from Figs.~2-5 that the released WMAP temperature maps contain considerable exposure dependent noise. It is no way  from a recovered temperature map to produce a corrected map in which the instrument-induced and exposure-dependent noise can be eliminated. Therefore, the anisotropy noise should contribute to large scale anomalies detected in existent CMB maps.  The exposure inhomogeneity of WMAP comes from its scan strategy, which can't be suppressed through accumulating observation time. To avoid the observation effect, it is needed to remake temperature maps from a uniform differential data set obtained by giving up partial observation data for pixels of high exposure. Comparing released WMAP maps and new maps from uniform data will help us to judge their origin of detected large scale anomalies, e.g. the low $l$ power issues detected in WMAP data, the unexplained orientation of large-scale patterns of CMB maps in respect to the ecliptic frame, the north-south asymmetry of temperature fluctuation power etc., and to see if the observational effect can also influence the angular power spectrum as well.  \\subsection{Avoiding the foreground-induced distortion} \\begin{figure}[p] \\label{f15} \\vspace{-18mm} \\begin{center} \\psfig{figure=f15a.ps,width=45mm,height=70mm,angle=270}\\\\ \\vspace{2mm}\\psfig{figure=f15b.ps,width=45mm,height=70mm,angle=270}\\\\ \\vspace{2mm}\\psfig{figure=f15c.ps,width=45mm,height=70mm,angle=270} \\caption{Differences between recovered $t$ and true temperature $t_0$. The one-dimensional temperature distribution $t_0$ consists of white noise and hot source between 240 - 260 pixel (mask region), as shown in the upper panel of Fig.~10. Differential data is obtained by simulation of one-dimensional scan for $t_0$ by a differential radiometer with beam-separation of 100 pixel. Recovered $t$ are calculated with map-making Eq.~\\ref{mm-w} after 50 iterations. {\\sl Upper panel}: $t$ recovered from all differential data (the lower panel of Fig.~10). {\\sl Middle panel}: $t$ recovered by iterations from all data for mask region and with excluding differences that contain hot source temperature for pixels out of mask. {\\sl Bottom panel}: $t$ recovered by iterations with initials estimated from the differential data by Eq.~\\ref{tj}. } \\end{center} \\end{figure}  The distortion by hot foreground sources on their $141\\degree$ rings in a WMAP map can not be removed with a foreground mask on the recovered temperature map. What's needed is to use the mask on the original differential data before map-making to avoid the foreground-induced error in the recovered map. The top panel of Fig.~15 shows the difference between the recovered and true temperature distributions (shown in the lower and upper panel of Fig.~10 respectively), where the distortion structure caused by the hot source on pixel 240 - 260 and beam separation of 100 pixel is clearly shown. We redo the temperature reconstruction with excluding the temperature differences that contain a beam side pointing to a pixel between 240 - 260 (``mask region'') during iterations for the pixels out of mask, the result is shown in the middle panel of Fig.~14, where the distortion structures are really suppressed.  A weakness of using mask in map-making process is decreasing the number of useful differential data. Another approach to avoid the distortion in recovered map by foreground emission is to properly set the initials of iteration for the foreground region. From Eq.~\\ref{t1} we see that the temperature deviation of the first iterative solution, $t_i^{(1)}-t_i$, will be suppressed if the temperature initials at pixels of hot source are set to be close to their true values to let $\\<t-t^{(0)}\\>_{ring}\\approx~0$. The initial $t_i^{(0)}$ of pixel $i$ can be taken as \\begin{equation} \\label{tj} t_i^{(0)}=\\frac{1}{N_i^\\prime}(\\sum_{k^+=i}d_k-\\sum_{k^-=i}d_k)~. \\end{equation} Where $\\sum_{k^+=i}$ means summing over the observations while the pixel $i$ is observed by the plus-horn and the pixel pointed by the minus-horn is out of mask, $\\sum_{k^-=i}$ means summing over the observations while the pixel $i$ is observed by the minus-horn and the pixel pointed by the plus-horn is out of mask, $N_i^\\prime$ is the total number of used observations. For the simulated differential data from the true temperatures shown in the upper panel of Fig.~10,  we make 50 iterations with Eq.~\\ref{mm-w} starting from initials calculated by Eq.~\\ref{tj}, the distortion structures are satisfactory suppressed in the resultant solution, as shown in the bottom panel of Fig.~15.  \\subsection{Remaking WMAP maps} We demonstrate in this paper that for existent CMB maps both the observation dependent noise and systematic error induced by foreground emission can not be neglected and both can produce large-scale anomalies and distort the angular power spectrum. These errors can not be completely excluded by performing noise suppressing or using foreground mask on temperature maps. We suggest to remake temperature maps from the original WMAP time-order-data by a modified algorithm with applying foreground mask in map-making to exclude mask pixels from use in iterations for CMB dominated region (or properly set temperature initials before iteration), and/or keeping used differential data uniform by giving up partial observation data for pixels of high exposure. New maps from modified map-making algorithm will help us to judge the origin of large scale anomalies detected in released WMAP maps, e.g. the low $l$ power issues, the unexplained orientation of large-scale patterns in respect to the ecliptic frame, the north-south asymmetry of temperature fluctuation power and the large non-Gaussian spots, and to see to what extent the statistical and systematical errors influence the angular power spectrum and the derived cosmological parameters. Believable conclusions on CMB anisotropy anomalies and precise temperature angular power spectrum from differential measurement should be based on temperature maps with homogeneous sky exposure and should avoid foreground-induced distortion during map-making.  This study is supported by the National Natural Science Foundation of China and the CAS project KJCX2-YW-T03. The data analysis in this work made use of the WMAP data archive and the HEALPIX software package."
    },
    "0806/0806.3616_arXiv.txt": {
        "abstract": "During last years a few massive binary systems have been detected in the TeV $\\gamma$-rays. This $\\gamma$-ray emission is clearly modulated with the orbital periods of these binaries suggesting its origin inside the binary system. In this paper we summarize the anisotropic IC $e^\\pm$ pair cascade model as likely explanation of these observations. We consider scenarios in which particles are accelerated to relativistic energies, either due to the presence of an energetic pulsar inside the binary, or as a result of accretion process onto the compact object during which the jet is launched from the inner part of the accretion disk, or in collisions of stellar winds from the massive companions. ",
        "introduction": "All detected binary systems at TeV $\\gamma$-rays belong to the class of high mass binaries in which compact object (neutron star or black hole) appears very close to the stellar surface at least during a part of its orbit. In fact, TeV $\\gamma$-rays have been expected on theoretical grounds from the well known Be binary system, PSR B1259-63/SS2883, containing 47.7 ms pulsar since it was well known that even isolated pulsars with similar parameters should emit TeV $\\gamma$-rays (e.g. the Crab pulsar and its nebula). Such emission has been discovered from this object during the periastron passage of the pulsar in 2004 (Aharonian et al.~2005a). A year later, emission from two compact binary systems, LS 5039 and LSI +61 303, classified as microquasars,  have been observed by HESS and MAGIC telescopes (Aharonian et al.~2005b, Albert et al.~2006). This emission is clearly modulated with the orbital periods of the binary systems (Aharonian et al.~2006, Albert et al.~2008). Moreover, there are evidences of detection of the TeV $\\gamma$-ray flare from another microquasar Cyg X-1 (Albert et al.~2007). Finally, the HESS discovered the TeV emission from the open cluster Wester\\-lund 2 which contains the most massive binary system of two WR type stars, WR 20a (Aharonian et al.~2007). This emission might also be related to this very unusual binary system.  TeV $\\gamma$-ray production in binary system containing energetic pulsar is usually interpreted in terms of the leptonic IC scattering model in which relativistic electrons scatter soft radiation coming from the massive star (e.g. see scenarios discussed by Maraschi \\& Treves~1981, Tavani \\& Arons~1997, Kirk et al.~1999).  Similar radiation processes have been also considered in more detail in the case of electrons accelerated in the jets of  microquasars (e.g. Bosch Ramon et al.  2006, Dermer \\& B\\\"ottcher~2006). In these calculations the absorption of $\\gamma$-ray photons on the stellar radiation is usually not taken into account or considered only in approximate way. More recently $\\gamma$-ray emission from one-directional IC $e^\\pm$ pair cascades (secondary leptons move in the direction of primary $\\gamma$-ray photons)  have been considered in a few papers (Aharonian et al. ~2006, Orellana et al.~2007, Khangulyan et al.~2008).  Here, we discuss the three dimensional IC $e^\\pm$ pair cascade scenario for the TeV $\\gamma$-ray origin in the massive binary systems which have been developed in the late 90-ties (e.g. Bednarek~2000) and later applied to these type of sources in a sequence of recent papers (Bednarek~2005, 2006a,b, 2007, Bednarek \\& Giovannelli~2007, Sierpowska-Bartosik~2007, Sierpowska \\& Bednarek~2005,  Sierpowska-Bartosik \\& Bednarek~2008, Sierpowska-Bartosik \\& Torres~2007, 2008). ",
        "conclusions": ""
    },
    "0806/0806.2801_arXiv.txt": {
        "abstract": "We report the results of an intense, spectroscopic survey of all 41 late-type, nitrogen-rich Wolf-Rayet (WR) stars in the Large Magellanic Cloud (LMC) observable with ground-based telescopes. This survey concludes the decade-long effort of the Montr\\'eal Massive Star Group to monitor every known WR star in the Magellanic Clouds except for the 6 crowded WNL stars in R136, which will be discussed elsewhere. The focus of our survey was to monitor the so-called WNL stars for radial-velocity (RV) variability in order to identify the short- to intermediate-period ($P \\la 200$ days) binaries among them. Our results are in line with results of previous studies of other WR subtypes, and show that the binary frequency among LMC WNL stars is statistically consistent with that of WNL stars in the Milky Way. We have identified four previously unknown binaries, bringing the total number of known WNL binaries in the LMC to nine. Since it is very likely that none but one of the binaries are classical, helium-burning WNL stars, but rather superluminous, hence extremely massive, hydrogen-burning objects, our study has dramatically increased the number of known binaries harbouring such objects, and thus paved the way to determine their masses through model-independent, Keplerian orbits. It is expected that some of the stars in our binaries will be among the most massive known. With the binary status of each WR star now known, we also studied the photometric and X-ray properties of our program stars using archival MACHO photometry as well as \\emph{Chandra} and \\emph{ROSAT} data. We find that one of our presumably single WNL stars is among the X-ray brightest WR sources known. We also identify a binary candidate from its RV variability and X-ray luminosity which harbours the most luminous WR star known in the Local Group. ",
        "introduction": "The optical spectra of Wolf-Rayet (WR) stars feature broad emission lines from highly-ionized elements. These emission lines arise in a fast, hot, and dense stellar wind which is generally optically thick in the inner part, thereby completely veiling the hydrostatic photosphere of the WR star. Depending on the elements they show in their optical spectra, WR stars are classified into three different subtypes. If a WR star displays predominantly helium (He) and nitrogen (N), which are formed during hydrogen burning via the CNO cycle, the star is classified as WN; if the WR star displays, in addition to He, predominantly carbon (C) or oxygen (O), which are the products of $3\\alpha$ He burning, it is classified as WC or WO.  Due to their chemical properties, it is now generally accepted that classical WR stars are evolved objects, namely the almost bare, hydrogen-depleted, helium-burning cores of stars whose initial mass on the main sequence (MS) was, at solar metallicity, above $\\sim 25 M_{\\sun}$, i.e. that started their lives as O stars (\\citealt{Lamers91}; \\citealt{MaedCont94}). Thus, the key question of WR-star formation is how a WR-star progenitor loses its outer, H-rich envelope to expose the CNO-enriched, deeper layers, and how WR-star formation depends on ambient metallicity. For single stars, three scenarios have been put forward, depending on the initial mass of the star. Stars with initial masses $25 M_{\\sun} \\la M_{\\rm i} \\la 40 M_{\\sun}$ are expected to become red supergiants (RSGs), as which they experience continuous mass loss through winds. Stars with initial masses $40 M_{\\sun} \\la M_{\\rm i} \\la 85 M_{\\sun}$ are believed to turn into luminous blue variables (LBVs), as which they experience outbursts of violent, eruptive shell-ejections (\\citealt{HD79}, \\citealt{HD94}). Stars with even higher initial masses ($M_{\\rm i} \\ga 85 M_{\\sun}$ at solar metallicity) are believed to reach the WN phase while they are still core-hydrogen burning, and to \\emph{directly} evolve into classical, He-burning WR stars (\\citealt{Conti76}) without going through the LBV phase.  In close binaries, Roche-lobe overflow (RLOF) is suspected to enhance WR-star formation by removing the H-rich envelope of the (more evolved) primary (\\citealt{KippWeig67}). \\citet{Pacz67} was the first to note that a thus stripped He-burning core would very likely resemble a WR star. Since RSGs descend from the initially least massive O stars (see above), it follows that RLOF might thus considerably contribute to the total WR numbers, in particular in low-metallicity environments, where radiatively-driven mass-loss rates are expected to be very low (see \\citealt{Vink00}; \\citealt{Vink01}). Indeed, older, non-rotating stellar-evolution models were unable to reproduce the observed WR populations at different $Z$ without either enhancing by a factor of two the then-higher mass-loss rates through stellar winds or increasing the fraction of interacting binaries (\\citealt{MaedMey94}). While updated models now include stellar rotation, and much better reproduce the observations without the need of increased binary interaction (\\citealt{MaedMey00a}), the influence of binarity on WR-star formation remains unclear.  From model calculations, \\citet{Vanbev98} found that close O+O binaries with initial periods $P \\la 1000$ d cannot escape RLOF if the primary star reaches the RSG stage, resulting in WR+O binaries with post-RLOF periods $\\la 200$d (\\citealt{Wellstein99}). Hence, in environments where RLOF is expected to become the increasingly important WR-star formation channel (i.e. at lower ambient $Z$), the fraction of WR binaries with present-day periods of up to 200 days should be higher. This prediction is accessible to observational tests. For such a study, the Magellanic Clouds are the ideal laboratory: $i$) The distances to both the SMC and the LMC are well established and $\\sim$ constant for all stars (e.g. \\citealt{KellerWood06}); $ii$) reddening towards the Clouds is low and fairly constant, contrary to the Galaxy (e.g. \\citealt{Niko04}), $iii$) the WR populations in both Clouds are nearly complete (e.g. \\citealt{MassDuff01}, but see \\citealt{Massey03}), $iv$) the WR population is large enough to allow for reasonable statistics. In total, the LMC harbours 132 WR stars (\\citealt{BAT99}, hereafter BAT99\\footnote{BAT99 lists 134 WR stars in the LMC, but see \\citet{Moff91} and \\citet{Niemela01} for the revised non-WR status of BAT99-4 and BAT99-6, respectively.}), while 12 WR stars are known in the SMC (\\citealt{MassDuff01}; \\citealt{Massey03}). To search for binaries, the total sample of 144 stars has been split into three distinct studies: \\citet{Bartz01} reported the results on the 25 Magellanic Cloud WC/WO stars, while Foellmi et al. (2003a,b) studied the 71 then-known, early-type WNE (=WN2-WN5) stars, with the exception of the H-rich WNE stars in and around the R136 cluster in the 30 Dor region. The observations of a 72nd WNE star, which was newly discovered by \\citet{Massey03} in the SMC, were reported by \\citet{Foell04}.  In the present paper, we describe our intense, spectroscopic survey of the remaining 41, late-type WNL (=WN5-11) stars in the LMC. The two WN6 stars in the SMC were already studied by \\citet{Foell03a}. Our survey includes those WNL stars in the periphery of R136 which could be observed with ground-based telescopes without adaptive optics (AO). For the 6 luminous, WN5-7ha stars in the very core of R136, AO-assisted, near-infrared spectroscopy using VLT/SINFONI was used; those results will be reported elsewhere (Schnurr et al., in preparation). The aim of our study is manifold. First and foremost, we will assess the binary status ($P \\la 200$ days) of each of our 41 program stars. This will conclude the decade-long effort of the Montr\\'eal Massive Star Group to study spectroscopically the entire WR population in the Magellanic Clouds, and pave the way to obtain a much clearer view of the role binarity plays in the evolution of massive stars at different metallicities. Secondly, binaries identified in this study can be used, in the future, to determine their respective masses by using model-independent, Keplerian orbits. Masses of WR stars are of greatest importance in the context of the calibration of both atmospheric and evolutionary models, in particular for the most massive stars. We have ample reason to believe that at least some H-rich WNL stars in our sample belong to the subgroup of very massive, possibly even the most massive stars known in the Local Group (\\citealt{Rauw96b}; \\citealt{Schweick99}; \\citealt{Rauw04}; \\citealt{Bonanos04}). Thirdly and as a side effect, we will be able to put publicly available, archival X-ray data from \\emph{Chandra} and \\emph{ROSAT} into context with the binary status of our stars, since massive binaries with colliding stellar winds are known to be strong X-ray emitters.  The paper is organized as follows: In Section \\ref{wnlsection2}, we will describe the observations of our program stars. In Section \\ref{wnlsection3}, we will briefly describe the data reduction. In Section \\ref{wnlsection4}, we will describe in detail how we analysed our spectroscopic and X-ray data and the results we obtained. These results will be discussed in Section \\ref{wnlsection5}. Section \\ref{wnlsection6} summarizes and concludes this paper. ",
        "conclusions": "\\label{wnlsection6}  We have carried out spectroscopic monitoring of all 41 WNL stars in the LMC that could be observed by conventional, ground-based observations. Measured RV curves were used to identify binaries with orbital periods from 1 to $\\sim200$ days, because these systems were expected to be post-RLOF candidates (cf \\citealt{Vanbev98}). Additionally, publicly available archive data from X-ray satellite missions were searched for our program stars to obtain X-ray luminosities. The results of our study can be summarized as follows (see also Table \\ref{binaries}):  \\begin{itemize}  \\item{We have identified four previously unknown binary systems: BAT99-12, 95, 99, and 113.}  \\item{We confirmed the previously known binaries BAT99-32, 77, 92. However, while we could reproduce the 2.76-day period that \\citet{M89} had reported for star BAT99-102, we did so for star BAT99-103. It presently remains unknown whether \\citet{M89} or we wrongly identified the binary.}  \\item{We also confirmed that BAT99-119 is a binary; however, we had to combine our RV data with those of \\citet{M89} to do so, and it required further combination with previously unpublished polarimetry to identify the 159-day period of the system. For reasons of completeness, we list the results here, but the complete study will be published elsewhere (Schnurr et al., in prep).}  \\item{One star, BAT99-107, had been suspected binary by \\citet{M89}, but we were unable to reproduced that results. Therefore, we consider 105 to be single (i.e., not a binary with an orbital period in the quoted range).}  \\item{Two binary candidates were identified from their RV variability and their X-ray luminosities, BAT99-116 and 118. Both systems merit a closer look, because 116 is one of the brightest X-ray sources among all WR stars, while 118 is the most luminous WR star and thus, the most luminous unevolved star known in the Local Group (cf. \\citealt{CroDess98}).}  \\item{One of our program stars, BAT99-92, a binary, was recognized to be a WNE and not a WNL star.}  \\item{Thus, our study brings the total number of known WNL binaries to 8, and the binary frequency among WNL stars in the LMC to 20\\%, which is fully consistent with the results for WC/WO stars (\\citealt{Bartz01}) and WNE stars (\\citealt{Foell03b}); thus, there is no statisticially significant differences between different WR populations in the LMC. However, the overall binary frequency is only half of what was predicted from model results by \\citet{MaedMey94}. The implications of this low binary frequency for massive-star evolution will be discussed in a forthcoming paper.}  \\end{itemize}   Remarkably, none of the WNL binaries contains a classical, hydrogen-deficient, helium-burning WR star; instead, the WN components are young, unevolved, objects (hot O3If/WN6 stars or more extreme WN5-7ha stars), which most likely are very luminous and hence very massive Of stars, and possibly even the most massive stars known. These binaries offer the tremendous opportunity to directly weigh these extreme stars using model-independent, Keplerian orbits. Follow-up observations have partly been obtained and are currently reduced, or are under way. The results of these observations will be published elsewhere."
    },
    "0806/0806.0804_arXiv.txt": {
        "abstract": "{Observations at 0\\farcs1 have revealed the existence of dark cores in the bright filaments of sunspot penumbrae. Expectations are high that such dark-cored filaments are the basic building blocks of the penumbra, but their nature remains unknown.} {We investigate the origin of dark cores in penumbral filaments and the surplus brightness of the penumbra. To that end we use an uncombed penumbral model.} {The 2D stationary heat transfer equation is solved in a stratified atmosphere consisting of nearly horizontal magnetic flux tubes embedded in a stronger and more vertical field. The tubes carry an Evershed flow of hot plasma.} {This model produces bright filaments with dark cores as a consequence of the higher density of the plasma inside the tubes, which shifts the surface of optical depth unity toward higher (cooler) layers. Our calculations suggest that the surplus brightness of the penumbra is a natural consequence of the Evershed flow, and that magnetic flux tubes about 250 km in diameter can explain the morphology of sunspot penumbrae.} {} ",
        "introduction": "At high angular resolution, penumbral filaments are observed to consist of a central dark lane and two lateral brightenings (Scharmer et al.\\ 2002; S\\\"utterlin et al.\\ 2004; Rouppe van der Voort et al.\\ 2004; Bellot Rubio et al.\\ 2005; Langhans et al.\\ 2007). The common occurrence of dark-cored filaments and the fact that their various parts show a coherent behavior have raised expectations that they could be the fundamental constituents of the penumbra. Their nature, however, remains enigmatic.  One possibility is that dark-cored filaments represent magnetic flux tubes carrying a hot flow. This would support the uncombed model proposed by Solanki \\& Montavon (1993), which describes the penumbra as a collection of nearly horizontal flux tubes embedded in a more vertical background field. The uncombed model is, by far, the most successful representation of the fine structure of the penumbra currently available. It explains the polarization profiles of visible and near-infrared lines observed in sunspots (e.g., Beck 2008), including their net circular polarization (NCP). This success is not trivial, since the behavior of the NCP depends on the details of the magnetic and velocity fields in a very subtle way (see M\\\"uller et al.\\ 2002; Borrero et al.\\ 2007; Tritschler et al.\\ 2007, and references therein). The uncombed model is supported not only by observations, but also by theoretical work. Schlichenmaeier et al.\\ (1998) and Schlichenmaier (2002) performed numerical simulations of penumbral flux tubes in the thin tube approximation. These calculations show filaments whose morphology and dynamics are very similar to those actually observed in the penumbra (Schlichenmaier 2003). Moreover, the simulations offer a natural explanation for the Evershed flow, the most conspicuous dynamical phenomenon of sunspots. In spite of these achievements, it is still not known whether magnetic flux tubes can also account for the existence of dark cores in penumbral filaments and, more importantly, for the surplus brightness of the penumbra. Schlichenmaier \\& Solanki (2003) suggested that hot upflows along magnetic flux tubes would indeed be able to heat the penumbra to the required degree if the tubes return to the solar interior after they have released their energy in the photosphere. At that time the existence of opposite-polarity field lines in the penumbra was unclear, but now it is a well-established observational fact: submerging flux tubes have been detected from Stokes inversions and even imaged directly by Hinode (Sainz Dalda \\& Bellot Rubio 2008). Thus, the Evershed flow remains the best candidate to explain the brightness of the penumbra.   Another possibility is that the dark-cored filaments are the manifestation of field-free gaps that pierce the sunspot magnetic field from below. The concept of a gappy penumbra was proposed by Spruit \\& Scharmer (2006) and Scharmer \\& Spruit (2006) as an alternative way to explain the surplus brightness of the penumbra, on the assumption that the Evershed flow is not sufficient. The gaps would sustain normal convection, carrying heat to the solar surface. Radiative transfer calculations need to be performed to show that a penumbra consisting of field-free gaps is able to explain the corpus of spectropolarimetric observations accumulated over the years. In its present form, however, the gappy model is bound to experience substantial difficulties when confronted with the observations (Bellot Rubio 2007).    \\begin{figure*}[t] \\begin{center} \\resizebox{.39\\hsize}{!}{\\includegraphics[bb=-8 8 275 191,clip]{fig1_B_vec.eps}} \\resizebox{.39\\hsize}{!}{\\includegraphics[bb=0 5 283 188,clip]{fig1_B.eps}} \\\\ \\resizebox{.39\\hsize}{!}{\\includegraphics[bb=0 0 283 203,clip]{fig1_gamma_map_40.eps}} \\resizebox{.39\\hsize}{!}{\\includegraphics[bb=0 0 283 203,clip]{fig1_deltaro.eps}} \\end{center} \\vspace*{-1.3em} \\caption{{\\em Top left:} magnetic field lines in the $xz$-plane. The circle centered at $(0,0)$~km represents the flux tube's boundary. Note the wrapping of the field lines around the tube. {\\em Top right:} field strength distribution.  {\\em Bottom left:} field inclination distribution.  {\\em Bottom right:} gas density distribution, for temperatures in the tube and background given by the cool model of Collados et al.\\ (1994). Shown are density differences with respect to the unperturbed atmosphere (Fig.~\\ref{modeloinicial}).} \\label{magneticfield} \\end{figure*}   Recently, Heinemann et al.\\ (2007) have presented first attempts to simulate the penumbra in 3D. The parameters governing the calculations are still far from those of the real sun and, as a consequence, the model sunspot does not show a typical penumbral pattern. Yet, an interesting result of the simulations is the existence of small blobs of plasma with weaker and more inclined fields than their surroundings. The magnetic properties of these structures are reminiscent of those of the flux tubes postulated by the uncombed model. Some of them show a dark lane similar to the dark cores of penumbral filaments. The dark lanes are produced by locally enhanced density and pressure that shift the $\\tau=1$ level to higher photospheric layers, where the temperature is lower. This effect was identified for the first time by Sch\\\"ussler \\& V\\\"ogler (2006) in magnetoconvection simulations of umbral dots. Interestingly, the parameter regime covered by those simulations is not the one relevant to sunspot penumbrae.  Our aim here is to shed some light on the origin of dark-cored penumbral filaments and the surplus brightness of the penumbra. To that end we solve the 2D stationary heat transfer equation in a stratified uncombed penumbra formed by magnetic flux tubes in a stronger background field (Sects.~\\ref{model} and \\ref{equations}). The tubes carry an Evershed flow of hot plasma. Our calculations show that one such tube would be observed as a dark-cored filament due to the higher density of the plasma within the tube (Sect.~\\ref{results}). We also find that the Everhed flow heats the background atmosphere very efficiently, increasing its temperature. This suggests that the surplus brightness of the penumbra is due to the Evershed flow (Sect.~\\ref{surplus}). Finally, we synthesize polarization maps using the model atmospheres resulting from the simulations and compare them with polarimetric observations of dark-cored filaments (Sect.~\\ref{stokes_maps}). ",
        "conclusions": "The heat transfer and radiative transfer calculations presented in this paper support the concept of a penumbra formed by small (but optically thick) magnetic flux tubes that carry hot flows, as deduced from high-resolution observations and spectropolarimetric measurements (see Solanki 2003 and Bellot Rubio 2004 for reviews). Tubes about 250~km in diameter explain not only the existence of dark-cored filaments, but also the surplus brightness of the penumbra; the Evershed flow efficiently heats the plasma outside the tubes, increasing its temperature to values compatible with the observations.  Further improvements of the model should include a more realistic treatment of the magnetic topology of the tubes and the external atmosphere, a better description of convection, and a full 3D solution of the heat transfer equation. In our opinion, however, these improvements will not change the conclusion that the uncombed model is the best representation of the penumbra at our disposal."
    },
    "0806/0806.2340_arXiv.txt": {
        "abstract": "We argue that, at least a fraction of the newly discovered population of ultra-faint dwarf spheroidal galaxies in the Local Group constitute the fossil relic of a once ubiquitous population of dwarf galaxies formed before reionization with circular velocities smaller than $v_{c}^{cr} \\sim 20$~km/s. We present several arguments in support of this model. The number of luminous Milky Way satellites inferred from observations is larger than the estimated number of dark halos in the Galaxy that have, or had in the past, a circular velocity $>v_{c}^{cr}$, as predicted by the ``Via Lactea'' simulation. This implies that some ultra-faint dwarfs are fossils. However, this argument is weakened by recent results from the ``Aquarius'' simulations showing that the number of Galactic dark matter satellites is 2.5 larger than previously believed. Secondly, the existence of a population of ultra-faint dwarfs was predicted by cosmological simulations in which star formation in the first minihalos is reduced -- but not suppressed -- by radiative feedback.  Here, we show the statistical properties of the fossil galaxies in those simulations are consistent with observations of the new dwarf population and with the number and radial distribution of Milky Way satellites as a function of their luminosity.  Finally, the observed Galactocentric distribution of dwarfs is consistent with a fraction of dSphs being fossils. To make our case more compelling, future work should determine whether stellar chemical abundances of simulated ``fossils'' can reproduce observations and whether the tidal scenarios for the formation of Local Group dwarf spheroidals are equally consistent with all available observations. ",
        "introduction": "Hierarchical formation scenarios predict that most of the galactic halos that formed before reionization (at $z=6-10$) had masses below $10^8-10^9$ M$_{\\odot}$.  Those that survived to the modern epoch, if they were able to form stars, would constitute a sub-population of dwarf satellites orbiting larger halos.  N-body simulations of cold dark matter (CDM) predict a number of dark matter halos around the Milky Way and M31 that is much greater than the number of observed luminous satellites \\citep{Klypinetal99, Mooreetal99}. This may indicate a problem with the CDM paradigm or that feedback processes are very efficient in suppressing star formation in the first small mass halos, which remain mostly dark.  Recent observational and theoretical advances require a re-visitation of the missing galactic satellite problem.  In addition, cosmological simulations of the formation of the first galaxies have shown most previously known dwarf spheroidals (dSphs) to have properties consistent with the surviving first galaxies, and predicted the existence of an undiscovered, lower surface brightness population of dwarfs \\citep{RicottiGnedin05} (hereafter, RG05). The recent discovery of a population of ultra-faint dwarfs confirms the aforementioned theoretical expectation and allows us to test in great detail cosmological simulations of the first galaxies, addressing important theoretical questions on feedback in the early universe and on the nature of dark matter.  The formation of the first dwarf galaxies - before reionization - is regulated by complex feedback effects that act on cosmological scales. These self-regulation mechanisms have dramatic effects on the number and luminosity of the first, small mass galaxies, and yet, are unimportant for the formation of galaxies more massive than $10^8$~M$_\\odot$, that may form before and after reionization. Galaxy formation in the high redshift universe is peculiar due to (i) the lack of important coolants, such as carbon and oxygen, in a gas of primordial composition and (ii) the small typical masses of the first dark halos.  The gas in halos with circular velocity smaller than $20$~km~s$^{-1}$ (roughly corresponding to a mass $M\\simlt 10^8$ M$_{\\odot}$ at the typical redshift of virialization) is heated to temperatures $\\simlt 10,000$ K during virialization.  At this temperature, a gas of primordial composition is unable to cool and initiate star formation unless it can form a sufficient amount of H$_2$ ($x_{H_2} \\simgt 10^{-4}$).  Although molecular hydrogen is easily destroyed by far ultraviolet (FUV) radiation (negative feedback), its formation can be promoted by hydrogen ionizing radiation emitted by massive stars, through the formation of H$^-$ (positive feedback) \\citep{HaimanReesLoeb96}.  It is difficult to determine the net effect of radiative feedback on the global star formation history of the universe before the redshift of reionization.  The effect of a dominant FUV background (at energies between $11.34$~eV and $13.6$~eV) is to destroy H$_2$, the primary coolant at the start of galaxy formation. The FUV radiation emitted by the first few Population~III stars is sufficient to suppress or delay galaxy formation in halos with circular velocities $v_{c}< v_c^{cr} \\sim 20$~km s$^{-1}$. According to this scenario, most halos with masses $<10^8-10^9$ ~M$_\\odot$ do not form stars and remain dark. Therefore, the number of pre-reionization fossils in the Local Group would be expected to be very small or zero.  However, this model does not take into account the effect of ionizing radiation and ``positive feedback regions'' \\citep{RicottiGnedinShull01, Ahnetal06, Whalenetal07}, that may have a dominant role in regulating galaxy formation before reionization \\citep{RicottiGnedinShull02a, RicottiGnedinShull02b}. Simulations including these processes show that star formation in the first small mass halos is inefficient, mainly due to winds produced by internal UV sources.  This produces galaxies that are extremely faint and have very low surface brightnesses.  However, according to the results of our simulations, a large number of these ultra-faint dwarfs (a few hundred galaxies per co-moving Mpc$^3$) form before reionization at $z \\sim 7-10$. Hence, according to this model, the Local Volume and the Local Group should contain hundreds of ultra-faint dwarf galaxies.  The small masses of the first minihalos have two other implications. First, the ionizing radiation emitted by massive stars can blow out most of the gas before SN-driven winds become important, further reducing star formation rates \\citep{RicottiGnedinShull08}.  Second, the increase in temperature of the intergalactic medium (IGM) to $10,000-20,000$ K due to \\HI~reionization, prevents the gas from condensing into newly virialized halos with circular velocities smaller than $10-20$~km~s$^{-1}$, with a critical value that depends on redshift \\citep{Gnedin-filteringmass00, OkamotoGT08}. It follows, that dwarf galaxies with $v_c < 10-20$~km~s$^{-1}$ stop forming stars after reionization. However, the value $v_c^{cr}\\sim 20$~km~s$^{-1}$ that we use to define a fossil is primarily motivated by the fundamental differences discussed above in cooling and feedback processes that regulate star formation in the early Universe and is not the critical value for suppression of star formation due to reionization.  Indeed, \\cite{Ricotti:08} argues that fossils dwarfs can have a late phase of gas accretion and star formation well after reionization, at redshift $z<1-2$. Thus, a complete suppression of star formation after reionization (about 12~Gyr ago) is not a defining property of a fossil dwarf.  Data on the velocity dispersion of the stars, $\\sigma$, the only observational measure of halo mass, shows that a typical dSph has $\\sigma <20$~km~s$^{-1}$. However, the circular velocity of the dark halo can be much larger than $\\sigma$ if typical radius of the stellar spheroid is much smaller than the dark halo radius. For a pre-reionization fossil, according to simulations in RG05, on average we measure $\\sigma/v_c \\sim 0.5$ at formation (see Fig.~2 in RG05).  RG05 compared the statistical properties of simulated pre-reionization galaxies to observations of Local Group dwarfs available in 2004-2005. Based on similarities between observed dSphs and simulated galaxies formed before reionization, they argued that many dSphs may be ``fossils'' of the first galaxies.  RG05 also predicted the existence of a yet undetected population of ultra-faint dSph galaxies. \\citep{GnedinKravtsov06} (hereafter, GK06) used results from RG05 to predict the radial distribution of fossils around the Milky Way and their Galactocentric luminosity function. They found a good agreement of simulations with observations for the most luminous fossils but, again, a lack of observed ultra-faint fossils, mainly in the outer Milky Way halo.  An ultra-faint population of dSphs has now been found in the Local Group.  As we will show in the present work, these new dwarfs have properties in agreement with our simulations of pre-reionization fossils. This discovery is certainly one of the most exciting developments in understanding galaxy formation in the early universe and has drawn renewed attention to ``near field cosmology'' as a tool to understand galaxy formation.  The new galaxies have been discovered by data mining the SDSS and surveys of the halo around M31, resulting in the discovery of 12 new ultra-faint Milky Way satellites \\citep{Belokurovetal06a, Belokurovetal07, Irwinetal07, Walshetal07, Willmanetal05ApJ, Willmanetal05AJ, Zuckeretal06a, Zuckeretal06b, Gehaetal08} and six new companions for M31 \\citep{Ibataetal07, Majewskietal07, Martinetal06}.  Here, we also argue that solely based on the observed number of Milky-Way satellites, at least a few of them must be a pre-reionization fossil, that means: it must have formed before reionization in a halo with circular velocity $<20$ km/s. However, some ultra-faint dwarfs may not be pre-reionization fossils. As we write, several works have been published that seem to show that the observed properties of the ultra-faint dwarf population can also be explained in the context of the tidal scenario, that assumes that these galaxies formed after reionization in halos that were much more massive and had different properties at the time of formation \\citep{PenarrubiaNM08}.  It is intriguing that both the tidal scenario and the pre-reionization fossil scenario are able to produce a population of ultra-faint dwarfs that follow very similar statistical trends in terms of size, surface brightness, mass to light ration and metallicity-luminosity relation. The jury is still out.  This paper is laid out as follows. In \\S~\\ref{sec:data}, we collect published data on the new dwarf population and, after correcting for completeness of the surveys, we estimate the total number of Local Group satellites (which increases from 32 to about 100). Using the results of published N-body simulations, we compare the observed number of luminous satellites to the estimated number of dark satellites that have or had in the past a circular velocity $>v_{c}^{cr}$, using the results of published N-body simulations, concluding that some ultra-faint dwarfs must be pre-reionization fossils.  In \\S~\\ref{sec:comp} we show that the properties of the new Milky Way and M31 dwarfs are in remarkable agreement with the theoretical data on the ``fossils'' from RG05 and with their Galactocentric distribution around the Milky Way calculated in GK06. In \\S~\\ref{sec:disc} we discuss the implications of the new dwarfs on the formation of the first galaxies and the missing galactic satellite problem. ",
        "conclusions": "\\label{sec:disc}  There are two main ideas for the origin of dSphs in the Local Group. Most importantly, these two ideas have very different implications for models of galaxy formation, and the minimum mass a dark halo needs to host a luminous galaxy. The ``tidal scenario'', holds the dwarfs we see today were once far more massive, having been stripped of most of their dark matter during interactions with larger galaxies \\citep[\\eg,][]{KravtsovGnedinKlypin04}.  In this model, we would expect the halos with original dark matter masses below $10^{8}$ ~M$_{\\odot}$ to be mostly dark at formation and at the modern epoch. The ``primordial scenario'', has dwarf galaxies starting with close to their current stellar mass of about $10^3-10^6$~M$_{\\odot}$ and, with several dark halos with mass at formation below the threshold of about $2 \\times 10^{8}$~M$_{\\odot}$ hosting a luminous galaxy.  Star formation in halos this small is possible only before reionization and is widespread if ``positive'' feedback plays a significant role in regulating star formation in the first galaxies \\citep{RicottiGnedinShull01,RicottiGnedinShull02a,RicottiGnedinShull02b}.  In this paper, we argue that the recent discovery of the ultra-faint dwarfs in the Milky Way and M31 supports the ``primordial scenario''. The existence of the ultra-faint dwarfs was predicted by simulations of the formation of the first galaxies (see RG05) and, as shown in the present work, the observed properties of this new population are consistent with them being the ``fossils'' of the first galaxies.  While tidal stripping can reproduce properties of an individual galaxy, it is unable to completely reproduce all the trends in the ultra-faint population.  This is primarily seen in the kinematics of the ultra-faint dwarfs.  Tidal stripping predicts a steeper than observed drop in $\\sigma$ with $L_V$ \\citep{PenarrubiaNM08}, while our simulations show primordial dwarfs which match the observed trends in $\\sigma$ extremely well. It has not been shown yet that star formation in dwarf galaxies more massive than $10^8-10^9$~M$_\\odot$ can reproduce the observed properties of ultra-faint dwarfs without requiring tidal stripping of stars.  The tidal model predicts that gas rich dIrr loose their gas and transform into dSphs due to tidal or ram pressure interaction with a host halo.  And XII, which shows a proper motion close to current published escape velocity of M31, may be on its first approach to the Local Group \\citep{Martinetal06, Chapmanetal07}.  A similar situation exists for And XIV.  With a dynamical mass of $M \\sim 3\\times10^7 M_{\\odot}$, And XIV has $v > v_{esc}^{M31}$, suggesting it is also just entering the Local Group \\citep{Chapmanetal07}.  In the tidal model \\citep{Mayeretal07,Mayeretal06}, And XII and And XIV would be expected to still harbor significant reservoirs of gas, however, observations show And XIV has $M_{HI} < 3\\times10^3 M_{\\odot}$ \\citep{Chapmanetal07} and And XII has no detected \\HI~ \\citep{Martinetal06}.  If neither of these dwarfs have undergone significant tidal interactions with their hosts, as their velocities suggest, how did they lose their gas?  Though its velocity is unknown, the recently discovered And XVIII \\citep{McConnachieetal08}, shows the same characteristics.  At a distance of $600$~kpc from M31 and 1.35 Mpc from the Milky Way, it is unlikely that And XVIII has undergone significant interaction with either Local Group spirals.  And XVIII is classified as a dSph with no detected \\HI~and is similar to the Cetus and Tucana dwarfs \\citep{McConnachieetal08}, both of which are good candidate fossil galaxies (RG05).  On the opposite end of the spectrum is the strange case of Leo~T, the properties of which are discussed in Section 2.1.1.  While, Leo~T has an \\HI~mass fraction typical of dIrr, its other properties are indistinguishable from the other newly discovered ultra-faint dwarfs \\citep{SimonGeha07}, all of which are dSph.  Leo~T's large distance from its host, \\HI~reservoir and low probability of recent tidal interactions \\citep{deJongetal08} make it a good candidate for a precursor to a dSph in the tidal scenario.  Particularly given that Leo T's dynamical mass within the stellar spheroid is small: $8.2\\times10^6$~M$_\\odot$ \\citep{SimonGeha07}, its gas is unlikely to survive a single tidal encounter intact.  Therefore, Leo T may have formed at or near its current mass, and the striking similarity of Leo T to its ultra-faint counterparts suggests that they too could have formed as primordial dwarfs at their current masses.  By our definition, pre-reionization fossils are dwarfs that form before reionization in dark halos with $v_{c}<v_c^{cr}\\sim 20$~km~s$^{-1}$, while non-fossils dwarfs form in halos with $v_{c}>v_c^{cr}$ before and after reionization. The value $v_c^{cr}\\sim 20$~km~s$^{-1}$ that we use to define a fossil is primarily motivated by fundamental differences in cooling and feedback processes that regulate star formation in these halos in the early Universe. This value of the circular velocity is also very close to estimates based on the suppression of star formation in dwarfs after reionization \\citep{Gnedin-filteringmass00, OkamotoGT08}. However, as argued in \\citep{Ricotti:08}, fossils dwarfs can have a late phase of gas accretion and star formation well after reionization, at redshift $z<1-2$. Thus, a complete suppression of star formation after reionization is not necessarily what defines a ``fossil galaxy''.  The number of Milky Way dark satellites that have or had in the past $v_{c}>v_c^{cr}$ can be estimated using the results of published N-body simulations (see \\S~\\ref{ssec:count}). We find that using the via Lactea N-body simulation there are approximately $N_{dark} \\approx 73 \\pm 16$ halos with $v_{c}>20$~km~s$^{-1}$ within the virial radius \\citep{Diemandetal07a}.  The new Aquarius simulations \\citep{Springeletal08}, however, show a factor of 2.5 increase in the number of halos with $v_{c}>20$~km~s$^{-1}$, \\ie, $N_{dark}\\sim 182 \\pm 40$ dark halos. Within a distance of $200$~kpc we estimate $N_{dark} \\approx 36 \\pm 8$ for the Via Lactea and $N_{dark} \\approx 91 \\pm 20$ for the Aquarius simulation.  If the number of observed dwarf satellites within the Milky Way (after applying completeness corrections) is larger than $N_{dark}$ we must conclude that some satellites are fossils.  Twelve new ultra-faint dwarfs have been discovered around the Milky Way by analyzing SDSS data in a region that covers about 1/5 of the sky. Applying a simple correction for the sky coverage we estimate that there should be about at least $85 \\pm 14$ Milky Way satellites. However, the data becomes incomplete for ultra-faint dwarf that are further than about $200$~kpc from the Galactic center. Comparing this number of luminous satellites to $N_{dark}$ within $200$~kpc we cannot conclusively conclude that some ultra-faint dwarfs are fossils because $N_{dark}$ for the Aquarius simulation is comparable to the estimate number of luminous satellites.  Once both sensitivity and survey area corrections are applied, \\cite{Tollerudetal08} estimates the existence of $300$ to $600$ luminous satellites within the virial radius ($R_{vir} \\sim 400$ kpc) of the Milky Way and $120$ within $200$~kpc. Comparing $N_{dark}$ to the \\cite{Tollerudetal08} estimates of the number of luminous Milky Way satellites implies that a significant fraction of them are fossils (regardless if we use the Via Lactea or the Aquarius simulations estimates for $N_{dark}$).  In Table~\\ref{tab:count} we have summarized the aforementioned results.  Another argument for the existence of fossils is provided by detailed comparison of the Galactocentric distribution of fossils in the Milky Way (GK06). Based on these comparison GK06 find that about $1/3$ of Milky Way dwarfs may be fossils.  In this paper, we show the GK06 theoretical results in comparison to updated observational data, including the new ultra-faint dwarfs found using SDSS data, and applying completeness correction due to the limited area surveyed by the SDSS (about 1/5 of the sky).  Assuming that the Local Group has a mass of $3 \\times 10^{12}$ M$_{\\odot}$, as in the GK06 simulation, we find that there are no ``missing galactic satellites'' with $L_V \\geq 10^5$~L$_{\\odot}$ within the virial radius of the Milky Way.  When the new dwarfs are included, the observed and predicted numbers of satellites agree near the Milky Way, however, for distances greater than $200$~kpc, it is clear that there is still a 'missing' population of dwarfs.  However, given that for $d >200$~kpc, dwarfs with $L_V \\sim 10^5$~L$_{\\odot}$ drop below SDSS detection limits (Koposov et al, 2007), the under-abundance of observed dwarfs at large distances is not surprising and likely due to the SDSS sensitivity limit.  A final comment regards the cosmological model. The RG05 and GK06 simulations use cosmological parameters from WMAP~1.  N-body simulations show that the number $N(M)$ of Milky Way dark matter satellites as a function of their mass is not overly sensitive to the cosmology, although there are some differences on the number of the most massive satellites \\citep{Madauetal08}.  However, $N(v_{max})$ should be sensitive to the cosmology \\citep{ZentnerBullock03}, and changes of $\\sigma_8$ and $n_s$ may affect the occupation number and Galactocentric distribution of luminous halos.  The collapse time of small mass halos in high density regions probably dominates the 20\\% variations in $\\sigma_8$ between WMAP~1 and WMAP~3, limiting effects due to the cosmology near large halos.  A decrease in luminous dwarf numbers, due to the lower $\\sigma_8$, could be evident in the distribution of the lowest mass luminous halos in the voids.  In conclusion, the number of Milky Way and M31 satellites provides an indirect test of galaxy formation and the importance of positive feedback in the early universe. Although the agreement of the SDSS and new M31 dwarfs' properties with predictions from the RG05 and GK06 simulations does not prove the primordial origin of the new ultra-faint dwarfs, it supports this possibility with quantitative data and more successfully than any other proposed model has been able to do so far. At the moment, we do not have an ultimate observational test that can prove a dwarf galaxy to be a fossil. Even a test based on measuring the SFH of the dwarf galaxies may not be discriminatory because, as has been recently suggested, fossil galaxies may have a late phase of gas accretion and star formation at $z<1-2$, during the last $9-10$~Gyrs \\citep{Ricotti:08}. The distinction between fossils and non-fossils galaxies thus is quite tenuous and linked to our poor understanding of star formation and feedback in dwarf galaxies. Arguments based on counting the number of dwarfs in the Local Universe probably provide the most solid argument to prove or disprove the existence of fossil galaxies. In the future, a possible test may be provided by deep surveys looking for ultra-faint or dark galaxies in the local voids. Some fossil dwarfs should be present in the voids if they formed in large numbers before reionization."
    },
    "0806/0806.0345_arXiv.txt": {
        "abstract": "We use sequences of images and magnetograms from Hinode to study magnetic elements in internetwork parts of the quiet solar photosphere. Visual inspection shows the existence of many long-lived (several hours) structures that interact frequently, and may migrate over distances $\\sim7~\\Mm$ over a period of a few hours. About a fifth of the elements have an associated bright point in G-band or \\CaIIH\\ intensity. We apply a hysteresis-based algorithm to identify elements. The algorithm is able to track elements for about $10~\\min$ on average. Elements intermittently drop below the detection limit, though the associated flux apparently persists and often reappears some time later. We infer proper motions of elements from their successive positions, and find that they obey a Gaussian distribution with an rms of $1.57\\pm0.08~\\kms$. The apparent flows indicate a bias of about $0.2~\\kms$ toward the network boundary. Elements of negative polarity show a higher bias than elements of positive polarity, perhaps as a result of to the dominant positive polarity of the network in the field of view, or because of increased mobility due to their smaller size. A preference for motions in $X$ is likely explained by higher supergranular flow in that direction. We search for emerging bipoles by grouping elements of opposite polarity that appear close together in space and time. We find no evidence supporting Joy's law at arcsecond scales. ",
        "introduction": "\\label{sec:introduction}  Magnetic elements have been extensively studied in network that partially outlines the boundaries of supergranular cells. They were first observed as ``magnetic knots'' \\citep{1968SoPh....4..142B} and as ``filigree'' \\citep{1973SoPh...33..281D}, before being resolved into strings of adjacent bright points by \\cite{1974SoPh...38...43M}. \\cite{1977SoPh...52..249M} and \\cite{1981SoPh...69....9W} showed that faculae, filigree, and bright points in wide-band \\CaIIH\\ filtergrams are manifestations of the same phenomenon. \\cite{1983SoPh...85..113M} introduced the name ``network bright point'', and subsequently initiated extensive studies of magnetic elements as G-band bright points \\citep{1984SoPh...94...33M}. Studies of bright points using high-resolution imaging \\citep[e.g.,][]{1995ApJ...454..531B, 1998ApJ...506..439B, 1998ApJ...495..973B, 2004A&A...428..613B, 1996ApJ...463..365B, 2001ApJ...553..449B, 2005A&A...435..327R, 2007A&A...466.1131R, 2007A&A...472..911I, 2007ApJ...661.1272B} have since established that network bright points are manifestations of small, kilogauss magnetic elements that form the magnetic network \\citep{1968SoPh....5..442C, 1969SoPh...10..294L, 1972SoPh...22..402H, 1972SoPh...27..330F, 1973SoPh...32...41S}.  Magnetic field in internetwork have been largely ignored until recently. However, it is currently being studied vigorously \\citep[e.g.,][]{2003ApJ...582L..55D, 2003ApJ...597L.177S, 2004ApJ...611.1139S, 2004ApJ...613..600L, 2004ApJ...616..587S, 2004Natur.430..326T, 2004ApJ...614L..89M, 2005A&A...436L..27K, 2006ApJ...636..496C, 2007A&A...476L..33R, 2007ApJ...657.1150S, 2007ApJ...659L.177H, 2007ApJ...670L..61O, 2008ApJ...672.1237L}. Many of these studies focus on determining the strength and distribution of flux. While there is some disagreement between results, it seems that field is ubiquitously present in internetwork at small scales.  Concentrations of flux that are sufficiently strong may form internetwork bright points. Their existence was already noted by \\cite{1983SoPh...85..113M}. Few recent studies have analyzed these internetwork magnetic elements. \\cite{2004ApJ...609L..91S} measured internetwork bright point density and lifetime, \\cite{2005A&A...441.1183D} reported that internetwork bright points trace locations of flux that may persist for periods of hours, \\cite{2007A&A...462..303T} analyzed morphology, dynamics, and evolution of bright points in \\CaIIK\\ in quiet sun, and \\cite{2007A&A...475.1101S} searched for photospheric foot points of transition region loops in quiet sun.  Magnetic elements were first modeled as ``flux tubes'' by \\cite{1976SoPh...50..269S}. Over the years, models grew increasingly complex \\citep[e.g.,][]{1988A&A...202..275K, 1990A&A...233..583K, 1992A&A...262L..29S, 1994A&A...285..648G, 1998A&A...337..928G, 1998ApJ...495..468S, 2005A&A...430..691S}. The explanation of photospheric brightness enhancement of faculae due to hot walls proposed by \\cite{1981SoPh...70..207S} was verified by MHD simulations by \\cite{2004ApJ...607L..59K} and \\cite{2004ApJ...610L.137C}. On disk, bright points are formed as a result of radiation escaping from deeper, hotter layers due to the fluxtube Wilson depression.  Some authors have noted that magnetic fields in internetwork areas appear to outline cells on mesogranular scales \\citep[e.g.,][]{2003ApJ...582L..55D, 2005A&A...441.1183D, 2007A&A...462..303T, 2008ApJ...672.1237L}, while \\citet{1998ApJ...495..973B} observed ``voids'' in active network. In addition, recent simulations indicate that field concentrates on boundaries of mesogranular cells \\citep{2006ApJ...642.1246S}. One would expect such a pattern to be set by granular motions, similar to supergranular flows that eventually advect internetwork field into network \\citep[e.g.,][]{2000SoPh..197...21L}. Perhaps magnetic elements form these patterns as a result of flows associated with ``trees of fragmenting granules'' \\citep{2004A&A...419..757R}, which were previously linked to mesogranules by \\citet{2003A&A...409..299R}. Flux is expunged by the sideways expansion of granular cells, and is collected in the downflows in intergranular lanes. In a ``tree of fragmenting granules'', these flows would be expected to drive flux not only to the edges of individual granules, but also to the edges of the tree.  In this paper, we present a study of the dynamics of magnetic elements in internetwork parts of the solar photosphere. This study is motivated by its relevance to the operation of a turbulent granular dynamo, the nature of quiet-sun magnetism, the generation of MHD waves that may propagate into the transition region and corona, and the coupling of internetwork field to the magnetic network. First, examples of magnetic elements are discussed in the context of fluxtube dynamics and lifetime (Sect.~\\ref{sec:visual}). Magnetic elements are compared with bright points in G-band and \\CaIIH\\ intensity in Sect.~\\ref{sec:comparebp}. A feature-tracking algorithm is applied in order to analyze the lifetime (Sect.~\\ref{sec:lifetime}) and the dynamics (Sects.~\\ref{sec:velocities} and~\\ref{sec:direction}) of magnetic elements. Finally, a search for emerging bipoles is presented in Sect.~\\ref{sec:emerging}. ",
        "conclusions": "\\label{sec:conclusion}  We have analyzed the dynamics of IMEs using a sequence of magnetograms. Visual inspection of the data shows the existence of many long-lived magnetic elements that have frequent interactions with other elements during their lifetime. We find that they may migrate over distances of $\\sim7~\\Mm$ in periods of several hours. Their interactions, migration, and the many short-lived concentrations of flux that appear in their vicinity make it cumbersome to uniquely identify an IME, or a set of IMEs. The IMEs sometimes drop below the detection level in our data, but commonly reappear some time later. IMEs appear to outline cells on scales of several Mm.  A manual inspection of IME locations in the G-band and \\CaIIH\\ filtergrams shows that only about a fifth of the IMEs have associated bright points. Visual inspection reveals no obvious correlation between IME morphology and the existence of bright points. There is a substantial correlation between the existence of a bright point in the G-band and \\CaIIH\\ intensities. Bright points are formed in the G band through weakening of molecular CH lines as a result of partial evacuation of the magnetic element, while they are caused in \\CaIIH\\ by influx of radiation from the hot walls of the Wilson depression. We therefore attribute the similarity in the appearance of bright points in these passbands to their common origin.  We identity magnetic elements using a hysteresis-based algorithm that is able to track IMEs for about $10~\\min$ on average. This is much shorter than the lifetime of several hours predicted by \\citet{2005A&A...441.1183D}. However, visual inspection shows that many elements intermittently drop below the detection limit of our algorithm, shortening their measured lifetime.  IMEs exhibit proper motions that resemble a Gaussian distribution with a slight overdensity of velocities near the origin and in the far wings. We measure an rms velocity of $1.57\\pm0.08~\\kms$.  The IMEs show a slight bias of about $0.2~\\kms$ for velocities toward the nearest network boundary. This bias persists up to timescales of at least $1200~\\s$. It is the likely cause of weak positive velocity autocorrelation at long delay times. In the data analyzed here, IMEs of negative polarity show a statistically significant higher drift to the nearest network boundary. It is tempting to assume that this difference is somehow related to the dominant positive polarity of the network in the field of view. Alternatively, elements of negative polarity may be more mobile compared to elements of positive polarity because of their smaller size.  We find a slight preference for velocities in $X$~direction. Elongation in $X$ of the main supergranular cell in the field of view suggests that supergranular flows are stronger in that direction. The observed preference in direction is thus in agreement with the classical picture that IMEs are advected by supergranular flows.  IMEs experience centrifugal accelerations that obey a Gaussian distribution with an average of $0.045\\pm0.013$ and an rms of $4.559\\pm0.011~\\mathrm{deg}\\;\\mathrm{km}/\\mathrm{s}^2$, in agreement with results of \\citet{2003ApJ...587..458N}.  We search for emerging bipoles by pairing elements of opposite polarity that appear nearby each other in space and time. There is no detectable preference in the orientation of a pair. We conclude, therefore, that there is no evidence to support Joy's law at arcsecond scales from these data."
    },
    "0806/0806.4616_arXiv.txt": {
        "abstract": "We present high-resolution (0.3$''$) Very Large Array (VLA) imaging of the molecular gas in the host galaxy of the high redshift quasar PSS\\,J2322+1944 ($z=4.12$).  These observations confirm that the molecular gas (CO) in the host galaxy of this quasar is lensed into a full Einstein ring, and reveal the internal dynamics of the molecular gas in this system. The ring has a diameter of $\\sim$1.5$''$, and thus is sampled over $\\sim$20 resolution elements by our observations.  Through a model-based lens inversion, we recover the velocity gradient of the molecular reservoir in the quasar host galaxy of PSS\\,J2322+1944. The Einstein ring lens configuration enables us to zoom in on the emission and to resolve scales down to $\\lesssim$1\\,kpc. From the model-reconstructed source, we find that the molecular gas is distributed on a scale of 5\\,kpc, and has a total mass of $M({\\rm H_2})=1.7 \\times 10^{10}\\,$\\,\\msol. A basic estimate of the dynamical mass gives $M_{\\rm dyn} = 4.4 \\times 10^{10}\\,\\sin^{-2}\\,i$\\,\\msol, that is, only $\\sim$2.5 times the molecular gas mass, and $\\sim$30 times the black hole mass (assuming that the dynamical structure is highly inclined).  The lens configuration also allows us to tie the optical emission to the molecular gas emission, which suggests that the active galactic nucleus (AGN) does reside within, but not close to the center of the molecular reservoir.  Together with the (at least partially) disturbed structure of the CO, this suggests that the system is interacting.  Such an interaction, possibly caused by a major `wet' merger, may be responsible for both feeding the quasar and fueling the massive starburst of 680\\,\\msol\\,yr$^{-1}$ in this system, in agreement with recently suggested scenarios of quasar activity and galaxy assembly in the early universe. ",
        "introduction": "A fundamental aspect in studies of galaxy formation and evolution is to understand the connection between AGN and starburst activity. The existence of a physical connection between both processes is suggested by the finding that present day galaxies show a strong relationship between the mass of their central supermassive black holes (SMBHs) and the mass and concentration of their stellar spheroids (Magorrian \\etal\\ \\citeyear{mag98}; Ferrarese \\& Merritt \\citeyear{fer00}; Gebhardt \\etal\\ \\citeyear{geb00}; Graham \\etal\\ \\citeyear{gra01}). If these relations were due to a coevolution of both components during the early assembly of a galaxy, high-redshift quasars and their associated host galaxies would be ideal objects to study the active formation of both SMBHs and bulge stars.  Studies of molecular gas (most commonly rotational transitions of CO), the prerequisite material that fuels star formation, have become an important tool to probe distant quasar host galaxies, and revealed large molecular gas reservoirs of $>$10$^{10}$\\,\\msol\\ in a number of these sources (see Solomon \\& Vanden Bout \\citeyear{sv05} for a general review).  These galaxies typically show huge far-infrared (FIR) luminosities in excess of 10$^{13}$\\,\\lsol, which are thought to be powered by starbursts (and possibly a central AGN; e.g., Omont et al.\\ \\citeyear{omo01}; Wang et al.\\ \\citeyear{wan07}). Observations of molecular gas trace the regions that can host massive starbursts. In addition, the velocity structure of molecular line emission has the potential to constrain the dynamical state of galaxies out to the earliest epochs.  Rotational molecular line emission typically emerges at FIR to radio wavelengths, i.e., in the limited wavelength regime where the AGN in distant quasars does not outshine all other emission.  However, the cosmological distances of high redshift quasars make it difficult to resolve the faint emission from their host galaxies a such long wavelengths.  The physical resolution of such observations is in some cases boosted by gravitational lenses acting as natural telescopes. The gravitational lensing effect also magnifies the observed flux of the background galaxy, in particular for systems in Einstein ring configurations. Due to the compactness of the AGN, optical quasars in Einstein ring lens configurations are rare.  Due to their greater extent, the host galaxies of quasars are much more likely to cross the inner Einstein ring caustic of a gravitational lens.  In this paper, we report on high (0.3$''$) angular resolution Very Large Array (VLA)\\footnote{The Very Large Array is a facility of the National Radio Astronomy Observatory, operated by Associated Universities, Inc., under a cooperative agreement with the National Science Foundation.}  observations of CO in the host galaxy of the $z$=4.12 quasar PSS\\,J2322+1944, one of only two known $z$$>$4 galaxies that are both gravitationally lensed and detected in molecular gas emission (the other being BRI\\,0952--0115 at $z$=4.43; Guilloteau \\etal\\ \\citeyear{gui99}). This galaxy was identified in a spectroscopic follow-up study of the Palomar Sky Survey (DPOSS; Djorgovski et al.\\ \\citeyear{djo00}), and found to be a strongly lensed optical quasar (S.~G.\\ Djorgovski, private communication). It was subsequently detected in hard X-ray (Vignali et al.\\ \\citeyear{vig05}), FIR dust (Omont et al.\\ \\citeyear{omo01}; Isaak et al.\\ \\citeyear{isa02}) and radio continuum emission (Carilli et al.\\ \\citeyear{car01}), as well as molecular line emission (Cox et al.\\ \\citeyear{cox02}; Carilli et al.\\ \\citeyear{car02a}). It follows the radio-FIR correlation of star-forming galaxies (Carilli et al.\\ \\citeyear{car01}; Beelen et al.\\ \\citeyear{bee06}), indicating that its FIR continuum emission is dominated by intense star formation.  In spite of the fact that this source shows only two unresolved quasar images in the optical, previous CO observations have shown that the molecular gas reservoir in its host galaxy is lensed into an Einstein ring (Carilli \\etal\\ \\citeyear{car03}; hereafter:\\ C03). These observations were also used to derive a first lensing model for this source.  Based on the dynamical structure revealed by our new, higher resolution observations of PSS\\,J2322+1944, we have developed a new lensing model, which enables us to reconstruct the velocity gradient in the spatially resolved gas reservoir.  We use a concordance, flat $\\Lambda$CDM cosmology throughout, with $H_0$=71\\,\\kms\\,Mpc$^{-1}$, $\\Omega_{\\rm M}$=0.27, and $\\Omega_{\\Lambda}$=0.73 (Spergel \\etal\\ \\citeyear{spe03}, \\citeyear{spe07}). ",
        "conclusions": "We have imaged and modeled a molecular Einstein ring of a galaxy at $z$=4.12. Our high resolution \\bco\\ maps of the lensed quasar host galaxy of PSS\\,J2322+1944 (a double image optical quasar) reveal spatially resolved structure that shows a clear velocity gradient in the CO emission line. By performing a model-based lens inversion of the Einstein ring that is consistent with the data, we are able to reconstruct the velocity structure of this distant quasar host galaxy. The gravitational lensing effect acts as a natural telescope, and allows us to zoom in on the molecular gas reservoir down to linear scales of only $\\lesssim$1\\,kpc, sufficient to reveal velocity structure over almost 10 resolution elements in the source plane. Our novel modeling of this system reveals how the molecular gas crosses the central caustic (causing the appearance of the Einstein ring) moving from the redshifted to the blueshifted molecular emission. We also find evidence that the optical quasar may be associated with the redshifted part of the molecular reservoir.  The full reservoir has a mass of $M({\\rm H_2})=1.7 \\times 10^{10}\\,$\\,\\msol\\ (corrected for lensing magnification).  The molecular gas mass alone could account for almost half of the dynamical mass in this system if the galaxy were to be seen close to edge-on ($M_{\\rm dyn}$\\,sin$^2\\,i$/$M({\\rm H_2})$ $\\simeq$ 2.5).  Due to the large spatial extent of the CO emission, and due to the fact that the AGN is probably largely offset from the center of the reservoir, we conclude that the molecular gas and dust are likely dominantly heated by star formation.  From the FIR luminosity of the source in the adopted cosmology ($L_{\\rm FIR} = 2.4 \\times 10^{13}\\,{\\mu_L}^{-1}$\\,\\lsol; Cox \\etal\\ \\citeyear{cox02}), we derive a star formation rate \\footnote{Assuming SFR=1.5$\\times$10$^{-10}\\,L_{\\rm FIR}$(\\msol\\,yr$^{-1}$/\\lsol) (Kennicutt \\citeyear{ken98a}). The dust temperature and spectral index of \\pss\\ indicate that the FIR continuum emission is dominated by star formation, in agreement with the finding that the source follows the radio-FIR correlation for star-forming galaxies (Beelen et al.\\ \\citeyear{bee06}).} (SFR) of 680\\,\\msol\\,yr$^{-1}$. At least part of the CO emission of the reconstructed source does not appear to follow a systemic trend in velocity.  In this picture, this structure may be due to interaction, possibly caused by a major merger.  Such an event could both feed the AGN and fuel the starburst, and thus be responsible for the coeval assembly of a supermassive black hole and the stellar bulge in this system.  Future observations of the FIR continuum at comparable spatial resolution may shed more light on this situation.  Motivated by these results, the dynamical mass derived from the molecular line observations in the host galaxy of the $z$=4.12 quasar \\pss\\ can be used in an attempt to constrain the relationship between the central SMBH mass and the stellar bulge mass ($M_{\\rm BH}$--$M_{\\rm bulge}$) in high-$z$ AGN galaxies. Such a relation has been proposed for different types of galaxies in the local universe, and appears to hold over more than three orders of magnitude in SMBH mass, essentially independent of galaxy type (predicting $M_{\\rm bulge} \\simeq 700\\,M_{\\rm BH}$, e.g., Kormendy \\& Gebhardt \\citeyear{kg01}). Assuming the optical lensing factor of $\\mu_L^{\\rm opt} = 4.7$ derived from our model of \\pss\\ and the Eddington limit derived in Section 3 gives $M_{\\rm BH} = 1.5 \\times 10^9\\,$\\msol. With the further assumption that the dynamical molecular structure is seen close to edge-on, and that it traces a major fraction of the gravitational potential that hosts the stellar bulge, we thus find that $M_{\\rm bulge} \\simeq 30\\,M_{\\rm BH}$ (subtracting out the black hole and gas masses would even give $M_{\\rm bulge} < 20\\,M_{\\rm BH}$). This value is by more than an order of magnitude offset from the local $M_{\\rm BH}$--$M_{\\rm bulge}$ relation, but in good agreement with results obtained for other high-$z$ quasars, which show similar or larger offsets (e.g., Walter et al.\\ \\citeyear{wal04}; Wei\\ss\\ et al.\\ \\citeyear{wei07}; Riechers et al.\\ \\citeyear{rie07}; see also Shields et al.\\ \\citeyear{shi06}).  \\begin{figure} \\epsscale{1.4} \\plotone{f6.eps} \\caption{Map of the detection probability $P$(pixel $>$ 0) in the source plane after lens inversion, as derived from the full MCMC parameter study. A region of 1.21$''$$\\times$1.21$''$ size is shown. \\label{f4}} \\end{figure}  \\begin{figure} \\epsscale{1.2} \\plotone{f7.eps} \\caption{Model-predicted differential gravitational magnification between the CO velocity channels shown in Figures \\ref{f3} and \\ref{f5}.  The magnification is shown in magnitudes (i.e., no magnification corresponds to $m_L=0$). The error bars include the modeling uncertainties. The numbers on top indicate the magnification in each velocity channel. The dashed line indicates the total magnification of $\\mu_L$=5.34 in the integrated emission line map. \\label{f6}} \\end{figure}   In the case of \\pss, one may attempt to account for this offset by assuming that the dynamical structure is seen close to face-on (i.e., $i < 12^\\circ$).  This would however predict a large intrinsic CO linewidth ($\\Delta v_{\\rm FWHM} > 1300$\\,\\kms ). Such a linewidth would significantly exceed the velocity dispersions observed in the spheroids of massive present day elliptical galaxies (which \\pss\\ will likely eveolve into), but is of the same order of magnitude as those observed toward some high-$z$ submillimeter galaxies (e.g., Carilli \\& Wang \\citeyear{cw06}).  Such large molecular linewidths are consistent with those predicted by simulations of the hierarchical buildup of massive quasar host galaxies at high redshift (e.g., Narayanan et al.\\ \\citeyear{nar07}), and thus compatible with the possible merger nature of \\pss. However, in interacting or merging systems, molecular lines are likely more broad due to the fact that the dynamical molecular structure is not fully virialized yet. The molecular line widths in these galaxies thus may be by a factor of a few higher than the actual virial velocity of the host halo, which would lead to an overprediction of the bulge mass.  It thus appears difficult to explain the offset from the local $M_{\\rm BH}$--$M_{\\rm bulge}$ relation by simply assuming a small inclination angle toward the line of sight.  Also, assuming that the AGN accretes at sub-Eddington rates, and/or taking into account that more than a third of the dynamical mass derived above is likely not stellar, but accounted for by gas, dust ($<$1\\% of the gas mass), and the black hole further increases this offset. Together with previous such examples, our results for \\pss\\ thus appear to indicate that the black holes in massive galaxies at high redshift assemble earlier than a large fraction of their stellar bulges.  The observations and modeling presented herein demonstrate the power of spatially and dynamically resolved molecular gas studies in strongly lensed, distant AGN-starburst systems to provide direct evidence for the scenarios of quasar activity and galaxy assembly in the early universe as suggested by recent cosmological simulations (e.g., Springel \\etal\\ \\citeyear{spr05}). The boost in line intensity and spatial resolution provided by Einstein ring lens configurations are currently the only means by which to probe the dynamical structure of the most distant star-forming galaxies at (sub-)kiloparsec resolution. Such observations provide an important foundation for future observations of molecular gas and dust in the early universe with the Atacama Large Millimeter/submillimeter Array (ALMA), which will be able to probe more typical galaxy populations at high redshift to comparable and higher physical resolution, even without the aid of gravitational lensing."
    },
    "0806/0806.3022_arXiv.txt": {
        "abstract": "We report new continuum observations of fourteen z$\\sim$6 quasars at 250 GHz and fourteen quasars at 1.4 GHz. We summarize all recent millimeter and radio observations of the sample of the thirty-three quasars known with $\\rm 5.71\\le z\\le 6.43$, and present a study of the rest frame far-infrared (FIR) properties of this sample. These quasars were observed with the Max Plank Millimeter Bolometer Array (MAMBO) at 250 GHz with mJy sensitivity, and $\\rm 30\\%$ of them were detected. We also recover the average 250 GHz flux density of the MAMBO undetected sources at $\\rm 4\\sigma$, by stacking the on-source measurements. The derived mean radio-to-UV spectral energy distributions (SEDs) of the full sample and the 250 GHz non-detections show no significant difference from that of lower-redshift optical quasars. Obvious FIR excesses are seen in the individual SEDs of the strong 250 GHz detections, with FIR-to-radio emission ratios consistent with that of typical star forming galaxies. Most 250 GHz-detected sources follow the $\\rm L_{FIR}$--$\\rm L_{bol}$ relationship derived from a sample of local IR luminous quasars ($\\rm L_{IR}>10^{12}L_{\\odot}$), while the average $\\rm L_{FIR}/L_{bol}$ ratio of the non-detections is consistent with that of the optically-selected PG quasars. The MAMBO detections also tend to have weaker $\\rm Ly\\alpha$ emission than the non-detected sources. We discuss possible FIR dust heating sources, and critically assess the possibility of active star formation in the host galaxies of the z$\\sim$6 quasars. The average star formation rate of the MAMBO non-detections is likely to be less than a few hundred $\\rm M_{\\odot}\\,yr^{-1}$, but in the strong detections, the host galaxy star formation is probably at a rate of $\\rm \\gtrsim10^{3}\\,M_{\\odot}\\,yr^{-1}$, which dominates the FIR dust heating. ",
        "introduction": "More than thirty quasars at z$\\sim$6 have been discovered (e.g., Fan et al. 2000, 2001a, 2003, 2004, 2006a; Jiang et al. 2007a; Willott et al. 2007). These objects are the earliest massive black hole accretion systems known (Jiang et al. 2006; 2007b), seen at an epoch close to the end of cosmic reionization (Fan et al. 2006b). They provide us with unique information on both the growth of supermassive black holes (SMBHs) and the formation of massive galaxies when the age of the universe was $\\rm \\lesssim1\\,Gyr$.  There are fundamental relationships between SMBH mass and bulge stellar mass/velocity dispersion in the local universe (e.g., Tremaine et al. 2002; Marconi \\& Hunt 2003; Hopkins et al. 2007), indicating that the formation of SMBHs and their spheroidal hosts are coupled. Active galactic nuclei (AGNs) buried in dusty starburst environments have been discovered in samples of low redshift ultraluminous infrared galaxies (ULIRGs), which are believed to be a transition phase between starburst mergers and typical optically bright AGNs (Sanders et al. 1988; Wu et al. 1998; Zheng et al. 2002). Studies of large samples of galaxies and AGNs also provide clues about the 'downsizing' process in galaxy and SMBH formation, whereby the formation of the most massive systems occur at early epochs (Cowie et al. 1996; Heckman et al. 2004; Kauffmann \\& Heckman 2005).  Examples of massive quasars hosted by interacting systems or bright submillimeter galaxies have been found at redshifts greater than 4, such as BR 1202-0725 at z=4.7 and BRI1335-0417 at z=4.4 (Omont et al. 1996a, 1996b; Benford et al. 1999; Beelen et al. 2006). These objects are all characterized by strong molecular CO line emission (Guilloteau et al. 1997, 1999; Carilli et al. 2002; Solomon \\& Vanden Bout 2005; Riechers et al. 2006) and FIR (Guilloteau et al. 1999) and radio (Momjian et al. 2005, 2007) emission originating from a starburst, with implied star formation rates of a few thousand $\\rm M_{\\odot}\\,yr^{-1}$. The results suggest that the galaxies are at an earlier evolution stage than are typical optically bright quasars in which the central AGN dominates the emission from X-ray to radio.  These studies have been extended to z$\\sim$6 quasar sample with sensitive submillimeter, millimeter, and radio telescopes (Priddey et al. 2003b; Bertoldi et al. 2003; Petric et al. 2003; Carilli et al. 2004; Wang et al. 2007). About 1/3 of these z$\\sim$6 quasars were detected at millimeter wavelengths, at mJy sensitivity (Priddey et al. 2003b; Wang et al. 2007). The millimeter detections imply FIR luminosities of $\\rm 10^{12}\\sim10^{13}\\,L_{\\odot}$ and dust masses of $\\rm \\gtrsim10^{8}\\,M_{\\odot}$ in the quasar host galaxies (Bertoldi et al. 2003a, Beelen et al. 2006). Such huge dust masses require rapid metal and dust enrichment via a starburst in the early evolution of these galaxies (Bertoldi et al. 2003a; Walter et al. 2003; Venkatesan et al. 2006; Li et al. 2007b; Dwek et al. 2007).  The heating sources of the FIR-emitting dust in the quasar systems at z$\\sim$6 have become a key question. The reprocessed emission from star formation-heated dust can provide a direct estimate of the star formation rate, thus constraining the bulge building stage in these quasar hosts. Observations at submillimeter wavelengths imply that the FIR emission in the two sources with the strongest millimeter detections, SDSS J114816.64+525150.3 at z=6.42 (hereafter J1148+5251) and SDSS J092721.82+200123.7 at z=5.77 (hereafter J0927+2001), is from dust components with temperatures of 50 to 60 K (Beelen et al. 2006; Wang et al. 2008a). Large amounts of highly excited molecular CO were also detected in the host galaxies of these two sources (Bertoldi et al. 2003b; Walter et al. 2003; Carilli et al. 2007), as well as strong [C {\\small II}] 158$\\,\\mu$m ISM gas cooling line emission in J1148+5251 (Maiolino et al. 2005). These results suggest that star formation plays an important role in the heating of FIR-emitting warm dust. The implied star formation rate is $\\rm \\gtrsim10^{3}\\,M_{\\odot}\\,yr^{-1}$, which argues for active bulge building in these two z$\\sim$6 quasars.  We have a long standing program to study the dust and gas emission from the z$\\sim$6 quasars (e.g., Wang et al. 2007, 2008a). In this paper, we report new observations and present an FIR emission study of the full sample of thirty-three z$\\sim$6 quasars discovered to date, aiming (i) to give a general view of the FIR emission properties of the sample, and (ii) to further constrain the dust heating and star forming activity in the quasar hosts. We describe the full quasar sample, the new observations, and summarize the current millimeter and radio results in section 2. The analysis of the full sample is given in Sections 3, 4, and 5. We present a discussion of star formation in Section 6, and give a brief summary in Section 7. We adopt a $\\rm \\Lambda$-model cosmology with $\\rm H_{0}=71km\\ s^{-1}\\ Mpc^{-1}$, $\\rm \\Omega_{M}=0.27$ and $\\rm \\Omega_{\\Lambda}=0.73$ throughout this paper (Spergel et al. 2007). ",
        "conclusions": "We study the SEDs of 33 quasars at z$\\sim$6 from FIR to radio wavelengths. We conclude that, when averaged with the whole sample and the 250 GHz undetected sources, no significant difference is seen between the mean FIR-to-radio SEDs and the templates of low-redshift optically selected quasars. In particular, we extrapolate the FIR emission of the templates with the typical quasar FIR-to-millimeter spectrum (i.e. $\\rm f\\sim\\nu^\\alpha,\\,\\alpha\\geq2$) from Sanders et al. (1989). This extrapolation is consistent with the average FIR emission of the 250 GHz undetected z$\\sim$6 quasars very well (see Section 3.1 and Figure 2). This fact suggests that these sources have a similar AGN dominated dust heating mechanism. The average star formation rate from the host galaxies of the quasars undetected by MAMBO is estimated to be less than a few hundred $\\rm M_{\\odot}\\,yr^{-1}$.  We detect a strong FIR excess in seven of the ten z$\\sim$6 quasars detected in the millimeter band. These FIR luminous sources are likely to be the high-mass counterparts of local IR luminous quasars which are in transition between the starburst phase and the mature quasar phase. Star formation at rates of $\\rm >10^{3}\\,M_{\\odot}\\,yr^{-1}$ may be the dominant mechanism to heat the dust that gives rise to the FIR emission. The millimeter detected sources also show weaker Lya emission compared to that of the non-detections, but the origin of this trend is not clear yet. Further observations of these FIR luminous quasars, especially spatially resolved line and dust studies, are required to further address questions of star formation and FIR dust heating in these extreme objects.  The current sample of z$\\sim$6 quasars is optically selected. There may be other quasar populations at z$\\gtrsim$6 which are still buried in their starburst environment and are obscured in the optical. Sensitive IR and millimeter facilities, such as Spitzer, ALMA, and Herschel, may discover such objects and provide a more complete view of galaxy and SMBH evolution in the early universe."
    },
    "0806/0806.1161_arXiv.txt": {
        "abstract": "Based on CORSIKA and REAS2 simulations, we investigate the dependence of geosynchrotron radio emission from extensive air showers on the energy of the primary cosmic ray and the depth of the shower maximum. It is found that at a characteristic lateral distance, the amplitude of the bandpass-filtered radio signal is directly proportional to the energy deposited in the atmosphere by the electromagnetic cascade, with an RMS uncertainty due to shower-to-shower fluctuations of less than 3\\%. In addition, the ratio of this radio amplitude and that at a larger lateral distance is directly related to the atmospheric depth of the shower maximum, with an RMS uncertainty of $\\sim15$--20$\\,$g$\\,$cm$^{-2}$. By measuring these quantities, geosynchrotron radio emission from cosmic ray air showers can be used to infer the energy of the primary particle and the depth of the air shower maximum on a shower-to-shower basis. ",
        "introduction": "During the last few years, the technique of radio detection of cosmic ray air showers has experienced an impressive renaissance \\citep{FalckeNature2005,HuegeCris2006,ArdouinBelletoileCharrier2005,vandenBergIcrc2007}. The activities are driven by the prospect of establishing a new observing technique with nearly 100\\% duty cycle and very good angular resolution, which would complement the existing particle detectors and air fluorescence telescopes. Due to the coherent nature of the radio emission, i.e.\\ an approximately quadratic scaling of the emitted radio power with the energy of the primary cosmic ray, the technique is particularly well-suited for the detection of ultra-high energy cosmic rays.  One important question is how the observables measured with the radio technique can be related to the energy and mass of the primary cosmic ray particle. In this article, we use the geosynchrotron model for radio emission from cosmic ray air showers \\citep{HuegeFalcke2003a,HuegeFalcke2005a,HuegeFalcke2005b,HuegeUlrichEngel2007a} to demonstrate how shower parameters could be derived from radio-only measurements that allow one to estimate the primary energy and composition. It is shown that the characteristics of the lateral distribution function of the radio emission can be exploited to estimate on a shower-to-shower basis both the energy deposited in the atmosphere and the depth of the shower maximum. The energy determination can be performed with a measurement at an optimum lateral distance which is independent of the energy of the primary particle in the considered energy range from $10^{18}$ to \\unit[$10^{20}$]{eV} --- it only depends on the air shower geometry and radio observing frequency. Combining this measurement with an additional measurement at a different lateral distance provides a handle on the depth of the shower maximum and consequently yields information related to the mass of the primary particle.  All results presented in this simulation study were derived with the geosynchrotron model as implemented in the REAS2 code \\citep{HuegeUlrichEngel2007a}. Additional mechanisms which can contribute to the radio signal, such as Cherenkov-like emission from a charge excess and contributions from the net charge variation during the air shower evolution \\citep{WernerScholten2008,MeyerVernetLecacheuxArdouin2008}, are the subject of current investigations and are not included in the analysis presented here. A direct comparison between the implementations of different emission models by various authors is difficult due to the very different calculational approaches. In contrast to other available models, however, the geosynchrotron model implemented in the REAS2 code allows us an absolute, parameter-free calculation of the emission from individual air showers, accounting for realistic shower-to-shower fluctuations and detailed spatial, angular and energy distributions of the shower particles. If radio emission from air showers is dominated by geomagnetic effects --- an assumption supported by experimental data showing strong correlations between the strength of the radio signal and the orientation of the air shower axis to the Earth's magnetic field \\citep{FalckeNature2005,HornefferArena2005,PetrovicApelAsch2006} --- the results of this analysis can be considered generally applicable.  In the following we first describe the methodology applied throughout the present analysis. We then investigate the case of air showers with 60$^{\\circ}$ zenith angle, where the effects are very prominent and could be well exploited experimentally. Afterwards we demonstrate that the same qualitative behaviour is fulfilled for air showers with 45$^{\\circ}$ zenith angle, where it would, however, be harder to use experimentally, followed by our conclusions. ",
        "conclusions": "Based on Monte Carlo simulations performed with CORSIKA and REAS2, we have shown that geosynchrotron radio emission from air showers of cosmic rays is sensitive to the primary particle energy and mass, and that this information content can be extracted from radio data on a shower-to-shower basis.  For a given observing frequency and zenith angle, there is a well-defined lateral distance region in which the filtered radio peak amplitude is directly proportional to the energy deposited in the atmosphere by the electromagnetic cascade of an air shower, $E_{\\mathrm{em}}$. Once the proportionality constant is known, $E_{\\mathrm{em}}$ can be derived from a radio measurement in this so-called {\\em flat region}. The RMS spread of the $E_{\\mathrm{em}}$ values due to shower-to-shower fluctuations amounts to less than 3\\% in case of an ideal radio signal measurement. Relating $E_{\\mathrm{em}}$ to the energy of the primary particle introduces an additional uncertainty of $\\sim 5$ to 7\\%.  The peak field strength in the {\\em flat region} can be combined with a field strength measured in a {\\em steep region}, found at a larger lateral distance and/or higher observing frequency. The ratio of these field strengths can then be directly related to the atmospheric depth of the air shower maximum, $X_{\\mathrm{max}}$, on a shower-to-shower basis. The RMS spread of the $X_{\\mathrm{max}}$ values around the corresponding fit function due to shower-to-shower fluctuations amounts to $\\sim 15$--$20\\,$g$\\,$cm$^{-2}$, not including any instrumental resolutions.  Exploitation of these signal characteristics with an experimental measurement requires antenna spacings dense enough to allow reconstruction of the electric field strengths in the {\\em flat region}. As the location of the {\\em flat region} shifts to larger lateral distances for increasing zenith angles, application of this strategy is particularly interesting for inclined showers with at least 45$^{\\circ}$, better 60$^{\\circ}$, zenith angle. (At smaller zenith angles, the {\\em flat} region will shift to distances very close to the shower core, and an extrapolation of the electric field strength to the core region will become necessary.) In the case of 60$^{\\circ}$ zenith angle and for an observing frequency of \\unit[32--64]{MHz}, the {\\em flat region} lies at a (ground-coordinate) distance of $\\sim 175\\,$m in the azimuthal direction defined by the shower axis. Decreasing the observing frequency, if technically possible, increases this distance considerably. In addition, observations at lower heights above sea level increase the distance, especially for zenith angles smaller than 60$^{\\circ}$.  Assuming that the geosynchrotron model accounts for the dominant part of radio emission from extensive air showers, exploitation of these characteristics could thus make radio detection a powerful tool for the determination of the primary energy of cosmic rays and the depth of shower maximum of air showers, which in turn is related to the mass of the primary particle.   \\begin{ack} The authors would like to thank T.\\ Pierog and D.\\ Heck for their outstanding support regarding many CORSIKA-related aspects of this work, S.\\ Laf\\`ebre for his efforts in simulating CORSIKA showers on the high-performance computer cluster ``Stella'' and A.\\ Haungs and H.\\ Falcke for very helpful discussions regarding the analysis and the manuscript. This research has been supported by grant number VH-NG-413 of the Helmholtz Association. \\end{ack}"
    },
    "0806/0806.1482_arXiv.txt": {
        "abstract": "We report the discovery of WASP-10b, a new transiting extrasolar planet (ESP) discovered by the WASP Consortium and confirmed using NOT FIES and SOPHIE radial velocity data.  A 3.09 day period, 29 mmag transit depth, and 2.36 hour duration are derived for WASP-10b using WASP and high precision photometric observations. Simultaneous fitting to the photometric and radial velocity data using a Markov-chain Monte Carlo procedure leads to a planet radius of 1.28R$_J$, a mass of 2.96M$_J$ and eccentricity of $\\approx$0.06. WASP-10b is one of the more massive transiting ESPs, and we compare its characteristics to the current sample of transiting ESP, where there is currently little information for masses greater than $\\approx$2M$_J$ and non-zero eccentricities. WASP-10's host star, GSC 2752-00114 (USNO-B1.0 1214-0586164) is among the fainter stars in the WASP sample, with V=12.7 and a spectral type of K5.  This result shows promise for future late-type dwarf star surveys. ",
        "introduction": "Photometric transit observations of extrasolar planets (ESP) are important because the transit strongly constrains their orbital inclination and allows accurate physical parameters for the planet to be derived. Their mass-radius relation allows us to probe their internal structure and is vital to our understanding of orbital migration and planetary formation. The radial velocity measurements which are used to confirm a candidate transiting ESP also provide more complete information  on the orbital eccentricity.  As wide field photometric transit surveys have collected additional sky and temporal coverage, and understood their noise components \\citep{cameron2007mcmc, smith07}, the number of transiting ESP has grown to over 50 in line with earlier predictions \\citep{h1}. Recently one such survey, SuperWASP \\citep{p1} published its first 5 confirmed ESP, all of which have periods of less than 3 days \\citep{A08, c4, P08, W08}, and reported an additional 10\\footnote{http://www.inscience.ch/transits/} \\citep{Hel08, Hebb08, yog08, West08}. SuperWASP is performing a \"shallow-but-wide\" transit search, designed to find planets that are not only sufficiently bright ($9 < V < 13$) for high-precision radial velocity follow-up to be feasible on telescopes of moderate aperture, but also for detailed studies such as transmission spectroscopy during transits. Details of the WASP project and observatory infrastructure are described in \\citet{p1}.    In this paper we present the WASP photometry of 1SWASP~J231558.30+312746.4 (GSC 2752-00114), higher precision photometric follow-up observations with the MERCATOR and Tenagra telescopes, and high precision radial velocity observations with the Nordic Optical Telescope new FIbre-fed Echelle Spectrograph (FIES) and the OHP SOPHIE collaboration. These observations lead to the discovery and confirmation of a new, relatively high mass, gas-giant exoplanet, WASP-10b.  \\begin{figure} \\psfig{file=fig1.eps,width=8.0cm} \\caption{ % (a.) The {\\it top} panel shows the SuperWASP light curve for WASP-10b (1SWASP~J231558.30+312746.4). All the data (apart from that from SuperWASP-N) were averaged in 300\\,second bins. The data were phased using the ephemeris, $T_O = 2454357.85808 ^{+ 0.00041}_{- 0.00036}$ $P = 3.09276$ days. % } \\end{figure} ",
        "conclusions": "Photometric surveys have now provided a large sample of transiting ESP that can be used to determine their mass-radius relation and provide constraints on their compositions. Here we presented the discovery of a new ESP with a mass of 2.96M$_J$, 1.28R$_J$ radius, and a significant eccentricity of $0.059 ^{ + 0.014}_{ -0.004}$. We now discuss  the properties of WASP-10b in relation to the current sample of transiting ESP, starting with its non-zero eccentricity.   Most of the current sample of published transiting ESP have orbits consistent with being circular and are fit with models using zero eccentricity as is expected for short-period planets in orbits with semi-major axes $<$ 0.2 AU.  Recent work \\citep{jack08, mard07} has investigated the effects of tidal dissipation on the orbits of short period ESP. The evolution of the orbital eccentricity appears to be driven primarily by tidal dissipation within the planet, giving a circularisation timescale substantially less than 1 Gyr for typical tidal dissipation parameter, Q$_p$ = 10$^5$ to 10$^6$. WASP-10 is a K dwarf with a spin period of 12 days and J$-$K=0.62 and is rotating more slowly than stars of comparable colour in the Hyades \\citep{tern00}. % This suggests a rotational age between 600 Myr and 1 Gyr. Thus, the persistence of substantial orbital eccentricity in WASP-10b is therefore surprising.  One plausible mechanism for maintaining the high eccentricity is secular interaction with an additional planet in the system. \\citet{AL06} explore the effects of dynamical interactions among planets in extrasolar planetary systems and conclude outer planets can cause the inner planet to move through a range of eccentricities over timescales that are short when compared to the lifetime of the system, but very long when compared to the current observational baseline. However, recently \\citet{mat08} have argued that an unseen companion driving short-Period systems is unlikely. They present an upper limit of 1 M$_{Neptune}$ for a possible unseen companion in the GJ 436 system and exclude this based on the current radial velocity upper limits of $\\leq$ 5 m/s. \\citet{mat08} also  present a range of tidal quality Q$_p$ timescales that could be as large as 10$^9$ years, and argue that this new class of eccentric, short period ESP are simply still in the process of circularizing. WASP-10b has not been extensively studied to rule out a putative outer plane that may be driving its eccentricity. Thus, the $\\approx$6\\% eccentricity of WASP-10b makes it an attractive target for future transit-timing variation studies, and for longer-term RV monitoring to establish the mass and period of the putative outer planet.    The majority of transiting ESP found have masses below 1.5M$_J$, although there are a few more massive ESP. HD 17156, and COROT-Exo-2 have similar masses to WASP-10b and although there are two more massive ESP, the nearly 9 M$_J$ HAT P-2 (HD~149026b) \\citep{bakos07} and 7.3 M$_J$ WASP-14b \\citep{yog08}, this higher mass region has been poorly explored. Additional transiting objects in the mass range are important for completing the current ESP mass-radius relations and constraining their compositions. The current sample of transiting extrasolar giant planets (ESP) reveals a large range of densities.  We derive a mean density for WASP-10b of $\\approx$1.89 g cm$^{-3}$ (1.42 $\\rho_J$) and it would lie along the higher density contour in a mass-radius plot \\citep{P08, sozz07}.  One ultimate goal of our transit-search programme is to provide the observational grist that will stimulate and advance refined models for the formation and evolution of the hot and very hot Jupiters \\citep{burr97, fortney2007, seag07}.  By thus constraining the underlying physics, we will have a richer context for the interpretation of the lower mass planets expected from missions such as COROT and Kepler."
    },
    "0806/0806.4744_arXiv.txt": {
        "abstract": "We calculate the nucleon sigma term in two-flavor lattice QCD utilizing the Feynman-Hellman theorem. Both sea and valence quarks are described by the overlap fermion formulation, which preserves exact chiral and flavor symmetries on the lattice. We analyse the lattice data for the nucleon mass using the analytical formulae derived from the baryon chiral perturbation theory. From the data at valence quark mass set different from sea quark mass, we may extract the sea quark contribution to the sigma term, which corresponds to the strange quark content. We find that the strange quark content is much smaller than the previous lattice calculations and phenomenological estimates. ",
        "introduction": "A piece of information on the nucleon structure can be extracted from its quark mass dependence. Nucleon sigma term $\\sigma_{\\pi N}$ characterizes the effect of finite quark mass on the nucleon mass. Up to non-analytic and higher order terms, the nucleon mass is written as $M_N=M_0+\\sigma_{\\pi N}$, where $M_0$ is the nucleon mass in the chiral limit. The exact definition of $\\sigma_{\\pi N}$ is given by the form of a scalar form factor of the nucleon at zero recoil as \\begin{equation} \\label{eq:piNsigma} \\sigma_{\\pi N}= m_{ud} \\left( \\langle N | \\bar{u}u+\\bar{d}d | N \\rangle -V\\langle 0 | \\bar{u}u+\\bar{d}d | 0 \\rangle \\right) \\end{equation} where $m_{ud}$ denotes degenerate up and down quark mass. The second term in the parenthesis represents a subtraction of the vacuum contribution, and $V$ is the (three-dimensional) physical volume \\footnote{ The nucleon state $|N(p)\\rangle$ is normalized as $\\langle N (p)|N (p^\\prime)\\rangle= (2\\pi)^3 \\delta^{(3)}(p-p^\\prime)$. In (\\ref{eq:piNsigma}) we omit the momentum argument for the nucleon, since we do not consider finite momentum insertion in this paper. }. For the sake of simplicity we represent the vacuum subtracted matrix element $\\langle N|\\bar{q}q|N\\rangle-V\\langle 0|\\bar{q}q|0\\rangle$ by $\\langle N|\\bar{q}q|N\\rangle$ in what follows. ($q$ represents a quark field: up ($u$), down ($d$), or strange ($s$).) Note that the sigma term is renormalization group invariant, since the renormalization factor cancels between the quark mass $m_q$ and the scalar operator $\\bar{q}q$.  While the up and down quarks contribute to $\\sigma_{\\pi N}$ both as valence and sea quarks, the strange quark appears only as a sea quark contribution. As a measure of the strange quark content of the nucleon, the $y$ parameter \\begin{equation} \\label{eq:y} y \\equiv \\frac{2\\langle N | \\bar{s}s  | N \\rangle} {\\langle N | \\bar{u}u +\\bar{d}d | N \\rangle}, \\end{equation} is commonly introduced. Besides characterizing the purely sea quark content of the nucleon, which implies a clear distinction from the quark model picture of hadrons, this parameter plays an important role to determine the detection rate of possible neutralino dark matter in the supersymmetric extension of the Standard Model \\cite{Griest:1988yr,Griest:1988ma,Bottino:1999ei,% Ellis:2003cw,Ellis:2005mb,Baltz:2006fm,Ellis:2008hf}. Already with the present direct dark matter search experiments one may probe a part of the MSSM model parameter space, and new experiments such as XMASS and SuperCDM will be able to improve the sensitivity by 2--3 orders of magnitude. Therefore, a precise calculation of the $y$ parameter (or equivalently another parameter $f_{T_s}\\equiv m_s\\langle N|\\bar{s}s|N \\rangle/M_N$) will be important for excluding or proving the neutralino dark matter scenario.  Phenomenologically, the sigma term can be related to the $\\pi N$ scattering amplitude at a certain kinematical point, {\\it i.e.} the so-called Cheng-Dashen point $t=+2m_\\pi^2$ \\cite{Cheng:1970mx}. Its value is in the range $\\Sigma_{CD}=70\\sim 90$~MeV \\cite{Pavan:2001wz}. After the corrections for the finite value of $t$, which amounts to $-15$~MeV \\cite{Gasser:1990ce}, one obtains $\\sigma_{\\pi N}=55\\sim 75$~MeV. On the other hand, the octet breaking of the nucleon mass, or the matrix element $\\langle N|\\bar{u}u+\\bar{d}d-2\\bar{s}s|N\\rangle$, can be evaluated from the baryon mass spectrum. At the leading order of Chiral Perturbation Theory (ChPT), the value of the corresponding sigma term is $\\hat{\\sigma}\\simeq 26$~MeV, while the heavy baryon ChPT (BChPT) gives $\\hat{\\sigma}=36\\pm 7$~MeV~\\cite{Borasoy:1996bx}. The difference between $\\sigma_{\\pi N}$ and $\\hat{\\sigma}$ is understood as the strange quark contribution; algebraically the relation is $\\sigma_{\\pi N} = \\hat{\\sigma}/(1-y)$. Then, one obtains a large value of $y$: y=0.3--0.6. (The value of $y$ is even larger than the estimate $y\\simeq 0.2$ in \\cite{Gasser:1990ce}, because of the more recent experimental data \\cite{Pavan:2001wz}.) For other phenomenological estimates, see, {\\it e.g.} \\cite{Bernard:2007zu}. Such large values of $y$ cannot be understood within the valence quark picture, hence raises a serious problem in the understanding of the nucleon structure. We note however that the analysis within chiral effective theories suffers from significant uncertainties of the low energy constants, especially at higher orders.  Using lattice QCD, one can in principle calculate the nucleon sigma term without involving any model parameters, since lattice calculation for a wide range of quark masses in the chiral regime offers essential information on the low energy constants which cannot be determined by experimental data alone. Furthermore, it is possible to determine the valence and sea quark contributions separately. A direct method to extract them is to calculate three-point functions on the lattice including an insertion of the scalar operator. It can also be done in an indirect way by analyzing the quark mass dependence of the nucleon mass for valence and sea quarks separately. Obviously, the dynamical fermion simulations are necessary to extract the disconnected contributions in the indirect method.  Previous lattice results were $\\sigma_{\\pi N}$ = 40--60~MeV, $y$ = 0.66(15) \\cite{Fukugita:1995ba}, and $\\sigma_{\\pi N}$ = 50(3)~MeV, $y$ = 0.36(3) \\cite{Dong:1996ec} within the quenched approximation, while a two-flavor QCD calculation \\cite{Gusken:1998wy} gave $\\sigma_{\\pi N}$ = 18(5)~MeV and $y$ = 0.59(13). There are apparent puzzles in these results: firstly the strange quark content due to the disconnected diagram (the value of $y$) is unnaturally large compared to the up and down contributions that contain the connected diagrams too. Secondly the values of the sigma term in the quenched and unquenched calculations are rather different, which might also imply significant effects of quark loops in the sea.  Concerning the first point, it was pointed out that using the Wilson-type fermions, which violate the chiral symmetry on the lattice, the sea quark mass dependence of the additive mass renormalization and lattice spacing can give rise to a significant uncertainty in the sea quark content \\cite{Michael:2001bv}. Unfortunately, after subtracting this contamination the unquenched result has large statistical error, $y=-0.28(33)$. In the present work, we remove this problem by explicitly maintaining exact chiral symmetry on the lattice for both sea and valence sectors, as described below.  The second puzzle may be resolved by incorporating an enhancement due to pion loops. Within BChPT at $O(p^3)$ or $O(p^4)$, a curvature is expected in the quark mass dependence of the nucleon, hence the sigma term, a derivative of $M_N$ in terms of $m_q$, increases towards the chiral limit.  An analysis using existing lattice data of two-flavor QCD with $m_\\pi>550$~MeV by CP-PACS \\cite{Ali Khan:2001tx}, JLQCD \\cite{Aoki:2002uc}, and QCDSF \\cite{Ali Khan:2003cu} yields $\\sigma_{\\pi N}=48\\pm 5 ^{+\\ 9}_{-12}$~MeV \\cite{Procura:2003ig,Procura:2006bj}, which is slightly smaller than but is still consistent with the phenomenological analysis. A more recent lattice data by the ETM Collaboration with $m_\\pi$ = 300--500~MeV in two-flavor QCD reported a higher value $\\sigma_{\\pi N}$ = 67(8)~MeV \\cite{Alexandrou:2008tn}. Such an analysis for the disconnected contribution to extract the strange content is yet to be done, which is another main point of this work.  In this work, we analyze the data of the nucleon mass obtained from a two-flavor QCD simulation employing the overlap fermion \\cite{Aoki:2008tq}. (For other physics results from this simulation, we refer \\cite{Matsufuru:2007uc} and references therein.) The overlap fermion \\cite{Neuberger:1997fp,Neuberger:1998wv} preserves exact chiral symmetry on the lattice, and there is no problem of the additive mass shift of the scalar density operator, that was a main source of the large systematic error in the previous calculations of the sigma term. We use the overlap fermion to describe both the sea and valence quarks. Statistically independent ensembles of gauge configurations are generated at six different sea quark masses; the nucleon mass is measured for various valence quark masses on each of those gauge ensembles. Therefore, we are able to analyze the valence and sea quark mass dependence independently to extract the connected and disconnected contributions. An estimate of the strange quark content can thus be obtained in two-flavor QCD. In the analysis, we use the partially quenched BChPT, which corresponds to the lattice calculations with valence quark masses taken differently from the sea quark masses. Therefore, the enhancement of the sigma term towards the chiral limit is incorporated. Since the two-flavor QCD calculation cannot avoid the systematic error due to the neglected strange sea quarks, our result should not be taken as a final result from lattice QCD. Nevertheless our study with exact chiral symmetry reveals the underlying systematic effects in the calculation of the nucleon sigma term, especially in the extraction of its disconnected contribution. It therefore provides a realistic test case, which will be followed by the 2+1-flavor calculations in the near future \\footnote{ For a very recent result from 2+1-flavor QCD, see \\cite{WalkerLoud:2008bp}. }.  Our paper is organized as follows. In Section \\ref{sec:Method}, we introduce the basic methods to calculate the nucleon sigma term. Our simulation set-up is described in Section~\\ref{sec:Simulation}. Then, in Section~\\ref{sec:BChPT}, we describe the BChPT fit to obtain the sigma term. In Section~\\ref{sec:PQChPT}, we study the sea quark content of the nucleon from PQChPT. In Section~\\ref{sec:Discussion}, we compare our results with previous calculations and discuss the origin of the discrepancy. Our conclusion is given in Section~\\ref{sec:Summary}. ",
        "conclusions": "\\label{sec:Discussion}  We found that the disconnected contribution to the sigma term is much smaller than the previous lattice calculations with the Wilson-type fermions $y\\simeq 0.36\\sim 0.66$ \\cite{Fukugita:1995ba,Dong:1996ec,Gusken:1998wy} (except for \\cite{Michael:2001bv} as explained below). The authors of \\cite{Michael:2001bv} found that the naive calculation with the Wilson-type fermions may over-estimate the sea quark contents due to the additive mass shift and the sea quark mass dependence of the lattice spacing. The key observation is that the additive mass shift is large depending significantly on the sea quark mass. Therefore, in order to obtain the derivative (\\ref{eq:deriv_sea}) one must subtract the unphysical contribution from the additive mass shift. This problem remains implicitly in the quenched calculations, since the derivative must be evaluated at the value of the valence quark mass even when the sea quark mass is sent to infinity. (There is of course the more fundamental problem in the quenched calculations due to the missing sea quark effects.)  Another problem is in the conventional scheme of setting the lattice scale in unquenched simulations. In many dynamical fermion simulations, the lattice spacing is set (typically using the Sommer scale $r_0$) at each sea quark mass, or in some cases, the bare lattice coupling $\\beta$ is tuned to yield a given value of $r_0$ independent of the sea quark mass. This procedure defines a renormalization scheme that is mass dependent, because the quantity $r_0$ could have physical sea quark mass dependence. Since the partial derivative (\\ref{eq:deriv_sea}) is defined in a mass independent scheme, {\\it i.e.} the coupling constant does not depend on the sea quark mass, one has to correct for the artificial sea quark mass dependence through $r_0$ when one calculates the nucleon sigma term. Combining these two effects, the authors of \\cite{Michael:2001bv} found that their unsubtracted result $y=0.53(12)$ is substantially reduced and becomes consistent with zero: $y=-0.28(33)$. The conclusion of this analysis is that the previous lattice calculations giving the large values of $y$ suffered from the large systematic effect, hence should not be taken at their face values.  Our calculation using the overlap fermion is free from these artifacts. The additive mass shift is absent because of the exact chiral symmetry of the overlap fermion. The lattice spacing is kept fixed in our analysis at a fixed bare lattice coupling constant. We confirmed that this choice gives a constant value of the renormalized coupling constant in the (mass independent) $\\overline{\\mathrm{MS}}$ scheme through an analysis of current-current correlators \\cite{Shintani:2008qe}. Therefore, the small value of $y$ obtained in our analysis (\\ref{eq:y_result}) provides a much more reliable estimate than the previous lattice calculations.   We study the nucleon sigma term in two-flavor QCD simulation on the lattice with exact chiral symmetry. Fitting the quark mass dependence of the nucleon mass using the formulae from Baryon Chiral Perturbation Theory (BChPT), we obtain $\\sigma_{\\pi N} = 53(2)(^{+21}_{-\\ 7})$~MeV, where our estimates of systematic errors are added in quadrature. This is consistent with the canonical value in the previous phenomenological analysis. Owing to the exact chiral symmetry, our lattice calculation is free from the large lattice artifacts coming from the additive mass shift present in the Wilson-type fermion formulations.  We also estimate the strange quark content of nucleon. From an analysis of partially quenched lattice data, we find that the sea quark content of the nucleon is less than 0.4 for the entire quark mass region in our study. The valence quark content is in fact the dominant contribution to the sigma term. Taking account of the enhancement of $\\langle N|(\\bar{u}u+\\bar{d}d)|N\\rangle$ near the chiral limit, the parameter $y$ is most likely less than 0.05 in contrast to the previous lattice calculations.  By directly calculating the disconnected diagram we may obtain further information. For instance, the effect of the strange quark loop on the dynamical configurations with light up and down quarks can be extracted. Such a calculation is in progress using the all-to-all quark propagators on the lattice. Another obvious extension of this work is the calculation including the strange quark loop in the vacuum. Simulations with two light and one strange dynamical overlap quarks are on-going \\cite{Hashimoto:2007vv}."
    },
    "0806/0806.2354_arXiv.txt": {
        "abstract": "Within the framework of Newtonian magneto-solid-mechanics, relying on equations appropriate for a perfectly conducting elastic continuous medium threaded by a uniform magnetic field, the asteroseismic model of a neutron star undergoing axisymmetric global torsional nodeless vibrations under the combined action of Hooke's elastic and Lorentz magnetic forces is considered with emphasis on a toroidal Alfv\\'en mode of differentially rotational vibrations about the dipole magnetic moment axis of the star. The obtained spectral equation for frequency is applied to $\\ell$-pole identification of quasi-periodic oscillations (QPOs) of X-ray flux during the giant flares of SGR 1806-20 and SGR 1900+14. Our calculations suggest that detected QPOs can be consistently interpreted, within the framework of this model, as produced by global torsional nodeless vibrations of quaking magnetar if they are considered to be restored by the joint action of bulk forces of shear elastic and magnetic field stresses. ",
        "introduction": "The recent detection of quasi-periodic oscillations (QPOs) on the light-curve tails of X-ray flaring SGR 1806-20 and SGR 1900+14 (Israel et al 2005; Watts \\& Strohmayer 2006) that have been ascribed to torsional seismic vibrations of quaking magnetar raises several questions of quite general interest for the asteroseismology of degenerate solid stars. Namely, whether these differentially rotational oscillations are predominately of an elastic nature, that is, restored by Hooke's force of shear solid-mechanical stresses or they should be thought of as the toroidal  magnetic Alfv\\'en mode of axisymmetric vibrations about dipole magnetic moment axis of the star and restored by the Lorentz force of magnetic field stresses (Glampedakis, Samuelsson \\& Andersson 2006). Also, it remains questionable whether these vibrations are of global character, that is, excited in the entire volume or they can be explained as locked in the peripheral finite-depth crustal region of a neutron star. These and related issues are currently the subject of intense theoretical investigations by classical theory methods of material continua  (Piro 2005; Levin 2007, Lee 2008; Bastrukov et al, 2007a; 2008a) and by general relativity  methods (Sotani, Kokkotas \\& Stergioulas 2008; Samuelssen \\& Andersson 2007).  In this article, continuing the above investigation, we study an asteroseismic model of a neutron star with uniform internal and dipolar external magnetic field undergoing quake-induced global differentially rotational, torsional, vibrations about the dipole magnetic moment axis of the star under the joint action of the elastic Hooke's and the magnetic Lorentz forces. In so doing we confine our attention to the regime of extremely long wavelength differentially rotational fluctuations of material displacements about motionless stationary state which are insensitive to core-crust compositional stratification of quaking neutron star. The characteristic feature of this regime is that conducting and highly robust to compressional distortions of solid-state material, of both core and crust, coupled by Maxwell stresses of uniform fossil magnetic field frozen in the star, sets in coherent axisymmetric differentially rotational vibrations with the nodeless toroidal field of material displacements identical to that for global torsional vibrations of a spherical mass of an elastic continuous medium capable of transmitting perturbation by transverse shear elastic waves generic to solid state of condensed matter, not a liquid one\\footnote{It is commonly agreed that the average speed of transverse wave of elastic shear in the crustal matter is $c_t=\\sqrt{\\mu/\\rho}\\sim 10^8$ cm s$^{-1}$ which holds for both the outer crust with the density $10^{8}< \\rho < 10^{11}$ g cm$^{-3}$ and the inner crust, whose density is ranged between the neutron drip line density  $\\rho \\sim 10^{11}$ g cm$^{-3}$ and density of overwhelming neutronization $\\rho\\sim 10^{13}$ g cm$^{-3}$ (e.g., Blaes et al 1990). Accordingly, the shear modulus of the crustal matter is ranged from the surface to the crust-core interface as $10^{24}< \\mu < 10^{29}$ dyn cm$^{-2}$. The much denser core matter compressed by self-gravity to nearly normal nuclear density  $\\rho=2.8\\cdot 10^{14}$ g cm$^{-3}$ is much harder to elastic shear distortions and its shear modulus is estimated as $\\mu \\sim 10^{33}$ dyn cm$^{-2}$ (e.g., Owen 2005, Bastrukov et al 2008b and references therein); in view of this the Fermi-degenerate neutron matter of atomic nuclei and neutron star cores is designated as an elastic Fermi-solid capable of transmitting shear elastic wave with the speed $c_t=\\sqrt{\\mu/\\rho}\\sim 0.2\\,c$ where $c$ being the speed of light.}. In the context of the above QPOs problem, such vibrations have been analyzed in some details in our recent works (Bastrukov et al 2007a, 2007b, 2008a), but the driving force was assumed to be of only elastic nature, that is, owing its origin to fluctuations in shear elastic stresses. The focus of this paper is placed, therefore, on the toroidal Alfv\\'en vibrational mode in which magnetic field and field of material displacement undergo coupled fluctuations restored by Lorentz force with  Amp\\'er's form of the conduction current density. Before proceeding to the details of calculations it worth noting that the liquid star model with axisymmetric poloidal uniform magnetic field inside and dipolar outside has been the subject of serious investigations in the past in the context of magnetic variables (e.g., Schwarzschild 1949; Chandrasekhar \\& Fermi 1953; Ledoux \\& Walraven 1958). The situation may be quite different for ultra strong internal magnetic fields frozen-in the super dense matter of the end products of stellar evolution, white dwarfs and neutron stars. In neutron stars self-gravity is counterbalanced by the degeneracy pressure of relativistic electrons in the crust, whose highly conducting matter is capable of sustaining persistent current-carrying flows and by degeneracy pressure of non-relativistic neutrons in the cores of neutron stars whose material can be in the state of paramagnetic magnetization caused by Pauli's mechanism of field-induced alignment of spin magnetic moment of neutrons along the frozen-in the star fossil magnetic field (e.g., Bastrukov et al; 2002a; 2002b). Of course, there is no compelling evidence that this is the case and there is no certainly general agreement how the magnetic fields of pulsars and magnetars are produced. Over the years a model of a neutron star with uniform internal and dipolar external field, pictured in Fig.1, has been invoked to discussions of evolution of magnetic dipole fields of isolated radio pulsars (Flowers \\& Ruderman 1977), whose surface dipolar magnetic fields are found to be highly stable to spontaneous decay (Bhattacharya \\& van den Heuvel 1991, Chanmugam 1994), as well as of origin and evolution of ultra strong magnetic fields of magnetars (Braithwaite \\& Spruit 2006; Geppert \\& Rheinhardt 2006; Spruit 2008), therefore the study of axisymmetric torsional vibrations within the framework of a neutron star model with uniform internal magnetic field seems amply justified. This interesting in its own right model describes situation when physically meaningful analytical solution of the eigenfrequency problem can be found and discloses mathematical difficulties one must confront when computing the frequency of Alfv\\'en vibrational modes in the neutron star models with non-homogeneous internal and dipolar external magnetic fields. An analytic example of such latter field is discussed in (Roberts 1981; Geppert \\& Rheinhardt 2006).  The paper is organized as follows. In Sec.2, the exact solution is given for the eigenfrequency problem of global torsional nodeless vibrations of a solid star with a uniform field inside and dipole outside driven  solely by Lorentz restoring force. The obtained frequency spectrum of toroidal and poloidal magnetic Alfv\\'en vibrational modes in such a star is then quantified numerically with input parameters, such as  mass, radius and magnetic field, characteristic to pulsars and magnetars. In Sec.3, the frequency spectrum of axisymmetric global torsional nodeless vibrations under the action of the combined forces of shear elastic and magnetic field stresses is obtained and applied to $\\ell$-pole identification of the QPOs frequencies in flaring SGR 1806-20 and SGR 1900+14. The newly obtained results are briefly summarized in Sec.4. ",
        "conclusions": "There is a common belief today that gross features of the asteroseimology of pulsars and magnetars can be understood on the basis of a solid star model presuming that quake-induced shear vibrations restored by bulk forces of intrinsic stresses of different physical nature are governed by described by solid mechanics or elastodynamics  (e.g., Hansen \\& van Horn 1979; McDermott, van Horn \\& Hansen 1988; Bastrukov, Weber \\& Podgainy 1999; Bastrukov et al 2007a). This point of view is quite different from theoretical approach to the asteroseismology of the main sequence stars at the base of which lies the liquid star model whose vibrations are treated within the framework of fluid-mechanical theory of continuous media, as is the case of helioseismology (e.g., Christensen-Dalsgaard 2002).  The main purpose of this work was  to examine the magneto-solid-mechanical variational method of the asteroseismology of neutron star by probing its interior with non-radial global differentially rotational, torsional, vibrations with nodeless toroidal field of material displacements, which are insensitive to compositional stratification of the star matter.  Bearing in mind that external magnetic fields of pulsars and magnetars are commonly thought of as produced by magnetic dipole moment of underlying neutron star we have considered a model of a neutron star with perhaps simplest, from the viewpoint of computational feasibility, imaginable configuration of the magnetostatic fossil magnetic field, pictured in Fig.1. Proceeding from this admittedly idealized model, the two-parametric spectral equation for the frequency of global torsional vibrations has been derived in analytic form showing the larger multipole degree $\\ell$ of torsional nodeless vibrations the higher is the frequency $\\omega(_0t_\\ell)$. The application of the obtained two-parametric spectral formula to modal analysis of QPOs during the flare of SGR 1806-20 and SGR 1900+14 shows that data on the QPOs frequencies with $\\ell$ from interval $2\\leq \\ell\\leq 20$ can be consistently interpreted as produced by global torsional nodeless vibrations restored by combined forces of shear elastic and magnetic field stresses. This inference is, of course, suggestive rather than conclusive, in view of adopted a highly idealized configuration of fossil magnetostatic field frozen in the star, and too much remains to be done to be at all certain of interpretations suggested for QPOs in the X-ray flux during the giant flares of the above magnetars."
    },
    "0806/0806.2162_arXiv.txt": {
        "abstract": "We report here the detection of a weak magnetic field of 50--100~G on the O9.7 supergiant $\\zeta$~Ori A, using spectropolarimetric observations obtained with NARVAL at the 2m T\\'elescope Bernard Lyot atop Pic du Midi (France). \\zetaori\\ is the third O star known to host a magnetic field (along with \\tori\\ and HD~191612), and the first detection on a 'normal' rapidly-rotating O star. The magnetic field of \\zetaori\\ is the weakest magnetic field ever detected on a massive star. The measured field is lower than the thermal equipartition limit (about 100~G). By fitting NLTE model atmospheres to our spectra, we determined that \\zetaori\\ is a 40~\\msun\\ star with a radius of 25~\\rsun and an age of about 5--6~Myr, showing no surface nitrogen enhancement and losing mass at a rate of about 2$\\times10^{-6}$~\\mspy.  The magnetic topology of \\zetaori\\ is apparently more complex than a dipole and involves two main magnetic polarities located on both sides of the same hemisphere; our data also suggest that \\zetaori\\ rotates in about 7.0~d and is about 40\\degr\\ away from pole-on to an Earth-based observer.  Despite its weakness, the detected magnetic field significantly affects the wind structure;  the corresponding Alfv\\'en radius is however very close to the surface, thus generating a different rotational modulation in wind lines than that reported on the two other known magnetic O stars.  The rapid rotation of \\zetaori\\ with respect to \\tori\\ appears as a surprise, both stars having similar unsigned magnetic fluxes (once rescaled to the same radius); it may suggest that the sub-equipartition field detected on \\zetaori\\ is not a fossil remnant (as opposed to that of \\tori\\ and HD~191612), but the result of an exotic dynamo action produced through MHD instabilities. ",
        "introduction": "Stellar magnetic fields have been detected across a large range of spectral types. In solar-type and essentially all cool, low-mass (i.e., mid F and later) stars, magnetic fields (and activity) are observed, often featuring a complex topology, and are thought to be due to dynamo processes occuring within the outer convective layers. In hotter, more massive stars with outer radiative zones, magnetic fields are also detected (with a significantly simpler topology though) but only in a small fraction of stars (e.g., the magnetic chemically peculiar stars among the A and late B stars).  The situation might be similar (though less well studied) among early B and O stars, with only two O stars (namely \\tori\\ and  HD~191612, \\citealt{Donati02, Donati06a}) and less than a handfull of early B-type stars (e.g., $\\tau$~Sco, $\\beta$~Cep, $\\zeta$~Cas, \\citealt{Donati06b, Donati01, Neiner03a}) yet known as magnetic.  Magnetic fields are nonetheless expected to play a significant role throughout the evolution of hot massive stars, by modifying the internal rotation, enhancing chemical transport and mixing, and producing enhanced surface abundances \\citep{Maeder03, Maeder04, Maeder05}.  Magnetic fields can also dramatically influence the way winds are launched \\citep[e.g.,][]{udDoula02} and the later phases of evolution \\citep[e.g., the collapse, ][]{Heger05};  a large number of observational phenomena (e.g., non-thermal radio emission, anomalous X-ray spectra, abundance anomalies, and H$\\alpha$ modulation) can also be explained (qualitatively at least) by the existence of a weak magnetic field. Yet, the origin of magnetism in massive stars is still an open question, with a lively debate between two classes of models. While some models assert that dynamo processes (either located in the convective core, e.g., \\citealt{Charbonneau01}, or acting within the radiative zone, e.g., \\citealt{Mullan05}) can produce the observed magnetic fields, some others claim that the field is fossil in nature \\citep{Ferrario05, Ferrario06}, being advected and amplified through the initial protostellar collapse.  The limited knowledge that we have about the existence and statistical properties of magnetic fields in massive O stars is mostly due to the fact that these fields are difficult to detect.  Absorption lines of O stars are both relatively few in number in the optical domain, and generally rather broad (because of rotation or to some other type of as yet unknown macroscopic mechanism, e.g., \\citealt{Howarth97}), decreasing dramatically the size of the Zeeman signatures that their putative fields can induce. The results obtained so far (on two stars only) suggest that magnetic O-type stars may be (i) slow rotators and (ii) may exhibit a peculiar spectrum with very regular temporal modulation. While this view may partly reflect an observational bias (magnetic detections being easier on slow rotators) or a selection effect (observations often concentrating on peculiar stars first), null results recently reported on intermediate and fast rotators argue that this effect may be real.  This question is nevertheless a key point for clarifying both the origin and evolutionary impact of magnetic fields in massive stars and therefore deserves being studied with great care.  With the advent of the new generation spectropolarimeters, such as ESPaDOnS at the Canada-France-Hawaii Telescope (CFHT) in Hawaii and NARVAL on the T\\'elescope Bernard Lyot (TBL) in southern France, studies of stellar magnetic fields have undergone a big surge of activity;  in particular, detecting magnetic fields of massive O stars (or providing upper limits of no more than a few tens of G) is now within reach. In this context, we recently initiated a search for magnetic fields in a limited number of 'normal' O stars, using NARVAL.  One of our targets is \\zetaori, a  O9.7~Ib supergiant \\citep{ma04} and the brightest O star at optical wavelengths. Evidence for azimuthal wind structuration (with an modulation timescale of about 6~d, compatible with the rotation period) is reported from both UV and optical lines \\citep[e.g.,][]{Kaper96, Kaper99} and possibly due to the presence of a weak magnetic field. \\zetaori\\ is also well-known for its prominent X-ray emission, $\\logLX=-6.74$  \\citep{berghofer97}.  The origin of this X-ray emission is however still controversial;  while \\citet{Cohen06} suggest that it is due to the classical wind-shock mechanism (with X-rays originating from cooling shocks in the acceleration zone), \\citet{Raassen08} invoke a collisional ionization equilibrium model and \\citet{Pollock07} argue for collisionless shocks controlled by magnetic fields in the wind terminal velocity regime.  For all these reasons, \\zetaori\\ is an obvious candidate for our magnetic exploration program.  In this paper, we report our spectropolarimetric observations of \\zetaori\\ and present the Zeeman detections we obtained (Sec.~\\ref{sec:obs}).  From the collected spectra, we reexamine the fundamental parameters of \\zetaori\\ and discuss the observed rotational modulation to attempt pinning down the rotation period (Sec.~\\ref{sec:spec}). We then carry out a complete modeling of the detected Zeeman signatures and describe the reconstructed magnetic topology (Sec.~\\ref{sec:magt}).  We finally summarise our results, discuss their implications for our understanding of massive magnetic stars, and suggest new observations to confirm and expand our conclusions (Sec.~\\ref{sec:disc}). ",
        "conclusions": "\\label{sec:disc}  We report in this paper the detection and the first modeling attempt of the weak large-scale magnetic field of the O9.7 supergiant \\zetaori.  We detect a field that corresponds to local surface magnetic fluxes of only a few tens of G. The field is everywhere lower than 100~G, making it (by far) the weakest magnetic field ever reported in a hot massive star \\citep{Donati90, auriere07}.  In particular, this magnetic field is weaker than the thermal equipartition limit, equal to about 100~G for \\zetaori;  this is the first sub-equipartition field unambiguously detected in a hot star. The magnetic chemically peculiar stars all show fields larger than their thermal equipartition limit \\citep{auriere07}.  This detection also brings the number of known magnetic O stars to three, with \\zetaori\\ thus joining \\tori\\ \\citep{Donati02} and HD~191612 \\citep{Donati06a}. This is also the first magnetic detection in a 'normal' rapidly rotating O star.  The detailed spectral modeling of \\zetaori\\ provides $\\teff=29,500\\pm1,000$~K and $\\logg=3.25\\pm0.10$ with normal abundances.  It follows that \\zetaori\\ is a 40~\\msun\\ star with a radius equal to about 25~\\rsun, seen from the Earth at an inclination angle of 40\\degr.  With an age of about 5--6~Myr, \\zetaori\\ essentially appears as an evolved counterpart of both \\tori\\ and HD~191612.  Given its evolutionary stage, \\zetaori\\ is expected to show significant N enrichment at its surface (as well as moderate C and O depletion); the normal nitrogen abundance that we measure is thus surprising.  It is tempting to suggest that magnetic fields may play a role in this process; this is however not what the first evolutionary models including magnetic field predict \\citep{Maeder05}.  More work (both on the observational and theoretical side) are required to investigate this issue further.  From a fit to H$\\alpha$, we estimate that the mass-loss rate is about 1.4--1.9$\\times10^{-6}$~\\mspy. From the temporal variability of spectral lines and the modulation of Zeeman signatures, we find that the period of \\zetaori\\ is about 7d. This is compatible with the \\vsini\\ that we measure (from the Fourier shape of the photospheric C~{\\sc iv} lines) and the radius that we derive (from the spectral synthesis), provided the star is view at intermediate inclinations ($i=40$ \\degr).  Given that \\zetaori\\ is typically 3 and 1.4 times larger in size than \\tori\\ and HD~191612 respectively, we find that its overall unsigned magnetic flux (i.e., the integral of the absolute value of the magnetic field over stellar surface) is slightly larger (by a factor of about 1.5) than that of \\tori\\ but much smaller than that of HD~191612 (by about an order of magnitude).  The extremely long rotation period of HD~191612 (about 538~d) suggests that the magnetic field is likely responsible for having dissipated (through confined mass loss) most of the angular momentum of HD~191612 \\citep{Donati06a}.  The slow rotation rate and extreme youth of \\tori\\ also suggests that primordial magnetic fields pervading the parent molecular cloud must have a strong impact onto the angular momentum dissipation throughout the cloud collapse, in qualitative agreement with what numerical simulations predict \\citep[e.g.,][]{Hennebelle08a}.  In this context, one would expect \\zetaori\\ to rotate, if not as slowly as HD~191612 (whose intrinsic magnetic flux is much higher), at least more slowly than \\tori\\ (whose intrinsic magnetic flux is similar) given its later evolution stage;  this is however not what we observe.  No more than speculations can be proposed at this stage. One possibility is that the magnetic field of \\zetaori\\ is not of fossil origin (as opposed to that of \\tori\\ and HD~191612) but rather dynamo generated, making the rotational evolution of \\zetaori\\ and \\tori\\ hardly comparable. The detected magnetic field is indeed much weaker than the critical limit above which MHD instabilities are inhibited (about six times the equipartition field or 600~G in the case of \\zetaori, \\citealt{auriere07}) and may thus result from exotic dynamo action;  the non-dipolar nature of the detected field could be additional evidence in favour of this interpretation, fossil fields being expected to have very simple topologies in evolved stars. Additional spectropolarimetric observations of \\zetaori\\ at different epochs (searching for potential variability of the large-scale field) and of similar 'normal' rapidly-rotating stars are obviously necessary to explore this issue in more details.  Computing the wind magnetic confinement parameter $\\eta_{\\ast}$ of \\citet{udDoula02} for \\zetaori\\ and taking $B\\simeq30-50$~G (at the magnetic equator), $R=25$~\\rsun, $\\Mdot=2\\times10^{-6}$~\\mspy\\ and $\\vinf=2,100$~\\kms\\ (see Secs.~\\ref{sec:spec} and \\ref{sec:magt}) yields $\\eta_{\\ast}\\simeq0.03-0.07$.  The magnetic field of \\zetaori\\ is therefore just strong enough (according to theoretical predictions) to start distorting the wind significantly \\citep{udDoula02}.  The observed rotational modulation in H$\\alpha$, H$\\beta$ and the C~{\\sc iii} lines confirms this first conclusion; the variation in mass loss rate that we measure, corresponding to a density contrast of $\\simeq$1.4, is compatible with what numerical simulations of magnetically confined winds predict \\citep[see, e.g., Fig~8 of][]{udDoula02}.  Note also that the observed line blueshift and asymmetries (cf. Sec.~\\ref{sec:rot} and Fig.~\\ref{fig:bal} for an illustration on the case of He~{\\sc i} 492~nm) are (rotation) phase dependant. The observed blueshift of most lines is maximum when the magnetic poles (i.e. the open field lines) cross the line of sight (at phase 0.8 and 0.45, see Fig.~\\ref{fig:bal}); more data (collected in particular over a longer baseline and densely sampling the rotation cycle) are of course needed to confirm this and to specify how exactly the line shifts and shape are varying with rotation phase (e.g. with 2 unevenly spaced maxima in the line blueshift per rotation period). It however suggests already ${\\it (i)}$ that the line profile variations reflect the varying conditions in which the wind form at the surface of \\zetaori\\ (as a result of the varying local field topology over the star) and ${\\it (ii)}$ that these variations can potentially be used to trace the density at the base of the wind and its variations with the local magnetic topology over the surface of the star.  On both \\tori\\ and HD~191612, H$\\alpha$ is rotationally modulated as a result of the magnetic obliquity (with respect to the rotation axis), with maximum emission occuring when the magnetic pole comes closest to the observer.  Similarly, maximum absorption in UV lines (with highest blueshifted velocities) are observed when the magnetic equator is crossing the line of sight. Extrapolating these results to \\zetaori, we would have first expected H$\\alpha$ in \\zetaori\\ to show maximum emission twice per rotation cycle, at phases of about 0.40 and 0.90 (see Fig.~\\ref{fig:map}), in contradiction with what we observe;  while Balmer emission indeed peaks at phase 0.95, phase 0.51 rather corresponds to minimum (rather than maximum) emission (see Fig.~\\ref{fig:bal}).  The analogy with \\tori\\ and HD~191612 can obviously not be directly applied to \\zetaori. Given the much weaker wind magnetic confinement parameter of \\zetaori\\ (roughly equal to 10 for both \\tori\\ and HD~191612), this is not altogether very surprising.  In particular, the Alfven radius is much closer to the surface of the star in \\zetaori, probably not further than 0.05--0.1~\\rstar\\ above the surface\\footnote{The corotation radius, i.e., the radius at which the Keplerian period equals the rotation period at the surface of the star, is equal to about 2~\\rstar\\ in \\zetaori, i.e., 1~\\rstar\\ above the surface of the star.} (as opposed to 1~\\rstar\\ above the surface for \\tori, \\citealt{Donati02}).  In the magnetically confined wind-shock model \\citep{bm97, Donati02}, the rotational modulation of H$\\alpha$ emission can be ascribed, in a generic way, to the varying aspect of the magnetic equatorial disc up to the Alfv\\'en radius;  in the case of \\zetaori, this variation is expected to be minimal.  We speculate that most of the redshifted H$\\alpha$ emission comes from a region located just above the photosphere (at the very base of the wind) and essentially reflects a difference between both magnetic poles (strongest emission being observed in conjunction with the weakest magnetic pole, see Fig.~~\\ref{fig:map}, bottom panel).  We also note that the excess absorption that both Balmer lines exhibit twice per rotation period in their distant blue wing (see Fig.~\\ref{fig:bal}, left and right panels) behave as UV absorption lines do in \\tori;  we propose that they reflect the magnetic equator crossing the line of sight (at rotation phases of about 0.20 and 0.75).  The maximum radial velocities associated to these absorption components (up to about 500~\\kms, i.e., less than 0.25\\vinf) confirm that they correspond to material located within the Alfv\\'en radius.  More data (densely sampled over several rotation cycles) are needed to investigate this issue more closely, and to pin down unambiguously the origin of the various H$\\alpha$ and H$\\beta$ components.  The 569.6~nm C~{\\sc iii} double-peak emission line is also a significant difference with respect to \\tori\\ and HD~191612 (where the line only features a single peak emission). The observed modulation is apparently related to the magnetic topology, with the red emission peaking at phases of maximum magnetic field and the blue emission peaking when the magnetic equator is crossing the line of sight (both phenomena occuring twice per rotation period).  The maximum velocities of both components (up to about 200~\\kms) also argue for the formation of this line within the Alfv\\'en radius and, therefore, it is likely a good indicator of the influence of the magnetic field on the launching of the wind. Further observational and theoretical studies are again required to examine how this line responds to a magnetised wind.  At the very least, our results demonstrate that the magnetic field of \\zetaori\\ has a significant impact on the wind despite being below pressure-equipartition and the weakest detected ever in a hot star.  Given the obvious importance of this result for our understanding of massive magnetic stars, we need to confirm and expand the present analysis with new data collected over several rotation periods of \\zetaori, i.e., over a minimum of 20 nights;  renewed observations will indeed allow us (i) to obtain an accurate measurement of the rotation period, (ii) to derive a fully reliable modeling of the large-scale magnetic topology and (iii) to estimate whether the field is  intrinsically variable as usual for dynamo topologies, e.g., on a typical timescale of 1~yr, and (iv) a detailed account of how wind lines (and in particular H$\\alpha$, H$\\beta$ and the 569.6~nm C~{\\sc iii} lines) are modulated with the viewing aspect of the magnetic topology."
    },
    "0806/0806.0217_arXiv.txt": {
        "abstract": "We used the Submillimeter Array (SMA) to image 860 \\micron\\ continuum and CO(3--2) line emission in the ultraluminous merging galaxy Arp 220, achieving a resolution of 0\\farcs23 (80 pc) for the continuum and 0\\farcs33 (120 pc) for the line. The CO emission peaks around the two merger nuclei with a velocity signature of gas rotation around each nucleus, and is also detected in a kpc-size disk encompassing the binary nucleus. The dust continuum, in contrast, is mostly from the two nuclei. The beam-averaged brightness temperature of both line and continuum emission exceeds 50 K at and around the nuclei, revealing the presence of warm molecular gas and dust. The dust emission morphologically agrees with the distribution of radio supernova features in the east nucleus, as expected when a starburst heats the nucleus. In the brighter west nucleus, however, the submillimeter dust emission is more compact than the supernova distribution. The 860 \\micron\\ core, after deconvolution, has a size of 50--80 pc, consistent with recent 1.3 mm observations, and  a peak brightness temperature of (0.9--1.6)$\\times 10^2$ K. Its bolometric luminosity is at least $2\\times10^{11}$ \\Lsol\\ and could be  \\about$10^{12}$ \\Lsol\\ depending on source structure and 860 \\micron\\ opacity, which we estimate to be of the order of  $\\tau_{860} \\sim 1$ (i.e.,  $N_{\\rm H_{2}} \\sim 10^{25}$ \\persquarecm). The starbursting west nuclear disk must have in its center a dust enshrouded AGN or a very young starburst equivalent to hundreds of super star clusters. Further spatial mapping of bolometric luminosity through submillimeter imaging is a promising way to identify the heavily obscured heating sources in Arp 220 and other luminous infrared galaxies. ",
        "introduction": "\\label{s.introduction} \\object{Arp 220} is the nearest ultraluminous infrared galaxy (ULIRG) and is an advanced merger ($\\Lir=10^{12.2} \\Lsol$; $D=75$ Mpc; 1\\arcsec = 361 pc; see Table \\ref{t.galparm}). Its vast luminosity is radiated almost entirely  between infrared and millimeter wavelengths as thermal dust emission. This is because the source, an extreme starburst, a quasar-class AGN(s), or both, is deeply embedded in dust  within the central few kpc \\citep{Soifer84,Joy86}. This merger has been a prime target to study the mechanism to generate the high luminosity of ULIRGs and to study galaxy evolution through major mergers \\citep{Sanders96, Genzel98}. Detailed studies of Arp 220 as well as other local ULIRGs may also help us to understand submillimeter-detected luminous galaxies at high redshift \\citep[e.g.,][]{Coppin06}, because the latter may be scaled-up versions of the former \\citep{Tacconi06}.    The Arp 220 system has two nuclei, presumably from the merger progenitors, with a projected separation of about 350 pc \\citep{Norris88,Graham90,Clements02}. Millimeter CO observations reveal a large amount of molecular gas (several $10^{9}$ \\Msol) in the central few kpc with peaks at the nuclei (\\citealt{Scoville91}; \\citealt{Scoville97} (\\altdef \\citest{Scoville97}); \\citealt{Downes98}; \\citealt{Sakamoto99} (\\altdef \\citest{Sakamoto99}); \\citealt{Downes07} (\\altdef \\citest{Downes07})). Sub-arcsecond resolution CO and \\ion{H}{1} observations show that a small gas disk ($r\\sim 50$ pc) rotates around each nucleus and that the two nuclear disks, counterrotating with respect to each other, are encompassed by a larger outer disk (\\citest{Sakamoto99}; \\citealt{Mundell01} (\\altdef \\citest{Mundell01}), in which one of the nuclear disks was modeled as a part of the outer disk.) The two nuclei dominate the dust emission at 1.3 mm \\citesp{Sakamoto99} and in mid-IR up to 25 \\micron\\ \\citep[$\\equiv$\\citest{Soifer99}]{Soifer99}, suggesting them to be a major source of the luminosity. Indeed, vigorous starburst activity in the nuclear disks was found from a concentration of radio supernovae and young supernova remnants within about 50 pc (\\about0\\farcs2) of each nucleus and from radio recombination lines in the nuclear disks \\citep{Smith98,Rovilos05,Lonsdale06,Anan00,Rodriguez05}. The presence of an AGN in Arp 220 has not been proven nor rejected based on X-ray observations because of high obscuration \\citep{Iwasawa01,Iwasawa05,Clements02}.   Arp 220 appears to represent a short luminous stage in merger evolution right before the coalescence of two nuclei, because there are luminous merging galaxies having similar structure. The advanced merger \\object{NGC 3256} ($\\Lir=10^{11.6} \\Lsol$) has twin gas peaks at two nuclei, possible rotation of the gas around each nucleus, and a large ($>$ kpc) gas disk surrounding the central region \\citep{Sakamoto06.n3256}. \\object{NGC 6240} ($\\Lir=10^{11.9} \\Lsol$) also has two nuclei within 1 kpc on the sky, with each having a rotating stellar core and being an AGN \\citep{Tecza00, Komossa03}. Studying Arp 220 is particularly important to understand this class of objects.    Accurate spatial distribution of the luminosity should tell us much about the hidden energy source(s) in Arp 220. The dichotomy of the nuclei (with the nuclear disks) and the outer disk has been proven useful for this approach. \\citest{Soifer99} showed using mid-IR imaging and the spectral energy distribution (SED) from infrared to millimeter wavelengths that the two nuclei  of \\about0\\farcs3 (\\about100 pc) extent have at least 40\\% of the total luminosity of Arp 220. They also pointed out that the outer disk, or a kpc-size extended component around the nuclei, could be another major component of the luminosity, contributing to the SED mainly at around 60--100 \\micron\\ but too faint in continuum to be detected at mid-IR and millimeter wavelengths. Further mid-IR spectroscopy and SED modeling suggested, however, that the less-obscured outer disk has a moderate starburst with up to 15\\% of the total luminosity, leaving a large majority of the luminosity to the heavily obscured nuclei \\citep{Soifer02,Spoon04,Gonzalez04}. The heavily embedded and compact source of luminosity was also suggested by \\citet{Dudley97} from mid-IR silicate absorption. Most of the luminosity in the nuclei is attributed to starburst by the groups that observed the radio supernovae and recombination lines. Still, there have been arguments for an AGN(s) \\citep[both based on (sub)millimeter line ratios favoring X-ray dominated regions]{Aalto07,Imanishi07} and even for its dominance in luminosity (\\citealt{Haas01}, based on deficit of 7.7 \\micron\\ PAH with respect to 850 \\micron\\ continuum; see \\citealt{Spoon04} for caveats).    Recently, \\citest{Downes07} obtained from their 0\\farcs3-resolution observations of 1.3 mm continuum a deconvolved size of $0\\farcs19 \\times 0\\farcs13$ and a deconvolved peak brightness-temperature of 90 K for the western nucleus. They attributed to the compact source a probable total luminosity of the nucleus, $9\\times10^{11}$ \\Lsol\\ in the nucleus-dominated model of \\citest{Soifer99}, and estimated the intrinsic dust temperature to be 170 K and the 1.3 mm opacity to be \\about0.7. They argued that the heating source must be an AGN on the basis of the high luminosity surface density. Although this model is possible, the minimum luminosity allowed by the 1.3 mm data is $3 \\times 10^{10}$ \\Lsol\\  in the blackbody limit; it is $1 \\times 10^{11}$ \\Lsol\\ if the 1.3 mm opacity is 1.3 as can be estimated from a power-law index of $\\beta=2$ for dust emissivity and the 2.6 mm dust flux in \\citest{Downes07}. Thus the majority of the luminosity in the west nucleus may come from the starburst in its nuclear disk rather than from its continuum core, and the source for the central luminosity may also be massive stars. Our quest for the accurate spatial distribution and the source of the luminosity in the merger is therefore not yet over.    In this paper, we report 0\\farcs2--0\\farcs3 resolution imaging of Arp 220 carried out with the SMA in the submillimeter CO(3--2) line and 860 \\micron\\ continuum. In submillimeter, where dust opacity is higher than in millimeter, we can better trace warm dust and better estimate the dust temperature and luminosity with less uncertainty from opacity correction. This is the first spatially-resolved imaging of the merger at a submillimeter wavelength and the second submillimeter interferometric observations of the galaxy following the original single-baseline experiment with the JCMT and CSO  \\citep{Wiedner02}. The SMA provided us with comparable or higher resolutions than previous millimeter observations. Our observations are described in \\S \\ref{s.obs} and data reduction in \\S \\ref{s.reduction}. The results are presented in \\S \\ref{s.result.co} for the CO line and in \\S \\ref{s.result.continuum} for continuum. We analyze the data for the dust opacity and temperature, and set a larger lower limit to the luminosity in the core of the west nucleus than the aforementioned limit from millimeter data by about an order of magnitude. The continuum data are also compared with the supernova distribution, and the line data are analyzed for gas dynamics and ISM properties in and around the nuclei. We discuss the properties of the hot nuclei and their energy source(s) in \\S \\ref{s.discuss}, showing that some luminosity-related parameters of the west nucleus are close to the highest for a starburst and that submillimeter observations may soon resolve the energy source issue. Our finding are summarized in \\S \\ref{s.summary}. ",
        "conclusions": "\\subsection{Mass and Dynamical Properties} \\label{s.nuclei.dynamical} The rotating disk around each nucleus is the currently accepted scheme, although there were alternative models attributing the velocity structure to velocity dispersion or non-circular motion \\citep[e.g.,][]{Downes98,Eckart01}. For the west nucleus, for example, HST/NICMOS images imply an inclined disk of absorption (i.e., gas and dust) elongated in the east-west direction \\citep{Scoville98}. The major axis and the extent of the disk agree with those of the observed gas velocity gradient. Radio supernovae and their remnants in the west nucleus are in a 160 pc by 70 pc region elongated in the east-west direction \\citep{Smith98,Lonsdale06,Parra07}, also consistent with the inclined and starbursting nuclear disk elongated in the same direction. \\citest{Downes07} also adopt a rotating disk.   The dynamical mass of each nucleus is \\about$10^{9}$ \\Msol\\ within \\about100 pc radius (see \\S\\ref{s.result.co.velocity} and Table \\ref{t.datasets}) as was previously derived from CO(2--1) and near-IR observations (\\citest{Sakamoto99}, \\citealt{Genzel01}). The estimate should be accurate to within a factor of a few, but it is difficult to tell in which direction it is likely to err. On one hand, the calculated mass underestimates the true mass because no inclination correction is applied, though the correction is likely small for the highly inclined west nuclear disk. On the other hand, the radius at which the rotational velocity peaks in each PV diagram is quite possibly overestimated because each nuclear disk is only marginally resolved in space; the dynamical mass will be overestimated by this error. We can not tell which of the two sources of bias is larger, and hence the calculated mass is neither a lower limit nor an upper limit. Still, the current data do not support a small mass of \\about$2\\times10^{8}/\\sin^2 i$ \\Msol\\ within about 100 pc from the west nucleus suggested by \\citest{Mundell01}. The small velocity gradient seen in \\ion{H}{1} absorption that lead to this mass estimate is incompatible with the large CO velocity range seen in our PV diagram (\\S \\ref{s.result.co.velocity}). Their model of merger configuration and history are partly based on the observation that the west nucleus is an order of magnitude lighter than the east nucleus, but that disparity is not supported by our data. If there is any disparity, the west nucleus {\\it appears} to be denser than the east nucleus as suggested from the larger velocity gradient of the former (see below), but that could be due to a lower inclination of the eastern disk.   The mean mass density calculated from the $V/R$ and the following formula, which does not assume a linearly-rising or any specific rotation curve, is $2\\times10^2$ \\Msol\\ \\percubicpc\\ for $r \\leq 110$ pc of the east nucleus and $2\\times10^3$ \\Msol\\ \\percubicpc\\ for $r \\leq 40$ pc of the west nucleus. \\begin{equation} \\left( \\frac{\\overline{\\rho}}{\\Msol\\, \\percubicpc} \\right) = 55 \\left( \\frac{V/R}{\\kms / \\rm pc} \\right)^2 \\end{equation} Elliptical galaxies and bulges have similar mass densities at the same scales \\citep{Lauer95} and our Galaxy has $\\overline{\\rho}\\, (r=12\\; {\\rm pc}) = 4 \\times10^3$ \\Msol\\ \\percubicpc\\ \\citep{Genzel96}. The mass densities that we derived are lower limits, because both the underestimate of the rotational velocity $V$ due to the lack of inclination correction and the overestimate of the radius $R$ due to insufficient spatial resolution make the ratio $V/R$ and hence the mean mass density smaller. For the same reason, we obtain  upper limits of 1--3 Myr to the dynamical time-scales (= rotation periods) of the nuclear disks from the velocity-radius pairs.     \\subsection{ISM Properties Suggested from Low CO(3--2) Equivalent Width} \\label{s.nuclei.ISMproperties_fromWeq} In addition to the high brightness temperature, our observations show that the nuclear disks have a much lower line-to-continuum ratio than the outer disk in the submillimeter, as was the case at 1.3 mm \\citesp{Sakamoto99}. Specifically, about 75\\% of the total (i.e., single-dish) 860 \\micron\\ continuum comes from the nuclear disks while only about 25\\% of the total CO(3--2) emission is from the nuclear disks. The rest of the emission is from the outer disk. To put it another way, the equivalent width for the CO(3--2) line is \\about$1\\times 10^3$ \\kms\\ in the nuclear disks, while that of the outer disk is several times larger, \\about$6\\times 10^3$ \\kms.   The larger equivalent width of the outer disk is in line with the observations of other IR bright galaxies and the small equivalent width (i.e., low line-to-continuum ratio) in the nuclear disks is exceptional. For example, \\citet{Seaquist04} showed that CO(3--2) emission typically comprises 25\\% of 850 \\micron\\ flux detected within a \\about$3\\times10^4$ \\kms\\ bandwidth at the central kiloparsecs of relatively IR-luminous galaxies ($\\log (L_{\\rm IR}/\\Lsol) \\approx$ 10--12) in the SLUGS survey \\citep{Dunne00,Yao03}. The fraction of 25\\% means a CO(3--2) equivalent width of $1\\times 10^4$ \\kms. Similarly, the equivalent widths in the central 15\\arcsec\\ (250 pc) of the starburst galaxies NGC 253 and M82 are $(1.5\\pm 0.3) \\times 10^4$ \\kms\\ and $(1.1\\pm 0.2) \\times 10^4$ \\kms, respectively, according to the 345 GHz observations of these galaxies \\citep{Alton99,Israel95,Tilanus91} and the prescription by \\citet{Seaquist04} to remove the CO(3--2) contribution from the 850 \\micron\\ bolometer flux.   A simple model tells us that the low equivalent width is basically due to either high temperature or high column density per velocity. Figure \\ref{f.lte} shows LTE calculations for CO(3--2) optical depth, excess brightness temperature, and equivalent width. The model assumes that gas and dust in the model region have the same uniform temperature, gas-to-dust mass ratio of 100, CO abundance of [CO/\\HH]=$10^{-4}$, and dust mass opacity coefficient at the CO(3--2) frequency of $\\kappa= 1$ \\unitofkappa, and also assumes that the dust emission is optically thin. Appendix \\ref{a.equiv-width} gives model formulae and the reason why the moderate dust optical depth of the nuclear disks does not affect the derived gas column density much. Figure 7 shows that there are two possible sets of  gas conditions for each set of equivalent width and brightness temperature. One has optically thin CO emission from warmer gas with lower column density per velocity, and the other has optically thick CO emission from cooler gas with higher column density per velocity.   The optically thin solution is unlikely for the nuclear region of Arp 220 because it needs a high temperature of about  800 K for the nuclear disks. The required temperature would be even higher for plausible non-LTE cases in which levels much higher than J=3 are less populated than in LTE because of their high critical densities. Such high temperatures are incompatible with the upper limit set on the dust temperature in \\S \\ref{s.result.continuum.tau_Td}. Moreover, the CO(3--2) to (2--1) intensity ratio $R_{(3-2)/(2-1)}$ is close to 1 at the two nuclei, consistent with optically thick emission. The ratio would be $\\frac{9}{4}$ if both lines were optically thin. Accepting that the CO(3--2) emission from the nuclear disks is optically thick, the bulk of the CO emission from the outer disk must also be optically thick because the intensity ratios between low-J CO lines are about 1 for the entire galaxy \\citep{Mauersberger99,Wiedner02}. Although the optically-thin solution of 800 K gas is unlikely, it provides a lower limit to the mean gas surface density in the nuclear disks, $\\Sigma_{\\rm H_{2}} > 4\\times 10^{23}\\, m_{\\rm H_{2}}\\, \\persquarecm = 6\\times 10^3$ \\Msol\\ \\persquarepc, because this solution has less gas surface density than the optically thick solution. This is a more stringent lower limit to the column density than the one in \\S \\ref{s.result.co.Tb_sigma}, thanks to the additional constraint from equivalent width.   In the preferred case of optically thick CO emission, the low equivalent width is due to a high column density per velocity width, $N_{\\rm CO}/\\Delta V \\approx 10^{18.5}$ \\persquarecm\\ \\perkms\\ for the equivalent width of $10^3$ \\kms. The \\HH\\ column density in the nuclear disks would be $\\Sigma_{\\rm H_{2}} = 1\\times 10^{25}\\, m_{\\rm H_{2}}\\, \\persquarecm = 2\\times 10^5$ \\Msol\\ \\persquarepc\\ if the line width in the disk is 300 \\kms, which is half of the total line width due to disk rotation and is an upper limit to the local line width in the disk. The high column density is in order-of-magnitude agreement with what we inferred from the dust emission alone in \\S\\ref{s.result.continuum.tau_Td}, as it should be. The high column density $N_{\\rm H_2} \\sim 10^{25}$ \\persquarecm\\ suggests that the volume density in the nuclear disks is high, $n_{\\rm H_2} \\gtrsim 10^{4.5}$ \\percubiccm\\ for a volume filling factor of 1 and a line-of-sight depth of $\\lesssim 100$ pc for the disks. The high density of the molecular gas is in agreement with the observations of molecular lines with high-critical densities \\citep[e.g.,][]{Wiedner02,Imanishi07} and is also in line with models of \\ion{H}{2} regions and supernova remnants \\citep{Anan00,Parra07}.   A more elaborate model that also employs optically thick CO emission and explains the lower CO equivalent width in the nuclear disks is that the ISM has a temperature gradient in such a way that the gas and dust are hotter inside and cooler near the surface of the disk. This reduces the CO emission through self-absorption and hence reduces the equivalent width of the line, because the CO line has higher optical depth than the continuum, $\\tau_{860} \\sim 1$. The absorption feature in the CO(2--1) line toward the west nucleus, observed by  \\citest{Downes07}, suggests that this is a part of the reason for the low equivalent width. It is quite possible that the disks have not only the temperature gradient but also a density gradient with the center having higher column and volume densities, as hinted by the compact dust emission.      \\subsection{Energy Source} \\label{s.discuss.energy} \\subsubsection{Sources in the Nuclear Region} Figure \\ref{f.vlbisources} shows the centimeter radio sources and the submillimeter continuum in the nuclear region. Relative positions of various centimeter sources in the figure are based on observations but the registration of centimeter and submillimeter data is made, as noted in \\S \\ref{s.result.continuum.positions}, with an assumption that the two nuclear sources seen in both wavelengths are the same. Specifically, the registration made the submillimeter peak coincide with the median centroid of the supernova-related sources (hereafter SNe) in the brighter west nucleus. One can alternatively use the peak of diffuse centimeter emission for the registration, because the offset between the two nuclei in that emission, ($\\Delta\\alpha, \\Delta\\delta$) = (0\\farcs96, \\minus0\\farcs15)  \\citep{Baan95}, also agrees with the offset measured in submillimeter. However, the western peak of the centimeter diffuse continuum (green circle in Fig. \\ref{f.vlbisources}) is only 40 mas from the center of the SNe distribution (red circle in Fig. \\ref{f.vlbisources}) according to the relative astrometry of the diffuse emission with respect to the SNe by \\citet{Rovilos03}. Thus our centimeter--submillimeter registration in Fig. \\ref{f.vlbisources} is probably accurate to \\about0\\farcs04.   Neither of the two OH megamasers in the west nucleus \\citep[shown in blue in Fig. \\ref{f.vlbisources}]{Lonsdale98} coincides with the submillimeter continuum peak, according to the registration made above. Of particular interest is the southern maser source called W2 in \\citet{Lonsdale98}. It has a large velocity gradient, $dV/dr \\approx 18.7$ \\kms\\ \\perpc\\ in a \\about6 pc diameter region, and may trace a heavily obscured AGN of \\about$1.7 \\times 10^{7}$ \\Msol\\ \\citep{Rovilos05}. Whether it is an AGN or not, our submillimeter map and its registration in Fig. \\ref{f.vlbisources} suggest that it is unlikely to be the main heating source of the submillimeter-emitting dust. We checked if the submillimeter data could be registered with the centimeter data using the OH megamasers and found it unlikely. Among the four OH megamasers in the nuclear region, the pair between the southern source in each nucleus has the separation that is closest to the one between the submillimeter peaks. This maser separation, ($\\Delta\\alpha, \\Delta\\delta$) = (0\\farcs89, \\minus0\\farcs15), is in worse agreement with the submillimeter separation (0\\farcs96, \\minus0\\farcs18) than the supernova-based separation (0\\farcs94, \\minus0\\farcs18) and the abovementioned separation in diffuse centimeter continuum. Moreover, the shift needed for the registration is larger if the southwest megamaser should coincide with the west submillimeter peak. Furthermore, the elongation in the northeast--southwest direction seen both in the submillimeter emission and the supernova distribution of the east nucleus suggests that these two, rather than the submillimeter and OH emission, should be correlated. Thus, although the final conclusion should be made from more accurate astrometric observations, we favor the registration in Fig. \\ref{f.vlbisources} and the SNe--submillimeter connection over the megamaser--submillimeter connection. The maser emission off the west nuclear disk may be collisionally excited by a bipolar wind or outflow from the disk.    \\subsubsection{Supernovae and Dust Emission} The spatial distribution of the submillimeter continuum is expected to follow the surface density of supernovae under certain circumstances that include the dominant energy source being a dust-enshrouded starburst (see Appendix \\ref{a.Id-Sigma_sn}). We show  the smoothed distribution of the SNe in Figure \\ref{f.smoothedSNe}. In Fig. \\ref{f.smoothedSNe}a  the  smoothing kernel is a 0\\farcs5 Gaussian that is about the same as the PSF of our low-resolution data (Fig. \\ref{f.low}) and the PSF of the mid-IR images of \\citest{Soifer99}. As noted earlier, the eastern nucleus has an oval shape elongated in the same NE-SW direction in the supernova distribution, 860 \\micron\\ continuum (see the deconvolved shape in Table \\ref{t.datasets}), and mid-IR continuum, among others. This agreement is suggestive of the dust being heated mainly by the starburst, although there may still be room for an accreting black hole to hide within or behind the starbursting disk. In contrast to the east nucleus, the brighter west nucleus shows different shapes in the submillimeter continuum and in the supernova distribution (see Figs. \\ref{f.smoothedSNe}b and  \\ref{f.vlbisources}). The dust continuum is more compact and the starburst is much more elongated in the east-west direction at the 0\\farcs23 resolution. This bright 860 \\micron\\ core has a minimum luminosity of $2\\times10^{11}$ \\Lsol, calculated in \\S \\ref{s.result.continuum.luminosity}.   The reason for the different spatial distribution between the submillimeter continuum and SNe must be either that many massive stars and their remnants are missed in the current VLBI observations at the densest region or that there is a dust-heating source other than massive stars in the center of the west nuclear disk. The former could be because of  free-free absorption, radio quiet supernovae, younger stellar population at the center of the W nucleus than the one around it and in the E nucleus, or more/less dust and gas in the nucleus/disk of Arp 220 W than its star formation implies. The non-stellar heating source, if it exists, must be an AGN.   Quantitatively, the west nucleus has three of the four new radio sources, i.e., radio supernovae, found by \\citet{Lonsdale06} in the merger over a period of one year. Two of them are in the center of the west nucleus, and \\citet{Parra07} pointed out that more core-collapse supernovae are probably left undetected because the detected ones are most likely type IIn that are radio-luminous and rare. It may therefore be possible that there is a stronger concentration of supernovae resulting from vigorous star formation at the center of the west nuclear disk than is visible in the existing VLBI observations. The starburst that caused the observed supernovae can provide the entire luminosity of the merger even before the correction for non-type IIn SNe \\citep{Smith98,Lonsdale06,Parra07}, although such an estimate involves assumptions on the stellar mass function and starburst history, which will be difficult to verify.   \\subsubsection{Constraints from Luminosity-to-Mass Ratio and Luminosity Surface Density} \\label{s.nulcei.energy_source.L/M_and_Sigma_L} Among quantitative parameters that could constrain the energy source from our data are luminosity-to-mass ratio and luminosity surface density. A lower limit to the luminosity-to-mass ratio within $r=40$ pc of the west nucleus is about $4\\times10^{2}$ \\Lsol/\\Msol\\ according to the parameters estimated in \\S \\ref{s.result.continuum.luminosity} and \\S \\ref{s.nuclei.dynamical}. The ratio would be as high as $10^3$ \\Lsol/\\Msol\\ if the nucleus has a luminosity of $10^{12}$ \\Lsol\\ as shown possible depending on its size and optical depth. The $L/M$ ratio, which depends on stellar mass function, may be as high as \\about$1\\times10^{3}$ \\Lsol/\\Msol\\ for a very young starburst  \\citep[age $< 10$ Myr, see, e.g.,][]{Leitherer99}. Its more robust upper limit is \\about$1\\times10^{4}$ \\Lsol/\\Msol\\ for a pure population of 100 \\Msol\\ stars. The high ratios in the range of ($10^{3}$--$10^{4}$) \\Lsol/\\Msol\\ are observed in super star clusters (SSCs) such as the one in NGC 5253 studied by \\citet{Turner00} and the Quintuplet and other clusters in \\citet{Figer99}. For comparison, an AGN fueled at the Eddington rate has a luminosity to black-hole mass ratio of $4\\times10^{4}$ \\Lsol/\\Msol. If the $L/M$ ratio is close to our lower limit, then it is quite possible, as far as this parameter is concerned, that the main luminosity source of the nucleus is an intense starburst producing many super clusters. The $L/M$ ratio for the $10^{12}$ \\Lsol\\ case, on the other hand, is close to the largest a starburst can have, considering that the observed dynamical mass includes everything in the region, such as young and old stars, massive black holes (if any), and the gas and dust of high column densities. We can therefore infer that the west nucleus, if it has a luminosity of \\about$10^{12}$ \\Lsol, must have either an energetically significant AGN or an extreme starburst that is compact ($r\\leq40$ pc), young (age $\\lesssim 10$ Myr), and maybe biased toward high mass stars.   The luminosity surface density of the west nucleus is $10^{7.6}$ \\Lsol\\ \\persquarepc\\ or larger in the central 80 pc. The lower limit is already among the highest for infrared-luminous starburst galaxies \\citep[][their Table 4]{Soifer01}. If the west nucleus has a luminosity of $10^{12}$ \\Lsol\\ within 50 pc diameter, then the parameter reaches $10^{8.7}$ \\Lsol\\ \\persquarepc. Obviously, a luminous AGN would have no problem to achieve this level of luminosity surface density. At the same time, there are known stellar systems that could also achieve the high values. Some of the youngest super star clusters have core radii of 0.1--1 pc, luminosities of \\about$10^{7}$--$10^{9}$ \\Lsol, and ages less than 10 Myr \\citep{Figer99, Turner00}. Taking a luminous SSC of $10^{9}$ \\Lsol\\ as a yard stick, there would be 200 and 1000 of them, respectively, for the ($L$, $d$)=($2\\times10^{11}$ \\Lsol, 80 pc) and ($10^{12}$ \\Lsol, 50 pc) cases mentioned above if the SSCs are the main luminosity source. Their mean spacing would be 11 pc and 4 pc, respectively. The clusters would probably overlap in their outskirts, since half of the \\about7000 O stars in the most luminous SSC in NGC 5235, used as the yard stick, are in the central 5 pc \\citep{Turner04}. Alternatively, many of the clusters may have already been disintegrated to form a single concentration of massive stars in the galaxy nucleus, considering the short dissolution time of SSCs \\citep{Fall05}. It could also be that  a smaller number of heavier cluster(s) was formed rather than hundreds of SSCs. This possibility cannot be excluded on the grounds of the number density of massive stars required to produce the luminosity surface density of Arp 220 W, since higher number densities are observed in SSCs.   \\subsubsection{Star Cluster v.s. AGN} As seen above, the parameters of the luminous west nucleus are in or close to the limit explicable by a star cluster or clusters. The starburst model may explain some of the observations that appear to favor an AGN but it certainly has some difficulties. The extreme luminosity surface density even for a ULIRG is expected to result in photodestruction of PAH particles as was observed around the model SSC in NGC 5253 \\citep{Beirao06}. This may be part of the reason for the merger's low PAH emission to 850 \\micron\\ continuum ratio that was taken as evidence for an AGN \\citep{Haas01}. The excess dust emission at the center of the west nucleus compared to the number of supernovae and young supernova remnants is reminiscent of the abovementioned SSC in NGC 5253 whose predominantly thermal radio spectrum is attributed to a very young starburst with few supernovae \\citep{Beck96}. On the other hand, as \\citet{Haas01} emphasized, a cluster of stars (or SSCs) with an extent of a few 10 pc or larger would be more difficult to enshroud with dust than an accreting black hole. The cluster model may therefore be incompatible with the deeply enshrouded but energetically dominant component used by \\citet{Spoon04} and \\citet{Dudley97} to model the mid-IR spectrum. Another difficulty in the cluster model is that it needs to form hundreds of SSCs or equivalent  high-mass stars at high volume density in the short time of $< 10$ Myr. The short time scale is necessary because the luminosity-to-mass ratio of a cluster declines once massive stars start to die and also because gas and dust in the nucleus would be dispersed by the energy injection from the starburst. Although SSCs tend to form in groups \\citep{Zhang01}, we do not yet know the mechanism for such a compact and intense burst of SSC formation or even whether it is possible.    A massive black hole of \\about$10^{8}$ \\Msol\\ can generate (0.2--1)$\\times10^{12}$ \\Lsol\\ via mass accretion at a sub-Eddington rate, as it is believed to do in quasars. The above-mentioned constraints on mass-to-luminosity ratio and luminosity surface density of the west nucleus can be easily met with a buried quasar. The weaker concentration of supernovae and supernova remnants toward the center of W than the concentration of the submillimeter continuum and bolometric luminosity can be explained if a large part of the luminosity is due to an AGN. There is enough material, $N_{\\rm H_2} \\sim 10^{25}$ \\persquarecm, to block the X-rays, and it is easier to enshroud an AGN than a million O stars. An AGN with a high-covering factor could explain the large equivalent width of the Fe K$\\alpha$ line of Arp 220 \\citep{Iwasawa05}. Regarding the formation mechanism, it is not surprising that the remnant nucleus of one or both of the merged galaxies has a massive black hole, since most if not all galaxies are thought to harbor a massive black hole at the center. The massive black hole in Arp 220 W would be at the center of the nearly edge-on nuclear disk of gas and dust, having abundant gas in its vicinity to accrete, and might have been significantly growing through the increased supply of fuel. The number of supernovae and the amount of free-free emission that appear to suggest a starburst large enough to provide the entire luminosity of Arp 220 is a problem for the luminous-AGN model. Our observations suggest that neither of the two OH megamasers about 50 pc off the luminosity center marks the main luminosity source even if they are associated with AGNs.    Taking everything above into account, a luminous AGN at the center of the west nucleus, proposed most recently by \\citest{Downes07}, seems to have fewer difficulties than the SSC-cluster model if the luminosity of the bright 860 \\micron\\ core is significantly larger than the lower limit we set. If, on the other hand, the luminosity is close to the observational lower limit, then the majority of the luminosity of the west nucleus is emitted from the nuclear disk rather than the core. In this case the heating source for the disk, as well as for the core, can be a starburst. In other words, the spatial distribution of luminosity within the west nucleus is found more centrally peaked than known before (including \\citest{Downes07}), but the parameters of the concentration obtained so far are not enough to conclude the existence of an energetically-dominant AGN. Really robust evidence is needed here considering the number of previous studies that have concluded one way or the other about this issue.    \\subsection{Prospects} Further observations will let us set tighter constraints on the energy sources and their spatial distributions. For example, higher resolution observations at around 860 \\micron\\ will better determine the luminosity-to-mass ratio. Such observations may soon become possible through the ongoing eSMA project that combines the SMA with the neighboring JCMT and CSO telescope. Observations at even shorter submillimeter wavelengths, where dust opacity is higher, will better constrain the dust temperatures in the nuclear and outer disks and will thereby determine their relative contributions to the total luminosity. They also let us search for a compact and bright core in the east nucleus. The luminosity estimate from  dust thermal emission in the submillimeter wavelengths complements other methods, because it bypasses the number of ionizing photons or supernovae to estimate luminosity and  also because it is less hindered by foreground extinction than infrared-based methods. In addition, various molecular lines in millimeter and submillimeter will tell us the physical properties of the ISM and alert us to any anomaly in the nuclei. Ongoing projects in these directions, further modeling of the structure and radiation of the ISM, and the eventual arrival of ALMA will tell us the heating sources not only in Arp 220 but also in ULIRGs in general. As an example, high-resolution photometry at a \\about10 pc resolution across the ALMA observing bands will allow us to locate a Compton-thick AGN of $10^{12}$ \\Lsol\\ as a source of a few 100 K, even if it is hidden in other wavelengths behind layers of absorbers. The dynamical mass estimated from molecular lines may reveal that its luminosity to mass ratio is too large for a starburst. Finding many such sources may prove the long-hypothesized formation of quasars in luminous mergers \\citep[and references therein]{Sanders88,Hopkins06}.      The ultraluminous merging galaxy Arp 220 has been imaged using the Submillimeter Array (SMA) at 0\\farcs2--0\\farcs3 (\\about100 pc) resolutions in the CO(3--2) line and 860 \\micron\\ continuum. We analyzed structure and properties of the submillimeter emission with particular attention to the spatial distribution of bolometric luminosity. The luminosity distribution allows us to constrain the energy source of Arp 220. Our main results are as follows.   1. Both the CO(3--2) line and the submillimeter continuum peak at or around the two nuclei of the merger. The continuum emission is mostly from compact (a few 0\\farcs1) regions at the two nuclei. The CO gas is more extended than the continuum at the nuclei, and it also has a component with a total extent of \\about3\\asec (1 kpc) that encompasses the double nucleus. The overall distribution of CO and dust emission is in agreement with the observations at lower frequencies.   2. The peak brightness temperature exceeds 50 K in both the CO(3--2) line and the submillimeter continuum emission, revealing the presence of warm gas and dust in the center of the merger. The west nucleus has the strongest continuum emission that is consistent with a Gaussian of FWHM 0\\farcs15 (50 pc) and peak brightness temperature $1.6\\times10^{2}$ K or a uniform-brightness disk of diameter 0\\farcs23 (80 pc) and a brightness temperature of  $0.9\\times10^{2}$ K.   3. The dust opacity at both nuclei is on the order of unity at 860 \\micron, according to the analysis of flux ratios to 1.3 mm and to 25 \\micron. Both the 860 \\micron\\ opacity and low equivalent widths of the CO(3--2) line at the two nuclei (\\about$10^{3}$ \\kms) suggest a high column density of gas and dust toward the nuclei, of the order of $N_{\\rm H_{2}} \\sim 10^{25}$ \\persquarecm. The equivalent-width analysis and the CO(3--2) to CO(2--1) intensity ratio close to 1 suggest that the CO(3--2) line is optically thick and the molecule is thermalized at least up to the J=3 level in the nuclear disks. The gas there must be also dense, $n_{\\rm H_2} \\gtrsim 10^{4.5}$ \\percubiccm.   4. The bolometric luminosity of the bright 50--80 pc core within the west nucleus is at least $2\\times10^{11}$ \\Lsol\\ and may be $10^{12}$ \\Lsol\\ or higher depending on the opacity and structure of the source. Our blackbody lower limit is an order of magnitude larger than that from existing 1.3 mm observations. The east nucleus has a lower limit to the bolometric luminosity of $5\\times10^{9}$ \\Lsol. It also may be more luminous if its 860 \\micron\\ emission is optically thin.   5. The sharp velocity shift across each nucleus is confirmed in the CO(3--2) velocity field, and so is the overall rotation of the larger (\\about kpc) gas disk. Each nucleus has a line-width of \\about500 \\kms\\ and a dynamical mass of \\about$1\\times10^{9}$ \\Msol\\ within a radius of \\about100 pc if the gas motion is due to rotation, as seems most likely. The mean mass densities in the nuclei are on the order of $10^2$--$10^3$ \\Msol\\ \\percubicpc, and are comparable to those of galaxy centers at the same scale. The overall merger system has a full CO line width of \\about1100 \\kms\\ and does not show extremely high velocity CO at our sensitivity. The overall structure and kinematics are consistent with what has been proposed, a binary system of counter-rotating nuclear disks encompassed in a larger disk.   6. Morphological agreement between dust continuum and surface density of supernova-related VLBI sources in the east nucleus is suggestive that the dust is mainly heated by a starburst. In contrast, the west nucleus has more compact dust emission than the supernova distribution. The compact hot core has a luminosity-to-mass ratio of $\\gtrsim 4\\times10^{2}$ \\Lsol/\\Msol\\ and a luminosity surface density of $\\geq 10^{7.6}$ \\Lsol\\ \\persquarepc\\ in the central 80 pc, with a possibility of actually having \\about5 times larger values than these lower limits. These parameters, at least the lower limits, are similar to those in young super star clusters and are close to the highest for a starburst. Hence we can not yet rule out an extreme starburst equivalent to hundreds of young SSCs. A buried luminous AGN accompanied by a starburst can also account for the luminosity and other observed parameters of the west nucleus, and it becomes more plausible in higher luminosity. Further observations as well as modeling are needed to determine the luminosity distribution and decide the source. High resolution imaging of submillimeter emission is a promising way to accomplish that. The OH megamaser in the west nucleus is off the peak of dust continuum and luminosity, and is hence unlikely to mark the position of a dominant heating source such as an AGN."
    },
    "0806/0806.2024_arXiv.txt": {
        "abstract": "We analyzed the eclipse light curve of the nova-like star RW Tri in its low luminosity state. During approximately 150 days, RW Tri was about one magnitude fainter than in its usual state. Our eclipse map shows that the brightness temperature in the disc ranges from 19000 K near the white dwarf to 8700 at the disc edge. For the inner parts of accretion disc, the radial temperature distribution is flatter than that predicted from the steady state models, and for the outer parts, it is close to the $R^{-3/4}$ law. Fitting of the temperature distribution with one for the steady state disc model gives a mean accretion rate of (3.85$\\pm$0.19)$\\cdot$10$^{-9}$ M$_\\odot$ year$^{-1}$. The hotspot in the disc is placed at a distance of 0.17$a$ from the white dwarf, where $a$ is the orbital separation. ",
        "introduction": "RW Tri is a bright well known eclipsing nova-like system. It was discovered by Protitch in 1937 \\citep{pro37}. \\citet{wal63} determined the orbital period to be 5.57~h. \\citet{afr78} found that eclipse timings demand that the ephemeris has a cyclic term with a period of 2777 or 4980 days. Different authors give different values of the system inclination angle $i$: 80$^\\circ$ \\citep{lon81}, 82$^\\circ$ \\citep{frk81}, 70.5$^\\circ$ \\citep{sma95}. \\citet{frk81} found that the disc size is about 0.4$a$ where $a$ is the orbital separation.  \\citet{hos85} performed the first eclipse mapping of the system. They found that the temperature of the inner part of accretion disc is about 40000 K. Also using the eclipse mapping technique, \\citet{rpt92} determined the mass accretion rate to be 3$\\cdot$10$^{-8}$ M$_\\odot$ year$^{-1}$.  \\citet{poo03} estimated the range for the primary and secondary stellar masses as 0.4  - 0.7 and 0.3 - 0.4  M$_\\odot$ respectively. \\citet{grp04} using spectrophotometric data found that the mass accretion rate is about 10$^{-8}$ M$_\\odot$ year$^{-1}$. Trigonometric parallax determination with the Hubble Space Telescope gave a distance of 341 pc to RW Tri \\citep{mca99}. ",
        "conclusions": "Using eclipse mapping techniques we calculated the radial brightness temperature distribution. For inner parts of accretion disc the slope of this distribution is close to the $R^{-1/2}$ law. For outer parts the temperature distribution corresponds to a steady state $R^{-3/4}$ law. We estimated the mass accretion rate in the system as $\\dot M$ = (3.85$\\pm$0.19)$\\cdot$10$^{-9}$ M$_\\odot$ year$^{-1}$. Our results show that even during the low luminosity phase, disc remains in the hot steady state."
    },
    "0806/0806.2815_arXiv.txt": {
        "abstract": "Using VLT/SINFONI, we have obtained repeated AO-assisted, NIR spectroscopy of the three central WN6ha stars in the core of the very young ($\\sim1$ Myr), massive and dense Galactic cluster NGC3603. One of these stars, NGC3603-A1, is a known 3.77-day, double-eclipsing binary, while another one, NGC3603-C, is one of the brightest X-ray sources among all known Galactic WR stars, which usually is a strong indication for binarity. Our study reveals that star C is indeed an 8.9-day binary, although only the WN6ha component is visible in our spectra; therefore we temporarily classify star C as an SB1 system. A1, on the other hand, is found to consist of two emission-line stars of similar, but not necessarily of identical spectral type, which can be followed over most the orbit. Using radial velocities for both components and the previously known inclination angle of the system, we are able to derive absolute masses for both stars in A1. We find $M_{\\rm 1} = (116 \\pm 31) M_{\\sun}$ for the primary and $M_{\\rm 2} = (89 \\pm 16) M_{\\sun}$ for the secondary component of A1. While uncertainties are large, A1 is intrinsically half a magnitude brighter than WR20a, the current record holder with 83 and 82 $M_{\\sun}$, respectively; therefore, it is likely that the primary in A1 is indeed the most massive star weighed so far. ",
        "introduction": "While models maintain that in the early Universe, the first-generation of stars were very massive and reached masses between 100 and 1000 M$_{\\odot}$ (e.g. \\citealt{NakaUme01}; \\citealt{Schaerer02}), it is generally accepted that under present-day conditions, relatively fewer massive stars are formed, i.e. that the initial-mass function (IMF) is much steeper and, more importantly, has a cut-off occurring around 150 M$_{\\odot}$ (\\citealt{WeidKroup04}; \\citealt{Figer05}). So far, however, whenever Keplerian orbits of binary systems are used to weigh stars -- the only way to obtain reliable, least model-dependent masses --, measured masses fall short by almost a factor of two with respect to the putative cut-off. Currently, stars with the highest known masses are both WN6ha components of the Galactic WR binary WR20a, with 83 and 82 M$_{\\odot}$, respectively (\\citealt{Rauw04}; \\citealt{Bonanos04}), and the O3f/WN6 star in the Galactic binary WR21a with a minimum mass of 87 M$_{\\odot}$ (\\citealt{Gamen08pp}). Significantly more massive stars, however, have so far remained elusive.  A most remarkable result that has emerged from the search for very massive stars is, however, that the highest Keplerian masses are \\emph{not} found among absorption-line O-type stars -- rather, masses of these stars remain below $\\sim60$ M$_{\\odot}$ (e.g. \\citealt{Lamontagne96}; \\citealt{Massey02}) --, but among Wolf-Rayet (WR) stars, more precisely among the so-called WN5-7ha (or WN5-7h) stars, an extremely luminous and hydrogen-rich subtype of the nitrogen-sequence WR stars. Contrary to classical WR stars, which are identified with bare, helium-burning cores of evolved massive stars, theoretical work confirms that these luminous WN5-7ha stars are hydrogen-burning, unevolved objects (\\citealt{deKoter97}; \\citealt{CroDess98}) which mimic the spectral appearance of WR stars because their high luminosities drive dense and fast winds (\\citealt{GraeHam08pp}).  If it is true that the most massive stars can be found in the most massive clusters (\\citealt{Bate02}; \\citealt{WeidKroup06}), then it can be expected that the most massive WN5-7ha stars are hosted by the most massive among the youngest, least evolved clusters known. NGC3603 is a Galactic example of such a cluster, and virtually a clone of its more famous LMC counterpart, the supermassive cluster R136, at the core of the giant HII region 30 Dor ({\\citealt{Moff94}). NGC3603's very core, itself denoted HD 97950, contains three extremely luminous WN6ha stars (\\citealt{Drissen95}), with stellar luminosities well in excess of $10^{6}$ L$_{\\odot}$ (\\citealt{deKoter97}; \\citealt{CroDess98}). \\citet{MoffNiem84} found from radial-velocity variations in unresolved spectra of HD 97950 that one of these WN6ha stars must be a binary with a period of 3.772 days, which was confirmed by \\citet{Moff85}. In a recent study, \\citet{Moff04} used HST-NICMOS J-band photometry to confirm that A1 is indeed a double-eclipsing binary with that period, while the other WN6ha stars B and C did not vary above the $\\sim0.05$ mag noise level. However, star C shows an extremely large X-ray luminosity ($L_{\\rm x} =$ several $10^{34}$ ergs$^{-1}$; \\citealt{Moff02}), which is a strong indication of binarity given that colliding winds in binary systems generate copious amounts of hard X-rays (e.g. \\citealt{Usov92}).  The double-eclipsing nature of A1 offers the rare opportunity to directly measure the mass of an extremely luminous binary system by simple, least model-dependent Keplerian orbits. Here we report the results of a spectroscopic monitoring campaign of the three central WN6ha stars in NGC3603. This Letter is organized as follows: In Section 2, we briefly describe the observations and data reduction. In Section 3, data analysis is described and results are presented. Section 4 summarizes the Letter. ",
        "conclusions": "We have obtained repeated, spatially resolved, AO assisted, near-IR spectroscopy of the three central WN6ha stars in HD 97950, the core of the young, unevolved and very massive Galactic cluster NGC3603. One of the stars, A1, is a previously known, double-eclipsing binary with an orbital period of 3.77 days (\\citealt{MoffNiem84}; \\citealt{Moff85}; \\citealt{Moff04}), while a second star, C, was newly identified as a binary in the present study. The third star, B, showed constant RVs over the observed time interval, and therefore is most likely not a binary, which is in line with its normal X-ray luminosity (\\citealt{Moff02}).  While in star C only the primary (WN6ha) component is visible -- the system is therefore classified as an SB1 binary --, A1 consists of two emission-line stars, most likely of similar, but hot identical spectral types. From the radial-velocity curves of the two components and the known inclination angle of the system (\\citealt{Moff04}), we derived component masses of $M_{\\rm 1} = (116 \\pm 31) M_{\\sun}$ for the primary and $M_{\\rm 2} = (89 \\pm 16) M_{\\sun}$ for the secondary, respectively. Despite the large uncertainties, we consider the primary WN6ha component of A1 to be the most massive star ever directly weighed.  While the primary component of C might have a mass similar to or even greater than that of A1's primary, it is possible that star B, single yet only slightly fainter than the combined binary system A1, is indeed the most massive member in NGC3603 and, therefore, the most massive main-sequence star known in the Galaxy."
    },
    "0806/0806.3434_arXiv.txt": {
        "abstract": "For a decade, N--body simulations have revealed a nearly universal dark matter density profile, which appears to be robust to changes in the overall density of the universe and the underlying power spectrum. Despite its universality, the physical origin of this profile has not yet been well understood.  Semi--analytic models by \\citet{Barnes05} have suggested that the density structure of dark matter halos is determined by the onset of the radial orbit instability (ROI).  We have tested this hypothesis using N--body simulations of collapsing dark matter halos with a variety of initial conditions.  For dynamically cold initial conditions, the resulting halo structures are triaxial in shape, due to the mild aspect of the instability.  We examine how variations in initial velocity dispersion affect the onset of the instability, and find that an isotropic velocity dispersion can suppress the ROI entirely, while a purely radial dispersion does not. The quantity $\\sigma^2/v_c^2$ is a criterion for instability, where regions with $\\sigma^2/v_c^2 \\lesssim 1$ become triaxial due to the ROI or other perturbations.  We also find that the radial orbit instability sets a scale length at which the velocity dispersion changes rapidly from isotropic to radially anisotropic.  This scale length is proportional to the radius at which the density profile changes shape, as is the case in the semi--analytic models; however, the coefficient of proportionality is different by a factor of $\\sim$2.5.  We conclude that the radial orbit instability is likely to be a key physical mechanism responsible for the nearly universal profiles of simulated dark matter halos. ",
        "introduction": "The universal density profile of dark matter halos has been widely observed in cosmological N--body simulations \\citep{Cole96,Moore99,Bullock2001,vdB02,Navarro04}; however, the physical origins are not yet understood.  Whether formed in an isolated collapse or by hierarchical merging, dark matter halo density profiles have a characteristic double-power law shape.  This shape is usually discussed in terms of the slope of the density profile $\\gamma \\equiv d(\\log{\\rho})/d(\\log{r})$.  The canonical halo profile has $\\gamma \\sim -1$ in the inner regions and $\\gamma \\sim -3$ in the outer regions \\citep{Dubinski91,NFW,Navarro97,Huss99,MacMillan06}. While there are disagreements on the exact values of $\\gamma$, the halo density profiles appear to be similar over decades in radius. There is also evidence that such profiles may accurately fit the luminous parts of early--type galaxies as well \\citep{Dalcanton01,Merritt05}.  Within dark matter halos, quantities other than density show universal profiles as well.  For example, \\citet{Taylor01} and later \\citet{Dehnen05} showed that the phase--space density, represented by $\\rho/\\sigma^3$, has a constant power-law slope of $\\alpha \\sim 1.9$ for over two and a half decades in radius in spite of a varying density profile (see also \\citet{Austin05} for an analytic exploration of this behavior).  There is also evidence for a universality in the relation between the slope of the density profile and the velocity anisotropy.  The anisotropy parameter, $\\beta$ is defined as  \\begin{equation} \\beta \\equiv 1 - \\sigma_{\\phi}^2/2\\sigma_r^2, \\end{equation}  \\noindent where $\\sigma_{\\phi}$ is the velocity dispersion in the tangential direction and $\\sigma_r$ is the velocity dispersion in the radial direction.  $\\beta$ describes the degree of anisotropy of the velocity dispersion in a spherical halo, which in simulations ranges from isotropic in the core ($\\beta=0$) to radially anisotropic at the outskirts ($\\beta=1$)\\citep{Cole96,Carlberg97,Fukushige01}.  \\citet{Hansen06} find a direct correlation between $\\beta$ and the density profile slope for halos with a wide variety of initial conditions (i.e. isolated collapse, mergers).  \\citet{Hanseneveryone06} find another universality, this time in the velocity distribution function, for a similar variety of initial conditions.  The radial orbit instability (ROI) may be the cause of the apparent universality of these dark matter halo properties.  The ROI occurs in anisotropic spherical systems composed of particles with predominantly radial orbits.  The instability arises when particles in precessing elongated loop orbits experience a torque due to a slight asymmetry. This torque causes them to lose some angular momentum and move towards the center of the system.  Particles with small enough angular momenta become trapped in box orbits, aligning themselves with the growing bar and reinforcing the initial perturbation.  As a result, the halo becomes triaxial in shape.  This phenomenon has been examined in detail by \\citet{Dejonghe88} using N--body simulations, by \\citet{ Weinberg91} using linear analysis, and by \\citet{Huss99} and \\citet{MacMillan06} for the specific case of isolated dark matter halo collapse.  These groups find that the onset of the ROI corresponds to a flattening of the central density cusp of the halo, and may be responsible for the double power-law shape of the halo.  The ROI is reviewed in Merritt (1999; \\S 6.2). \\nocite{Merritt99}  Semi-analytic models have also been used to examine the spherical collapse of dark matter halos, beginning with \\citet{Gunn72}.  These models \\citep{Gott75, Gunn77, Bertschinger85} include only radial motions and produce single power-law density profiles.  When non--radial motions are included, halo density profiles range from power-laws \\citep{Ryden87} to NFW--like \\citep{Hiotelis02,LeDelliou03,Ascasibar04}.  Recently, however, \\citet{Barnes05} extended the semi-analytic model of \\citet{Williams04} to include a physical representation of the ROI. The resulting halos have a double power-law density slope similar to those of N--body simulations.  \\citet{Barnes05} concluded that the density scale length and the anisotropy radius are correlated, and that the ROI is directly responsible for the shape of the density profile.  While \\citet{Barnes05}'s work is suggestive, the complexities of a dynamical instability are not easily captured by analytic or semi--analytic techniques.  We therefore turn to studying the link between the ROI and halo structure by using high resolution N--body simulations.  In this paper, we examine in detail the effect of random motions on the onset of the ROI and the final structure of collapsed dark matter halos, and attempt to verify the relation between the scale length and the anisotropy radius reported by \\citet{Barnes05}.  We analyze a set of N--body simulations of isolated, collapsing dark matter halos with a variety of initial velocity dispersions to study the evolution of halo properties.  Considering a range of velocity dispersions allows us to supress the ROI in some cases, helping us to isolate its physical effects.  In \\S2 we describe our simulations and show that they are robust to resolution and softening effects.  We describe the properties of the halos in \\S3, including the shape and anisotropy evolution, the effects of velocity dispersion, and the scale-length -- anisotropy radius relation.  We summarize our results in \\S4. ",
        "conclusions": "We have conducted N--body simulations of a variety of isolated, collapsing dark matter halos to investigate the role of the radial orbit instability on the final halo structure.  The ROI plays a role in determining the eventual shape, density profile, and phase--space density profile of a halo.  The ROI is suppressed in halos with isotropic initial velocity dispersions, but halos with purely radial velocity dispersions appear to undergo the instability.  Many numerical simulations have produced triaxial dark matter halos \\citep{Frenk88,Warren92, Thomas98, Jing02}.  Observational evidence for triaxial halos is also appearing, from low surface brightness and dwarf galaxies \\citep{Bureau99, Simon05, Hayashi06} and from galaxy clusters \\citep{Oguri03, Lee04}. Such evidence for the existence of triaxial halos strengthens our hypothesis that dynamically cold halos undergo the ROI during collapse and therefore obtain a triaxial shape. A warmer system would not become unstable and would have a final shape that is too spherical to explain the observed structures.  While the linear overdensities in the early universe are themselves triaxial, not spherical \\citep{BBKS}, and hence also contribute to the triaxiality of the resulting halos, the ROI significantly affects the inner, observable parts of the halo.  The characteristic anisotropy profile seen in all of our halos, and in others in the literature, is independent of the ROI, and appears to be a result of the general process of halo collapse.  However, in cases where the ROI does occur, it is likely to be the cause of such profiles.  In these cases, there is a link between the density scale radius and the anisotropy radius, indicating a causal relation between the two.  Further semi--analytic and N--Body simulations may help to clarify this potential connection.  While simulations of isolated collapsing halos are useful for trying to understand the effects of the ROI, it is well known that in realistic cosmological settings, halos form hierarchically.  Their evolution is marked by periods of gentle accretion, typically in the form of minor mergers, occasionally punctuated by major mergers.  To understand the role and relevance of the ROI in halos evolving in this way, we are presently in the process of carrying out and analyzing simulations of halos subject to controlled major and minor mergers (cf. \\citep{Poole07}).  At this early stage, we are finding that the complex dynamics at play during major mergers makes it difficult to ascertain straightforwardly whether the ROI plays a significant role during the subsequent relaxation process; a priori, we would speculate that it does not.  However, there are indications that the ROI is operational during periods of quiescent accretion, by acting upon nearly radial tidal streams associated with disrupting subhalos, which become isotropic.  We speculate that the weak tides induced by these weak structures in fact both seeds and reinforces the instability."
    },
    "0806/0806.3744_arXiv.txt": {
        "abstract": "This paper aims to show the possibility of the existence of super-massive compact objects with radii less than the Schwarzschild one, which is one of the principal consequences of the author\u2019s geodesic-invariant gravitation equations (Ann. Phys. (Berlin), 17 (2008) 28). The physical interpretation of the solutions of the equations is based on the conclusion that only an aggregate \u201cspace-time geometry + used reference frame\u201d has a physical sense. ",
        "introduction": "\\label{intro} In Einstein's theory of gravitation space-time is relative in the sense that the metric depends on the distribution of matter. However,  long before the Einstein  theory   Poincar\\'{e} showed that  geometry of space  depends also on the properties of measuring instruments.  Only an aggregate  \"geometry + measuring instruments\" has a physical meaning, verifiable by  experience. After   Minkowski  this can be also said  about  geometry of space-time. Some results of the attempt to actualize these ideas, and a generalization of  the vacuum Einstein's  gravitation  equations are considered in  \\cite{Verozub08a}. Such approach allows to consider gravitation both as a field in flat space-time  and as a space-time curvature. The equations do not contradict the existing observations data. However, the physical consequences resulting from them are radically  different from the ones of general relativity at  distances of the order of the Schwarzschild radius or less than that from a dot mass.  It is very important that this fact provides a natural explanation of  modern data of   the Universe expansion. However, this fact  leads also to another  important physical consequence, which still has no confirmation.  Observations  give evidences for the existence of supermassive compact  cold objects in  galactic centers \\cite{Genzel}.  Standard  conditions of the equilibrium of selfgravitating degenerate Fermi-gas  forbid the  existence of very massive objects. For this reason  they are usually identified with black holes.  However lack of the evidence  of the existence of an event  horizon  admits  also other explanations of the nature of such objects.  In \\cite{Verozub96} the possibility of the  existence of supermassive  equilibrium  configurations of the degenerate  Fermi gas with radii less than the Schwarzschild one  has been considered. Such objects  have no event horizon  and are an alternative to the hypothesis of the existence of black holes. It is later,  in \\cite{Verozub06a},  some observable consequences of  the existence of such object in the  Galaxy center have been considered.  However, it is seems that  a  theoretical justification of  the existence of such objects  is insufficiently convincing for observers as  it is based on conclusions resulting from  the author's  gravitation equations considered  only in a brief note  \\cite{Verozub91}. In the present paper, being based on the recent paper \\cite{Verozub08a} a simple  and a clear justification  of  possibility  of the existence of such objects is given . ",
        "conclusions": "Despite the fact that the above calculations  yield  rather qualitative than quantitative results, they show clearly that the according to equations of gravitation  (\\ref{MyVacuumEqs}), which do not contradict available observation data,  there are  stable supermassive configurations of degenerate  Fermi-gas with radiuses less than the Schwarzschild radius. Just such object can be located in the Galaxy center  \\cite{Verozub06a}."
    },
    "0806/0806.4678_arXiv.txt": {
        "abstract": "\\vskip 2cm \\hskip 6cm {\\Large \\bf Abstract} \\vskip 0.5cm We study the formation and evolution of several molecules in a collapsing interstellar cloud using a reasonably large reaction network containing more then four hundred atomic and molecular species. We employ a time dependent, spherically symmetric, hydrodynamics code to follow the hydrodynamic and chemical evolution of the collapsing cloud. The flow is assumed to be self-gravitating. We use two models to study the hydrodynamic evolution: in the first model, we inject matter into an initially low density region and in the second model, we start with a constant density cloud and let it collapse due to self-gravity. We study the evolution of the central core for both the cases. We include the grain chemistry to compute the formation of molecular hydrogen and carried out the effect of gas and grain chemistry at each time step. We follow the collapse for more than $10^{14}$s (about $3$ million years) and present the time evolution of the globally averaged abundances of various simple but biologically important molecules, such as  glycine, alanine etc. We compare our results with those obtained from observations found that for lighter molecules the agreement is generally very good. For complex molecules we tend to under predict the abundances. This indicates that other pathways could be present to form these molecules or more accurate reaction rates were needed. \\\\\\\\  Keywords: hydrodynamics; star formation; ISM; chemical evolution\\\\  PACS No.:  95.30.Lz; 97.10. Bt;98.38.-j; 98.62.Bj ",
        "introduction": "More than $125$ species of molecules have been observed in the interstellar clouds and star forming regions. Among them, over half are organic. Serious efforts have been made over the years to investigate the formation of such molecules in cool interstellar clouds in frigid conditions (Hasegawa et al., 1992; Hasegawa and Herbst, 1993; Leung et al., 1984; Prasad and Huntress, 1980a, 1980b). It is now quite certain that the most important building block, namely, the molecular hydrogen ($H_2$) and some of the other lighter molecules must be produced in the presence of grains (Gould and Salpeter, 1963; Hollenbach and Salpeter, 1971; Hollenbach et al., 1971). Several analytical and numerical works have successfully shown how the molecular hydrogen may have been produced (Biham et al., 2001). A number of results are present in the literature where hydrodynamic and chemical evolutions have been attempted simultaneously. For example, Shalabiea and Greenberg (1995) used the pseudo as well as partially real time-dependent models for the hydrodynamical evolution. In the pseudo time-dependent method, they assumed a constant density and temperature of the cloud using which the chemical evolution was computed. In their time-dependent model, they included the density and temperature variations throughout the cloud. However, in their initial approach to the time-dependent modelling, they assumed a constant temperature but allowed only the density to vary. Ceccarelli et al. (1996) used the ``inside-out\", isothermal, spherical collapse model of Shu (1977) and coupled it with a time-dependent chemical evolution code. They included the heating and the cooling processes with an emphasis on the line emission. Shematovich et al. (1997) used Zeus 2D code which included the heating and the cooling. They present 1D hydrodynamic and chemo-dynamical evolution of the proto-stellar cloud illuminated by the diffused interstellar UV radiation. They solved the equations of chemical kinetics, hydrodynamics and thermal balance simultaneously. In Lim et al. (1999) 2D numerical code was developed using the adaptive grid technique. Here, $454$ reactions among $42$ atomic and molecular chemical species were taken including the basic elements like $H$, $He$, $C$, $N$, $O$ and a representative low ionization potential metal $Na$. At each grid point, the chemical evolution was followed by a calculation of the reaction rates using the local conditions obtained from the hydrodynamical flow. They primarily concentrated on the diffused clouds and emphasized the interfaces of the interstellar media and resulting dynamical mixing. Aikawa et al. (2005) studied time-dependent evolution of Bonner-Ebert spheres by assuming clouds having a specific parameter $\\alpha$ which is the ratio of the gravitational force to the pressure force. Recently, Acharyya et al. (2005) solved the Master equations and rate equations of Biham et al. (2001) for various cloud parameters and followed the evolution of $H_2$ as a function of time. Both this work and the earlier works of Chakrabarti and Chakrabarti (2000a) employed steady state matter distribution and assumed that the density and the temperature distributions at a given radial distance do not change with time. Chakrabarti and Chakrabarti (2000a) used a large number of species and the reaction rates were taken from the UMIST data base. Some of the reaction rates which were not available in the literature were assumed to be similar to other two body reactions. Subsequently, these new and assumed reaction rates were parametrized (with reaction rates up to a thousand times smaller compared to Chakrabarti and Chakrabarti (2000a) to include the effect of the size of the reactant molecules (Chakrabarti and Chakrabarti, 2000b). It was shown that even under frigid and tenuous conditions of the interstellar media, a significant and perhaps a detectable amount of simple amino acids and even important ingredients of DNA molecule (such as adenine) may form. Ceccarelli et al. (2000) estimated the upper limit of the abundance of glycine to be about $10^{-10}$ (cooler outer cloud) to $7\\times 10^{-9}$ (hot core). Kuan et al. (2003) estimated the fractional abundance of glycine to be $2.1 \\times 10^{-10}$ for Sgr B2, $1.5 \\times 10^{-9}$ for Orion, and $2.1 \\times 10^{-10}$ for W51. These numbers are comparable to what was predicted in Chakrabarti and Chakrabarti (2000a), however, there are clearly some debate on the possible pathways for the formation of glycine with the route followed in the Chakrabarti and Chakrabarti (2000a, 2000b). Similarly, there are also some debate on whether glycine is actually observed (Hollis et al., 2003; Snyder et al., 2005). A few other relevant results which may be mentioned in passing are as follows. Tarafdar et al. (1985) presented a model of chemical and dynamical evolution of isolated, initially diffused and quiescent interstellar clouds. A semi-empirically derived dependence of the observed cloud temperatures on the visual extinction and density was used in this work. Sorrell (2001) outlined a theoretical model for the formation of the interstellar amino acids and sugars. In this model, first ultraviolet photolysis creates a high concentration of free radicals in the mantles and the heat input due to the grain-grain collision causes radicals to react chemically with another to build complex organic molecules. Bernstein et al. (2002) reported a laboratory demonstration that the complex bio-molecules like glycine, alanine etc. are naturally formed from the ultraviolet photolysis of the interstellar grains. Munoz Caro et al. (2002) report the detection of amino acids even at room temperature on an interstellar ice analogue that was irradiated with ultraviolet light in a high vacuum at 12K. Altogether sixteen amino acids were identified. The results demonstrate that a spontaneous generation of amino acids in the interstellar medium is possible, supporting the suggestion that prebiotic molecules could have been delivered to the early earth by cometary dust, meteorites or interplanetary dust particles.  In this backdrop of this observational and experimental status, we carry out our investigation by improving earlier work by first incorporating the accurate grain chemistry as elaborated in Acharyya et al. (2005) and then by actually combining the results of a time dependent hydrodynamics code with the chemical evolution code to see how the abundances vary with the grid locations. We also chose initial conditions very much different from the earlier studies. We used two different but realistic models. In one model (Model A), the cloud matter is injected through the grid boundary at a constant speed and the cloud as well as the core are allowed to form {\\it ab initio}. In the other model (Model B), the computational grid area is chosen to be the central part of a much larger spherical cloud of constant density and temperature. The rate at which the interstellar matter enters into the grid (i.e., the rate at which the larger cloud is evacuated) depends on the gravitational pull between the inner cloud within computatioal grid and the outer cloud outside of the computational grid. In following the chemical evolution, we used the available  standard reaction rates for most of the reactions, but the rates of some of the very complex molecule formation are still very much uncertain and as such we present the evolution of the mass fractions only for simpler bio-molecules. In the next section, we present the hydrodynamic equations which govern the cloud collapse. In \\S 3, we present the hydrodynamical and chemical evolution. In \\S 4, we discuss in detail the nature of the cloud models which we simulate and how the chemical evolution code is used in conjunction with the results of the hydrodynamic simulation. In \\S 5, we present the results. Finally, in \\S 6, we make concluding remarks. ",
        "conclusions": "In this paper, we presented the preliminary results on the chemical evolution inside a collapsing interstellar cloud. Our models are distinctly different from the other models used by several authors (Aikawa et al., 2005; Ceccarelli et al., 1996; Lim et al., 1999; Shalabiea and Greenberg, 1995; Shematovich et al., 1997). We do not include the heating and cooling processes while determining the dynamics of the cloud, and thus our model is not totally self-consistent. While in the outer edge of a diffused cloud the heating and cooling time scales are comparable to the infall time scales and should have been included, deep inside, the infall time scale is much shorter and cooling can be ignored. On the other hand, the outermost shell is also of very low density and thus the heating is low. The cooling is also negligible as the reactions rates are low. Thus we believe that even if the heating and cooling were included, the result would not have differed significantly.  Unlike the previous study Chakrabarti and Chakrabarti (2000a, 2000b) we have incorporated the grain chemistry of $H_2$ formation self-consistently. This major improvement gave the most realistic abundances of $H_2$ molecules in the grain and the gas phases. We find that our computed average abundances generally agree with the observed abundances (see, e.g., Allen and Knapp, 1978; Friberg et al., 1988; Irvine and Hjalmarson, 1983; Matthews et al., 1985; Ohishi et al., 1995; T\\\"olle et al., 1981) except when the molecules are complexes. We always seem to underestimate  them as compared to the observed values or `upper limits'. It is not clear at this moment how close the results of the bio-molecules are in comparison to the actual values, since the reaction rates we used are similar to the neutral-neutral rates which need not be accurate. If, for instance, the reaction rate at each step of glycine formation were higher by a factor of ten, the resulting glycine abundance would have been $10^3$ times higher, which is closer to observed claims. Given that the pathways to produce glycine may itself be different, perhaps one needs to look into laboratory experiments which simulate interstellar clouds for guidance (see, e.g., Elsila, 2007).  In future, we plan to improve the hydrodynamic model by including the angular motion and shock formation in the flow. We expect that jets and outflows would form and a part of this will fall back on the disk and the matter would be recycled. We also plan to improve the grain chemistry to include the formation of $CH_3OH$, $CO$, $NH_3$, $OH$ on the grains themselves. Thus we anticipate that the chemical abundance will be strongly affected by such recycling of matter and such incorporation of newer species on the grain surfaces.  A. Das acknowledge the support from an ISRO project.  \\newpage \\vskip 2cm \\hskip 6cm {\\Large \\bf APPENDIX} \\vskip 0.5cm \\hskip 5cm {\\bf SOLUTION PROCEDURE}\\\\ The equations are solved on a spherical grid extending from $r_{in}$ to $r_{out}$ composed of $N$ equal logarithmically spaced grids along the radial direction. The code is customized to take care of the collapsing spherical hydrodynamical flow which is strictly one dimensional. That is, no back flow is possible. Thus, to solve Eqs. (1-2) we use the first order upwind differencing method. After appropriately splitting, the Eqs. (1-2) become, $$ \\rho _{i}^{j+1}=\\rho_{i}^{j} - \\frac{dt}{{r_i}^2(r_{i+1}-r_i)} (\\rho _{i+1}^{j}v{_r}_{i+1}^{j} r_{i+1}^2-\\rho _{i}^{j}v{_r}_{i}^{j}r_{i}^2), \\eqno (A.1) $$ $$ \\rho_{i}^{j+1}v{_r}_{i}^{j+1}= \\rho_{i}^{j}v{_r}_{i}^{j} - \\frac{dt}{r_{i}^2 (r_{i+1}-r_i)} $$ $$ ({\\rho _{i+1}^{j}}^2v{{_r}_{i+1}^{j}}^2r_{i+1}^2/\\rho_{i+1}^{j}- {\\rho _{i}^{j}}^2{v{_r}_{i}^{j}}^2r_{i}^2/\\rho_{i}^{j}) $$ $$ -\\frac{dt}{(r_{i+1}-r_i)} [\\rho_{i}^{j}(\\phi_{i+1}^{j}-\\phi_{i}^{j})+ (p_{i+1}^{j}-p_{i}^{j})] \\eqno (A.2) $$ Here, $i$ denotes the index for the radial grid and $j$ denotes the index for the time. To avoid the instability in the code we chose the time step by using the Courant-Friedrichs-Lewy stability criterion, which gives, $$ {|v|\\Delta t}/{\\Delta r} \\leq 1, \\eqno (A.3) $$ i.e., $$ \\Delta t \\sim \\Delta r/|v|, $$ where, $|v|$ is the magnitude of velocity, $\\Delta t$ is the time step, and $\\Delta r$ is the grid spacing along the radial direction. We always advance the time step after ensuring that the Courant condition is satisfied. To be on the safer side, we chose time step $dt=\\Delta t/2$. Even though we use self-gravitating flow, we do not solve Poisson equation to get the potential $\\phi_i (r)$ here, since we are dealing with a spherical flow. Potential at any point is computed as a sum of two terms, one coming from the cloud itself [$\\phi_{cloud}= -GM_{cloud}(r)/r$, where, $M_{cloud}(r)$ is the mass of the cloud $r_{in}<r<r_{out}$ within the grid] and the other is due to the cloud core $\\phi_{core}=-GM_{core}/r$, where, $M_{core}$ is the mass of the core within $r<r_{in}$, the inner edge of the computational grid. The value of $\\phi (r)$ that is required during the simulation is dynamically calculated from the mass within $r$, i.e., $\\phi=-G(M_{cloud}+M_{core})/r$. Here, $M_{cloud}$ at $j$th time step is calculated by adding contributions from each spherical shell, $$ {M_{cloud}}^j=\\Sigma_i 4\\pi {r_i}^2 dr_i \\rho_{i}^{j}, \\eqno (A.4) $$ where, $i=1,2,..... N$ and the densities at the $j$th time step are used. We put $r=r_{out}$ to get the potential at the outer  boundary."
    },
    "0806/0806.2574_arXiv.txt": {
        "abstract": "{The formation of the first stars (Population III; PopIII) marks the end of the dark ages of the universe, a subject of lively scientific debate. Not (yet) accessible to direct observations, this early stage of the universe is mostly studied via theoretical calculations and numerical simulations. An indirect window is provided by integrated present day observables such as the metal abundance or the diffuse extragalactic photon fields.} {We aim to derive constraints on the properties of the PopIII and low metallicity Population II (LM PopII) stars utilizing limits on the density of the extragalactic background light (EBL), recently derived from very-high-energy (E$>$100\\,GeV; VHE) observations.} {A model calculation for the evolving EBL density produced by PopIII/LM PopII stars is presented. The model utilizes stellar population spectra (SPS) for zero and low metallicity stars and accounts for the changing emission of an aging stellar population. Emission from the dense HII regions surrounding the stars (nebula) is included. The resulting EBL density for different scenarios (metallicity, star formation rate, initial mass function) is compared to the limit on the EBL density. The potential for detecting a cut-off in HE/VHE spectra is discussed.} {Assuming a maximum contribution from PopIII/LM PopII stars to the EBL density of 5\\,nW\\,m$^{-2}$\\,s$^{-1}$ at 2\\,$\\mu$m a limit on the star formation rate (SFR) of the first stars of 0.3 to 3\\,M$_\\odot$\\,Mpc$^{-3}$\\,yr$^{-1}$ in the redshift range $7 - 14$ is derived. The limit depends on the assumed shape of the SFR and metallicity.} { The EBL can be used as a probe to investigate the properties of PopIII/LM PopII stars. Limits on the EBL density derived from VHE observations can provide constraints on the parameters of the these stars, in particular the star formation rate.} ",
        "introduction": "\\label{Sec:Introduction}  The end of the dark ages of the universe - the epoch of reionization - is a field of great interest (e.g. \\citealt{barkana:2001a,ciardi:2005a}). This epoch is associated with the formation of the first stars (Population III; PopIII)% \\footnote{See \\citet{oshea:2008a} for a discussion on naming conventions.}  (e.g. \\citealt{bromm:2004a,glover:2005a}), which are believed to start the reionization of the universe at redshift of about $z= 10 - 30$. PopIII stars form in a pristine environment, in clouds of hydrogen and helium with little or no heavy elements (primordial composition). Due to the absence of heavy elements, the cooling of such collapsing gas clouds is likely dominated by H$_2$ cooling through molecular emission lines. Numerical simulations of collapsing clouds with primordial composition predict very massive stars (100-1000\\,M$_\\odot$) with high effective temperatures  ($\\sim 10^{5}$\\,K) and short lifetimes ($\\sim10^6$\\,yrs) (e.g. \\citealt{bromm:1999a,bromm:2002a,abel:2002a}). Such hot massive stars produce copious amount of ionizing photons \\citep{schaerer:2002a} and can therefore reionize the universe. The formation of lower mass stars is also possible, if e.g. the cloud cooling is driven by hydrogen-deuterium (HD) and atomic hydrogen (H) cooling \\citep{uehara:2000a, johnson:2006a}. Other processes including turbulent fragmentation \\citep{klessen:2005a}, magnetically-regulated fragmentation \\citep{silk:2006a} and dust cooling at very high densities \\citep{omukai:2005a} could also explain stars with lower masses $< 100$\\,M$_\\odot$% \\footnote{For a more complete discussion on formation of the first stars see e.g. the 2008 updated version of \\citealt{ciardi:2005a}, astro-ph/0409018.}.  PopIII stars produce the first heavier elements, paving the way for the second generation of stars. When the star forming cloud reaches a critical metallicity ($Z_{\\mathrm{CR}} \\sim 10^{-6} - 10^{-4}$\\,Z$_{\\odot}$ e.g. \\citealt{schneider:2006b,omukai:2005a}) cooling through line emission from heavier atoms (C,O) and molecules (H$_2$O, CO, O$_2$) starts to dominate. Thereby, the second generation of stars with (likely) lower masses and \"Salpeter-like\" initial mass function start to form (Population II; PopII) (see e.g. \\citealt{bromm:2001a,schneider:2002a,schneider:2003a,bromm:2003a,schneider:2006a}). The transition from dominant PopIII to PopII star formation could already happen at early times (e.g. $z \\gg 7$), since pair-instability and core-collapse supernova explosion from PopIII stars can effectively enrich their environments with metals \\citep{schneider:2002a,scannapieco:2003a,bromm:2003a,tornatore:2007a}.  Direct observations of this early period of the universe are challenging: halo stars with extremely low metallicities have been detected in our galaxy \\citep{christlieb:2002a}, but the observation of a true PopIII star with zero-metallicity is still pending. The upcoming satellite experiment James Webb Space Telescope (JWST)\\footnote{http://ngst.gsfc.nasa.gov}, expected to be launched in 2013, with high sensitivity in the 1-10\\,$\\mu$m near-infrared (NIR) band is aiming to detect the redshifted ultraviolet (UV) to optical (O) emission from source at high redshifts $z > 10$.  Other constraints on the PopIII stars can be derived from integrated properties like e.g. the number of baryons bound in stars or the number of ionizing photons produced. If the contribution from other sources to these integrated properties are reasonably well known, the contribution from PopIII stars can be derived. This can then be compared with model calculations for different PopIII scenarios. \\citet{tumlinson:2006a} simulated the formation of PopIII stars using galactic chemical evolution models and compared the model output with the present day metallicity distribution function (MDF) of the Galaxy. They found that, while not yet formally conclusive, the MDF could best be described by a PopIII initial mass functions which includes lower mass stars.  \\cite{nagamine:2006a} used several integrated properties including the extragalactic background light (EBL; see next paragraph) to derive constraints on the cosmological star formation history of PopII stars.  In the optical to near-infrared (O-NIR) wavelength regime of the diffuse meta-galactic photon field (extragalactic background light; EBL)% \\footnote{We will use the term extragalactic background light (EBL) to denote the diffuse meta-galactic photon field in the UV to IR wavelength regime.} stars are the main contributors to the EBL density. Luminous PopIII stars can leave a distinct signature in the EBL density  \\citep{bond:1986a}. In particular, their contribution may exceed significantly the EBL density inferred from low redshift ($z < 5$) sources. Direct measurements of the EBL are difficult due to dominant foregrounds in our planetary system (zodiacal light) and the Galaxy \\citep{hauser:1998a}. Nevertheless, the discovery of such a NIR background excess (NIRBE) with high significance has been claimed by \\citet{matsumoto:2005a}, while other data showed a marginal excess (see \\citealt{hauser:2001a} for a review). The nature of this excess is still under debate. \\citet{dwek:2005c} find that it is likely a foreground artifact from zodiacal light and not of extragalactic origin. \\footnote{In addition, \\citet{mattila:2006a} argued that the claimed discontinuity in the EBL at UV-O wavelengths, which has been interpreted as a signature for the first stars, is also an artifact of foreground subtraction.} A possible PopIII origin of the NIRBE has been investigated by many authors \\citep{santos:2002a,salvaterra:2003a,dwek:2005c,madau:2005a,salvaterra:2006a,fernandez:2006a}. While \\citet{dwek:2005c} and \\citet{madau:2005a} argue that the number of stars required to produce such an excess would overproduce todays metallicity and would lead to a too high number of baryons in stars, \\citet{fernandez:2006a} (FK06) find that, if accounting for the final stage of the first stars in more detail, a PopIII origin of the NIRBE seems possible.  An indirect way of deriving constraints on the EBL comes from the measurement of very high energy (VHE) $\\gamma$-ray spectra from distant sources \\citep{stecker:1992a}. VHE $\\gamma$-rays interact with low energy photons from the EBL via pair-production \\citep{nikishov:1962a,gould:1967a}. The cross-section of the pair-production is strongly peaked, so this process leaves an energy dependent attenuation signature in the measured VHE spectra. With assumption about the source physics, upper limits on the EBL density can be derived (e.g. \\citealt{dwek:2005a,aharonian:2006:hess:ebl:nature,mazin:2007a}). \\citet{dwek:2005b} considered the effect of a high NIRBE on the spectra of distant blazars and concluded that such a high density as reported by \\citet{matsumoto:2005a} seemed unlikely. Recently, strong limits on the EBL density in the NIR have been derived (e.g. \\citealt{aharonian:2006:hess:ebl:nature,mazin:2007a}), which exclude the claimed NIRBE with high significance, and are only a factor $\\sim 2$ above the lower limits derived from source counts \\citep{madau:2000a}.  In this paper these limits on the diffuse EBL density are used to derive constraints on the properties of the PopIII/LM PopII stars. Results from a detailed model calculation of the EBL for different PopIII/LM PopII star scenarios are compared with recent limits on the EBL density. Our model accounts for the time evolution of the emissivity of a stellar population, which, for the case of low mass stars with long lifetimes, has profound implications for the resulting EBL.  The paper is organized as follows: In Sec.~\\ref{Sec:EBL_Model} the model calculations for the EBL density from  PopIII/LM PopII stars are described. In Sec.~\\ref{Sec:Constrains} the resulting EBL density for different sets of  PopIII/LM PopII star parameters is calculated and compared with recent limits. Limits on cosmological star formation rate (SFR) are derived and the detectability of a cut-off in high energy spectra resulting from EBL attenuation are discussed. In Sec.~\\ref{Sec:Discussion} the derived limits are compared with previous results and the consequences for the PopIII/LM PopII star properties are discussed. We summarize our results in Sec.~\\ref{Sec:Summary}.  Throughout this paper flat Friedman cosmology is adopted with $\\Omega_0 = 0.25$ , $\\Omega_\\Lambda=0.75$  and a Hubble constant of $H_0=70$~km s$^{-1}$ Mpc$^{-1}$. ",
        "conclusions": "\\label{Sec:Summary}  We investigate how limits on an integrated present-day observable, the EBL, can be used to constrain the parameters of the early stars. A detailed model for the PopIII/LM PopII star emission from a large range of different scenarios is used to calculate the evolving EBL from these stars, taking into account the time evolution of the emissivity and the emission from reprocessed ionizing photons (nebula). Recent limits on the EBL density derived by \\citet{aharonian:2006:hess:ebl:nature} from the detection of hard VHE $\\gamma$-ray spectra from distant sources together with lower limits from source counts \\citep{madau:2000a,totani:2001a} suggest a maximum PopIII EBL contribution of $\\sim 5 $\\,nW\\,m$^{-2}$\\,sr$^{-1}$ at $1-2\\,\\mu$m. Comparing this contribution with our model calculations, a limit on the co-moving SFR of PopIII stars of 0.3 to 3\\,M$_\\odot$\\,Mpc$^{-3}$\\,yr$^{-1}$  is derived for the redshift range $7 - 14$. This limit depends on the redshift, on the exact shape of the SFR and on the assumed scenario for the early star formation: if the early star formation is dominated by second generation stars with low metallicity, the limit is factor two lower than in the case of zero metallicity stars. The SFR limit directly scales with the assumed PopIII EBL contribution, e.g., if the EBL limit is lowered by a factor 2 the corresponding SFR limit is also lowered by the same factor.  The SFR at redshift $> 5$ is difficult to access via direct observations. A few measurements and limits exist, generally favoring a lower SFR in the range of 10$^{-3}$ to 10$^{-2}$\\,M$_\\odot$\\,Mpc$^{-3}$\\,yr$^{-1}$, but the spread and uncertainties are large (Fig.~\\ref{Fig:SFR_results}). Recent measurements of the SFR in the redshift range $z = 3 - 6$ favor a flat (or even increasing) SFR of the order of 0.1\\,M$_\\odot$\\,Mpc$^{-3}$ \\citep{yuksel:2008a,faucher-giguere:2008a:astro-ph}, which is in the range of our best limit (0.3\\,M$_\\odot$\\,Mpc$^{-3}$\\,yr$^{-1}$ at $z = 7$ for LM, SFR(7, 4, 0)). Stringent constraints on the properties of the first stars come from reionization studies, which combine complex semi-analytical modeling with limits on the reionization history. Predictions for the peak SFR for PopIII and (LM) PopII  are in the range of $10^{-4} - 10^{-3}$\\,M$_\\odot$\\,Mpc$^{-3}$\\,yr$^{-1}$ and $\\sim 10^{-2} - 10^{-1}$\\,M$_\\odot$\\,Mpc$^{-3}$\\,yr$^{-1}$, respectively \\citep{choudhury:2006a, greif:2006a}.\\footnote{A higher SFR is also possible, see Sect.~\\ref{Sec:Discussion}.} While for PopIII stars the limits derived in this paper are 1 to 5 orders of magnitude above the SFR expected from the best fit models from these studies, for (LM) PopII stars the limits are close (factor 3 to 10) to these predictions.  Pair-creation of VHE photons from distant sources ($z > 1$) with the low energy photons from the EBL results in a sharp cut-off in energy spectra $\\gtrsim 30$\\,GeV, which should be detected by the FERMI experiment. To derive constraints on the PopIII/LM PopII stars from the detection of such a cut-off is challenging since (a) the photon statistics will likely be low, (b) attenuation from the PopIII/LM PopII EBL competes with the attenuation due to the EBL from PopII stars, which is likely the dominant contribution to the total attenuation, and (c) the general problem to discriminate between source intrinsic effects and attenuation from the EBL.  Constraints on the EBL can provide additional insides in the star formation processes of the early universe. Though the limits are not (yet) strongly constraining, they provide an independent probe for the star formation at redshift $z > 5$. With the current limits on the EBL in the NIR it is not possible to distinguish between different PopIII IMFs or metallicity scenarios. In the future, the Cherenkov Telescope Array (CTA)\\footnote{http://www.cta-observatory.org/} will provide sensitive measurements in the $\\sim$20\\,GeV to 100\\,TeV energy range, which will result in strong constraints on the EBL in a wide wavelength range. Together with direct detections and deep source counts from upcoming satellite and ground-based telescopes this will enable to resolve many of the contributors to the EBL and thereby tighten the limits on the PopIII/LM PopII stars properties derived from the EBL."
    },
    "0806/0806.0011.txt": {
        "abstract": "A technical and metho\\-do\\-logical comparison of liquid noble gas experiments is presented and the low energy physics application of double phase noble gas detectors in direct Dark Matter investigations is discussed. ",
        "introduction": "\\label{Int}  Liquid noble gas detectors have been proposed and used in different fields of research for many years \\cite{ref1,Bar}. The use of liquid xenon as a pure scintillator in direct Dark Matter particle investigations dates back to 1990 with the pioneering work by DAMA/LXe  \\cite{dam1} and later on by the UK coll. \\cite{zep0}. Other authors have also expressed an interest in using liquid argon or xenon, see e.g. ref. \\cite{Cli1}.  There has recently been a great deal of interest in double phase noble gas detectors, i.e. detectors exploiting the primary scintillation in the liquid phase and the secondary proportional scintillation in the gas phase, as their use might be possible in one approach to direct Dark Matter investigations, the detection of WIMP (Weakly Interacting Massive Particles) candidates \\cite{GW} in the case of their interaction would  only produce recoil nuclei.  Different kinds of WIMP interaction have been studied: interactions with the target nuclei giving rise to elastic scattering (with spin independent and/or spin dependent coupling), inelastic scattering, preferred inelastic interaction (with the Dark Matter particle that goes into an excited state instead of the nucleus) \\cite{SW},  the impact of a partial electromagnetic contribution due to ionisation and excitation of bound atomic electrons induced by the presence of the recoiling atomic nucleus (Migdal effect) \\cite{Mig}, etc.. It has also been indicated that electromagnetic interaction with atomic electrons may occur, when WIMP interaction with the nucleus is inhibited \\cite{ele}.  Many other particles and types of interaction have been considered and Dark Matter could even be of a multicomponent nature. Varying lines of research also suggest that different kinds of interaction require suitable levels of sensitivity for their detection. For example a  large number of papers on axion-like particles have indicated that although they share similar phenomenology with ordinary matter like the axion, they have significantly different mass and coupling constants. Hence in this case, detection is based on the total conversion of the absorbed bosonic mass into electromagnetic radiation \\cite{bos}. In the framework of warm dark matter, light dark matter candidates have also been considered \\cite{ldm}, inelastic scattering channels on the electron or on the nucleus have been studied for axino, sterile neutrinos and even MeV-scale LSP in Susy theories. After inelastic interaction, a lighter particle is produced and the target recoil energy is the detectable quantity. The target can be an electron or a nucleus, so the detectable quantity is due to a different mechanism and has different features.     On the other hand, experiments exploiting the dual phase noble gas detection considered in this paper, such as those based on the bolometer + ionisation/scin\\-til\\-lation detection method, only focus on WIMPs and one type of WIMP interaction: the elastic scattering process on the target nuclei. However, when the target only has isotopes with even nuclei, such as is the case for natural argon and neon (the latter is practically even as $^{21}$Ne is only 0.27\\%) they are not sensitive to spin dependent WIMP-nucleus inte\\-ra\\-ctions at all\\footnote{Even for nuclei sensitive to spin dependent interaction different sensitivities are expected among odd-nuclei having an unpaired proton (for example, $^{23}$Na and $^{127}$I) and odd-nuclei having an unpaired neutron (for example, the odd Xe and Te isotopes and $^{73}$Ge). This is due to uncertainties concerning the choice of  the nuclear potential,  the effective WIMP-nucleon coupling strengths and form and spin factors.}. Hence nuclear recoil is the only process that can be considered in these experiments and the only feature that might be detected. As a result, an exclusion plot in the plane cross section versus WIMP mass is usually produced in a fixed single model framework of experimental and phenomenological parameters without taking into account any uncertainty or alternative choices, and assuming that ideal results have been obtained using the several applied procedures. This is a significant limitation to any investigation that wishes to explore Dark Matter effectively. Due to these and other arguments no model independent comparison can be drawn from results obtained using different target materials and/or different experimental approaches.  Three gases have been discussed with regard to two phase detectors: xenon (XENON, ZEPLIN), argon (WARP) and neon (SIGN) and recently measurements have been published for the ZEPLIN-II, ZEPLIN-III, XENON10 and WARP prototypes.  Although these state of art prototypes are at the R\\&D stage, successors on ton scale are already being discussed (ArDM, ELIXIR, LUX).  This paper looks at the low energy experimental application of liquid noble gases and compares features of the apparata such as design and  performance, purity and radiopurity, trigger, calibration, data reduction, rejections and analyses. The data  published up to now and examined here  correspond to an exposure of 136 kg $\\times$ day for XENON10, 225 kg $\\times$ day for ZEPLIN-II, 847 kg $\\times$ day for ZEPLIN-III and 96.5  kg $\\times$ day for WARP. This paper compares the performance of apparata as published by the authors or presented at international conferences. It is worth underlining that technical aspects are extremely important as they can affect experimental results and should be studied accurately before high experimental sensitivity is claimed.\\\\ It is hoped that the discussion on double phase detectors will prove useful for commissioned  projects, projects under construction and future projects.  In order to avoid any misunderstanding, any published sentences quoted in this paper are written in italics. ",
        "conclusions": "\\label{Con}  This paper has analysed and compared the very low energy application of double phase noble gas detectors. The main technical aspects of the existing experimental applications have been discussed and some implications have been outlined. The main topics to be addressed in further research and developments have also been presented."
    },
    "0806/0806.4132_arXiv.txt": {
        "abstract": "{} {We analysed eight XMM-Newton observations toward the Small Magellanic Cloud (SMC), performed between October 2006 and June 2007, to investigate high mass X-ray binary systems.} {We produced images from the European Photon Imaging Cameras (EPIC) and extracted X-ray spectra and light curves in different energy bands from sources which yielded a sufficiently high number of counts for a detailed temporal and spectral analysis. To search for periodicity we applied Fourier transformations and folding techniques and determined pulse periods using a Bayesian approach. To identify optical counterparts we produced X-ray source lists for each observation using maximum likelihood source detection techniques and correlated them with optical catalogues. The correlations were also used for astrometric boresight corrections of the X-ray source positions.} {We found new X-ray binary pulsars with periods of 202~s (XMMU\\,J005929.0-723703), 342~s (XMMU\\,J005403.8-722632), 645~s (XMMU\\,J005535.2-722906) and 325~s (XMMU\\,J005252.1-721715), in the latter case confirming the independent discovery in Chandra data. In addition we detected sixteen known Be/X-ray binary pulsars and six ROSAT-classified candidate high mass X-ray binaries. From one of the candidates, RX\\,J0058.2-7231, we discovered X-ray pulsations with a period of 291~s which makes it the likely counterpart of XTE\\,J0051-727. From the known pulsars, we revise the pulse period of CXOU\\,J010206.6-714115 to 967 s, and we detected the 18.37~s pulsar XTE\\,J0055-727 (=~XMM\\,J004911.4-724939) in outburst, which allowed us to localise the source. The pulse profiles of the X-ray pulsars show a large variety of shapes from smooth to highly structured patterns and differing energy dependence. For all the candidate high mass X-ray binaries optical counterparts can be identified with magnitudes and colours consistent with Be stars. Twenty of the Be/X-ray binaries were detected with X-ray luminosities in the range 1.5\\ergs{35} - 5.5\\ergs{36}. The majority of the spectra is well represented by an absorbed power-law with an average power-law index of 0.93. The absorption (in addition to the Galactic foreground value) varies over a wide range between a few \\ohcm{20} and several \\ohcm{22}. An overall correlation of the absorption with the total SMC \\Hone\\ column density suggests that the absorption seen in the X-ray spectra is often largely caused by interstellar gas.} {} ",
        "introduction": "Together with the Milky Way, the Small Magellanic Cloud (SMC) is the galaxy with the highest number of known high mass X-ray binaries (HMXBs). In these systems a compact object is accreting mass from an early type companion star. With only one exception (the supergiant system SMC\\,X-1), the vast majority of HMXBs in the SMC comprises systems with a Be star as mass donor and a neutron star as compact object that is usually recognised as an X-ray pulsar. Depending on the eccentricity of the binary orbit, the mass accretion rate can be enhanced around the time of periastron passage when the neutron star approaches the circumstellar disc of the Be star \\citep[see, e.g.,][]{2001A&A...377..161O}. These outbursts usually last a few days. Longer outbursts with durations of weeks (type II) can be caused by a drastic expansion of the disc around the Be star.  Many of the Be/X-ray binaries in the SMC were discovered during X-ray outburst. Then pulsations could be found as periodic modulation of the X-ray flux. During quiescence the sources were often not detected at all. It is not clear how many of these X-ray transients we are still missing, but present day X-ray observations with Chandra, the Rossi X-ray Timing Explorer (RXTE) and \\xmm\\ keep finding new Be/X-ray binaries in the SMC with a rate of a few per year. Currently, in the SMC nearly 50 Be/X-ray binary pulsars are known \\citep{2004A&A...414..667H,2005MNRAS.356..502C}. In addition, more than two dozen Be/X-ray binary candidates are discussed in the literature which are either classified as such from their X-ray properties and/or are optically identified with a Be star counterpart. In these cases, however, uncertain X-ray positions and missing X-ray pulsations make the cases less clear.  \\begin{table*} \\caption[]{XMM-Newton EPIC observations of SMC fields in 2006/2007.} \\begin{tabular}{lllllrrc} \\hline\\hline\\noalign{\\smallskip} \\multicolumn{1}{c}{Observation} & \\multicolumn{2}{c}{Pointing direction} & \\multicolumn{1}{c}{Sat.} & \\multicolumn{1}{c}{EPIC$^{(a)}$} & \\multicolumn{1}{c}{Start time (UT)} & \\multicolumn{1}{c}{End time (UT)}\\\\  \\multicolumn{1}{c}{ID} & \\multicolumn{1}{c}{R.A.} & \\multicolumn{1}{c}{Dec.} & \\multicolumn{1}{c}{Rev.} & \\multicolumn{1}{c}{Instrument} & \\multicolumn{1}{c}{} & \\multicolumn{1}{c}{} \\\\  \\multicolumn{1}{c}{} & \\multicolumn{2}{c}{(J2000.0)} & \\multicolumn{1}{c}{} & \\multicolumn{1}{c}{configuration} & \\multicolumn{1}{c}{} & \\multicolumn{1}{c}{} \\\\ \\noalign{\\smallskip}\\hline\\noalign{\\smallskip} 0404680101 & 00 47 36.0 & -73 08 24 & 1249 & PN FF thin      & 2006-10-05 00:44:47 & 2006-10-05 06:51:09 \\\\ &            &           &      & M1/M2 FF medium &            00:22:05 &            06:50:54 \\\\ 0404680201 & 00 52 26.4 & -72 52 12 & 1263 & PN FF thin      & 2006-11-01 01:18:41 & 2006-11-01 09:58:46 \\\\ &            &           &      & M1/M2 FF medium &            00:55:59 &            09:58:31 \\\\ 0403970301 & 00 47 39.4 & -72 59 31 & 1329 & PN EFF thin     & 2007-03-12 21:02:47 & 2007-03-13 06:53:05 \\\\ &            &           &      & M1/M2 FF thin   &            20:01:50 &            06:52:50 \\\\ 0404680301 & 00 51 00.7 & -73 24 17 & 1344 & PN FF thin      & 2007-04-11 20:00:39 & 2007-04-12 02:15:32 \\\\ &            &           &      & M1/M2 FF medium &            19:37:57 &            02:15:17 \\\\ 0404680501 & 01 07 42.3 & -72 30 11 & 1344 & PN FF thin      & 2007-04-12 03:29:37 & 2007-04-12 09:44:42 \\\\ &            &           &      & M1/M2 FF medium &            03:06:55 &            09:44:27 \\\\ 0501470101 & 00 59 41.8 & -71:38:15 & 1371 & PN FF thin      & 2007-06-04 09:22:05 & 2007-06-04 18:20:14 \\\\ &            &           &      & M1/M2 FF thin   &            08:59:23 &            18:19:59 \\\\ 0500980201 & 01 00 00.0 & -72 27 00 & 1372 & PN FF thin      & 2007-06-06 09:14:26 & 2007-06-06 16:52:51 \\\\ &            &           &      & M1/M2 FF medium &            08:51:44 &            16:52:36 \\\\ 0500980101 & 00 53 02.4 & -72 26 17 & 1380 & PN FF thin      & 2007-06-23 06:13:53 & 2007-06-23 13:03:10 \\\\ &            &           &      & M1/M2 FF medium &            05:51:11 &            13:02:55 \\\\ \\noalign{\\smallskip}\\hline\\noalign{\\smallskip} \\end{tabular}  $^{(a)}$ FF: full frame CCD readout mode with 73~ms frame time for PN and 2.6~s for MOS; EFF: PN `extended FF' with 200~ms frame time; thin and medium optical blocking filters. \\label{tab-obs} \\end{table*}  We used the European Photon Imaging Cameras (EPIC) on \\xmm\\ to investigate HMXBs in the SMC. Here we report on results from eight observations - seven from our own dedicated programs to investigate candidate HMXBs and supersoft X-ray sources and one from the public archive - which cover six candidate Be/X-ray binaries and sixteen known Be/X-ray binary pulsars. We present a temporal and spectral analysis of the majority of these objects and additional newly discovered Be/X-ray binary pulsars in the SMC. ",
        "conclusions": "We detected a large sample of twenty-six Be/X-ray binaries in eight \\xmm\\ observations of the SMC between Oct. 2006 and June 2007. After astrometric boresight corrections the accurate X-ray positions allowed us either to confirm previously proposed optical counterparts or, in cases of the newly discovered systems, to identify their counterparts. The properties of all optical counterparts (like brightness, colours and long-term temporal behaviour) are consistent with Be/X-ray binaries. No candidate for a supergiant HMXB was found.  Twenty of the sources were observed with luminosities above $\\sim$\\oergs{35}, providing X-ray spectra and light curves for a detailed, systematic study. Most of the X-ray spectra can be modelled by an absorbed power-law, typical for HMXBs. The average power-law index is 0.93 with 90\\% of the values between 0.71 and 1.27, which is in agreement with the peak of the distribution at 1.0 found by \\citet{2004A&A...414..667H} for a smaller sample. \\citet{2007MNRAS.376..759M} report harder spectra from four Be/X-ray binary pulsars detected in the Chandra SMC wing survey and discuss possibilities for systematic differences between pulsars in the wing and the bar of the SMC. We see a general trend of harder power-law spectra when the sources are brighter (e.g. XMMU\\,J004723.7-731226, Sect.~\\ref{sect-lxharda}; CXOU\\,J010712.6-723533, Sect.~\\ref{sect-lxhardb}) which suggests that the harder spectra in the wing pulsars might at least partly be due to a brightness selection effect. From the short ($\\sim$10~ks) Chandra observations, spectra could only be derived when the sources were bright.  The highest luminosities were reached by the 18.37~s pulsar XMMU\\,J004911.4-724939 and the new 202~s pulsar XMMU\\,J005929.0-723703 with 5.5\\ergs{36} and 4.5\\ergs{36}, respectively, while all other sources with EPIC spectra were observed with luminosities between \\oergs{35} and \\oergs{36}. Six Be/X-ray binaries were detected only as faint sources (for four of them we estimate their luminosity from the count rates, the other two are located at the edge of the field of view) and about 5-10 Be/X-ray binaries (uncertain positions and unclear classifications make the numbers uncertain) were not detected in our eight fields. I.e. we now know about 31-36 Be/X-ray binaries in the eight fields from which 26 (about 70-85\\%) are detected above a luminosity of 4\\ergs{34} (the weakest detected sources have count rates down to 1\\ct{-3} for EPIC-PN which converts to 4\\ergs{33} for a typical HMXB power-law spectrum).  For twelve of the Be/X-ray binaries covered by our \\xmm\\ observations, orbital ephemeris obtained from RXTE X-ray data are available \\citep{2008arXiv0802.2118G}. All our clear detections are from orbital phases near zero (which defines maximum X-ray intensity as seen by RXTE). Consistently, three pulsars observed around phase 0.4 were found to be either faint or not detected. The 74.7~s pulsar AX\\,J0049-729, the 59~s pulsar XTE\\,J0055-724 and the 82.4~s pulsar XTE\\,J0052-725 were relatively faint during the \\xmm\\ observations close to phase 0, but the uncertainties in their ephemeris do not allow an exact determination of the orbital phase. Therefore, it is not clear if the outbursts were missed by the \\xmm\\ observations or if they did not happen. Overall, the luminosity states in which we found the Be/X-ray binaries during the \\xmm\\ observations are all consistent with the orbital variations determined by RXTE. This also suggests, that we did not observe any type II outburst with \\xmm.  Systematic deviations from the power-law model are seen in the spectra of a few of our investigated Be/X-ray binaries (Fig.~\\ref{fig-spectra}). In particular CXOU\\,J010712.6-723533 and SMC\\,X-3 yield bad power-law fits (in terms of reduced $\\chi^2$ as can be seen from Table~\\ref{tab-spectra}). Additional cases might be RX\\,J0050.8-7316 and XMMU\\,J005252.1-721715. However, spectra with better statistical quality are required to prove this. Expanding the model with a blackbody component yields acceptable fits for  CXOU\\,J010712.6-723533 and SMC\\,X-3. The inferred blackbody temperatures are around 1.2 keV and the sizes of the emitting areas are relatively small (radius $\\sim$0.7 km). These spectral characteristics are similar to those of RX\\,J0146.9+6121 and X\\,Persei \\citep{2006A&A...455..283L,2007A&A...474..137L}. These authors interprete the blackbody component as emission from the hot polar caps of the accreting neutron star \\citep[for a discussion of the origin of soft emission excesses see][]{2004ApJ...614..881H}.  We discovered X-ray pulsations from four new transient Be/X-ray binaries \\citep[for the 325~s pulsar XMMU\\,J005252.1-721715 an independent discovery from Chandra data was reported by ][]{2008arXiv0803.3941C} and found pulsations from two previously known Be/X-ray binaries. The pulse periods range between 202~s and 967~s, the latter being the second longest known from SMC pulsars. The high statistical quality of the EPIC data allowed us to investigate pulse profiles in several different energy bands, revealing a large variety of pulse shapes and different energy dependence.  All of our investigated pulsars with pulse periods longer than 150~s and known period history over at least a few years show secular spin-up of the neutron star. They all exhibit more or less regular outburst activity as seen by RXTE \\citep{2008arXiv0802.2118G}. At shorter spin periods at least two of the pulsars show long-term spin-down: the 138~s pulsar CXOU\\,J005323.8-722715, which was detected only twice by RXTE at high brightness level \\citep{2008arXiv0802.2118G} and the 7.78~s pulsar SMC~X-3. The periods determined for the 18.37~s pulsar XMMU\\,J004911.4-724939 are consistent with no change. This fits in the general picture of the spin evolution of accretion powered pulsars, in which the neutron star is spun-up by accretion until an equilibrium period is reached \\citep[see, e.g. equations 8 and 9 in ][]{1997ApJS..113..367B}. The fact that the majority of long-period pulsars in the SMC show secular spin-up, suggests however a much longer timescale for spinning up the neutron star, as was also seen for HMXB pulsars in the Milky Way. High resolution BATSE measurements of the long-term spin evolution of HMXB pulsars have shown that the small long-term spin-up rates are a consequence of frequent transitions between spin-up and spin-down which can occur on timescales of less than 10 days \\citep{1997ApJS..113..367B}. To confirm such a behaviour for SMC Be/X-ray binary pulsars would require a monitoring program with \\xmm\\ with frequent observations for several weeks.  With the new discoveries, we now know of several pairs of SMC pulsars with very similar periods. In particular two 202~s pulsars exist with an angular distance of only 13.8\\arcmin. RX\\,J0059.3-7223 was discovered with a period of 201.9$\\pm$0.5~s in October 2000, which decreased to 200.5$\\pm$0.3~s, and XMMU\\,J005929.0-723703 was discovered with 202.52$\\pm$0.02~s. A similar case is XMMU\\,J005403.8-722632, newly discovered in this work with a period of 341.87$\\pm$0.15~s, and SAX\\,J0103.2-7209 with the earliest measured period of 348.9$\\pm$0.1~s \\citep[in ASCA data from May 1996; ][]{1998IAUC.7009....3Y}, which decreased to 341.2$\\pm$0.5~s in October 2000 \\citep{2004A&A...414..667H}. These two pulsars are located 45.2\\arcmin\\ apart. Due to the overlapping ranges of spin periods which evolve with time and the projected proximity of many Be/X-ray binaries in the SMC some of these pulsar pairs my not be differentiated by non-imaging instruments which needs to be considered when interpreting the spin period evolution.  The SMC absorption component varies over a wide range between a few \\ohcm{20} and several \\ohcm{22}. It is evident, that the sources in field 0500980101 show similar values of low absorption, while those of field 0404680101, near the emission nebula N19, show all very high absorption. For further investigation we list the total \\Hone\\ column density of the SMC in the direction of each source in the last column of Table~\\ref{tab-spectra}. In Fig.~\\ref{fig-absorption} the X-ray measured column density \\nh\\ is plotted as functions of the B-V colour index of its optical counterpart (from MCPS in Table~\\ref{tab-ids}) and the total SMC column density as derived from \\Hone\\ measurements. In both cases a correlation is indicated, although with considerable scatter in \\nh. The scatter is expected as 1) the sources are located at different depth in the SMC and 2) a large part of the absorption can originate locally in the Be/X-ray binary systems \\citep[suggested by strong variations of \\nh\\ over longterm timescales as seen, e.g., from RX\\,J0103.6-7201; ][]{2005A&A...438..211H}. The general increase of the X-ray measured \\nh\\ with the total SMC column density suggests that a significant fraction of the X-ray absorption seen in the X-ray spectra of Be/X-ray binaries arises in the interstellar medium of the SMC.  \\begin{figure} \\resizebox{0.98\\hsize}{!}{\\includegraphics[angle=-90,clip=]{mcps-NH.ps}} \\resizebox{0.98\\hsize}{!}{\\includegraphics[angle=-90,clip=]{NH-HI.ps}} \\caption{Equivalent hydrogen column density derived from the X-ray spectra as a function of the optical colour index B-V (top) and as a function of the total \\Hone\\ column density of the SMC. The line in the bottom panel marks \\nh\\ = \\Hone.} \\label{fig-absorption} \\end{figure}"
    },
    "0806/0806.1521_arXiv.txt": {
        "abstract": "We examine a physical process that leads to the efficient formation of gas giant planets around intermediate mass stars. In the gaseous protoplanetary disks surrounding rapidly-accreting intermediate-mass stars we show that the midplane temperature (heated primarily by turbulent dissipation) can reach $\\gtrsim$1000~K out to 1 AU. Thermal ionization of this hot gas couples the disk to the magnetic field, allowing the magneto-rotational instability (MRI) to generate turbulence and transport angular momentum.  Further from the central star the ionization fraction decreases, decoupling the disk from the magnetic field and reducing the efficiency of angular momentum transport.  As the disk evolves towards a quasi-steady state, a local maximum in the surface density and in the midplane pressure both develop at the inner edge of the MRI-dead zone, trapping inwardly migrating solid bodies. Small particles accumulate and coagulate into planetesimals which grow rapidly until they reach isolation mass.  In contrast to the situation around solar type stars, we show that the isolation mass for cores at this critical radius around the more massive stars is large enough to promote the accretion of significant amounts of gas prior to disk depletion.  Through this process, we anticipate a prolific production of gas giants at $\\sim 1$ AU around intermediate-mass stars. ",
        "introduction": "The discovery of a plethora of extra-solar planets around solar-type main sequence stars has established that planet formation must be a common process, not a peculiarity of our own solar system. As observational techniques for planetary detection have become more sophisticated, the discovery domain has expanded to include host stars with a wide range of masses. While on the main sequence, intermediate mass stars (stars with $1.5 M_{\\sun} \\lesssim M_* \\lesssim 3 M_{\\sun}$) make poor radial velocity (RV) survey candidates as they have few spectral lines which also tend to be rotationally broadened (\\citet{Griffin.etal.2000} but see \\citet{Galland.etal.2006}). However, once these stars evolve off the main sequence, their relatively cool and slowly rotating outer layers make them more suitable candidates for high precision spectroscopic studies. Recent RV surveys targeting evolved intermediate mass stars suggest that they differ from solar-type stars as planetary hosts in at least two respects. First, the total frequency of giant planets (with periods less than a few years) appears be higher around intermediate mass stars. Second, the planets have different statistical properties.  Their semi-major axis distribution is concentrated at 1-2 AU and there is an apparent lack of short-period (days to months) planets, despite observational selection effects favoring their discovery \\citep{LovisMayor.2007}.  In this paper, we propose a common explanation for prolific gas giant formation with semi-major axes comparable to 1~AU and for the rarity of close-in planets around intermediate-mass stars.  As it is unlikely that all planets within 1 AU have been engulfed or had their orbits disrupted by the current expanded envelope of the host stars \\citep{Johnson.etal.2007}, we attribute both properties to the formation and early evolutionary processes rather than to post-main-sequence evolution.  We begin by examining the physical properties of circumstellar disks which may affect the probability of forming giant planets. In the core-accretion model of planet formation (cf. \\citet{Bodenheimer.Pollack.1986}), the emergence of Jupiter-like gas giants requires that a population of solid cores form within a gaseous protoplanetary disk. These cores grow through cohesive collisions with planetesimals, with a growth rate determined by the velocity dispersion of the planetesimal swarm. The magnitude of this velocity dispersion is set by a balance between excitation by gravitational perturbations and damping by gas drag. In the gas-rich environment of typical protostellar disk, gas drag dominates so that field planetesimals only attain relatively small equilibrium velocity dispersion. As a result the most massive protoplanetary embryos can access only those building blocks within their gravitational feeding zones \\citep{Kokubo.Ida.1998}.  When these embryos have collected all the planetesimals within about five times their Roche radius on either side of their orbits, their growth stalls. This maximum embryo mass, a function of planetesimal surface density and distance from the central star, is referred to as the embryo's isolation mass ($\\Miso$). Gas giants can only form if the embryos' $\\Miso$ is sufficiently large for the cores to begin accreting gas prior to the depletion of their nascent disks \\citep{IdaLin2004}.  Although the gravity of lunar-mass embryos is adequate to accrete disk gas with temperature $<10^3$ K, \\emph{efficient} dynamical gas accretion is only possible for cores with masses greater than some critical value ($M_{\\rm crit}$). In a minimum mass solar nebula \\citep{Hayashi.1981} with an interstellar grain size distribution, $M_{\\rm crit} \\sim 10 M_\\oplus$ at a semi-major axis $a\\sim 5$~AU \\citep{Pollack.etal.1996}, although this critical mass decreases both with lowered grain opacity \\citep{Ikoma.etal.2000, Hubickyj.etal.2005} and with increased density of the ambient gas \\citep{Bodenheimer.Pollack.1986, Papaloizou.Terquem.1999}. Gas giant formation therefore requires that the heavy-elements in the disk can be efficiently assembled into massive cores with mass greater than $M_{\\rm crit}$. In order to understand the spatial distribution of the gas giant planets we must understand how the building blocks of these cores migrate and are retained in gaseous disks.  In protoplanetary disks solid retention first  becomes an issue once grains grow beyond a few cm in size. In most regions of protostellar disks, the midplane pressure ($P_{\\rm mid}$) decreases with distance from the central star ($r$) so that the gas is slightly pressure supported, resulting in a sub-Keplerian azimuthal velocity. Grains larger than a few cm are decoupled from the gas and move at Keplerian speeds.  Consequently, grains typically experience head winds and undergo orbital decay \\citep{Weidenschilling1977}.  However if $P_{\\rm mid}$ does not monotonically decrease with $r$ then immediately interior to a local pressure maximum the gas attains super-Keplerian velocities. This motion introduces a tail wind on the decoupled grains and causes them to drift outwards towards local pressure maxima \\citep{Bryden.etal.2000, HaghihipourBoss2003}.  Solid retention again becomes an issue once planetesimals grow into earth-mass embryos and tidal interactions with the gaseous disk become important.  Before embryos are sufficiently massive to open up gaps in the disks \\citep{LinPapaloizou1986}, they can exchange angular momentum with the gas via their Lindblad and co-rotation resonances \\citep{Goldreich.Tremaine.1980}.  A geometric bias causes an imbalance between the Lindblad resonances which generally leads to a loss of angular momentum and orbital decay for the embryos \\citep{Ward.1986, Ward1997}. However, embryos will gain angular momentum through their co-rotation resonances if there is a positive $P_{\\rm mid}$ gradient \\citep{Tanaka.etal.2002, Masset.etal.2006}. Numerical models which take into account these physical effects have reproduced the observed $M_p-a$ distribution around solar type stars \\citep{IdaLin2008}.  Several physical processes can lead to local maxima in $P_{\\rm mid}$. Various authors have explored the potential accumulation of grains at transient pressure maxima formed by turbulent fluctuations \\citep{Johansen.etal.2006b,Fromang.Nelson.2005} or spiral waves \\citep{Rice.etal.2006}. These mechanisms, while likely extremely important for forming planetesimals at a large range of radii, are still quite ``leaky'' as a significant fraction of the solid material simply undergoes a slightly slower random walk towards the central star. However, longer lived pressure maxima may also exist due to large scale changes in the disk viscosity \\citep{Kretke.Lin.2007}.  We expect radial variations in viscosity if turbulence caused by the magneto-rotational instability (MRI; \\citet{BalbusHawley1991}) is the primary mechanism for transporting angular momentum. These variations result from changes in the ionization fraction at different radii in the disk since free electrons are needed to couple the gas to the magnetic field. The disk is thermally ionized in the hot inner regions, but further out stellar x-rays and diffuse cosmic rays ionize only the surface layers, resulting in a viscously active turbulent surface sandwiching an inactive ``dead zone'' \\citep{Gammie1996}. At the critical radius marking the inner edge of the dead zone ($\\acrit$), the effective viscosity decreases with increasing distance from the central star. In a quasi-steady state situation (expected to develop rapidly in the inner regions of the disk) this decrease in viscosity leads to a local increase in the magnitude of $\\Sigma_g$ and hence of $P_{\\rm mid}$ with radius.  This disk structure provides a promising barrier to the orbital decay of both boulders and embryos. The radial location of $\\acrit$ depends on the stellar mass and the mass accretion rate which we argue explains the observed differences between the statistical distributions of planets around solar-type stars and around intermediate mass stars.  In this paper, we present a model for the formation of planets at the inner edge of the dead zone and argue why this process is more relevant for intermediate mass than for solar mass stars. In $\\S$\\ref{sec:model} we describe our quantitative numerical model for the evolution of solids in the disk, based upon the work of \\citet{Garaud.2007}. An important aspect of this model is the location of $\\acrit$ (the inner edge of the dead zone)  which we derive in $\\S$\\ref{sec:acrit} as a function of stellar mass and mass accretion rate. In $\\S$\\ref{sec:results} we present the model results for a $2~M_\\sun$ star and estimate how this planet formation mechanism scales with stellar mass. In $\\S$\\ref{sec:summary} we summarize our conclusions. ",
        "conclusions": "\\label{sec:summary} While planets are likely to form by the same basic mechanisms regardless of the environment in which they form, the properties of their host stars and the detailed structure of their nascent disks will strongly affect the statistical outcome of the formation process.  Observations hint that planets may form systematically more efficiently around intermediate mass stars, and that there is a statistically significant lack of giant planets on orbits with semi-major axes much smaller than 1 AU.  In this paper we propose a mechanism which forms giant planets preferentially around intermediate mass stars with radial distributions roughly consistent with these observations. In this model the gaseous protoplanetary disk evolves due to MRI-driven turbulence, creating a pressure maximum at the inner edge of the dead zone ($\\acrit$) which traps solid material. In order for the cores formed at this location to grow large enough to seed giant planets, the inner edge of the dead zone must be sufficiently far from the host star. We demonstrate that, as $\\acrit$ is roughly proportional to $M_*$, this condition is only likely to be met around intermediate mass stars.  The amount of solids which accumulates near $\\acrit$ is limited by momentum feedback on the gas by the solids. This is interesting as it means that the surface density of solids at $\\acrit$ depends on the gas surface density and \\emph{not} on the initial fraction of solids, except in extremely metal poor disks. This suggests that while there may be a critical metallicity required in order to form planets by the mechanism described in this paper, beyond that critical value the frequency of giant planets should depend only weakly on stellar metallicity.  The prolific production of gas giants near $\\acrit$ can also promote the emergence of additional gas giants at larger distances from the same host stars.  We may expect the fraction of intermediate-mass stars with multiple Jupiter-mass planets is likely to be larger than that around solar-type stars.  Nevertheless, we anticipate the peak in the planets' semi-major axis distribution to be around 1 AU.  This corresponds to the location of the original pressure maximum at the inner edge of the dead zone.  Quantitative verification of this expectation requires population synthesis which will be carried out and presented elsewhere.  Observational confirmation of this peaked period distribution will provide clues and constraints on the outstanding issue of magnetic turbulent transport in protostellar disks.  \\vspace{1em} \\noindent ACKNOWLEDGMENTS.  We thank D. Fischer, S. Ida, and P. Bodenheimer for many useful conversations.  This work is supported by NASA (NAGS5-11779, NNG06-GF45G, NNX07A-L13G, NNX07AI88G), JPL (1270927), and NSF(AST-0507424).  \\appendix"
    },
    "0806/0806.1717_arXiv.txt": {
        "abstract": "s{ IceCube is a cubic kilometer neutrino telescope under construction at the South Pole. The primary goal is to discover astrophysical sources of high energy neutrinos. We describe the detector and present results on atmospheric muon neutrinos from 2006 data collected with nine detector strings.} ",
        "introduction": "The IceCube detector is a cubic kilometer  neutrino telescope under construction at the South Pole\\cite{icecube}. The main goal of IceCube is to detect cosmic neutrinos of all flavors in a wide (100 GeV to 100 EeV) energy range\\footnote{So far the only observed extra-terrestrial neutrinos are low energy neutrinos from the Sun\\cite{Sun} and SN1987A\\cite{SN}.}. IceCube will search for  point sources of extra-terrestrial muon neutrinos and diffuse fluxes of extra-terrestrial neutrinos of all flavors. Possible high energy neutrino sources are active galactive nuclei (AGNs), gamma-ray bursters (GRBs) and supernova remnants (SNRs). Other IceCube physics topics include searches for WIMP anihilation in the Earth and Sun, signatures of supersymmetry in neutrino interactions, and exotica like magnetic monopoles or extra dimensions\\cite{icecube}. ",
        "conclusions": "The detection of atmospheric neutrinos with $9$ string configuration established IceCube as a neutrino telescope\\cite{icecube-first,atmospheric-9strings}. A significant improvement of the sensitivity for both point-like neutrino sources as well as diffuse neutrino fluxes is expected from 2007 data taken with $22$ strings. Analyses of these data are in progress. The IceCube detector continues to grow.  We expect  that an integrated exposure of 1km$^3$ $\\cdot$ year will be reached in 2009 and the first extra-terrestrial neutrino signal may be detected. Stay tuned!  \\newline"
    },
    "0806/0806.2302_arXiv.txt": {
        "abstract": "We present a multiwavelength study of the formation of massive stellar clusters, their emergence from cocoons of gas and dust, and their feedback on surrounding matter.  Using data that span from radio to optical wavelengths, including Spitzer and Hubble ACS observations, we examine the population of young star clusters in the central starburst region of the irregular Wolf-Rayet galaxy IC~4662.  We model the radio-to-IR spectral energy distributions of embedded clusters to determine the properties of their HII regions and dust cocoons (sizes, masses, densities, temperatures), and use near-IR and optical data with mid-IR spectroscopy to constrain the properties of the embedded clusters themselves (mass, age, extinction, excitation, abundance). The two massive star-formation regions in IC~4662 are excited by stellar populations with ages of $\\sim 4$ Ma and masses of $\\sim 3 \\times 10^5$ M$_\\odot$ (assuming a Kroupa IMF).  They have high excitation and sub-solar abundances, and they may actually be comprised of several massive clusters rather than the single monolithic massive compact objects known as Super Star Clusters (SSCs).  Mid-IR spectra reveal that these clusters have very high extinctions, A$_{\\rm V} \\sim 20-25$ mag, and that the dust in IC~4662 is well-mixed with the emitting gas, not in a foreground screen. ",
        "introduction": "\\label{intro}  Massive stars in starbursts are thought to form predominantly in Super Star Clusters (SSCs), which are massive, compact, young stellar clusters that are candidates for present-day analogues of young globular clusters because of their similar sizes and inferred masses \\citep[e.g.][]{whitmore93,schweizer96}.  Although the great majority of SSCs could be short-lived, destined to dissolve within ~$\\sim10$ megayears (Ma) and merge into the field star population \\citep{fall05,mengel05}, it is clear that their massive stars dominate feedback in starbursts.  They return enriched matter  and energy to the interstellar medium (ISM) via stellar winds and supernovae, and their hard radiation fields heat and excite gas and destroy molecules and dust grains.  A high surface density of star formation produces galactic-scale superwinds that can expel the ISM from a galaxy \\citep[e.g.][]{heckman00}.  SSCs are probably formed within giant molecular cloud complexes, so observations that can probe embedded dusty regions are required to detect the youngest SSCs. By the time SSCs have blown away enough of their natal material to be detected at visible wavelengths, they are usually at least several million years old, so their most massive individual stars are expected to have already evolved off of the main sequence.  At this stage it is difficult to infer much about SSC formation environments, because the massive stars have significantly altered them. Infrared and radio wavelengths are best suited to studying the youngest SSCs, and radio interferometry has revealed a population of heavily obscured, compact thermal radio sources with inferred nebular densities n$_{\\rm e} > 10^3$ cm$^{-3}$ and radii of a few parsecs that appear to be powered by 100s-1000s of massive stars.  These sources were dubbed ultra-dense HIIRs \\citep[UDHIIs,][]{kobulnicky99} in analogy with Galactic ultra-compact HIIRs \\citep[UCHIIs, e.g.][]{Wood89}, which are excited by one or a few massive stars on subparsec scales, with  n$_{\\rm e} > 10^4$ cm$^{-3}$.  UDHIIs are thought to be young, embedded SSCs that have not yet disrupted and expelled the majority of their natal and surrounding material.  Their hard intrinsic emission is absorbed and reprocessed by surrounding dense gas and dust to emerge as a thermal blackbody component in the mid-IR, where it peaks at $\\sim 60-100\\ \\mu$m depending on the temperature and distribution of the dust \\citep[e.g.][]{vacca02}. In the radio regime, UDHIIs are identified by a thermal bremsstrahlung spectrum that is self-absorbed below a few cm \\citep[e.g.][]{johnson03}. These features dominate the spectral energy distribution (SED) of UDHIIs and can be used along with any escaping visible light emission to constrain the temperature, density, and geometry of the gas and dust excited by the SSC, as well as the age and stellar content of the cluster itself \\citep[e.g.][]{martinhernandez05,martinhernandez06}.  Spectroscopic measurements provide further constraints on these parameters as well as the excitation and composition of the emitting material \\citep[e.g.][]{verma03}.   \\begin{figure}[t] \\begin{center} \\includegraphics[width=2.5in]{irac368_3cm_label.pdf} \\caption{ IRAC image of IC~4662 (with red, green, blue in 3.3, 5.8, 8 $\\mu$m bands respectively) features three dusty regions that are coincident with the optically identified star-formation regions A2, A1, and B (from left to right) of \\citet{heydarimalayeri90}.  Bright thermal radio sources are shown in ATCA 3 cm contours  \\citep[levels start at 3$\\sigma$, steps are 3$\\sigma$,][]{johnson03}} \\label{fig:irac} \\end{center} \\end{figure} ",
        "conclusions": "The nearby dwarf irregular galaxy IC~4662 harbors two sites of recent massive star formation that have thermal radio spectra, bright dust and H$\\alpha$ emission, and rich mid-IR spectra that feature nebular fine-structure lines, a hint of PAHs at 11.3 $\\mu$m, and broad Hu$\\alpha$ emission.  Our emission-line analysis indicates that this starburst has high excitation and low abundances like many UCD galaxies and in agreement with literature values.  However, we find much higher extinctions (A$_{\\rm V} \\sim 20-25$ mag) than previous shorter-wavelength studies, which can only be reconciled with the optical observations by a mixed geometry for the gas and dust in these regions. SED fitting of the radio-to-near-IR SEDs of the UDHIIs in regions A1 and A2 suggests that they are more evolved toward the SSC stage than similar embedded objects (e.g. in He $2-10$ and NGC~5253):  A1 and A2 have lower gas masses in their HIIRs and dust masses surrounding them, and this is consistent with their older ages of about 4 Ma, inferred from radio and optical/near-IR data.  The clusters that power the UDHIIs have masses of about $3 \\times 10^5$ M$_\\odot$ (assuming a full Kroupa IMF), but they may not be massive, compact, monolithic SSCs like those found in many starbursts, but rather clusters of $10^3-10^4$ M$_\\odot$ clusters that fill a larger volume of $~\\sim 50$ pc diameter."
    },
    "0806/0806.4063.txt": {
        "abstract": "Mass loss and axial rotation are playing key roles in shaping the evolution of massive stars. They affect the tracks in the HR diagram, the lifetimes, the surface abundances, the hardness of the radiation field, the chemical yields, the presupernova status, the nature of the remnant, the mechanical energy released in the interstellar medium,  etc...  In this paper, after recalling a few characteristics of mass loss and rotation, we review the effects of these two processes at different metallicities. Rotation probably has its most important effects at low metallicities, while mass loss and rotation deeply affect the evolution of massive stars at solar and higher than solar metallicities. ",
        "introduction": "Radiation triggers mass loss through the line opacities in hot stars. It may also power strong mass loss through the continuum opacity when the star is near the Eddington limit. For cool stars, radiation pressure is exerted also on the dust.  For hot stars, typical values for the terminal wind velocity, $\\upsilon_\\infty$ is of the order of 3 times the escape velocity, {\\it i.e.} about 2000-3000 km/s,  mass loss rates are between 10$^{-8}$-10$^{-4}$ M$_\\odot$ per year increasing with the luminosity and therefore the initial mass of the star (Vink et al. 2000; 2001). The comparison of mass loss rates for O-type stars obtained by different technics shows sometimes very important differences. For instance, Fullerton et al. (2006) using UV line of P$^{+4}$ obtained mass loss rates reduced by  a factor ten or more with respect to mass loss determination from radio or H$\\alpha$ determination. Bouret et al. (2005) obtained qualitatively similar results to Fullerton et al. (2006) but with considerable lower reduction factor (about 3). Such reduction of the mass loss rates during the O-type star phase may have important consequences. Typically a 120 M$_\\odot$ loses during its lifetime of 2.5 Myr about 50 M$_\\odot$ whith mass loss rates of the order of 2 10$^{-5}$ M$_\\odot$ per year. Dividing this mass loss rate by 10, would imply that in the same period, the star would lose only 5 M$_\\odot$! Unless stars are strongly mixed (by e.g. fast rotation), or that all WR stars originate in binary systems, it would be difficult to understand how WR stars form with such low mass loss rates.  Stars with initial masses below about 30 M$_\\odot$ at solar metallicity evolve to the red supergiant stage where mass loss is enhanced with respect to the mass loss rates in the blue part of the HR diagram (see for instance de Jager et al. 1988). In this evolutionary stage, determination of the mass loss is more difficult than in the blue part of the HR diagram due in part to the presence of dust and to various instabilities active in red supergiant atmospheres (e.g. convection becomes supersonic and turbulent pressure can no long be ignored). An illustration of the difficulty comes from the determination of red supergiant mass loss rates by van Loon et al. (2005). Their study is based on the analysis of optical spectra of a sample of dust-enshrouded red giants in the LMC, complemented with spectroscopic and infrared photometric data from the literature. Comparison with galactic AGB stars and red supergiants shows excellent agreement for dust-enshrouded objects, but not for optically bright ones. This indicates that their recipe only applies to dust-enshrouded stars. If applied to objects which are not dust enshrouded, their formula gives values which are overestimated by a factor 3-50! In this context the questions of which stars do become dust-enshrouded, at which stage, for how long, become critical to make correct implementations of such mass loss recipes in models.  Stars with initial masses above about 30 M$_\\odot$ at solar metallicity may evolve into a short Luminous Blue Variable (LBV) phase. LBV stars show during outbursts mass loss rates as high as 10$^{-4}$-10$^{-1}$ M$_\\odot$ per year. For instance $\\eta$ Carinae ejected near the middle of the eighteenth century between 12 and 20 M$_\\odot$ in a period of 20 years, giving an average mass loss rate during this period of 0.5 M$_\\odot$ per year. Such a high mass loss cannot be only radiatively driven according to Owocki et al. (2004). These authors have shown that the maximum mass loss rate that radiation can drive is given by $\\dot M\\sim 1.4 \\times 10^{-4} L_6 {\\rm M}_\\odot {\\rm yr}^{-1},$ with $L_6$ the luminosity expressed in unit of 10$^6$ L$_\\odot$. This means that for $L_6=5$ (about the case of $\\eta$Car) the maximum mass loss rate would be less than 10$^{-3}$ M$_\\odot$ per year, well below the mass loss during the outbursts. These outbursts, which are more shell ejections than steady stellar winds, involve other processes in addition to the effects of the radiation pressure. Among the models proposed let us mention the geyser model by Maeder (1992b), or the reaching of the $\\Omega\\Gamma$-limit (Maeder \\& Meynet 2000a).  After the LBV phase, massive stars evolve into the Wolf-Rayet phase, also characterized by strong mass loss rates. Many recent grids of stellar models use the recipe given by Nugis \\& Lamers (2000) for the WR mass loss rates. These authors deduced the mass loss rates from radio emission power and accounted for the clumping effects.  %\\begin{figure} %\\includegraphics[width=2.6in,height=2.6in,angle=0]{jm85z00S800.eps} %\\hfill %\\includegraphics[width=2.6in,height=2.6in,angle=0]{om85z00S800.eps} %\\caption{\\textit{Left:} profile of the specific angular momentum $j_\\mathrm{m}$ inside our 85 $M_\\odot$ at the end of hydrostatic core Si-burning (continuous line). The dotted line is $j_{\\rm K}=r_{\\rm LSO}\\,c$ \\cite[ p. 428]{ST83}, where the radius of the last stable orbit, $r_{\\rm LSO}$, is given by $r_{\\rm ms}$ in formula (12.7.24) from \\cite[p. 362]{ST83} for circular orbit in the Kerr metric. $j_{\\rm K}$ is the minimum specific angular momentum necessary to form an accretion disc around a rotating black hole. $j_\\mathrm{Schwarzschild}=\\sqrt{12}Gm/c$ (long-dashed line) and $j_\\mathrm{Kerr}^\\mathrm{max}=Gm/c$ (short-dashed line) are the minimum specific angular momentum necessary for a non-rotating and a maximally-rotating black hole, respectively. \\textit{Right:} profile of $\\Omega_\\mathrm{m}$ inside the same model. %}\\label{OME} %\\end{figure}  \\subsection{Metallicity dependence of the stellar winds}  In addition to the intensity of the stellar winds for different evolutionary phases, one needs to know how the winds vary with the metallicity. This is a key effect to understand the different massive star populations observed in regions of different metallicities. This has also an important impact on the nature of the stellar remnant and on the chemical yields expected from stellar models at various metallicities .  Current wisdom considers that very metal-poor stars lose no or very small amounts of mass through radiatively driven stellar winds. This comes from the fact that when the metallicity is low, the number of absorbing lines is small and thus the coupling between the radiative forces and the matter is weak. Wind models impose a scaling relation of the kind $ \\dot M(Z)=\\left({Z \\over Z_\\odot} \\right)^\\alpha\\dot M(Z_\\odot), $ where $\\dot M(Z)$ is the mass loss rate when the metallicitity is equal to $Z$ and $\\dot M(Z_\\odot)$ is the mass loss rate for the solar metallicity, $Z$ being the mass fraction of heavy elements. In the metallicity range from 1/30 to 3.0 times solar, the value of $\\alpha$ is between 0.5 and 0.8 according to stellar wind models (Kudritzki et al. 1987; Leitherer et al. 1992; Vink et al. 2001). Such a scaling law implies for instance that a non-rotating 60 M$_\\odot$ with $Z=0.02$ ends its stellar life with a final mass of 14.6 M$_\\odot$, the same model with a metallicity of $Z=10^{-5}$ ends its lifetime with a mass of 59.57 M$_\\odot$ (cf. models of Meynet \\& Maeder 2005 and Meynet et al. 2006 with $\\alpha=0.5$).   During the red supergiant stage, at the moment there is no commonly accepted rule to account for a possible metallicity dependence of the winds. let us just mention here that according to van Loon et al. (2005),  dust-enshrouded objects mass loss appears to be similar for objects in the LMC and the Galaxy. This may have very important consequences for our understanding of metal-poor red supergiant stars. For the LBV's, there is also no real knowledge on how mass loss can depend on metallicity. If the mechanism is mainly triggered by continuum opacity, we can expect that there is only a weak or may be no dependence on the metallicity.  Until very recently, it was considered that the WR mass loss rates did not depend on the initial metallicity {i.e.} that a WN stars in the SMC, LMC and in the Galaxy would lose mass at the same rate provided they have the same luminosity and the same actual surface abundances. This view has been challenged by Vink \\& de Koter (2005) who find that the winds of WN stars are mainly triggered by iron lines. They suggest a dependence of mass loss on $Z$ (initial value) similar to that of massive OB stars. According to these authors, the winds of WC stars depends also on the iron abundance, but in this case, the metallicity dependence is less steep than for OB stars. Their results apply over a range of metallicities given by 10$^{-5} \\le (Z/Z_\\odot) \\le 10$. Very interestingly, they find that once the metal abundance drops below $(Z/Z_\\odot) \\sim 10^{-3}$, the mass loss of WC stars no longer declines. This is due to an increased importance of radiative driving by intermediate-mass elements, such as carbon. These results have profound consequences for the evolution of stars at low metallicity, affecting the predicted Wolf-Rayet populations (Eldridge \\& Vink 2006), the evolution of the progenitors of collapsars and long soft Gamma Ray Bursts (Yoon \\& Langer 2005; Woosley \\& Heger 2006ab; Meynet \\& Maeder 2007). ",
        "conclusions": "%Massive stars are a key piece in the processes linking stellar and galactic evolution. %Understanding their formation, evolution and impact on their surrounding at various metallicities, %represent thus a necessary stage for improving our knowledge of the evolution of galaxies. %Probably at very low metallicity their evolution is changed significantly by their axial rotation %making them ``spinstars'', {\\it i.e.} stars undergoing strong rotational mixing and may be also important mass loss induced by rotation. At high metallicity, mass loss by line driven winds %add its effects producing nearly evaporating stars (at least in the upper mass range) having %a deep impact on the dynamical and chemical evolution of the interstellar medium."
    },
    "0806/0806.2134_arXiv.txt": {
        "abstract": "This work provides analytical solutions describing the post-shock structure of radiative shocks growing in astrophysics and in laboratory. The equations including a cooling function $\\Lambda \\propto \\rho^{\\epsilon} P^{\\zeta} x^{\\theta}$ are solved for any values of the exponents $\\epsilon$, $\\zeta$ and $\\theta$. This modeling is appropriate to astrophysics as the observed radiative shocks arise in optically thin media. In contrast, in laboratory, radiative shocks performed using high-power lasers present a radiative precursor because the plasma is more or less optically thick. We study the post-shock region in the laboratory case and compare with astrophysical shock structure. In addition, we attempt to use the same equations to describe the radiative precursor, but the cooling function is slightly modified. In future experiments we will probe the PSR using X-ray diagnostics. These new experimental results will allow to validate our astrophysical numerical codes. ",
        "introduction": "In this work both radiative shocks (RS) arising in astrophysics and those generated in laboratory are studied. As these RS are involved in all stages of stellar evolution (accretion shocks, pulsating stars, supernovae and interstellar medium), accurate modeling is needed and, therefore, experimental results are used to validate our codes. Until now our team performed RS ex\\-periments~\\cite{Michaut:ASS:07,Koenig:POP:06,Bouquet:PRL:04} using the high-power laser at Laboratoire d'Utilisation des Lasers Intenses (\\'Ecole Poly\\-technique). In laboratory, the plasma can be considered at local thermodynamical equilibrium and optically thick or intermediate~\\cite{Drake:Book:06}. With specific radiative hydrodynamics codes~\\cite{Michaut:ASS:07}, we study these RS structures since at high-Mach numbers ($M$) they exhibit precursor.\\\\ On the other hand in astrophysics, the plasma is often optically thin~\\cite{Bertschinger:APJ:86} and radiation escapes without interaction with the surrounding material. Its role can be modeled~\\cite{Bertschinger:APJ:86} by a cooling function $\\Lambda(\\rho,P)$. The main objective of analytical modeling suggested here is to predict the extension of the optically thin cooling zone behind the RS front. We consider the equations presented in \\cite{Chevalier:APJ:82}, but in this paper they are solved analytically for any $\\Lambda$ proportional to a power law of $\\rho$ and $P$. In this case only the post-shock region (PSR) is structured by cooling like for Polars where a RS arises for magnetic white dwarf accreting neighboring star material.\\\\ In addition, we attempt to calculate the precursor length of steady laboratory RS. Based on Drake's work~\\cite{Drake:Book:06}, we switch the cooling function by a equivalent system which represents the radiation flux propagating towards the precursor. In next experiments scheduled in '08 on LIL (Bordeaux, France), we plan to probe PSR using X-ray diagnostics. With analytical solutions of the above astrophysical model, we can predict the structure of PSR for these experiments. In the same way, experimental results of the post-shock cooling will allow to validate our astrophysical codes, since by confrontation with analytical results we can determine the main physical process {\\it i.e.} the value of the exponents in $\\Lambda$.\\\\ ",
        "conclusions": "We have generalized the analytical solutions of the model for any cooling function and compared them with previous numerical work~\\cite{Wu:APJ:94} and with some trivial analytical solutions~\\cite{Chevalier:APJ:82,Laming:PRE:04,Imamura:APJ:85}. Although we have recovered already known results, we have derived additional classes of solutions. This work provides directly pieces of information relevant to search for astrophysical objects or phenomena suitable for comparisons with (rescaled) data obtained experimentally. Based on recent Drake's work~\\cite{Drake:Book:06}, we can apply same calculation to evaluate the length of the radiative precursor of laboratory radiative shocks. Moreover with equations of the post-shock zone we can prepare future experiments on lasers even more powerful than LIL as LMJ or NIF, under condition to reach a stationary shock regime. As a result, we can emphasize the occurring physical processes through the exponent determination in the cooling function. As a conclusion, we strengthen the connection between experimental and numerical studies on radiative shocks by introducing our analytical predictions."
    },
    "0806/0806.2072_arXiv.txt": {
        "abstract": "The time scale for cooling of  the gravitationally bound  gaseous  intracluster medium (ICM)  is not determined by radiative processes alone. If the  ICM is in quasi-hydrostatic equilibrium in the fixed gravitational field of the dark matter halo then energy losses incurred by the  gravitational potential energy of the gas should also be taken into account. Here, the virial theorem is  applied to gas in quasi-hydrostatic equilibrium in an external gravitational field, neglecting the gravity of the gas. For a standard NFW form of halo profiles and for a finite gas density, the response of the gas temperature to changes in the total energy is significantly delayed. The effective cooling time could be prolonged by more than an order of magnitude inside the  scale  radius ($\\rs$) of the halo. Gas lying at a distance twice the scale radius, has negative heat capacity so that the temperature increases as a result of energy losses. Although external heating (e.g. by AGN activity)  is still required  to explain the lack of cool ICM near the center,  the analysis here may circumvent the need for heating in farther out regions where the effective cooling time could be prolonged to become larger than the cluster age and also explains the increase of temperature with radius in these regions. The  results may be relevant for large elliptical galaxies. ",
        "introduction": "Clusters of galaxies are the most massive virialized objects observed in the Universe. Their potential depths correspond to virial temperatures of $1-10\\; \\rm keV$ $(10^7-10^8 \\; \\rm K$) and the baryon number density in the inner regions could be as high as $0.1 \\rm cm^{-3}$ (e.g. Vikhlinin \\etal 2005; Pointecouteau, Arnaud \\& Pratt 2005).  For these temperatures and densities, radiative losses are expected to  the  bring the temperature in the central regions down to $\\gtsim 10^4 \\; \\rm K$ within the available time. Yet in none of the observed clusters does the temperature  drop to the level dictated by cooling alone. The absence of significant amounts of cold gas in the cores of massive clusters is a major puzzle posed by X-ray observations of massive clusters (e.g. Peterson \\etal 2001). Hence,   efficient heating mechanisms must operate at the cores of all cooling clusters. The most popular mechanism for suppressing cooling is energy released by an  AGN in the central cluster galaxy (cf.  Quilis \\etal 2001, Babul \\etal 2002, Kaiser \\& Binney 2003, Dalla Vecchia \\etal 2004, Roychowdhury \\etal 2004, Voit \\& Donahue 2005, Nipoti \\& Binney 2005, and references therein) or by multiple AGN activity in all galaxies in the cores of clusters (Nusser, Silk \\& Babul 2006; Eastman et. al. 2007; Nusser \\& Silk 2008). Over-pressurized  ejecta from the AGN transform into hot bubbles that eventually reach pressure equilibrium with the ICM and proceed to rise buoyantly away from the center. These bubbles could heat the ICM by means of shock waves generated as they expand to reach the ICM pressure (Nusser, Silk \\& Babul 2006), and by drag forces when they become buoyant (e.g. Churazov et. al. 2001). Mechanical activity near the center could also generate  sound waves which are believed to eventually dissipate their energy in the ICM (Pringle 1989, Ruszkowski \\etal 2004, Heinz \\& Churazov 2005, Fujita \\& Suzuki 2005, Sanders \\& Fabian 2007). To balance cooling in a cluster of   X-ray luminosity of $L_{\\rm x}\\sim 10^{44}\\rm erg \\; s^{-1}$, a central AGN must produce $\\sim 10^{60} \\rm erg$ over the entire life-time of the cluster. For the most massive clusters (potential depths corresponding to velocity dispersions $>500 \\; \\rm km/s$) the required heating  could be more than an order of magnitude  larger than the observed range of AGN energy output in galaxy clusters, based on the $pV$ content of X-ray cavities (e.g. Best \\etal 2007). This is not too worrying since weak shocks could certainly compensate for the missing energy needed to balance cooling. For less massive clusters the $pV $ energy is sufficient to balance cooling (e.g. B\\^irzan \\etal 2004). The challenge, however, is to arrange for  efficient energy transport from the AGN over the entire cooling core, or  out to distances of up to $\\sim 100\\;  \\rm kpc$.  The current work is motivated by the following observation. The temperature in the inner regions increases gradually as we move away from the center. At first, this behavior may seem reasonable since the radiative cooling becomes  more efficient nearer to the  center. But, the cooling time is significantly shorter than the cluster age over a significant part of the inner regions and the ICM had ample opportunity to cool to very low temperatures (e.g. Fig. 12 in Wise, McNamara \\& Murray  2004). So why is there not a temperature plateau extending over the region where the cooling time is shorter than the cluster age? One explanation might be that,   on account of the lower density, heat conduction is more significant as we  move away from the center.  However, heat conduction is not universally  important in these regions (e.g. Wise, McNamara \\& Murray 2004). Here we offer an explanation for this behavior of the temperature profile. The explanation may also help alleviate other problems associated with observations of the ICM. In a nutshell we will show that the cooling time could significantly be modified when the potential energy of the ICM in the dark halo is taken into account. We will use a version of the virial theorem to show that the potential energy will absorb some of the  energy loss incurred by the system. In some cases  the potential energy will decrease by an amount larger than the actual loss, forcing the system to compensate the energy difference by increasing  its  thermal energy. This is the case of negative heat capacity. ",
        "conclusions": "\\label{sec:Conclude}  We used a simplified model of the ICM to study the its gravo-thermodynamical properties. In quasi-hydrostatic equilibrium, the inclusion of the change in the potential energy prolongs the response of the   gas temperature  in the inner regions lying within $r_{_{\\rm th}}\\approx  2 r_{_{\\rm s}}=2R_{_{\\rm v}}/c$. Outside this radius, the form of dark halo gravitational potential is such that  the temperature is increased as a result of energy loss, i.e. the gas heat capacity is negative. The boosting of the cooling time in the inner region formally diverges at $r_{_{\\rm th}}$ reaching  a  factor of $3$ as at  $r=0$.   Our results may circumvent the need for heating the ICM in regions where the standard radiative cooling time is an order of magnitude shorter than the life-time of the cluster.  Those regions lie at a significant fraction of $r_{_{\\rm s}}$ where  the prolonged effective cooling time could be larger than  the cluster age. as is the case for example for the cluster A1068. For this cluster, $r_{_{\\rm s}}\\sim 400 \\; \\rm kpc $ (Pointecouteau, Arnaud \\& Pratt 2005) and  the ratio of  standard radiative cooling time to the cluster age  is  $\\sim 0.1- 1$ over the region between $\\sim 70 \\; \\rm kpc$ to $\\sim 300\\;  \\rm kpc$. Our results do not eliminate the need for heating of the ICM in central regions ($r<< r_{_{\\rm th}} $) since the effective cooling time there is still shorter than the cluster age.    As mentioned before, the approach taken here aims at addressing a specific point related to the heat capacity of the ICM. The analytic methods used  here could be followed only by invoking a simplified (perhaps oversimplified) of the ICM.  A more thorough a analysis should include  a variety of effects such as large scale motions, buoyant bubbles, the gravity of the gas and the back-reaction of the halo profile as a result of the variations  in the gas distribution."
    },
    "0806/0806.3074_arXiv.txt": {
        "abstract": "We provide computationally convenient expressions for all marginal distributions of the polarization CMB power spectrum distribution $P(C_{\\ell}|\\sigma_{\\ell})$, where $C_{\\ell} = \\{C_{\\ell}^{\\textrm{TT}}, C_{\\ell}^{\\textrm{TE}}, C_{\\ell}^{\\textrm{EE}}, C_{\\ell}^{\\textrm{BB}}\\}$ denotes the set of ensemble averaged polarization CMB power spectra, and $\\sigma_{\\ell} = \\{\\sigma_{\\ell}^{\\textrm{TT}}, \\sigma_{\\ell}^{\\textrm{TE}}, \\sigma_{\\ell}^{\\textrm{EE}}, \\sigma_{\\ell}^{\\textrm{BB}}\\}$ the set of the realization specific polarization CMB power spectra. This distribution describes the CMB power spectrum posterior for cosmic variance limited data. The expressions derived here are general, and may be useful in a wide range of applications.  Two specific applications are described in this paper. First, we employ the derived distributions within the CMB Gibbs sampling framework, and demonstrate a new conditional CMB power spectrum sampling algorithm that allows for different binning schemes for each power spectrum. This is useful because most CMB experiments have very different signal-to-noise ratios for temperature and polarization. Second, we provide new Blackwell-Rao estimators for each of the marginal polarization distributions, which are relevant to power spectrum and likelihood estimation. Because these estimators represent marginals, they are not affected by the exponential behaviour of the corresponding joint expression, but converge quickly. ",
        "introduction": "During the last few decades cosmology has evolved from a data starved branch of astrophysics, into a data driven high-precision science in which theories may be subjected to stringent observational tests. This revolution has to a large extent been driven by steadily improving observations of the cosmic microwave background (CMB), allowing cosmologists to have a close-up look at the very young universe. Two leading experiments were the COBE-DMR \\citep{smoot:1992} and WMAP \\citep{bennett:2003} satellite missions, while the third generation experiment, Planck, will be launched late this year.  As observations continue to improve, increasingly demanding requirements are imposed on the data analysis. While rather crude approximations may be acceptable when interpreting low signal-to-noise data, the situation is very different in the mid and high signal-to-noise regime. Here, even ``small'' effects become clearly visible, and may potentially compromise any cosmological conclusion. Using accurate methods in this regime is critical. Some real-world issues relevant to the CMB problem are non-cosmological foregrounds, improper noise and/or beam characterization, and sub-optimal likelihood approximations.  In 2004, a new approach to CMB analysis was proposed and implemented by \\citet{jewell:2004}, \\citet{wandelt:2004} and \\citet{eriksen:2004}. Rather than taking the traditional approximate Monte Carlo approach \\citep[e.g.,][]{hivon:2002}, this new method employs the Gibbs sampling algorithm to facilitate exact (in the maximum-likelihood sense), global and efficient analysis of even high-resolution data sets.  Equally important, the Gibbs sampling framework has unique capabilities for error propagation, as it allows for easy marginalization over virtually any auxiliary stochastic field. One important example is that of non-cosmological foregrounds.  Since then, the method has been generalized to handle polarized CMB data \\citep{larson:2007} and joint foreground and CMB analysis \\citep{eriksen:2008a}, and has been applied most successfully to the WMAP data \\citep{odwyer:2004, eriksen:2007a, eriksen:2007b, eriksen:2008b}. Some useful examples of issues correctly identified by the Gibbs sampler, but missed by other techniques, are 1) the first-year WMAP likelihood bias at $\\ell \\lesssim 30$ \\citep{eriksen:2007a, hinshaw:2007}, 2) foreground residuals in the 3-year WMAP polarization sky maps \\citep{eriksen:2007b}, and 3) residual monopole and dipole components in the 3-year temperature sky maps \\citep{eriksen:2008b, hinshaw:2008}. Following up on these methodological advances, the WMAP team adopted the Gibbs sampler as a central component in their analysis of the 5-year data, and, in fact, their default low-$\\ell$ temperature likelihood module is precisely the Gibbs-based Blackwell-Rao code written and published by \\citet{chu:2005}.  While WMAP has done an excellent job on characterizing the large-scale CMB temperature fluctuations, the current frontier in CMB science is polarization. In just a few years, full-sky high-sensitivity data will be available from Planck. And then, very likely, the situation will be quite analogous to the one WMAP experienced in the temperature case: Having robust, exact methods that allows for proper characterization and propagation of systematics will be essential in the mid to high signal-to-noise regime. The Gibbs sampler is among the leading candidates to serve such a purpose.  Unfortunately, the Gibbs sampler, as currently described in the literature, has two major limitations that needs to be resolved before this promise can be fulfilled. First, the direct Gibbs sampler is inherently inefficient in the low signal-to-noise regime, because the step size between two consecutive samples is determined by cosmic variance alone, whereas the full posterior width is determined by noise.  Second, it is non-trivial to establish a full likelihood approximation from the samples produced by the Gibbs sampler, because of the dimensionality of the underlying space. Both of these issues are currently under development, and reports are expected in the near future (Jewell et al. 2008; Rudjord et al. and Eriksen et al., in preparation).  In the present paper, we take a small but important first step towards resolving these issues, by considering the marginal and conditional densities of the probability distribution $P(C_{\\ell}|\\sigma_{\\ell})$, where $C_{\\ell}$ is the ensemble averaged CMB power spectrum, and $\\sigma_{\\ell}$ is the observed power spectrum of one given CMB realization. This distribution plays a crucial role within the CMB Gibbs sampling framework. On the one hand, it forms one of the two conditionals in the main sampling scheme. On the other, it is the kernel of the Blackwell-Rao estimator. Being able to describe this analytically in different forms is therefore very useful. Two specific applications are demonstrated in this paper, namely 1) a $C_{\\ell}$ sampling algorithm that allows for different binning schemes in each of the polarization components, and 2) Blackwell-Rao estimators for each of $P(C_{\\ell}^{TT}|C_{\\ell}^{TE}, C_{\\ell}^{EE},\\mathbf{d})$, $P(C_{\\ell}^{TE}|C_{\\ell}^{EE},\\mathbf{d})$, and $P(C_{\\ell}^{EE}|\\mathbf{d})$. Further applications will be demonstrated in the papers mentioned above.  We also note that these expressions are completely general, and may prove useful for any other method that considers both $C_{\\ell}$ and the CMB sky signal $\\mathbf{s}$ as free variables. One such example is the CMB Hamiltonian sampler recently developed by \\citet{taylor:2007}. ",
        "conclusions": "\\label{sec:conclusions}  We have derived computationally convenient expressions for all marginals of $P(C_{\\ell}|\\sigma_{\\ell})$. These expressions may be useful to any CMB analysis method that considers both the CMB sky signal $\\mathbf{s}$ and the power spectrum $C_{\\ell}$ as unknown parameters. One prominent example is the CMB Gibbs sampler.  We have also presented two specific applications of these expressions. First, we demonstrated a new sampling algorithm for $P(C_{\\ell}|\\sigma_{\\ell})$ that supports different binning schemes for each polarization component. This is useful because most experiments have very different signal-to-noise ratio to temperature and polarization.  Second, we have provided explicit expressions for the Blackwell-Rao estimators for $P(\\ct|\\cx,\\ce,\\mathbf{d})$, $P(\\cx|\\ce, \\mathbf{d})$ and $P(\\ce|\\mathbf{d})$. Together, these three can be used to map out the joint distribution $P(C_{\\ell}|\\mathbf{d})$ for a single multipole in terms of univariate distributions alone. Further, we note that any of the distributions listed in Table \\ref{tab:distributions} give rise a separate, and potentially useful, Blackwell-Rao estimator."
    },
    "0806/0806.3597_arXiv.txt": {
        "abstract": "Astrophysical constraints of new physics are often limited to weakly interacting light particles, such as axions, the Kaluza-Klein (KK) gravitons from the ADD model, sterile neutrinos and unparticles. We discuss the possibility for an astrophysical scenario to (dis)confirm new physics for heavy particles beyond $\\mathrm{TeV}$ energy scale. In our scenario, the KK protons (the KK excited quarks/gluons within protons) within the framework of universal extra dimensions (UEDs), are produced by high energy $p + p$ collisions in Fermi accelerated environments, with protonic isotropic spectrum $d N / d E \\propto E^{-2}$ up to at least $10^{18}\\,\\mathrm{eV}$. Thus, because they are also electrically charged, they should be re-accelerated by mechanism similar to normal protons. The KK states (no matter whether they have already decayed to the lightest KK particle or not) should contaminate $10^{-5}$ to $10^{-2}$ of cosmic-ray events for some fixed energy $E$ (within some suitable assumptions). Hence, if we have techniques to identify them from air shower data, we can constrain UEDs scenario. Our method is an ``existence proof'' that we can constrain new physics beyond $\\mathrm{TeV}$ scale or much higher by classical astrophysical scenarios, which can also be generalized to supersymmetric models, the bulk Standard Model fields within the RS model, and the endlessly emerging new models. Moreover, it can exploit domains which have no possibility to be studied in terrestrial experiments.  \\vspace{0.5cm} \\noindent \\emph{Key words}: cosmic-rays; Fermi mechanism; Kaluza-Klein states; models beyond the Standard Model; universal extra dimensions\\\\ \\emph{PACS}: 12.60.-i, 04.50.-h, 95.85.Ry, 95.30.Qd, 98.35.Eg ",
        "introduction": "}  Brane-world scenarios, such as the Arkani-Hamed-Dimopoulos-Dvali (ADD) model~\\cite{ArkaniHamed:1998rs,Antoniadis:1998ig,ArkaniHamed:1998nn} and Randall-Sundrum (RS) model~\\cite{Randall:1999ee,Randall:1999vf}, give an alternative framework for solving the hierarchy problem. Within the ADD model, the Standard Model (SM) fields are confined to a $3$-brane ($(3+1)$-dimensional spacetime) while gravitons propagate freely in a torus compactified bulk space (large extra dimensions), and the gravitational coupling constant $G = 1/M_{\\mathrm{pl}}^2$ observed in our $(3+1)$-dimensional world is an effective one. So it is natural to understand why $G$ is so small. The RS model solves the same problem by a slice of $AdS_5$ spacetime. Some superstring-inspired descriptions make these models more attractive.  Some lineage scenarios of the ADD model also allow the SM particles propagating, to some extent, in the bulk spaces. Beside gravitons, such kind of SM particles can also be Kaluza-Klein (KK) excited in the extra dimensions, thus give abundant physical phenomena. A natural extension of the ADD model is let the brane with a finite thickness and complex substructures~\\cite{DeRujula:2000he}. An example is given in Ref.~\\cite{ArkaniHamed:1999dc}. In this scenario, quarks and leptons are confined to different branes, while the Higgs and SM gauge fields are sandwiched in, hence the possibility for proton decay can be exponentially suppressed. The KK excited gauge particles (which may be called branons) can be either baryophobic or leptophobic, because they feel nontrivial on the brane substructures. Another natural extension is asymmetrical compactification, which has two (as minimum, maybe more) separate compactification scales. The ``very large'' extra dimensions let only graviton propagate, just as what the ADD model says, but the $\\mathrm{TeV}^{-1}$ scale ``large'' extra dimensions, may have the SM fields extending~\\cite{Lykken:1999ms}. Of course we can on the other hand keep all compact dimensions at $\\mathrm{TeV}^{-1}$ scale rather than make the compactification asymmetric. However, the advantage of solving the hierarchy problem in the ADD model is bereft.  In the case of that the SM fields also extend to some extra dimensions, to obtain chiral fermions in the 4-dimensional effective theory, we have only two ways to go: (i) to confine fermions to branes only~\\cite{Dicus:2000hm}, or (ii) to impose bulk fermions orbifold boundary conditions~\\cite{Georgi:2000wb}. Universal extra dimensions (UEDs)~\\cite{Appelquist:2000nn} scenario is an example of the second approach. In UEDs scenario, all SM fields can propagate in these ``universal'' extra dimensions, and conservation of momentum in the universal dimensions turns to conservation of the KK number in our $(3+1)$-dimensional world. For two or more universal extra dimensions, the na\\\"{i}ve KK mode sums diverge when the KK tower number $N_\\mathrm{KK} \\rightarrow \\infty$. So, let us consider only one universal extra dimension as minimal universal extra dimensions (MUEDs) in this context. In this case, $S^1 / \\mathbb{Z}_2$ orbifold compactification is assumed, and the KK mass eigenvalues have a simple form $M_n^\\mathrm{KK} = n / R$, where $R$ is the compactification scale. For the reason that only loop diagrams can contribute electroweak observables by the KK number conservation, the experimental bound for UEDs is only $M_1^\\mathrm{KK} = 1/R \\geq 300\\,\\mathrm{GeV}$. In the tree level, the mass spectrum of the KK excited SM particles has the form $M_{\\mathrm{SM},\\,n} = \\sqrt{(M_n^\\mathrm{KK})^2 + M_\\mathrm{SM}^2}$, where $M_\\mathrm{SM}$ is the zero-mode on-shell mass of the corresponding SM particles. Hence, their masses are level-by-level highly degenerated when $M_\\mathrm{SM} \\ll M_n^\\mathrm{KK}$. However, when radiative corrections are concerned~\\cite{Georgi:2000ks,Cheng:2002iz}, the mass degeneration is broken, to some extend, as $M_{\\mathrm{g},\\,n} > M_{\\mathrm{Q},\\,n} > M_{\\mathrm{q},\\,n} > M_{\\mathrm{W^\\pm},\\,n} \\sim M_{\\mathrm{Z^0},\\,n} > M_{\\mathrm{L},\\,n} > M_{\\mathrm{l},\\,n} > M_{\\mathrm{\\gamma},\\,n} \\sim M_n$, where $g$ denotes gluon, $Q$ ($L$) denotes weak-doublet quark (lepton), $q$ ($l$) denotes weak-singlet quark (lepton), and $\\gamma$ denotes photon. The KK number conservation breaks down to a KK parity that the even and odd KK numbers cannot transform to each other. The correction scale depends on some unknown parameters; however, some reasonable choice of parameters shows that the largest correction $\\Delta M_{\\mathrm{g},\\,n}$ may be as large as $10\\%$~\\cite{Cheng:2002iz}. Hence, heavier KK excited states should cascade decay (by the KK conserving or even violating interactions) to the lightest KK particle (LKP) $\\gamma_1$ which is stable~\\cite{Cheng:2002ab}, by emitting soft SM particles. When considering the possibilities of experimental discovery, one always assumes that the lifetimes of heavier KK excited states are sufficiently short, thus the states can decay within the collider; however, it is not supposed to do so. The total width cannot in fact be calculated by the failure of reconstructing the Breit-Wigner resonance~\\cite{Datta:2005zs}. Notice that the decay rate of an unstable particle $d\\Gamma \\propto 1/m_\\mathcal{A}$ in the phase space formula, the lifetime $\\tau \\propto m_\\mathcal{A}$ where $m_\\mathcal{A}$ is the mass of the decay particle. If the $\\Delta M$s are smaller for a set of parameters different from in~\\cite{Cheng:2002iz}, or soft SM cascade processes are suppressed by other reasons, the lifetimes of heavier KK excited states should be even longer. Specific calculations for whether the not-very-short-lived KK excited states can affect other more mature scientific scenarios, such as disturb predictions of Big-bang nucleosynthesis, or distort the cosmic microwave background, are need; however, they are beyond our scope of our paper. Some na\\\"{i}ve considerations show that all of them are not very crucial, because we do not really need longer-lived KK excited states (although the long-lived ones are also possibilities we shall consider in the identification section in \\S\\ref{(sec)_air_shower_identification}), but some not-very-short-lived KK excited states to suffer the time scale of Fermi acceleration (which is maybe $\\sim \\mathrm{s}$ or much shorter), which is much shorter than the time scale of the scenarios we mentioned above. So we assume that the lifetimes of heavier KK excited states are long enough to suffer the astrophysical scenario we draw in this paper.  Astrophysical constraints of new physics are often limited to weakly interacting light particles, such as axions~\\cite{Dicus:1979ch,Fukugita:1982gn,Iwamoto:1984ir,Dearborn:1985gp,Frieman:1987ui,Raffelt:1987yu,Raffelt:1987yt,Turner:1987by,Burrows:1988ah,Haxton:1991pu,Raffelt:1994ry,Keil:1996ju}, the KK gravitons from the ADD model~\\cite{ArkaniHamed:1998nn,Barger:1999jf,Cullen:1999hc,Cassisi:2000hy,Hanhart:2000er,Biesiada:2001iy,Hannestad:2001xi,Hannestad:2003yd}, sterile neutrinos~\\cite{Kusenko:1997sp,Kusenko:1998bk,Hidaka:2006sg,Fryer:2005sz} and unparticles~\\cite{Hannestad:2007ys}. We want to construct an astrophysical scenario to (dis)confirm new physics for heavy particles beyond $\\mathrm{TeV}$ energy scale. Notice that in Fermi accelerated environments, protons in a power law spectrum up to at least $10^{18}\\,\\mathrm{eV}$ should be produced (even if we have already derived an overall Lorentz factor $\\Gamma \\sim 300$), thus $p + p$ collisions up to a tremendously large energy should happen there, which we cannot even imagine in terrestrial experiments. However, we have to brain storm to know their happenings. In this paper, we construct a scenario which may (dis)confirm UEDs by high energy observation of cosmic-rays. This scenario may or may not have opportunities to give stronger bounds than colliders, because of the large uncertainties in our estimations, and the technical details of lots of synergic scientific domains (which we cannot discuss at length in this paper). However, it is at least an ``existence proof'' for this kind of methodology. It can also explore domains which have no possibility to be studied in terrestrial experiments. Some similar scenario in Ref.~\\cite{Anchordoqui:2004bd}, also suggested the production of some kind of strongly interacting massive particles in $p + p$ collisions in astrophysical environments; however, our scenario have a lot of advantages than theirs. The advantages rise mainly because (i) one of the protons in their scenario stays at rest, but both of the protons in our scenario are Fermi accelerated, and (ii) our KK excited states suffer an additional accelerated process. Detail comparisons are given in \\S\\ref{(sec)_discussion_outlook}.  In our scenario, the KK protons (with either KK excited quarks or gluons in it) are produced by $p + p$ collisions in Fermi accelerated environments. Both the original Fermi mechanism or diffusive shock accelerating model have an isotropic spectrum $d N / d E \\propto E^{-2}$ up to at least $10^{18}\\,\\mathrm{eV}$, hence they are okay for our purpose. The KK protons should also be accelerated just as normal protons by the same mechanism; however, they should have different properties than normal ones. Beside being discovered one by one from air shower data directly, they may be accelerated to energies normal protons cannot be accelerated to, or they (or their decayed final state) may contaminate significant amount of ultra-high-energy cosmic-ray (UHECR) events because of an overall energy shift, both of which may make them a discovery. In \\S\\ref{(sec)_Producing}, we calculate the cross section and production rate of the KK protons. We show that the production rate may be large enough to make meaningful scientific constraints. In \\S\\ref{(sec)_accelerating_propagating}, we discuss the accelerated property of them in Fermi accelerated environments, making a comparison with normal protons. We notice that the KK states should contaminate $10^{-5}$ (for special sources) to $10^{-2}$ (for diffuse flux) of cosmic-ray events for some fixed energy $E$ (if assuming the optical depth $\\tau_\\mathrm{pp} = 1$), which are not too small a sample to be discovered by air shower detection. We also discuss the propagating properties of them related to soft photon interactions. In \\S\\ref{(sec)_air_shower_identification}, we consider the probabilities to identify them (or their decayed final state) from other cosmic-ray particles from air shower data. If it can be done so, our method can have larger possibilities to give smaller parameter space for UEDs than other methods. In \\S\\ref{(sec)_neutrino_detectors}, we calculate the possible constraints of the KK cosmic-ray flux from neutrino detectors; however, the constraints are very loose for current scientific equipments to affect our former estimations. We discuss our results and draw the possibilities to generalize our method to other new physics models in \\S\\ref{(sec)_discussion_outlook}. ",
        "conclusions": "}  Phenomenologies link ambitious theoretical physical models to reality, thus make physics a \\emph{science}~\\footnote{At least in the philosophy of Karl Popper or previously.}. Terrestrial experiments are one of the ordinary methods to constraint new physics models; however, they have limited power because of our finite energy sources on earth. Cosmology can open an extraordinary window for new physics studies; however, Big Bang (and the extreme physical environment it has) happened only once in our universe, thus makes re-enactment impossible. Astrophysical constraints always have larger scope than terrestrial experiments (because they do not have the upper threshold of maximum achievable energy); however, meaningful scenarios presently known are always only available for weakly interacting light particles. The motivation of this paper is to search for another way to construct new physics beyond the SM.  In this paper, we construct an astrophysical scenario to (dis)confirm new physics for heavy particles beyond $\\mathrm{TeV}$ energy scale. In our scenario, the KK protons are produced by $p + p$ collisions in Fermi accelerated environments, and they themselves are accelerated by the same environments. Because they may change the compositions and proterties of cosmic-ray events, air shower experiments can give a constraint to their properties.  To know whether our scenario can give meaningful constraints to UEDs (and maybe other new physics models in later researches), we make some quantitative estimations. We first investigate whether enough KK excited states can be produced by $p + p$ collisions. We calculate the overall KK cross sections by precalculated Feynman rules and amplitude-squared, with also the CTEQ6.6m proton PDFs. Subprocesses $g + g \\rightarrow g_n^\\ast + g_n^\\ast$ and $q + q \\rightarrow q_n^\\ast + q_n^\\circ$ may be most important in quark level. Because of color confinement, the KK excited quarks and gluons should form the KK excited protons. We then calculate what percentage of the KK protons can be produced by an isotropic and power law distributed proton spectrum. For the spectral index $\\alpha = 2.0$, $n_\\mathrm{p_\\mathrm{KK}} / n_\\mathrm{p} \\sim 10^{-10}$ as an overall contamination after one time of $p + p$ collision. However, when considering that the KK protons are also accelerated by the Fermi mechanism, the scene is much different. It is reasonable to believe that the energy spectrum of the KK protons should shift to higher energy than protons by an amount of $m_\\mathrm{KK} / m_\\mathrm{p}$, hence for some fixed energy $E$, the KK states should contaminate $10^{-5}$ (for some special astrophysical sources with the spectral index $\\alpha = 2.0$) to $10^{-2}$ (for diffuse flux) of cosmic-ray events. So, if we can identify them from other cosmic-ray particles from air shower data, our method is capable of giving meaningful constraints. We notice that the GZK cutoff energy also shifts by a factor of $m_\\mathrm{KK} / m_\\mathrm{p}$ to higher energy, if it still exists. Hence, observations of cosmic-ray particles with energy much above the GZK cutoff, are given a reasonable explanation by the KK particles; however, the possibilities of this kind of observation is really small, even if the GZK cutoff does not exist. We also investigate the possibilities to identify the KK cosmic-ray events by air shower data. The investigation is still superficial, because quantitative simulations are needed; however, the LKP $\\gamma_1$ may be easy to identify, because they may interact with geomagnetic fields just like normal photons, thus make the air shower tomography much different from that of protons. Finally, we calculate the possible constraints of the KK cosmic-ray flux from neutrino detectors; however, the constraints are very loose for current scientific equipments to affect our former estimations.  It is appropriate to regard our calculations (outlined in this paper) as an ``existence proof'' for this kind of methodology. In fact, any charged particles beyond the SM, which are neither too light~\\footnote{Of course, this scenario is also suitable for lighter particles. However, we can restrict the parameter space (of the endlessly emerging new models) tighter by other astrophysical/terrestrial methods. Hence, this scenario may specialize in new physics particles beyond $\\mathrm{TeV}$ energy scale, especially just above the energy scale the best colliders can in touch.} nor having a too short lifetime to suffer a Fermi acceleration, are suitable for our scenario. One immediate example is $W_1^\\pm$ in UEDs, which we do not discuss in this paper because the extension is really straightforward. $W_1^\\pm$ can cascade decay to $L_1$ or $\\nu_1$~\\cite{Cheng:2002ab}; however, because of their analogous masses, the lifetimes of $W_1^\\pm$ should be much longer than $W^\\pm$ in the Glashow-Weinberg-Salam theory of weak interactions. In fact, $W_1^\\pm$ is also an intermediate state of our $g_1$ decay in this paper. Some lineage scenarios of the ADD model, which allow bulk bosons rather than bulk fermions~\\cite{DeRujula:2000he,ArkaniHamed:1999dc,Dicus:2000hm}, can also excite Kaluza-Klein $W^\\pm$ which suffer our accelerating. However, a careful calculation of production rates and lifetimes is absent. There are also a lot of lineage scenarios of the RS model which allow the bulk SM fields~\\cite{Goldberger:1999wh,Chang:1999nh,Gherghetta:2000qt}. These models are more reasonable, because the RS model has an inherent orbifold configuration (to obtain chiral fermions), and the bulk SM fields can help us to understand some stiff physical problems like fermion mass hierarchy. Charged sparticles in supersymmetric (SUSY) models are also good candidates for these kind of scenarios~\\cite{Martin:1997ns}. In order to explain the non-baryonic dark matter, the lightest sypersymmetric particle (LSP) is preferred to be electrically neutral; however, it is not supposed to do so. Even if the LSP is really neutral, they can also be accelerated by the Fermi mechanism if the lifetimes of the charged ones decaying to them are not very short. Whereas different from the case of the LKP $\\gamma_1$, if the LSP is gravitino or the lightest neutralino, they may hardly cause air shower processes because of their relatively small cross sections, even if they bang into the earth with tremendous energy. Gluino (which may exist as the form of gluino-containing hadron, compare to our KK proton) in split supersymmetry is also a good idea. There has already been one paper in Ref.~\\cite{Anchordoqui:2004bd}, in which the gluinos are produced by astrophysical $p + p$ collisions; however, (the astrophysical aspect of) our scenario has a lot of advantages than theirs, include: (i) One of the protons in their scenario stays at rest, hence the center-of-mass energy $\\sqrt{s} = \\sqrt{m_\\mathrm{p} E_\\mathrm{p}}$ of $p + p$ collisions is at least $10^{14-15}\\,\\mathrm{eV}$, which is no more than $p + \\mathrm{(air)}$ (for UHECRs to collide with the atmosphere hadrons) center-of-mass energy here in earth; however, because both of the protons in our scenario are Fermi accelerated, the center-of-mass energy of our $p + p$ collisions should be at least $10^{18}\\,\\mathrm{eV}$ (we have already derived an overall Lorentz factor $\\Gamma \\sim 300$), which is impossibility for any other scenarios to achieve near earth. Despite the fact that we do not know whether our ideas of (renormalizable) quantum field theory or new physics nowadays are suitable for such a huge center-of-mass energy $\\sqrt{s}$, we know that something should happen there. (ii) The maximum energy a gluino-containing hadron can achieve in their scenario is only $10^{13.6}\\,\\mathrm{eV}$~\\cite{Anchordoqui:2007pn}; however, the maximum energy of our exotic cosmic-ray particles, can be even larger than the GZK cutoff (see \\S\\ref{(subsec)_GZK_cutoff} for a detailed discussion). Hence, because of the fact that the cosmic-ray spectrum itself has a large negative power law index of about $- \\bar{\\alpha} \\sim -2.7$ to $-3.0$ (below or above the ``knee'' energy), for gluino-containing hadrons of energy $E_\\mathrm{exo}$ produced by protons with energy $E_\\mathrm{p}$, their content of cosmic-rays with definite $E$ has an additional inhibitory factor of $(E_\\mathrm{exo} / E_\\mathrm{p})^{\\bar{\\alpha}}$, which has an order of magnitude of $(10^{13.6}/E_\\mathrm{knee})^{2.7} \\times (E_\\mathrm{knee}/E_\\mathrm{GZK})^{3} \\sim 7.7 \\times 10^{-19}$. If most of the UHE protons did not produce gluino-containing hadrons, the inhibition should be even stronger. Thus even if the gluino-containing hadrons are recorded by our scientific equipment (e.g., the cosmic-ray observatories), they are very difficult to be found out by such a lot of events with similar energy. However, because our charged exotic particles (KK protons in the context) are also accelerated by Fermi acceleration, their content in the UHE region should be as large as $10^{-5}$ to $10^{-2}$, hence not very hard to be identified. (iii) We do not really need the charged exotic particles to be longeval enough to suffer the travel from the source to earth; it is enough that their longevities are long enough to suffer an astrophysical Fermi acceleration. Hence, we do not have to worry about some adolescent (comparison with the``baby universe'' era) cosmological bounds, such as the predictions of Big-bang nucleosynthesis, or cosmic microwave background. Another very interesting particle candidate for our scenario is the charged massive particles (CHAMPs)~\\cite{DeRujula:1989fe}. It is an ambitious dark matter constituent as yet (hence it is longeval to suffer a Fermi acceleration). The association of our astrophysical scenario drawing in this context and the CHAMPs, is a very interest issue; however, we leave it to the later publications."
    },
    "0806/0806.3768_arXiv.txt": {
        "abstract": "We discuss the testable predictions of a phenomenological model in which the accelerated expansion of the universe is the result of the action of a non-gravitational force field, rather than the effect of a negative-pressure dark-energy fluid or a modification of general relativity.  We show, through the equivalence principle, that in such a scenario the cosmic acceleration felt by distant standard candles like SNIa (type Ia Supernovae (SNe)) depends  on the mass of the host system, being larger in isolated galaxies than in rich clusters. As a consequence, the scatter in the observed SNIa Hubble diagram has mostly a physical origin in this scenario: in fact, the SNIa distance modulus is increasing, at fixed redshift, for SNe that are hosted in isolated galaxies with respect to the case of SNe hosted in rich galaxy clusters. Due to its strong dependence on the astrophysical environments of standard candles, we conclude that alternative non-gravitational mechanisms for the observed accelerated expansion of the universe can be interestingly contrasted against the standard metric interpretation of the cosmological acceleration by means of an environmental analysis of the cosmic structures in which SNIa are found. The possible absence of such environmental effects would definitely exclude non-gravitational mechanisms being responsible for the accelerated cosmological expansion and will therefore reinforce a metric interpretation. ",
        "introduction": "\\label{intro}  The unprecedent convergence of observational results that we are currently witnessing has narrowed down the region of the cosmological parameter space which is compatible with all the different and independent probes of cosmology: Supernovae \\cite{Riess98} \\cite{Perlmutter99}, CMB \\cite{deb00, Spergel07} and Large Scale Structures \\cite{Tegmark06, Guzzo08, Mar08}. Robustly growing evidence suggests that {\\it i)} ordinary matter is a minority ($\\sim 1/6$)  of all the matter content of the universe, {\\it ii)} matter -- mostly dark, non-baryonic matter -- is a minority ($\\sim 1/4$) of all the cosmological mass--energy density, mostly contributed by an obscure form of energy referred to as `dark energy', {\\it iii)} the 3D spatial geometry of the universe is flat and {\\it iv)} the expansion of the cosmic metric has been accelerating for the last $\\sim 7$ Gyrs of our universe lifetime.  Even though the big picture is in place, the two dominant contributions to the stress-energy tensor -- i.e. dark energy and dark matter -- still have a hypothetical nature and they have  not been discovered yet. While there is widespread evidence for the existence of the non-baryonic dark matter component producing the potential wells of large-scale structures \\cite{Clowe06}, as yet no persuasive theoretical explanation has been able to elucidate the physical nature of the dark energy component \\cite{pebrat}. As a matter of fact, unveiling the nature of dark energy and its role in cosmology and gravitation is a difficult and subtle challenge. In such a context, it should not be overlooked that the large roaming from model to model, and the abundance of weakly predictive theories, might eventually limit the possibility of discriminating between different competitors proposed so far for explaining the observed dynamics of the accelerating universe.  In the absence of a compelling theoretical explanation for the dark energy component, and in a minimal, zero-order approach, we explore here  the possibility of preserving the standard metric interpretation of the accelerated expansion against possible alternative physical scenarios. To this end, we first evaluate and then discuss the observable consequences of local, non-gravitational mechanisms which could in principle  accelerate matter in our Hubble patch of the universe. We assume here that the universe is described by general relativity, that it is dominated by components which satisfy the usual energy conditions (according to which the universe can only decelerate) and that the onset of recent accelerated expansion  is the result of the presence of a hypothetical non-gravitational force field. Such an alternative explanation is rather conservative, since it assumes neither a cosmological constant (or negative-pressure fluid) nor a modification of general relativity. Accordingly, we first work out a self-consistent, non-geometric model for the cosmic acceleration that is able to reproduce the current observations of standard candles (i.e. SNIa) and then we discuss a falsifiability procedure aimed at testing its observational predictions.  The motivation behind this work is to put strong limits on a hypothetical (or non usually considered) physics that is possibly missing in our picture of the universe, and, in turn, to strengthen the evidence supporting the standard paradigm with which we are currently explaining its past history, its present stage and its future fate. ",
        "conclusions": ""
    },
    "0806/0806.2857_arXiv.txt": {
        "abstract": "We consider bound geodesic orbits of test masses in the exterior gravitational field of a rotating astronomical source whose proper angular momentum varies linearly with time. The linear perturbation approach of Lense and Thirring is herein extended to the nonstationary case. In particular, we investigate the instability of Lense-Thirring precessing orbits due to the slow temporal variation of the gravitomagnetic field of the source. ",
        "introduction": "\\label{s1}  Nine decades ago, Lense and Thirring considered the motion of a free test mass in the stationary exterior gravitational field of a rotating astronomical source within the framework of general relativity \\cite{1,2}. They treated the influence of the gravitomagnetic field on the particle orbit via the Lagrange planetary equations and showed by means of linear perturbation theory that---when averaged over the fast Keplerian motion---the orbit keeps its shape and slowly precesses. Specifically, they found that the semimajor axis $a$ of the osculating ellipse does not change, while its eccentricity $e$ and orbital inclination $i$ contain periodic terms that vanish on the average; moreover, the osculating ellipse precesses.  This occurs both within the orbital plane of the osculating ellipse and without, as the orbital plane precesses about the axis of rotation of the central source. The frequency of both precessions can be described by \\begin{equation}\\label{eq:1} \\boldsymbol{\\omega}_{LT}=\\frac{2G}{a^3(1-e^2)^{3/2}} [\\mathbf{J}_0-3 (\\mathbf{J}_0\\cdot \\hat{\\mathbf{n}}) \\hat{\\mathbf{n}} ].\\end{equation} Here, $\\mathbf{J}_0$ is the constant angular momentum of the source, $\\hat{\\mathbf{n}}$ is a unit vector parallel to the orbital angular momentum of the osculating ellipse and $\\hat{\\mathbf{J}}_0 \\cdot \\hat{\\mathbf{n}}=\\cos i$. Thus the Runge-Lenz vector and the orbital angular momentum vector of the osculating ellipse both precess with the Lense-Thirring frequency \\eref{eq:1}.  Astronomical bodies in general rotate; however, the magnitude of the proper angular momentum is seldom constant. In two recent papers \\cite{3,4}, the gravitational physics around a rotating central source whose spin angular momentum vector is fixed in space but varies linearly in time has been explored. In particular, it has been shown in \\cite{3} that sufficiently far from such a source, the spacetime metric is given by \\begin{equation}\\label{eq:2} ds^2=-c^2 \\left( 1-2\\frac{\\Phi}{c^2}\\right) dt^2-\\frac{4}{c}(\\mathbf{A}\\cdot d\\mathbf{x}) dt+\\left( 1+2\\frac{\\Phi}{c^2} \\right) \\delta_{ij}dx^idx^j, \\end{equation} where \\begin{equation}\\label{eq:3} \\Phi =\\frac{GM}{r},\\quad \\mathbf{A}=\\frac{G}{c} \\frac{\\mathbf{J}(t)\\times \\mathbf{x}}{r^3} \\end{equation} are the gravitoelectric and gravitomagnetic potentials, respectively. Here, $r=|\\mathbf{x}|$, $M$ is the mass of the source and its angular momentum is given by \\begin{equation}\\label{eq:4} \\mathbf{J}(t)=(J_0+J_1t)\\hat{\\mathbf{z}}; \\end{equation} moreover, $\\Phi\\ll c^2$ and $|\\mathbf{A}|\\ll c^2$. Thus $r\\gg GM/c^2$, $r\\gg J/(Mc)$ and all $O(c^{-4})$ contributions to the metric tensor have been neglected in this linear post-Newtonian approach to general relativity. As explained in \\cite{3,4}, we simply ignore the processes by which the variation of angular momentum is turned on and off and assume that equation \\eref{eq:4} holds throughout the temporal interval of interest; furthermore, all radiative effects are neglected.  The motion of a free test particle in the gravitational field of the source is given by the geodesic equation in the spacetime with metric \\eref{eq:2}. This equation, as shown in \\cite{4}, can be written in its reduced form \\begin{eqnarray}\\label{eq:5} \\nonumber \\fl \\frac{d\\mathbf{v}}{dt} +\\frac{GM\\mathbf{x}}{r^3} =& \\frac{GM}{c^2r^3} [4(\\mathbf{x}\\cdot \\mathbf{v} )\\mathbf{v}-v^2\\mathbf{x}] + \\frac{2G}{c^2} \\frac{\\dot{\\mathbf{J}} \\times \\mathbf{x}}{r^3} -\\frac{2}{c}\\mathbf{v} \\times \\mathbf{B}\\\\ &{}-\\frac{6GJ(t)}{c^4r^5} [\\hat{\\mathbf{J}} \\cdot (\\mathbf{x}\\times \\mathbf{v})] (\\mathbf{x}\\cdot \\mathbf{v})\\mathbf{v}, \\end{eqnarray} where an overdot represents differentiation with respect to time $t$ and $\\mathbf{B}=\\boldsymbol{\\nabla}\\times \\mathbf{A}$ is the gravitomagnetic field given by \\begin{equation}\\label{eq:6} \\mathbf{B}=\\frac{G(J_0+J_1t)}{cr^5} (3z\\mathbf{x} -r^2\\hat{\\mathbf{z}}). \\end{equation}  The right-hand side of \\eref{eq:5} contains all of the linear post-Newtonian contributions that arise from potentials given in \\eref{eq:3}. It turns out, however, that in a general treatment to $O(c^{-2})$, the nonlinear gravitoelectric term $4G^2M^2\\mathbf{x}/(c^2r^4)$, which is quadratic in $\\Phi$ and hence absent in our linear treatment, should also be taken into account. In the present work, we explore further the influence of the temporal variation of $J$ on motion around a central rotating source to first post-Newtonian order, namely, $O(c^{-2})$. Thus instead of \\eref{eq:5}, we consider \\begin{eqnarray}\\label{eq:7} \\frac{d\\mathbf{v}}{dt}+\\frac{GM\\mathbf{x}}{r^3} =\\mathbf{F},\\\\ \\mathbf{F} =\\frac{GM}{c^2r^3} [4(\\mathbf{x}\\cdot \\mathbf{v})\\mathbf{v}-v^2\\mathbf{x}] + \\frac{4G^2M^2}{c^2r^4} \\mathbf{x}+\\frac{2G}{c^2} \\frac{\\dot{\\mathbf{J}}\\times \\mathbf{x}}{r^3} -\\frac{2}{c}\\mathbf{v}\\times \\mathbf{B}. \\label{eq:8}\\end{eqnarray}  As demonstrated in \\cite{3,4}, equation \\eref{eq:2} represents the metric of a nonstationary linearized Kerr spacetime. The geodesic equation in Kerr spacetime is completely integrable \\cite{5}; more recent results are contained, for instance, in \\cite{6} and references therein. ",
        "conclusions": "\\label{s4}  We have studied the instability of bound Keplerian orbits induced by a time-varying gravitomagnetic field in the post-Newtonian approximation. Circular and elliptical orbits have been treated separately in sections~\\ref{s2} and \\ref{s3}, respectively. The results are expected to be of interest in the study of variable collapsed astrophysical systems.  \\ack C. Chicone was supported in part by the grant NSF/DMS-0604331."
    },
    "0806/0806.2585_arXiv.txt": {
        "abstract": "s{MOND-- modified Newtonian dynamics-- may be viewed as an algorithm for calculating the distribution of force in an astronomical object from the observed distribution of baryonic matter.  The fact that it works for galaxies is quite problematic for Cold Dark Matter.  Moreover, MOND explains or subsumes systematic aspects of galaxy photometry and kinematics-- aspects that CDM does not address or gets wrong.  I will present evidence here in support of these assertions and claim that this is effectively a falsification of dark matter that is dynamically important on the scale of galaxies. } ",
        "introduction": "Modified Newtonian dynamics, or MOND, was proposed by Milgrom~\\cite{mil83} as an alternative to dark matter.  Over the past 25 years a considerable lore has grown up around this idea, and now the very word seems to provoke strong reactions-- pro or con-- depending upon ones preconceptions or inclinations.  Here I want to provide a minimalist definition of MOND-- a definition which is as free as possible from emotive charge of the idea; therefore, I will avoid terms like modified inertia or modified gravity.  {\\it MOND is an algorithm that permits one to calculate the distribution of force in an object from the observed distribution of baryonic matter with only one additional fixed parameter having units of acceleration.}  This algorithm works very well on the scale of galaxies.  The fact that it works is problematic for Cold Dark Matter (CDM), because this is not something that dark matter can naturally do.  Moreover, MOND explains or subsumes systematic aspects of galaxy photometry and kinematics-- aspects which CDM does not address or gets wrong.  Several of these systematics were not evident at the time that MOND was proposed, so this constitutes a predictive power going beyond the ability to explain observations {\\it a posteriori}.  I will present the evidence in favor of these assertions, so this will be a discussion primarily of the phenomenology.  I will, however, draw the conclusion which to me is also minimal and quite obvious: standard CDM is falsified by the existence of this successful algorithm. ",
        "conclusions": "Although eq.\\ 1 predicts the detailed distribution of force in galaxies from the observed distribution of baryonic matter, it appears to break down in clusters of galaxies.  Applying the MOND formula in the hydrostatic gas equation, we find that, for X-ray emitting clusters, MOND reduces the mass discrepancy by a factor of two, but there still remains a factor of two or three more mass than is directly observed in hot gas and stars in galaxies.  Formally, this is not a falsification because we may always find more mass in clusters (it would be a falsification if MOND predicted {\\it less} mass than is observed), but this is seen by some as devastating for a proposed alternative to dark matter.  I take quite the opposite point of view.  The existence of an algorithm which precisely predicts the force in galaxies from the observed distribution of baryonic matter is devastating for dark matter which clusters on the scale of galaxies, CDM. In fact, it constitutes a falsification.  To explain the MOND phenomenology with dark matter would require an intimate dark matter-baryon coupling which is totally at odds with the proposed nature of CDM.  Baryons behave quite differently from CDM: they dissipate and collapse to the center of a system; they are blown out by supernovae; they are left behind in collisions between galaxies or clusters of galaxies. The intimate connection of dark matter and baryons implied by the phenomenology of rotation curves is incomprehensible in terms of CDM.  In the context of CDM, global scaling relations, such as the Tully-Fisher or Faber-Jackson relation, have their origin in aspects of galaxy formation.  Yet, galaxy formation, as emphasized by Milgrom~\\cite{mil08}, is quite a haphazard process with each galaxy having its own unique history of formation, interaction, and evolution.  It is difficult to imagine that the ratio of baryonic to dark mass would be a constant in galaxies, or even vary systematically with galaxy mass.  And yet this is required, in a very precise way, to explain the baryonic Tully-Fisher relation-- an exact correlation between the baryonic mass and the asymptotic rotation velocity which supposedly is a property of the dark matter halo. Any initial intrinsic velocity-mass relation of proto-galaxies would surely be erased in the stochastic process of galaxy formation.  To believe that vague processes such as ``feedback'' or ``self-regulation'' can restore even tighter correlations is equivalent to faith in the tooth fairy.  And finally, there is the ubiquitous appearance of $a_0\\approx cH_0$. How do CDM halos, which embody no intrinsic acceleration scale, account for the facts that $a_0$ is the acceleration at which the discrepancy appears in galaxies, that $a_0$ determines the normalization of the Tully-Fisher relation for spiral galaxies and the Faber-Jackson relation for hot systems, that $a_0$ is the characteristic internal acceleration of spheroidal systems ranging from sub-galactic objects to clusters of galaxies, that $a_0$ defines a critical surface brightness below which the discrepancy is present.  This body of evidence cannot be ignored and constitutes a profound case against CDM.  Moreover, this phenomenology implies that there is something essentially correct about MOND.  Although I have avoided the subject here, the implications are far-reaching."
    },
    "0806/0806.0852_arXiv.txt": {
        "abstract": "{The majority of observed mass-to-light ratios of globular clusters are too low to be explained by `canonical' cluster models, in which dynamical effects are not accounted for. Moreover, these models do not reproduce a recently reported trend of increasing $M/L$ with cluster mass, but instead predict mass-to-light ratios that are independent of cluster mass for a fixed age and metallicity.} {This study aims to explain the $M/L$ of globular clusters in four galaxies by including stellar evolution, stellar remnants, and the preferential loss of low-mass stars due to {{energy equipartition}}.} {Analytical cluster models are applied that account for stellar evolution and dynamical cluster dissolution to samples of globular clusters in Cen A, the Milky Way, M31 and the LMC. The models include stellar remnants and cover metallicities in the range $Z=0.0004$---$0.05$.} {Both the low observed mass-to-light ratios and the trend of increasing $M/L$ with cluster mass can be reproduced by including the preferential loss of low-mass stars, {{assuming}} reasonable values for the dissolution timescale. This leads to a mass-dependent $M/L$ evolution and increases the explained percentage of the {{observations}} from 39\\% to 92\\%.} {This study shows that the hitherto unexplained discrepancy between observations and models of the mass-to-light ratios of globular clusters can be explained by dynamical effects{{, provided that the globular clusters exhibiting low $M/L$ have dissolution timescales within the ranges assumed in this Letter}}. Furthermore, it substantiates that $M/L$ cannot be assumed to be constant with mass at fixed age and metallicity.} ",
        "introduction": "\\label{sec:intro} The mass-to-light ratios of globular clusters (GCs) have been given a lot of attention recently \\citep[e.g.,][]{mclaughlin05,rejkuba07,mieske08a,dabringhausen08}. \\citet{rejkuba07} have observed an $M/L$ trend with cluster mass above a certain cluster mass, as more massive clusters appear to have higher $M/L$ than low-mass clusters \\citep[see also][]{mandushev91}. This is an observation contrary to fundamental plane studies of GCs \\citep[e.g.,][]{mclaughlin00} and also in strong disagreement with the constant $M/L$ for fixed age that is commonly assumed in observational and theoretical GC studies \\citep[e.g.,][]{harris06,mora07,bekki07}. Moreover, for Galactic GCs \\citet{mclaughlin00} find $M/L_V=1.45~\\ml$, whereas Simple Stellar Population models \\citep[e.g.,][]{bruzual03,andersfritze03} predict significantly higher values of $M/L_V=2$---$4~\\ml$ for typical GC metallicities. Given the important role of GCs in galactic astronomy, it is essential to explain these apparent contradictions.  In numerical and analytical studies of dynamical effects in clusters \\citep[e.g.,][]{baumgardt03,lamers06,kruijssen08} it has become clear that the dynamical evolution of clusters strongly affects cluster luminosity, colour and mass-to-light ratio. In \\citet[hereafter KL08]{kruijssen08} it is shown how the evolution of these observables changes due to dynamical effects such as the preferential loss of low-mass stars and the retain of stellar remnants, but also due to the stellar initial mass function and metallicity. It is shown that $M/L$ cannot be assumed to be constant for a fixed age and metallicity, but {instead} depends on cluster mass when dynamical effects are accounted for.  In this Letter, the analytical cluster models from KL08 are applied to explain the observations of GCs in several galaxies from \\citet{rejkuba07} and \\citet{mieske08}. In Sect.~\\ref{sec:model} I first summarise the models presented in KL08, which is applied to the observations in Sect.~\\ref{sec:rejkuba}. In Sect.~\\ref{sec:mlplane} the effect of metallicity and the cluster dissolution timescale on cluster evolution in the \\{$M,M/L_V$\\}-plane is investigated. The observations are discussed in Sect.~\\ref{sec:obs} and are compared to the models in Sect.~\\ref{sec:expl}. A discussion of the results and the conclusions are presented in Sect.~\\ref{sec:concl}. ",
        "conclusions": "\\label{sec:concl} In this Letter, I have shown that the hitherto unexplained discrepancy between observations and models of the mass-to-light ratios of globular clusters can be explained by dynamical effects. The preferential loss of low-mass stars {due to energy equipartition} gives rise to $M/L$ evolution that depends on cluster mass, contrary to what is assumed in canonical cluster models. This is confirmed by the application of models that include dynamical effects to the GC populations of Cen A, the Milky Way, M31 and the LMC.  Without {the preferential loss of low-mass stars}, current stellar population models cannot explain mass-to-light ratios below 2~\\msun~L$^{-1}_\\odot$ for metallicity $Z=0.0004$ and below 2.8~\\msun~L$^{-1}_\\odot$ for $Z=0.004$. As becomes clear from Fig.~\\ref{fig:expl}, this would leave half of the cluster sample in Cen A and most of the Milky Way sample unexplained. Accounting for the effects of {energy equipartition} increases the explained percentage of the observations from 39\\% to 92\\%.  The dissolution timescales required to explain the observed GC samples lie within the physically reasonable range of $t_0=10^5$---10$^8$~yr. Observed trends of decreasing dissolution timescale with galaxy mass and metallicity are as expected when considering the strength of tidal dissolution and the radial metallicity gradient in galaxies.  The dependence of $M/L$ on cluster mass (and thus on luminosity) implies that photometrically derived masses using canonical models may be strongly overestimated (KL08). The results presented here underline the importance of accounting for dynamical effects when modeling clusters or interpreting observations of (globular) clusters."
    },
    "0806/0806.0255_arXiv.txt": {
        "abstract": "We present a novel design of a waveguide to microstrip or coplanar waveguide transition using a unilateral finline taper. The transition from the unilateral finline mode to the TEM microstrip mode is done directly, avoiding the antipodal finline tapers that have commonly been employed. This results in significant simplification of the design and fabrication, and shortening of the chip length, thereby reducing insertion loss. In this paper we shall present designs at 90~GHz that can be employed in superconducting tunnel junction mixers or Transition Edge Sensor bolometers, and scale-model measurements at 15~GHz. ",
        "introduction": "A high performance astronomical millimetre wave receiver consists of an array of horns which couples power from the sky to cryogenic detectors. The detectors are usually fabricated in superconducting planar circuits whose components are in most cases miniature microstrip lines, hence an efficient transition from waveguide to microstrip is needed. In previous publications \\citep{Yassin:1997,Yassin:2000,Kittara:2004} we have reported the successful operation of SIS (Superconductor-Insulator-Superconductor) mixers at frequencies ranging from 220-700~GHz using antipodal finline tapers. The antipodal finline taper \\citep{Yassin:2000} transforms the waveguide mode into the TEM microstrip mode using overlapping superconducting Nb films, separated by 400~nm of SiO oxide. The taper is deposited on a $\\sim100$~$\\mu$m quartz substrate which supports the structure in the E-plane of a rectangular waveguide. Before the fins overlap, the taper acts as a unilateral since the oxide is much thinner than the quartz substrate. When the fins start to overlap it behaves like an antipodal finline, and when the overlap becomes larger than the oxide thickness the transition to microstrip is performed using a semicircular taper (see Fig. \\ref{fig:combined_pic}). The supporting substrate, usually quartz or silicon, has a relatively high dielectric constant; here the matching between the unloaded waveguide and waveguide loaded with substrate is important. Broad band matching can be achieved by a 2-step reduction in the substrate width, shown in the photographs in Fig. \\ref{fig:combined_pic}. This transformer can be optimised to give return losses of 15-20~dB across a wide ($\\>30$\\%) band. As we shall see later, the performance of the taper is largely determined by the substrate mismatch.  \\begin{figure} \\caption[Photograph of scale model replicas of the antipodal (top) and direct coupling to microstrip (bottom) transitions]{Photograph of scale model replicas of the antipodal (top) and direct coupling to microstrip (bottom) transitions} \\centering \\includegraphics[width=8.6cm]{CombinedPhoto2.jpg} \\label{fig:combined_pic} \\end{figure}  Finline tapers have several advantages, including broad band operation and ease of detector block fabrication. The substrate dimensions are large relative to the microstrip width, which allows elegant integration of additional circuits of the receiver on a single chip. For example, a microstrip line with a 400~nm insulating layer of SiO2 and an impedance of approximately 20~$\\Omega$ is 3~$\\mu$m wide \\citep{Yassin:1995}, much smaller than the width of the substrate. The length of the taper is at least one free space wavelength long, so important circuits such as bandpass filters, balanced and image separating mixer circuits can easily be integrated \\citep{Kerr:1996}. Recently, this design was used in conjunction with Transition Edge Sensors (TES) and delivered a coupling efficiency well above 90\\% \\citep{Audley:2008}. It is however evident that the antipodal section with overlapping fins is difficult both to design and to fabricate, particularly when the lateral separation between the fins is very small, since at that point the field is significantly influenced by both the oxide that separates the fins and the supporting substrate. This makes the computation complicated and requires a large amount of memory. We have also learned that a lot of care is needed when fabricating the overlapping fin sections in order to avoid shorts between the very closely spaced fins as they begin to overlap. ",
        "conclusions": "We have presented a new type of finline transition from waveguide to microstrip or to CPW. The transition is particularly suited to millimetre wave detector applications where the lateral dimensions become very close to the limits of what can be fabricated using standard photolithography. The transition to microstrip was greatly simplified by replacing the overlapping fins section with direct coupling from slotline to microstrip. In the case of transition to CPW, no additional layers of deposition are required. Scale model measurements agree well with simulated results."
    },
    "0806/0806.2911_arXiv.txt": {
        "abstract": "We investigate the sensitivity of the Gamma-ray Large Area Space Telescope (GLAST)  to indirectly detect weakly interacting massive particles (WIMPs) through the $\\gamma$-ray signal that their pair annihilation produces. WIMPs are among the favorite candidates to explain the compelling evidence that about 80\\% of the mass in the Universe is non-baryonic dark matter (DM). They are serendipitously motivated by various extensions of the standard model of particle physics such as Supersymmetry and Universal Extra Dimensions (UED). With its unprecedented sensitivity and its very large energy range (20 MeV to more than 300 GeV) the main instrument on board the GLAST satellite, the Large Area Telescope (LAT), will open a new window of discovery. As our estimates show, the LAT will be able to detect an indirect DM signature for a large class of WIMP models given a cuspy profile for the DM distribution. Using the current state of the art Monte Carlo and event reconstruction software developed within the LAT collaboration, we present preliminary sensitivity studies for several possible sources inside and outside the Galaxy. We also discuss the potential of the LAT to detect UED via the electron/positron channel. Diffuse background modeling and other background issues that will be important in setting limits or seeing a signal are presented. ",
        "introduction": "The Gamma-ray Large Area Space Telescope (GLAST) \\cite{Atwood:1993zn,Michelson:1999,Michelson:2007zz,Meegan:2007zz} is a satellite-borne \\gray \\ detector launched on 11 June 2008. The Large Area Telescope (LAT) is the main instrument on GLAST, which will also host the Gamma-ray Burst Monitor (GBM) \\footnote{For a summary of overall design properties, general performance, planned operation modes, data analysis tools and strategies of the LAT, as well as a full list of participating scientists, see the web site of the collaboration \\protect\\cite{glast_website}.}. In this paper, we investigate the potential of the LAT for confirming, or constraining, the most interesting models of particle dark matter  of the Universe, WIMP (weakly interacting massive particle) models. The model of the lightest supersymmetric particle is the most studied template, which we use for most of our analysis. For general reviews on supersymmetric and other models of particle candidates for dark matter, see \\cite{reviews}. We also study some aspects of a completely different class of models, so-called Kaluza-Klein models of Universal Extra Dimensions (see \\cite{KK} and references therein). Although estimates of \\gray \\ signals exist in the literature (see \\cite{reviews}), the new feature in the present paper is that the theoretical predictions are fed through the experimental response function of the LAT and the software analysis chain that will be used for the actual data analysis after launch. Thus we give  good estimates for the potential of GLAST for detecting or limiting dark matter models. We also discuss the most important astrophysical and instrumental backgrounds.\\\\  \\noindent The LAT will have more than an order of magnitude better sensitivity in the 20 MeV to 10 GeV region than its predecessor, EGRET onboard the Compton Gamma-ray Observatory \\cite{EGRET}, and furthermore will extend the high energy region to roughly 300 GeV. Therefore, the LAT emerges as an instrument that is well suited to search for signals from dark matter  annihilation, which in the case of WIMPs should be populate just this  energy range. As the Large Hadron Collider (LHC) at CERN will also start taking data by the end of 2008, there is a non negligible probability to detect a good WIMP candidate, and verify through the \\gray \\ signal that such particles constitute the dark matter halo of the Milky Way (or neighboring galaxies or sub-halos, which we also treat). For a thorough discussion of the interplay between discovery at LHC and detection of dark matter through other methods, in particular \\gray \\ observations, see \\cite{baltz}.\\\\   \\noindent It is now established beyond reasonable doubt from a combined study of the cosmic microwave background radiation \\cite{Spergel:2006hy}, supernova cosmology \\cite{sn} and large Galaxy redshift surveys \\cite{Tegmark:2006az,Sanchez:2005pi} that non-baryonic dark matter is needed, while other models such as modifications of the laws of gravity have problems explaining the wealth of observations that now are in place. For example, it seems that the combined X-ray and optical observations of the ``Bullet Cluster'' essentially exclude explanations not involving dark matter \\cite{Bradac:2006er}.\\\\  \\noindent Technically denoted CDM, Cold Dark Matter, the particles constituting the cosmologically required dark matter have to be moving non-relativistically at the epoch of structure formation to reproduce the observed structure of the Universe, especially at small scales. This property is always fulfilled by particles with masses in the GeV to TeV range that interact with the weak interaction strength, i.e. WIMPs. They will have velocities which have redshifted since the time of thermal decoupling in the early Universe, and they will now move with typical Galactic velocities $v/c\\sim 10^{-3}$ in the Milky Way halo. This is in contrast with massive neutrinos, Hot Dark Matter, which give an observationally disfavored top-down structure formation scenario with relatively little structure on small scales. In fact, the agreement of the cosmological power spectrum with that of CDM allows us to put stringent bounds on the neutrino mass (see, e.g, \\cite{Hannestad:2003xv}). On the other hand, the non-zero neutrino masses inferred from measured neutrino oscillations, although not enough to explain more than a few percent of the dark matter, constitute a first demonstration that non-baryonic dark matter indeed is likely to exist.\\\\  \\noindent We thus have an excellent class of particle candidates, WIMPs, which behaves as the cosmologically needed CDM, and which could have been thermally produced in the early Universe to give the required relic density \\cite{Spergel:2006hy} $\\Omega_{CDM}h^2\\sim 0.1$, where $h$ is the Hubble constant in units of $100$ kms$^{-1}$Mpc$^{-1}$. The ability of WIMPs to give the measured relic density from readily computed  thermal processes in the early Universe without much fine tuning is sometimes termed  the ``WIMP miracle''.\\\\  \\noindent Of course it has to be kept in mind that dark matter does not necessarily have to be WIMPs in the mass range detectable by GLAST. The Warm Dark Matter model with thermal-relic particle masses above 2 keV may also explain the present observations in both sky surveys and N-body simulations. In the case when, say, a gravitino is the LSP, no signal would be observable with GLAST. Here we work, however, with the assumption that the dark matter particle has detectable \\gray \\ couplings and present the discovery potential for GLAST.\\\\    \\noindent In Figure 1, we sketch how \\gray s are produced from DM annihilation, and also show a schematic of a typical simulation and analysis chain as used in this paper. Table \\ref{tab:searches} shows the various approaches to the search for WIMP dark matter signals in \\gray\\ data explored in this paper. A ``smoking gun'' would be the detection of line emission in WIMP annihilation \\cite{Bergstrom:1997fj}, through the loop-induced annihilation into two photons, which for slow-moving dark matter particles would give rise to a striking, almost monoenergetic, photon signal.  However, the branching ratio for the annihilation into lines is typically about $10^{-3}$ or less in most models, as WIMPs turn out  to be electrically neutral, and thus do not couple directly to photons. There are, however, some exceptions to this estimate \\cite{Gustafsson:2007pc}. In addition to considering signatures in \\gray s we also illustrate the LAT capability to detect electron/positrons, which could provide signatures of Kaluza-Klein particle Dark Matter.\\\\       \\begin{table}[t] \\centering \\caption{ \\it The various venues GLAST will explore in its search for WIMPs, and the advantage/disadvantage of each method.} \\vskip 0.1 in \\begin{tabular}{|l|l|l|} \\hline Search  &  Advantages & Challenges \\\\ &             &            \\\\ \\hline \\hline Galactic     & Good            & Source   confusions\\\\ center      & statistics       & Uncertainty in  \\\\ &                  & diffuse background prediction \\\\ \\hline Satellites & Low background,            & Astrophysical           \\\\ & good source identification & Uncertainties  \\\\   \\hline Galactic   & Very good           & Uncertainties            \\\\ halo       & statistics      & in Galactic diffuse      \\\\ &                 & background prediction    \\\\ \\hline Extra      & very good          & Uncertainties in Galactic \\\\ galactic   & statistics      & diffuse contribution  \\\\ &                 & Astrophysical        \\\\ &                 & uncertainties        \\\\ \\hline Spectral   & No astrophysical & Potentially low         \\\\ lines      & uncertainties    & statistics  \\\\ & ``Smoking gun'' signal  &       \\\\   \\hline \\end{tabular} \\label{tab:searches} \\end{table}  \\begin{figure} \\begin{center} \\includegraphics[height=8.5cm,width=7.5cm]{Conrad_fig1.eps} \\includegraphics[height=5.5cm,width=7.0cm]{Conrad_fig1b.eps} \\caption{A diagrammatic flow of how gamma rays are produced by annihilation of dark matter and elements of the analysis chain used by the GLAST collaboration to detect them. The double question mark in the simulation chain indicates high uncertainty in the models of dark matter density and the new particle theories discussed in the paper. The single question mark over the cosmic ray propagation and interaction models indicates lesser, although significant, uncertainty in those models that generate backgrounds to the potential dark matter gamma ray signal. In this paper {\\sffamily GALPROP} ({\\sref{sec:bckgd}})  is used to estimate those backgrounds. In the next step, \\gray \\ detection is simulated using standard detector simulation packages (GEANT 4). Finally,these simulated LAT events are treated by various analysis software programs (event reconstruction and statistical analysis) to generate the results presented in this work. The same procedure is applied to the smoking gun signal of $\\chi \\chi \\rightarrow \\gamma \\gamma$, except that in this case hadronization does not have to be taken into account.} \\label{schematics} \\end{center} \\end{figure}   \\noindent The  paper is organized as follows: In \\sref{sec:LAT} we give a description of the LAT instrument and the software used for the analyses presented in this paper. In \\sref{sec:calculation} the calculations of the WIMP signal flux are described and a discussion of the considered \\gray \\ background is given. In \\sref{sec:sens} we summarize the sensitivities to generic WIMP annihilation signals achievable by the LAT for the search channels currently pursued by the LAT collaboration. Specific models are studied in \\sref{sec:specific}.\\\\ ",
        "conclusions": "Using the current state-of-the-art Monte Carlo and event reconstruction software developed within the LAT collaboration, we present preliminary  sensitivity calculations for several astrophysical searches of DM annihilation. In particular, we present sensitivities for detecting DM annihilation from the GC, Galactic and extragalactic diffuse emission, Galactic known and unknown satellites, point sources and a dedicated search for the 2$\\gamma$ line signal.  We exemplify the possibility to constrain specific particle physics scenarios (especially within in mSUGRA and UED frameworks) on the search for DM annihilation \\gray \\ emission from the GC and by using the LAT not as a \\gray \\ detector, but employing its very good electron/positron detection capabilities at very high energy.\\\\  \\noindent The sensitivities presented here are based on analyses which are idealized in the sense that systematic uncertainties in the instrument performance estimates are neglected and that uncertainties in the background estimate are quantitatively taken into account only  in a preliminary manner. For a given particle physics model, the dominant uncertainty in the sensitivities presented here is due to the lack of knowledge on the DM density distribution. Depending on which DM halo profile and/or substructure is assumed sensitivities can easily vary by one or two orders of magnitude. For most examples we consider ``standard'' DM halo structure, i.e. we are conservative in the sense that we do not include density enhancements that might boost the expected annihilation signals.\\\\  \\noindent Using simplifying assumptions, the uncertainty in the background prediction (most crucial in the search for DM signal from the halo) leads to a decrease in sensitivity by between 25 \\% and 45 \\% depending on the mass of the WIMP. For the GC analysis, the  dominant background will be from sources in the vicinity of the GC, removal of which will lead to systematic uncertainties. Assessing those without the GLAST data does not make much sense, given uncertainties in extrapolating into the GLAST energy region. For the diffuse extragalactic and high-latitude source searches the charged particle background and uncertainties therein are potentially important. Prior to launch, the levels of the charged particle background and its uncertainty are very difficult to assess. Estimates have to rely on poorly constrained Monte Carlo simulations. The charged particle background included at the level of roughly 10 \\% of the extragalactic background, which is in compliance with specifications, leads to about  a 10 \\% decrease in sensitivity for a signal in the EGRB and  in the number of detectable satellites. Uncertainties in the charged particle background are negligible for the sensitivity for both signals (assuming they are at the level of $\\sim$ 20 \\%).\\\\   \\noindent For the Galactic and extragalactic diffuse background the range of possible backgrounds is illustrated by assuming several representative models, which are discussed in detail. The sensitivities are preliminary in the sense that estimates of the instrument performance, analysis methods and estimates of the expected backgrounds are being continuously improved. In context of specific particle physics scenarios, also the estimates of the signals are continuously updated: For example, calculations of the \\gray \\ flux for Supersymmetric Dark Matter annihilation incorporating QED corrections, indicate that for part of the parameter space \\gray \\ yields might be boosted by three or four orders of magnitude and and lead to distinct spectral signatures \\cite{Bringmann:2007nk}.\\\\   \\noindent We conclude that the LAT has good potential to discover DM annihilation for a significant fraction of interesting parameter space, i.e. for values of annihilation cross-section of between $<\\sigma v> \\simeq 10^{-26}$ cm$^3$s$^{-1}$  and $<\\sigma v> \\simeq 10^{-24}$ cm$^3$s$^{-1}$ depending on WIMP masses in the range between 40 and 500 GeV. For less conservative assumptions on the Dark Matter density (for example additional substructure or a Moore profile) the sensitivity improves by one to two order of magnitudes. Our conclusions are consistent with previous work that employed cruder representations of the GLAST response  and/or less thorough treatments of the backgrounds (see e.g \\cite{reviews},\\cite{Bergstrom:1997fj}, \\cite{Dodelson:2007gd} and references therein). If indeed a significant DM signal is present, GLAST will be able to image the DM structure in our Galaxy.\\\\     \\ack{ We thank the members of the LAT collaboration for providing many interesting discussions and elements of the software used in this analysis. Sergio Colafrancesco, Steve Ritz and Julie McEnery are thanked for careful reading of the manuscript. Numerous discussions with Jeff Scargle are gratefully acknowledged. This work was supported by the U.S. Department of Energy contract number DE-AC02-76SF00515, the U.S. Department of Energy contract number, DE-FG02-91ER40690,  NASA, Vetenskapsr{\\aa}det and the Swedish Space Board. IM acknowledges partial support by the NASA APRA program. The LAT is being developed by an international collaboration with primary hardware and software responsibilities at (in alphabetical order) The Agenzia Spaziale Italiana, Centre National de la Recherche Scientifique / Institut National de Physique Nucl\\'eaire et de Physique des Particules, Commissariat \\`a l'Energie Atomique, Goddard Space Flight Center, Hiroshima University, Istituto Nazionale di Astrofisica, Istituto Nazionale di Fisica Nucleare, Naval Research Laboratory, Ohio State University, Kalmar University, Royal Institute of Technology - Stockholm, Stanford Linear Accelerator Center, Stanford University, University of California at Santa Cruz, and University of Washington. The LAT project is managed by the Stanford Linear Accelerator Center, which is also the location of the Instrument Science Operations Center, and the LAT Principal Investigator is Peter Michelson at Stanford University.Other institutions that have made significant contributions to  the instrument development include the Institute of Space and Astronautical Science, Stockholm University, University of Tokyo, and Tokyo Institute of Science and Technology.\\\\   }"
    },
    "0806/0806.1229_arXiv.txt": {
        "abstract": "Two consecutive transits of planetary companion OGLE-TR-111b were observed in the I band.  Combining these observations with data from the literature, we find that the timing of the transits cannot be explained by a constant period, and that the observed variations cannot be originated by the presence of a satellite. However, a perturbing planet with the mass of the Earth in an exterior orbit could explain the observations if the orbit of OGLE-TR-111b is eccentric. We also show that the eccentricity needed to explain the observations is not ruled out by the radial velocity data found in the literature. ",
        "introduction": "The observations of transiting extrasolar planets have produced some of the most interesting results in the study of other planetary systems. Their orbital configuration have permitted the first direct measurements of radius, temperature, and composition \\citep[and references therein]{swain2008,harrington2007}, all of which are critical to constraining the interior and evolution models of extrasolar planets \\citep[e.g.][]{fortney2008}.  It has been further realized that the presence of variations in the timing of transits can be attributed to otherwise undetectable planets in the system~\\citep[see, for example,][]{miralda-escude2002,holmanmurray2005,agol2005,heylgladman2007,fordholman2007,simon2007}. \\citet{deeg2008} and \\citet{ribas2008} reported indirect detections of unseen companions by monitoring eclipse timing of the binary stellar system \\object{CM~Draconis} (1.5 M$_J$ to 0.1 M$_\\odot$ candidate) and variations in the orbital parameters of the planetary system around \\object{GJ~436} (5 $M_\\oplus$ companion), respectively. However, this last case has been recently argued against by \\citet{alonso2008}. Besides, recently-discovered transiting planets \\citep{tr182,tr211} exhibiting shifts in their radial velocities are promising new candidates to search for variations in the timing of their transits. On the other hand, \\citet{steffenagol2005} found no evidence of variations in the timing of transits of the \\object{TrES-1} system, after analysing data for 12 transits. Also, after monitoring 15 transits of the star \\object{HD~209458}, \\citet{miller-ricci2008} were able to set tight limits to a second planet in the system.  Here we report a significant detection of variability in the timing of the transits of extrasolar planet \\object{OGLE-TR-111b} \\citep{udalski2002,pont2004} and discuss its possible causes, including a second unseen planet OGLE-TR-111c.  \\defcitealias{tlc111}{W07}  In a previous work \\citep{minniti2007} we reported a single transit observed in the V band which occurred around 5 minutes before the expected time obtained using the ephemeris of \\citet[][hereafter W07]{tlc111}, but the result was inconclusive since it had a 2.6-$\\sigma$ significance.  In the present work we analyse data of two consecutive follow-up transits of the same planet.  Section \\ref{datasec} presents the new data and the reduction procedures, in Section \\ref{measusec} we describe the technique used to measure the central times of the transits. Finally, in Section \\ref{ressec} we present our results and discuss their implications. ",
        "conclusions": "\\label{ressec}  \\begin{figure}[t] \\centering \\plotone{f3.eps} \\caption{Observed-minus-calculated times (in minutes) for the transits of planet OGLE-TR-111b in front of its host star. The filled circles are the new transits presented in this work, the empty circles are from \\citetalias{tlc111} and the empty square is the transit presented by \\citet{minniti2007}, which has been reprocessed for this work.} \\label{oc} \\end{figure}  We fitted a straight line to the central times of the two transits together with those from \\citetalias{tlc111} and \\citet{minniti2007}. The central time of this last transit ($T_{c,VIMOS}$) has been remeasured using the procedure described above and the result is shown in Table~\\ref{param}. In this way we obtained a new ephemeris for the transit times: \\begin{eqnarray} T_c &=& 2454092.80607 \\pm 0.00029\\; \\mathrm{[HJD]}\\\\ P &=& 4.0144540 \\pm 0.0000038\\; \\mathrm{days} \\label{eq:P}\\; \\; , \\end{eqnarray} with correlation coefficient $\\rho = 0.785$. The reduced $\\chi^2$ is 9.04, indicating a poor fit. Note that the value of the period is consistent with the value reported by \\citetalias{tlc111}. The fit was repeated including a point for the OGLE data, and we also obtained the period from a simultaneous fit to all the available photometry (OGLE, \\citetalias{tlc111}, \\citet{minniti2007} and this work). In both cases the obtained value is in excellent agreement with the one reported above.  In Fig.~\\ref{oc} we plot the residuals of the fit. It is clear that the observed-minus-computed (O-C) values are not consistent with a constant period since the VIMOS transit, one of the transits from \\citetalias{tlc111} and one of the FORS transits lie -3.29-$\\sigma$, 2.79-$\\sigma$ and -2.52-$\\sigma$ away from zero, respectively. However, the data available to date are not enough to determine the nature of these variations. Nevertheless, we have been able to discard a few possibilities and study some others. We present some preliminary results here and defer a more detailed study for a future work.  First, the hypothesis of an exomoon seems unlikely, since the mass needed to produce the observed O-C amplitude is at least one twenty-sixth of the planetary mass if the moon is at a Hill radius from the planet. However, at this distance the moon system is expected to be unstable. For moons closer to the planet, the needed mass increases. These are extreme values when compared with the Solar System, where this ratio never exceeds $2.5\\times 10^{-4}$ \\citep{allen}.  On the other hand, several planetary system configurations reproduce the observed trend. The equations of motion for the three-body problem were solved with the Bulirsch-Stoer algorithm implemented in the Mercury package \\citep{chambers99} using different sets of orbital parameters for the perturbing planet, and the results were compared with the observations. A particularly interesting solution is that an exterior Earth-mass planet near the 4:1 resonance produces the observed amplitude and periodicity in the O-C times, if the orbit of TR111b is eccentric ($e = 0.3$). On the other hand, the mass of the perturber planet must be at least around 4 $M_{Jup}$ if the orbit of the interior planet is nearly circular. This shows the importance of accurately measuring the ecentricity of the interior planet through RV data or measurements of the planet occultation \\citep[see][]{deming2007}.  \\begin{figure} \\epsscale{0.95} \\centering \\plotone{f4.eps} \\caption{Radial velocity measurements from \\citet{pont2004} together with the best fit (solid line), and the corresponding $\\pm 1\\sigma$ curves (dotted lines). Also shown is the fit for $e=0$ (dashed line).} \\label{rvfit} \\end{figure}  In the discovery paper by \\citet{pont2004}, the orbital solution was obtained by fixing the eccentricity of TR111b to zero. Although this is reasonable for a single planet in a close orbit to the star, since circularization is very effective in those conditions \\citep[see, for example,][]{zahn77}, a second planet can perturb the orbit of the first one, increasing its eccentricity. Therefore, we reanalysed the radial velocity (RV) data from \\citet{pont2004}, in order to constrain the possible eccentricity of the system. We found that the data are compatible with an eccentricity of 0.3, with a reduced $\\chi^2$ of about 0.4 (for 5 degrees of freedom, see Fig.~\\ref{rvfit}) compared to the value of 0.7 for a circular orbit, as reported in the original paper.  Additionally, note that the 1.55-$\\sigma$ difference between the transit length presented in Table~\\ref{param} and that reported by \\citetalias{tlc111} might indicate a change in the inclination angle of OGLE-TR-111b \\citep[see][]{ribas2008,miralda-escude2002} which could in principle help constrain the parameters of the perturber planet.  Future observations are warranted in order to pinpoint the origin of the variation in the period of this interesting planet."
    },
    "0806/0806.3689_arXiv.txt": {
        "abstract": "Dark matter (DM) annihilations in the Galaxy may produce high energy neutrinos, which can be detected by the neutrino telescopes, for example IceCube, ANTARES and Super-Kamiokande. The neutrinos can also arise from hadronic interaction between cosmic ray and atmosphere around the Earth, known as atmospheric neutrino. Current measurements on neutrino flux is consistent with theoretical prediction of atmospheric neutrino within the uncertainties. In this paper, by requiring that the DM annihilation neutrino flux is less than the current measurements, we obtain an upper bound on the cross section of dark matter annihilation $ \\left\\langle {\\sigma v} \\right\\rangle$. Compared with previous investigations, we improve the bound by including DM substructure contributions. In our paper, two kinds of substructure effects are scrutinized. One is the substructure average contribution over all directions. The other is point source effect by single massive sub-halo. We found that the former can improve the bound by several times, while the latter can improve the bound by $ 10^1 \\sim 10^4$ utilizing the excellent angular resolution of neutrino telescope IceCube. The exact improvement depends on the DM profile and the sub-halo concentration model. In some model, IceCube can achieve the sensitivity of $ \\left\\langle {\\sigma v} \\right\\rangle \\sim 10^{ - 26} cm^3 s^{ - 1} $. ",
        "introduction": "Many astronomical observations indicate that most of the matter in our universe is dark (see e.g. Ref. \\cite{Jungman:1995df}). The evidences come mainly from the gravitational effects of the dark matter (DM), such as the rotation curves of spiral galaxies \\cite{Begeman:1991iy,Persic:1995ru}, the gravitational lensing \\cite{Tyson95} and the dynamics of galaxy clusters \\cite{White93}. The studies such as primordial nucleosynthesis \\cite{Peebles71} and cosmic microwave background (CMB) \\cite{Spergel03} show that the DM is mostly non-baryonic. Combining recent cosmological measurements, for example from the Wilkinson Microwave Anisotropy Probe (WMAP), one could deduce precisely the relic density of DM, namely $\\Omega_{DM}h^2=0.1143\\pm 0.0034$ \\cite{Hinshaw:2008kr}. However, the nature of dark matter is still unclear. In the literature there is a ``zoo'' of particle candidates for DM \\cite{Bertone04}, among which the most popular candidate at present is the weakly interacting massive particle (WIMP) such as the lightest supersymmetric particle (LSP), lightest Kaluza-Klein particle (LKP) $etc$.  Search for WIMP in particle physics experiments is very important to pin down the properties of the DM. Besides searching missing energy signals at accelerator-based experiments, there are usually two classes of methods to detect WIMP, namely direct and indirect ones. The former method detects WIMP by measuring the recoil of heavy nucleus in the detector and gives the most strong evidence for the existence of DM. The latter one detects the DM self-annihilation signals, which include neutrinos, photons, anti-matter particles and so on. Among them neutrinos are one of the most attractive signals. For the {\\em low energy} neutrinos (say much less than 100 GeV), their interactions with matter are highly suppressed by a factor at least $Q^2/m_W^2$ with $Q$ the typical energy scale of the interaction. The neutrinos are hardly energy loss and trajectory deflection during their propagation, therefore they may carry the information of the nature and distribution of the DM. However due to the same reason, it is hard to capture such kind of {\\em low energy} and relatively low flux neutrinos. For the {\\em high energy} neutrinos (say around 100 GeV or higher), the interactions among neutrinos and matter become much stronger. These neutrinos may keep the information of the DM, and it is relatively easy to observe them experimentally.  One proposal of detecting the high energy neutrino signals is to explore the locations close to us such as the center of the Sun or the Earth to get enough neutrino flux. The DM particles are gravitational trapped in the center of the Sun or Earth and produce neutrinos by annihilation \\cite{Liu:2008kz}. If the annihilation and capture processes are in equilibrium, the neutrino flux are mainly determined by the cross section of the DM and nuclei. Another proposal is to detect the neutrino signals from DM annihilation in the Milky Way (MW). Though the sources in the MW are farther than the Sun, it is natural to expect that the larger amount of DM can compensate the distance. Moreover the neutrino flux depends on the DM annihilation rate and number density square, therefore the regions with high density in the MW, such as the Galactic Center (GC) or sub-halos, should be potentially excellent observational targets.  The GC is conventionally thought to be source-rich astrophysical laboratory and has attracted many attention of astronomers. It is also true for the DM indirect searches. Due to the weak interaction of DM particles, the DM density at the GC is highly accumulated as shown by detail simulations, which makes the GC a bright source of DM annihilation. However, the complicated astrophysical environment and various kinds of astrophysical sources  make the GC a high background site. In addition, the overlapping with bayonic matter objects (e.g., the central massive black hole \\cite{Gondolo:1999ef}) may also affect the DM distribution and increase the uncertainties. It should be emphasized that the galactic sub-halos may be good candidates as DM sources. Since the self-annihilation of DM is square-dependent on the number density, the clump of substructure is expected to effectively enhance the annihilation signal and plays a role of the so-called ``boost factor'' \\cite{bi06a,bi06b,yuan07}. Furthermore as indicated by simulations, the spatial distribution of DM sub-halos is tend to be spherical symmetric in the MW halo, which may locate at a low-background site and effectively avoid the source confusion in the galactic plane. The effects of DM sub-halos on the flux of induced neutrino is what we try to investigate in this work.  The neutrinos detected by high-energy neutrino telescopes such as Super-Kamikande \\cite{Desai:2004pq}, AMANDA \\cite{Ahrens:2003fg} etc. are thought to be mainly from atmospheric neutrinos. Here they originate from the decay of hadrons which are produced by the strong interactions of cosmic rays with atmosphere. Experimentally no obvious excess has yet been observed. Then the measurements on neutrino flux can be utilized to set bounds on the DM annihilation cross section. As neutrinos are the most difficult to detect in the DM annihilating final states, the authors of Ref. \\cite{Beacom:2006tt} and \\cite{Yuksel:2007ac} assumed that the DM annihilate solely into neutrinos. They calculated the extra-galactic and the galactic neutrino fluxes, compared them with atmospheric neutrino flux, and set an upper bound on the DM annihilation cross section $\\langle\\sigma v\\rangle$. However, the DM may annihilate into final states other than neutrinos. In most of the DM models there are several annihilation channels. Moreover high energy neutrinos from the DM annihilation will lead to gauge bosons bremsstrahlung \\cite{Kachelriess:2007aj,Bell:2008ey} even in the standard model (SM). The electromagnetic final states through higher-order corrections are also inevitable \\cite{Dent:2008qy}. Therefore the assumption that DM annihilate into only neutrinos gives the most conservative bound on DM annihilation cross section.  In this paper we calculated the neutrino flux from the DM annihilations in the MW including the contributions from sub-halos by assuming that the DM annihilate into neutrinos only \\cite{Beacom:2006tt,Yuksel:2007ac}. By comparing the predicted flux with the available atmospheric neutrino measurements, we set a very strict constraint on the DM total annihilation cross section. Compared to the previous studies, in this work we utilize the angular resolution of the neutrino telescope to derive the stricter constraints. Here the massive sub-halos can be treated as the point-like sources. Based on our analysis we may observe the high energy neutrino flux provided that precise angular resolution data from ANTARES \\cite{Pradier:2008iv} and IceCube \\cite{Ahrens:2003ix} is available. On the other hand if no excess flux out of atmospheric neutrino is observed, an improved upper-bound of the annihilation cross section and/or the exclusion of certain sub-halo models can be obtained.  This paper is organized as following. In Sec. II, we describe the sub-halo models according to the N-body simulation results. In Sec. III, we give the constraints on the dark matter annihilation cross section. The conclusions and discussions are given in Sec. IV. ",
        "conclusions": "In this paper, by requiring the DM induced neutrino flux less than the measured ones, we give the improved upper bounds on the DM annihilation cross section $\\langle\\sigma v\\rangle$ with the DM substructure effects included. Here we assume the DM particles annihilate into neutrinos solely following the previous works. The observed neutrino flux depends on the particle physical and astrophysical factors. Thus we first investigate the astrophysical factor. Several different DM profiles and sub-halo concentration models are adopted based on the numerical simulations. Our studies show that at the anti-GC direction, the enhancement factor of sub-halos for B01$\\times2$ model is about 3 and  25 for NFW and Moore profiles respectively. While the whole-sky average ($halo\\ average$, with cone half-angle $180^{\\circ}$) does not have prominent enhancement. The best case of our adopted models is only $\\sim4$ times larger than the smooth ones (see bottom-right figure in Fig. \\ref{jpsieps}). If we take the $30^{\\circ}$ angular average ($halo\\ angular$ in Ref. \\cite{Yuksel:2007ac}), there is almost no enhancement, as can be seen in Fig. \\ref{jpsieps}. This is because the enhancement from sub-halos is spatially dependent on the MW halo, instead of a universal one \\cite{bi06b}. On one hand, the smooth component increases more rapidly and dominates the annihilation flux near the GC; on the other hand, the tidal disruption on sub-halos is most effective close to the GC. Thus the effect of substructures is not significant near the GC. In the cases of $halo\\ average$ and $halo\\ angular$, the GC contributions are included and play a dominant role in the total flux, therefore no remarkable enhancement from sub-halos is found.  In this paper we emphasize the important role of the massive sub-halos (e.g., $M_{sub}>10^6$M$_{\\odot}$). Since the number of massive sub-halo is small, it should be regarded as the point-like source. For the massive sub-halos, the high angular resolution of neutrino detector can be utilized to suppress the atmospheric neutrino background, thus the constraints on $\\langle\\sigma v\\rangle$ are expected to be improved. The angular resolution $\\sim 1^{\\circ}$ for energy greater than $50$ GeV of the forthcoming experiment IceCube is employed \\footnote{ For Super-Kamiokande, the resolution angle is $\\sim 3^{\\circ}$ for $E>100$ GeV and the background is $9$ times greater.}. The neutrino signal flux in a cone with half-angle $1^{\\circ}$ is calculated and compared with the atmospheric background in the same cone. We found that the constraints on $\\langle\\sigma v\\rangle$ are indeed improved significantly. Note that the constraints are model-dependent. For the moderate case B01+NFW, we find the upper bound of $\\langle\\sigma v\\rangle$ is about $10^{-23}$ cm$^3$ s$^{-1}$. While for the concentration model B01-subhalo, the bound can reach $10^{-26}$ cm$^3$ s$^{-1}$, which is even lower than $\\langle\\sigma v\\rangle$ for the natural scale which can induce the correct relic density of DM.  It should be noted that DM can annihilate into final states other than neutrinos. Thus the assumption that DM annihilates into only neutrinos gives the most conservative bound on the DM annihilation cross section and is independent of the particle properties of DM particle.  Neutrinos are  thought to be an important complementary particles for DM indirect searches besides  $\\gamma$-rays and charged anti-particles. It is shown that the detectability of $\\gamma$-rays from DM sub-halos on GLAST is optimistic \\cite{kuhlen08}. The effects of subhalos on positrons \\cite{cumber07, lavalle07} and antiprotons \\cite{lavalle08} are also investigated. The combination and cross check of different kinds of signals will be very crucial to identify the DM sources and investigate the properties of DM particles."
    },
    "0806/0806.3640_arXiv.txt": {
        "abstract": "{}{The line profile variability and photometric variability of the O9.5\\,Vp star \\object{HD\\,93521} are examined in order to establish the properties of the non-radial pulsations in this star.}{Fourier techniques are used to characterize the modulations of the He\\,{\\sc i} $\\lambda\\lambda$\\,5876, 6678 and H\\,$\\alpha$ lines in several spectroscopic time series and to search for variations in a photometric time series.}{Our spectroscopic data confirm the existence of two periods of 1.75 and 2.89\\,hr. The line profiles, especially those affected by emission wings, exhibit also modulations on longer time scales, but these are epoch-dependent and change from line to line. Unlike previous claims, we find no unambiguous signature of the rotational period in our data, nor of a third pulsation period (corresponding to a frequency of 2.66\\,d$^{-1}$).}{HD\\,93521 very likely exhibits non-radial pulsations with periods of 1.75 and 2.89\\,hr with $l \\simeq 8 \\pm 1$ and  $l \\simeq 4 \\pm 1$ respectively. No significant signal is found in the first harmonics of these two periods. The 2.89\\,hr mode is seen at all epochs and in all lines investigated, while the visibility of the 1.75\\,hr mode is clearly epoch dependent. Whilst light variations are detected, their connection to these periodicities is not straightforward.} ",
        "introduction": "} Time resolved spectroscopy of O-type stars has shown that absorption line profile variability at the level of a few per cent is a common feature (e.g.\\ Fullerton et al.\\ \\cite{FGB}). Various mechanisms, including magnetic fields, stochastic structures at the base of the wind and non-radial pulsations have been proposed to explain this variability. Despite their rather subtle signature, the diagnostic potential of these phenomena is considerable.  Especially in the case of non-radial pulsations, the emerging discipline of asteroseismology offers the possibility to gain insight into the interiors of early-type stars. However, to characterize the nature of the phenomenon requires a rather long and well sampled time series of spectra with a high resolution and a high S/N (see e.g.\\ Aerts et al.\\,\\cite{nuEri} for the case of the $\\beta$\\,Cephei variable $\\nu$\\,Eri). Up to now, such detailed studies have therefore been restricted to a few, rather bright and well-known O-stars such as $\\zeta$\\,Pup (Baade et al.\\ \\cite{Baade}) and $\\zeta$\\,Oph (Kambe et al.\\ \\cite{Kambe}).\\\\  In this context, the high Galactic latitude O-star HD\\,93521 ($l_{\\rm II} = 183.14^{\\circ}$, $b_{\\rm II} = 62.15^{\\circ}$) is an extremely interesting target. Based on optical spectroscopy, Fullerton et al.\\ (\\cite{FGB2}) and Howarth \\& Reid (\\cite{HR}) found HD\\,93521 to exhibit bumps at the 1\\% level moving from blue to red across the profiles of the He\\,{\\sc i} lines, whilst no variability was detected in the He\\,{\\sc ii} lines. Fullerton et al.\\ (\\cite{FGB2}) as well as Howarth \\& Reid (\\cite{HR}) accordingly interpreted these features as the signature of non-radial pulsations with a period of 1.8\\,hours. The He\\,{\\sc i} -- He\\,{\\sc ii} dichotomy was interpreted as arising from the substantial gravity darkening that favours He\\,{\\sc i} line formation near the cooler equatorial regions where the pulsational amplitude attains a maximum (Townsend \\cite{Townsend}). The existence of non-radial pulsations was subsequently confirmed by Howarth et al.\\ (\\cite{HTC}) using {\\it IUE} spectra and three different periods were identified. Recently, Rzaev \\& Panchuk (\\cite{RP}) reported on the existence of slightly different variability patterns between the strong and weak He\\,{\\sc i} lines. However, since the Rzaev \\& Panchuk (\\cite{RP}) data set covered only 2.7\\,hours, it yields hardly any constraint on the properties of the pulsations.\\\\  HD\\,93521 has one of the largest rotational velocities known among O-stars (341\\,km\\,s$^{-1}$, Penny \\cite{Penny}; 432\\,km\\,s$^{-1}$, Howarth et al.\\ \\cite{HSHP}; 390\\,km\\,s$^{-1}$, see below in this paper). The stellar wind has an apparently low terminal velocity and is likely heavily distorted by rotation (e.g.\\ Bjorkman et al.\\ \\cite{Bjorkman}). In the optical, the wind produces emission features in the wings of the H$\\alpha$ line, although this line has so far never been reported as a pure emission feature. \\begin{table*} \\caption{Summary of our observing runs (see text for the meaning of the different columns). The green and red wavelength ranges stand respectively for 5500 -- 5920 and 6530 -- 6710\\,\\AA, whilst SPM stands for the echelle spectra taken at San Pedro M\\'artir. The last column yields the mean signal-to-noise ratio of the spectra gathered during the corresponding campaign. The last line refers to the photometric monitoring campaign that took place in coordination with the April 2005 spectroscopic campaign.\\label{journal}} \\begin{center} \\begin{tabular}{l c c c c c c c} \\hline \\vspace*{-3mm}\\\\ Epoch & Domain & $\\Delta{\\rm T}$ (days) & N & $\\overline{\\Delta t}$ (days) & $\\Delta\\,\\nu_{\\rm nat}$ (d$^{-1})$ & $\\nu_{\\rm max}$ (d$^{-1})$ & $\\overline{\\rm S/N}$\\\\ \\hline February 1997 & red & 4.155 & 35 & $2.49 \\times 10^{-2}$ & 0.241 & 20.1 & 165\\\\ April 2005 & red & 3.809 & 55 & $1.72 \\times 10^{-2}$ & 0.263 & 29.1 & 280 \\\\ April 2005 & SPM & 5.096 & 90 & $0.84 \\times 10^{-2}$ & 0.196 & 59.6 & 250 \\\\ February 2006 & green & 6.173 & 75 & $1.60 \\times 10^{-2}$ & 0.162 & 31.3 & 650 \\\\ April 2006 & green & 0.180 & 13 & $1.49 \\times 10^{-2}$ & 5.55 & 33.5 & 580 \\\\ April 2007 & green & 6.191 & 56 & $1.46 \\times 10^{-2}$ & 0.162& 34.2 & 550 \\\\ \\hline April 2005 & $U\\,B\\,B_1\\,B_2\\,V\\,V_1\\,G$ & 28.221 & 378 & $0.31 \\times 10^{-2}$ & $0.035$ & 161.3 & \\\\ \\hline \\end{tabular} \\end{center} \\end{table*}   \\begin{figure*}[t!] \\begin{minipage}{6.0cm} \\resizebox{6.0cm}{!}{\\includegraphics{0002f1a.ps}} \\end{minipage} \\begin{minipage}{6.0cm} \\resizebox{6.0cm}{!}{\\includegraphics{0002f1b.ps}} \\end{minipage} \\begin{minipage}{6.0cm} \\resizebox{6.0cm}{!}{\\includegraphics{0002f1c.ps}} \\end{minipage} \\caption{Average spectrum (top) and temporal variance spectrum (TVS, bottom) computed from the green, red and SPM spectra taken in February + April 2006 (left panel), April 2005 (middle) and April 2005 (right) respectively. The dotted lines yield the 99\\% significance level for the variability evaluated following the approach of Fullerton et al.\\ (\\cite{FGB}).\\label{average}} \\end{figure*}  While the optical spectrum of HD\\,93521 leads to an O9.5\\,V classification, the nature of this star has been subject to debate over many years. In fact, assuming a typical absolute magnitude for a Population I O9.5\\,V star, HD\\,93521 is located about 1.4\\,kpc above the Galactic plane (Irvine \\cite{Irvine}), far away from any site of recent massive star formation. While there is still some uncertainty concerning the motion of the star (towards or away from the Galactic plane, see Gies \\cite{Gies}), it seems unlikely that HD\\,93521 could have formed in the plane and subsequently moved to its current position. Furthermore, based on the rather low stellar wind terminal velocity of HD\\,93521 and assuming that $v_{\\infty} = 3 \\times v_{esc}$, Ebbets \\& Savage (\\cite{ES}) concluded that this star was likely a low-mass evolved Population II object. However, Irvine (\\cite{Irvine}) showed that the relation between $v_{\\infty}$ and $v_{esc}$ does not hold for late O-type stars such as HD\\,93521 and he proposed that the star is in fact a normal main sequence star that has formed in the Galactic halo (another example of a massive star born in the halo can be found in Heber et al.\\ \\cite{Heber}). Lennon et al.\\ (\\cite{Lennon}) measured the equivalent widths (EWs) of several metal lines in the spectrum of HD\\,93521. Though these lines are washed out by rotational broadening and the EWs are affected by large uncertainties, their strengths are inconsistent with Population II metal abundances. Finally, Massa (\\cite{Massa}) showed that the other peculiar features of this star (its unusually low UV continuum, its abnormally strong wind lines in the UV, its low excitation photospheric lines) can all be accounted for by the effect of gravity darkening in a `normal' Population I star rotating at 90\\% of its breakup velocity and seen nearly equator on.\\\\  In this paper, we present the results of a spectroscopic and photometric monitoring campaign of HD\\,93521. The aim of this project was to check the long-term stability of the periodicities identified in previous investigations. ",
        "conclusions": "\\subsection{General properties} A summary of the amplitudes of all the frequencies detected in the spectroscopic part of this study is provided in Table\\,\\ref{tabampl}.  A first conclusion from this table is that the $\\nu_1$ frequency is not always detected: while it clearly dominates in the April 2005 data of the He\\,{\\sc i} $\\lambda$\\,6678 line, it is absent from our February 1997 time series of the same line. Also, this frequency is never detected in the H$\\alpha$ line profile variations. In the He\\,{\\sc i} $\\lambda\\lambda$\\,5876 and 6678 lines, this frequency is associated with a maximum semiamplitude of modulation of 0.004 -- 0.006 in units of the continuum. On the other hand, $\\nu_2$ or its aliases are detected in each line and each observing campaign. The maximum semiamplitudes of this mode are 0.003 (He\\,{\\sc i} $\\lambda\\lambda$\\,5876, 6678 in 2006 and April 2005), 0.006 (H$\\alpha$, February 1997) and 0.008 (He\\,{\\sc i} $\\lambda$\\,6678, February 1997) in units of the continuum.  Another conclusion is the lack of any significant signal at the $\\nu_3 = 2.66$\\,d$^{-1}$ frequency ($P_3 = 9.0 \\pm 1.2$\\,hr) that was reported by Howarth et al.\\ (\\cite{HTC}) from their analysis of 103 {\\it IUE} spectra with a median sampling of 0.60\\,hr. This frequency is thus apparently absent from our data, except perhaps for the detection of the 2.51\\,d$^{-1}$ signal in the April 2005 SPM H$\\alpha$ data.  Generally speaking, the (low-level) periodic variations with frequencies $\\nu_1$ and $\\nu_2$ are not the dominant source of line profile variations in any of the lines investigated here (except for the He\\,{\\sc i} $\\lambda$\\,6678 line in April 2005). In the majority of the cases, the most important variations are actually characterized by low frequencies. These low-frequency variations affect both the core of the lines as well as the wings and emission lobes (in the case of the H$\\alpha$ and He\\,{\\sc i} $\\lambda$\\,5876 lines). Apparently, these modulations do not occur with a single stable clock: none of the low frequencies is detected in more than one line and for more than one observing campaign (except perhaps for the frequency around 1.10\\,d$^{-1}$). It seems thus more appropriate to talk about time scales than periods for these longer term variations. This result casts doubt on a rotational modulation as the origin of the long-term variations: Howarth \\& Reid (\\cite{HR}) reported the analysis of 21 optical echelle spectra of HD\\,93521 acquired over two nights in February 1992. From the variations of the emission wings of the H$\\alpha$ and He\\,{\\sc i} $\\lambda$\\,5876 lines, these authors estimated a rotational period of 35\\,hr ($\\nu = 0.69$\\,d$^{-1}$). None of the lines investigated here displays a significant signal at this frequency. The closest detections are found for the 1997 OHP H$\\alpha$ data (the alias of the highest peak at 0.70\\,d$^{-1}$) and for the same line observed in 2005 (the alias of the highest peak at 0.78\\,d$^{-1}$). Therefore, our data do not confirm the existence of a rotational period of about 35\\,hr and no unambiguous rotational period can be identified observationally for this star.  In the remaining subsections we will focus on the interpretation of the variations seen at the $\\nu_1$ and $\\nu_2$ frequencies. \\subsection{$\\nu_1$ and $\\nu_2$ as non-radial pulsations} In this section, we assume that the line profile variability at the frequencies $\\nu_1$ and $\\nu_2$ is due to two different non-radial pulsation modes. For multi-mode pulsations, one expects to observe variability with the genuine pulsation frequencies and their harmonics, as well as with their sums and beat frequencies (see e.g.\\ the case of the O9.5\\,V star $\\zeta$\\,Oph, Kambe et al.\\ \\cite{Kambe}, Walker et al.\\ \\cite{Walker}). In HD\\,93521, we find no significant signal at the first harmonics of $\\nu_1$ and $\\nu_2$ nor at the sum or beat frequencies. Schrijvers \\& Telting (\\cite{ST}) argue that for non-radial pulsations with an amplitude of order 10\\% of the mean line depth, the absence of a first harmonic is an indication that the line profile variability is mainly due to temperature effects (rather than to the Doppler-redistribution of flux). In HD\\,93521, the amplitudes of the $\\nu_1$ and $\\nu_2$ modes in the He\\,{\\sc i} $\\lambda$\\,6678 line are of order 5 -- 10\\% of the maximum line strength. While these amplitudes might be somewhat low for the harmonics to be detected, we nevertheless note that temperature effects could play a significant role in the line profile variability observed in HD\\,93521.  \\begin{table*}[htb!] \\caption{Properties of the frequencies detected in the 2-D Fourier analyses of the line profile variability of HD\\,93521. The $l$ values were derived from the blue-to-red phase differences discussed in Sect.\\,\\ref{sect: fourier} using the formula of Telting \\& Schrijvers (\\cite{TS}).\\label{tabspec}} \\begin{center} \\begin{tabular}{r c c c c l} \\hline $\\nu$ (d$^{-1}$) & P (hr) & Detection & $\\Delta\\,\\phi/\\pi$ & $l$ & Comment \\\\ \\hline 14.61 & 1.64 & He\\,{\\sc i} $\\lambda$\\,5876 (2007) & 6.5 & 7.2 & likely alias of $\\nu_1$ \\\\ 13.68 & 1.75 & He\\,{\\sc i} $\\lambda$\\,6678 (04/05) & 7.0 & 7.7 & $= \\nu_1$ \\\\ 12.63 & 1.90 & He\\,{\\sc i} $\\lambda$\\,5876 (2006) & 7.0 & 7.7 & likely alias of $\\nu_1$ \\\\ 8.31 & 2.89 & He\\,{\\sc i} $\\lambda$\\,6678, He\\,{\\sc i} $\\lambda$\\,5876 (2006,2007) & 4,3.5 & 4.5,3.9 & $= \\nu_2$ \\\\ 7.30 & 3.29 & He\\,{\\sc i} $\\lambda$\\,6678 (02/97) & 3.0 & 3.4 & likely alias of $\\nu_2$ \\\\ \\hline \\end{tabular} \\end{center} \\end{table*}  Telting \\& Schrijvers (\\cite{TS}) and Schrijvers \\& Telting (\\cite{ST}) derived linear formulae relating the observable phase differences between the blue and red line wings to the degree $l$ and the absolute value of the azimuthal order $|m|$ of the pulsation mode. These relations are applicable to non-zonal ($m \\neq 0$) spheroidal and toroidal modes and are valid also for multi-mode pulsations including those cases where temperature effects dominate over radial velocity effects. The blue-to-red phase difference of the mode yields $l$, whilst the phase difference for the first harmonic leads to the value of $|m|$ (Telting \\& Schrijvers \\cite{TS}).  For the $\\nu_1$ pulsation mode of HD\\,93521, the phase $\\phi_1(\\lambda)$ is a monotonic function of wavelength (see Figs.\\,\\ref{phase6678april2005}, \\ref{fourmean5876} and \\ref{four5876}), i.e.\\ there are no huge changes in the slope over the interval where the significant line profile variations are detected. Therefore, this mode is unlikely to be an `outlier' in the sense defined by Telting \\& Schrijvers (\\cite{TS}) and the relation between $l$ and the blue-to-red phase difference $\\Delta\\phi$ inferred by the latter authors should thus be applicable to these modes\\footnote{Telting \\& Schrijvers (\\cite{TS}) caution however that their relation was established for models with moderate rotation and hence stars that are not significantly flattened by rotation.}. Applying these relations to the phase differences given in Table\\,\\ref{tabspec} yields $l$ values of $8 \\pm 1$ for $\\nu_1$ and $4 \\pm 1$ for $\\nu_2$.  Howarth \\& Reid (\\cite{HR}) interpreted the $\\nu_1$ modulation as the signature of sectoral mode non-radial pulsations with $l = -m \\simeq 9$. They further inferred a horizontal to radial velocity variation amplitude of $k < 0.3$. Later on, Howarth et al.\\ (\\cite{HTC}) derived $l \\simeq 10 \\pm 1$ and $6 \\pm 1$ for $\\nu_1$ and $\\nu_2$ respectively, with $m + l \\leq 2$. Whilst our values of $l$ are in rough agreement with those of Howarth \\& Reid (\\cite{HR}) and Howarth et al.\\ (\\cite{HTC}), the lack of a significant power in the first harmonics prevents us from deriving the value of $|m|$ from the simple scaling relations for this frequency. We note that the low amplitudes (or actually the upper limits on the amplitude) of the first harmonics compared to those of the genuine frequencies, suggest indeed that $|m|$ is likely close to $l$ (Schrijvers et al.\\,\\cite{schrijvers}). However, we stress that in a rapidly rotating early-type star, such as HD\\,93521, the combined effects of a concentration of the pulsation towards the equator and of a non-uniform temperature distribution due to gravity darkening leads to more complicated amplitudes as a function of wavelength (e.g.\\ Townsend \\cite{Townsend}). A detailed line profile modelling is therefore required to derive the values of $m$ and we defer this to a forthcoming paper.  A priori, the shape of the semiamplitude of the modes as a function of wavelength (see Figs.\\,\\ref{phase6678april2005}, \\ref{fourmean5876} and \\ref{four5876}) suggests that the modes have a rather modest ratio between the amplitudes of the horizontal and radial velocity variations ($k \\leq 0.5$). However, Schrijvers \\& Telting (\\cite{ST}) showed that the typical double-peaked shape of the amplitude found for velocity-dominated line profile variability with high $k$ values vanishes when temperature effects become important. Therefore, since we cannot exclude that the pulsation modes in HD\\,93521 might be affected by temperature effects (see above), we cannot make a secure statement about $k$.  \\begin{figure}[htb!] \\resizebox{9.0cm}{!}{\\includegraphics{0002f11.ps}} \\caption{Top panel: normalised Fourier power spectrum of the He\\,{\\sc i} $\\lambda$\\,6678 line as observed in April 2005. Middle panel: normalised Fourier power spectrum of the $B_1$ photometric data obtained in April 2005. Bottom panel: product of the two normalised power spectra. \\label{product}} \\end{figure}  Our photometric data do not reveal significant variability at the $\\nu_1$ and $\\nu_2$ frequencies. However, because the $\\nu_1$ mode is of high degree, the integrated flux variability due to these pulsations is indeed expected to be rather modest a priori and it could be that these variations just have too low an amplitude to be detected. To see whether or not this explanation is plausible, we have used the method outlined in Aerts et al.\\ (\\cite{deltaCeti}). First, we have normalized the periodogram of the April 2005 He\\,{\\sc i} $\\lambda$\\,6678 time series by dividing it by the power of the highest peak. The same procedure was then applied to the periodograms (for each filter) of the photometric data, also taken in April 2005. The normalized periodograms of the spectroscopic and photometric data were finally multiplied two by two with the rationale that frequencies that are present in both data sets should be dominant in the product of the periodograms (see Aerts et al.\\ \\cite{deltaCeti}). In this way, we find that the $\\nu_2$ frequency is clearly seen in all the products, whilst it had a significantly lower amplitude than $\\nu_1$ in the spectroscopic periodogram (see e.g.\\ the case of the $B_1$ filter in Fig.\\,\\ref{product}). The $\\nu_1$ frequency is seen with a strength slightly larger than the $\\nu_2$ frequency only in the product of the spectroscopic and photometric periodograms of the $U$ and $G$ filters. The fact that $\\nu_2$ appears more prominently in these products than $\\nu_1$ would be consistent with our conclusion that the former mode has a lower $l$ value than the latter. Whilst this test is not a proof for the presence of the spectroscopic frequencies in the photometric variations, it nevertheless suggests that obtaining an extensive photometric time series with a lower noise level (e.g.\\ using a space-borne observatory) would definitely be worth the effort.  We have further found several other possible periodicities that could be present in the photometric data. If real, these modulations might correspond to either radial or lower degree non-radial pulsations. Such modes are difficult to detect in rotationally broadened line profiles, but might well produce an observable signature in the photometric data. Walker et al.\\ (\\cite{Walker}) reported on {\\it MOST} photometry of $\\zeta$\\,Oph: they found that the light curve is dominated by a 4.633\\,hr period with a semiamplitude of 7.3\\,mmag, whilst the other modes have semiamplitudes below 2.25\\,mmag. These authors suggested that the light variations of $\\zeta$\\,Oph mainly result from radial pulsations. In the case of HD\\,93521, the vast majority of the frequencies detected in the periodogram of the time series have semiamplitudes that are significantly lower than what was found in the case of $\\zeta$\\,Oph. Nevertheless, HD\\,93521 would be a very interesting target for an intensive photometric monitoring from space. \\subsection{The need for an alternative interpretation?} Whilst multi-frequency non-radial pulsations offer an attractive interpretation of the line profile variability observed in the spectrum of HD\\,93521, we must nevertheless ask the question whether there could be alternative explanations. We stress that the stability of the frequencies over many years (when they are detected) likely implies that the profile variations are produced by one or several stable clocks such as pulsations (considered above), rotation, orbital motion... Generally speaking, the distinction between these different mechanisms is a non-trivial issue (see e.g.\\ the discussion in Uytterhoeven et al.\\,\\cite{Uytterhoeven}). For instance, in the case of  $\\zeta$\\,Oph, Harmanec (\\cite{Harmanec}) proposed a different scenario where the moving bumps in the absorption lines actually stem from rotating inhomogeneities of the circumstellar material rather than from non-radial pulsations.  In our case, there are also a number of reasons to consider alternative scenarios. Indeed, it is clear that currently a theoretical model to predict the line profile variations produced by non-radial pulsations in a rapidly rotating massive star such as HD\\,93521 is still lacking. Therefore, the interpretation of these features within the framework of the available models requires an extrapolation that might be difficult to justify a priori.  Another feature that is certainly puzzling is the fact that the line profile variability is significantly detected only in lines that are potentially affected by residual emission possibly associated with a circumstellar disk (or flattened wind). Whilst the lack of a TVS signal in the purely photospheric O\\,{\\sc iii} $\\lambda$\\,5592 line (see Fig.\\,\\ref{average}) might be interpreted as this line forming near the hotter poles of the star where the pulsational amplitude is lower (as for He\\,{\\sc ii} lines), the same explanation cannot hold for the photospheric features that occur at temperatures and gravities that are typical of those of the equatorial region (see the C\\,{\\sc ii}, N\\,{\\sc ii} and Si\\,{\\sc iii} lines discussed in Sect.\\,\\ref{spec}). The latter features produce at best a marginal signal in the TVS (see Fig.\\,\\ref{average}). However, regarding the possibility that the variability stems from rotating features in the circumstellar material, we first note that the frequencies $\\nu_1$ and $\\nu_2$ are not detected in the emission wings of the lines analysed in this work. If the line profile variations were coming from the flattened wind, one would also expect them to affect the emission wings. We further note that all the lines where we have detected line profile variability are quite strong. Actually, the rather low amplitude of the profile variations might render them undetectable in the weaker equatorial lines.  In relation to this, we note that a recent study of HD\\,60848 revealed evidence for the existence of rather short (3.51 and 3.74\\,hr) periods in the radial velocities derived from the emission lines of this O9.5\\,IVe star (Boyajian et al.\\,\\cite{Boyajian}). These authors argued that these features might result from {\\it changes in the disk density or illumination caused by non-radial pulsations in the underlying star}. Therefore, it seems that a disk origin for short period variations cannot be ruled out a priori, although it would also reflect the existence of pulsations in this specific example.  Since HD\\,93521 was considered as a possible runaway object (Gies \\cite{Gies}, but see also the discussion in Sect.\\,\\ref{intro}), it could host a compact companion which could then trigger a periodic structure in the inner part of the disk that is seen projected against the photosphere. However, the presence of a compact companion should make HD\\,93521 a bright X-ray source, at least episodically when the compact object crosses the equatorial wind (similar to Be-type high-mass X-ray binaries). We thus checked the {\\it ROSAT} All Sky Survey images as well as several other X-ray catalogues. To the best of our current knowledge, there is no indication that HD\\,93521 is, or has ever been, a bright or even moderate X-ray emitter. Another point is that the detected frequencies $\\nu_1$ and $\\nu_2$ are much too high to correspond to the orbital period ($\\geq 13.5$\\,hr and likely of order several days) of such a putative companion. A compact companion scenario appears therefore rather unlikely.  Another point concerns the fact that both frequencies are too high to correspond to the rotational period of the star. In fact, whilst many of the physical parameters of HD\\,93521 remain unknown or poorly determined, we can obtain a rough estimate of the rotational frequency by considering typical parameters of a late O-type star. Let us assume that HD\\,93521 is an O8.5 star\\footnote{The observational classification as an O9.5 star is likely biased towards later spectral types as a result of gravity darkening due to the high rotational velocity and the nearly equator-on orientation of the star.} with a polar radius of 8\\,R$_{\\odot}$ and a mass of 20\\,M$_{\\odot}$. If the star is actually seen equator-on, the resulting rotational frequency would be 0.81\\,d$^{-1}$. In general, we can say that for any reasonable assumption on the stellar parameters, we find that $\\nu_{\\rm rot} \\leq 1$\\,d$^{-1}$. Now, if the line profile variations actually stem from a regular pattern of moving spokes in the circumstellar disk, $\\nu_1$ and $\\nu_2$ would not necessarily have to correspond to $\\nu_{\\rm rot}$, but could rather be some harmonics of the latter (see Uytterhoeven et al.\\,\\cite{Uytterhoeven}). We would thus have to look for a ``super'' frequency such that $\\nu_1 = n_1\\,\\nu_{\\rm sup}$ and $\\nu_2 = n_2\\,\\nu_{\\rm sup}$, where $n_1$ and $n_2$ would be integer numbers. Within the uncertainties of our period determinations, a candidate for such a frequency would be $\\nu_{\\rm sup} = 2.75$\\,d$^{-1}$ (with $n_1 = 5$ and $n_2 = 3$). Again, the latter frequency is much too high to correspond to the rotational frequency. The super-frequency could be in agreement with the likely value of the rotational frequency if $n_1$ and $n_2$ were muliplied by 3. However this would imply a large number (15 and 9) of co-rotating structures around the star which seems rather difficult to explain.  In summary, we conclude that, all the alternative scenarios envisaged here fail in explaining the modulations at the $\\nu_1$ and $\\nu_2$ frequencies. Hence, despite some difficulties, the multi-period non-radial pulsations model remains currently the most plausible explanation for the line profile variations seen in HD\\,93521.   \\acknowledgement{We thank the referee, Dr.\\ D.\\ Gies for his very helpful report. The Li\\`ege group acknowledges financial support from the FRS-FNRS (Belgium), as well as through the XMM and INTEGRAL PRODEX contract (Belspo). The travels to OHP were supported by the `Communaut\\'e Fran\\c caise' (Belgium). The Mercator telescope and its operations are funded by the Catholic University of Leuven, the Flemish community and the Fund for Scientific Research of Flanders (FWO). The Mercator observations were performed by the Leuven team in the framework of the FWO project G.0178.02. PE acknowledges support through CONACyT grant 67041.}"
    },
    "0806/0806.4256_arXiv.txt": {
        "abstract": "{The standard picture of accretion is a steady flow of matter from the disc onto the young star - a concept which assumes the star-disc system to be completely isolated. However, in a dense cluster environment star-disc systems do interact gravitationally.} {The aim here is to estimate the encounter-induced accretion rate in an ONC-like environment.} {Combining simulations of the cluster dynamics with simulations of the effect of encounters on star-disc systems we determine the likelihood and degree of encounter-triggered accretion processes.} {We show that accretion bursts triggered by encounters of star-disc systems are common in young dense clusters like the ONC leading in the outburst phase to typical accretion rates of 10$^{-7}$-10$^{-4}$ \\Msun/yr. Up to a third of stars presently in the Trapezium region accreted at least 1\\% of their disc mass via this mechanism in the last 1Myr. Accretion of over 6-7\\% of the disc material can occur in a single encounter. Despite losing their discs quickly, the total percentage of disc matter accreted per star is largest for the massive stars.} {Supplementing the steady accretion flow there exist episodic periods of high accretion in dense cluster environments. Due to their high accretion rate these processes should be observable even now in some of the low-mass stars in the ONC. }  \\titlerunning{Accretion bursts in young stars} \\authorrunning{Pfalzner et al.} ",
        "introduction": "The theory of accretion has been extensively studied in the context of the standard accretion disc model \\cite{pringle:72,shakura:73,pringle:81}. The main idea is that of gas moving inward due to radial turbulent transport of angular momentum to the outer disc regions. Currently the magneto-rotational instability \\cite{balbus:91} seems to be able to provide the required degree of turbulence in accretion discs.   The disc evolution that follows is not completely understood. In the later phases accretion happens exclusively from the disc material, so linking disc development and accretion directly. The often-used  theoretical models of viscous disc evolution \\cite{lynden:74,hartmann:98} treat the disc as an isolated smooth axisymmetric structure that evolves due to an unspecified source of temporally independent, turbulent viscosity. These standard models assume a possibly decreasing with time but relatively {\\em steady accretion} process. Recently  models were suggested where accretion is caused by gravitational instabilities in the disc \\cite{vorobyov:06,boley:06} leading to alternating phases of high and low accretion. However, for these instabilities to occur the disc has to be relatively massive ($m_d >$ 0.1 $M^*$), which means it will predominantly occur very early in star formation.  From observations the general picture emerges that the disc fraction is a strong function of age, decreasing from $\\sim$ 80\\% for clusters at 1 Myr down to few discs for clusters older than 10 Myr \\cite{haisch:01,hernandez:07}. The typical accretion rate for a young star is $\\sim$ 10$^{-8}$- 10$^{-10}$\\Msun/yr \\cite{hillenbrand:95,gullbring:98,haisch:01}. However, the development of the disc and the accretion rate of individual sources in the same environment can differ considerably \\cite{flaherty:08}. The accretion rate depends on the stellar age, mass and environment. As expected from viscous disc models a decrease in accretion rates with stellar age has been found \\cite{hartmann:98, muzerolle:00}. A matter of debate is the dependence on the stellar mass of the form $M^\\alpha$ \\cite{calvet:04,sicilia:06,natta:06}, possibly with  $\\alpha \\sim $ 2. The effect of the environment possibly manifests itself in that  Herbig Ae stars have on average higher accretions rates \\cite{lopez:06} than low-mass stars in Taurus and Ophiuchus.  Here we investigate an accretion process that will be present in dense cluster environments {\\em in addition} to the steady accretion - the passage of a star inducing a simultaneous transport of disc matter outwards and inwards via spiral arms, leading to accretion. We will show that the consequences are short bursts of high accretion (10$^{-7}$ - 10$^{-4}$ \\Msun/yr), inevitably occuring in dense clusters and observable due to these high accretion rates.  Although this process has properties in common with accretion induced by gravitational instabilities  \\cite{vorobyov:06,boley:06} and in wide binaries \\citep{bonnell:92}, the cluster dependence is unique to the process treated here.  For encounters to play an important role the cluster has to be dense. A typical example for such an environment is the Orion Nebula Cluster(ONC). In an earlier paper \\cite{pfalzner:06a}, we demonstrated that encounters cause a 3-5\\% specific angular momentum loss in the ONC  rising to 10-15\\% in the dense inner Trapezium region. Since specific angular momentum loss is a prerequisite for accretion we suggested that in the final star formation stages an additional growth mechanism for massive stars exist (cluster-assisted accretion) as massive stars lose their specific angular momentum to a higher degree than low-mass stars. Here we study  encounter-induced accretion directly by measuring the amount of matter kicked into the inner disc areas. Obviously this is not strictly accretion, but represents a often-used method in circumstances where computational expense does not allow to further resolve the inner disc \\cite{bate:02,vorobyov:05}. In the following the term ``accreted'' will be used in this sense but the limitations of this approach should be kept in mind. ",
        "conclusions": "There is one observed systems that displays the link between encounter and accretion rate. The morphology and kinematics \\cite{cabrit:06} of RW Aurigae A and B strongly suggest tidal stripping of the primary disc by a recent fly-by that is occurring. This system displays a high accretion rate of $\\sim$ 2-10 $\\times 10^{-7}$ \\Msun/yr \\cite{basri:89,hartigan:95} while having a particularly low disc mass of $m_d^{c} \\sim$ 3 $\\times$ 10$^{-4}$ \\Msun . Using these parameters our simulations show the likely total accreted mass to be 0.3 - 4 \\% of the pre-encounter mass, equivalent to a maximum induced accretion rate of 1$\\cdot 10^{-6}$ - 1.2 $\\cdot 10^{-5}$ \\Msun/yr assuming a pre-encounter disc mass $ m_d^{pre} < 10 \\times m_d^{c}$) and that most mass is accreted within $\\sim$ 10 years (although high accretion for several 100 years exist). This accretion rate is consistent with the observed accretion rate of $\\sim 2 - 10 \\cdot 10^{-7}$ \\Msun/yr. So encounter-induced accretion is probably occuring in this system.   \\begin{figure} \\resizebox{\\hsize}{!}{\\includegraphics[angle=-90]{enc_time_av.ps}} \\caption{Average number of encounters in any time interval of 10 000 yrs that lead to accretion as a function of cluster age.} \\label{fig:enc_time_av} \\end{figure}   However, RW  Aurigae is not located in a dense cluster, so either the rare case of an encounter in a low-density region or a strongly elliptical wide binary is being observed. Even in the latter case, the results described in Section 3 are still applicable as the effect of a parabolic encounter on the induced accretion is very similar to that in a strongly elliptical wide binary. Episodes of rapid accretion due to tidal disturbances to the disc induced by a binary companion were suggested by Bonnell and Bastien (1992) in the context of FU Orionis flareups, which are believed to be powered by episodes of exceptionally rapid accretion (up to 10$^{-4}$ \\Msun /yr) \\cite{hartmann:96}. Alternatively, such outbursts could be due to thermal instability \\cite{bell:94,bell:95}, or to gravitational instability in the disc \\cite{vorobyov:06}. While many of the properties (accretion rate, duration of burst) of the accretion bursts  described here agree with those of FUors, the strongest argument against it is the lack of FUors observed in dense stellar regions. The arguments for and against a connection of encounter-triggered accretion bursts to FU Orionis objects will be discussed in more detail in Pfalzner (2008).  Accretion bursts might be even more important in the earlier phases where star-disc systems are still in their formation process \\cite{larson:mnras02}. The simulation of the formation of a small cluster of stars by Bate (2002)  clearly shows the formation of many circumstellar discs. The complex dynamics of the system includes frequent protostellar interactions disrupting the discs by encounters. However new discs soon form from gas that continues to fall inwards, so that in this phase protostars may gain much of their mass in discrete episodes of encounter-triggered accretion.  The observed accretion rates in Class I protostars are orders of magnitude smaller than those needed to form a star during the lifetime of a Class I object \\cite{mckee:07}. Kenyon et al. (1990) suggested two solutions to this problem - either significant accretion continues into the T Tauri stage or  most of the accretion occurs in the embedded stage. The former appears to be ruled out because such stars accrete very slowly with no significant disc or envelope mass reservoir that they can draw on. In the latter case accretion could be episodic, so that the median accretion rate is much smaller than the mean.  Another fact that can be explained if the accretion occurs in short bursts is that the observed luminosities of many protostars are lower than is predicted for models with steady accretion \\cite{kenyon:95,calvet:04}. The jet-like Herbig-Haro outflows probably powered by rapid accretion onto forming stars at early stages of evolution, are episodic or pulsed,  suggesting that the accretion process is itself episodic \\cite{reipurth:01}. Remarkably the jet sources are frequently (in 85\\% of cases) found to have close stellar companions \\cite{larson:mnras02}."
    },
    "0806/0806.2700_arXiv.txt": {
        "abstract": "We investigate {\\em p}-mode absorption in a sunspot using SOHO/MDI high-resolution Doppler images. The Doppler power computed  from a three and a half hour data set is used for studying the absorption in a sunspot. The result shows an enhancement in absorption near the umbral-penumbral boundary of the sunspot. We attempt to relate the observed absorption with the magnetic field structure of the sunspot. The transverse component of the potential field is computed using the observed  SOHO/MDI line-of-sight magnetograms. A comparison of the power map  and the computed potential field shows enhanced absorption near the umbral-penumbral boundary where the computed transverse field strength is higher. ",
        "introduction": "The interaction of solar oscillations with magnetic field has been reported by many authors. Two of the important findings are the reduced {\\em p}-mode power in regions of strong magnetic field and the enhancement of power in higher frequencies surrounding the regions of strong fields (Woods and Cram, \\citeyear{woods}; Tarbell {\\it et al}., \\citeyear{tarbel}; Brown {\\it et al}., \\citeyear{brown}). The  magnetic field reduces {\\em p}-mode power in active regions in the 3 mHz band while enhances power of these modes in the 5 mHz band (Hindman and Brown, \\citeyear{hindman}; Venkatakrishnan, Kumar, and Tripathy, \\citeyear{pvk}). The mechanisms of {\\em p}-mode absorption were reviewed by Spruit (\\citeyear{spruit}) and he suggested a promising mechanism for reduced power as the conversion of {\\em p} mode into a downward propagating slow mode along the magnetic flux tubes (Spruit and Bogdan, \\citeyear{spruit1}; Cally and Bogdan, \\citeyear{cally1}; Cally, Bogdan, and Zweibel, \\citeyear{cally2}). The waves in the flux tube, once excited by the sound wave, can carry energy out of the {\\em p} modes through the  wave guide into the convection zone, thereby  producing a reduction in observed power. Simulations by Cally (\\citeyear{cally}) show that the enhanced absorption  takes place primarily in the more inclined magnetic field regions towards the edge of the spot. A comparison of the spatial distribution of Doppler power and the magnetic-field configuration in a sunspot could reveal the variation of Doppler power with magnetic-field strength and the field inclination. In this paper, we study {\\em p}-mode absorption in magnetic field concentrations and compare that with the longitudinal and the computed transverse field configuration in a sunspot. ",
        "conclusions": "We analyzed the spatial distribution of Doppler power in a sunspot observed near disk center. We find a structured ring like absorption pattern in Doppler power  near the umbral-penumbral boundary. The computed transverse field is higher at those locations where the peak depression in Doppler power is observed. The computed inclination angle ranges between 35$^{\\circ}$ and 45$^{\\circ}$ at these locations. In order to understand the exact dependence of magnetic field strength and inclination on the Doppler power absorption, full vector magnetic-field observations are required. Also, it is preferred to have Doppler observation in magnetically insensitive lines to avoid any cross-talk between the line-of-sight velocity measurement and Zeeman splitting of the spectral line.  \\begin{acks} I would like to thank B. Ravindra for providing the code for potential field calculation and the anonymous referee for their valuable comments which helped to improve the paper substantially. SOHO is a project of international cooperation between ESA and NASA. \\end{acks}"
    },
    "0806/0806.4727_arXiv.txt": {
        "abstract": "{The MHD instabilities can generate complex field topologies even if the initial field configuration is a very simple one.} {We consider the stability properties of magnetic configurations containing a toroidal and an axial field. In this paper, we concentrate mainly on the behavior of non-axisymmetric perturbations in axisymmetric magnetic configurations.} {The stability is treated by  a linear analysis of ideal MHD equations.} {In the presence of an axial field, it is shown that the instability can occur for a wide range of the azimuthal wavenumber $m$, and its growth rate increases with increasing $m$. At given $m$, the growth rate  is at its maximum for perturbations with the axial wave-vector that makes the Alfv\\'en frequency approximately vanishing. We argue that the instability of magnetic configurations in the ideal MHD can typically be dominated by perturbations with very short azimuthal and axial wavelengths.} {} ",
        "introduction": "A wide variety of MHD instabilities can occur in magnetized astrophysical bodies where they play an important role in the evolution and formation of various structures and in enhancing transport processes, among others. The onset of instabilities can be caused both by hydrodynamic motions (for instance, differential rotation) or properties of the magnetic configuration. Even magnetic fields with a relatively simple topology (for example, a purely toroidal field) can be subject to instability. Magnetic fields generated by the dynamo action or stretched by hydrodynamic motions are topologically more complex and can cause this sort of instability as well. Which field strength and topology can sustain a stable magnetic configuration is still rather uncertain despite all the extensive work already done (see Borra et al. 1982; Mestel 1999 for review).  The simplest and best-studied magnetic configuration is most likely a purely toroidal one. {This has been known since the paper by Tayler (1957), where stability properties of the toroidal field $B_{\\varphi}$ are determined by the parameter $\\alpha = d \\ln B_{\\varphi} / d \\ln s$  where $s$ is the cylindrical radius. The field is unstable to axisymmetric perturbations if $\\alpha  > 1$ and to non-axisymmetric perturbations if $\\alpha> -1/2$.} The growth time of instability is close to  the time taken for an Alfv\\'en wave to travel around the star on a toroidal field line. Numerical modeling by Braithwaite (2006) confirms that the toroidal field with $B_{\\varphi} \\propto s$ or $\\propto s^2$ is unstable to the $m=1$ mode ($m$ is the azimuthal wave number) as  predicted by Tayler (1957, 1973). However, even a purely toroidal field can be stable in the region where it decreases rapidly with $s$. {A purely toroidal field cannot be stable through the whole star because the stability condition for axisymmetric modes ($\\alpha < 1$) is incompatible with the condition that the electric current in the $z$ direction has no singularity at $s \\rightarrow 0$, which implies $\\alpha > 1$.} The stability of the toroidal field in rotating stars has been considered by Kitchatinov \\& R\\\"udiger (2007), who argue that the magnetic instability is essentially three-dimensional and that the finite thermal conductivity creates a strong destabilizing effect. Terquem \\& Papaloizou (1996) and Papaloizou \\& Terquem (1997) considered the stability of an accretion disk with the toroidal magnetic field and found that the disks containing a purely toroidal field are always unstable and calculated the spectra of unstable modes in the local approximation.  The stability properties of purely poloidal magnetic fields are also well-studied. It has been understood since the papers by Wright (1973) and Markey \\& Tayler (1973, 1974) that the poloidal field is subject to dynamical instabilities in the neighborhood of points (or lines) { where the poloidal field is vanishing (neutral points/lines).} These authors recognized first that the magnetic field in the neighborhood of a neutral line resembles that of a toroidal, pinched discharge, which is known to be unstable. The instability of a poloidal field is also rather fast: its growth time can reach a crossing time of few Alfv\\'en times (Van Assche et al. 1982; Braithwaite \\& Spruit 2006) that is very short, for example, compared to the time-scales of stellar evolution. However, the instability of a poloidal field can be suppressed by the addition of a toroidal field in the neighborhood of neutral points (Markey \\& Tayler 1973; Wright 1973).  { Conversely}, the addition of even a relatively weak poloidal field alters  the stability properties of the toroidal field substantially. { For example, if the poloidal field is uniform and relatively weak, the instability condition of axisymmetric modes reads $\\alpha > -1$,  at variance with the condition of instability for a purely toroidal field (see, e.g., Knobloch 1992; Dubrulle \\& Knobloch 1993), which predicts that an unstable toroidal field configuration has $\\alpha > -1/2$. Therefore, a weak poloidal field has a destabilizing effect. However, a  strong enough poloidal field can suppress the instability of the toroidal field.} It turns out that configurations containing comparable toroidal and poloidal fields are more stable than purely toroidal or purely poloidal ones (Prendergast 1956; Tayler 1980) and, generally, the possibility exists that there are configurations containing mixed fields, which have no instabilities arising on a dynamical time-scale. In his study of unstable magnetic configurations Tayler (1980) has not found any instability if the axial field $B_z > 0.3 B_{\\varphi}$ for instance, even though such configurations can be unstable for a wide range of the azimuthal wavenumber $m$ if $B_z$ is weaker. With numerical simulations Braithwaite \\& Nordlund (2006) studied the stability of a random initial field in the stellar radiative zone and found that the stable magnetic configurations generally have the form of tori with comparable poloidal and toroidal field strengths.  In this paper, we consider in detail the stability properties of magnetic configurations containing the toroidal and axial magnetic fields with respect to non-axisymmetric perturbations. We show that the instability may occur for a wide range of the azimuthal wavenumber $m$, and the growth rate is typically higher for higher $m$. Unstable modes with large $m$ have a very short vertical lengthscale, so it can be hard to resolve them in numerical calculations. Depending on the profile $B_{\\varphi}(s)$ and the ratio $B_z/B_{\\varphi}$, the instability can occur in two regimes that have substantially different growth rates.  The remainder of this paper is arranged as follows. In Sect. 2, we derive the equation that governs the eigenfunctions and eigenvalues of the magnetic field. We describe the numerical procedure and present the results of calculations in Sect. 3. A brief discussion of the results is given in Sect. 4. ",
        "conclusions": "We have considered the linear stability of magnetic configurations containing the toroidal and axial fields, assuming that the behavior of small perturbations is governed by equations of the non-dissipative incompressible magnetohydrodynamics. This approximation is  justified if the magnetic field is subthermal and the Alfv\\'en velocity is low compared to the sound speed. The stability of magnetic configurations is a key issue for understanding the properties of various astrophysical bodies such as peculiar A and B stars, magnetic white dwarfs and neutron stars. The magnetic instability can alter qualitatively the properties of configurations generated, for example, by dynamo in stars. Many dynamo models predict that the toroidal field should typically be stronger than the poloidal one, but such configurations can be unstable if the generated toroidal field does not decrease enough rapidly with $s$. The instability generates large- and small-scale motions that should alter the geometry of a generated magnetic field.  Even though the poloidal field is weaker than the toroidal one in a number of dynamo models, its effect on the stability properties cannot be neglected. This  particularly concerns the behavior of the nonaxysimmetric perturbations considered in the present paper. If $B_z$ is relatively weak ($B_z < B_{\\varphi}$) then, typically, there exists a wide range of the azimuthal wave-numbers $m$ for which the instability may occur. For any given $m$, only perturbations within some particular range of the vertical wave-vectors $k_z$ can be unstable. The growth rate is maximal for perturbations with $k_z$ of about \\begin{equation} k_z \\sim - \\frac{m}{s} \\; \\frac{B_{\\varphi}}{B_z}. \\end{equation} Equation~(19) is approximately equivalent to the condition that the Alfv\\'en frequency is vanishing somewhere within the cylindrical layer (see Eq.~(17)). {It should be noted that the ratio $B_z/ B_{\\varphi}$ can be rather low in stars, and the maximum growth rate at given $m$ corresponds to very short axial wavelengths, \\begin{equation} \\lambda_z \\sim 2 \\pi \\; \\frac{s}{m} \\; \\frac{B_z}{B_{\\varphi}}. \\end{equation} Taking, for example, $B_z/B_{\\varphi} \\sim 10^{-2}$ and assuming that $s$ is comparable to the stellar radius, $s \\sim 10^{11}$ cm, we find that the most rapidly growing modes should have the axial wavelength $\\sim 6 \\times 10^9/ m$ cm. From our results it follows that the maximum growth rate increases slowly with increasing $m$ and, therefore, perturbations with a very short azimuthal wavelength (very large $m$) should dominate the development of instability. For instance, the most rapidly growing mode with $m=100$ corresponds to the axial wavelength $\\sim 6 \\times 10^7$ cm, which  is very short compared to the radius.} Therefore, the instability of magnetic configurations  can often be determined by the modes with very large $m$ and an extremely short wavelength in the $z$-direction. This fact can cause problems in numerical modeling of the instability because simulations will require a very high resolution in the $\\varphi$- and $z$-directions.  Depending on the profile of the toroidal field and the strength of the axial field, the instability can arise in two essentially different regimes. In the case of a weak axial field, $B_{\\varphi 0} \\gg B_z$, the value of $\\alpha$ that distinguishes between the regimes is $\\approx -1/2$. If $\\alpha > -1/2$, then the instability grows on the Alfv\\'en timescale determined by the toroidal field and is rather fast. In this case, the growth time is \\begin{equation} \\tau \\sim 0.1 \\rho_{-4}^{1/2} s_{11}B_{\\varphi 3}^{-1} \\;\\;\\; {\\rm yrs}, \\end{equation} where $s_{11} = s /10^{11}$cm. If $\\alpha < -1/2$, then the growth time is given by same expression (18) but where $B_{\\varphi 3}$ should be replaced by $B_{z 3}= B_z/ 10^3$ G. Since $B_{\\varphi} \\gg B_z$, the instability is slower and grows on the timescale determined by the axial field in this case. The transition between two regimes occurs at larger $\\alpha$ if the axial field increases.  It is fairly difficult to compare our results obtained for a simple model with the available numerical simulations, which usually use completely different basic magnetic configurations. For example, in calculations by Braithwaite (2006, 2007), the basic configuration was assumed to be either purely toroidal or purely poloidal, and the stability properties of such configurations differ qualitatively from those considered in this paper. Recently, Braithwaite \\& Nordlund (2006) and Braithwaite (2008) have considered stability of the magnetic configuration with random initial fields. A vector potential was set up as a random field containing spatial scales up to a certain value. This random field was then multiplied by some screening function, so that the field strength in the atmosphere was negligible.  This initial configuration contains both the toroidal and poloidal fields but is very different from our simple model. Nevertheless, some features seem to be in common even for such different models. { Braithwaite \\& Nordlund (2006) and Braithwaite (2008) find that the instability can lead to different equilibrium configurations depending on the screening function. If the screening function for random fields decreases slowly or does not decrease at all, then the final equilibrium magnetic configuration is essentially non-axisymmetric. In contrast, an equilibrium configuration is closer to axisymmetry (but not axisymmetric) if the screening function decreases rapidly. This dependence on equilibrium configurations obtained in numerical calculations can reflect the regimes of ``strong'' and ``weak'' instabilities that correspond to different growth rates depending on the value of $\\alpha$. In accordance with our analysis, the growth rate at given $\\varepsilon$ is higher for higher $\\alpha$ (compare, e.g., Figs.~3 and 7). Our parameter $\\alpha$ mimics to some extent the screening parameter $p$ introduced by Braithwaite (2008) with decreasing $\\alpha$ corresponding to increasing $p$. Therefore, we can expect from our analysis that non-aximetric instabilities should be more efficient for the screening function with $p=0$ than with $p=1$ and that the final configuration exhibits stronger departures from axisymmetry for smaller $p$. This conclusion seems to be in qualitative agreement with the results of Braithwaite (2008).}  Our simple model does not take into account the stratification that can be important in many astrophysical applications. Basically, stratification provides a stabilizing influence if the temperature gradient is sub-adiabatic. However, this influence is small if perturbations have a relatively short wavelength in the axial direction, $\\lambda < \\lambda_c$, such that inequality (11) is satisfied. Our results are related to this case. The case when $\\lambda > \\lambda_c$ and stratification is important will be considered elsewhere.  \\vspace{0.5cm}  \\noindent  {\\it Acknowledgments.} This research project was  supported by a Marie Curie Transfer of Knowledge Fellowship of the European Community's Sixth Framework Program under contract number MTKD-CT-002995. VU also thanks  INAF-Ossevatorio Astrofisico di Catania for hospitality."
    },
    "0806/0806.3331_arXiv.txt": {
        "abstract": "We consider to what extent the long-term dynamics of cyclic solar activity in the form of Grand Minima can be associated with random fluctuations of the parameters governing the solar dynamo. We consider fluctuations of the alpha-coefficient in the conventional Parker migratory dynamo, and also in slightly more sophisticated dynamo models, and demonstrate that they can mimic the gross features of the phenomenon of the occurrence of Grand Minima over a suitable parameter range. The temporal distribution of these Grand Minima appears chaotic, with a more or less exponential waiting time distribution, typical of Poisson processes. In contrast however, the available reconstruction of Grand Minima statistics based on cosmogenic isotope data demonstrates substantial deviations from this exponential law. We were unable to reproduce the non-Poissonic tail of the waiting time distribution either in the framework of a simple alpha-quenched Parker model, or in its straightforward generalization, nor in simple models with feedback on the differential rotation. We suggest that the disagreement may only be apparent and is plausibly related to the limited observational data, and that the observations and results of numerical modeling can be consistent and represent physically similar dynamo regimes. ",
        "introduction": "The solar cycle is believed to be a result of dynamo action occurring somewhere inside the solar convective envelope. According to the classical Parker (1955) model, this dynamo action can be envisaged as follows. Differential rotation $\\Omega$ produces toroidal magnetic field from poloidal, while the \"$\\alpha$-effect\" associated with the helicity of the velocity field produces poloidal magnetic field from toroidal. According to this scheme, the solar cycle length is identified with the dynamo time-scale, which can be estimated from the product of the amplitudes of the $\\alpha$-effect and rotational shear ($\\partial \\Omega / \\partial r$, $r$ being the radial coordinate), appropriately normalized with the turbulent diffusion coefficient (these yielding the dimensionless dynamo number), and with the turbulent diffusion time. The Parker model results in a periodic process in the form of propagation of a toroidal field pattern in the latitudinal direction (the \"butterfly diagram\"). A suitable choice of governing dynamo parameters gives equatorward pattern propagation as well as allowing the cycle period to agree with observations. More realistic dynamo models try to demonstrate that some plausible choice of parameters is compatible with available observational information from, say, helioseismology, or elaborates this simple scheme by various additional details, such as meridional circulation (see Petrovay, 2000; Charbonneau, 2005, for reviews).  In fact, the solar cycle is far from being a strictly periodic phenomenon. The amplitude of solar cycles varies substantially in time and reaches unusually large levels during the so-called Grand Maxima, one of which is now believed to be occurring. From time to time the level of solar cyclic activity becomes extremely low if not disappearing completely. Such minima of the cyclic activity are known as Grand Minima, the most well-known example being the Maunder Minimum, which occurred in the middle 17-th - beginning of 18-th centuries. The statistics of Grand Minima (and Maxima) can be to some extent reconstructed from data  on cosmogenic isotope $^{14}$C in tree rings (Usoskin {\\it et al.}, 2007). Quantification of the sequence of such events is still a contentious topic.  It is important that the isotopic data provides a much longer record of Grand Minima/Maxima than do the sunspot observations. Moreover, the sequence of Grand Minima (and/or Maxima) appears to be random, rather  than a periodic process.  It is known that simple deterministic numerical dynamo models of the solar cycle, which essentially develop the ideas of the Parker migratory dynamo, can give events comparable with  Grand Minima/Maxima ({\\it e.g.} Brandenburg {\\it et al.}, 1989), even showing behaviour which is irregular and chaotic in time (see {\\it e.g.} Jennings and Weiss, 1991; Jennings, 1991; Tobias {\\it et al.}, 1995; Covas {\\it et al.}, 1998  -- see also Moss and Brooke, 2000 in a more complex model). The presence of a long-term dynamics needs however an explanation. The most straightforward idea here is to recognize that the $\\alpha$-effect, being the result of the electromotive force averaged over turbulent vortices, can contain a fluctuating contribution (Hoyng, 1993; Hoyng {\\it et al.}, 1994; Ossendrijver and Hoyng, 1996).  The idea can lead to events similar to the Maunder Minimum on the timescale of centuries (see {\\it e.g.} Tworkowski {\\it et al.}, 1998; also  Brandenburg and Spiegel, 2008).  The aim of this paper is to investigate the long-term dynamics of solar  activity by confronting the predictions of a Parker migratory dynamo model containing a random contribution to the $\\alpha$-coefficient with the available data concerning the sequence of Grand Minima and Maxima, as inferred from the isotopic data. We also consider, more briefly, a more sophisticated dynamo model. Our general conclusion is that the fluctuations in  the dynamo governing parameters can lead to phenomena similar to the Grand Maxima and Minima, in that the temporal distribution of the events appears chaotic.  We recognize a disagreement between the observational data and numerical simulations of our dynamo model, in that the statistics of the waiting time distributions of Grand Minima appear to have exponential tails, in contrast to the isotopic data  in which the temporal distribution of Grand Minima and Maxima demonstrate a substantial deviation from exponential statistics. We argue that the disagreement plausibly is only apparent, and is connected with the limited extent in time of the observations, and that observations and modeling may represent physically similar dynamo regimes.  In this paper we concentrate on the distribution of Grand Minima. ",
        "conclusions": "\\label{disc}   We have demonstrated that the phenomenon of the occurrence of solar Grand Minima can be simulated as an effect of fluctuations in the governing parameters in a simple model of solar dynamo at least in the framework of the interpretation of observation suggested. We stress that the limited nature of the observations available does not make it possible to compare the results of simulations and observations in complete detail ; however we do not see in the observational data anything that is basically incompatible with the simulations. Thus, simulations in regimes marked as \"exp*\" in the tables look close to the observed phenomenology  and might be regarded to be not inconsistent with the observations.  However, since the results cannot be directly compared in the statistical sense, other possibilities which exist to explain the seeming disagreement between simulations and observations deserve consideration. First of all, the phenomena of Grand Minima and Maxima may be associated with fine details of the solar dynamo (for example, the exact shape of the solar rotation law), rather than being a general property of nonlinear dynamos in a spherical shell for a suitable parameter range. (Our experiments with a realistic rotation law in a 2D dynamo model, although not encouraging, were too limited to rule this out completely.) A further possibility is that the dynamo mechanism itself produces a Poisson-like sequence of Grand Maxima/Minima, but there are also long-term trends in solar hydrodynamics (on the scale of thousand years) which affect the timescale of the weighting time and mimics the non-Poisson behaviour. This could be, for example, via the Reynolds stresses that drive the differential rotation. Another option is that the non-Poissonic nature of the observed sequence of Grand Minima/Maxima is an artefact of the limited statistics.   We stress once more that we present here a development of a very strong suggestion that the solar dynamo engine does not contain any specific mechanism that produces Grand Minima (and Maxima), but that they are rather a result of random fluctuations in the dynamo governing parameters. On one hand, the ability to convert random noise into a sequence of clearly separated events looks  an intriguing feature of the dynamo. This ability however can be considered as an example of intermittency, which is a known property of various nonlinear systems  where it can produce various spatial or/and temporal structures from random noise (see {\\it e.g.} Zeldovich {\\it et al.}, 1991). On the other hand, it looks more than plausible that the solar dynamo does possess something specific that allows fine tuning of WTD of Grand Minima, which produces non-Poisson tails. However we can not identify this feature of the dynamo engine at the moment, and realize that the physical mechanism behind  the occurrence of Grand Minima may be only partly or not at all related to random fluctuations ({\\it e.g.} Petrovay, 2007).   We appreciate that the problem considered here belongs to the general topic of the influence of noise, which is addressed in many fundamental papers. Our ability to exploit the deep methods suggested in this area ({\\it e.g.}  Abarbanel {\\it et al.}, 1993) is however restricted by the limited nature of the available observational data. Note however that we have incorporated some memory effects into models III, IV and V  so that, in principle at least, we are going slightly beyond studying the effects of random noise and the expected associated chaos. Also that it is not a priori altogether obvious how these random inputs will appear after passing through the dynamo \"machine\".   Our work is based on the tacit assumption that the sunspot number is linearly related to the magnetic energy in the dynamo. This is a common asssumption in solar cycle modelling, but it is quite possible that the number and size of the active regions appearing on the surface might be more plausibly taken to be proportional to the toroidal flux, rather than to the energy. However this option leads to similar conclusions to those presented above. On the other hand, it also seems possible that there could be a threshold effect at play here, so that active regions only emerge if the toroidal field strength exceeds some minimal value ({\\it e.g.} Ruzmaikin, 2001). This would introduce a marked nonlinearity into the relationship between activity indices and toroidal field parameters. It is clear that the results obtained can be sensitive to this nonlinearity which, in principle, we could investigate. Accordingly, our investigation of the Grand Minima phenomenon has to be considered to some extent as illustrative until the influence of the nonlinearity is resolved. A further possibility in a 2D model is to use the toroidal flux or energy in the immediate sub-surface region as a proxy for surface activity.  Finally, we mention the regime with so-called ``dynamo outbursts'' observed in the VKS dynamo experiment (Ravelet {\\it et al.}, private communication) as presenting one further topic that may be relevant to our investigation. We note that episodes from the time series of magnetic field evolution taken from a sensor in this experiment look very similar to that presented in Figure~4.  We recognize that our results are not positive in the sense of answering in a clear-cut manner the key question of whether a simple model, such as we have considered, is able to reproduce the observed statistics of the occurrence of solar Grand Minima. However we do feel that we may have provided some insight into the question of the statistical stability of the observations, and to have provided information and guidance for future investigations.            \\begin{acks}  D.S. is grateful  to the  Royal Society for supporting his visit to Manchester, and also to RFBR for financial support under grant 07-02-00127. D.M. and D.S. thank the Finnish Academy of Science and Letters (V\\\"ais\\\"al\\\"a foundation) for supporting their visits to the University of Oulu.   \\end{acks}"
    },
    "0806/0806.4382_arXiv.txt": {
        "abstract": "We argue that the cosmological constant is exponentially suppressed in a candidate ground state of loop quantum gravity as a nonperturbative effect of a holographic Fermi-liquid theory living on a two-dimensional spacetime. Ashtekar connection components, corresponding to degenerate gravitational configurations breaking large gauge invariance and CP symmetry, behave as composite fermions that condense as in Bardeen--Cooper--Schrieffer theory of superconductivity. Cooper pairs admit a description as wormholes on a de Sitter boundary. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0806/0806.2047_arXiv.txt": {
        "abstract": "From our radio continuum and polarization observations of a sample of spiral galaxies with different morphological types, inclinations, and star formation rates (SFR) we found that galaxies with low SFR have higher thermal fractions/ smaller synchrotron fractions than those with normal or high SFR.  Adopting an equipartition model, we concluded from our observations that the nonthermal radio emission and the \\emph{total magnetic field} strength grow nonlinearly with SFR.  We also studied the magnetic field structure and disk thicknesses in highly inclined (edge-on) galaxies. We found in five galaxies that -despite their different radio appearance- the vertical scale heights for both, the thin and thick disk/halo, are about equal (0.3/1.8kpc), independently of their different SFR. They also show a similar large-scale magnetic field configuration, parallel to the midplane and X-shaped further away from the disk plane, independent of Hubble type and SFR in the disk. Hence we conclude that the amplification and formation of the \\emph{large-scale} magnetic field structure is independent of SFR. ",
        "introduction": "Radio observations of the continuum emission in the cm-wavelength regime are the best way to study magnetic fields in galaxies. Magnetic fields consist of regular and turbulent components. The total magnetic field strength in a galaxy can be estimated from the nonthermal radio emission under the assumption of equipartition between the energies of the magnetic field and the relativistic particles (the so-called {\\em energy equipartition}) as described in Beck \\& Krause (2005).\\\\ The linear polarization reveals the strength and structure of the resolved regular magnetic field. Its analysis and the observational results of the field strengths and patterns of nearby face-on galaxies are summarized e.g. by Beck (this volume).\\\\ ",
        "conclusions": ""
    },
    "0806/0806.0042_arXiv.txt": {
        "abstract": "We present a simple and efficient anisotropic generalization of the semi-isotropic (two-integral) axisymmetric Jeans formalism which is used to model the stellar kinematics of galaxies. The following is assumed: (i) a constant mass-to-light ratio $M/L$ and (ii) a velocity ellipsoid that is aligned with cylindrical coordinates $(R,z)$ and characterized by the classic anisotropy parameter $\\beta_z=1-\\overline{v_{z}^2}/\\overline{v_{R}^2}$. Our simple models are fit to \\sauron\\ integral-field observations of the stellar kinematics for a set of fast-rotator early-type galaxies. With only two free parameters ($\\beta_z$ and the inclination) the models generally provide remarkably good descriptions of the shape of the first ($V$) and second ($V_{\\rm rms}\\equiv\\sqrt{V^2+\\sigma^2}$) velocity moments, once a detailed description of the surface brightness is given. This is consistent with previous findings on the dynamical structure of these objects. With the observationally-motivated assumption that $\\beta_z\\ga0$, the method is able to recover the inclination. The technique can be used to determine the dynamical mass-to-light ratios and angular momenta of early-type fast-rotators and spiral galaxies, especially when the quality of the data does not justify more sophisticated modeling approaches. This formalism allows for the inclusion of dark matter, supermassive black holes, spatially varying anisotropy, and multiple kinematic components. ",
        "introduction": "According to the theory that best reproduces the observations, galaxy formation proceeds in a hierarchical fashion, driven by gravity, in a Universe dominated by dark matter of unknown nature \\citep[e.g.][]{springel05nat}. The hierarchy of merging is accompanied by changes in galaxy structure and morphology. In particular early-type galaxies (Es and S0s) are thought to form by the gas-rich merging of spiral galaxies or by gas starvation of spirals, followed by subsequent collisionless mergers \\citep[e.g.][]{faber07}.  Three key global parameters can be used to characterize galaxies structure while studying this sequence of merging of galaxies and dark matter: (i) The angular momentum, which varies during mergers and increases with the amount of gas dissipation, (ii) the stellar population, which records the history of star formation events during the gas-rich mergers, and (iii) the mass-to-light ratio, which is affected by both the population and by the dark-matter fraction.  The large majority of the galaxies in the Universe are to first order axisymmetric (except for bars) and posses disks. This includes fast-rotator early-type galaxies \\citep{emsellem07,cappellari07} and spiral galaxies. Both the fast-rotator early-types \\citep{gerhard01,rusin05,cappellari06,koopmans06,thomas07} and the spiral galaxies \\citep{persic96,palunas00,bell01,kassin06} appear dominated by the stellar matter inside one half-light radius. Observations suggest that they have a dynamical structure characterized by a flattening of the velocity ellipsoid in the $z$ direction parallel to the galaxy symmetry axis \\citep{gerssen97,gerssen00,shapiro03,cappellari07,Noordermeer08}.  The goal of this paper is to include the knowledge of the structure of the fast-rotator and spiral galaxies, into a realistic but simple dynamical modeling method, which can be applied to the measurement of both the mass-to-light ratio and the amount of rotation (for which the inclination is needed) in the central regions of these galaxies, while also allowing for the inclusion of dark matter and the study of multiple kinematical components. The success of the adopted model's assumptions in describing the kinematics of real galaxies is verified against integral-field observations of the stellar kinematics obtained with \\sauron\\ \\citep{bacon01}.  In \\refsec{sec:solving_jeans} we briefly review the theory and past applications of the Jeans equations and of the shape of the velocity ellipsoid in galaxies. In \\refsec{sec:anisotropic_jeans_solutions} we describe our new anisotropic Jeans formalism, which we apply and test in \\refsec{sec:tests}. Finally our results are summarized in \\refsec{sec:summary}. ",
        "conclusions": "\\label{sec:summary}  We present a generalization of the widely used semi-isotropic (two-integral) axisymmetric Jeans modeling method to describe the stellar dynamics of galaxies. Our method uses the powerful Multi-Gaussian Expansion (MGE) technique to accurately parameterize the galaxies photometry. It represents an anisotropic extension of what was presented in the semi-isotropic case by \\citet{emsellem94}, and it maintains its simplicity and computational efficiency. We assume (i) a constant mass-to-light ratio and (ii) a velocity ellipsoid which is aligned with the cylindrical $(R,z)$ coordinates and has a flattening quantified by the classical $z$-anisotropy parameter $\\beta_z=1-\\overline{v_{z}^2}/\\overline{v_{R}^2}$, where $z$ is the galaxy symmetry axis.  We test the technique using \\sauron\\ integral-field observations of the stellar kinematics \\citep{emsellem04} for a small set of fast-rotator galaxies with a variety of kinematical properties. For galaxies that are constrained by the photometry to be close to edge-on we find that, although in the semi-isotropic limit ($\\beta_z=0$) the models do not provide a good description of the data, the variation of the {\\em single} global anisotropy $\\beta_z$ is generally sufficient to accurately predict the shape of both the first ($V$) and second ($V_{\\rm rms}=\\sqrt{V^2+\\sigma^2}$) velocity moments, once an detailed MGE parametrization of the photometry is given. An accurate description of the photometry, including ellipticity variations and disky isophotes, appears crucial to reproduce in detail the features of the kinematics.  In all cases we find that $\\beta_z\\ga0.05$, while generally $\\sigma_{\\phi}^2\\approx\\overline{v_{R}^2}$ with good accuracy. This confirms previous findings on the dynamical structure of these galaxies, showing that their velocity ellipsoid tends to be oblate \\citep{cappellari07}. The anisotropy derived with our anisotropic Jeans dynamical modeling method agrees within the errors with the one previously measured using a more general Schwarzschild approach.  For fast-rotator galaxies that are {\\em not} constrained by the photometry to be close to edge-on, we find that in general the inclination $i$ (or the corresponding galaxy shape) and the anisotropy $\\beta_z$ are highly correlated and cannot be independently determined. However, if we introduce the observationally-motivated constraint $\\beta_z\\ga0.05$, the inclination becomes constrained to a narrow range of values and it agrees with independent determinations when those are available.  We are applying this method to determine the inclination, the mass-to-light ratio and the amount of rotation of a large sample of galaxies for which integral field kinematics are available. For galaxies close to edge-on, the global anisotropy or the dynamical structure of different galaxy subcomponents (e.g.\\ bulge and disk) can also be investigated. We are using this method to test independent determinations of the masses of supermassive black holes in the nuclei of fast-rotator galaxies. This technique is ideal to study the dark matter content and the anisotropy of disks of spiral galaxies."
    },
    "0806/0806.0618_arXiv.txt": {
        "abstract": "We have used the IRAM Plateau de Bure millimeter interferometer and the UKIRT 1--5\\,$\\mu$m Imager Spectrometer (UIST) to test the connection between the major phases of spheroid growth and nuclear accretion by mapping CO emission in nine submillimetre-detected QSOs at $z=1.7$--2.6 with black hole (BH) masses derived from near-infrared spectroscopy. When combined with one QSO obtained from the literature, we present sensitive CO(3--2) or CO(2--1) observations of 10 submillimetre-detected QSOs selected at the epoch of peak activity in both QSOs and submillimetre (submm) galaxies (SMGs). CO is detected in 5/6 very optically luminous ($M_\\mathrm{B}\\sim -28$) submm-detected QSOs with BH masses $M_\\mathrm{BH}\\simeq10^9$--10$^{10}$\\,M$_\\odot$, confirming the presence of large gas reservoirs of $M_\\mathrm{gas}\\simeq3.4\\times10^{10}$\\,M$_\\odot$.   Our BH masses and dynamical mass constraints on the host spheroids suggest, at face value, that these optically luminous QSOs at $z=2$ lie about an order of magnitude above the local BH-spheroid relation, $M_\\mathrm{BH}/M_\\mathrm{sph}$, although this result is dependent on the size and inclination of the CO-emitting region. However, we find that their BH masses are $\\sim30$ times too large and their surface density is $\\sim300$ times too small to be related to typical SMGs in an evolutionary sequence.  Conversely, we measure weaker CO emission in four fainter ($M_\\mathrm{B}\\sim-25$) submm-detected QSOs with properties, BH masses ($M_\\mathrm{BH}\\simeq5\\times10^8$\\,M$_\\odot$), and surface densities similar to SMGs.  These QSOs appear to lie near the local $M_\\mathrm{BH}/M_\\mathrm{sph}$ relation, making them plausible `transition objects' in the proposed evolutionary sequence linking QSOs to the formation of massive young galaxies and BHs at high-redshift.  We show that SMGs have a higher incidence of bimodal CO line profiles than seen in our QSO sample, which we interpret as an effect of their relative inclinations, with the QSOs seen more face-on. Finally, we find that the gas masses of the four fainter submm-detected QSOs imply that their star formation episodes could be sustained for $\\sim10$\\,Myr, and are consistent with representing a phase in the formation of massive galaxies which overlaps a preceding SMG starburst phase, before subsequently evolving into a population of present-day massive ellipticals. ",
        "introduction": "It has been established that every massive, local spheroid harbours a supermassive black hole (SMBH) in its centre whose mass is proportional to that of its host (e.g.\\ \\citealt{Magorrian98}; \\citealt{Gebhardt00}). This suggests that the black holes (BHs) and their surrounding galaxies were formed synchronously.  This hypothesis has found support from hydrodynamical simulations of galaxy formation, which use feedback from winds and outflows from active galactic nuclei (AGN) to link the growth of the SMBH to that of its host (e.g.\\ \\citealt{DiMatteo05}; \\citealt{Hopkins05}; \\citealt{Bower06}). Thus these models support a picture, first presented by \\citet{Sanders88}, where a starburst-dominated ultra-luminous infrared galaxy (ULIRG), arising from a merger, evolves first into an obscured QSO and then into an unobscured QSO, before finally becoming a passive spheroid.  The high-redshift population of ULIRGs in this proposed evolutionary cycle are the submillimetre (submm) galaxies (SMGs; \\citealt{SIB97}; \\citealt{Chapman05}; \\citealt{Coppin06}).  These systems have ULIRG-like bolometric luminosities, $L_\\mathrm{IR}\\geq 10^{12}$\\,L$_\\odot$ (\\citealt{Kovacs06}; \\citealt{Coppin07}), and they have many of the properties expected for gas-rich mergers (\\citealt{Swinbank04,Swinbank06}; \\citealt{Tacconi06}).  This population evolves rapidly out to a peak at $z\\sim 2.3$, crudely matching the evolution of QSOs \\citep{Chapman05} and providing additional circumstantial evidence for a link between SMBH growth and spheroids.  Two further results have shed light on the evolutionary link between SMGs and QSOs.  Firstly, a modest fraction of optically luminous QSOs at $z\\sim 2$ are detected in the submm/mm ($\\sim(25\\pm10)$\\%; \\citealt{Omont03}) showing that the QSO- and SMG-phases do not overlap significantly, given the lifetime estimates of the two populations (QSOs make up $\\sim 4$\\% of flux-limited samples of SMGs; \\citealt{Chapman05}).  But when a QSO {\\it is} detected in the submm/mm then it could to be in the transition phase from an SMG to an unobscured QSO, making its properties a powerful probe of the evolutionary cycle (e.g.\\ \\citealt{Stevens05}; \\citealt{Page04}).  Secondly, the evolutionary state of the SMBHs within SMGs can also be judged using the 2-Ms {\\it Chandra} Deep Field North observations \\citep{Alexander03} to derive accurate AGN luminosities and hence lower limits on the BH masses ($M_\\mathrm{BH}$) in those SMGs with precise redshifts in this region (\\citealt{Alexander05nat,Alexander05}; \\citealt{Borys05}).  These studies suggest that the AGN in typical SMGs are growing almost continuously -- but that the SMBHs in these galaxies appear to be several times less massive than seen in comparably massive galaxies at $z\\sim 0$ \\citep{Alexander07}.  Together these results argue for a fast transition from a star-formation-dominated SMG-phase to the AGN-dominated QSO-phase \\citep{Page04}. The latter phase will result in the rapid BH growth necessary to account for the present-day relation between spheroid and SMBH masses (e.g.\\ \\citealt{Magorrian98}; \\citealt{Gebhardt00}). {\\it Can we confirm this and more generally test the proposed evolutionary link between SMGs and QSOs at the peak of their activity at $z\\sim 2$?}  This evolutionary cycle has been tested in the local Universe by comparing the properties of QSOs and ULIRGs (e.g.\\ \\citealt{Tacconi02}). However, both of these populations are $1000\\times$ less abundant in the local Universe than they were at the era of their peak activity at $z\\sim2$ and so we have to be cautious about extrapolating from local examples to the high-redshift progenitors of the bulk of today's massive spheroids (\\citealt{Genzel03}; \\citealt{Swinbank06}). Thus, to properly test the validity of this cycle for typical spheroids we have to compare QSOs and ULIRGs at the era where their populations peaked:  $z\\sim2$.  The critical pieces of information needed to test the link between SMGs and QSOs are the relative dynamical, gas and SMBH masses of these two populations. In principle the dynamical masses can be derived from optical or near-infrared observations of emission line gas in the SMGs or QSOs (see \\citealt{Swinbank04,Swinbank05,Swinbank06}). However, there is the problem of removing the QSO emission in these observations, as well as the effects of dust obscuration and outflows. In contrast, molecular CO emission line profiles are relatively immune to the effects of obscuration and outflows, while at the same time yielding additional constraints on the relationship between QSOs and SMGs from their gas masses.  There is currently a lack of sensitive CO observations of QSOs at $z\\sim 1$--3, with data published on only eight sources (e.g.~\\citealt{Frayer98}; \\citealt{Guilloteau99}; \\citealt{Beelen04}; \\citealt{Hainline04}). Instead the focus has been on CO studies of QSOs at $z\\gsim4$ (e.g.\\ \\citealt{Omont96}; \\citealt{Walter04}; \\citealt{Riechers06}), although these QSOs have little overlap with the redshift range where SMGs are typically detected. The paucity of CO constraints for $z>1$ QSOs also reflects the difficulty in determining their systemic redshifts with sufficient precision to guarantee that the CO emission falls within the bandwidth of typical millimetre (mm) correlators.  However, sensitive near-infrared spectroscopy of the C{\\sc iv}, Mg{\\sc ii}, and [O{\\sc iii}]5007 emission lines in QSOs can provide redshifts with required precision as well as H$\\alpha$ or H$\\beta$ fluxes and line widths to yield $M_\\mathrm{BH}$ estimates (\\citealt{Takata07}; \\citealt{Alexander07}).  We have carried out a quantitative test of the proposed link between SMGs and QSOs at $z\\sim 2$ where both populations are most common.  We have obtained precise systemic redshifts from near-infrared spectroscopy of potential transition QSOs (i.e.~submm/mm-detected QSOs) and then used the IRAM Plateau de Bure Interferometer (PdBI) to search for CO emission.  We relate their dynamical, gas and SMBH masses to SMGs from the PdBI CO survey \\citep{Greve05}. We test: a) whether the cold gas masses in these QSOs are similar to those in SMGs; b) whether the line widths and dynamical masses of these two populations are comparable; and c) how the ratio of SMBH to dynamical masses for these submm-detected QSOs relate to the estimates for SMGs and those for optically luminous QSOs (which lie on the present-day $M_\\mathrm{BH}$--$M_\\sigma$ relation; \\citealt{McLure04}). Together these observations can constrain the proposed evolutionary sequence which links QSOs to the formation of massive young galaxies and SMBHs at high redshift.  We describe the sample selection, observations and data reduction in \\S \\ref{obsdr}. The results of the near-infrared and mm CO spectra are given in \\S \\ref{results}. The CO properties of the submm-detected QSOs are compared and contrasted with SMGs in \\S \\ref{discuss}, and we discuss the evolutionary status of the submm-detected QSOs in \\S \\ref{evolution}.  Our conclusions are given in \\S \\ref{concl}. We adopt cosmological parameters from the \\textit{WMAP} fits in \\citet{Spergel03}: $\\Omega_\\Lambda=0.73$, $\\Omega_\\mathrm{m}=0.27$, and $H_\\mathrm{0}=71$\\,km\\,s$^{-1}$\\,Mpc$^{-1}$.  All quoted magnitudes are on the Vega system.  \\setcounter{figure}{0}  \\begin{figure} \\epsfig{file=Rband.ps,width=0.5\\textwidth} \\caption{$R$-band apparent magnitude versus redshift for our sample of submm-detected QSOs.  We show that these lie within the interquartile range (shaded region) of the redshift distribution of SMGs (median $<\\!z\\!>\\simeq2.3$; dashed line) from \\citet{Chapman05}.  This shows that we are probing a wide range in optical luminosities (and potentially a wide range in $M_\\mathrm{BH}$) at the epoch where the QSO and SMG populations peak. } \\label{fig:selection} \\end{figure} ",
        "conclusions": "We have carried out a millimetre interferometry survey of nine submm-detected QSOs at $z=1.7$--2.6 in order to test the link between SMGs and QSOs at the era where these two important populations were most numerous.  We include in our analysis comparable observations of a similarly selected QSO from the literature to provide a final sample of ten submm-detected QSOs. To support this survey we obtained near-infrared spectroscopy of these QSOs to derive accurate systemic redshifts needed to tune the millimetre receivers to the correct frequencies.  The near-infrared spectra also provide H$\\alpha$ fluxes and line widths needed to derive reliable BH mass estimates for the QSOs.  Our main findings are:  \\smallskip (1) We detect CO emission from six of the ten submm-detected QSOs in our sample, confirming that they contain a significant amount of molecular gas and that a large fraction of the mm emission is from starbursts.  The median gas mass of our sample (including non-detections) is $(2.5\\pm 0.7)\\times 10^{10}$\\,M$_\\odot$, similar to that found for $z\\sim 2$ SMGs and to $z\\gsim4$ QSOs. The star formation efficiencies of our QSOs are also comparable to those measured for SMGs, $250\\pm100$\\,L$_\\odot\\,\\mathrm{(K\\,km\\,s^{-1}\\,pc^{2})^{-1}}$, suggesting that the gross properties of the star formation in the QSOs are like those seen in SMGs.   Adopting a 2\\,kpc scale size for the gas distribution in the QSOs and a typical inclination of 20$^\\circ$ we derive a median dynamical mass of M$(<2$\\,kpc$)\\sim(2.1\\pm1.4)\\times 10^{11}$\\,M$_\\odot$, similar to SMGs (assuming an inclination angle appropriate for random inclinations). We find a lower incidence of double-peaked CO line profiles in the QSOs, compared to SMGs, which we believe results from a selection bias towards lower average inclination angles for the QSOs.  \\smallskip (2) Our near-infrared spectroscopy indicates a median black hole mass in our QSO sample of $(1.8\\pm1.3)\\times10^{9}$\\,M$_\\odot$. Combined with our dynamical estimates of the spheroid mass, these yield $M_\\mathrm{BH}/M_\\mathrm{sph}\\sim 9\\times10^{-3}$.  This $M_\\mathrm{BH}/M_\\mathrm{sph}$ ratio for this sample of submm-detected QSOs at $z=2$ is an order of magnitude larger than the local ratio, although $M_\\mathrm{sph}$ suffers from large uncertainties due to the unknown CO radii and inclination angles.  This ratio is also significantly above that seen for SMGs at $z\\sim 2$.  However, this comparison masks a broad range in BH masses within our QSO sample and so we split the sample into two subsets based on their BH masses.  \\smallskip (3) Looking at the optically luminous submm-detected QSOs in our sample we find that we detect CO emission in 5/6 of these QSOs. However, the estimated BH masses for these QSOs, $M_\\mathrm{BH}\\simeq10^{9}$--10$^{10}$\\,M$_{\\odot}$, are too large (and their number densities too small) for them to be related to typical SMGs in a simple evolutionary cycle.  We propose that the progenitors of these most massive QSOs are a rare subset of SMGs with $M_\\mathrm{gas}>4\\times10^{11}$\\,M$_\\odot$ with a number density of $\\simeq10\\,\\mathrm{deg}^{-2}$ which will be possible to detect with future SCUBA-2 surveys.  \\smallskip (4) For the optically less luminous ($\\sim L^{\\star}$) submm-detected QSOs, we marginally detect one source in CO and obtain sensitive limits for three further QSOs.  The BH masses for these systems are $M_\\mathrm{BH}\\simeq10^{8}$\\,M$_\\odot$, similar to the estimates for BHs in SMGs.  These submm-detected QSOs are consistent with being `transition' objects between SMGs and submm-undetected QSOs, as we show it is feasible to link their BH masses to those of SMGs by Eddington limited growth for a period comparable to the gas depletion timescale of the QSOs, $\\sim 10$\\,Myrs.  The space density of these QSOs is also in rough agreement with that expected for the descendents of SMGs given current estimates of the relative lifetimes of QSOs and SMGs.  We conclude that these $\\sim L^{\\star}$, $M_\\mathrm{BH}>10^{8}$\\,M$_\\odot$ submm-detected QSOs are consistent with being in a very brief prodigious star formation phase, and that they simply do not possess sufficiently large gas reservoirs to sustain the SFR (which is why these might be less often detected in CO), although a larger sample of CO observations of submm-detected QSOs with these BH masses is required for confirmation.  \\bigskip To make further progress on understanding the evolutionary links between SMGs and QSOs requires a larger survey of the submm and CO emission from typical QSOs ($M_\\mathrm{B}\\approx -25$).  In addition, measurements of other CO transitions for the submm-detected QSOs (e.g.~from IRAM 30-m, ALMA, EVLA, and SKA) are required to place better constraints on the temperature and density of the molecular gas and thus provide a more accurate determination of the line luminosity ratios and hence total gas masses of these systems. Similarly, higher resolution CO observations are essential to put strong constraints on the reservoir sizes and inclination angles, and hence $M_\\mathrm{dyn}$, needed to compare the two populations.  Finally, better measurements of the far-infrared SEDs (with SABOCA, SCUBA-2 or \\textit{Herschel}) will yield more accurate measures of $L_\\mathrm{FIR}$ and $T_\\mathrm{dust}$ for submm-detected QSOs to constrain the contribution from an AGN component."
    },
    "0806/0806.4936_arXiv.txt": {
        "abstract": "{} {We aim at directly detecting the presence of optically thin circumstellar dust emission within the terrestrial planetary zone around main sequence stars known to harbour cold debris discs. The present study focuses on a sample of six bright A- and early F-type stars.} {High-precision interferometric observations have been obtained in the near-infrared $K$ band with the FLUOR instrument installed on the CHARA Array. The measured squared visibilities are compared to the expected visibility of the stellar photospheres based on theoretical photospheric models taking into account rotational distortion. We search for potential visibility reduction at short baselines, a direct piece of evidence for resolved circumstellar emission.} {Our observations bring to light the presence of resolved circumstellar emission around one of the six target stars ($\\zeta$\\,Aql) at the $5\\sigma$ level. The morphology of the emission source cannot be directly constrained because of the sparse spatial frequency sampling of our interferometric data. Using complementary adaptive optics observations and radial velocity measurements, we find that the presence of a low-mass companion is a likely origin for the excess emission. The potential companion is characterised by a $K$-band contrast of four magnitudes. It has a most probable mass of about $0.6 M_{\\odot}$ and is expected to orbit between about 5.5\\,AU and 8\\,AU from its host star assuming a purely circular orbit. Nevertheless, by adjusting a physical debris disc model to the observed Spectral Energy Distribution of the $\\zeta$\\,Aql system, we also show that the presence of hot dust within 10\\,AU from $\\zeta$\\,Aql, producing a total thermal emission equal to $1.69\\pm 0.31$\\% of the photospheric flux in the $K$ band, is another viable explanation for the observed near-infrared excess. Our re-interpretation of archival near- to far-infrared photometric measurements shows however that cold dust is not present around $\\zeta$\\,Aql at the sensitivity limit of the IRS and MIPS instruments onboard Spitzer, and urges us to remove $\\zeta$\\,Aql from the category of bona fide debris disc stars.} {The hot debris disc around Vega (Absil et al.\\ 2006) currently remains our only secure resolved detection within the context of this survey, with six genuine early-type debris disc stars observed so far. Further observations will be needed to assess whether $\\zeta$\\,Aql also belongs to this hot debris disc category.} ",
        "introduction": "Debris discs are optically thin, gas-poor dust discs around main sequence (MS) stars. The presence of circumstellar dust around stars with ages above $\\sim$10~Myr is attributed to populations of planetesimals that were neither used to make up planets nor ejected from the system by the time the nebular gas was dispersed \\citep{Mann06}. These leftovers produce dust by mutual collisions and comet-type activity. Being continuously replenished by small bodies, the disc can then persist over much of the star's lifetime. Due to its large total cross-section area, dust is much easier to observe than planets, not to speak of planetesimals. On the other hand, distributions of dust respond to the presence of planetary perturbers, reflect distributions of the parent bodies and bear important memory of the planetary formation process in the past. Hence debris discs can be used as sensitive tracers of planets, as well as small body populations, and should reflect evolutionary stages of planetary systems. This explains the substantial effort invested in the observation and modelling of debris discs over the last two decades.  Early-type stars, and A-type stars in particular, have been the most successfully studied targets in the context of debris disc studies so far. This is due to a great extent to their intrinsic brightness, which efficiently lights up their debris discs up to large distances (Earth-like temperatures occur at about 5\\,AU from such stars). Additionally, their MS lifetimes ($\\sim$800~Myr) are long enough to encompass the main evolutionary stages of a typical planetary system. Far-infrared surveys carried out with space satellites have been particularly successful in detecting cold dust around nearby A-type stars. For instance, Spiter/MIPS observations of 160~A-type MS stars have shown that about 33\\% of such stars possess significant excess emission at 24 and 70~$\\mu$m \\citep{Su06}, which represents a considerably higher excess rate than what has been found for old solar analogs and M dwarfs \\citep{Bryden06}.  Unlike these successful detections of cold debris in the far-infrared, the discovery of warm ($\\sim$300\\,K) or hot ($\\sim$1000\\,K) dust through mid- and near-infrared observations has been limited to a very small number of targets so far. Furthermore, the rare detections have generally been obtained for young MS stars showing strong mid-infrared silicate features, such as $\\beta$~Pic \\citep[$\\sim$12\\,Myr,][]{Pantin97}, HD\\,145263 \\citep[$\\sim$8\\,Myr,][]{Honda04}, $\\eta$~Tel \\citep[$\\sim$12\\,Myr,][]{Chen06} or HD\\,172555 \\citep[$\\sim$12\\,Myr,][]{Chen06}. At such ages, planetary systems are still expected to be in the process of forming planets, especially in the inner part of the disc ($<10$\\,AU) where rocky planets may take up to one hundred Myr to accrete most of the small bodies on nearby orbits and reach their final mass \\citep[see][for a solar-mass star]{Kenyon06}. The observed large mid-infrared excesses are therefore expected to be related to the end of the planet building phase rather than to the dust produced by the collisional grinding of an evolved planetary system, similar to the zodiacal dust in our solar system. For ``mature'' A-type stars, the lack of near- to mid-infrared excess has generally been associated with a dearth of warm dust grains, a suggestion generally confirmed by the presence of inner holes in resolved images around famous objects such as $\\alpha$~PsA~\\citep{Kalas05} or $\\alpha$~Lyr~\\citep{Su05}.  Incidentally, the inner part of debris discs is a very interesting region to study, as it directly probes the location where planets are supposed to have formed and evolved---conversely, the outer part of the disc ($>20$\\,AU), similarly to the Kuiper belt in our solar system, only bears a memory of the outermost massive planets through gravitational interactions such as mean motion resonances \\citep[see e.g.][]{Reche08}. The characterisation of inner debris discs is therefore crucial for an understanding of the formation, evolution and dynamics of planetary systems (including our own solar system), as well as to set the scene for the emergence of life on rocky planets.  Reaching a better sensitivity to the inner part of debris discs has thus been an important challenge during the past years, and the detection of small amounts of hot dust around mature early-type stars has finally been enabled by the advent of high-precision near-infrared stellar interferometry. Using the VLTI/VINCI instrument, \\citet{DiFolco04} derived upper limits to the $K$-band dust emission around five stars. The first robust resolved detection of hot dust was obtained for the bright A0V-type star $\\alpha$~Lyr (Vega) by \\citet{Absil06} with CHARA/FLUOR on an optimised set of interferometric baselines, showing the presence of extended emission accounting for $1.29 \\pm 0.19$\\% of the photospheric flux in the $K$ band. This result was already suggested by \\citet{Ciardi01}, although with a large uncertainty on the flux ratio.  This recent detection has raised a number of questions regarding the nature and origin of inner dust grains. In particular, a scenario involving the presence of star-grazing comets, injected into the inner planetary system by dynamical perturbations caused by migrating planets, has been proposed, inspired by the Late Heavy Bombardment (LHB) that happened in the early solar system~\\citep{Gomes05}. A similar scenario is also proposed by \\citet{Wyatt07a} to explain the presence of warm dust ($\\sim$300\\,K) detected by the Spitzer Space Telescope around a few solar-type stars. We have thus decided to initiate a near-infrared interferometric survey of nearby debris disc stars to assess the occurrence of hot excesses around MS stars. The first paper of this series \\citep[][hereafter \\citetalias{DiFolco07}]{DiFolco07} focused on two solar-type stars, showing the presence of hot dust around $\\tau$\\,Cet. In this paper, we discuss early-type stars, which hold a privileged position in our survey because their brightness makes them well suited for an interferometric study. ",
        "conclusions": "In this paper, we have investigated the close neighbourhood ($<1 \\arcsec$) of six nearby A- and early F-type MS stars, in search for hot counterparts to the cold debris discs detected by mid- and far-infrared spectro-photometric space-based observations. The high-accuracy squared visibilities collected with the CHARA/FLUOR interferometer, combined with semi-empirical models of stellar photospheres including rotational distortion, has allowed us to reach dynamic ranges ranging from 1:175 to 1:625 at $1\\sigma$ around the target stars. At this level of precision, the presence of a resolved $K$ band emission has been identified around only $\\zeta$\\,Aql, with a estimated $K$-band excess of $1.69\\pm 0.31$\\%. This detection adds to our previous results on $\\alpha$\\,Lyr \\citep{Absil06} and $\\tau$\\,Cet \\citepalias{DiFolco07}, giving an overall near-infrared excess detection rate of $3/9$ for the MS stars surveyed so far, among which $2/7$ are early-type stars. The healthy statistical behaviour of the five non-detections in the present sample and the confirmation of the excess emission around $\\alpha$\\,Lyr with an improved photospheric model and a new version of the FLUOR Data Reduction Software demonstrate the robustness of our approach for hot debris disc detection.  Our near-infrared interferometric measurements are not sampling the Fourier frequency plane in a sufficiently dense manner to derive the morphology of the excess emission source. In particular, both a point-like source and an extended circumstellar emission can reproduce our observations. While in the cases of $\\alpha$\\,Lyr and $\\tau$\\,Cet, the presence of a bound or unbound companion to the target stars within the small FLUOR field-of-view could be rejected with a high confidence, we cannot rule out the presence of a low-mass companion in the close vicinity of $\\zeta$\\,Aql to reproduce the measured $K$-band excess. The combination of the astrometric stability of $\\zeta$\\,Aql measured by {\\sc Hipparcos}, the variability of the radial velocities measured with HARPS and the absence of off-axis companion in PUEO observations restricts the parameter space of the high-contrast binary scenario to masses in the range $0.6$ to $0.65\\,M_{\\odot}$ and semi-major axes between 5.5 and 8\\,AU (i.e., about 200 and 300\\,mas). The $K$-band contrast between the primary and its companion would then be $\\Delta K \\sim 4$, making it one of the closest high-contrast companions resolved around MS stars so far.  Besides a low-mass companion, the presence of hot circumstellar dust grains producing a significant thermal emission in the $K$ band is another viable explanation of the observed excess emission. In particular, we show that the debris disc model that \\citet{Absil06} have proposed in the context of $\\alpha$\\,Lyr is consistent with both our $K$-band detection and archival near- and mid-infrared spectro-photometric measurements. However, our re-interpretation of archival Spitzer/MIPS and Spitzer/IRS data clearly shows that the presence of an outer debris disc, suggested by \\citet{Chen05}, can be firmly ruled out at the sensitivity level of MIPS and IRS. In the absence of significant amounts of cold dust, the hot debris disc scenario is not favoured to explain our CHARA/FLUOR measurements.  The statistics of the ``hot debris disc'' phenomenon presently remains poorly constrained: for early-type stars, a hot debris disc was found around only one star ($\\alpha$\\,Lyr) out of six bona fide debris disc stars observed so far. The case of $\\alpha$\\,Lyr could therefore be rather unusual. Our interferometric survey of debris disc stars will be extended to larger sample in the coming years to improve our statistics. The case of $\\zeta$\\,Aql must be considered separately, since we show that it is not surrounded by cold dust at the sensitivity level of the Spitzer instruments, and since a close companion is a likely explanation to the observed $K$-band excess. This star will deserve special attention in the future, and further observations will be performed to determine the actual nature of the near-infrared excess emission that we have resolved. In particular, infrared aperture masking experiments on large telescopes have the potential to reveal the true nature of the observed excess."
    },
    "0806/0806.1452_arXiv.txt": {
        "abstract": "We study the perturbative behaviour of topological black holes. We calculate both analytically and numerically the quasi-normal modes of scalar perturbations. In the case of small black holes we find discontinuities of the quasi-normal modes spectrum at the critical temperature and we argue that this is evidence of a second-order phase transition. ",
        "introduction": "The knowledge of the spectrum of the quasi-normal modes (QNMs) is a powerful tool in the study of the late time behaviour of black holes. The QNMs are the complex frequencies by which a black hole responds if it is initially perturbed, and they do not depend on the details of the initial perturbation but rather on the intrinsic features of the black hole itself. The radiation associated with these modes is expected to be seen with gravitational wave detectors in the coming years, giving valuable information on the properties of black holes.  The QNMs of black holes in asymptotically flat spacetimes have been extensively studied and their spectrum was computed numerically and in many cases also analytically (for reviews, see~\\cite{KS,N}). The advances in string theory and mainly the Anti-de Sitter - conformal field theory (AdS/CFT) correspondence has renewed the interest of computing the QNMs of black holes in asymptotically AdS spacetimes~\\cite{CM,HH,CL,WLA,kokkotas,Konoplya}. Recent results from the Relativistic Heavy Ion Collider \\cite{RHIC} show that a thermal quark-gluon plasma (QGP) is formed which is strongly coupled. The AdS/CFT correspondence provides the connection between the QGP and string black holes. In \\cite{Friess:2006kw} the quasi-normal modes of AdS$_{5}$-Schwarzschild black hole have been calculated and it was shown that they provide a dual description of the fluctuations of the QGP. Recently, in a more phenomenological approach, the AdS/CFT correspondence was applied to condensed matter physics~\\cite{condensed_Matter}. It was shown~\\cite{Hartnoll:2008vx} that fluctuations of the metric and the background electromagnetic field determines the conductivity of the boundary theory.    According to the AdS/CFT correspondence, a large static black hole in AdS corresponds to an (approximately) equilibrium thermal state in the CFT~\\cite{cft}. In ref.~\\cite{HH} it was shown that the QNMs for the scalar perturbations of large Schwarzschild-AdS black holes scaled with the temperature and it was argued that the perturbed system in the dual description will approach to thermal equilibrium of the boundary conformal field theory. However, when the black hole size is comparable to the AdS length scale there is a clear departure from this behaviour. It was then conjectured that this behaviour may be connected with a Hawking-Page phase transition \\cite{phase1, phase2} which occurs when the temperature lowers.  For small black holes the behaviour of QNMs is not very well understood. For the case of Schwarzschild-AdS black holes the QNMs do not scale linearly with the temperature any more and their connection with the boundary conformal field theory is not clear. In~\\cite{Koutsoumbas:2006xj,Koutsoumbas:2008pw} we studied the QNMs of electromagnetic and gravitational perturbations of topological black holes (TBHs) in AdS spacetime coupled to a scalar or an electromagnetic field. For small black holes, compared to the length scale  of the AdS space, we found that the QNMs behave quite differently from the QNMs of large black holes. Near the critical temperature purely dissipative modes appeared in the spectrum and the QNMs change slope around that point. We attributed this behaviour to a second order phase transition of the charged TBH towards the AdS vacuum solution.  In the literature there is a discussion of possible connections between the classical and thermodynamical properties of black holes~\\cite{Reall:2001ag}. In particular the question whether the knowledge of the QNM spectrum can give information about thermodynamical phase transitions in a wider class of black holes has recently gained considerable interest. It was suggested in~\\cite{Jing:2008an} that the Dirac and Rarita-Schwinger perturbations are related to thermodynamic phase transitions of charged black holes. This has been criticised in~\\cite{Berti:2008xu}, the argument being that the relation between the QNMs and the phase transition had not been properly formulated. In this direction, the discontinuities observed in the heat capacity of charged Kaluza-Klein black holes with squashed horizons, were connected with the quasi-normal spectrum~\\cite{He:2008im}. Further evidence of a non-trivial relation between the thermodynamical and dynamical properties of black holes was provided in~\\cite{Shen:2007xk}.  In this work we calculate both analytically and numerically the QNMs of scalar perturbations of topological-AdS black holes in $d=4,5$ and $6$ dimensions. In all dimensionalities considered, we find a discontinuity of the QNMs spectrum at the critical point. This provides further evidence of a second order phase transition at the critical point observed in~\\cite{Koutsoumbas:2006xj,Koutsoumbas:2008pw}. In our numerical calculations we used the method developed in~\\cite{HH}. We conjecture that this method is problematic for large $n$ because it breaks down, and a regularization scheme is proposed.  Our work is organized  as follows: in section 2 the analytical calculation is presented, section 3 contains comments on the numerical method used, and in section 4 the numerical results may be found. Finally the conclusions are presented in section 5. ",
        "conclusions": "\\label{sec5}  We have studied the perturbative behaviour of the topological-AdS black holes. We have calculated both analytically and numerically the QNMs of scalar perturbations of these black holes.  Analytical calculations show that for small black holes  at any dimension there is a critical point (at $r_+=1$) below which the real part decreases with $n$, having a positive slope, whereas above the critical point the oscillatory modes increase with a negative slope. We also found that below the critical point there is a critical value of $\\xi$ below which there are purely decaying modes while above the critical point there are only oscillatory modes for any $\\xi$. In five dimensions the QNMs of scalar perturbations of TBH-AdS can be obtained explicitly from the Heun function which solves the five-dimensional wave equation.  These results are also supported by  numerical investigations of the QNMs. The numerical results show clearly a change of slope of QNMs around a critical temperature for dimensions $d=4,5,6.$  For larger dimensions the root finding algorithm is difficult to be implemented. This is connected with the observation that for higher dimensional theories the regularization needed for the convergence of the series becomes less efficient."
    },
    "0806/0806.3457_arXiv.txt": {
        "abstract": "We present an integrated study of star formation and galactic stellar mass assembly from $z=0.05-1.5$ and galactic metallicity evolution from $z=0.05-0.9$ using a very large and highly spectroscopically complete sample selected by rest-frame NIR bolometric flux. Our NIR (rest-frame $0.8-2.4~\\mu$m) sample consists of 2634 galaxies with fluxes in excess of $2\\times 10^{-15}$~ergs~cm$^{-2}$~s$^{-1}$ in the GOODS-N field. It probes to a complete mass limit of $10^{10}$~M$_\\odot$ for $z=0.05-0.9$ and includes all Milky Way mass galaxies for $z=0.05-1.5$. We have spectroscopic redshifts and high-quality spectra from $4500-10000$~\\AA\\ for 2020 (77\\%) of the galaxies. Our 13-band photometric redshift estimates show that most of the spectroscopically unidentified sources in the above redshift ranges are early-type galaxies. We assume a Salpeter IMF and fit Bruzual \\& Charlot (2003) models to the data to compute the galactic stellar masses and extinctions. We calibrate the star formation diagnostics internally using our $z=0.05-0.475$ sample. We then derive the galactic stellar mass assembly and star formation histories. We compare our extinction corrected UV-based star formation rate densities with the combination of the star formation rate densities that we compute from the 24~$\\mu$m fluxes and the extinction uncorrected \\oii\\ luminosities. We determine the expected formed stellar mass density growth rates produced by star formation and compare them with the growth rates measured from the formed stellar mass functions by mass interval. We show that the growth rates match if the IMF is slightly increased from the Salpeter IMF at intermediate masses ($\\sim 10$~M$_\\odot$). We investigate the evolution of galaxy color, spectral type, and morphology with mass and redshift and the evolution of mass with environment. We find that applying extinction corrections is critical when analyzing the galaxy colors. As an example, prior to correcting for extinction, nearly all of the galaxies in the green valley are 24~$\\mu$m sources, but after correcting for extinction, the bulk of the 24~$\\mu$m sources lie in the blue cloud. We also compute the metallicities of the sources between $z=0.05-0.9$ that have well-detected \\hb, \\oii~$\\lambda3727$, and \\oiii~$\\lambda5007$ emission lines using the R23 diagnostic ratio. At $z<0.475$ we use the R23, \\nii/\\oii, and \\nii/\\ha\\ diagnostic ratios. We find an evolution of the metallicity-mass relation corresponding to a decrease of $0.21\\pm0.03$~dex between the local value and the value at $z=0.77$ in the $10^{10}-10^{11}$~M$_\\odot$ range. We use the metallicity evolution to estimate the gas mass of the galaxies, which we compare with the galactic stellar mass assembly and star formation histories. Overall, our measurements are consistent with a galaxy evolution process dominated by episodic bursts of star formation and where star formation in the most massive galaxies ($\\gtrsim 10^{11}$~M$_\\odot$) ceases at $z<1.5$ because of gas starvation. ",
        "introduction": "\\label{secintro}  One of the fundamental goals of modern cosmology is to understand the formation and evolution of the galaxy population as a whole. We shall refer to this as the cosmic galaxy formation problem. There has been spectacular progress in addressing the cosmic galaxy formation problem over the last twenty years, beginning with the determination of the star formation history (e.g., Cowie et al.\\ 1995; Lilly et al.\\ 1996; Madau et al.\\ 1996; Steidel et al.\\ 1999; Haarsma et al.\\ 2000; Barger et al.\\ 2000; Le Floc'h et al.\\ 2005; P{\\'e}rez-Gonz{\\'a}lez et al.\\ 2005; Hopkins \\& Beacom 2006; Wang et al.\\ 2006; Reddy et al.\\ 2008). This has been followed more recently by efforts to measure the galactic stellar mass assembly history (e.g., Brinchmann \\& Ellis 2000; Cole et al.\\ 2001; Bell et al.\\ 2003, 2007; P{\\'e}rez-Gonz{\\'a}lez et al.\\ 2003, 2008; Dickinson et al.\\ 2003; Rudnick et al.\\ 2003, 2006; Fontana et al.\\ 2003, 2004, 2006; Drory et al.\\ 2004, 2005; Bundy et al.\\ 2005, 2006; Conselice et al.\\ 2005, 2007; Borch et al.\\ 2006; Pannella et al.\\ 2006; Elsner et al.\\ 2008) and the evolution of metallicity with galaxy mass and redshift (e.g., Kobulnicky et al.\\ 2003; Lilly et al.\\ 2003; Kobulnicky \\& Kewley 2004; Tremonti et al.\\ 2004; Liang et al.\\ 2004; Savaglio et al.\\ 2005). However, ideally what one wants is a comprehensive analysis of the history of star formation, the growth of galactic stellar mass and metals content, and the changes in morphology with redshift, galaxy mass, and the environment for a large, mass-selected galaxy sample that could be compared in detail with local galaxy properties and cosmological simulations of galaxy evolution. In particular, such an analysis could yield clear explanations for the migration of star formation to lower mass galaxies at later cosmic times and the simultaneous quenching of star formation in the most massive galaxies (the downsizing of Cowie et al.\\ 1996), as well as for the color bimodality of galaxy populations (e.g., Strateva et al.\\ 2001; Baldry et al.\\ 2004).  Up until now such an analysis has not been possible since existing data sets are either visually selected, have limited color information, and are poorly suited to a metals analysis because of the spectroscopic wavelength coverage (e.g., the DEEP2 survey); mass selected but based on photometric redshifts (e.g., Combo17/GEMS); or mass selected and spectroscopically observed but based on a relatively small sample (e.g., the Gemini Deep Deep Survey).  In this paper we present, for the first time, an integrated, mass-based analysis made possible by the availability of a large, homogeneous, near-infrared (NIR) selected and spectroscopically observed galaxy sample in the Great Observatories Origins Deep Survey-North (GOODS-N; Giavalisco et al.\\ 2004) field. We have obtained extremely deep, wide-field NIR images (Keenan et al.\\ 2008, in preparation) and highly complete spectroscopic identifications of the sources in this field (Barger et al.\\ 2008, in preparation). We are therefore able to use, for the most part, spectroscopic redshifts to make our determinations of the galactic stellar mass assembly and star formation histories, as well as high-quality measurements of line fluxes to obtain the metallicity history.  However, we caution that even with such an excellent data set there are many complicating factors in relating the star formation history to the stellar mass assembly history and the formation of metals in galaxies, even at late cosmic times. (Here we shall take late cosmic times to be $z<1.5$.) At the conceptual level, methods of measuring star formation rates use diagnostics which are sensitive to the high-mass end of the stellar initial mass function (IMF), while stellar mass measurements are dominated by lower mass stars. Therefore, while the shape of the sub-solar IMF only enters as a normalization factor, the shape of the IMF at higher masses is critical in relating the star formation rates to the stellar masses. Thus, we must be concerned about the uncertainties in the IMF shape and the potential variations in the IMF shape between different types of galaxies. In principle we could minimize this problem by considering the growth of the stellar mass in metals rather than the growth of the total stellar mass, since the metals are produced by the same high-mass stars that are measured by the star formation diagnostics (Cowie 1988). However, even this is subject to uncertainties in the yields and would require the measurement of not only the total stellar mass evolution but also the metals evolution in both stars and gas, which would be very challenging to do.  Measurements of the star formation rates, stellar masses, and metals are also complicated by other factors. Extinction reradiates light from the rest-frame UV to the far-infrared (FIR), and we must determine total star formation rates over a wide range of galaxies with radically different morphologies and dust column densities. Conversions even of NIR light to stellar mass are complicated by ongoing active star formation, and there are still major uncertainties in the stellar modeling of the galaxy populations. Finally, determinations of the metals throughout the redshift range of interest can only be made for the gaseous baryons in the star-forming galaxies and depend on the notoriously uncertain conversions of the strong oxygen and nitrogen emission lines to metallicities.  Cosmic variance is also a significant issue in a field size as small as the GOODS-N (e.g. Somerville et al. 2004) and can affect our analysis of the evolution of quantities such as the galaxy mass density and the universal star formation rates.  These problems must be borne in mind throughout any work of the present type, and we attempt at all points to work forward as self-consistently as possible from the raw information (NIR luminosities, galaxy line strengths, raw star formation diagnostics, etc.) to inferences about the evolution of derived quantities, such as stellar masses, star formation rates, and metallicities. Also, wherever possible, we have used multiple independent methods to determine the sensitivity of the derived quantities to our underlying assumptions. We attempt to self consistently estimate the effects of cosmic variance within the data set and also to estimate the effects which analytic error estimates of the variance could introduce in our analysis. Finally we compare our results throughout to other recent work using different, and in some cases much larger, fields to check for consistency in these portions of the paper.  The outline of the paper is as follows. In \\S\\ref{secsample} and \\S\\ref{seclum} we describe the basic data and the sample selections. In \\S\\ref{secfit} we fit Bruzual \\& Charlot (2003) models to the data to determine galactic stellar masses and extinctions in the galaxies. In \\S\\ref{secspectral} we measure equivalent widths and line fluxes from the spectra. In \\S\\ref{secrelext} we compare measurements of the continuum and line extinctions. In \\S\\ref{secsfr} we derive self-consistent calibrations of the various star formation rate diagnostics. In \\S\\ref{secha} and \\S\\ref{seco3} we derive the metallicities with mass and redshift using various metallicity diagnostics. In \\S\\ref{secabs} we consider the galaxies missing from the metals analysis. This is a long paper, and some readers may wish to skip much of the detail and move to the discussion (\\S\\ref{secdisc}) and summary (\\S\\ref{seccon}), which we have tried to make separately readable and which contain the high-level interpretation of the data, including the derivation of the stellar mass assembly history with redshift, the evolution of the mass-metallicity and mass-morphology relations with redshift and environment, and the use of the metals evolution to derive an estimate of the baryonic gas mass reservoir in the galaxies. We find that all of our measurements provide a broad, self-consistent picture of a galaxy evolution process dominated by episodic bursts of star formation and where star formation in the most massive galaxies is terminated at later cosmic times as a consequence of gas starvation.  We adopt the $-1.35$ power-law Salpeter IMF (Salpeter 1955) extending from 0.1 to 100~M$_\\odot$ for ease of comparison with previous results. Most importantly, this allows us to compare directly with the local mass function computed by Cole et al.\\ (2001; hereafter, Cole01) for this IMF. The Salpeter IMF only differs significantly from the current best IMFs (Kroupa 2001; Chabrier 2003) below 1~M$_\\odot$, and thus these three IMFs differ only in the normalization of the galactic stellar mass and star formation rate determinations. We can convert the total mass formed into stars prior to stellar mass loss (which we will refer to as the formed stellar mass to distinguish it from the present stellar mass, which is the stellar mass present at any given time) from the Salpeter IMF to the Chabrier IMF by dividing by 1.39 and to the Kroupa IMF by dividing by 1.31. The exact conversion when considering the present stellar masses rather than the formed stellar masses depends on the average evolutionary stage of the galaxies. However, this dependence is relatively weak, and we may approximately convert the present stellar mass from the Salpeter IMF to the Chabrier IMF by dividing by 1.70 and to the Kroupa IMF by dividing by 1.54. Note that these latter conversion factors have been computed for the distribution of ages in our ensemble of galaxies. The present stellar mass for the Salpeter IMF is roughly 0.74 of the formed stellar mass.  We assume $\\Omega_M=0.3$, $\\Omega_\\Lambda=0.7$, and $H_0=70$~km~s$^{-1}$~Mpc$^{-1}$ throughout. All magnitudes are given in the AB magnitude system, where an AB magnitude is defined by $m_{AB}=-2.5\\log f_\\nu - 48.60$. Here $f_\\nu$ is the flux of the source in units of ergs~cm$^{-2}$~s$^{-1}$~Hz$^{-1}$. We assume a reference value of the solar metallicity of $12+\\log({\\rm O/H})=8.66$ and a conversion to the mass fraction of metals of $Z=0.0126$ (Asplund et al.\\ 2004). This conversion is weakly dependent on the assumed chemical composition relative to the oxygen abundance.   \\begin{inlinefigure} \\figurenum{1} \\centerline{\\psfig{figure=f1.ps,angle=90,width=3.5in}} \\vskip -0.2cm \\figcaption[]{ The observed area in the GOODS-N. The area is centered on RA(2000) and Dec(2000) coordinates (189.2282, 62.2375) with corners at (189.5435, 62.2749), (188.9137, 62.2000), (189.3090, 62.3824), and (189.1482, 62.0909). The covered area is 145~arcmin$^2$ ($9\\farcm3$ by $15\\farcm7$). The NIR-selected sample is shown with black dots, the 663 24~$\\mu$m detected sources with red squares, the 229 X-ray detected sources with blue diamonds, and the 97 20~cm detected sources with green open triangles. The concentration of the X-ray sources to the field center reflects the variation in the sensitivity of the X-ray image over the field. \\label{acs_nirbol_sample} } \\end{inlinefigure} ",
        "conclusions": "\\label{secdisc}  We now bring together the results from the previous sections to form an integrated picture of galactic stellar mass assembly, star formation, and metallicity evolution. In \\S\\ref{secgsmf} and \\S\\ref{secmassassembly} we measure the galactic stellar mass functions and the growth of the stellar mass densities with redshift. In \\S\\ref{secsfh} we measure the star formation rate densities as a function of redshift and galaxy mass. In \\S\\ref{secgrowth} we compute the expected formed stellar mass density growth rates produced by star formation and show that they match those measured from the formed stellar mass functions if the IMF is slightly increased from the Salpeter IMF at intermediate masses ($\\sim10~$M$_\\odot$). In \\S\\ref{secstarspec} we determine the instantaneous specific star formation rates, which give a quantitative description of the range of behaviors in the galaxies. We show that only galaxies below about $10^{11}~$M$_\\odot$ are growing substantially over $z=0.05-1.5$. In \\S\\ref{seccolor} we analyze the distributions of galaxy colors, equivalent widths, and 4000~\\AA\\ break strengths and find that star formation in all but the lowest mass galaxies in our sample is occurring in bursts with characteristic time intervals of about $4\\times10^9$~yr. We also find that most of the growth in the mass density is in the red sequence galaxies, whether these are chosen by color or from the rest-frame EW(\\oii). In \\S\\ref{secmassmorph} we show that as the redshift decreases, the galaxy types smoothly migrate from spirals to S0s and elliptical galaxies at all masses. The mass build-up is primarily in the E/S0 galaxies, which are also the red sequence galaxies, and there is little change in the galactic stellar mass function of the spirals. In \\S\\ref{secenv} we find that although the masses are a strongly increasing function of environment, there is little redshift dependence in this relation. Unlike local results, the fraction of galaxies in the red sequence shows little environmental dependence and appears to depend primarily on the galaxy mass. Finally, in \\S\\ref{secmetev} we compare the metallicity evolution in the present sample with local and high-redshift metallicity measurements. In \\S\\ref{secgasmass} we combine the measured increases in metallicity with redshift with the metals returned from the measured SFRs to make a crude estimate of the galaxy gas mass reservoir, and we compare this to the stellar mass density as a function of galaxy mass.   \\subsection{Galactic Stellar Mass Functions} \\label{secgsmf}  In Figure~\\ref{fig_mfun} we show the galactic stellar mass functions {\\em (black squares)\\/} in three redshift intervals [(a) $z=0.9-1.5$, (b) $z=0.475-0.9$, and (c) $z=0.05-0.475$] computed using the $1/V$ methodology (Felten 1976) and compared with the local mass function of Cole01 adjusted to the present cosmology {\\em (purple crosses)\\/}. Both the Cole01 and our own mass functions are computed with the Salpeter IMF assumed throughout.  For each mass function we have computed the errors in two ways. First, we assigned $1\\sigma$ errors based on the number of sources in each bin {\\em (red error bars)\\/}. These errors dominate at the high-mass end where there are small numbers of sources in each bin. Second, we estimated the effect of cosmic variance in a simple empirical way. In each redshift interval we eliminated the strongest redshift sheet from the sample. For example, in the $z=0.475-0.9$ redshift interval we removed all of the sources lying between $z=0.845$ and $z=0.86$ (see Figure~\\ref{mass_byz}). Typically the strongest single sheet will contain about 10\\%$-20$\\% of the galaxies in the redshift interval. We then recomputed the mass function corresponding to the new volume and used the difference as an error estimate ({\\em black error bars\\/}). These error estimates dominate at the low-mass end. The internal error estimate from this method are similar in size to the analytic estimates of Somerville et al. 2004 which would give systematic error of 0.25, 0.15 and 0.1 dex in the $z=0.05-0.475$, $z=0.475-0.9$ and $z=0.9-1.5$ redshift intervals.   \\begin{inlinefigure} \\figurenum{37} \\centerline{\\psfig{figure=f37a.ps,angle=90,width=3.5in}} \\vskip -0.6cm \\centerline{\\psfig{figure=f37b.ps,angle=90,width=3.5in}} \\vskip -0.6cm \\centerline{\\psfig{figure=f37c.ps,angle=90,width=3.5in}} \\vskip -0.2cm \\figcaption[]{ Galactic stellar mass functions in the redshift intervals (a) $z=0.9-1.5$, (b) $z=0.475-0.9$, and (c) $z=0.05-0.475$ {\\em (black squares)\\/}. The red error bars are $1\\sigma$, based on the number of sources in each bin, while the black error bars are derived from excluding the strongest velocity sheet in each redshift interval (see text for details). The purple crosses and associated $1\\sigma$ error bars are the same in all three redshift intervals and show the local mass function of Cole01 for the Salpeter IMF adjusted to the present cosmology. The cyan curves show the best-fit Schechter functions obtained using the Sandage et al.\\ (1979) maximum likelihood method. \\label{fig_mfun} } \\end{inlinefigure}   In order to provide parametric fits to the data, we have assumed a Schechter (1976) form, \\begin{equation} \\phi(M) = \\phi_{\\star}\\Biggl({M\\over {M_{\\star}(z)}}\\Biggr)^{\\alpha(z)} {e^{-M/{M_{\\star}(z)}}\\over M_{\\star}(z)} \\,, \\label{eqnshec} \\end{equation} for the mass function, where $\\phi(M)$ is the number of galaxies per unit mass per Mpc$^{3}$ at mass $M$. We used the Sandage et al.\\ (1979) maximum likelihood method to determine the power-law index $\\alpha(z)$ and the characteristic mass $M_\\star(z)$ for each of the three redshift intervals. The best-fit functions are shown in Figure~\\ref{fig_mfun} {\\em (cyan curves)\\/}, and the derived $\\alpha$ and $M_\\star$ for each redshift interval are shown in Figure~\\ref{mlf_all} together with the 68\\% and 95\\% confidence error ellipses. We obtained the normalizations $\\phi_\\star(z)$ by normalizing to the number of objects in each redshift interval. We used our variance estimates (which dominate the error budget) in computing the error on this quantity. We summarize the maximum likelihood fits in each redshift interval in Table~\\ref{sty_fits}, together with the 68\\% confidence ranges.   \\begin{inlinefigure} \\figurenum{38} \\centerline{\\psfig{figure=f38.ps,angle=90,width=3.5in}} \\vskip -0.2cm \\figcaption[]{ Sandage et al.\\ (1979) fits to $\\alpha(z)$ and $M_\\star(z)$ for the three redshift intervals $z=0.05-0.475$ {\\em (red diamond)\\/}, $z=0.475-0.9$ {\\em (black square)\\/}, and $z=0.9-1.5$ {\\em (blue triangle)\\/}. In each case the inner contour shows the 68\\% confidence range and the outer contour shows the 95\\% confidence range for the symbol of the same color. \\label{mlf_all} } \\end{inlinefigure}   A number of analyses of galactic stellar mass functions have recently been undertaken (see references in \\S\\ref{secintro}), and our results are broadly consistent. We may make the most straightforward comparison with Fontana et al.\\ (2006; hereafter Fontana06), who analyze their data in a very similar way and who also use the Salpeter IMF. While their GOODS-MUSIC sample relies heavily on photometric redshifts, in other respects it is comparable to the present sample. Their Figure~4 bears a striking resemblence to our Figure~\\ref{fig_mfun}. Conselice07 use a much larger field size with almost twenty times as many objects and also draw a similar conclusion. The agreement between fields is reassuring that our results are robust and are not being dominated by cosmic variance.  The effects of downsizing can be clearly seen in both Fontana06's Figure~4 and our Figure~\\ref{fig_mfun}: the galaxy number densities at the low-mass end are still rising down to the lowest redshift interval, while the number densities at the high-mass end ($>10^{11}$~M$_\\odot$) are changing much more slowly over $z=0.05-1.5$.  Over this redshift range Fontana06 find a relatively invariant $\\alpha(z)\\sim -1.2$, but their $M_\\star(z)$ rises by about 0.12~dex from $z=0$, where $\\log M_\\star =11.16$ from Cole01, to $z=1.35$. Our $\\alpha(z)$ values are slightly greater than this, though consistent with a constant $\\alpha(z)=-1.2$ within the $2\\sigma$ errors (see Fig.~\\ref{mlf_all}). Our slightly greater $\\alpha(z)$ values are a consequence of Fontana06 having a turn-up at the low-mass end (see their Fig.~4). It is quite likely that this is a result of their use of photometric redshifts, which are more problematic for low-mass galaxies because such galaxies are predominantly blue and harder to fit, but it could also be a measure of the cosmic variance.  However, we do see a significant evolution in the mean mass (a rise in either $M_\\star$ or $\\alpha$ with increasing redshift) in our data. In particular, the fitted values in the highest redshift interval are not consistent at the $3\\sigma$ level with those in the lowest redshift interval (see Fig.~\\ref{mlf_all}). We can see this evolution most clearly by adopting a fixed $\\alpha(z)=-1.18$ from the local Cole01 fit and computing $M_\\star(z)$ for this fixed slope. This is given as the final column in Table~\\ref{sty_fits}, and it shows a rise of about 0.12~dex between $z=0$ (Cole01) and our highest redshift interval, an identical result to that of Fontana06.   \\begin{inlinefigure} \\figurenum{39} \\centerline{\\psfig{figure=f39.ps,angle=90,width=3.5in}} \\vskip -0.2cm \\figcaption[]{ Stellar mass density per logarithmic galaxy mass interval in the redshift range $z=0.8-1.2$ {\\em (black squares)\\/} compared with the local distribution from Cole01 {\\em (purple crosses)\\/}. The error bars are as in Figure~\\ref{fig_mfun}. \\label{mass_den_plot} } \\end{inlinefigure}   The change in $M_\\star(z)$ is a numerical consequence of the build-up of the low-mass region of the galactic stellar mass function relative to the high-mass region. In Figure~\\ref{mass_den_plot} we compare the stellar mass density per logarithmic galaxy mass interval in the redshift range $z=0.8-1.2$ {\\em (black squares)\\/} with that from Cole01 locally {\\em (purple crosses)\\/}. Both distribution functions are strongly peaked with most of the mass density lying in galaxies with masses close to the peak value. However, the Cole01 function is broader and extends to lower masses.  The peak of the mass density per logarithmic galaxy mass, $M_\\star(z) (2\\alpha(z)+1)$, provides an alternative way to characterize the Schechter function (see Baldry et al.\\ 2006). The evolution of the peak is better defined than the evolution of either $\\alpha(z)$ and $M_{\\star}(z)$ separately. The location of the peak increases from $11.07\\pm0.01$~M$_\\odot$ in Cole01 and 10.88 ($10.77-11.09$)~M$_\\odot$ in our lowest redshift interval to $11.18\\pm0.03$~M$_\\odot$ in our highest redshift interval.   \\subsection{Stellar Mass Density Growth with Redshift} \\label{secmassassembly}  In Figure~\\ref{mden_byz} we plot the stellar mass density versus redshift for various mass intervals. In both panels we use black solid squares and 68\\% confidence limits to show the mass density evolution for all of the galaxies in our sample above $10^{10.5}$~M$_\\odot$. We also obtained the local mass density above $10^{10.5}$~M$_\\odot$ by integrating the Cole01 function. We denote this by a black open square, which we show extended to all redshifts {\\em (black dashed line)\\/} for ease of comparison. $10^{10.5}$~M$_\\odot$ is the lowest mass to which we can measure the mass density over our entire $z=0.05-1.5$ redshift range.   \\begin{inlinefigure} \\figurenum{40} \\centerline{\\psfig{figure=f40a.ps,angle=90,width=3.5in}} \\vskip -0.6cm \\centerline{\\psfig{figure=f40b.ps,angle=90,width=3.5in}} \\vskip -0.2cm \\figcaption[]{ Universal stellar mass density vs. redshift. In both panels the black solid squares show the evolution of the mass density for all sources above $10^{10.5}$~M$_\\odot$, and the black open square (and black dashed line) shows the local mass density above $10^{10.5}$~M$_\\odot$ obtained by integrating the Cole01 function. (a) The evolution in the mass intervals $10^{11.5}-10^{12}$~M$_\\odot$ {\\em (blue triangles, highest-mass interval)\\/} and $10^{11}-10^{11.5}$~M$_\\odot$ {\\em (red diamonds, high-mass interval)\\/}. (b) The evolution in the mass intervals $10^{10.5}-10^{11}$~M$_\\odot$ {\\em (cyan triangles, intermediate-mass interval)\\/} and $10^{10}-10^{10.5}$~M$_\\odot$ {\\em (purple diamonds, low-mass interval)\\/}. The low-mass interval is only shown at $z<0.9$ where it is complete. Errors are the maximum of the numerical and variance errors discussed in the text. The least-square polynomial fits of $\\log \\rho$ vs. $\\log(1+z)$ are shown with the solid curves. The local mass densities for each of the mass intervals obtained by integrating the Cole01 function {\\em (open symbols, color corresponds to the mass interval)\\/} are also shown extended to all redshifts {\\em (dashed lines)\\/}. \\label{mden_byz} } \\end{inlinefigure}   In Figure~\\ref{mden_byz}a we show the mass density evolution for two high-mass intervals: $10^{11.5}-10^{12}$~M$_\\odot$ {\\em (blue triangles; hereafter, our highest-mass interval)\\/} and $10^{11}-10^{11.5}$~M$_\\odot$ {\\em (red diamonds; our high-mass interval)\\/}. In Figure~\\ref{mden_byz}b we show the mass density evolution for two lower mass intervals: $10^{10.5}-10^{11}$~M$_\\odot$ {\\em (cyan triangles; our intermediate-mass interval)\\/} and $10^{10}-10^{10.5}$~M$_\\odot$ {\\em (purple diamonds; our low-mass interval)\\/}. We show the evolution in the low-mass interval only to $z=0.9$ where the sample is complete. We obtained the local mass density for each of the mass intervals by integrating the Cole01 function over those intervals {\\em (open symbols)\\/}. For ease of comparison, we show those local mass densities extended to all redshifts {\\em (dashed lines)\\/}.  Locally almost equal mass densities lie in the high {\\em (open red diamond in Fig.~\\ref{mden_byz}a)\\/} and intermediate {\\em (open cyan triangle in Fig.~\\ref{mden_byz}b)\\/} mass intervals. However, the high-mass interval material is changing more slowly than the intermediate-mass interval material, which is growing smoothly over the whole $z=0.05-1.5$ range. The low-mass interval material is also growing smoothly over $z=0.05-0.9$. The highest-mass interval material has a slow evolution (and indeed is consistent with having no evolution) over the entire $z=0.05-1.5$ range, though the error bars are large and the amount of material contained in this mass interval is small.  The mass-sliced data may be most easily compared with Conselice07, who present their data in a similar way. While the Conselice07 sample is a large one (almost 50,000 objects), the data are not of the same quality as that in Fontana06 or in the present paper. Only about 22\\% of their sample have spectroscopic redshifts from the DEEP2 observations, and their photometric redshifts were calculated using only 4 ($BRIK$; half the sample) or 5 ($BRIJK$) band colors. These are too few colors to get reliable measurements of all of the desired quantities (i.e., extinction, photometric redshifts, mass, age/metallicity, and evolutionary model). For comparison, recall that over the redshift range $z=0.05-1.5$ we have spectroscopic redshifts for 84\\% of the galaxies in our uniform NIR flux-limited sample (see \\S\\ref{secspec}). Moreover, for the very small fraction of sources where we used photometric redshifts, they were derived from 13  band colors ($UBVRIz'JHK_s$ and the four {\\em Spitzer\\/} IRAC bands). Fontana06 used 14 band colors to determine their predominantly photometric redshifts.  Despite these limitations, the agreement is good. We can compare our Figure~\\ref{mden_byz} with Figure~4 of Conselice07 if we convert our stellar masses and stellar mass densities to the Chabrier (2003) IMF used by Conselice07 by dividing by 1.70, as discussed in \\S\\ref{secintro}. Our results extend to lower masses than Conselice07 since our NIR sample is deeper. However, if we compare the region of mass and redshift where the two studies overlap, then we find good agreement between the two results in both absolute value and shape. Both analyses show slow growth at the higher masses and more rapid growth at the lower masses. The one quantitative difference is that we do not reproduce the sharp drop in the mass density which Conselice07 see in their $10^{11}-10^{11.5}$~M$_\\odot$ interval above $z=1.2$. It is not clear whether this is simply a result of cosmic variance (neither of the surveys are large enough for cosmic variance not to matter), or whether it is related to Conselice07's use of 4 or 5 band photometric redshifts. Since the present work is based on a much deeper and far more spectroscopically complete sample than that of Conselice07, our results should be more reliable if this is the explanation.  In order to quantify the results of Figure~\\ref{mden_byz}, we made least-square polynomial fits to the logarithmic stellar mass densities [including the local values obtained by integrating the Cole01 function] versus the logarithmic cosmic time. We show these in Figure~\\ref{mden_byz} with solid curves. For the total mass above $10^{10.5}$~M$_\\odot$ we find \\begin{equation} \\log {\\rm \\rho_\\ast(>10.5~{\\rm M_\\odot})} = 8.56 + (0.73\\pm0.16) \\log(t/t_{0})  \\,, \\label{massint1} \\end{equation} and for the mass intervals we find \\begin{equation} \\log {\\rm \\rho_\\ast(10^{10}-10^{10.5}~{\\rm M_\\odot})} = 7.89 + (1.11\\pm0.28) \\log(t/t_{0})  \\,, \\label{massint2} \\end{equation} \\begin{equation} \\log {\\rm \\rho_\\ast(10^{10.5}-10^{11}~{\\rm M_\\odot})} = 8.20 + (0.95\\pm0.10) \\log(t/t_{0})  \\,, \\label{massint3} \\end{equation} \\begin{equation} \\log {\\rm \\rho_\\ast(10^{11}-10^{11.5}~{\\rm M_\\odot})} = 8.22 + (0.65\\pm0.12) \\log(t/t_{0})  \\,, \\label{massint4} \\end{equation} \\begin{equation} \\log {\\rm \\rho_\\ast(10^{11.5}-10^{12}~{\\rm M_\\odot})} = 7.57 + (0.59\\pm0.37) \\log(t/t_{0})  \\,. \\label{massint5} \\end{equation} The fit for Equation~\\ref{massint2} is over the $z=0.05-0.9$ redshift range, while the remaining fits are over the $z=0.05-1.5$ range. The low-mass ranges ($<10^{11}$~M$_\\odot$) are growing approximately linearly with time, while the high-mass ranges are growing more slowly. Thus, the low-mass galaxies have roughly constant stellar mass density growth rates, and it is the drop in the growth rates in the high-mass galaxies that is responsible for the overall drop in the growth rates seen in Equation~\\ref{massint1}.  We also computed the total stellar mass density evolution by integrating to the limiting mass in each redshift interval and extrapolating to estimate the contribution from lower mass galaxies. This correction is not large since the total mass density is dominated by galaxies near $M_\\star(z)$. Locally about 72\\% of the mass density lies in galaxies above $10^{10.5}$~M$_\\odot$ and about 89\\% in galaxies above $10^{10}$~M$_\\odot$. If we use our best-fit Schechter function, then in the $z=0.9-1.5$ interval about 88\\% of the mass density lies in galaxies above $10^{10.5}$~M$_\\odot$ and about 97\\% in galaxies above $10^{10}$~M$_\\odot$. If we instead force-fit to $\\alpha=-1.18$, then the percentages are 77\\% and 91\\%, respectively. When we include these corrections, the total stellar mass density has a slightly steeper dependence on redshift than does the mass density above $10^{10.5}$~M$_\\odot$. Using, respectively, our best-fit Schechter function and the fit where $\\alpha$ is forced to $-1.18$, we find \\begin{equation} \\log {\\rm \\rho_\\ast(total)} = 8.77 + (0.91\\pm0.15) \\log(t/t_{0})  \\,, \\label{massinttot1} \\end{equation} and \\begin{equation} \\log {\\rm \\rho_\\ast(total)} = 8.77 + (0.80\\pm0.15) \\log(t/t_{0})  \\,. \\label{massinttot2} \\end{equation}   \\subsection{Star Formation History} \\label{secsfh}  We tested our empirical SFR calibrations (\\S\\ref{seccal}) by calculating the universal SFH from $z=0.05-1$ for all of the galaxies in our NIR sample. We show this in Figure~\\ref{ha_rho_star_bymass}, where we use red open (solid) squares to denote the star formation rate densities (SFRDs) calculated from the \\hb\\ (UV) luminosity densities after applying our extinction corrections. As is well known, the extinction corrections are substantial, typically factors of five in the various redshift bins. At $z=0.9$ about 80\\% of the light is dust reradiated. Our extinction corrected values at $z>0.3$ agree broadly with the many measurements in the literature (see the references in \\S\\ref{secintro}). We explicitly compare with the values derived from radio and submillimeter data from the stacking analysis of Wang et al.\\ (2006) {\\em (black squares)\\/}. These should be a good measure of the total star formation history, including highly-obscured sources. The good agreement suggests that we are seeing most of the higher redshift star formation in our present sample.   \\begin{inlinefigure} \\figurenum{41} \\centerline{\\psfig{figure=f41.ps,angle=90,width=3.5in}} \\vskip -0.2cm \\figcaption[]{ Star formation histories calculated from the rest-frame \\hb\\ {\\em (open)\\/} and UV {\\em (solid)\\/} luminosity densities of the NIR sample using the calibrations of \\S\\ref{seccal}. The red symbols show the extinction corrected SFRDs. The purple symbols show the extinction uncorrected SFRDs. The black solid squares were derived by Wang et al.\\ (2006) from radio and submillimeter data and agree well with our extinction corrected values. The blue symbols show the SFRDs directly seen at rest-frame UV wavelengths for UV selected galaxies. We only show the {\\em GALEX\\/} determinations of Wyder et al.\\ (2005; local) and Schiminovich et al.\\ (2005) {\\em (downward pointing triangles)\\/} and the ground-based determinations of Wilson et al.\\ (2002) {\\em (upward pointing triangles)\\/}, since these are the most accurate measurements near $z=1$. All agree reasonably well with our extinction uncorrected values. The formal errors are mostly smaller than the symbol sizes. \\label{ha_rho_star_bymass} } \\end{inlinefigure}   However, the SFRDs from the NIR sample drop steeply at lower redshifts. At these lower redshifts most of the star formation is instead seen as direct UV emission from lower mass galaxies, which have very small extinctions. Thus, in Figure~\\ref{ha_rho_star_bymass} we also show the SFRDs derived from our extinction uncorrected data. Here we use purple open (solid) diamonds to denote the SFRDs calculated from the \\hb\\ (UV) luminosity densities. Since the extinction correction is smaller at \\hb, the \\hb\\ points lie above the UV points. We can compare the UV data with the literature results for rest-frame UV flux measurements of UV selected galaxies {\\em (blue symbols)\\/}, such as the ground-based observations of Wilson et al.\\ (2000) {\\em (upward pointing triangles)\\/} and the more recent {\\em GALEX\\/} determinations of Wyder et al.\\ (2005; local) and Schiminovich et al.\\ (2005) {\\em (downward pointing triangles)}. We again find close agreement.  In Figure~\\ref{sfr_comp_bymass} we show our extinction corrected SFRDs {\\em (red squares)\\/} calculated from the rest-frame UV luminosity densities versus redshift for the mass intervals (a) $10^{10}-10^{10.5}$~M$_\\odot$, (b) $10^{10.5}-10^{11}$~M$_\\odot$, (c) $10^{11}-10^{11.5}$~M$_\\odot$, and (d) $10^{11.5}-10^{12}$~M$_\\odot$. As another check of our UV extinction corrections, we also show the SFRDs {\\em (blue diamonds)\\/} obtained by adding those computed from the $24~\\mu$m fluxes for the obscured star formation with those computed from the extinction uncorrected \\oii\\ luminosities for the unobscured star formation (see Conselice07). This method has no dependence on the UV extinction corrections and shows extremely similar results to the UV-based method.   \\begin{figure*} \\figurenum{42} \\centerline{\\psfig{figure=f42a.ps,angle=90,width=3.5in} \\psfig{figure=f42b.ps,angle=90,width=3.5in}} \\vskip -0.6cm \\centerline{\\psfig{figure=f42c.ps,angle=90,width=3.5in} \\psfig{figure=f42d.ps,angle=90,width=3.5in}} \\vskip -0.2cm \\figcaption[]{ Star formation rate densities from two different methods vs. redshift for the mass intervals (a) $10^{10}-10^{10.5}$~M$_\\odot$, (b) $10^{10.5}-10^{11}$~M$_\\odot$, (c) $10^{11}-10^{11.5}$~M$_\\odot$, and (d) $10^{11.5}-10^{12}$~M$_\\odot$. The red squares show the extinction corrected SFRDs calculated from the UV luminosity densities using our empirical calibrations. The blue diamonds show the completely independent calculation of the SFRDs from the 24~$\\mu$m fluxes and the extinction uncorrected \\oii\\ fluxes. \\label{sfr_comp_bymass} } \\end{figure*}   In Figure~\\ref{final_sfr_dist} we show our UV-based extinction corrected SFRDs per unit logarithmic mass versus the logarithmic mass in the redshift intervals $z=0.05-0.475$ {\\em (red diamonds)\\/}, $z=0.475-0.9$ {\\em (black squares)\\/}, and $z=0.9-1.5$ {\\em (blue triangles)\\/}. In each redshift interval we only show the SFRDs over the mass range where the NIR sample is complete. Just for a shape comparison, we also show on the figure the stellar mass density distribution function for $z=0.475-0.9$ divided by the age of the universe at $z=0.7$ {\\em (black curve)\\/}. The contrast between this and the SFRD distribution is striking, with the latter being much more strongly weighted to low-mass galaxies. At all redshifts there is very little star formation at masses $>10^{11.1}$~M$_\\odot$, but at high redshifts the star formation peaks in the interval $10^{10.5}-10^{11.1}$~M$_\\odot$. It is the drop in star formation for galaxies in this mass interval that results in the drop in the overall SFH. At lower masses there is relatively little change over $z=0.05-0.9$.   \\begin{inlinefigure} \\figurenum{43} \\centerline{\\psfig{figure=f43.ps,angle=90,width=3.5in}} \\vskip -0.2cm \\figcaption[]{ Extinction corrected star formation rate density per unit logarithmic mass vs. logarithmic mass in the redshift intervals $z=0.05-0.475$ {\\em (red diamonds)\\/}, $z=0.475-0.9$ {\\em (black squares)\\/}, and $z=0.9-1.5$ {\\em (blue triangles)\\/}. The solid black curve shows the corresponding shape of the stellar mass density distribution function for $z=0.475-0.9$ divided by the age of the universe at $z=0.7$. \\label{final_sfr_dist} } \\end{inlinefigure}   \\subsection{Comparison of Growth Rates} \\label{secgrowth}  In this section we want to compare the stellar mass density growth rates produced by star formation with those measured from the stellar mass functions. Other groups have made this comparison before (e.g., Borch et al.\\ 2006; P{\\'e}rez-Gonz{\\'a}lez et al.\\ 2008), but generally they have not done so by mass interval, nor have they done it self-consistently, since they rely on other groups' determinations of the star formation. Thus, they have to assume that the galaxies producing the star formation have masses in the same range as the galaxies in their stellar mass function analysis. A comparison by mass interval is a much more powerful way to analyze this type of data.  First, to obtain the growth rates from the stellar mass functions, we took the derivative of the least-square polynomial fits of Equations~\\ref{massint2}$-$\\ref{massint5} after multiplying the mass densities by the average correction of 1.35 given in \\S\\ref{secintro} to convert to formed stellar mass densities (i.e., the total mass density formed into stars prior to stellar mass loss; it is the formed mass density growth rates which are directly related to the SFRDs). We show these growth rates {\\em (purple solid lines)\\/} in Figure~\\ref{nuv_rho_star_bymass} versus redshift for the mass intervals (a) $10^{10}-10^{10.5}$~M$_\\odot$, (b) $10^{10.5}-10^{11}$~M$_\\odot$, (c) $10^{11}-10^{11.5}$~M$_\\odot$, and (d) $10^{11.5}-10^{12}$~M$_\\odot$.   \\begin{figure*} \\figurenum{44} \\centerline{\\psfig{figure=f44a.ps,angle=90,width=3.5in} \\psfig{figure=f44b.ps,angle=90,width=3.5in}} \\vskip -0.6cm \\centerline{\\psfig{figure=f44c.ps,angle=90,width=3.5in} \\psfig{figure=f44d.ps,angle=90,width=3.5in}} \\vskip -0.2cm \\figcaption[]{ Mass formation rate densities from two different methods vs. redshift for the mass intervals (a) $10^{10}-10^{10.5}$~M$_\\odot$, (b) $10^{10.5}-10^{11}$~M$_\\odot$, (c) $10^{11}-10^{11.5}$~M$_\\odot$, and (d) $10^{11.5}-10^{12}$~M$_\\odot$. The red squares show the formed stellar mass density growth rates that the extinction corrected SFRDs calculated from the UV luminosity densities using our empirical calibrations would produce in each mass interval, as calculated from Eq.~\\ref{eqnmasscon2}. The blue diamonds show the same but this time based on the SFRDs calculated from the 24~$\\mu$m and extincton uncorrected \\oii\\ fluxes. The purple solid lines show the formed stellar mass density growth rates as obtained from the derivative of the least-square polynomial fits to the stellar mass density history (Eqs.~\\ref{massint2}$-$\\ref{massint5}) after multiplying by a factor of 1.35 to convert the present stellar masses to formed stellar masses (see \\S\\ref{secintro}). The purple dashed lines show the range given by the $1\\sigma$ errors on the fits. The black dotted lines show the 68\\% range in the specific star formation rates from Noeske et al. 2007. \\label{nuv_rho_star_bymass} } \\end{figure*}   Next, to compute the expected formed stellar mass density growth rates produced by star formation, we need to allow for the fact that the mass of a galaxy will grow as it forms stars. Consequently, the mass that formed in one mass interval will eventually end up in another mass interval. We can describe this movement with the conservation equation \\begin{equation} \\dot{\\rho}(M) = s(M) - {d\\over dM}(\\dot{M}\\rho(M)) \\,, \\label{eqnmasscon} \\end{equation} where $\\rho(M)$ is the stellar mass density per unit mass interval and $s(M)$ is the SFRD per unit mass interval. Since $\\dot{M}$ is related to $s(M)$ through $\\dot{M}=s(M)/\\phi(M)$, where $\\phi(M)$ is the number of galaxies per unit mass interval, and $M$ is related to $\\rho(M)$ through $M=\\rho(M)/\\phi(M)$, we can rewrite the equation as \\begin{equation} \\dot{\\rho}(M) = -{ds(M)\\over d\\ln(M)} \\,. \\label{eqnmasscon2} \\end{equation} Integrated over the total stellar mass function, Equation~\\ref{eqnmasscon2} simply says that the rate of increase in the formed stellar mass density per unit time is equal to the SFRD.  Merging can also redistribute mass from low to high masses in the galactic mass function (e.g., Conselice07), but this is much harder to quantify. As we shall see below, Equation~\\ref{eqnmasscon2} provides a good description of the changes seen in the formed stellar mass function with mass and time without including substantial merging. Conselice07 reach a similar conclusion.  As a simple intuitive example of the meaning of Equation~\\ref{eqnmasscon2}, we may consider the case where the specific star formation rate is constant (i.e., $s\\propto \\rho$). In this case the mass distribution function per unit logarithmic mass, as shown in Figure~\\ref{mass_den_plot}, simply moves to the right in the x-axis. Then the mass density in a given interval grows for high masses above the peak in the mass distribution function, where $d\\rho/dln(M)$ is negative ($\\dot{\\rho}$ positive), and the mass density in a given interval drops at low masses, where $d\\rho/dln(M)$ is positive ($\\dot{\\rho}$ negative).  In Figure~\\ref{nuv_rho_star_bymass} we show the formed stellar mass density growth rates produced by star formation (as calculated from Equation~\\ref{eqnmasscon2}) for both the UV-based {\\em (red squares)\\/} and the $24~\\mu$m$+$\\oii-based {\\em (blue diamonds)\\/} methods. We can now compare these with the growth rates found earlier from the formed stellar mass functions {\\em (purple lines)\\/}. The agreement between the shapes with both redshift and mass is amazingly good. However, there is a slight normalization difference, with the Equation~\\ref{eqnmasscon2} UV-based measurements being higher by about 0.1~dex in the $10^{10}-10^{10.5}$~M$_\\odot$ interval and higher by about 0.2~dex in the $10^{10.5}-10^{11}$~M$_\\odot$ and $10^{11}-10^{11.5}$~M$_\\odot$ intervals. (The offset is $-0.2$~dex in the $10^{11.5}-10^{12}$~M$_\\odot$ interval, but the uncertainties are large.) The offset is only slightly reduced if we exclude X-ray selected AGNs from the star formation calculation. Most previous analyses have also found the measurements based on star formation to be higher than the measurements based on the formed stellar mass functions (e.g., Fardal et al.\\ 2007 and references therein). However, given the completely different wavelength ranges used (UV versus NIR), the uncertainties in the extinction corrections and in the stellar models, and the effects of cosmic variance, some discrepancy would inevitably be expected. In this sense, the agreement is remarkably good, even for the normalization.  The offset is unlikely to be due to relative uncertainties in the local mass function, which generally lie at the 20\\%$-$30\\% level (e.g., compare Cole01, Bell et al.\\ 2003, and Eke et al.\\ 2005 with each other). Cole01 lies at the high end of the mass density estimates, and reducing the local mass density would increase the discrepancy. The offset is also unlikely to be due to overestimation of the UV extinction corrections, since our $24~\\mu$m$+$\\oii-based method for calculating the SFRDs is completely independent of the UV extinction corrections and shows extremely similar results to the UV-based method both in the direct SFRDs (Fig.~\\ref{sfr_comp_bymass}) and in the formed stellar mass density growth rates computed from Equation~\\ref{eqnmasscon2} (Fig.~\\ref{nuv_rho_star_bymass}).  However, the issue of the uncertainty in the population synthesis models and the recent treatments of TP-AGB stars, which we discussed in \\S\\ref{secfit}, is more complicated. If the BC03 masses which we are using are uniformly too high, then this would increase the discrepancy. However, if the local masses are close to the BC03 values and the high-redshift values are lower than the BC03 values, then we will increase the gradients in Equations~\\ref{massint1}$-$\\ref{massint5}, and this will reduce the discrepancy. We may estimate a maximum effect by assuming that the mass densities in the $z=1.2-1.5$ interval are only 60\\% of the BC03 values (the maximum correction in Bruzual 2007). This only slightly increases the formed stellar mass density growth rates in the various mass intervals, generally by less than 0.1~dex. We can see the reason for this by inspecting Figure~\\ref{mden_byz}. Typically we are building a large fraction of the local mass density over $z=0.05-1.5$ and changing the starting point mass density downward has a relatively limited effect on the required growth rate. Thus, it appears that the mass density growth rates cannot be raised enough to explain the discrepancy with the rates inferred from star formation.  A possible explanation for the offset is that the IMF is slightly different from the assumed Salpeter form. We can resolve the problem by changing the index of the IMF to $-1.10$. This is well within the uncertainties in the slope of the high-end IMF (Kroupa 2001; Chabrier 2003) and close to the Baldry \\& Glazebrook (2003) value of $-1.15$. However, in order to obtain our observed \\ha\\ to UV ratio, we must turn this over at high masses. Using the $-1.10$ index to 10~M$_\\odot$ and an index of $-1.6$ from 10~M$_\\odot$ to 100~M$_\\odot$ resolves both problems. Fardal et al.\\ (2007) argued that this type of mid--mass-weighted IMF, which they describe as paunchy, can also help in providing a consistent description of the extragalactic background light.   \\subsection{Specific Star Formation Rates} \\label{secstarspec}  We now consider the distribution of specific SFRs (SSFRs) in the individual galaxies. Although in computing the instantaneous SSFRs we are simplifying the effects of the time history in the individual galaxies, it is still useful to have a quantitative description of the range of behaviors in the galaxies. We shall return to a study of the time history of the galaxies in \\S\\ref{seccolor}.   \\begin{inlinefigure} \\figurenum{45} \\centerline{\\psfig{figure=f45.ps,angle=90,width=3.5in}} \\vskip -0.2cm \\figcaption[]{ Specific star formation rates (SFR per unit mass in the galaxy) vs. mass in the redshift intervals $z=0.05-0.475$ {\\em (red diamonds)\\/}, $z=0.475-0.9$ {\\em (black squares)\\/}, and $z=0.9-1.5$ {\\em (blue triangles)\\/}. The high-redshift sample is only shown above the limiting $3\\times10^{10}$~M$_\\odot$ to which it is complete. The large symbols show the mean values for each redshift interval. The solid lines show the inverse age of the universe at redshifts $z=0.05$ {\\em (red)\\/}, $z=0.475$ {\\em (black)\\/}, and $z=0.9$ {\\em (blue)\\/}. Only galaxies with an average SSFR above the inverse age of the universe at the redshift of the galaxy can undergo a significant change in mass. The red ($z=0.2-0.7$) and black ($z=0.85-1.1$) dotted lines show the ``main sequence'' range of Noeske et al.\\ (2007). They claim that 68\\% of the galaxies should lie within this range based on their DEEP2 observations. \\label{mass_mstar} } \\end{inlinefigure}   In Figure~\\ref{mass_mstar} we show the SSFRs versus galaxy mass in the redshift intervals $z=0.05-0.475$ {\\em (red diamonds)\\/}, $z=0.475-0.9$ {\\em (black squares)\\/}, and $z=0.9-1.5$ {\\em (purple triangles)\\/}. There is clearly a wide spread at all redshifts and masses. However, only galaxies with SSFRs larger than the inverse age of the universe at the redshift of the galaxy can change their mass significantly if those rates are maintained over the full time interval. We shall refer to such galaxies as strong star formers. Note that if the star formation is episodic, then the mass change in the galaxies will be smaller. Thus, the number of strong star formers represents an upper bound on the fraction of galaxies that may grow significantly at a given redshift. The mean SSFRs {\\em (large symbols)\\/} reproduce the results of \\S\\ref{secsfh}. That is, they show that, on average, only galaxies with masses $\\lesssim10^{11}$~M$_\\odot$ grow significantly in any of the redshift intervals. Since the high-redshift blue triangles cross the growth line at slightly higher masses than the black squares and red diamonds of the lower redshift intervals, the typical mass at which growth is taking place is downsizing in the later redshift intervals. In the lowest redshift interval where we can measure the masses below $10^{10}$~M$_\\odot$, we see that the mean SSFRs finally flatten out (at a high level). There also appears to be a rough maximum to the SSFRs of about $3\\times10^{-9}$~yr$^{-1}$.  In Figure~\\ref{dist_specific_bymass} we show the distribution functions of the SSFRs for the mass intervals (a) $10^{11}-10^{11.5}$~M$_\\odot$, (b) $10^{10.5}-10^{11}$~M$_\\odot$, and (c) $10^{10}-10^{10.5}$~M$_\\odot$. In each panel the redshift intervals $z=0.9-1.5$ {\\em (blue triangles)\\/}, $z=0.475-0.9$ {\\em (black squares)\\/}, and $z=0.05-0.475$ {\\em (red diamonds)\\/} are shown. We see little evolution in the distribution functions over the observed redshift range, but they do have very different shapes in the different mass intervals.   \\begin{inlinefigure} \\figurenum{46} \\centerline{\\psfig{figure=f46a.ps,angle=90,width=3.5in}} \\vskip -0.6cm \\centerline{\\psfig{figure=f46b.ps,angle=90,width=3.5in}} \\vskip -0.6cm \\centerline{\\psfig{figure=f46c.ps,angle=90,width=3.5in}} \\vskip -0.2cm \\figcaption[]{ Distribution functions of the specific star formation rates for logarithmic mass intervals (a) $11-11.5$~M$_\\odot$, (b) $10.5-11$~M$_\\odot$, and (c) $10-10.5$~M$_\\odot$. In each panel the red diamonds denote $z=0.05-0.475$, the black squares $z=0.475-0.9$, and the blue triangles $z=0.9-1.5$. The error bars show the 68\\% confidence limits. The blue triangles and black squares have been slightly displaced in the x-axis (by plus and minus 0.03, respectively) to allow the error bars to be distinguished. The solid vertical lines show the log of the inverse age of the universe at redshifts $z=0.05$ {\\em (red)\\/}, $z=0.475$ {\\em (black)\\/}, and $z=0.9$ {\\em (blue)\\/}. \\label{dist_specific_bymass} } \\end{inlinefigure}   In the highest mass interval (Fig.~\\ref{dist_specific_bymass}a) most of the galaxies have very low SSFRs. There are only a small fraction of stong star formers at any redshift. Overall only about 10\\% of the galaxies can be growing significantly. In the two lower mass intervals (Figs.~\\ref{dist_specific_bymass}b,c) the number of sources with strong SSFRs increases. In fact, in the lowest mass interval (Fig.~\\ref{dist_specific_bymass}c) galaxies with SSFRs above the $10^{-10}$~yr$^{-1}$ dominate the population. The percentages of strong star formers in all three mass intervals are given in Table~\\ref{tabsf}. In the lowest mass range we give the values only in the two lower redshift intervals where the sample is complete.  We may conclude from this that what star formation is occurring in the galaxies in the $10^{11}-10^{11.5}$~M$_\\odot$ interval is spread over many galaxies, and there are very few galaxies in this mass interval that are undergoing significant growth. The situation is less clear in the lower mass intervals. By the time we reach the $10^{10}-10^{10.5}$~M$_\\odot$ interval, the distribution is roughly evenly split between galaxies undergoing strong star formation and galaxies with weak star formation (see Table~\\ref{tabsf}). This could be a distinction between two populations: one with strong ongoing star formation and one with weak or little growth. Alternatively, it could be that there is a high frequency of bursting relative to steady star formation at these redshifts with all galaxies undergoing significant star formation on average. Regardless of this point, a substantial number of the low-mass galaxies have SSFRs that, if maintained over the time frame, would change their mass significantly.  We cannot easily compare our results with previous analyses of the evolution of the SSFRs over this redshift interval, such as Brinchmann \\& Ellis (2000) or Bauer et al.\\ (2005), since they did not include extinction corrections, which make substantial increases in the SSFRs. However, we can compare our results with a recent analysis by Noeske et al.\\ (2007), who used a portion of the DEEP2 sample with $K$-band and 24~$\\mu$m observations to analyze the SSFRs. They used 24~$\\mu$m plus emission line estimates of the SFRs. Noeske et al.\\ (2007) claim that the SSFRs lie within a rather tightly defined range as a function of mass. The normalization of this range increases with redshift, with the SSFRs increasing by roughly a factor of three from $z=0.3$ to $z=1$. They argue that this implies a smooth evolution in the galaxy SFRs, a result which would be inconsistent with our subsequent analysis of the star formation histories in the more massive galaxies using Balmer lines and colors. We show their ranges in Fig.~\\ref{mass_mstar} with the dotted lines {\\em (red: $z=0.2-0.7$; black: $z=0.85-1.1$)\\/}, where we have corrected their Kroupa masses to Salpeter. It is clear that the present results are inconsistent with the Noeske et al.\\ (2007) analysis. While their upper bound corresponds roughly to the maximum values seen in the present SSFRs, we see a much larger scatter in the values for the high-mass galaxies above $10^{10}$~M$_\\odot$. Our data include many galaxies with low SSFRs. The result is not dependent on the method we used to calculate the SFRs. We find the same effect using the 24~$\\mu$m plus emission line estimates of the SFRs. The result is also not a simple consequence of the optical magnitude selection used in DEEP2 ($R=24.1$), since nearly all of the high-mass galaxies would be included by such a selection, as Noeske et al.\\ (2007) discuss and we self-consistently find in the present data. The difference may lie in more subtle effects of the spectroscopic completeness versus color and optical magnitude or in the limited photometry of the DEEP2 sample.  We can also compare our results with the local analysis of Brinchmann et al.\\ (2004), and, in particular, with their Figure~24. (Their masses are based on the Kroupa IMF and must be increased by a factor of 1.54 to match ours.) Brinchmann et al.\\ (2004) only included star forming galaxies in their analysis, so their distribution is truncated at low SSFRs and the means are slightly higher. Nevertheless, the overall shape and normalization, including the roughly constant SSFRs at low mass, the decline in the SSFRs above $10^{10}$~M$_\\odot$, the upper bound on the SSFRs, and the change in the distribution of SSFRs at high mass are all in extremely good agreement with the present results.   \\subsection{Galaxy Colors, Equivalent Widths, and the 4000~\\AA\\ Break} \\label{seccolor}  Rest-frame galaxy colors and features in the spectra, such as the equivalent widths of the emission lines and the strengths of the 4000~\\AA\\ break, provide a measure of the SSFRs convolved with the recent time history of the star formation. Since the colors and the various spectral features are sensitive to different stellar mass ranges, they can provide information on the time history of the star formation and how smooth or episodic it is. (See, e.g., the Kauffmann et al.\\ 2003a analysis of the SDSS sample using the 4000~\\AA\\ break and the H$\\delta$ line.) Thus, the combination of photometric and spectroscopic information in the present sample provides a powerful tool to investigate the nature of the star formation.  Locally the SDSS results have shown that the galaxy colors are bimodal and divide into a red sequence of galaxies that are not currently undergoing star formation and a blue cloud of galaxies with active star formation (e.g., Strateva et al.\\ 2001; Baldry et al.\\ 2004). The red sequence dominates above $3\\times 10^{10}$~M$_\\odot$ (Kauffmann et al.\\ 2003b), and the blue cloud dominates below. This color bimodality is also seen in higher redshift, optically-selected samples (e.g., Bell et al.\\ 2004; Weiner et al.\\ 2005; Giallonga et al.\\ 2005; Willmer et al.\\ 2006). However, these analyses did not correct for internal extinction, which, as we shall show below, is important. Moreover, the bimodality appears to be at least partially a consequence of the optical selection and is not so strong in our mass-selected samples at the higher redshifts.  In Figure~\\ref{color_redshift} we show the rest-frame UV$-$blue (AB3400$-$AB4500) colors uncorrected for extinction for our full NIR sample versus redshift. (Note that if we change to the Vega-based magnitudes used by Willmer et al.\\ 2006 in their DEEP2 analysis, then we find a nearly identical range of colors as they.) We see almost no evolution in the color distribution over the $z=0.05-1.5$ redshift range. The precise split between the blue cloud and the red sequence depends on the mass or luminosity, but we have shown the rough split with the red dashed line. This is based on the average value of the relation given by van Dokkum et al.\\ (2000) in the appropriate luminosity range. It is not easy to see evidence for strong bimodality in this figure. Rather, we see a uniform spread of colors stretching from the blue cloud to the red sequence. Nearly all of the galaxies in the intermediate color range (sometimes referred to as the green valley) are 24$~\\mu$m sources {\\em (green triangles)\\/}.   \\begin{inlinefigure} \\figurenum{47} \\centerline{\\psfig{figure=f47.ps,angle=90,width=3.5in}} \\vskip -0.2cm \\figcaption[]{ Rest-frame $3400-4500$~\\AA\\ color vs. redshift for our NIR sample. Green triangles denote $24~\\mu$m sources, and black squares denote sources which are not detected at $24~\\mu$m. Sources with X-ray luminosities implying the presence of an AGN are enclosed in red diamonds. The red dashed horizontal line shows the approximate separation between the blue cloud and the red sequence based on the average value of the relation given by van Dokkum et al.\\ (2000) in the appropriate luminosity range. \\label{color_redshift} } \\end{inlinefigure}   The spread in colors may be more clearly seen in histogram form. In Figure~\\ref{ub_hist}a we show the colors prior to applying any extinction corrections. We see that the 24~$\\mu$m sources lie in the green valley {\\em (green dashed line)\\/}. While there is a hint of bimodality in the total sample {\\em (black solid line)\\/}, it is quite weak with sources present at all colors. In Figure~\\ref{ub_hist}b we show the colors after correcting for extinction. Now the bulk of the 24~$\\mu$m sources lie in the blue cloud, and there is a more clearly bimodal distribution in the total sample. Even with the extinction correction there are still many sources (both $24~\\mu$m and non-$24~\\mu$m) in the intermediate color region, but it is clear that applying extinction corrections is critical when analyzing the galaxy colors. Many of the sources seen in the green valley and the red sequence prior to correcting for extinction are, in fact, dusty sources with intrinsically blue colors. Quantitatively, 801 out of 2254 sources (35\\%) are in the red sequence (defining this as AB3400$-$AB4500$>1.03$) prior to the extinction correction, but nearly half of these are dusty blue galaxies. After applying the extinction correction, the number in the red sequence gets reduced to 466 out of 2254 sources, or roughly 20\\%.  The extinction corrections have a mass dependence since higher mass galaxies, with their larger column densities of gas and dust, reprocess more of their UV light. (The very highest mass galaxies will generally have lower extinctions because they are gas deficient.) The extinction corrections may also have a redshift dependence due to the evolution in the metallicity and gas content of the galaxies. Thus, any analyses of the color versus mass or color versus luminosity relations that do not apply extinction corrections will be biased.  In Figure~\\ref{color_mass_dered} we show the dereddened colors versus galaxy mass for the redshift intervals (a) $z=0.9-1.5$, (b) $z=0.475-0.9$, and (c) $z=0.05-0.475$. The red sequence is clearly seen in all of the intervals. Rather than attempt to measure the slope of the color-mass relation from the present data, we have assumed the locally determined slope of 0.08~mag per dex in mass determined by van der Wel et al.\\ (2007). We then normalized this slope to match the red sequence galaxies in the $z=0.9-1.5$ redshift interval with masses above $3\\times10^{10}$~M$_\\odot$ to obtain the red sequence relation with mass, $M$, \\begin{equation} {\\rm AB3400}-{\\rm AB4500} = 1.26+0.08(\\log M-9) \\,. \\label{sepeq} \\end{equation} We show this relation with the solid black line in Figure~\\ref{color_mass_dered}. In contrast to previous results from optically-selected and extinction uncorrected data (e.g., Bell et al.\\ 2004; van der Wel et al.\\ 2007), we see no change in the position of the red sequence with redshift. This suggests that the effect they observed was primarily a consequence of reddening. The intrinsic colors of the reddest galaxies are not changing over this redshift interval.   \\begin{inlinefigure} \\figurenum{48} \\centerline{\\psfig{figure=f48a.ps,angle=90,width=3.5in}} \\vskip -0.6cm \\centerline{\\psfig{figure=f48b.ps,angle=90,width=3.5in}} \\vskip -0.2cm \\figcaption[]{ Distribution of rest-frame $3400-4500$~\\AA\\ color for our NIR sample (a) before and (b) after correcting for internal extinction. The black histogram shows the distribution of the total sample. The green dashed (red dotted) line shows the distribution of $24~\\mu$m (non-24~$\\mu$m) sources. The blue vertical line shows the approximate separation between the blue cloud and the red sequence based on the average value of the relation given by van Dokkum et al.\\ (2000) in the appropriate luminosity range. \\label{ub_hist} } \\end{inlinefigure}   While it is clear from Figure~\\ref{ub_hist} that there is no precise split between the red sequence and the blue cloud, we may approximately separate them with a cut lying about 0.25~mag below the track of the red sequence, \\begin{equation} {\\rm AB3400}-{\\rm AB4500} = 1.01+0.08(\\log M-9) \\,. \\label{sepeqn2} \\end{equation} This allows for the spread in colors in the red sequence itself. We show this relation in Figure~\\ref{color_mass_dered} with the solid blue line. The high-mass, blue cloud galaxies are nearly all $24~\\mu$m sources {\\em (green triangles)\\/}. Moreover, most of the $24~\\mu$m sources lie in the blue cloud, though at the highest masses we also see some $24~\\mu$m sources that move into the red sequence region.   \\begin{inlinefigure} \\figurenum{49} \\centerline{\\psfig{figure=f49a.ps,angle=90,width=3.5in}} \\vskip -0.6cm \\centerline{\\psfig{figure=f49b.ps,angle=90,width=3.5in}} \\vskip -0.6cm \\centerline{\\psfig{figure=f49c.ps,angle=90,width=3.5in}} \\vskip -0.2cm \\figcaption[]{ Rest-frame $3400-4500$~\\AA\\ color corrected for extinction vs. logarithmic mass for our NIR sample in the redshift intervals (a) $z=0.9-1.5$, (b) $z=0.475-0.9$, and (c) $z=0.05-0.475$. The colors are only shown above the mass at which the sample in the given redshift interval is complete. Green triangles denote $24~\\mu$m sources, and black squares denote sources which are not detected at $24~\\mu$m. Sources with X-ray luminosities implying the presence of an AGN are enclosed in red diamonds. The solid black line shows the red sequence (with locally determined slope 0.08~mag per dex in mass) normalized to match the $z=0.9-1.5$ interval. The blue line shows this relation offset by 0.25~mag, which we adopt as the split between the red sequence and the blue cloud. \\label{color_mass_dered} } \\end{inlinefigure}   Although we can see that the red sources dominate at the high masses and that nearly all low-mass galaxies below $10^{10}$~M$_\\odot$ are blue cloud galaxies, there does not appear to be a clear transition mass in any of the redshift intervals. Rather, the fraction of galaxies with high SSFRs drops as we move to higher masses in all the redshift intervals. We also note that the split between the blue cloud and the red sequence becomes more pronounced at low redshifts, while at higher redshifts there are a considerable number of intermediate color sources.   \\begin{inlinefigure} \\figurenum{50} \\centerline{\\psfig{figure=f50a.ps,angle=90,width=3.5in}} \\vskip -0.6cm \\centerline{\\psfig{figure=f50b.ps,angle=90,width=3.5in}} \\vskip -0.2cm \\figcaption[]{ (a) EW(\\oii) and (b) the 4000~\\AA\\ break vs. the extinction corrected rest-frame AB3400$-$AB4500 color for the mid-$z$ sample ($z=0.05-0.9$) with masses greater than $10^{10}$~M$_\\odot$. The black squares (green triangles) show sources without (with) $24~\\mu$m detections. Sources containing AGNs based on their X-ray luminosities are enclosed in large red diamonds. The blue dashed lines show the rough division between the blue cloud and the red sequence. A color selection of AB3400$-$AB4500$>1.07$, which would separate the red sequence from the blue cloud at $10^{11}$~M$_\\odot$, roughly corresponds to an EW(\\oii)$<10$~\\AA\\ or a 4000~\\AA\\ break strength greater than 1.5. \\label{tests} } \\end{inlinefigure}   We may also use the EW(\\oii) or the 4000~\\AA\\ break to separate the galaxies. The EW(\\oii) is independent of the extinction correction, and the 4000~\\AA\\ break is nearly independent so these provide an invaluable check of our analysis of the colors. In particular this removes any dependence on our BC03 fitting. In Figure~\\ref{tests} we show (a) the EW(\\oii) and (b) the 4000~\\AA\\ break versus the extinction corrected rest-frame AB3400$-$AB4500 color. While there is not a perfect one-to-one relation, a color selection of AB3400$-$AB$4500>1.07$ {\\em (blue dashed vertical line)\\/}, which would separate the red sequence from the blue cloud at $10^{11}$~M$_\\odot$, roughly corresponds to an EW(\\oii)$<10$~\\AA\\ {\\em (blue dashed horizontal line in a)\\/} or a 4000~\\AA\\ break strength greater than 1.5 {\\em (blue dashed horizontal line in b)\\/}. Thus, these cuts may also be used to separate red sequence galaxies from blue cloud galaxies.   \\begin{inlinefigure} \\figurenum{51} \\centerline{\\psfig{figure=f51a.ps,angle=90,width=3.5in}} \\vskip -0.6cm \\centerline{\\psfig{figure=f51b.ps,angle=90,width=3.5in}} \\vskip -0.2cm \\figcaption[]{ Mass density per unit log mass functions for (a) the red sequence galaxies and (b) the blue cloud galaxies in the redshift intervals $z=0.9-1.5$ {\\em (blue triangles)\\/}, $z=0.457-0.9$ {\\em (black squares)\\/}, and $z=0.05-0.475$ {\\em (red diamonds)\\/}. \\label{evol_mass} } \\end{inlinefigure}   In Figure~\\ref{evol_mass} we show the mass density per unit log mass distribution functions separated by color using our dividing line between the blue cloud and the red sequence given in Equation~\\ref{sepeqn2}. There are almost equal amounts of mass in the two color-selected samples, though the red sequence is highly peaked at $10^{11}$ M$_\\odot$ while the blue sequence has a substantial contribution from lower mass galaxies. In the $z=0.475-0.9$ redshift interval the red sequence contains $1.29\\times10^{8}$ M$_\\odot$ Mpc$^{-3}$ and the blue cloud  $1.45\\times10^{8}$ M$_\\odot$ Mpc$^{-3}$ to the $10^{10}$ M$_\\odot$ completeness limit at this redshift. Figure~\\ref{evol_mass}a shows that growth in the mass density is occurring in red galaxies with masses in the interval $10^{10.5}-10^{11}$~M$_\\odot$. However, it must be noted that this result is only based on the difference between the $z=0.475-0.9$ and $z=0.9-1.5$ redhift intervals in the one mass bin and therefore the conclusion is rather weak. In contrast, Figure~\\ref{evol_mass}b shows little apparent change in the mass distribution of the blue cloud with redshift. Bell et al.\\ (2004), who first noted this effect, argued that since the star formation, and hence the mass build-up, is primarily occurring in the blue galaxies, the blue galaxies must be shifting to the red sequence at all times in order to leave the blue mass function invariant. However, Bundy et al. (2006) and Borch et al. (2006) both show evidence for a decline in the massive blue galaxies with cosmic time.   \\begin{inlinefigure} \\figurenum{52} \\centerline{\\psfig{figure=f52a.ps,angle=90,width=3.5in}} \\vskip -0.6cm \\centerline{\\psfig{figure=f52b.ps,angle=90,width=3.5in}} \\vskip -0.6cm \\centerline{\\psfig{figure=f52c.ps,angle=90,width=3.5in}} \\vskip -0.2cm \\figcaption[]{ Mass density per unit log mass functions for rest-frame (a) EW(\\oii$)<4$~\\AA, (b) 4~\\AA$~<~$EW(\\oii$)<10$~\\AA, and (c) EW(\\oii$)>10$~\\AA\\ in the redshift intervals $z=0.9-1.5$ {\\em (blue triangles)\\/}, $z=0.457-0.9$ {\\em (black squares)\\/}, and $z=0.05-0.475$ {\\em (red diamonds)\\/}. \\label{evol_o2} } \\end{inlinefigure}   We can make a finer division using the EW(\\oii). In Figure~\\ref{evol_o2} we show the mass function split into three classes according to the strength of the \\oii\\ line. In (c) we show sources with rest-frame EW(\\oii$)>10$~\\AA, which, as we have discussed above, roughly corresponds to the blue cloud color selection. As for Figure~\\ref{evol_mass}b, we see few signs of evolution. (Note that cosmic variance may be causing us some problems with the lowest redshift sources.) We then split the remaining galaxies into (b) weak emitters with $4$~\\AA$~<~$EW(\\oii$)<10$~\\AA\\ and (a) passive sources with EW(\\oii$)<4$~\\AA. With this division we can see growth occurring in both the passive sources and in the sources with weak star formation signatures, suggesting that the red sequence contains both a truly passive population with no signs of recent star formation and a population which has more recently evolved off the blue cloud and still contains signatures of recent star formation.  We may consider this further by plotting the rest-frame AB3400$-$AB8140 color versus the rest-frame EW(\\hb) to investigate the star formation history. Both are measures of the SSFRs in the galaxies and are therefore correlated with one another, but the EW(\\hb) is produced by higher mass stars and thus fades more rapidly than the AB3400$-$AB8140 color, providing a well-known age signature. In Figure~\\ref{ewhb_color} we plot the extinction corrected rest-frame AB3400$-$AB8140 color versus the rest-frame EW(\\hb) for the $z=0.05-0.9$ sample with masses (a) $10^{11}-10^{12}$~M$_\\odot$, (b) $10^{10.5}-10^{11}$~M$_\\odot$, and (c) $10^{10}-10^{10.5}$~M$_\\odot$. All but one of the galaxies in the highest mass interval (Figure~\\ref{ewhb_color}a) have little \\hb\\ emission. Most of the sources are very red, but there is a tail which extends to bluer colors. The 24~$\\mu$m detected galaxies {\\em (blue triangles)\\/} are preferentially bluer than the non-24~$\\mu$m galaxies {\\em (black squares)\\/}. We compare the observations with some simple evolutionary tracks from the BC03 models for galaxies with exponentially declining SFRs of $10^8$~yrs {\\em (black solid)\\/} (essentially a burst model), $5\\times 10^8$~yrs {\\em (red dashed)\\/}, and $5\\times 10^9$~yrs {\\em (green dotted)\\/}. The positions of the galaxies in Figure~\\ref{ewhb_color}a are not consistent with the green dotted line, where the smoothly declining star formation history would continue to produce \\hb\\ emission at intermediate AB3400$-$AB8140 colors. Rather, we appear to be seeing episodic bursts of star formation that have moved the galaxy off the red sequence, but where the massive stars powering the \\hb\\ emission have already burned away. Considering the burst model shown by the black line in Figure~\\ref{ewhb_color}a, about a quarter of the galaxies have colors which would require a burst to have occurred in the last $3\\times10^8-10^9$~yrs {\\em (thin portion)\\/}. It is these galaxies, which still have substantial UV flux, that are inferred to have high SSFRs in Figure~\\ref{dist_specific_bymass}, while the more evolved galaxies lie in the low SSFR portion of this diagram. (This emphasizes again that the SSFRs in Figure~\\ref{dist_specific_bymass} are not necessarily measures of the instantaneous star formation in the galaxy but are a time convolution of the SFR history.) In order to have this fraction of galaxies in the portion of the evolutionary track we would require bursts to occur in all galaxies about every $4\\times10^9$~yrs. If some galaxies are totally passive and do not participate in this cycling between the red sequence and the blue cloud then the remaining galaxies must have more frequent bursts.  Differential extinction, where the very massive stars producing the \\hb\\ line are more extinguished than the stars producing the UV continuum, could increase the EW(\\hb) strength. However, this effect would have to be very large to move the blue sources with very weak \\hb\\ to the smooth star formation curves, and, as we have discussed in \\S\\ref{secrelext}, we see no signs of this effect in our comparison of the Balmer line ratios and the SED derived extinctions. Errors in the UV extinction corrections may also introduce scatter in the y-axis and place some sources at bluer locations than they should have, but, again, this effect cannot be large enough to move the sources onto the smooth star formation curves. Finally, truncation of the SFRs in the smooth models can move the tracks laterally over onto the burst model on short timescales, but it would not account for the bluest sources in the figure, which can only be reproduced with short bursts.   \\begin{inlinefigure} \\figurenum{53} \\centerline{\\psfig{figure=f53a.ps,angle=90,width=3.5in}} \\vskip -0.6cm \\centerline{\\psfig{figure=f53b.ps,angle=90,width=3.5in}} \\vskip -0.6cm \\centerline{\\psfig{figure=f53c.ps,angle=90,width=3.5in}} \\vskip -0.2cm \\figcaption[]{ Extinction corrected rest-frame AB3400$-$AB8140 vs. rest-frame EW(\\hb) for the $z=0.05-0.9$ sample in the logarithmic mass intervals (a) $11-12$~M$_\\odot$, (b) $10.5-11$~M$_\\odot$, and (c) $10-10.5$~M$_\\odot$. The black squares (blue triangles) show sources without (with) $24~\\mu$m detections. Sources containing AGNs based on their X-ray luminosities are denoted by large red diamonds. The curves show the tracks expected from the BC03 models for galaxies with exponentially declining star formation rates of $10^8$~yrs {\\em (black solid)\\/}, $5\\times10^8$~yrs {\\em (red dashed)\\/}, and $5\\times10^9$~yrs {\\em (green dotted)\\/}. The black curve is divided into ages of $3\\times10^{8}-10^9$~yrs {\\em (thin portion)\\/} and greater than $10^9$~yrs {\\em (thick portion)\\/}. \\label{ewhb_color} } \\end{inlinefigure}   The intermediate mass galaxies in Figure~\\ref{ewhb_color}b have a larger fraction of galaxies with blue AB3400$-$AB4500 colors, suggesting that the burst frequency is higher in these galaxies. For the lowest mass galaxies shown in Figure~\\ref{ewhb_color}c there appears to be a distinction between the 24~$\\mu$m sources {\\em (blue triangles)\\/}, which still preferentially lie along the burst track, and the non-24~$\\mu$m sources {\\em (black squares)\\/}, which are more consistent with smooth ongoing star formation. This suggests that it is the burst process which results in the dusty galaxies producing the 24~$\\mu$m emission. In the higher mass galaxies bursting is the dominant process and all the galaxies with high SFRs are 24~$\\mu$m sources.   \\subsection{Galaxy Morphologies} \\label{secmassmorph}  Galaxy morphologies, while closely related to the colors and spectral properties of the galaxies, provide an alternative view of the evolution. In particular, we would like to see in which types of galaxies the star formation and mass evolution is occurring and use this information to clarify the relationship between star formation and stellar growth in the mass density.  The Ellis morphological classifications that we are using (see \\S\\ref{secmorph}) are based on the {\\em HST\\/} F850LP images. Therefore, in order to avoid biasing in type by observing the galaxies in the rest-frame UV rather than in the rest-frame optical, we will restrict the NIR sample to only galaxies at $z<1.2$ for this section. This ensures that the F850LP band corresponds to rest-frame wavelengths above 4000~\\AA.  We first compare the morphological typings with the spectral characteristics of the galaxies. In Figure~\\ref{morph_split} we show the distribution of the galaxy types in both rest-frame EW(\\hb) and 4000~\\AA\\ break strength. The black symbols show the E/S0 galaxies (classes $0-2$), the green symbols show the Sab and S galaxies (classes $3-4$), and the small red symbols show the Scd and Irr galaxies (classes $5-6$). The large red squares show the galaxies classified as Mergers (class 8). Both the spectral and the morphological typings place nearly all of the massive galaxies ($10^{11}-10^{12}$~M$_\\odot$; Figure~\\ref{morph_split}a) into the E/S0 or Sab-S categories, while the lower mass galaxies ($10^{10}-10^{11}$~M$_\\odot$; Figure~\\ref{morph_split}b) show a much wider distribution, including many Irr. Figure~\\ref{morph_split} may be compared with Figure~4 of Barbaro \\& Poggianti (1997) for a local sample. The positions of the morphological types in the 4000~\\AA\\ break-EW(\\hb) plane match closely the positions of the local values. There is no change in this distribution over the redshift range $z=0.05-0.9$.  In Figure~\\ref{ew_morph} we quantitatively show the distribution of the EW(\\hb) by morphological type for the mass intervals (a) $10^{11}-10^{12}$~M$_\\odot$, (b) $10^{10.5}-10^{11}$~M$_\\odot$, and (c) $10^{10}-10^{10.5}$~M$_\\odot$. We divide the galaxies into the broader classes of E/S0s (classes $0-2$) {\\em (solid black line)\\/}, Spirals (classes $3-5$) {\\em (dashed red line)\\/}, and Peculiars (classes 6 and 8) {\\em (dotted cyan line)\\/}. The distribution of equivalent widths is very similar for the Spirals and the Peculiars, but it is a strong function of mass. Most of the massive Spirals are only weak star formers, while the lower mass Spirals have a much wider distribution of equivalent widths, and the mean equivalent width is much larger.   \\begin{inlinefigure} \\figurenum{54} \\centerline{\\psfig{figure=f54a.ps,angle=90,width=3.5in}} \\vskip -0.6cm \\centerline{\\psfig{figure=f54b.ps,angle=90,width=3.5in}} \\vskip -0.2cm \\figcaption[]{ Distribution of morphological types in both 4000~\\AA\\ break strength and rest-frame EW(\\hb) in the mid-$z$ sample ($z=0.05-0.9$) for the logarithmic mass intervals (a) $11-12$~M$_\\odot$ and (b) $10-11$~M$_\\odot$. The black symbols show the E/S0 galaxies (solid squares = class 0, solid triangles = class 1, and open downward-pointing triangles = class 2); the green symbols show the Sab and S galaxies (diamonds = class 3, leftward-pointing triangles = class 4); the small red symbols show the Scd and Irr galaxies (open squares = class 5, rightward-pointing triangles = class 6); and the large red squares show the Mergers (class 8). \\label{morph_split} } \\end{inlinefigure}   In Figure~\\ref{mass_morph} we show galaxy morphological type versus mass for the redshift intervals $z=0.05-0.475$ {\\em (red diamonds)\\/}, $z=0.475-0.9$ {\\em (black squares)\\/}, and $z=0.9-1.2$ {\\em (blue triangles)\\/}. For each redshift range we only show galaxies above our mass completeness level of $2\\times10^9$~M$_\\odot$, $10^{10}$~M$_\\odot$ and $2\\times10^{10}$~M$_\\odot$ respectively. Confirming a well-known result, we see a strong correlation between the galaxy morphology and the galaxy mass, with many of the most massive galaxies being E/S0s (classes $0-2$). The large symbols show the median morphological types by mass and by redshift. A strong evolution with redshift in the mass-morphology relation is also evident. As an example, the typical $10^{11}$~M$_\\odot$ galaxy has moved from being an Sb-like (class 4) galaxy in the $z=0.9-1.2$ redshift interval {\\em (blue)\\/} to being an S0 (class 2) galaxy in the $z=0.05-0.475$ redshift interval {\\em (red)\\/}, but this kind of morphological type evolution is present across the entire mass range.   \\begin{inlinefigure} \\figurenum{55} \\centerline{\\psfig{figure=f55a.ps,angle=90,width=3.5in}} \\vskip -0.6cm \\centerline{\\psfig{figure=f55b.ps,angle=90,width=3.5in}} \\vskip -0.6cm \\centerline{\\psfig{figure=f55c.ps,angle=90,width=3.5in}} \\vskip -0.2cm \\figcaption[]{ Distribution of rest-frame EW(\\hb) for galaxies in the redshift interval $z=0.05-0.9$ and in the logarithmic mass intervals (a) $11-12$~M$_\\odot$, (b) $10.5-11$~M$_\\odot$, and (c) $10-10.5$~M$_\\odot$. In each panel the solid black line shows the E/S0 galaxies (classes $0-2$), the red dashed line shows the Spirals (classes $3-5$), and the dotted blue line shows the Peculiars (classes 6 and 8). \\label{ew_morph} } \\end{inlinefigure}   We may now examine the morphological type distribution in which the stellar growth in the mass density is occurring. As we showed in \\S\\ref{seccolor} (Fig.~\\ref{evol_o2}), the mass build-up in the redshift range $z=0.05-0.9$ is primarily in the passive (EW(\\oii$)<4$~\\AA) and weakly active galaxies (4~\\AA$~<~$EW(\\oii$)<12$~\\AA) in the $10^{10.5}-10^{11}$~M$_\\odot$ interval. In Figure~\\ref{ewo2_morph_dist} we show the distribution of morphological types in this redshift and mass interval split into passive galaxies {\\em (black solid line)\\/}, weakly active galaxies {\\em (red dashed line)\\/}, and strong emission line galaxies {\\em (blue dotted line)\\/}. It can be seen that the weakly active galaxies primarily lie in the spiral and S0 classes, though they are more strongly weighted to S0s than are the strong emitters, while the passive galaxies predominantly lie in the E/S0 classes. (We note in passing that we have visually checked the emission line galaxies that are morphologically classified as E, and these classifications are generally robust.)   \\begin{inlinefigure} \\figurenum{56} \\centerline{\\psfig{figure=f56.ps,angle=90,width=3.5in}} \\vskip -0.2cm \\figcaption[]{ Galaxy morphological types (classes $0-8$) vs. mass for the redshift intervals $z=0.05-0.475$ {\\em (red diamonds)\\/}, $z=0.475-0.9$ {\\em (black squares)\\/}, and $z=0.9-1.2$ {\\em (blue triangles)\\/}. For each redshift range we only show galaxies above our mass completeness level of $2\\times10^9$~M$_\\odot$, $10^{10}$~M$_\\odot$ and $2\\times10^{10}$~M$_\\odot$ respectively. In each redshift interval the large symbols show the median values for that mass interval with 68\\% confidence limits. The error bars are generally one morphological class or less which can result in an asymmetrical appearance. \\label{mass_morph} } \\end{inlinefigure}   \\begin{inlinefigure} \\figurenum{57} \\centerline{\\psfig{figure=f57.ps,angle=90,width=3.5in}} \\vskip -0.2cm \\figcaption[]{ Distribution of morphological types for galaxies in the $z=0.05-0.9$ redshift range with logarithmic masses $10.5-11$~M$_\\odot$. The black solid line shows passive galaxies with EW(\\oii$)<4~$\\AA, the red dashed line shows weakly active galaxies with 4~\\AA$~<~$EW(\\oii$)<12$~\\AA, and the blue dotted line shows strong emission line galaxies with EW(\\oii$)>12$~\\AA. \\label{ewo2_morph_dist} } \\end{inlinefigure}   \\begin{figure*} \\figurenum{58} \\centerline{\\psfig{figure=f58a.ps,angle=90,width=3.5in} \\psfig{figure=f58b.ps,angle=90,width=3.5in}} \\centerline{\\psfig{figure=f58c.ps,angle=90,width=3.5in} \\psfig{figure=f58d.ps,angle=90,width=3.5in}} \\vskip -0.2cm \\figcaption[]{ Build-up of the mass density per unit logarithmic mass for the redshift intervals $z=0.9-1.2$ {\\em (blue triangles)\\/}, $z=0.475-0.9$ {\\em (black squares)\\/}, and $z=0.05-0.475$ {\\em (red diamonds)\\/} and for the morphological classes (a) E (class 0), (b) S0 (classes $1-2$), (c) Spirals (classes $3-5$), and (d) Peculiars (classes 6 and 8). \\label{evol_morph_mass} } \\end{figure*}   In Figure~\\ref{evol_morph_mass} we show the build-up of the mass density per unit logarithmic mass as a function of redshift and galaxy type. We have separated the E/S0 class into E (class 0) and S0 (classes $1-2$) classes, though if we instead separate it into E (class $0-1$) and S0 (class 2) classes, it does not change our conclusions. Here it can be seen that both the E (Fig.~\\ref{evol_morph_mass}a) and S0 (Fig.~\\ref{evol_morph_mass}b) classes are building up strongly in the $10^{10.5}-10^{11}$~M$_\\odot$ interval, while the Spirals (Fig.~\\ref{evol_morph_mass}c) and Peculiars (Fig.~\\ref{evol_morph_mass}d) are not changing significantly. When we further separate the E/S0 galaxies by the EW(\\oii), we find that all of the growth is in the passive or weakly active galaxies.  Thus, it appears that the mass build-up is primarily moving into elliptical and S0 galaxies. In other words, mass formation occurs in the spiral galaxies, and the spiral galaxies gradually move into the E/S0 class with decreasing redshift as the overall star formation dies away. This results in the mass function of the strong emitters and spiral galaxies being relatively invariant and the mass build-up being primarily seen in the passive and weakly active E/S0 galaxies.   \\subsection{Galaxy Environments} \\label{secenv}  As described in \\S\\ref{secgalenv}, we use the projected nearest neighbor density, $\\Sigma_3$, to characterize the galaxy environment. In order to provide a uniform sample with a sufficiently high density to minimize edge effects, we use the galaxy sample with masses greater than $2\\times10^{10}$~M$_\\odot$ to compute $\\Sigma_3$. This restricts our analysis to the redshift range $z=0.3-1.2$. The lower redshift bound is set by the size of the field and our edge restriction (objects must be more than 1~Mpc from the field edge), and the upper redshift bound is set by the mass limit of the sample. The average density is $\\Sigma_3=0.96$~Mpc$^{-2}$, and very few objects have densities less than 0.3~Mpc$^{-2}$, where edge effects begin to enter. (A $4\\times10^{10}$~M$_\\odot$ sample gives an average $\\Sigma_3=0.31$\\ Mpc$^{-2}$, where this issue would be more significant.)   \\begin{inlinefigure} \\figurenum{59} \\centerline{\\psfig{figure=f59.ps,angle=90,width=3.5in}} \\vskip -0.2cm \\figcaption[]{ Galaxies in the $z=1.0156$ sheet in the GOODS-N region are shown by the red squares (large symbols correspond to galaxies with masses above $2\\times10^{11}$~M$_\\odot$; small symbols to lower mass galaxies). The full region of the field is shown by the black rectangle. The galaxies with measured $\\Sigma_3$ are shown with small black squares. This region is smaller than the full field because of the edge constraint. Galaxies in the velocity slice with $\\Sigma_3>3$~Mpc$^{-2}$ are shown enclosed in green squares. The large blue symbol shows the position of the diffuse X-ray emission from Bauer et al.\\ (2002). \\label{goods_group} } \\end{inlinefigure}   Nearly all of the sources with $\\Sigma_3>5$\\ Mpc$^{-2}$ lie in just four velocity sheets at $z=0.4851$,\\ 0.8472,\\ 0.9367,\\ and 1.0156. The strongest of these is the well-known structure at $z=1.0156$, which we illustrate in Figure~\\ref{goods_group} {\\em (red squares)\\/}. This feature is dominated by a fairly substantial cluster lying at the southern end of the GOODS-N region, together with a smaller concentrated group to the north. We have identified 29 galaxies in the southern cluster with masses above $2\\times10^{10}$~M$_\\odot$. The total stellar mass of these 29 galaxies alone is $4\\times10^{12}$~M$_\\odot$. The velocity dispersion is 470~km~s$^{-1}$. There is associated diffuse X-ray emission centered on one part of the cluster (Bauer et al.\\ 2002) {\\em (blue square)\\/}. Within the region where the density parameter can be measured, nearly all of the galaxies at $z=1.0156$ lie in substantially overdense regions with $\\Sigma_3>3~$Mpc$^{-2}$ {\\em (green squares)\\/}. The other sheets are weaker and, in some cases (e.g., the $z=0.8472$ structure), more diffuse.  Because the redshift intervals have different mass limits, it is most natural to plot the density parameter versus mass. In Figure~\\ref{mass_dens} we show this relation for the redshift intervals $z=0.9-1.2$ {\\em (blue triangles)\\/}, $z=0.475-0.9$ {\\em (black squares)\\/}, and $z=0.3-0.475$ {\\em (red diamonds)\\/}. The mass-density relation is clearly seen in the median values {\\em (large symbols)\\/}. At lower masses (below $10^{11}$~M$_\\odot$) the dependence on the environment is very weak. However, nearly all of the most massive galaxies lie in higher density regions. There is relatively little evolution in the mass-density relation over the $z=0.05-0.9$ redshift range.   \\begin{inlinefigure} \\figurenum{60} \\centerline{\\psfig{figure=f60.ps,angle=90,width=3.5in}} \\vskip -0.2cm \\figcaption[]{ $\\Sigma_3$ density parameter vs. stellar mass for the redshift intervals $z=0.9-1.2$ {\\em (blue triangles)\\/}, $z=0.475-0.9$ {\\em (black squares)\\/}, and $z=0.3-0.475$ {\\em (red diamonds)\\/}. In each case the large symbols show the median values for the mass intervals with 68\\% confidence limits. \\label{mass_dens} } \\end{inlinefigure}   While this relationship is well known in general terms, it is not easy to compare it precisely with other results, either because the environmental parameters are expressed in different ways, or because the increases are measured for optical luminosity rather than for mass (e.g., Croton et al.\\ 2005; Hoyle et al.\\ 2005; Cooper et al.\\ 2007, 2008). However, Baldry et al.\\ (2006) provide a local analysis of the galaxy masses using a similar environmental parameter. They characterize their Schechter function fits to the galaxy mass functions in different environments with the mass at which the contribution to the local galaxy mass density per dex in galaxy mass peaks. They show in their Figure~8d how this peak mass depends on $\\Sigma$. We have measured this quantity for our $z=0.6-1.2$ sample and show the result in Figure~\\ref{mass_dens_plot}. In order to compare with Baldry et al.\\ (2006) {\\em (blue solid line)\\/}, we have adjusted their Kroupa masses to Salpeter masses, and we have matched their median $\\Sigma$ parameter to ours. We see from Figure~\\ref{mass_dens_plot} that the mass versus $\\Sigma_3$ relation has a very similar slope at both redshifts but that the mass in a given environment is about 0.3~dex higher at $z=0.9-1.2$ than it is now. This is expected, since the high-mass galaxies are already in place at the higher redshifts, while the low-mass galaxies are still forming. However, it shows that this relative growth of the low-mass galaxies is occurring across our measured density range and is not a strong function of environment.   \\begin{inlinefigure} \\figurenum{61} \\centerline{\\psfig{figure=f61.ps,angle=90,width=3.5in}} \\vskip -0.2cm \\figcaption[]{ Mass at which contribution to the mass density per dex in galaxy mass peaks vs. $\\Sigma_3$ density parameter. The black squares show the results from the present data in the $z=0.9-1.2$ redshift interval. The blue line shows the corresponding local result derived by Baldry et al.\\ (2006). \\label{mass_dens_plot} } \\end{inlinefigure}   However, unlike the local analysis of Baldry et al.\\ (2006), we do not see an environmental dependence for the distribution of galaxies between the blue cloud and the red sequence at $z=0.6-1.2$. In Figure~\\ref{ub_mass_dens} we show the distribution of colors for galaxies with $z=0.6-1.2$ separated by both mass and environment {\\em (black histograms)\\/}. We compare this with the distribution of all galaxies in each mass interval {\\em (green curves)\\/} normalized to the number of galaxies in that particular sample {\\em (number in upper right corner)\\/}. The blue histograms show sources detected at $24~\\mu$m which dominate the blue cloud at these masses, and the red histograms show sources which are not detected at $24~\\mu$m which dominate the red sequence. It can be seen that the distributions are essentially invariant with environment, while the fraction of red galaxies increases with mass. Thus, the environmental dependence of the red fraction seen in the local sample must have been imprinted over the $z=0-1$ redshift interval. This suggests that the star formation switch-off may have been more rapid in the higher density environments.   \\begin{figure*} \\figurenum{62} \\centerline{\\psfig{figure=f62a.ps,angle=90,width=3.5in} \\psfig{figure=f62b.ps,angle=90,width=3.5in}} \\centerline{\\psfig{figure=f62c.ps,angle=90,width=3.5in} \\psfig{figure=f62d.ps,angle=90,width=3.5in}} \\figcaption[]{ Distribution of the AB$3400-$AB4500 colors separated by both mass and environment for galaxies with $z=0.6-1.2$ {\\em (solid black histogram)\\/}. In (a) and (b) we show the distributions for the mass interval $2\\times10^{10}-10^{11}$~M$_\\odot$. In (c) and (d) we show the distributions for the mass interval $10^{11}-10^{12}$~M$_\\odot$. (a) and (c) correspond to $\\Sigma_3<1$~Mpc$^{-2}$, and (b) and (d) correspond to $\\Sigma_3$ greater than this value. The green curves show the distributions of all galaxies in the given mass interval normalized to the number of galaxies in that particular sample, which is shown in the upper right corner. The red (blue) histograms show the distribution of galaxies detected (undetected) at 24~$\\mu$m, respectively. \\label{ub_mass_dens} } \\end{figure*}   \\subsection{Metal Evolution} \\label{secmetev}  We show the evolution of the metallicity-mass relation with redshift in Figure~\\ref{tremonti}. In Figure~\\ref{tremonti}a we compare the locally derived metallicity-mass relation of Tremonti04 {\\em (black solid curve)\\/} with our relations derived from the R23 method (using the Tremonti04 calibration) for $z=0.05-0.475$ {\\em (red)\\/} and $z=0.475-0.9$ {\\em (cyan)\\/}. Since the Tremonti04 masses are computed for the Kroupa (2001) IMF, we had to increase them by a factor of 1.54 to make the comparison (see \\S\\ref{secintro}). As we have discussed in \\S\\ref{secintro}, the adopted IMF does not otherwise affect the results. We show both the polynomial fits to the metallicity-mass relations {\\em (colored lines)\\/} and the median values and errors in various mass bins {\\em (symbols)\\/}.   \\begin{inlinefigure} \\figurenum{63} \\centerline{\\psfig{figure=f63a.ps,angle=90,width=3.5in}} \\vskip -0.6cm \\centerline{\\psfig{figure=f63b.ps,angle=90,width=3.5in}} \\vskip -0.2cm \\figcaption[]{ (a) Metallicity-mass relations in the $z=0.05-0.475$ {\\em (red triangles)\\/} and $z=0.475-0.9$ {\\em (cyan diamonds)\\/} redshift intervals derived from the R23 method using the Tremonti04 calibration and compared with Tremonti04's local metallicity-mass relation {\\em (black curve)\\/}. The colored lines show the polynomial fits to the individual data points. The solid symbols show the median values in the mass bins with 68\\% confidence limits. The medians lie slightly higher than the average values represented by the fits. The dotted red line shows the solar abundance. The black dashed (dotted) line shows the local relation reduced by 0.14~dex (0.22~dex) to match the data in the $z=0.05-0.475$ ($z=0.475-0.9$) interval. (b) Metallicity-mass relation in the $z=0.05-0.475$ redshift interval {\\em (red solid line and solid triangles)\\/} computed using the \\nii/\\ha\\ diagnostic and compared with the Shapley et al.\\ (2004) measurements of the LBG sample computed using the same method {\\em (solid black region)\\/}. The red dashed line shows the red solid line reduced by 0.32~dex to match the LBGs. The black curve, the red open triangles, and the cyan open diamonds are all taken from (a) and show the R23-based results. \\label{tremonti} } \\end{inlinefigure}   We summarize the decrease in the metallicities with increasing redshift and decreasing mass in Table~\\ref{metal_drop}. In column~2 we give the average of the Tremonti04 values for galaxies in the given mass interval, and then in columns~3 and 4 we give the drop from Tremonti04 to each of our average values, respectively, for $z=0.05-0.475$ and $z=0.475-0.9$. The errors are 68\\% confidence limits.  The data are not adequate to determine if the shape of the local relation has changed with redshift. There are hints at the $2\\sigma$ level that there is less evolution at the high-mass end, but within the accuracy that can be obtained with our data, the shape of the metallicity-mass relation over the observed mass range could be invariant from $z=0.05-0.9$. However, if we normalize the local relation to our data at $z=0.05-0.475$ {\\em (black dashed curve)\\/} and at $z=0.475-0.9$ {\\em (black dotted curve)\\/}, then the best fits show that the metallicity in the $10^{10}-10^{11}$~M$_\\odot$ interval is lower by $0.10\\pm0.04$~dex than the local value at the median redshift of $0.44$ (low-redshift interval) and lower by $0.21\\pm0.03$~dex than the local value at the median redshift of $0.75$ (high-redshift interval).  The conclusion that there is a decrease in the metallicity at a given mass with increasing redshift from $z=0.05-0.9$ is consistent with previous work, though the precise values have varied considerably. The most recent work of Savaglio et al.\\ (2005) gives a considerably larger change in the normalization and also finds a steeper slope at $z\\sim0.7$, namely $12+\\log({\\rm O/H})=8.84+(0.48\\pm0.06)\\log M_{10}$, which may be compared with our relation of $12+\\log({\\rm O/H})=8.70+(0.17\\pm0.05)\\log M_{10}$ (Eq.~\\ref{eqr23hizrel_tr}). The Savaglio et al.\\ (2005) relation is fitted over a much wider mass range (down to masses below $10^{9}$~M$_\\odot$) using a sample that is substantially incomplete and biased towards star formers at the lower masses. This weights the lower mass bins to lower metallicities and steepens the fit.  In Figure~\\ref{tremonti}b we compare the median metallicities {\\em (red solid triangles)\\/} and the least-square polynomial fit {\\em (red solid line)\\/} computed from the \\nii/\\ha\\ method in the $z=0.05-0.475$ redshift interval with the $z\\sim2.1$ Lyman break galaxy (LBG) sample of Shapley et al.\\ (2004), which was also computed using this diagnostic {\\em (black solid region)\\/}. The R23-based measurements from Figure~\\ref{tremonti}a are also shown {\\em (black curve and colored open symbols)\\/}. The R23-based red open triangles for the same redshift interval are not significantly different than the \\nii/\\ha-based red solid triangles. The LBG galaxies lie about 0.32~dex lower {\\em (dashed red line)\\/} than the median $z=0.44$ galaxies {\\em (red solid line)\\/} at the same mass and are similar in metallicity to local galaxies that are almost an order of magnitude lower in mass.   \\subsection{Gas Masses} \\label{secgasmass}  In the simplest closed-box model for metal evolution, the metallicity $Z$(O) (the fraction by mass of O) in the gas is simply related to the oxygen yield $y$(O) by the well-known relation \\begin{equation} Z({\\rm O}) = y({\\rm O})~\\log(M_{g}/M_{T})  \\,, \\label{meteq} \\end{equation} where $M_{g}$ is the gas mass and $M_{T}$ is the sum of the gas mass and the stellar mass. The change in metallicity for a change in stellar mass is \\begin{equation} \\delta Z({\\rm O}) = y({\\rm O})~\\delta M_{star}/M_{g}  \\,. \\label{deriveq} \\end{equation} Thus, the derivative of $Z$(O) with respect to the stellar mass measures the quantity $y$(O)/$M_{g}$, and, if we assume a value for the yield, we can derive the gas mass. Essentially we are measuring the gas reservoir required to dilute the metals returned from the known star formation to match the observed metal evolution.  The effective yields have been empirically measured for local galaxies and are found to be approximately independent of galaxy mass for masses above $10^{9.5}$~M$_\\odot$ (Garnett 2003). Tremonti04 find a weak dependence on mass in this mass range, but their results depend on using star formation as a proxy for gas mass. We will assume a time-independent and mass-independent value of $\\log y({\\rm O})=-1.9$, which is probably a reasonable approximation given the uncertainties.  We roughly compute the gas reservoir mass densities using the mass change between $z=0.05$ and $z=0.77$ from the least-square fits of Equations~\\ref{massint2}$-$\\ref{massint4} and the corresponding change in $Z$(O) from Table~1. For the lowest logarithmic mass interval $9.5-10$~M$_\\odot$ we made the fit over $z=0.05-0.44$, where the mass sample is complete. If the local mass density is lower than the Cole01 estimate this would reduce $\\delta M_{star}$ and hence the inferred gas mass. Deriving the values from the star formation instead gives broadly similar results, with the largest change being an increase by almost a factor of two in $M_g$ in the logarithmic mass interval $10-10.5$~M$_\\odot$.   \\begin{inlinefigure} \\figurenum{64} \\centerline{\\psfig{figure=f64.ps,angle=90,width=3.5in}} \\vskip -0.2cm \\figcaption[]{ Mass density of the gas reservoirs inferred from the metal evolution compared with the stellar mass density history. The mass density of the gas reservoirs is shown with colored lines, and the stellar mass density history is shown with corresponding colored symbols in the mass intervals $10^{9.5}-10^{10}$~M$_\\odot$ {\\em (blue line and triangles)\\/}, $10^{10}-10^{10.5}$~M$_\\odot$ {\\em (black dashed line and open squares)\\/}, $10^{10.5}-10^{11}$~M$_\\odot$ {\\em (black solid line and solid squares)\\/}, and $10^{11}-10^{11.5}$~M$_\\odot$ {\\em (red line and diamonds)\\/}. The gas mass densities are computed over $z=0.05-0.77$ for the higher mass intervals and over $z=0.05-0.44$ for the lowest mass interval. The local mass densities for each of the mass intervals were obtained by integrating the Cole01 function. \\label{mgas_byz} } \\end{inlinefigure}   We compare the gas mass densities {\\em (colored lines)\\/} with the stellar mass density history {\\em (corresponding colored symbols)\\/} for the mass intervals $10^{9.5}-10^{10}$~M$_\\odot$ {\\em (blue line and triangles)\\/}, $10^{10}-10^{10.5}$~M$_\\odot$ {\\em (black dashed line and open squares)\\/}, $10^{10.5}-10^{11}$~M$_\\odot$ {\\em (black solid line and solid squares)\\/}, and $10^{11}-10^{11.5}$~M$_\\odot$ {\\em (red line and diamonds)\\/} in Figure~\\ref{mgas_byz}. (The gas and stellar mass densities are derived using the Salpeter IMF and thus will change consistently if we use an alternate IMF.) In the highest mass interval {\\em (red)\\/} where the stellar mass is growing only slowly, the sum of the gas plus stellar mass densities at $z=0.44$ is roughly comparable to the local stellar mass density obtained by integrating the Cole01 function. By contrast, the galaxes in the lower mass intervals where assembly is still progressing have much larger gas reservoirs. Thus, these results suggest that star formation is terminating as a consequence of gas depletion.  We estimate the total gas mass density over all galaxies larger than $10^{9.5}$~M$_\\odot$ at $z=0.44$ to be $\\sim7\\times10^{8}$~M$_\\odot$~Mpc$^{-3}$. We can compare this with the current mass density in stars over the same mass range, namely $4.7\\times10^{8}$~M$_\\odot$~Mpc$^{-3}$ obtained by integrating Cole01. This suggests that there is still a significant amount of gas mass to be converted into stars. From Figure~\\ref{mgas_byz} we can see that this principally lies in galaxies with masses below $10^{11}$~M$_\\odot$. We may very crudely infer that at the present time the highest mass galaxies will have very little gas mass remaining, while the lower mass galaxies will have comparable amounts of gas and stars. However, the uncertainties in this are very large, both from the observations and from the overly simple modeling."
    },
    "0806/0806.4041_arXiv.txt": {
        "abstract": "We present a 3-dimensional model of supernova remnants (SNRs) where the hydrodynamical evolution of the remnant is modeled consistently with nonlinear diffusive shock acceleration occuring at the outer blast wave. The model includes particle escape and diffusion outside of the forward shock, and particle interactions with arbitrary distributions of external ambient material, such as molecular clouds. We include synchrotron emission and cooling, bremsstrahlung radiation, neutral pion production, inverse-Compton (IC), and Coulomb energy-loss. Boardband spectra have been calculated for typical parameters including dense regions of gas external to a $1000$ year old SNR. In this paper, we describe the details of our model but do not attempt a detailed fit to any specific remnant. We also do not include magnetic field amplification (MFA), even though this effect may be important in some young remnants. In this first presentation of the model we don't attempt a detailed fit to any specific remnant. Our aim is to develop a flexible platform, which  can be generalized to include effects such as MFA, and which can be easily adapted to various SNR environments, including Type Ia SNRs, which explode in a constant density medium, and Type II SNRs, which explode in a pre-supernova wind. When applied to a specific SNR, our model will predict cosmic-ray spectra and multi-wavelength morphology in projected images for instruments with varying spatial and spectral resolutions. We show examples of these spectra and images and emphasize the importance of measurements in the hard X-ray, GeV, and TeV gamma-ray bands for investigating key ingredients in the acceleration mechanism, and for deducing whether or not TeV emission is produced by IC from electrons or \\pion\\ from protons. ",
        "introduction": "Supernovae (SNe) are the only known sources capable of providing the energy needed to power the bulk of the galactic cosmic rays (CRs) with energies below the spectral feature called the ``knee'' around $3 \\times 10^{15}$\\,eV  \\citep[e.g.,][]{Drury83}. If SNe are the main sources of Galactic CRs, the acceleration mechanism must be efficient so that $\\gtrsim 10$\\% of the total SN explosion energy in our Galaxy ends up in cosmic rays \\citep[e.g.,][]{Hillas2005}. Observational evidence that the outer blast wave shock accelerates electrons to ultra-relativistic energies in some young SNRs \\citep[e.g.,][]{KoyamaSN1006_95}, and the existence of a well-developed model of particle acceleration at shocks, i.e., diffusive shock acceleration (DSA) \\cite[e.g.,][]{Drury83,BE87,JE91} support the above contention.  When confronting observations with theoretical models, however, there remain a number of important ambiguities and uncertainties from both the observational and theoretical perspectives. Resolution of these ambiguities and uncertainties by new telescopes will be essential to claim evidence for the pion-decay feature in the GeV-TeV emission from SNRs.  The Gamma-ray Large Area Space Telescope (GLAST), to be launched in 2008, will probe this crucial energy range with unprecedented sensitivity and resolution.  Fundamental questions for CR origin also concern the spectral shape and maximum ion energy a given SNR can produce. Electron energy spectra inferred from young SNRs vary and can be substantially harder than CR electron spectra observed at Earth, even after correction for propagation in the galaxy \\citep[e.g.,][]{BV2006}. The maximum CR ion energy SNRs actually produce will remain uncertain until a firm identification of \\pion\\ emission is obtained and gamma-ray emission is detected past a few 100~TeV, the maximum possible electron energy in SNRs.  There remain other basic questions concerning the DSA mechanism. For instance, is DSA efficient enough for nonlinear effects, such as shock smoothing and magnetic field amplification, to become important in young SNRs? How does particle injection occur and how does injection and acceleration vary between electrons and protons? While the galactic CR electron-to-proton ratio, $\\epRatio$, of $0.01$--$0.0025$ observed at Earth at \\rel\\ energies is often used to constrain the ratio in SNRs, this ratio has not been observed outside of the heliosphere.\\footnote{We note that while energetic electrons and protons are observed from solar flares and at low Mach number heliospheric shocks, these observations provide limited help for understanding the high Mach number shocks expected in young SNRs and other astrophysical sourses where a large fraction of the shock energy is put into \\rel\\ particles.} The $\\epRatio$ ratio is crucial in deciding whether the $\\gamma$-ray emission from different SNRs, or observed in different parts of an individual SNR, is of hadronic or leptonic origin.  The recent discovery of spatially thin, hard X-ray filaments in some young SNRs \\citep[e.g.,][]{BambaEtal2003,Uchiyama_J1713_2007} supports previous suggestions \\citep[e.g.,][]{Cowsik80,BL2001,RE92} that the particle acceleration process can amplify the ambient magnetic field by large factors. If magnetic field amplification in DSA is as large as now appears to be the case \\citep[e.g.,][]{BKV2003}, it will have far-reaching consequences not only for understanding the origin of Galactic CRs, but for interpreting \\syn\\ emission from shocks throughout the universe.  Since shocks and related superthermal particle populations exist in diverse environments, the knowledge gained from studying SNRs will have wide applicability.  The advent of new space- and ground-based telescopes will result in observations of SNRs at many different wavelengths with greatly improved sensitivity and resolution. It is even conceivable that features in the CR spectrum observed at Earth might be associated with nearby SNRs with future observations \\citep[e.g.,][]{EW1999,KobayashiEtal2001}.  In order to take full advantage of current and future observations, and to improve our understanding of the DSA mechanism, the data must be analyzed with consistent, broadband photon emission models including nonlinear effects. This has prompted us to develop a three-dimensional model of young SNRs where the evolution of the remnant is coupled to nonlinear diffusive shock acceleration (NL-DSA) \\citep[e.g.,][]{EDB2004,EC2005}, in an environment with an arbitrary mass distribution. We focus on radiation from CR electrons and protons and leave the modeling of heavier ions for future work. In this preliminary study, we also ignore other possible acceleration processes, most notably second-order stochastic acceleration, and do not include magnetic field amplification.  We believe our work is a significant advance over previous work for several reasons.  Of particular importance is that we include ``escaping'' particles \\SCly. In NL-DSA, a sizeable fraction of the SN explosion energy can be put into very energetic CRs that escape the forward shock and stream into the surrounding ISM. These particles will produce detectable radiation if they interact with dense, external material.  Another advantage is that we have a ``coherent'' model, easily expandable to include more complex effects, where the various environmental and theoretical parameters can be straightforwardly varied and the resulting radiation can be compared directly with observations. This is important because all SNe and SNRs are different and complex with many poorly constrained parameters. It is essential that the underlying theory consistently model broad-band emission from radio to TeV $\\gamma$-rays taking into account individual characteristics of the remnants and their environments.  In Sections~\\ref{section:DSA} and \\ref{section:Diff_Interaction} we give a brief general description of nonlinear diffusive shock acceleration and describe the environmental and model parameters required for a hydrodynamical solution. We place a time-dependent, spherically symmetric, hydrodynamic calculation of a SNR, including NL-DSA, in a three-dimensional box consisting of $51 \\times 51 \\times 51$ cells.\\footnote{The resolution of the 3-D box is, of course, adjustable and limited only by computational considerations.} The energetic particles produced by the outer blast wave shock propagate through the simulation box where they interact with an arbitrary distribution of matter placed external to the outer shock. The energetic particles in the box, including those within the SNR, suffer energy losses and produce broad-band continuum emission spectra by interacting with the magnetic field, photon field, and matter density of each cell. In Section~\\ref{section:results} we show some examples including line-of-sight projections of the emitted radiation which are suitable for comparison with observations.  There are a number of young SNRs under active investigation, including SNR RX~J1713 \\citep[e.g.,][]{AharonianJ1713_2006,Uchiyama_J1713_2007}, Vela Jr. \\citep[e.g.,][]{AharonianVela2005}, RCW~86 \\citep[][]{Hoppe_etal_RCW86_2007,Ueno_etal_RCW_2007,RhoEtal2002}, IC~443 \\citep[][]{Albert_etal_IC443,Humensky_etalIC443} and W~28 \\citep[][]{AharonianW28}. However, here we concentrate on a general study using various parameters typical of young, shell Type Ia SNRs and leave detailed modeling of individual remnants for future work. ",
        "conclusions": ""
    },
    "0806/0806.4088_arXiv.txt": {
        "abstract": "{Our Galaxy hosts at its dynamical center Sgr\\,A*, the closest supermassive black hole. Surprisingly, its luminosity is several orders of magnitude lower than the Eddington luminosity. However, the recent observations of occasional rapid X-ray flares from Sgr\\,A* provide constraints on the accretion and radiation mechanisms at work close to its event horizon. } {Our aim is to investigate the flaring activity of Sgr\\,A* and to constrain the physical properties of the X-ray flares.} {In Spring 2007, we observed Sgr\\,A* with XMM-Newton with a total exposure of $\\sim$230\\,ks. We have performed timing and spectral analysis of the new X-ray flares detected during this campaign. To study the range of flare spectral properties, in a consistent manner, we have also reprocessed, using the same analysis procedure and the latest calibration, archived XMM-Newton data of previously reported rapid flares. The dust scattering was taken into account during the spectral fitting. We also used Chandra archived observations of the quiescent state of Sgr\\,A* for comparison. } { On April 4, 2007, we observed for the first time within a time interval of roughly half a day, an enhanced incidence rate of X-ray flaring, with a bright flare followed by three flares of more moderate amplitude. The former event represents the second brightest X-ray flare from Sgr A* on record. This new bright flare exhibits similar light-curve shape (nearly symmetrical), duration ($\\sim$3\\,ks) and spectral characteristics to the very bright flare observed in October 3, 2002 by XMM-Newton. The measured spectral parameters of the new bright flare, assuming an absorbed power law model taken into account dust scattering effect,  are $N_{\\rm H}$=12.3$^{+2.1}_{-1.8}\\times10^{22}$\\,cm$^{-2}$ and $\\Gamma\\sim$2.3$\\pm$0.3 calculated at the 90\\% confidence level. The spectral parameter fits of the sum of the three following moderate flares, while lower ($N_{\\rm H}$=8.8$^{+4.4}_{-3.2}\\times10^{22}$\\,cm$^{-2}$ and $\\Gamma\\sim$1.7$^{+0.7}_{-0.6}$), are compatible within the error bars with those of the bright flares. The column density found, for a power-law model taking into account the dust scattering, during the flares is at least two times higher than the value expected from the (dust) visual extinction toward Sgr\\,A* ($A_{\\rm V}\\sim25$ mag), i.e., 4.5$\\times10^{22}$\\,cm$^{-2}$. However, our fitting of the Sgr\\,A* quiescent spectra obtained with Chandra, for a power-law model taking into account the dust scattering, shows that an excess of column density is already present during the non-flaring phase. } { The two brightest X-ray flares observed so far from Sgr\\,A* exhibited similar soft spectra. } ",
        "introduction": "Located at the center of our Galaxy, \\object{Sgr A*} is the closest supermassive black hole to the solar system at a distance of about 8\\,kpc \\citep{Reid93,Eisenhauer03,Eisenhauer05}. Its mass of about 3--4$\\times$10$^{6}$\\,$M_{\\odot}$ has been determined thanks to the measurements of star motions \\citep[e.g.,][]{Schoedel02,Ghez03,Ghez05}. Amazingly, this source is much fainter than expected from accretion onto a supermassive black hole. Its bolometric luminosity is only about 3$\\times$10$^{-9}$ L$_{\\rm Edd}$ \\citep{Melia01,Zhao03}. In particular, its 2--10\\,keV X-ray luminosity is  only about 2.4$\\times$10$^{33}$\\,erg\\,s$^{-1}$ within a radius of 1.5$^{\\prime\\prime}$ \\citep{Baganoff03}. Thus, Sgr\\,A* radiates in X-rays at about 11 orders of magnitude less than its corresponding Eddington luminosity. This has motivated the development of various radiatively inefficient accretion models to explain the dimness of the Galactic Center black hole, e.g., Advection-Dominated Accretion Flows (e.g., \\citealt{Narayan98}), jet-disk models (e.g., \\citealt{Falcke00}), Bondi-Hoyle with inner Keplerian flows (e.g., \\citealt{Melia00}). The recent discovery of X-ray flares from Sgr\\,A* has provided new exciting perspectives for the understanding of the processes at work in the Galactic nucleus.  The first detection of such events was found with {\\sl Chandra} in October 2000. This flare had a duration of about 10\\,ks, with a flare peak luminosity of about 1.0 $\\times$ 10$^{35}$\\,erg\\,s$^{-1}$, in the 2--10\\,keV energy range, i.e., about 45 times the quiescent state \\citep{Baganoff01}. Further on several other X-ray flares were detected by {\\sl XMM-Newton} \\citep{Goldwurm03,P03c,Belanger05} and {\\sl Chandra} \\citep{Baganoff03b,Eckart04,Eckart06,Eckart08,Marrone08}. The majority of X-ray flares detected up to now have moderate flux amplitude with factor of about 10--45 compared to the quiescent state. Only one very bright flare with a flux amplitude of about 160 was observed in October 2002 \\citep{P03c}. Remarkably, its peak luminosity of $\\sim$3.6$\\times$10$^{35}$\\,erg\\,s$^{-1}$ was comparable to the bolometric luminosity of Sgr\\,A* during its quiescent state. The light curve of the X-ray flares can exhibit short (e.g., 600\\,s, \\citealt{Baganoff01}; 200\\,s, \\citealt{P03c}) but deep drops close to the flare maximum. This short-time scale could indicate that the X-ray emission is emitted from a region as small as 7$R_{\\rm S}$ ($\\sim$ 13\\,$R_{\\odot}$).  In contrast to the near-IR (NIR) flares which appear to be present for up to 40\\% of the time (e.g., \\citealt{Genzel03,Ghez04,Yusef06}, Yusef-Zadeh et al.\\ 2008 in prep.), the X-ray flaring has a much lower duty cycle, of typical 1-5\\% \\citep{Baganoff03b,Belanger05,Eckart06}. This means either that the majority of NIR flares have no X-ray counterpart or that the X-ray-to-NIR ratio of some of the flares is too small to allow an X-ray detection above the strong, diffuse X-ray emission from the central parsec. When both NIR and X-ray flares are detected simultaneously, they show similar morphology in their light curves as well as no apparent delay between the peaks of flare emission \\citep{Eckart04,Eckart06,Eckart08,Yusef06}. The current interpretation is that both flares come from the same region.  We report here the results of our Sgr\\,A* observation campaign performed with {\\sl XMM-Newton} from March 30 to April 4, 2007. The whole results of this multi-wavelength campaign (VLA, CSO, VLT/NACO-VISIR, HST/NICMOS, Integral) will be reported elsewhere (Yusef-Zadeh et al.\\ 2008 in prep.; Dodds-Eden et al.\\ 2008, in prep.). We also observed during the three {\\sl XMM-Newton} observations, two bright transient sources in outburst \\citep{P07b}. The source located at about 90$^{\\prime\\prime}$--SW from Sgr\\,A* has been  associated with the eclipsing X-ray burster \\object{AX J1745.6-2901}. Seven deep eclipses were observed as well as type-I bursts. The second source is located at about 10$^{\\prime}$--NW from Sgr\\,A* and has been associated with the neutron star low-mass X-ray binary \\object{GRS\\,1741.9-2853} (a.k.a.\\ \\object{AX J1745.0-2855}). The data analysis of these two sources will be reported in forthcoming papers.  We report, here, for the first time a high level of flaring activity with four X-ray flares, one bright and three moderate, detected in half a day. The bright flare is the second brightest X-ray flare detected so far from Sgr\\,A*.  In $\\S$\\ref{sec:data} we describe the observations and data reduction procedure used in this work. In $\\S$\\ref{sec:results} we report the timing analysis of Sgr\\,A*, and the spectral analysis of the bright flare and the sum of the three following moderate flares observed during this Spring 2007 campaign. In $\\S$\\ref{sec:comparison}, we perform a homogeneous and self-consistent comparison of spectral properties of the X-ray flares of this campaign with the reprocessed data of X-ray flares previously observed with {\\sl XMM-Newton}. Finally, in $\\S$\\ref{sec:discussion} we summarize our main results and discuss their possible implications. ",
        "conclusions": "\\label{sec:discussion}  Here, we reported the data analysis of the {\\sl XMM-Newton} April 2007 campaign observation of Sgr\\,A* (three observations with a total of exposure of $\\sim$230\\,ks). We observed four X-ray flares occurring within only half a day on April 4: one bright flare ($\\#2$, with a detection level of $\\sim$21$\\sigma$) followed by three moderate ones ($\\#3$, $\\#4$, and $\\#5$, with a detection level of $\\sim$6$\\sigma$, $\\sim$6$\\sigma$, and $\\sim$8$\\sigma$, respectively). As well a weak flare ($\\#1$, with a detection level of $\\sim$3$\\sigma$) was observed on April 2, 2007. The flare $\\#2$ is the second brightest X-ray flare detected so far from Sgr\\,A* and has a similar duration ($\\sim$3\\,ks) to that of the October 2002 brightest flare \\citep{P03c}. Its light curve is almost symmetrical but without no significant short-time scale drop (i.e., about 50\\% of flux decrease) contrary to that reported for the October 2000 flare \\citep{Baganoff01} and the October 2002 flare \\citep{P03c}. However, for the flare $\\#$2 we cannot rule out a moderate drop in the X-ray light curve. A similar group of three moderate flares were already observed with {\\sl XMM-Newton} on 2004 March 31 \\citep{Belanger05}, but no such preceding bright flare was observed.  This is the first time that a such level of X-ray flaring activity from Sgr\\,A*, both in amplitude and frequency, is reported. Observations such as those reported in this paper can eventually constrain flare models in Sgr A*, perhaps even ruling some out. For example, the quick succession of several events separated by only a few hours, might argue against a disruption mechanism that relies on the temporary storage of mass and energy, if all the corresponding accretion energy is released during the flare. The accretion rate $\\dot M$ in this system (some estimates place it as low as $10^{16}-10^{17}$ g s$^{-1}$; see \\citealt{Melia07}) might not produce a transient accumulation of mass $\\Delta M$ between flares of sufficient magnitude for $\\eta GM\\Delta M/ 3R_S$ to provide the observed outburst power. In this expression, $R_S$ is the Schwarzschild radius, and $\\eta$ is the (poorly known) radiative efficiency in Sgr A*, believed to be at most a few percent. If instead the flares are due to a magneto-rotational instability, then the energy liberated during the flare must also be accumulated over the short inter-burst period. The low $\\dot M$, from which the energy is derived, might argue against this type of mechanism as well.  On the other hand, if the flare arises from the infall of a clump of gas \\citep[e.g.,][]{Liu06,Tagger06} then there would be less restriction on how often these could come in. \\cite{Genzel03} have shown that the total energy release $\\geq$ 10$^{39.5}$ erg during a flare requires a gas accreted mass of a few times 10$^{19}$\\,g (assuming a radiation efficiency of $\\sim$10$\\%$), i.e.\\ comparable to that of a comet or a small asteroid. Recently, \\cite{Cadez06} have argued that the flares could be produced by tidal captures and disruptions of such small bodies. The comet/asteroid/planetesimal idea for depositing the additional mass and energy to initiate a flare is attractive. The distance from Sgr A* at which such a small body would get tidally disrupted is the Roche radius, which is for a rigid body: \\[ \\frac{R_{\\cal R}}{R_{\\rm S}} = 13.2 \\times \\left(\\frac{M_{\\rm BH}}{4 \\times 10^6 M_\\odot}\\right)^{-2/3} \\times \\left(\\frac{\\rho_p}{{\\rm 1\\,g\\,cm^{-3}}}\\right)^{-1/3}\\, , \\] where $M_{\\rm BH}$ is the black hole mass and $\\rho_p$ is the density of the rigid body. Thus, for a black-hole mass of $4\\times 10^6\\;M_\\odot$ and a density of 1\\,g\\,cm$^{-3}$, this corresponds to $13.2\\,R_{\\rm S}$, in good agreement with the size of the region where the flares are thought to occur. The flaring rates would then depend on processes occurring much farther out, so the fact that so many flares are seen so close together on some days, and much less frequently at other times, would simply be due to stochastic events. However, it would still be difficult to distinguish between a compact emission region and emission within a jet (see, e.g., \\citealt{Markoff01}), since these disruption events could still end up producing an ejection of plasma associated with the flare itself.\\\\   Four of the five X-ray flares observed during this campaign have a simultaneous NIR flare counterpart: flares $\\#1$, $\\#4$, and $\\#5$ were observed with HST/NICMOS (Yusef-Zadeh et al.\\ 2008, in prep.); and flare $\\#2$ was observed with VLT/NACO (Dodds-Eden et al.\\ 2008, in prep.). The flare $\\#3$ has not been covered by a simultaneous NIR observation. This strengthens the relationship observed up to now between X-ray and NIR flares when there is a simultaneous X-ray/NIR observation coverage: all X-ray flares have an NIR flare counterpart, while all NIR flares are not each time associated with an X-ray flare counterpart \\citep[e.g.,][]{Eckart04,Eckart06,Yusef06,Hornstein07,Marrone08}.   We have made the first detailed comparison of X-ray flare properties observed with {\\sl XMM-Newton} (this April 2007 campaign and previous observations) based on a fully self-consistent analysis approach, namely we use the same SAS version with the latest calibration, the same definition of the flare interval (preventing from significant contribution of the non-flaring level), an optimized extraction region for each instrument, and we took into account the dust scattering effect. We found that the physical parameters of the flare $\\#2$, $N_{\\rm H}$=12.3$^{+2.1}_{-1.8}$$\\times$ 10$^{22}$\\,cm$^{-2}$ and $\\Gamma$=2.3$\\pm$0.3, are well constrained and are very similar to that of the October 2002 (brightest) flare, $N_{\\rm H}$=12.3$^{+1.6}_{-1.5}$$\\times$10$^{22}$\\,cm$^{-2}$ and $\\Gamma$=2.2$\\pm$0.3. The spectral parameter fit values of the sum of the three moderate flares following the bright flare observed on April 4, 2007, $N_{\\rm H}$=8.8$^{+4.1}_{-3.2}$$\\times$ 10$^{22}$\\,cm$^{-2}$ and $\\Gamma$=1.7$^{+0.7}_{-0.6}$, while lower, are compatible within the error bars with those of the bright flares. The column density found during the bright flares is at least two times higher than the value expected from the (dust) visual extinction toward Sgr\\,A* ($A_{\\rm V}\\sim25$ mag; \\citealt{Eisenhauer05}), i.e., 4.5$\\times10^{22}$\\,cm$^{-2}$. However, our fitting of the Sgr\\,A* quiescent spectra obtained using four {\\sl Chandra} observations for a total observation time of $\\sim$ 230\\,ks, taking into account the dust scattering, shows that a column density excess is already present during the non-flaring phase. One possible explanation, as already pinpoint by \\cite{Maeda02}, is that such column density excess in the line-of-sight of Sgr\\,A* could be due to the warm (dust free) gas associated with the ionized gas halo, \\object{Sgr\\,A West} extended, suggested by the turnover absorption observed at 90\\,cm, embedded the Sgr\\,A complex \\citep[e.g.,][]{Pedlar89,Anantharamaiah99,Yusef00}. \\cite{Anantharamaiah99} inferred that this ionized halo has a dimension of about 4$\\arcmin$ ($\\sim$9pc) and an electron density of about 100--1000\\,cm$^{-3}$. For a column excess of about 2.9$\\times$10$^{22}$\\,cm$^{-2}$ as found for the Sgr\\,A* quiescent state, this would correspond to a consistent electron density of about 1300\\,cm$^{-3}$. However, in case the genuine continuum shape is curved compared to a power law and bremsstrahlung models, as for example a black body like shape, no significant column excess would be required.  On the basis of our current analysis we conclude that the dichotomy of the spectral index between moderate and bright flares (i.e., a correlation between photon index and flux) suggested by \\cite{Belanger05} is not confirmed. However, to test this more thoroughly we need further observations so as to extend the sampling of the flaring activity seen from Sgr A* (and in particular those instances where good signal to noise can be achieved through the co-adding of weak to moderate flaring events).   In conclusion, this study establishes that the two brightest X-ray flares observed so far from Sgr\\,A* exhibited similar (well constrained) soft spectra. Therefore, any model proposed to explain the flaring behavior of Sgr\\,A* must take into account this observational X-ray spectral property."
    },
    "0806/0806.1387_arXiv.txt": {
        "abstract": "{} {To investigate the effect of gravitational lensing of supernovae in large ongoing surveys.} {We simulate the effect of gravitational lensing magnification on individual supernovae using observational data input from two large supernova surveys. To estimate the magnification due to matter in the foreground, we simulate galaxy catalogs and compute the magnification along individual lines of sight using the multiple lens plane algorithm. The dark matter haloes of the galaxies are modelled as gravitational lenses using singular isothermal sphere or Navarro-Frenk-White profiles. Scaling laws between luminosity and mass, provided by Faber-Jackson and Tully-Fisher relations, are used to estimate the masses of the haloes. } { While our simulations show that the SDSSII supernova survey is marginally affected by gravitational lensing, we find that the effect will be measurable in the SNLS survey that probes higher redshifts. Our simulations show that the probability to measure a significant ($3\\sigma$) correlation between the Hubble diagram residuals and the calculated lensing magnification is $\\ga 95\\%$ in the SNLS data. Moreover, with this data it should be possible to constrain the normalisation of the masses of the lensing galaxy haloes at the $1\\sigma$ and $2\\sigma$ confidence level with $\\sim 30\\%$ and $\\sim 60\\%$ accuracy, respectively. } {} ",
        "introduction": "Type Ia supernovae (hereafter SNe~Ia) are exceptionally useful tools for cosmological investigations. Ongoing large supernova surveys, such as ESSENCE \\citep{miknaitis07,wv07}, SNLS \\citep{astier06} and SDSSII \\citep{frieman07} harvest hundreds of distant supernovae with high precision.  Several effects will alter the luminosity of these distant sources, and thus dilute the cosmological signal. In this paper we investigate gravitational lensing (de)magnification of the light from these explosions. Although this effect is hardly large enough to dominate the uncertainties in current experiments, as has already been addressed by several groups on statistical grounds \\citep[e.g.,][]{riess04,hl05,astier06,wv07}, the question if it can be corrected for on an individual supernova basis remains. This has been investigated in a series of papers \\citep{gunnarsson06,jonsson06} and a tentative detection of a correlation between calculated lensing magnification and Hubble residuals was recently found using the very high-$z$ supernovae in the GOODS field \\citep{jonsson07}.  In this paper we investigate to what extent other surveys can be affected. The GOODS supernovae \\citep{riess04,strolger04,riess07} remain unique in that their large distances clearly make them more likely to be lensed. On the other hand, the ongoing ground based surveys will measure significantly more supernovae with much better sampling, and the improved statistics may well compensate for the smaller distances. To find out which surveys are most likely to display a lensing signal we have performed simulations based on our previous work. The results are presented below, and show that the prospects of detecting lensing are very good.  In Sect.~\\ref{surveys} we briefly introduce the surveys we have investigated and the data from these surveys that are used as input for the simulations. A short summary of gravitational lensing of SNe~Ia is given in Sect.~\\ref{lensing}. Section~\\ref{simulations} describes the simulations and the results are presented in Sect.~\\ref{results}, where we also discuss how a detection of a lensing signal could be used to obtain information about the lensing matter. Finally, we summarise our results in Sect.~\\ref{discussion}. ",
        "conclusions": "\\label{discussion} We have simulated the effect of gravitational lensing in two major ongoing supernova surveys. For the relatively nearby SDSSII supernova search, the effect of gravitational lensing will be small. For the more distant supernovae in the SNLS survey, we predict that the signal from gravitational lensing will be observed with high confidence. Our simulations indicate that a correlation between Hubble diagram residuals and magnification for individual supernovae will be present at high (at least $3\\sigma$) significance level. This could be used both to somewhat reduce the scatter in the Hubble diagram, and to learn about the properties of the lensing material. A project to investigate this effect in the SNLS data is underway. We also note that the prospects of using weak lensing of supernovae to constrain the matter distribution in the Universe with future surveys such as SNAP look very promising. A satellite like SNAP would provide not only high redshift SNe~Ia, but also deep observations in many filters which would allow reliable photometric redshifts of the galaxies to be obtained."
    },
    "0806/0806.4277_arXiv.txt": {
        "abstract": "We present the formulation of a new infinite family of self-consistent stellar models, designed to describe axisymmetric flat galaxies. The corresponding density-potential pair is obtained as a superposition of members belonging to the generalized Kalnajs family, by imposing the condition that the density can be expressed as a regular function of the gravitational potential, in order to derive analytically the corresponding equilibrium distribution functions (DF). The resulting models are characterized by a well-behaved surface density, as in the case of generalized Kalnajs discs. Then, we present a study of the kinematical behavior which reveals, in some particular cases, a very satisfactory behavior of the rotational curves (without the assumption of a dark matter halo). We also analyze the equatorial orbit's stability and Poincar\\'e surfaces of section are performed for the 3-dimensional orbits. Finally, we obtain the corresponding equilibrium DFs, using the approaches introduced by \\cite{kal} and \\cite{dej}. ",
        "introduction": "The obtention of density-potential pairs (PDP) corresponding to idealized thin discs is a problem of great astrophysical relevance, motivated by the fact that the main part of the mass in many galaxies is concentrated in an stellar flat distribution, usually assumed as axisymmetric (\\cite{BT}). Once the potential-density pair (PDP) is formulated as a model for a galaxy, usually the next step is to find the corresponding distribution function (DF). This is one of the fundamental quantities in galactic dynamics specifying the distribution of the stars in the phase-space of positions and velocities. Although the DF can generally not be measured directly, there are some observationally accesible quantities that are closed related to the DF: the projected density and the light-of-sight velocity, provided by photometric and kinematic observations, are examples of DF's moments. Thus, the formulation of a PDP with its corresponding equilibrium DFs establish a self-consistent stellar model that can be corroborated by astronomical observations.  Now then, there is a variety potential-density pairs for such flat stellar models, e.g. \\cite{WM}; \\cite{KUZ}; \\cite{Schmith}; \\cite{T1,T2}; \\cite{BB}; \\cite{KAL}; \\cite{GR}. In particular, \\cite{T1,T2} formulated a generalized family of models whose first member is the one introduced by \\cite{KUZ}. This family represents a set of discs of infinite extension, derived by solving the Laplace equation  in cylindrical coordinates subject to appropriated boundary conditions on the discs and at infinity.  Analogously, \\cite{GR} obtained a family of finite thin discs (generalized Kalnajs discs) whose first member corresponds precisely to the well-known model derived by \\cite{KAL}. Such family was derived by using the Hunter's method (\\citeyear{HUN1}), which is based in the obtention of solutions of Laplace equation in terms of oblate spheroidal coordinates, by imposing some appropriate conditions on the surface density. So, by requiring that the surface density behaves as a monotonously decreasing function of the radius, with a maximum at the center of the disk and vanishing at the edge, detailed expressions for the gravitational potential and the rotational velocity were obtained as series of elementary functions. Also, some two-integral distribution functions for the first four members of this family were recently obtained by \\cite{PRG}. Now then, as the generalized Kalnajs models correspond to discs of finite extension, they can be considered as more realistic descriptions of flat galaxies than the Toomre's family.  In the present paper we formulate a new infinite set of finite thin discs, obtained by superposing the members of the generalized Kalnajs family in such a way that the resulting density surface can be expressed as a well behaved function of the gravitational potential. As it was pointed out by some authors, this is a fundamental requirement for the searching of equilibrium distribution functions (DF) describing such axisymmetric systems (see for example, \\cite{fricke}, \\cite{Hunter}, \\cite{jiang}). Thus, the new family formulated here has the advantage of easily providing the corresponding two-integral DFs.  Furthermore, the models have two additional advantages. In one hand, the mass surface density is well-behaved, as in the case of generalized Kalnajs discs, having a maximum at the center and vanishing at the edge. Moreover, the mass distribution of the higher members of the family is more concentrated  at the center. On the other hand, the rotation curves are better behaved than in the Kalnajs discs. We found that, in some cases, the circular velocity increases from a value of zero at the center of the disc, then reaches a maximum at some critical radius and, after that, remains approximately constant. As it is known, such behavior has been observed in  many disclike  galaxies.  Now, apart from the circular velocity, there are two important quantities concerning to the interior kinematics of the models: the epicyclic and vertical frequencies, which describe the stability against radial and vertical perturbations of particles in quasi-circular orbits. We found that the models formulated here are radially stable whereas vertically unstable, which is a characteristic inherited from the generalized Kalnajs family (\\cite{ramos}). When we deal with three dimensional orbits, there are also a common feature between the new models and the Kalnajs family: the phase space structure of disc-crossing orbits, that can be viewed through the Poincar\\'e surfaces of section, is composed by shape-of-ring KAM curves and prominent chaotic zones. However, there are certain situations in which the chaoticity disappears and the 3-dimensional motion of test particles is completely regular. In such cases, one can suggest the existence of a third integral of motion, as in the case of St\\\"{a}ckel and Kuzmin potentials.  Finally, in order to formulate the new family as a set of self-consistent stellar models, we shall deal with the problem of obtaining the corresponding equilibrium DFs. By Jeans's theorem, they are functions of the isolating integrals of motion that are conserved in each orbit. Some authors have shown that, if the density can be written as a function of the gravitational potential, it is possible to find such kind of two-integral DFs (see  Eddington (1916); Fricke (1952); \\cite{kal}; Jiang and Ossipkov (2007)). In this paper we shall adopt the approach introduced by \\cite{kal}, which fits quite well to axisymmetric disc-like systems. Then, starting from the DFs derived from this method, a second kind of DFs is obtained by using the formulae introduced by \\cite{dej}, that takes into account the principle of maximum entropy. These DFs describe stellar systems with a preferred rotational state.  Accordingly, the paper is organized as follows. First, at section \\ref{sec:newfam}, we obtain the potential-density pairs for the new family of thin disc models. Then, at section \\ref{sec:kin}, we study the motion of test particles around these new galactic models and, at section \\ref{sec:DFs}, we derive the distribution functions associated with the models. Finally, at section \\ref{sec:conc}, we summarize our main results. ",
        "conclusions": "\\label{sec:conc}  We have obtained a set of models for axisymmetric flat galaxies, by superposing members belonging to the generalized Kalnajs discs family. The mass distribution of each model (labeled through the parameter $m=2,3,\\ldots$), described by (\\ref{dennew}), is maximum at the center and vanishes at the edge, in concordance with a great variety of galaxies. Moreover, the mass density can be expressed as a function of the gravitational potential (see equation (\\ref{dennew2})), which makes possible to derive, analytically, the equilibrium DFs describing the statistical features of the models.  These models have also interesting features concerning with the interior kinematical behavior. On one hand, we showed that for some values of $B_{1}$, the circular velocity has a behavior very similar to that seen in many discoidal galaxies. This is a very relevant fact, which suggests that it is not always necessary to introduce the hypothesis of dark matter halos (or MOND theories) in order to describe adequately a variety of rotational curves. On the other hand, the analysis of epicyclic and vertical frequencies, associated to quasi-circular orbits, reveals that the models are stable under radial perturbations but unstable under vertical disturbances. With regard to the motion of test particles around the models formulated here, we found that the behavior of disc-crossing orbits is similar to that seen in the generalized Kalnajs family. However, for certain values of the parameter $B_{1}$, the Poincar\\'e surface of section reveals that one can suggest the existence of a (non analytical) third integral of motion.  On the other hand, we find two kinds of equilibrium DFs for the models. Such two-integral DFs can be formulated, at first, as functionals of the Jacobi's integral, as it was sketched in the formalism developed by \\cite{kal}. This class of DFs essentially describes systems which rotational state, in average, behaves as a rigid body. Then, we use the procedure introduced by \\cite{dej}, obtaining DFs which represents systems with a mean rotational state consistent with the maximum entropy principle and, therefore, more probable than the first ones. The statements exposed above suggest that the family presented here, can be considered as a set of realistic models that describes satisfactorily a great variety of galaxies."
    },
    "0806/0806.3382_arXiv.txt": {
        "abstract": "\\medskip \\noindent We show that a contribution to the total curvature perturbation may be due  to the presence of flat directions in supersymmetric models. It is generated at the first oscillation of the flat direction condensate when the latter relaxes to the minimum of its potential after the end of inflation. We also point out that, if the contribution to the total curvature perturbation from supersymmetric flat direction is the dominant one, then a significant level of non-Gaussianity in the cosmological perturbation is also naturally expected. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0806/0806.0662_arXiv.txt": {
        "abstract": "We compare the cosmic evolution of star formation rates in galaxies with that of their neutral hydrogen densities. We highlight the need for neutral hydrogen to be continually replenished from a reservoir of ionized gas to maintain the observed star formation rates in galaxies. Hydrodynamic simulations indicate that the replenishment may occur naturally through gas infall, although measured rates of gas infall in nearby galaxies are insufficient to match consumption. We identify an alternative mechanism for this replenishment, associated with expanding supershells within galaxies. Pre-existing ionized gas can cool and recombine efficiently in the walls of supershells, molecular gas can form {\\em in situ\\/} in shell walls, and shells can compress pre-existing molecular clouds to trigger collapse and star formation. We show that this mechanism provides replenishment rates sufficient to maintain both the observed HI mass density and the inferred molecular gas mass density over the redshift range $0\\le z\\lesssim 5$. ",
        "introduction": "\\label{int}  Our understanding of the cosmic star formation history (SFH) of galaxies has progressed significantly over the past decade \\citep[e.g.,][]{Hop:04,HB:06}. In the same time the space density of neutral hydrogen gas has been measured over the majority of cosmic history \\citep[see Figure~8 of][]{Lah:07}. The evolution of the atomic hydrogen (HI) in the universe will be comprehensively determined within the next few years by extremely sensitive surveys with the next generation of radio telescope instrumentation \\citep[e.g.,][]{vdH:04,Raw:04,Joh:08}, and it is timely to consider mechanisms associated with this evolution.  The space density of HI in galaxies appears to evolve surprisingly little from $z\\approx 5$ to $z\\approx 0.2$ \\citep{Lah:07}, a span of roughly 10\\,Gyr, the latter half of which sees a decline in the space density of star formation rate (SFR) in galaxies by almost an order of magnitude \\citep[e.g.,][]{HB:06}. Given the SFR density it is easy to show that the HI plus molecular gas at high redshift would be exhausted on timescales of a few Gyr if it were not continually replenished. \\citet{Erb:08} presents a model incorporating gas infall, outflows and consumption by star formation, to explain both replenishment and the mass-metallicity relation in high-redshift ($z\\approx 2$) galaxies.  Hydrodynamic simulations advocating hot and cold modes of accretion indicate that the infall rate closely tracks the star formation rate \\citep[e.g.][]{Ker:05,Bir:07}, with star formation moderated by the rate of infall. The simulations, however, neglect gas outflows from galaxies, which are a significant component of gas depletion. The quantitative infall rates predicted are thus insufficient to maintain a constant HI density in galaxies. Observed rates of gas infall in local galaxies, also, are only about 10\\% of the star formation rate \\citep{San:08}. The difficulties in explaining replenishment through infall leave the physical mechanism of this replenishment as a critical open question in galaxy evolution.  In this Letter we suggest a mechanism directly associated with the SFR in galaxies that can provide the necessary replenishment of neutral gas to maintain an essentially unevolving, or slowly evolving, HI mass density. We infer the density of gas required to reproduce the observed SFH in \\S\\,\\ref{data}. In \\S\\,\\ref{deltarho} we present a number of models for the replenishment of this gas, and show that a replenishment proportional to the SFR density can reproduce the necessary gas mass density. A replenishment mechanism associated with galactic supershells is detailed in \\S\\,\\ref{disc}, and the results are summarised in \\S\\,\\ref{summ}. Throughout this analysis we adopt the ``737\"\\footnote{Thanks to Sandhya Rao \\citep{Rao:06} for this terminology.} cosmology with $H_0=70\\,$km\\,s$^{-1}$\\,Mpc$^{-1}$, $\\Omega_M=0.3$, $\\Omega_{\\Lambda}=0.7$ \\citep[e.g.,][]{Spe:03}. ",
        "conclusions": "\\label{disc} \\subsection{A physical mechanism for replenishment} \\label{physmech} The ISM in galaxies contains expanding supershells or superbubbles associated with previous generations of star formation. We propose that the neutral and molecular gas replenishment in the walls of supershells is sufficient, and of the appropriate form, to provide a natural mechanism explaining the relatively flat evolution in the HI mass density.  Supershells are large scale expanding shells of gas driven by supernovae (SNe) and stellar winds from OB star clusters \\citep{Oey:96,OS:98,McCG:01}. Supershells have long been suggested to have a triggering effect on subsequent generations of star formation \\citep{McC:87,Elm:98,Har:01,Ber:04,Oey:05}, and are effective at replenishing star-forming gas through several mechanisms. First, supershells are efficient at cooling and recombining ionized gas through radiative cooling in shell walls \\citep{Koo:92}. This may be critical in converting gas shock-heated by previous generations of SNe within a galaxy, or new hot-mode infall gas, to a potentially star-forming state, as the $10^6\\,$K gas may otherwise never cool to support subsequent star formation. Second, molecular gas can form from neutral gas {\\em in situ\\/} in shell walls, where compression and the development of instabilities leads to sufficiently high neutral gas densities to allow for cooling and self-shielding on timescales of $10 - 20\\,$Myr \\citep{Ber:04,Hen:08}. Finally, they can compress pre-existing molecular material to trigger molecular cloud collapse and star formation \\citep{Elm:98}. The timescales for these processes are shorter than, or comparable to the supershell lifetime ($\\sim20\\,$Myr), which is in turn short compared to the global gas consumption timescale of several Gyr.  To establish whether the replenishment achievable in supershells is sufficient to make this mechanism feasible, we first convert the replenishment rates given in \\S\\,\\ref{deltarho} into a replenished mass per SN event. A replenishment rate proportional to the SFR density is also proportional to the rate of supernova type II (SNII\\footnote{Here and throughout we assume the inclusion of all core-collapse supernovae: types II, Ib, and Ic}). Converting a proportionality to SFR density into one depending on the SNII rate, $\\dot{\\rho}_{\\rm SNII}$, depends on the assumed IMF. From \\citet{HB:06} $\\dot{\\rho}_{\\rm SNII}=(0.00915/M_{\\odot})\\,\\dot{\\rho}_*$ for the SalA IMF. The replenishment rate $K(t)=1.6\\dot{\\rho}_*$ becomes $K(t)=174.9\\dot{\\rho}_{\\rm SNII}\\,M_{\\odot}$.  The other extreme choice of IMF consistent with the normalization of the SFH \\citep{HB:06} is that of \\citet[hereafter the BG IMF]{Bal:03}. For the BG IMF $\\dot{\\rho}_{\\rm SNII}=(0.0132/M_{\\odot})\\,\\dot{\\rho}_*$. The recycled fraction is $R=0.56$ \\citep{HB:06}, changing the consumption term in Equation~(\\ref{theeqn}) to $-1.44\\dot{\\rho}_*$. The corresponding replenishment rate is $K(t)=109.1\\dot{\\rho}_{\\rm SNII}\\,M_{\\odot}$. These extremes imply that, depending on the IMF, sufficient gas replenishment to maintain a constant HI mass density with redshift would be achieved if each SN event caused the recombination and cooling of $\\approx 110-180\\,M_{\\odot}$ of gas. These IMFs are the extrema given the SFH normalization limits, and most reasonable IMFs should result in masses within this range.  Detailed measurements to confirm molecular gas formation within supershells are observationally challenging. We use the limited data currently available to assess the replenishment rates associated with supershells, and to establish whether at least one well-studied supershell achieves the required rate.  \\citet{McCG:05} and \\citet{Daw:08} have shown explicit cases of molecular clumps along the edges of supershells, suggestive of some degree of {\\em in situ\\/} formation, with a significant amount of molecular material associated with the supershell walls. The supershell investigated by \\citet{Daw:08} is associated with about $2\\times10^5\\,M_{\\odot}$ of molecular gas, of which those authors estimate that $80\\%$ likely comes from a pre-existing giant molecular cloud. Of the remaining $\\approx 4\\times10^4\\,M_{\\odot}$ of molecular gas it is difficult to determine how much is pre-existing and how much has been cooled and recombined by the expansion of the shell. We can use $\\approx 4\\times10^4\\,M_{\\odot}$ as an upper limit to the replenishment rate. About 30 stars with stellar mass $M_*> 7\\,M_{\\odot}$ are required to form this supershell, including stars that may not yet have gone supernova. This gives $\\lesssim 1300 - 2000\\,M_{\\odot}$ of molecular mass replenished per SN event, a limit comfortably encompassing the required rate. This upper limit could change significantly depending on the fraction of pre-existing molecular material and also on the fraction of stars that have not yet gone supernova.  Not all SNe lie within supershells, although \\citet{HL:05} estimate that a minimum of $65\\%$ of SNII should occur in superbubbles, increasing to $\\approx 80-90\\%$ when the spatial and temporal correlations of stellar clusters are considered. If $80\\%$ of SNII are associated with supershells, for example, this would increase the required replenishment rate per SN to $\\approx 140-230\\,M_{\\odot}$. But even if as few as 10\\% of all SNII contribute in this way to the replenishment, the rate implied by the results of \\citet{Daw:08} would still be sufficient. This confirms that the necessary replenishment rates are likely to be achievable within supershells.  The observed decline by a factor of two in the HI mass density may be a natural consequence of a replenishment rate about 95\\% of that required to match consumption, as shown by the heavy solid line in Figure~\\ref{fig:gasmodels}. If the actual replenishment rate from supershells lies somewhere between the required rate and our derived upper limit, though, there may in fact be too much newly replenished gas to allow any decline in the neutral gas mass density. A possible resolution in this scenario would be increasing the proportionality between the gas outflow rates and the SFR as redshift decreases. This is not unreasonable, as the SFH is becoming progressively more dominated by lower-mass galaxies with decreasing redshift \\citep{Jun:05,Pan:07,Mob:08}. Galaxies with stellar masses $M_* \\lesssim 10^{10}\\,M_{\\odot}$ dominate the SFH at $z\\lesssim 1$ \\citep{Mob:08}. Such low-mass galaxies lose more mass in gas outflows in proportion to their SFR than high-mass galaxies, simply due to the former's shallower potential wells \\citep[e.g.,][]{DS:86,MF:99,FT:00}. This effect may contribute to the slow decline in the HI mass density.  \\subsection{Limitations of the proposed mechanism} \\label{limitations} We have treated a number of complex physical processes in very general terms. While being cautious of oversimplification, we have attempted to capture the essential interactions between star formation, recycling from stellar evolutionary processes, ISM processes of heating and ionization, recombination, cooling and molecule formation, together with infall from the IGM, and outflow of ISM material. Most of this complexity is concealed within the replenishment factor $K(t)$.  One issue is that stellar winds and SNe contribute to all components of the ISM rather than solely to $\\rho_{\\rm SFG}$. In a ``galactic fountain\" \\citep{SF:76,HB:90}, infalling gas will contribute to, and outflowing gas will strip from, all components. If recycled gas includes a component that never subsequently forms stars (such as some recycled gas in the ionized phase being ejected from the galaxy before contributing to star formation), the factor $+R\\dot{\\rho}_*(t)$ in Equation~(\\ref{theeqn}) will be reduced, and $K(t)$ will need to be increased to compensate.  Our quantitative results strongly depend on the assumed gas outflow rate. Variations by a factor of two or so in either direction will still result in a constant or slowly varying HI mass density, as long as a proportionality with the SFR of the host galaxies remains \\citep[as suggested by][]{Vei:05}. The chosen outflow rate is an effective average over all star forming galaxies, and is consistent with observed trends \\citep[e.g.,][]{Mar:99,Pet:00,Vei:05}. While individual galaxies show a large observed scatter between outflow rates and SFRs, for the ensemble properties of the total population this assumption should be robust.  The proposed replenishment through the supershell mechanism is not inconsistent with some simultaneous replenishment through infall. Metallicity considerations, which we do not address here, do require infall of some low metallicity gas \\citep{Erb:08}, and gas infall in local galaxies is well established \\citep[e.g.,][]{Bla:07,San:08}, although the observed infall rate is insufficient to match consumption."
    },
    "0806/0806.4102.txt": {
        "abstract": "We present an exhaustive numerical investigation of the optical caustics in gravitational lensing by a spinning black hole for an observer at infinity. Besides the primary caustic, we examine higher order caustics, formed by photons performing one or several loops around the black hole. Our investigation covers the whole parameter space, including the black hole spin, its inclination with respect to the line of sight, the source distance, and the caustic order. By comparing our results with the available analytical approximations, we find perfect agreement in their respective domains of validity. We then prove that all caustics maintain their shape (a tube with astroidal cross-section) in the entire parameter space without suffering any transitions to different caustic shapes. For nearly extremal spin, however, higher order caustics grow so large that their cross-sections at fixed radii wind several times around the black hole. As a consequence, for each caustic order, the number of images ranges from 2 to $2(n+1)$, where $n$ is the number of loops spanned by the caustic. As for the critical curves, we note that for high values of the spin they develop a small dip on the side corresponding to prograde orbits. ",
        "introduction": "If General Relativity is the correct theory of gravity, the Kerr solution describes the spacetime metric outside spinning black holes \\cite{Kerr}. Therefore it is currently utilized in all models trying to reconstruct the phenomena observed around observed astrophysical black holes, either remnants of stellar collapse or supermassive black holes lying in the central regions of several galaxies \\cite{Cha,Mel}. A crucial step in the comprehension of physics in such extreme environments is the complete understanding of the phenomenology related to the bending of photon trajectories caused by spacetime curvature. In order to get reliable predictions on any observables, it is necessary to keep in mind that a Kerr black hole acts as a very strong gravitational lens, generating an infinite number of images of the same source \\cite{HasPer}. The total flux of a source close to a Kerr black hole, (such as the accretion disk itself \\cite{Accret,BecDon}, an isolated bright spot on it \\cite{KVP,BroLoe} or a star orbiting the black hole \\cite{CunBar}) gets a significant contribution from the secondary image and from higher order images \\cite{BecDon}. This is also true for the details of fine structures in the profile of spectral lines, such as the iron $K\\alpha$ line in the X-ray domain, which are strong indicators of the presence of an intrinsic angular momentum of the black hole \\cite{Tan}.  On the other hand, the progresses in radio \\cite{Kri} and infrared band \\cite{IR} interferometry, along with the projects of X-ray interferometry in space (MAXIM, http://maxim.gsfc.nasa.gov), foreshadow that resolved pictures of the nearest supermassive black hole (Sgr A* in the center of the Milky Way) will be feasible in a not so far future \\cite{FMA}. This will represent a spectacular advance in our knowledge of black hole physics. In particular, the contributions of different images of the same source will be identified and studied separately. Higher order images will then provide a huge amount of independent information on the inner portions of the accretion disks of supermassive black holes and will become precious witnesses of the strong gravitational field just outside the horizon.  Gravitational lensing theory states that the multiplicity of the images of a given source depends on the source-lens-observer configuration. In a given spacetime metric, for an observer in a particular spacetime point, the multiplicity only depends on the source position. If the metric is stationary (as in the case of the Kerr metric) and the observer is static, the caustic can be defined in the 3-dimensional subspace at constant time as a 2-dimensional surface separating regions of space in which a source would give rise to a different number of images. When a point-like source crosses a caustic, a pair of additional images with infinite magnification is created or destroyed (the finite source size acts as a cut-off for real sources) \\cite{SEF}.  It can be easily guessed that the study of the shape of the caustics is of fundamental importance for a reliable and complete description of the whole phenomenology related to the environment of astrophysical black holes. In fact, the multiplicity of the images and their brightness is essentially determined by the position of the source within the caustic structure. Even temporal variations in the observed overall luminosity may be due to caustic crossing of bright features around the black hole \\cite{RauBla}.  Surprisingly, 45 years after the discovery of the Kerr metric, the complete structure of the caustics of a spinning black hole has not yet been derived. Indeed, the complexity of the metric prevents from finding simple analytical solutions for the caustic surfaces. Even numerical studies are very challenging and not straightforward. The first indication of the existence of non-degenerate caustics came from the work of Cunningham and Bardeen \\cite{CunBar}, who traced the light curves of a source star orbiting a spinning black hole. They noticed that the magnification of the primary and secondary images diverged at some particular points, signaling that caustic crossings were occurring. In spite of the huge number of ray-tracing codes in Kerr spacetime developed in so many years, the only comprehensive study of the caustic surfaces has been performed by Rauch and Blandford \\cite{RauBla}. They have explicitly shown that the primary caustic is a tube with a cross-section having the shape of an astroid (a closed curve with four cusps), which is very typical in gravitational lensing theory as soon as the spherical symmetry of the lens is broken by an external or internal perturbation. They have shown several pictures of the primary caustic and derived some simple asymptotic behaviors for its size. Besides the primary caustic, the authors have mentioned the existence of higher order caustics, but they have not shown any pictures of them, leaving several questions about the size and the shape of these caustics open. Later on, Sereno and De Luca have found an analytical approximation for the primary caustic valid for large source distances \\cite{SerDeL}. In a series of papers based on the strong deflection limit approximation \\cite{Boz1}, we have derived a perturbative analytical approximation describing the higher order caustics (but not the primary caustic) \\cite{BDSS,BDS,BozSca}. These approximations show that at low spin values the higher order caustics are still tubes with astroidal cross sections with increasing size. Basically, this is all we know about caustics in the Kerr spacetime.  All these studies provide several hints about the caustic structure of the Kerr black hole lens in several limits. Yet, the fate of the higher order caustics at high values of the spin still remains unclear. As they become larger and larger with increasing spin, do they undergo any transitions to different caustic shapes? Does their size stay finite? Do they merge? The present work provides clear answers to these and other questions of theoretical and observational relevance, clarifying the whole panorama of the caustic structure of the Kerr spacetime. We present a thorough numerical analysis of the caustics generated by a Kerr black hole at all caustic orders, studying their dependence on the black hole spin, its inclination and the source distance. The reliability of our results is also double-checked against former studies and all analytical approximations available up to now.  The paper is organized as follows. Section II traces the methodology followed for the generation of the caustics, referring to the appendix for a detailed explanation of all steps. Section III deals with the primary caustic. Section IV discusses the dependence of higher order caustics on the spin and its inclination. Section V focuses on the caustics of extremal black holes. Section VI is devoted to critical curves in the observer's sky. Section VII contains the conclusions.  %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% ",
        "conclusions": "%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%  In all applications of gravitational lensing, the study of critical curves and caustics of specific lens models has always represented a fundamental step in the comprehension of the whole phenomenology. As it can be easily imagined, the derivation of caustic surfaces in a full general relativistic context is much more involved than in classical lens models analyzed under the weak deflection paradigm. The simplest general relativistic lens is the Schwarzschild black hole. In this model, however, the caustics are degenerate and are therefore trivially tractable. The first general relativistic lens with a non-trivial caustic structure is the Kerr black hole. Going beyond the results of previous works, focused on particular limits of the caustics \\cite{RauBla,SerDeL,BDSS,BDS,BozSca}, this paper contains a complete investigation of the full caustic structure of the Kerr metric. This represents a considerable step forward for gravitational lensing phenomenology in full General Relativity.  In summary, we have shown 3-dimensional pictures of the primary and higher order caustic surfaces. We have analyzed the dependence of the cross-sections of the caustic surfaces on the source distance, the black hole spin, the spin inclination and the caustic order. We have shown that the caustic surfaces always wind an infinite number of times around the horizon following a logarithmic law. The size of the caustics always remains finite in the whole parameter space and there are no transitions to different kinds of caustic singularities in the lens mapping. For extremal spin values, the size of higher order caustics increases exponentially with the caustic order. This implies that the number of higher order images also grows exponentially, as the same caustic can be crossed an exponential number of times. We have compared our results with previous analytical and numerical approximations, finding perfect agreement in the respective limits of validity of the approximations. We have also shown some critical curves in the observer's sky, focusing on their dependence on the source distance.  In addition, the code developed in this paper for the calculation of the caustics is particularly well-suited for the study of higher order images, as being inspired by the Strong Deflection Limit methodology. An interesting future development of this code might be the implementation of an efficient resolution algorithm of the Kerr lens equation including higher order images.  Besides the purely theoretical interest of the analysis presented in this paper, we can easily imagine that the knowledge of the complete caustic structure of the Kerr black hole in all ranges of parameters will stand as an extremely helpful guide in future astrophysical applications involving very strong light bending by the gravitational field of spinning black holes."
    },
    "0806/0806.0512_arXiv.txt": {
        "abstract": "From extensive radiative transfer calculations we find that clumpy torus models with \\No\\ \\about\\ 5--15 dusty clouds along radial equatorial rays successfully explain AGN infrared observations. The dust has standard Galactic composition, with individual cloud optical depth \\tV\\,\\about\\ 30--100 at visual. The models naturally explain the observed behavior of the 10\\mic\\ silicate feature, in particular the lack of deep absorption features in AGN of any type. The weak 10\\mic\\ emission feature tentatively detected in type 2 QSO can be reproduced if in these sources \\No\\ drops to \\about 2 or \\tV\\ exceeds \\about 100. The clouds angular distribution must have a soft-edge, e.g., Gaussian profile, the radial distribution should decrease as $1/r$ or $1/r^2$. Compact tori can explain all observations, in agreement with the recent interferometric evidence that the ratio of the torus outer to inner radius is perhaps as small as \\about 5--10. Clumpy torus models can produce nearly isotropic IR emission together with highly anisotropic obscuration, as required by observations. In contrast with strict variants of unification schemes where the viewing-angle uniquely determines the classification of an AGN into type 1 or 2, clumpiness implies that it is only a probabilistic effect; a source can display type 1 properties even from directions close to the equatorial plane. The fraction of obscured sources depends not only on the torus angular thickness but also on the cloud number \\No. The observed decrease of this fraction at increasing luminosity can be explained with a decrease of either torus angular thickness or cloud number, but only the latter option explains also the possible emergence of a 10\\mic\\ emission feature in QSO2. X-ray obscuration, too, has a probabilistic nature. Resulting from both dusty and dust-free clouds, X-ray attenuation might be dominated by the dust-free clouds, giving rise to the observed type 1 QSO that are X-ray obscured. Observations indicate that the obscuring torus and the broad line region form a seamless distribution of clouds, with the transition between the two regimes caused by dust sublimation. Torus clouds may have been detected in the outflow component of H$_2$O maser emission from two AGN. Proper motion measurements of the outflow masers, especially in Circinus, are a promising method for probing the morphology and kinematics of torus clouds. ",
        "introduction": "Recent VLTI interferometric observations in the 8--13~\\mic\\ wavelength range by \\cite{Tristram07} confirm the presence of a geometrically thick, torus-like dust distribution in the nucleus of Circinus, as required by unification schemes of Seyfert galaxies. Several aspects of their data require that this torus is irregular, or clumpy, in agreement with the earlier prediction of \\cite{Krolik88}.  We have recently developed the first formalism for handling clumpy AGN tori and presented initial results \\citep{NIE02, Elitzur04, Elitzur06, Elitzur07}. The reported clumpy models have since been employed in a number of observational studies, including the first analysis of Spitzer observations by the GOODS Legacy project \\citep{Treister04}. Our clumpy torus models were also employed in the analysis of spatially-resolved, near-diffraction-limited 10~\\mic\\ spectra of the NGC~1068 nucleus \\citep{Mason06}. The geometry and kinematics of both water maser \\citep{Greenhill97, Gallimore01} and narrow-line emission \\citep{Crenshaw00} indicate that the NGC~1068 torus and accretion disk are oriented nearly edge-on. The \\citeauthor{Mason06} clumpy model for IR emission is the first to correctly reproduce the observed near-IR flux with an edge-on orientation. In contrast, smooth-density models require viewing angles 22\\deg--30\\deg\\ above the equatorial plane in order to bring into view the warm face of the torus backside \\citep{Granato97, Gratadour03, Fritz06}. Clumpiness is also essential for understanding the puzzling interferometry result that dust temperatures as different as $\\ga$ 800 K and \\about 200--300\\,K are found at such close proximity to each other \\citep{Schartmann05}. The mounting observational evidence in favor of clumpy, rather than smooth, dust distribution in AGN tori has sparked additional modeling efforts by \\cite{Dullemond05} and \\cite{Hoenig06}.  This two-paper series expands the analysis of \\cite{NIE02}. In its first part \\cite[part I hereafter]{AGN1} we develop the full formalism for continuum emission from clumpy media and construct the source functions of dusty clouds---the building blocks of the AGN torus. Here we assemble these clouds into complete models of the torus, and study the model predictions and their implications to IR observations. In comparing the predictions of any torus model with observations one faces a difficult problem---the overwhelming majority of these observations do not properly isolate the torus IR emission. Starburst emission is increasingly recognized as an important component of the IR flux measured in many, perhaps most, AGN \\citep[e.g.,][]{Netzer07}.  In addition to this well known contamination, even IR from the immediate vicinity of the AGN may not always originate exclusively from the torus, further complicating modeling efforts. A case in point is the \\cite{Mason06} modeling of NGC1068. All flux measurements with apertures $< 0.5''$ are in good agreement with the model results, but the flux collected with larger apertures greatly exceeds the model predictions at wavelengths longer than \\about 4\\mic. This discrepancy can be attributed to IR emission from nearby dust outside the torus. Mason et al show that the torus contributes less than 30\\% of the 10 \\mic\\ flux collected with apertures $\\ge 1''$ and that the bulk of the large-aperture flux comes at these wavelengths from dust in the ionization cones; while less bright than the torus dust, it occupies a much larger volume \\citep[see also][]{Poncelet07}. On the other hand, the torus dominates the emission at short wavelengths; at 2 \\mic, more than 80\\% of the flux measured with apertures $\\ge 1''$ comes from the torus even though its image size is less than 0.04$''$ \\citep{Weigelt04}.  These difficulties highlight a problem that afflicts all IR studies of AGN. The torus emission can be expected to dominate the AGN observed flux at near IR because such emission requires hot dust that exists only close to the center. But longer wavelengths originate from cooler dust, and the torus contribution can be overwhelmed by the surrounding regions. Unfortunately, there are not too many sources like NGC1068. No other AGN has been observed as extensively and almost no other observations have the angular resolution necessary to identify the torus component, making it impossible to determine in any given source which are the wavelengths dominated by torus emission. There are no easy solutions to this problem. One possible workaround is to forgo fitting of the spectral energy distribution (SED) in individual sources and examine instead the observations of many sources to identify characteristics that can be attributed to the torus signature. One example for the removal of the starburst component is the \\cite{Netzer07} composite SED analysis of the \\emph{Spitzer} observations of PG quasars. Netzer et al identify two sub-groups of ``weak FIR'' and ``strong FIR'' QSOs and a third group of far-IR (FIR) non-detections. Assuming a starburst origin for the far-IR, they subtract a starburst template from the mean SED of each group. The residual SEDs are remarkably similar for all three groups, and thus can be reasonably attributed to the intrinsic AGN contribution, in spite of the many uncertainties. However, while presumably intrinsic to the AGN, it is not clear what fraction of this emission originates from the torus as opposed to the ionization cones. An example of sample analysis that may have identified the torus component is the \\cite{Hao07} compilation of \\emph{Spitzer} IR observations. In spite of the large aperture of these measurements, Seyfert 1 and 2 galaxies show a markedly different behavior for the 10\\mic\\ feature, both in their mean IR SEDs and in their distributions of feature strength. Furthermore, Ultraluminous IR Galaxies (ULIRG) that are not associated with AGN show yet another, entirely different behavior, indicating that the observed mean behavior of Seyfert galaxies is intrinsic to the AGN. Accepting the framework of the unification scheme, the differences Hao et al find between the appearances of Seyfert 1 and 2 can be reasonably attributed to the torus contribution; the ionization cones dust is optically thin, therefore its IR emission is isotropic and cannot generate the observed differences between types 1 and 2.  Here we invoke both approaches in comparing our model predictions with observations. We start by assembling dusty clouds into complete models of the torus, as described in \\S\\ref{sec:Model}. Our model predictions for torus emission and the implications to IR observations are presented in \\S\\S\\ref{sec:SEDs}--\\ref{sec:SPP}, while in \\S\\ref{sec:Others} we discuss aspects of clumpiness that are unrelated to the IR emission, such as the torus mass, unification statistics, etc. In \\S\\ref{sec:Discussion} we conclude with a summary and discussion. ",
        "conclusions": "\\label{sec:Discussion}  We have developed a formalism for handling radiative transfer in clumpy media and applied it to the IR emission from the AGN dusty torus. In the calculations we execute only the first two steps of the full iteration procedure outlined in \\S3.2, paper I,  and the moderate total number of clouds considered here validates this procedure. When that number increases, the probability for unhindered view of the AGN decreases, the role of indirectly heated clouds becomes more prominent and eventually requires higher order iterations. Our current calculations employ some additional simplifying approximations: The grain mixture is handled in the composite-grain approximation, all dust is in clouds without an inter-cloud medium and all clouds are identical. We have already begun work on removing these assumptions and will report the results in future publications.  In contrast with the smooth-density case, the clumpy problem is not well defined because clouds can have arbitrary shapes, and any given set of parameters can have many individual realizations. Our formalism invokes a statistical approach for calculating an average behavior, and it is encouraging that other approaches produce similar results. \\cite{Dullemond05} conduct ``quasi-clumpy'' calculations in which the torus is modeled as a set of axisymmetric rings, and compare the results with the smooth-density case. In agreement with our conclusions they find that only smooth-density models can produce very deep absorption feature while clumpy dust produces stronger near-IR, broader SED and much more isotropic IR emission. \\cite{Hoenig06} employ 3D Monte carlo calculations that bypass some of our approximations. They also treat different cloud realizations for the same global parameters, allowing them to show the intrinsic scatter in SED due to the stochastic nature of the problem. Their results are in agreement with ours, validating our approach and the approximations we employ. Since the dust properties in their calculations are from \\cite{Draine84}, the 10 \\mic\\ feature reaches somewhat larger strengths than in our calculations, which employ the \\cite{OHM} ``cool\" dust \\citetext{but are similar to our original results in \\citealt{NIE02}, which also employed \\citeauthor{Draine84} dust}. In spite of these differences, \\cite{Hoenig06} too find that the silicate absorption feature is never as deep as expected for a uniform dust distribution, and obtain qualitatively similar behavior of the silicate emission feature and overall SED shape.  The models presented here show that clumpy torus models are consistent with current AGN observations if they contain \\No\\ \\about 5--15 dusty clouds along radial equatorial rays, each with an optical depth \\tV\\ \\about 30--100.  The cloud angular distribution should decline smoothly toward the axis, for example, a Gaussian profile centered on the equatorial plane. Power-law radial distributions $r^{-1}$ -- $r^{-2}$ produce adequate results. Dust grains with optical properties of the standard Galactic mixture provide satisfactory explanation to the IR observations. The behavior of the 10\\mic\\ silicate feature, in particular the lack of any deep absorption features, is reproduced naturally without the need to invoke any special dust properties. Several suggestions that the abundance or composition of AGN dust might differ from its Galactic counterpart can be discarded because of subsequent developments. \\cite{Risaliti99} note that, assuming standard dust abundance, the large column densities discovered in X-ray absorption imply torus masses in excess of the dynamical mass, posing a problem for the system stability. However, their mass estimates were based on the uniform mass distribution and large torus sizes derived from smooth-density models. The compact sizes and steep density distributions of clumpy models eliminate the problem (see \\S\\ref{sec:Mass}). \\cite{Maiolino01} suggested that the widely different UV and X-ray extinctions they found in individual sources could imply low dust abundance, but the subsequent discovery of rapid variations shows that X-ray obscuration by dust-free clouds is the more likely explanation (see \\S\\ref{sec:X-rays}). They also invoked the lack of prominent 10\\mic\\ absorption features as an indication that AGN dust is different from Galactic, but this is a natural consequence of clumpy dust distributions (see \\S\\ref{sec:sil}). Intrinsic extinction curves deduced from spectral analysis of type 1 sources \\citetext{see \\citealt{Czerny07} for a recent review and a comprehensive discussion of uncertainties} generally indicate a depletion of small grains, as could be expected: the obscuration in type 1 sources is dominated by the dusty clouds closest to the center and these clouds contain predominantly large grains, which survive at the smallest distances from the AGN (see \\S\\ref{sec:sublimation}). There is no compelling evidence for significant differences between the properties of AGN and Galactic dust. Other dust compositions are not ruled out, but nothing in the current data requires major departures from the dust grains we use.  The close proximity of dust temperatures as different as $\\ga$ 800 K and \\about 200--300\\,K found in interferometry around 12\\mic\\ cannot be explained by smooth-density models even when they account for the individual temperatures of grains with different sizes \\citep{Schartmann05}. Clumpiness resolves this puzzling observation because the dust on the dark side of an optically thick cloud is much cooler than on the bright side. Thanks to the mixture of different dust temperatures at the same radial distance, clumpy models naturally explain the torus compact size. In spite of the high anisotropy of its obscuration, the torus emission is observed to be nearly isotropic at $\\lambda \\ga$ 12 \\mic. Clumpy models resolve this puzzle too, since the emission from a torus with radial thickness $Y$ = 10 varies little with viewing angle. The variation is especially small if the radial distribution is $1/r^2$ or steeper, and such steep radial profiles maintain a nearly isotropic emission even at larger torus sizes.  In addition to IR observations, clumpiness significantly impacts the analysis of other data, in particular obscuration statistics. The fraction $f_2$ of obscured sources is controlled not only by the torus angular thickness $\\sigma$, as in all analyses to date, but also by the cloud number \\No. With \\No\\ = 5, a 70\\% fraction of type 2 AGN implies $\\sigma$ \\about\\ 30\\deg\\ instead of the standard 45\\deg. Observations indicate that increasing the bolometric luminosity from the Seyfert to the quasar regime induces (1) a decrease of $f_2$ and (2) a switch to emission feature at 10 \\mic\\ for both type 1 and some type 2 AGN. Both trends can be explained with a change in a single torus parameter---\\No\\ decreases from \\about\\ 5 in Seyfert galaxies to \\about\\ 2 in QSO (see figures \\ref{Fig:Sil10a} and \\ref{Fig:f2}). Decreasing $\\sigma$, the scenario known as the receding torus model, explains the first trend but has no effect on the second. The emergence of the 10 \\mic\\ in emission would require in this case the additional increase of individual clouds optical depth to \\tV\\ $\\ga$ 100 in QSO.  The decreasing-\\No\\ scenario provides the simplest explanation for the trends observed when $L$ is increasing, but that does not guarantee its validity. This demonstrates the difficulties in deducing the model parameters from observations that cannot yet resolve the torus basic ingredient---the individual dusty clouds. The problem is compounded by the lack of angular resolution that hinders clean separation of the torus component from the flux measured at most IR wavelengths and by the degeneracies of the radiative transfer solutions that prevent decisive, one-to-one associations between model parameters and observable quantities. The only practical way around these difficulties is to match trends identified in the data with similar general properties of the models.  Our main conclusions can be summarized as follows: \\begin{itemize} \\parskip -3pt  \\item The torus angular distribution has to be soft edged  \\item Clumpy models can produce nearly isotropic IR emission together with extremely anisotropic obscuration  \\item Clumpy models can explain all current observations with compact torus sizes; SED fitting is a poor constraint on the size  \\item Standard interstellar dust describes adequately AGN observations; there does not seem to be a need for any major modifications of grain properties  \\item Clumpy sources never produce a very deep silicate feature; apparent optical depth, obtained from $I = e^{-\\tau_{\\rm app}}$ where $I$ is the residual intensity at maximum absorption, is a poor indicator of the actual optical depth  \\item The probability for direct line-of-sight to the AGN at large viewing angles is small, but not zero  \\item The statistics of obscured sources depend on both the torus angular thickness and the number of clouds along radial rays  \\item The torus and the BLR are the dusty (outer) and dust-free (inner) regions in a continuous cloud distribution; a more appropriate designation for the torus is Toroidal Obscuration Region (TOR)  \\item X-ray obscuration comes from both TOR and, predominantly, BLR clouds  \\end{itemize}  As long as IR observations are incapable of resolving individual torus clouds, one must rely on the combined evidence for clumpy structure instead of on a ``smoking gun''. Individual TOR clouds seem to have been resolved in observations of outflow water masers in Circinus and NGC 3079. Proper motion measurements and comparison of these masers with their disk counterparts provide the most promising method for probing the TOR structure and kinematics. The Circinus AGN, whose dust emission has been resolved in VLTI observations, is an especially attractive target for studying the dusty and molecular content of TOR clouds."
    },
    "0806/0806.0038_arXiv.txt": {
        "abstract": "Herbig-Haro (HH) jets are commonly thought of as homogeneous beams of plasma traveling at hypersonic velocities. Structure within jet beams is often attributed to periodic or ``pulsed'' variations of conditions at the jet source.  Simulations based on this scenario result in knots extending across the jet diameter. Observations and recent high energy density laboratory experiments shed new light on structures below this scale and indicate they may be important for understanding the fundamentals of jet dynamics. In this paper we offer an alternative to ``pulsed'' models of protostellar jets.  Using direct numerical simulations we explore the possibility that jets are chains of sub-radial clumps propagating through a moving inter-clump medium.  Our models explore an idealization of this scenario by injecting small ($r<r_{jet}$), dense ($\\rho>\\rho_{jet}$) spheres embedded in an otherwise smooth inter-clump jet flow. The spheres are initialized with velocities differing from the jet velocity by $\\sim15$\\%. We find the consequences of shifting from homogeneous to heterogeneous flows are significant as clumps interact with each other and with the inter-clump medium in a variety of ways.  Structures which mimic what is expected from pulsed-jet models can form, as can previously unseen ``sub-radial'' behaviors including backward facing bow shocks and off-axis working surfaces. While these small-scale structures have not been seen before in simulation studies, they are found in high resolution jet observations. We discuss implications of our simulations for the interpretation of protostellar jets with regard to characterization of knots by a ``lifetime'' or ``velocity history'' approach as well as linking observed structures with central engines which produce the jets. ",
        "introduction": "Herbig-Haro objects have been the subject of significant analytical, observational, and numerical attention since their discovery. Observations using optical \\citep[e.g.][]{bally2002hh1/2} and IR \\citep[e.g.][]{velusamy2007hh46/47} techniques reveal that these jets typically show striking large scale collimation extending out to parsec distances combined with features appearing on a range of smaller scales.  Structure along the jet beam (``knots'' or ``clumps'') have, in particular, received considerable attention.  The origin of knots remains a subject of debate. Early studies focused on clumpyness of the HH bow shocks; \\cite{norman1979} postulated the existence of single ``interstellar bullets,'' while \\cite{schwartz1978} attributed the structures to stationary clumps being overrun by a wind. Stationary crossing shocks due to an overpressured jet beam expanding and then re-collimating were an early possibility that was considered for knots along the beam \\citep{buehrke1988,raga1990internal}. More recently, \\cite{rubini2007obliqueshocks} have suggested oblique shock focusing as a natural mechanism for hydrodynamic knot formation, though the presence of magnetic fields \\citep{hartigan2007bfields}, precession \\citep{masciadri2002precession}, and interactions with the environment \\citep{raga2002hh110, dalpino1999, yirak2008} all offer other means by which dense clumps might be created.  While considerable work has gone into these scenarios, currently the most favored model for the knots are internal working surfaces where shocks are driven down the beam by pulsation at the jet source. This ``pulsation'' model was first proposed by \\cite{rees1978} and was extensively explored by Raga and collaborators \\citep{raga1990ws, biro1994, raga1993}. In pulsed jet simulations the density and velocity cross-sectional profiles $\\rho_j(r)$ \\& $v_j(r)$ in the jet-launching region are kept fixed, while the magnitude of the velocity varies sinusoidally \\citep{raga1990ws, volker1999}. The pulsation scenario has become so dominant that even when attempting to address questions unrelated to clump formation, periodic inflow variations are frequently employed \\citep[e.g.][]{suttner1997}.  A variety of observational signatures can be recovered via pulsed jet models through careful choice of specific jet physical parameters and sinusoidal variability. In \\cite{raga2002hh34/111}, for example, a two-mode launching model was proposed using velocity histories extracted from observations of HH~34 and HH~111. Using these pulsation modes axisymmetric hydrodynamic simulations provided a convincing match to the location of the leading bow shock and the location of bright knots in the beam. These and similar results provide strong support for pulsed jet models.  A detailed examination of jets observed at the highest spatial resolution however shows features which do not fit into the pulsed jet paradigm.  In particular a number of ``archetypal'' jets show features at scales below the jet radius ($r<r_j$) which are distinctly displaced from the jet axis. In the case of HH~47 the jet clearly shows a non-axisymmetric morphology in the form an apparent helical bending of the beam \\citep{hartigan2005hh47}.  The beam itself is defined by a sequence of quasi-periodic knots with displacements to either side of the nominal jet axis. Explanations for this bending have included impacts with objects \\citep{raga2002hh110, dalpino1999}, magnetic fields \\citep{hartigan2007bfields} and precession of the jet source \\citep{masciadri2002precession}. We consider the presence of sub-radial, non-axisymmetric features to be a challenge to the pulsed jet paradigm.  Here, we consider an alternative to the pulsed jet model and investigate the consequences of intrinsic density heterogeneity in jet formation and evolution.  We are motivated to explore this model both by observations and new high energy density laboratory astrophysics (HEDLA) plasma experiments.  Using pulsed power wire array technologies \\cite{lebedev2005labjet}, \\cite{ciardi2007labjet}, and Ciardi et al.~(2008, in preparation) have presented experiments that track the evolution of fully magnetized, hypersonic, radiative jets.  The stability of hydro and MHD jets has long been a topic of debate, and these experiments shed some light on the real dynamics of 3-D systems \\citep{xu2000jetstability3d}. The experiments show that kink mode instabilities strongly affect the jet.  As the kink mode grows into the non-linear regime its disrupts but does not destroy the jet. The saturation of the instability transforms the jet into a sequence of collimated chains of knots which propagate with a range of velocities. Similar fragmented chains have been seen in other pulsed power experiments \\citep{golingo2005zpinch}.  Thus we propose a model in which the velocity and density profile of the jet are variable in time and space on scales less than the jet age and jet radius, respectively. Specifying conditions on these scales allows the model to achieve complex structures not seen before in simulations. Our simulations utilize AMR techniques and so are able resolve clumps in the jet beam at acceptable levels.  In the following sections, we describe the model in detail, offer dynamical and observational signatures, and briefly discuss the results. ",
        "conclusions": "We have performed 3-D simulations using a new ``collimated clumps'' scenario for protostellar jets.  Our model offers a fundamentally different paradigm for understanding jet origins and dynamics in the sense that heterogeneity is seen as being intrinsic and links the jet morphology on ``meso-scales'' to the processes (such as instabilities) occurring on ``micro-scales'' near the central engine.  While the pulsed-jet model has been successful at interpreting some aspects of jets, it may be misleading if used too generally. In particular, the assumption of sinusoidal pulsations can limit the interpretation of HH object observations. In our simulations, which had no sinusoidal variation in time, we nonetheless were able to recover (erroneous) sinusoidal behavior using an analysis similar to that which has been carried out in the past on observations. While one part of this behavior (the mean velocity $v_{j,0}$) fit the initial conditions well, the assumption of periodic pulsation allows the false conclusion that the structures in the beam arose due to periodic ejection behavior. We conclude therefore that care should be taken when attributing observations of apparent sinusoidal velocity in protostellar jets to corresponding sinusoidal behavior of a central engine.  In contrast to pulsed-jet models, our model offers two attractive features: first, a natural mechanism (knot-knot interactions) helps explain small-scale features along the jet axis. The idea that knots or bow shocks in HH objects are evolving clumpy structures has been discussed before in the context of observations \\citep[e.g.][]{reipurth2002hh34}, with ongoing HST observations giving continued support to this idea \\citep[e.g.][]{hartigan2005hh47}. The interaction of distinct sub-radial knots with each other and with an overall bow shock offers a simple explanation for such evolution. The second feature the model offers is the presence of unique observational characteristics in the form of forward- and reverse-facing bow shocks and spur shocks at the edges of the jet beam. Our scenario provides a simple mechanism for the formation of sub-radial non-axisymetric features via multiple dense, clumps which have a nonzero velocity dispersion. The ease with which non-axisymmetric features like spur shocks \\citep{heathcote1996hubble} develop in our models is attractive.  What would be the origin of the entrained clumps? The issues of the stability of jets particularly at ``micro-scale'' regions near the launch region remains unresolved \\citep{konigl2004diskwinds, ouyed2003jetstability, xu2000jetstability3d, micono1998khjets}. In addition, issues of magnetic field strength and velocity perturbations near the jet launch region also present problems \\citep{hartigan2007bfields}. Although we do not explicitly consider details of disk/jet launching here, plausible scenarios present themselves. It has long been considered that the interaction between the forming star and accretion disk---via magnetic fields---is responsible for jet launching \\citep{konigl1982}. Further, anisotropy of the magnetized disk during the YSO phase is a distinct possibility, as suggested by \\cite{combet2008}. One could therefore theorize the occurrence of a non-cylindrically-symmetric accretion burst as a possible explanation for non-axial density enhancements in the jet beam.  Recent high energy density laboratory astrophysics (HEDLA) investigations provide an unique window into the behavior of fully 3-D radiative hypersonic MHD jets.  These experiments demonstrate that magnetized jet beams in the lab may break up into a sequence of quasi-periodic knots due to the kink instability \\citep{ciardi2007labjet}. These knots may be displaced slightly to the side of the nominal jet axis and may propagate with varying velocities. This results in morphologies qualitatively reminiscent of HH-jet beams. It should be noted that the present simulation does not employ magnetic fields; however, it remains an open question whether magnetic fields remain dynamically important on the length scales of consideration here \\citep{hartigan2007bfields, ostriker2001molecularclouds}. It seems plausible that a similar process could occur in the astrophysical context, beginning with a smooth beam near the central engine which then becomes disrupted owing to the kink or sausage instabilities on intermediate scales. This would result in a series of knots which continue to evolve as they propagate away from the central engine. Such a scenario would also explain the observed velocity differences between knots, attributable to the particulars of each knot's formation. The present simulation is an idealization of this model.  The degree to which the results of the simulation presented here depend on the values chosen for parameters such as $\\chi_{cj}$ and $\\Delta v_c$ requires further investigation. In particular, $t_{launch}<t^*$ was not satisfied for all of the instantiated clumps, implying that some should disperse before reaching the jet head. The present simulation did not progress far enough to witness this behavior (that is, $t_{cc}>t_{sim}$). Also, using a $\\Delta v_c\\neq0$ and varying from clump to clump should result in an injection of considerable vorticity into the jet beam. Future work should therefore explore the long term evolution of the clumped jet. The aspect ratio (length/width) of the jet's final state in the present simulation is less than that for HH objects, implying again that simulations which progress for longer times would be of interest. While increased resolution would of course be of benefit, extracting details of clump-clump interaction within the jet-beam environment is probably outside the scope of this study and would be appropriate for a separate investigation (Dennis et al. 2008, in preparation). Finally, the inclusion of magnetic fields to study scenarios of knot formation and the resulting observational consequences would be of additional interest."
    },
    "0806/0806.2818_arXiv.txt": {
        "abstract": "Forty new low mass members with spectral types ranging from M4$-$M9 have been confirmed in the Orion Molecular Cloud 2/3 region.  Through deep, \\izjhk photometry of a 20$^{\\prime}$ $\\times$ 20$^{\\prime}$ field in OMC 2/3, we selected brown dwarf candidates for follow-up spectroscopy.  Low resolution far-red and near-infrared spectra were obtained for the candidates, and 19 young brown dwarfs in the OMC 2/3 region are confirmed.  They exhibit spectral types of M6.5$-$M9, corresponding to approximate masses of 0.075$-$0.015 $M_{\\odot}$ using the evolutionary models of \\citet{bcah98}.  At least one of these {\\it bona fide} young brown dwarfs has strong H$\\alpha$ emission, indicating that it is actively accreting.  In addition, we confirm 21 new low mass members with spectral types of M4$-$M6, corresponding to approximate masses of 0.35$-$0.10 $M_{\\odot}$ in OMC 2/3.  By comparing pre-main sequence tracks to the positions of the members in the H-R diagram, we find that most of the brown dwarfs are less than 1 Myr, but find a number of low mass stars with inferred ages greater than 3~Myr.  The discrepancy in the stellar and substellar ages is due to our selection of only low luminosity sources; however, the presence of such objects implies the presence of an age spread in the OMC 2/3 region.  We discuss possible reasons for this apparent age spread. ",
        "introduction": "The highly active star forming region known as the Orion Molecular Cloud 2/3 region (hereafter, OMC 2/3) is located in a molecular filament \\citep{b87} in the northernmost part of the Orion A molecular cloud at a distance of 450 pc \\citep{gs89}.  This region contains one of the richest concentrations of protostars and prestellar cores known.  Thirty-four submillimeter sources have been detected in OMC 2/3 at 1300 $\\mu$m and 350 $\\mu$m, evidence that it is a site of vigorous star formation activity \\citep{chini97,lis98}.  All six condensations in OMC 3, the northernmost part of OMC 2/3, fit the definition of Class 0 object, L$_{bol}$/L$_{submm}$ $<$ 200, and thus are in the earliest stages of stellar evolution \\citep{chini97}.  In addition, at least 80 knots of H$_2$ $v = 1 - 0$ $S(1)$ emission, signifying jets and outflows, have been detected in OMC 2/3, confirming the region's extreme youth \\citep{ybd97}.  Although OMC 2/3 has such a high density of star-forming cores, as shown by the richness of submillimeter sources, the population of pre-main sequence (PMS) stars and young brown dwarfs has not been studied as extensively or to similar depth in the infrared as the Orion Nebula Cluster (ONC), which is directly south of OMC 2/3 \\citep[e.g.][]{shc04,lrah01}.  As part of a study to survey the PMS star population in OMC 2/3, we present the results of a survey of young brown dwarfs.  With the existence of brown dwarfs firmly established, attention has shifted towards understanding the formation mechanism.  \\citet{rc01} proposed that brown dwarfs are formed by the ejection of accreting ``stellar embryos'' from dynamically unstable systems of protostars, resulting in the premature termination of accretion and substellar masses.  Our study of brown dwarfs in OMC 2/3 was initially motivated by this proposal; the detection of a large number of brown dwarfs around Class 0 sources would be evidence for ejection at slow velocities.  \\citet{rc01} suggest that given the amount of star formation activity, this would be the best place to find both very young brown dwarfs and potentially proto-brown dwarfs, and hence test models of their formation.  A more detailed discussion of brown dwarf formation in OMC 2/3 will be discussed in a subsequent paper on {\\it Spitzer} colors of the confirmed young brown dwarfs presented in this paper.  In order to obtain a complete census of the young brown dwarfs in OMC 2/3 (and eventually determine whether they are likely to have circumstellar disks), near-infrared observations are critical for identification.  We performed near-infrared observations that are two to three magnitudes deeper than the Two Micron All Sky Survey \\citep[2MASS;][]{2mass06} in order to identify the most deeply embedded substellar objects in OMC 2/3, down to 25 \\mjup, based on 1~Myr isochrones \\citep{bcah98}.  Before 2MASS, near-infrared observations of OMC 2/3 reached $K-$band magnitudes of 12.20 magnitudes \\citep{jgjhg90}; as a comparison, for our near-infrared observations we reach a magnitude limit of $K =$ 17.9.  Near-infrared and visible wavelength imaging and spectroscopy provide the information necessary for the selection and confirmation of substellar objects in star-forming regions.  Young brown dwarfs, with ages much less than 10~Myr, are much more luminous and thus easier to detect than field brown dwarfs (typically older than 1~Gyr) at distances greater than 100 pc. However, the near-infrared colors of young brown dwarfs are similar to those of very low mass stars making it difficult to identify brown dwarfs from near-infrared photometry alone.  Therefore techniques developed by \\citet{l00} are applied to select brown dwarf candidates in OMC 2/3 from combined, deep near-infrared and visible wavelength photometry.  Once candidate brown dwarfs are photometrically selected, far-red (0.6$-$1.0 $\\mu$m) and near-infrared (0.8$-$2.5 $\\mu$m) spectroscopic observations are used to confirm the candidates as {\\it bona fide} brown dwarfs.  We present a multiwavelength analysis of the OMC 2/3 region to identify the population of young brown dwarfs.  In Section \\ref{sec:observations}, we introduce the ground-based imaging of the OMC 2/3 region.  Next, in Section \\ref{sec:selection}, we describe the techniques used for selecting candidate young brown dwarfs for spectroscopy using the multiwavelength photometry. Then, we describe the spectroscopic observations (Section \\ref{sec:specobs}) and the spectral type classification methods used as well as the spectral types determined from the spectra (Section \\ref{sec:spectral_typing}). Finally, we discuss the age of OMC 2/3 in Section \\ref{sec:omc23age}. ",
        "conclusions": "Forty new low mass members with spectral types of M4$-$M9 have been spectroscopically confirmed in OMC 2/3.  Using deep, \\izjhk photometry of a 20$^{\\prime}$ $\\times$ 20$^{\\prime}$ field in OMC 2/3, we selected candidates with luminosities and colors consistent with brown dwarfs for follow-up spectroscopy.  Low resolution, far-red and near-infrared spectra were obtained for a sample of candidates and 19 objects with spectral types of M6.5$-$M9 and 21 objects with spectral types of M4$-$M6 were confirmed by fitting these spectra to standard spectra, and presented.  These spectral types correspond to masses of 0.35$-$0.015 $M_{\\odot}$ using the evolutionary models of \\citet{bcah98} and we estimate an age for OMC 2/3 from the substellar members of approximately 1~Myr.  However, a significant fraction of selected low luminosity sources are pre-main sequence stars with ages ranging from 1$-$10~Myr.  We discussed several reasons for the age spread we see in OMC 2/3.  First, that the older stars are part of the OB~1c association in the foreground to the Orion A cloud.  Although this possibility has not been ruled out, we find it unlikely because the older stars seem to be embedded in the Orion molecular cloud, as was found for the older stars in the ONC \\citep{shc04}.  Second, that the measurement of the ages are effected by uncertainties in the luminosity of the sources.  Our data exhibit some evidence that the presence of disks can lead to overestimating the age of the sources.  For example, the apparent old ages of the edge-on disk d216-0939 and the strongly accreting candidate 21 may be explained by the presence of an edge-on or nearly edge-on flared disk which is obscuring the star.  Finally, it is possible that we are observing a {\\it bona fide} age spread in OMC 2/3."
    },
    "0806/0806.0730_arXiv.txt": {
        "abstract": "{ We develop a method of multifractal analysis of $N$-body cosmological simulations that improves on the customary counts-in-cells method by taking special care of the effects of discreteness and large scale homogeneity.  The analysis of the Mare-Nostrum simulation with our method provides strong evidence of self-similar multifractal distributions of dark matter and gas, with a halo mass function that is of Press-Schechter type but has a power-law exponent $-2$, as corresponds to a multifractal.  Furthermore, our analysis shows that the dark matter and gas distributions are indistinguishable as multifractals.  To determine if there is any gas biasing, we calculate the cross-correlation coefficient, with negative but inconclusive results. Hence, we develop an effective Bayesian analysis connected with information theory, which clearly demonstrates that the gas is biased in a long range of scales, up to the scale of homogeneity.  However, entropic measures related to the Bayesian analysis show that this gas bias is small (in a precise sense) and is such that the fractal singularities of both distributions coincide and are identical. We conclude that this common multifractal cosmic web structure is determined by the dynamics and is independent of the initial conditions.} ",
        "introduction": "The large scale structure of the Universe can be described as a ``cosmic web'' formed by matter sheets, filaments and nodes.  This type of structure was initially proposed in connection with simplified but insightful models of the cosmic dynamics \\cite{Shan-Zel} and has been since confirmed by galaxy surveys and $N$-body cosmological simulations \\cite{Rien}.  Cosmological simulations have been especially helpful in testing models of structure formation. In a sense, they have been complementary to observations, since observations are biased towards the luminous matter, while simulations have fully considered the evolution of the dark matter, which is actually the dominant component. In fact, many simulations {\\em only} consider dark matter, in particular, non-baryonic cold dark matter, whose dynamics is simplest to simulate and gives rise to cosmic structure that is in accord with observations.  However, due to the advances in parallel computing, the development of efficient codes, and the availability of more powerful computers, the scope of $N$-body simulations has recently changed: now it is possible to simulate the combined dynamics of the non-baryonic dark matter and the baryon gas in large cosmological volumes and with relatively good resolution.  We analyse here the data output of a recent large cosmological simulation of the combined dark matter and gas dynamics, namely, a simulation of the cosmic evolution of $1024^3$ dark-matter particles and an equal number of gas particles carried out by the Mare-Nostrum supercomputer in Barcelona. This dataset has already been analysed by the researchers in charge of the Mare-Nostrum universe project \\cite{Gott1,Falten,Gott2}.  Here, we are interested in a particular aspect of the dark matter and gas distributions: their geometry and, specifically, their fractal geometry.  Fractal geometry \\cite{Mandel} is the geometry of sets or distributions that have noticeable geometrical features on ever decreasing scales. It is related to scale invariance and indeed appears in nonlinear dynamical systems in which the dynamics is characterized by the absence of reference scales. This is the case of the dynamics of collision-less cold dark matter (CDM), only subjected to the gravitational interaction.  Therefore, the cosmic web produced by this type of dynamics has fine structure and it is, arguably, statistically self-similar. We can reasonably assume that the cosmic web is a multifractal attractor of the gravitational dynamics.  This model is supported by the results of CDM simulations \\cite{Valda,Colom,Yepes,I4,I5}.  Although the gas dynamics is more complex (due to the gas pressure, etc), the gas takes part in the nonlinear dynamics of structure formation and can also have a multifractal attractor. Indeed, scaling laws in the distribution of galaxies have a long history, which has been reviewed in Refs.~\\cite{Borgani,Jones-RMP,Sylos-Pietro}.  Therefore, it is interesting to compare the scaling laws in the distribution of gas with the scaling laws in the distribution of dark matter.   Fractal models of the cosmic structure can only be valid in a range of scales, whose upper cutoff is the scale of homogeneity. Its value has been the subject of considerable debates and still is controversial \\cite{Jones-RMP,Sylos-Pietro}. In contrast, the lower cutoff to scaling has attracted less attention.  In fact, the CDM gravitational dynamics does not introduce any small reference scale that can play the r\\^ole of a lower cutoff, but the gas dynamics introduces the Jeans length.  This length is not a fixed reference scale, for it depends on the local thermodynamical parameters.  In any event, one should expect that the lower cutoff to scaling in the dark matter distribution is smaller than the lower cutoff appropriate for the distribution of galaxies.  However, the opposite seems to be true if one compares galaxy surveys with the results of cosmological simulations, since the latter exhibit reduced scaling ranges, even in dark matter only simulations.  Peebles has included this problem in his list of anomalies in standard cosmology \\cite{Pee}.  In his words: ``scale-dependent biasing seems an awkward way to account for the power-law forms of the low order galaxy position correlation functions.''  One can be inclined to place more trust in the scaling range found in galaxy surveys: cosmological simulations allow one to obtain better statistics but they are not free of systematic errors that affect an important range of the smaller scales. Indeed, it has been long known that $N$-body simulations are not fully reliable on scales smaller than the mean particle spacing $N^{-1/3}$ \\cite{KMS,KMSS}.  In spite of the ever-growing value of $N$, the range of scales between the scale $N^{-1/3}$ and the homogeneity scale is still rather small.  In the Mare-Nostrum universe, this scale range spans a factor of 30 (see Sects.~\\ref{anal} and \\ref{MF}).  Our goal is to demonstrate multifractality of the dark matter and gas distributions in the valid scale range.  Furthermore, given that this scale range is small, we devise a method to correct for discreteness effects and thus extend the valid range to smaller scales, obtaining a reasonable scaling range.  We also intend to test if the dark matter and gas distributions constitute a unique distribution or to what extent they differ. Hence, we make a model of {\\em fractal biasing}.  We describe our method of coarse multifractal analysis by counts in cells and define the basic objects (halos) in Sect.~\\ref{anal}.  In our method, the scale of homogeneity is explicitly introduced to calculate the multifractal spectrum (Sect.~\\ref{anal_w_homo}).  In Sub-sect.~\\ref{features}, we show how to obtain the main features of this spectrum and how they are influenced by discreteness and large scale homogeneity.  In Sect.~\\ref{MF}, we apply our method to the zero-redshift particle distributions of the Mare-Nostrum universe: (i) we obtain the halo mass functions and discuss its relation to the Press-Schechter mass function in Sect.~\\ref{mfun}; (ii) we obtain the multifractal spectra and discuss their relevance in regard to other geometrical studies of the cosmic web in Sect.~\\ref{MFsp}; and (iii) we demonstrate scaling and compute sound values of the correlation dimensions in Sect.~\\ref{frac_dist}.  The similarity of the results corresponding to the gas and the dark matter suggests that both distributions are identical and shows the need of precise statistical methods to discriminate between them (Sect.~\\ref{bias}).  Since the cross-correlations cannot give a definite answer (Sub-sect.~\\ref{cross}), we develop an effective Bayesian analysis (Sub-sect.~\\ref{Bayes-sect}) which we apply to various cell distributions (Sub-sect.~\\ref{appl_Bayes}). This analysis connects with the thermodynamic entropy of mixing (Sub-sect.~\\ref{G-DM_entropy}).  Therefore, we study the application of entropic measures to discriminating between mass distributions, and we study the connection of entropies in the continuum limit with the multifractal spectrum (Sect.~\\ref{entropies}).  Finally, we discuss our results (Sect.~\\ref{discuss}).  A note on notation: we use frequently the asymptotic signs $\\sim$ and $\\approx$; for example, $f(x) \\sim g(x)$ or $f(x) \\approx g(x)$ (often without making explicit the independent variable $x$). The former means that the limit of $f(x)/g(x)$ is finite and non-vanishing when $x$ approaches some value (which can be zero or infinity), while the latter means, in addition, that the limit is one.  We also use the sign $\\simeq$, which only refers to imprecise numerical values (with unspecified errors). ",
        "conclusions": "\\label{discuss}  We have improved the method of coarse multifractal analysis based on counts in cells by devising a procedure for extracting from a sample of a distribution the maximal information about its multifractal properties. The procedure is based on a clear understanding of the r\\^ole of the upper and lower cutoffs to scaling, which are, respectively, the homogeneity and discreteness scales. The homogeneity scale is used in the definition of coarse multifractal exponents [Eq.~(\\ref{ctauq})], while the discreteness scale is crucial to understand and quantify the effects of under-sampling.  We have employed our procedure to analyse the gas and dark matter distributions in the Mare-Nostrum universe at redshift $z=0$.  The only intrinsic scale present in an $N$-body simulation is actually the discreteness scale $V= N^{-1}$ (besides the size of the simulation cube, which we take as the reference scale).  The homogeneity scale is present as well but it is dynamical and grows with time.  Between these two scales the matter distribution can be considered continuous and representative of the nonlinear dynamics.  The discreteness scale $V= N^{-1}$ defines what we call the master cell distribution, which best resolves the overall mass distribution. The mass function of objects at this scale (halos) adopts a power-law form with a large-mass cutoff, similar to the Press-Schechter mass function. However, its power-law exponent is $-2$, which would correspond to an initial power spectrum with index $n=-3$ in the Press-Schechter theory, whereas the actual value in the Mare-Nostrum universe is $n=1$. In conclusion, the Mare-Nostrum mass function confirms the form of the mass function found in Ref.~\\cite{I4} and its independence of the initial power spectrum.  Of course, the Press-Schechter theory and the consequent mass function are not applicable to equal-size objects.  However, Vergassola et al \\cite{V-Frisch}, in their study of the adhesion model (described in Ref.~\\cite{Shan-Zel}), also define coarse-grained objects of equal size and, nevertheless, they find a power-law mass function with exponent depending on the initial spectral index and with an exponential large-mass cutoff, like in the Press-Schechter theory. On the other hand, Vergassola et al \\cite{V-Frisch} show that the adhesion model gives rise to a multifractal cosmic-web structure (see also Ref.~\\cite{Bou-M-Parisi}).  In this regard, it is especially interesting to compare our results with theirs, and to emphasize that the power-law exponent $-2$ is unrelated to the initial power spectrum, unlike their power-law exponent.  The dependence of their power-law exponent on the initial power spectrum is surely due to the nature of the Zel'dovich approximation, in which the dynamics is trivial before the formation of singularities.  In contrast, the real gravitational dynamics is {\\em chaotic}.  Therefore, the multifractal attractor of the real dynamics is independent of the initial conditions and must arise even when the initial conditions do not have a scale invariant power spectrum.  The mass function power-law exponent $-2$ is, in fact, naturally associated with the multifractal mass concentrate. Furthermore, we find that the precise form of the exponential large-mass cutoff suggests that the power law is actually an approximation of a lognormal mass function, as expected in a multifractal \\cite{I4} and found in the Mare-Nostrum universe on larger scales.  Our first direct test of scale invariance consists in calculating the coarse multifractal spectrum in a range of nonlinear scales, namely, from $l=2^{-12}$ up to $2^{-7}$.  For this, we use the improved definition of coarse exponents~(\\ref{ctauq}), which includes the scale of homogeneity (estimated through the condition $\\mu_2 =1.1$). This improvement is necessary when the scale of homogeneity is considerable smaller than the box size.  The resulting multifractal spectra (Fig.\\ \\ref{MFspec}) agree in their respective ranges (except near $\\a_{{\\rm max}}$). Moreover, the spectra corresponding to the dark matter and to the gas are almost identical.  However, the introduction of the scale of homogeneity produces an anomalous extension of the multifractal spectrum: it gives rise to {\\em negative} fractal dimensions. They can be understood as representing improbable matter fluctuations that can be ignored.  From the multifractal spectra, we deduce two important dimensions, namely, the dimension of the mass support $D_0 = 3$ and the dimension of the mass concentrate $D_1 \\simeq 2.4$.  Both dimensions provide information on the type of multifractal cosmic-web structure.  The former dimension shows that this multifractal is non-lacunar while the latter shows that it is not very concentrated.  The overall weak concentration indicated by $D_1 \\simeq 2.4$ can be due to the dominance of surface singularities (``pancakes'') but can also be due to the clustering of lower dimensional singularities, namely, filaments or nodes.  Cosmic web singularities are difficult to define in galaxy or $N$-body samples, but can be partially unveiled with appropriate algorithms \\cite{S4,vW-Sch}.  At any rate, one must notice that a non-lacunar cosmic web structure has a very complex geometry \\cite{I5}.  Of course, this geometry is determined by the dynamics of gravitational collapse and, in particular, by its type of anisotropy; but further discussion of this question is beyond the scope of this work (the r\\^ole of anisotropic collapse in the formation of the cosmic web is discussed in Ref.~\\cite{Rien}, for example).   Our study of the multifractal spectra on decreasing scales from $l=2^{-7}$ to $2^{-12}$, including the discreteness scale $l= N^{-1/3}= 2^{-10}$, allows us to discern the progressive influence of discreteness. The most obvious change is, of course, the shrinking range of $\\a$, namely, the reduction of $\\a_{{\\rm max}}$ caused by lack of mass resolution: depleted small cells must be empty. Furthermore, the mass distribution is under-sampled in cells with few particles, altering the ends of the spectra near $\\a_{{\\rm max}}$.  We can measure these deviations, for we can compare small scale spectra with the complete spectra at $l=2^{-7}$. Actually, the spectra are almost complete at $l=2^{-8}$. For $l>2^{-7}$, there appear early signs of the transition to homogeneity.  It is interesting to connect our results about the influence of discreteness, which only concern the statistical properties of the redshift $z=0$ distributions, with the studies by Kuhlman, Melott \\& Shandarin \\cite{KMS} and Splinter et al \\cite{KMSS} of the {\\em dynamical} effects of discreteness. Those authors conclude that these effects are the more important the less converging the particle motion is. Thus, we have, on the one hand, that expanding volume elements give rise to voids, with local dimension $\\a > 3$, which are only well represented in the multifractal spectra corresponding to scales considerably larger than $l = N^{-1/3}$. On the other hand, collapsing volume elements give rise to mass concentrations with the smaller dimension the larger is the number of independent axis along which they collapse.  These mass concentrations can be well represented in the spectra corresponding to $l < N^{-1/3}$.  For example, isotropic collapse gives rise to the smallest dimension concentrations, which are the most robust against the effects of undersampling; and, in fact, the low-$\\a$ end of the multifractal spectrum is essentially correct even for scales $l < 2^{-12}$. However, the strong singularities with low $\\a$ do not represent the full cosmic web structure.  Our second and most direct test of scale invariance is made in the standard way, namely, by studying the dependence of the second order moment $M_2$ on the scale $l$: we calculate $M_2(l)$ from $l=2^{-12}$ to $2^{-2}$, a broad range that includes the discreteness and homogeneity scales. On the smaller scales, we correct for the effect of discreteness by suppressing under-sampled cells, according to the information provided by the already computed spectra. We find two well-defined scaling ranges: the fractal range, spanning from $l=2^{-12}$ to $2^{-6}$, and the homogeneous range, from $l=2^{-4}$ upwards. The transition to homogeneity takes place between $l=2^{-6}$ and $l=2^{-4}$.  For definiteness, we choose as homogeneity scale $l_0=2^{-5}$, which in physical units is 16 $h^{-1}$ Mpc.  The fractal correlation dimensions are $D_2=1.26$, for the dark-matter, and $D_2=1.30$, for the gas, in accord with conventional values of the galaxy correlation dimension \\cite{Jones-RMP,Sylos-Pietro}.  To find out if the equivalence of the gas and dark matter distributions goes beyond their scaling properties, we have undertaken a detailed statistical study of the relation between these distributions.  Since we employ the method of counts in cells, we have specified two kinds of comparison: (i) the two cell distributions, defined by their respective sets of occupation numbers $\\{n_i\\}$, are compared as if they were two discrete probability distributions with respective probabilities $\\{p_i = n_i/N\\}$; (ii) the two cell distributions $\\{n_{{\\rm m}\\,i}\\}$ and $\\{n_{{\\rm g}\\,i}\\}$ are compared to decide if it is likely that they are samples from the same multinomial distribution (given by some coarse distribution $\\{p_i\\}$).  The first kind of comparison leads us to measures discriminating between discrete probability distributions (and between their continuum limits).  We have considered firstly the cross-correlation coefficient and lastly entropic distances (or ``divergences''), actually motivated by our method of deciding if two cell distributions are samples of the same multinomial distribution.  Since there are many (pseudo)distances to discriminate between discrete probability distributions, the comparison based on one of them has no absolute value. However, all the measures that we employ to discriminate between the coarse gas and dark matter distributions tell us that they are very close.  To decide if it is likely that the two cell distributions $\\{n_{{\\rm m}\\,i}\\}$ and $\\{n_{{\\rm g}\\,i}\\}$ are samples from the same multinomial distribution, we develop a Bayesian method of analysis.  The two distributions are compared by means of the Bayes information about the equality $p_{{\\rm m}} = p_{{\\rm g}}$, namely, by means of the logarithm of the corresponding Bayes factor~(\\ref{Bf1}).  The Bayes information corresponding to a set of massive cells can be expressed as a sum of negative cell terms, proportional to the entropy of mixing per particle minus one, added to a positive global term. The application of this formula to the master cell distributions, starting from the most massive halos, demonstrates gas biasing.  In particular, the gas is less concentrated in massive halos.  The bias is attenuated on larger scales but only disappears at $l=2^{-4}$, namely, at the scale of full homogeneity.  Naturally, it is to be expected that there is no bias at homogeneity, for it essentially preserves the initial conditions.  However, we do not have any argument that forbids that the bias vanishes at a smaller scale, so the fact that it vanishes only at homogeneity could be coincidental.  Since the Bayesian analysis can be formulated in terms of the entropy of mixing, we have studied in detail the entropic comparison of continuous distributions.  We must assume that the R\\'enyi entropies of the compared distributions have well defined continuum limits, which amounts to assuming that the distributions are multifractal (including regular distributions with $D_q=3$).  Thus, the first element of comparison is the spectrum of R\\'enyi dimensions or, equivalently, the multifractal spectrum.  As regards their multifractal spectra, the dark matter and gas distributions in the Mare-Nostrum universe are indistinguishable.  However, the multifractal spectrum gives the sizes of the sets of dark-matter or gas concentrations (or depletions) with equal strength but is insensitive to the location of those sets.  In fact, the R\\'enyi dimensions only contain partial information about a continuous distribution.  In particular, $D_1$ represents only one part of its entropy.  Another part of the entropy is of relational nature and can be expressed as a relative entropy or as a statistical entropic distance equal to (the square root of) the {\\em neg-entropy} of mixing, proportional to one minus the entropy of mixing per particle.  The high entropy of mixing or small entropic distance between the gas and dark-matter distributions is due to the fact that their respective singularities actually coincide, namely, the respective singularities at the same positions have equal local dimensions.  The appearance of common singularities in the gas and in the dark matter surely has a physical origin, despite the differences between the dynamics of each component.  It is natural to conjecture that the common multifractal structure is due to the fact that the gas and the dark matter are both dominated, on a long range of scales, by the gravitational interaction, which produces common power-law singularities.  The differences in the dynamics are the cause of gas biasing but do not interfere with the essential multifractal features of the distributions (except on very small scales).  In fact, the Mare-Nostrum universe is not based on a very realistic model of gas dynamics, insofar as it does not consider thermal radiation or conduction. Nevertheless, if the cosmic web singularity structure is due to gravity only, the analysis of future simulations will corroborate that the gas biasing does not alter that structure.  Then, we can speak of a kind of {\\em universality}: the cosmic dynamics has a unique type of cosmic web multifractal attractor, independent of the initial conditions.  In particular, the multifractal spectrum obtained here from the Mare-Nostrum universe or before from the GIF2 simulation \\cite{I4} must be characteristic of the cosmic web."
    },
    "0806/0806.2673_arXiv.txt": {
        "abstract": "s{ Directional detection can provide unambiguous observation of Dark Matter interactions even in presence of insidious backgrounds. The DM-TPC collaboration is developing a detector with the goal of measuring the direction and sense of nuclear recoils produced in  Dark Matter interactions. The detector consists of a  Time Projection Chamber with optical readout filled with CF$_4$ gas at low pressure. A collision between a WIMP and a gas molecule results in a nuclear recoil of 1-2 mm. The measurement of the energy loss along the recoil allows us to determine the sense and the direction of the recoil. Results from a prototype detector operated in a low-energy neutron beam clearly demonstrate the suitability of this approach to measure directionality. A cubic meter prototype, which is now being designed, will allow us to set competitive limits on spin-dependent Dark Matter interactions using a directional detector. } ",
        "introduction": " ",
        "conclusions": "Directional detection may hold the key to the unambiguous observation of Dark Matter in presence of backgrounds, and allows us to discriminate between models that predict Dark Matter to come from different directions in our galaxy.  The DM-TPC collaboration is developing a detector to achieve this goal. The device consists of a low-pressure TPC filled with CF$_4$ gas read out by an array of CCD cameras. Our prototypes proved the detector concept and demonstrated its ability to reconstruct both the sense and direction of nuclear recoils above 100 keV. A larger detector is being designed for underground operations in 2009 with the goal of obtaining competitive results on spin-dependent interactions using directional information. The success of this device will lay the foundation for a large Dark Matter experiment that will be able to detect the direction of WIMPs and discriminate between DM models in our galaxy."
    },
    "0806/0806.1670_arXiv.txt": {
        "abstract": "The Arecibo Legacy Fast ALFA (ALFALFA) survey is a program aimed at obtaining a census of HI-bearing objects over a cosmologically significant volume of the local universe. When complete in $\\sim$3--4 years, it will cover 7000 square degrees of high latitude sky using the 305~m telescope and the seven-beam Arecibo L-band feed array (ALFA). As of May 1, 2008, almost 60\\% of the required observations are complete and a catalog exists in preliminary form for 25\\% of the final sky area. ALFALFA is detecting about twice as many HI sources as predicted based on previously published HI mass functions and should deliver a final catalog of $> 25000$ extragalactic HI sources. ALFALFA will detect hundreds of galaxies with HI masses less than $10^{7.5} M_\\odot$ and similarly large numbers greater than $10^{10.3} M_\\odot$. Its centroiding accuracy allows for the immediate identification of highly probably optical counterparts to each HI detection. Fewer than 3\\% of all extragalactic HI sources, and $<$ 1\\% of ones with $M_{HI} > 10^{9.5} M_\\odot$ cannot be identified with a stellar counterpart. The hundreds of HI sources with observed line widths of $20-30$ km~s$^{-1}$ include a population of optically faint dwarf galaxies. The objects with highest HI masses exhibit a range of morphologies, optical colors and surface brightnesses, but most appear to be massive disk systems. The latter represent the population likely to dominate future studies of HI at high redshift. ",
        "introduction": "ALFALFA, the Arecibo Legacy Fast ALFA Survey, is a two-pass spectral line survey to cover 7000 deg$^{2}$ of high galactic latitude sky\\cite{Gio05a} with $\\sim$eight times the sensitivity, four times the angular resolution, three times the spectral resolution, and 1.6 times the total bandwidth of the HI Parkes All-Sky Survey (HIPASS)\\cite{Barnes01}. The ALFALFA survey strategy has been designed specifically to exploit Arecibo's superior sensitivity, angular resolution and digital technology to conduct a census of the local HI universe over a cosmologically significant volume\\cite{Gio08}\\cite{Hay08}. The effective integration time is $\\sim$40 sec per beam area, yielding approximately 2.2 mJy at 10 km~s$^{-1}$ resolution (after Hanning smoothing). The survey is intended to map, with complete 2-pass coverage, the region from 0$^\\circ$ to +36$^\\circ$ in declination and from $22^h <$R.A.$< 3^h$  and $7^h30^m <$R.A.$< 16^h30^m$. The fixed azimuth, ``minimum intrusion'' observing technique\\cite{Gio05b}\\cite{Gio07} delivers extremely high data quality and observing efficiency (99\\% ``open shutter'' time). Because of its wide areal coverage and photometric accuracy, ALFALFA is providing a legacy dataset for the astronomical community at large, serving as the basis for numerous studies of the local extragalactic Universe. The survey was initiated in February 2005; as of May 1, 2008, $\\sim 60$\\% of the survey observations have been completed.  The ALFALFA team is an open collaboration of more than 60 researchers from 34 institutions in 13 countries. Anyone with an interest in the science that can be done with the ALFALFA dataset and the willingness to contribute to the collective effort is welcome to join. Guidelines for joining ALFALFA can be found at \\url{http://egg.astro.cornell.edu/alfalfa/joining.php}. The ALFALFA survey is also serving as the backbone for student research projects at both the graduate and undergraduate level. One Ph.D. thesis\\cite{AS07t} based on ALFALFA is already complete; currently ten graduate students from six different institutions are working on Ph.D. dissertations centered on ALFALFA research. Current ALFALFA team projects are summarized at \\url{http://egg.astro.cornell.edu/alfalfa/projects/teamprojects.php}. The undergraduate ALFALFA program supports the participation of faculty and students at 14 institutions engaged in research, observing and educational exchange. An undergraduate workshop \\url{http://egg.astro.cornell.edu/alfalfa/ugradteam/ugradj08.php} is held at Arecibo each year. Eight ALFALFA senior theses have been undertaken so far.  Data processing for ALFALFA makes us of the IDL-based ALFALFA pipeline developed at Cornell and exported successfully to more than 16 institutions running both Linux and MacOS. Heavy use is made of Virtual Observatory protocols and web services for real-time cross-correlation with public multiwavelength databases. The identification of HI sources in the gridded data is performed using a Fourier domain matched filter signal extraction technique \\cite{AS07}.  Simultaneously, the most probable optical counterpart, where such exists, is identified using public OIR survey datasets. In addition to the most reliable, high S/N detections, sources of lower S/N but which coincide in both position and redshift with known optical galaxies are also included, but flagged as such, in the catalog of detections. Data catalogs and products are available at \\url{http://arecibo.tc.cornell.edu/hiarchive/alfalfa/}.  Two catalogs of HI sources extracted from 3-D spectral data cubes have been published\\cite{Gio07}\\cite{AS08} and a third has been submitted\\cite{Kent08}. Several additional publications should be submitted this summer, including a completed catalog of the $\\sim$1600 square degrees in the region $7^h30^m <$R.A.$< 16^h30^m$, $+04^\\circ <$ Dec $< +16^\\circ$. In the latter area, ALFALFA detects $> 6200$ high quality sources versus 290 for HIPASS. ",
        "conclusions": ""
    },
    "0806/0806.3443_arXiv.txt": {
        "abstract": "Recent work on abundance gradients have focussed not only on their magnitudes, but also on their spatial and temporal variations. In this work, we analyze the behaviour of radial abundance gradients in the galactic disk giving special emphasis on these variations. The data used includes planetary nebulae and objects in different age brackets, namely open clusters, HII regions, cepheid variables and stars in OB associations. We find evidences for a space variation of the radial gradients as measured  for element ratios such as O/H, S/H, Ne/H, Ar/H and [Fe/H], in the sense that the gradients tend to flatten out at large galactocentric distances. Moreover,  near the bulge-disk interface a steep decrease in the abundances is observed. The time evolution of the gradients is also evaluated on the basis of approximate ages attributed to the central stars of planetary nebulae and open cluster stars. It is concluded that the available data is consistent with a time flattening of the gradients during the last 6 to 8 Gyr, a time interval in which the age determinations are probably more accurate. ",
        "introduction": "Radial abundance gradients in the galactic disk are among the main constraints of chemical evolution models. They can be derived from a variety of objects, such as HII regions, planetary nebulae (PN), open clusters and young stars. Recent work on abundance gradients have focussed on (i) their magnitudes, (ii) their spatial variations and (iii) their temporal evolution, so that the number of derived observational constraints has been considerably increased. In this work, we analyze the behaviour of radial abundance gradients using a large sample of PN for which the ages of the central stars (CSPN) have been estimated. Such data is complemented with other objects in different age brackets, such as open clusters, HII regions, cepheid variables and stars in OB associations. We focus on the possible space variations of the gradients as measured  for elements such as O, S, Ne, Ar and Fe at large galactocentric distances and near the bulge-disk interface, as well as on the time evolution of the gradients based on approximate ages attributed to the CSPN and open cluster stars. ",
        "conclusions": ""
    },
    "0806/0806.3219_arXiv.txt": {
        "abstract": "We study, via numerical experiments, the r\\^ole of bound states in the evolution of cosmic superstring networks, being composed by $p$ F-strings, $q$ D-strings and $(p,q)$ bound states. We find robust evidence for scaling of all three components of the network, independently of initial conditions.  The novelty of our numerical approach consists of having control over the initial abundance of bound states. This indeed allows us to identify the effect of bound states on the evolution of the network. Our studies also clearly show the existence of an additional energy loss mechanism, resulting to a lower overall string network energy, and thus scaling of the network. This new mechanism consists of the formation of bound states with an increasing length. ",
        "introduction": "Cosmic superstrings can be formed as a result of brane annihilations in the context of brane-world scenario~\\cite{Sarangi:2002yt,Jones:2003da,dvalietal}. In the brane-world models, brane inflation takes place~\\cite{dvali-tye99} while two branes move towards each other, and their annihilation releases the brane tension energy that heats up the universe to start the hot big bang era. Typically, strings of all sizes and types may be produced during the collision. Considering a IIB string theory in a (9+1)-dimensional space-time, interactions of Dirichlet (D) branes leads to the {\\sl unwinding} and subsequent {\\sl evaporation} of higher dimensionality branes, with the survival of three-dimensional (D3) branes embedded in a (9+1)-dimensional bulk --- one of which could play the r\\^ole of our universe --- and D1-branes (D-strings)~\\cite{Durrer:2005nz}. Large Fundamental (F) strings and D-strings that survive the cosmological evolution become cosmic superstrings.  They are of cosmological size and could play the r\\^ole of cosmic strings~\\cite{Vilenkin_shellard,ms-cs07}, false vacuum remnants formed generically at the end of hybrid inflation within Grand Unified Theories~\\cite{jrs03,ms-cs08}. Cosmic superstrings have gained a lot of interest, particularly since it is believed that they may be observed in the sky, providing both, a means of testing string theory and a hint for a physically motivated inflationary model (for a recent review, see e.g. Ref.~\\cite{sakellariadou2008}).  Brane collisions lead also to the formation of bound states, $(p,q)$-strings, which are composites of $p$ F-strings and $q$ D-strings~\\cite{cmp04,lt04}.  The presence of stable bound states implies the existence of junctions, where two different types of string meet at a point and form a bound state leading away from that point.  Thus, when cosmic superstrings of different types collide, they can not intercommute, instead they exchange partners and form a $Y$-junction, as a consequence of charge conservation at the junction of colliding $(p, q)$-strings.  The evolution of F, D strings and their bound states is a rather complicated problem, which necessitates both numerical as well as analytical investigations. Junctions may prevent the network from achieving a {\\sl scaling} solution, invalidating the cosmological model leading to their formation.  In a number of studies, cosmic superstring evolution has been addressed via numerical experiments~\\cite{Sakellariadou:2004wq,Avgoustidis:2004zt,Copeland:2005cy,Hindmarsh:2006qn,Rajantie:2007hp,Urrestilla:2007yw}. The formation of three-string junctions and kinematic constraints for their collisions have been also investigated analytically~\\cite{cks06,Copeland:2006if,Copeland:2007nv}.  In what follows, we address the question of the effect of junctions in the evolution of a cosmic superstring network, being composed by three components: $p$ F-strings, $q$ D-strings and their $(p,q)$ bound states. The results of our numerical investigations can be summarised as follows. Firstly, there is clear evidence for {\\sl scaling} of all three components of the network, independently of the chosen initial configurations. Secondly, the existence of bound states effects the evolution of the network. Thirdly, for $(p,q)$ strings there is a supplementary energy loss mechanism, in addition to the chopping off of loops, and it is this new mechanism that allows the network to scale. More precisely, the additional energy loss mechanism is the formation of bound states, whose length increases, lowering the overall energy of the network. We note that in our simulations we can have control over the initial population of bound states, and this renders our novel results particularly important. ",
        "conclusions": "In the context of brane-world models, brane collisions lead to the formation of large Fundamental (F) strings and/or D1-branes (D-strings). Those who survive the cosmological evolution become cosmic superstrings, playing the r\\^ole of their solitonic analogues. By observing strings in the sky, we may be able to test for the first (and maybe only) time string theory. To know the observational consequences of cosmic superstrings, it is essential to master the properties and evolution of such networks.  Performing numerical experiments, we have studied the dynamics and overall properties of superstring networks.  More precisely, we have performed field theory simulations of $p$ F- and $q$ D-strings, and their $\\pq$ bound states. We have investigated the effect of bound states and in particular the approach to {\\sl scaling}, which is crucial for the cosmological consequences of cosmic superstring networks. Defects which do not scale are cosmologically undesired, since they may over-close the universe, leading to the {\\sl old} monopole problem.  Our studies have clearly shown that the three components of the network scale, independently of the chosen initial conditions. In addition, by having control over the initial abundance of bound states, we were able to identify the effect of bound states on the overall network evolution.  Finally, we have found that there is an additional energy loss mechanism, beyond the chopping off loops. This new mechanism consists on the formation of bound states, resulting to a lower  overall energy of the network, thus leading to scaling.   \\ack  The work of M.S. is partially supported by the European Union through the Marie Curie Research and Training Network {\\sl Universenet} (MRTN-CT-2006-035863).   \\vskip1.truecm"
    },
    "0806/0806.2609_arXiv.txt": {
        "abstract": "The cluster lens Cl~0024+1654 is undoubtedly one of the most beautiful examples of strong gravitational lensing, providing five large images of a single source with well-resolved substructure. Using the information contained in the positions and the shapes of the images, combined with the null space information, a non-parametric technique is used to infer the strong lensing mass map of the central region of this cluster. This yields a strong lensing mass of $1.60 \\times 10^{14} M_\\odot$ within a 0.5$'$ radius around the cluster center. This mass distribution is then used as a case study of the monopole degeneracy, which may be one of the most important degeneracies in gravitational lensing studies and which is extremely hard to break. We illustrate the monopole degeneracy by adding circularly symmetric density distributions with zero total mass to the original mass map of Cl~0024+1654. These redistribute mass in certain areas of the mass map without affecting the observed images in any way. We show that the monopole degeneracy and the mass-sheet degeneracy together lie at the heart of the discrepancies between different gravitational lens reconstructions that can be found in the literature for a given object, and that many images/sources, with an overall high image density in the lens plane, are required to construct an accurate, high-resolution mass map based on strong-lensing data. ",
        "introduction": "Due to the gravitational deflection of light, a galaxy or cluster of galaxies can affect the light that we receive from background sources. On larger scales, this leads to slight deformations of the shapes of the background sources, but close to the center of the deflecting object, the gravitational lens, more elaborate deformations are possible. When a background source is sufficiently well aligned with the gravitational lens, this strong lens effect can even cause multiple images of said source to appear. One of the most spectacular examples of strong gravitational lensing can be seen in the cluster lens Cl~0024+1654. Using recent ACS observations, one can easily see that five well resolved images depict a single source, but even before these five images were identified, it was clear that three arc segments were caused by a gravitational lens effect \\citep{1988lsmu.book..513K}.  This strong lensing information was first used in \\citet{1992ApJ...400...41K}. The authors of this work noted that these arc segments do not obey the so-called length theorem \\citep{1990LNP...360...16K}, implying that no simple elliptical lens model can be used. They show that if perturbations by cluster members are added, the observed arc lengths can indeed be reconstructed. In \\citet{1995ApJ...441...58W}, a more advanced reconstruction technique was used, consisting of a smooth lens model perturbed by some smaller galaxies and a non-parametric source model. Whereas previous work suggested that the main cluster potential was offset from the largest galaxy, these authors find that these positions, in fact, agree well.  After the first HST images clearly revealed the presence of five images, more lensing studies followed. The new, well resolved images were used in \\cite{1996ApJ...461L..83C} to study the source itself, a blue galaxy containing some interesting dark features and a bar-like structure. \\citet{1998ApJ...498L.107T} use the images to find the parameters describing elaborate lens and source models. Their algorithm constructs the complete image plane based on a set of source and lens parameters and compares the result with the HST observations. They find that the mass distribution is dominated by a smooth dark matter component with a considerable core radius, centered at a position near the largest cluster member.  Much of the earlier mass uncertainties originated from the poorly established source redshift. \\citet{2000ApJ...534L..15B} finally measured a spectroscopic redshift of 1.675 and used this information in their own inversion. They found that the image positions can be accurately reproduced using a model which traces the locations of the brightest cluster members. In \\citet{2007ApJ...661..728J}, a non-parametric method is used to invert the lens, using both strong and weak-lensing data. In the strong lensing region, the retrieved mass profile closely resembles the result of \\citet{2000ApJ...534L..15B}, but according to \\citet{2000ApJ...542L...1S}, the associated velocity dispersion is too high to correspond to the measured value of 1150 km s$^{-1}$ \\citep{1999ApJS..122...51D}.  In the present paper, we employ a non-parametric method to infer the mass map of Cl~0024+1654. This is the first time that only the information about the images themselves as well as the null space -- i.e. the region where no images are observed -- is used to reconstruct the mass distribution of this cluster in the strong lensing region. No information about the positions of cluster members is used. Clearly, many other possibilities have already been presented in the past, but it is not our intention to add to the confusion. Instead, the reconstruction is used to explain how the different previous inversions are related to each other.  Below, we will first briefly review the non-parametric technique that is worked out in detail in previous articles. In section \\ref{sec:inversion} this method is applied to reconstruct the mass distribution of Cl~0024+1654, and this result is used in section \\ref{sec:monopole} to illustrate the importance of the monopole degeneracy. The implications of these observations are discussed in section \\ref{sec:conclusion}.  \\begin{figure*} \\centering \\includegraphics[width=0.98\\textwidth]{fig1.eps} \\caption{The image parts of the five images of source A which were used in the reconstruction, labeled in the same way as in the work of \\citet{2007ApJ...661..728J}. Due to the extended nature of these images, several corresponding features are easily identified. The images shown here are not displayed on the same scale.} \\label{fig:inputA} \\end{figure*}  \\begin{figure*} \\centering \\subfigure{\\includegraphics[width=0.48\\textwidth]{fig2left.eps}} \\subfigure{\\includegraphics[width=0.48\\textwidth]{fig2right.eps}} \\caption{Left panel:~after averaging 28 individual reconstructions, this is the resulting mass map for Cl~0024+1654 predicted by our procedure. The positions of the input images of source A are also indicated in this figure. The critical density used in this figure corresponds to a redshift $z=3$. Right panel:~the standard deviation of the individual reconstructions shows that the different solutions tend to disagree about the exact shape in the central part of the mass distribution. In particular, this figure suggests that the mass peak around $(0.2', -0.2')$ in the left panel should not be regarded as an actual feature.} \\label{fig:sol1} \\end{figure*} ",
        "conclusions": "\\label{sec:conclusion}  In this article we have applied a previously described non-parametric inversion method to the cluster lens Cl~0024+1654. The method uses both the information from the extended images and the null space and can easily be adapted to incorporate other available constraints. It requires the user to specify a square shaped area in which the algorithm should search for the mass distribution and it is assumed that no mass resides beyond the boundaries of this region, but no other bias is present. Different runs of the inversion method can produce results that differ somewhat, depending on the amount of constraints available. This allows one to inspect which features are common in all inversions and which aspects tend to differ.  Using this inversion procedure we obtained an averaged mass map which clearly displays features that can also be seen in the ACS images. The most recent gravitational lensing study of this lens, is that of \\citet{2007ApJ...661..728J}, which used both strong and weak lensing data. The strong lensing mass of $1.60\\times 10^{14} M_\\odot$ is less than the value of $(1.79 \\pm 0.13)\\times 10^{14} M_\\odot$ found in their study, but it is still in good agreement. We mentioned earlier that our size estimate for source A is higher than found in other works. This is a well-known consequence of a generalized version of the mass-sheet degeneracy, for which the name steepness degeneracy is more appropriate. As we showed in \\citet{Liesenborgs3}, this steepness degeneracy is very hard to break for lensing systems with only a handful of sources, even if these have different redshifts. As the original mass-sheet degeneracy, the generalized degeneracy leaves the observed images identical but the reconstructed sources are scaled versions of the original ones while the density profile of the lens becomes less steep.  The relation with the inversion of \\citet{2007ApJ...661..728J} can be revealed by comparing the predicted source sizes. The size of source A in our inversion is five times larger than in their work, thereby identifying the scale factor in the mass-sheet degeneracy.  When we downscale our mass reconstruction by a factor of five and add a constant sheet of mass in such a way that the strong lensing mass is unaffected, the circularly averaged density profile in Fig.~\\ref{fig:sol1profiledegen} is obtained (thick black line). This clearly shows much resemblance to the profile found in \\citet{2007ApJ...661..728J} in the strong lensing region. Note that since our reconstruction procedure only looks for mass in a region which is 1.3$'$ $\\times$ 1.3$'$ in size, the profile will quickly drop to zero beyond the range shown in the figure.  \\begin{figure} \\centering \\includegraphics[width=0.48\\textwidth]{fig8.eps} \\caption{The density profile in the circular region indicated in Fig.~\\ref{fig:sol1overlay} is described by the dotted curve. If this profile is scaled down by a factor of five and a mass sheet is added to keep the strong lensing mass constant, the the profile described by the thick black line is obtained. In the strong lensing region, it clearly resembles the profile shown in the work of \\citet{2007ApJ...661..728J}, suggesting that the results shown here differ mainly by the mass-sheet degeneracy. This figure also clearly shows that the strong lensing mass estimate from this work differs from the one in \\citet{2007ApJ...661..728J}.} \\label{fig:sol1profiledegen} \\end{figure}  When the monopole degeneracy was applied to the case of Cl~0024+1654, a simple optimization routine was used to remove substructure from the previously obtained mass distribution. However, there is no general rule as to how the mass map may be modified. For example, with some extra effort the existing mass map could have been transformed into one which followed the light more closely, or which corresponded better to the available X-ray data \\citep{2004ApJ...601..120O}. The only constraints which matter in this respect are the absence of unobserved images and possibly dynamic measurements. Image positions, fluxes and time delays are completely unaffected by this type of degeneracy, which allows you to redistribute matter in any number of ways. This freedom is illustrated in Fig.~\\ref{fig:degendiff}, which depicts the differences between the two mass distributions shown in this article.  The monopole degeneracy seems to be under-appreciated: the only direct application that can be found is in the work of \\citet{2006ChJAA...6..141Z}, where circularly symmetric modifications of power-law models for PG~1115+080 were explored. Yet from the discussion above it is clear that the degeneracy is an important aspect of any gravitational lens inversion, as it can be used to introduce or remove many kinds of features. The explanation in section \\ref{sec:monopole} illustrates the importance of the distance between the images, implying that the resolution that can be obtained when inverting a strong gravitational lens system is determined by the local density of the images. This fact is also mentioned in the work of \\citet{LensPerfect}, but was not linked to the monopole degeneracy. It is also interesting to note that when the total mass in the region indicated in Fig.~\\ref{fig:sol1overlay} is calculated for the degenerate solution, one finds the slightly larger value of $1.62 \\times 10^{14} M_\\odot$.  This indicates how this degeneracy may be responsible for differences in strong lensing masses in different studies.  Using both the generalized mass-sheet degeneracy and the monopole degeneracy as described in this work, it seems likely that the majority of differences between existing lens models can be explained. The first thing that one needs to do, is look at the predicted source sizes. This readily identifies the presence of the mass-sheet degeneracy. When this is compensated for, the remaining differences can then be minimized by redistributing the mass using the monopole degeneracy (which can also easily alter the steepness of the mass distribution). Clearly, if an accurate mass map is required without additional assumptions about the shape of the distribution, a large amount of images are needed. Without sufficient coverage by images, a fundamental and large uncertainty exists in the regions between the images. This can only be resolved by identifying additional multiply imaged sources, and not by more detailed observations of the existing images (although this can reveal small-scale substructure in the vicinity of these images). This fundamental uncertainty in the overall lens equation also implies that care must be taken when using an existing lens model in trying to identify new multiply imaged sources.  \\begin{figure} \\centering \\includegraphics[width=0.48\\textwidth]{fig9.eps} \\caption{This plot shows the differences between the mass distributions in Fig.~\\ref{fig:sol1} and Fig.~\\ref{fig:degen}. Clearly, the structure has been altered in a way which does not display any particular symmetry.} \\label{fig:degendiff} \\end{figure}"
    },
    "0806/0806.3749_arXiv.txt": {
        "abstract": "We have designed and tested an in-plane \\'{e}chelle spectrograph configured to investigate precision radial velocities from ground-based near-infrared observations. The spectrograph operates across the spectral range of $0.9$-$1.7$~\\mm\\ at a spectral resolution of $R = 50,000$, and uses a liquid nitrogen-cooled  {\\sc hawaii} 1K detector. Repeated measurements of the Earth's rotation via integrated Sunlight with two different instrument arrangements in the near infrared Y band have produced radial velocities with $\\sim 10$~\\ms\\thinspace RMS over a period of several hours. The most recent instrument configuration has achieved an unbinned RMS of $7$ \\ms\\thinspace and suggests that infrared radial velocity precisions may be able to approach those achieved at optical wavelengths. ",
        "introduction": "The achievement of meter-per-second radial velocity precision is one of the major technological breakthroughs of recent years. Although this effort is frequently viewed as driven by the search for extra-solar planets following the discovery of the first such system in the early 1990s (Wolszczan \\& Frail 1992), the quest for highly accurate radial velocity measurements was actually under development for many decades beforehand. The path from solar studies (Becker 1976; Koch \\& Woehl 1984) to pioneering stellar observations using HF gas cells (Campbell \\& Walker 1979; Campbell, Walker \\& Yang 1988) to today's standards of long-term several \\ms\\thinspace (or short-term sub-\\ms) RMS measurements with Keck/HIRES (e.g., Vogt et al. 1994, Butler et al, 1996), ESO~3.6m/HARPS (e.g., Pepe et al. 2000), HET/HRS (e.g., Tull 1998), VLT-UVES (Dekker et al. 2000), and AAT/UCLES (Tinney et al. 2001) is one that leads through several spectrographs, detectors, and calibration methodologies.  To date, the only astronomical programs that have yielded sub-16~\\ms\\thinspace radial velocities are based on optical ($500$- to $600$-nm) \\'{e}chelle spectroscopy, with velocity calibration based upon I$_2$ absorption or Th-Ar emission line references. While this wavelength region is optimal for investigating solar-like stars, it suffers when investigating cooler objects (especially in comparison with the infrared region).  Main sequence M stars are particularly inviting targets for precision radial velocity work, both because these objects represent the most numerous class of stars and their low mass makes them promising objects for detection of Earth-mass planets in the ``habitable zone\" (Kasting \\& Catling, 2003).  This is one of the top priorities of the Exoplanet Task Force~\\footnote{http://www.nsf.gov/mps/ast/aaac/exoplanet\\_task\\_force/reports/aaac\\_draft.pdf}.   No instruments have published radial velocity results in the near infrared with precisions approaching those in the optical, although exciting science results have been obtained in this wavelength region with hundreds of \\ms\\thinspace precisions. For example, Stassun, Mathieu \\& Valenti (2006) characterized the orbital parameters of the brown dwarf binary \\hbox{{\\sc 2mass} J05352184$-$0546085} using the {\\sc phoenix} spectrograph (Hinkle et al. 2003) on the Gemini South telescope. Using six nights of observations with {\\sc nirspec} (McLean et al. 1998) at the Keck telescope, Martin et al. (2006) surpassed the precision achieved by previous optical radial  velocity work on late-type stars; these infrared radial velocity measurements of the nearby brown dwarf LP944-20 report an RMS precision of 360~\\mbox{m s$^{-1}~$}. Blake et al. (2007) used CO bands and telluric features to reach precisions of approximately 300~\\ms\\thinspace around L dwarfs.  Moving beyond the CCD range to a wavelength regime where there is no tradition of precision ($< 20$ \\ms) radial velocity studies is challenging because there is no legacy work to build upon. In time it should be possible to increase the sensitivity of tellurically-calibrated programs down to mean atmospheric velocities of 10-25~\\ms\\thinspace that have now been achieved in the optical (\\eg Johnston, 2006 and Gray \\& Brown, 2006).  However, more precise calibration techniques are required if infrared observations are to reach the \\ms\\thinspace accuracy of current state-of-the-art optical projects.  This paper describes our experiments with a laboratory-based, high-resolution infrared spectrograph configured to measure precision radial velocities. This instrument, designated the Precision Radial Velocity Spectrograph Pathfinder (hereafter, Pathfinder), was designed and constructed in the Department of Astronomy and Astrophysics at Penn State. The goal of this project was to measure the rotation velocity of the Earth by obtaining infrared spectra of integrated Sunlight to explore the limitations of making precision radial velocities in the $0.9$-$1.7$~\\mm\\thinspace spectral region. The most serious concern is the lack of a proven calibration technique.  NIR arrays also have issues not present in CCDs, such as inter-pixel capacitance, multiplexer cross-talk, persistence, and sensitivity to thermal background radiation.  Telluric contamination is also a major issue in the NIR, as is modal noise, since the number of modes in optical fibers decreases with wavelength.  We carried out initial observations in August 2006 and, after a series of upgrades and intermediate tests, conducted another set of observations in November 2007.  \\S~2 presents the design of the instrument and the properties of its various components.  The data processing and analysis are described in \\S~3.  Our conclusions, which point to the challenges that lie ahead, are given in \\S~4. ",
        "conclusions": "For the first time in the NIR, we have demonstrated RMS radial velocity precisions below $10$ \\ms\\thinspace over a period of several hours, using $10$-s observations of the Sun. There are a number of improvements that we hope to implement when additional funding is secured.  Foremost is to install a R4 \\'{e}chelle appropriately sized and a VPH or standard cross-disperser grating optimized for the $1-1.3$~\\mm region.  These two items in themselves will lead to close to an order of magnitude increase in throughput.   In addition, we will update our calibration system to allow us to move from precision acceleration measurements over a short period of time to precise velocity measurement over weeks to months.    The upgraded calibration system will include both Th-Ar and U-Ar hollow cathode lamps as well as a gas cell that is illuminated by a Q-I lamp.  While we do not intend to use a gas cell simultaneously in line with the object observations as pioneered by Butler et al (1996) in the visible, it will allow a better absolute velocity determinations as well as a means to track systematics in the hollow cathode lamps.   We are implementing the capability to rapidly switch between the target spectrum and the gas cell to allow calibrations between target spectra.  We will initially test HF and water vapor gas cells in the 1.0 to 1.25 micron range but note that no single gas has a large number of deep lines over the desired wavelength range.  Adding the U-Ar hollow cathode lamp will increase the number of lines at least a factor of two and allow less dependence on Ar lines.  We also will investigate the effects persistence may have on the achievable precision.  With the noted improvements we will be ready to move this instrument to the Hobby-Eberly Telescope for tests on M dwarfs to firmly establish the viability of NIR precision radial velocity spectroscopy in these objects."
    },
    "0806/0806.3113_arXiv.txt": {
        "abstract": "The energy levels of hydrogen and helium atoms in strong magnetic fields are calculated in this study. The current work contains estimates of the ground and first few excited states of these systems that are improvements upon previous estimates. The methodology involves computing the eigenvalues and eigenvectors of the generalized two-dimensional Hartree-Fock partial differential equations for these one- and two-electron systems in a self-consistent manner. The method described herein is applicable to calculations of atomic structure in magnetic fields of arbitrary strength as it exploits the natural symmetries of the problem without assumptions of any basis functions for expressing the wave functions of the electrons or the commonly employed adiabatic approximation. The method is found to be readily extendable to systems with more than two electrons. ",
        "introduction": "Introduction}  The motivation to study atoms in magnetic fields of strength beyond the perturbative regime was in a large part due to the discovery of such fields being present in white dwarf stars \\cite{Kemp1970, Angel1978, Angel1981} and neutron stars \\cite{Trumper1977,Trumper1978}. The most commonly observed neutron stars - pulsars, have been observed to have magnetic fields on the order of \\begin{math}10^{11}\\end{math} - \\begin{math}10^{13}\\end{math}G \\cite{Ruder94}. Magnetars \\cite{DT1992}, which are strongly magnetized neutron stars, can have magnetic field strengths well in excess of \\begin{math}10^{13}\\end{math}G. White dwarf stars on the other hand have somewhat less extreme fields, albeit still high, \\begin{math}\\sim10^{6}\\end{math} - \\begin{math}10^{8}\\end{math}G \\cite{Ruder94}. At the high field strengths observed in these compact objects, the electron cyclotron energy of an atom becomes greater than the corresponding Coulomb potential energy, i.e. \\begin{math}\\hbar \\omega_B > Ze^2/r\\end{math} \\cite{Ruder94}. Here, \\begin{math} \\omega_B \\end{math} is the cyclotron frequency. Thus, a Zeeman-type perturbative treatment of the field \\cite{Landau} is not possible. The structure of atoms in both instances however, is considerably altered from the low field case. At high field strengths, spherical symmetry of the atom is broken and the atom is stretched along the field, however, azimuthal symmetry remains intact and thus it has been observed that it is more convenient to model the atom in cylindrical coordinates \\cite{Ruder94}.  Since the 1970's this problem has been tackled by various researchers using different methods. Hitherto, the most tractable and accurate approaches have relied upon an assumed basis of functions for expressing the electron wave functions of an atom in a strong magnetic field. Using such an assumption the problem of the hydrogen atom in a strong magnetic field was tackled using either a variational approach \\cite{CK1972} or by attempting to solve the Schr\\\"{o}dinger equation directly \\cite{Praddaude1972, SV1978, Friedrich1982, WR1980, RWRH1983, RWRH1984, Ivanov1988}. Initial attempts for estimating the energies and wave functions of different electronic states of the helium atom were based upon a purely variational approach \\cite{Cohen1970, Henry1974, Mueller1975, Banerjee1974, CRS1974, GK1975, Glasser1975, Larsen1979, VB1989} or $Z-$dependent perturbation theory \\cite{Gadiyak1982}. The majority of these studies however were analytical and were therefore limited in their applicability to a problem that was inherently more tractable numerically. The most accurate and reliable solutions to the helium atom in a strong magnetic field, thus far, involved using the Hartree-Fock (HF) technique \\cite{Hartree}. Employing the HF technique researchers in the past have calculated energy levels and wave functions of the ground and first few excited states of the helium atom and helium-like species \\cite{Virtamo1976, Proschel1982, Thurner1993}. These treatises employed Landau orbitals \\cite{Landau} to describe the motion of the electron perpendicular to the field. The electrons were required to reside in the ground Landau state, thereby simplifying calculations somewhat by restricting the wave functions to the so called adiabatic approximation \\cite{Proschel1982}. Ivanov \\cite{Ivanov1994} in 1994 obtained similar results for the helium atom in strong magnetic fields using an unrestricted HF technique. Elsewhere, Quantum-Monte-Carlo methods were employed successfully for determining the ground and first few excited states of the helium atom in low to strong magnetic fields \\cite{Jones1996, Jones1997, Jones1999}. In the treatises described above \\cite{Virtamo1976, Proschel1982, Thurner1993, Ivanov1994, Jones1996, Jones1997, Jones1999}, usually an approximation was employed for calculating both the direct and the exchange interactions between the electrons. Such approximations generally involved finding appropriate expansions that mimicked the behaviour of the inter-electron terms in the Hamiltonian. Necessarily, such an approach is limited by the accuracy of the expansions employed. In addition, this increases the complexity of the computational problem. Heyl \\& Hernquist \\cite{HH1998} in 1998 described both an analytical as well as a numerical approximation for evaluating the effective inter-electronic potentials that was not only intuitive but also computationally less expensive. Their numerical technique extended the idea of adopting a basis of functions for the directions both transverse and parallel to the field. They constructed the wave function in the axial direction with the help of harmonic oscillator Hermite polynomials. This method was seen to yield accurate results for high magnetic fields \\begin{math}\\beta_Z \\geq 10^3 \\end{math} for hydrogen and helium. They employed a method of calculating the total energy of the system given an assumed set of wave functions and then proceeded to minimise this energy by varying the free parameters of the wave functions. This was seen to yield accurate results consistent with other work \\cite{Ruder94}. One of the advantages over other methods \\cite{Virtamo1976, Proschel1982, Thurner1993, Ivanov1994, Jones1996, Jones1997, Jones1999} was observed to be the significantly lesser amount of computation involved. More recently, Mori \\& Hailey \\cite{MH2002} and Mori \\& Ho \\cite{MH2007} adopted a perturbative approach to treat the exchange terms and higher Landau states with success for finding upper bounds for the energies for the ground and first few excited states of helium and other mid-Z atoms in high magnetic field strengths.  In the literature there are numerous studies outlining accurate estimates of the energies of different states of the hydrogen atom, however the number of investigations of the helium atom and helium-like species in the strong or intermediate magnetic field regime is rather small. The estimates of the energy levels of these species are only moderately accurate, and the computational expense is rather high. However, for most observable neutron stars and many white dwarf stars the magnetic field strengths lie in the intermediate field regime \\cite{Ruder94}. In order to facilitate a proper understanding of the spectra of neutron stars and white dwarf stars one must necessarily have more stringent bounds on the energy levels of atoms in the atmospheres of these compact objects in the intermediate regime of magnetic field strengths. This is the aim of the current work. The work described herein extends previous work \\cite{HH1998} and presents a numerical treatment of hydrogen and helium atoms in magnetic fields, yielding accurate results for the eigenvalues and eigenvectors of the first few low-lying states over a wide range of field strengths in the intermediate field regime. The calculated energy eigenvalues are seen to be improvements upon previous estimates. The procedures described herein do not make any assumptions of basis of functions and neither are they restricted to the adiabatic approximation. The direct and exchange interactions of the electrons are computed using a novel method and do not rely upon any approximations. The method is readily extendable to many-electron systems and arbitrary field strengths. The overall method is also computationally straightforward to implement. ",
        "conclusions": "Conclusions }  The work described here was motivated by the need to have accurately determined values for the upper bounds for the energy levels of atoms in strong magnetic fields. As was discussed earlier, this need has arisen due to the presence of strong magnetic fields in neutron stars and white dwarf stars. The most commonly present atoms in the atmospheres of these compact objects, hydrogen and helium were studied here with the intention of obtaining accurate estimates of the energy levels of the first few low lying states in strong magnetic fields. We described a method adopting a physically motivated approach governed by the inherent symmetries of the problem. We simultaneously circumvented the need for adopting a definite basis of functions to describe the wave functions of the electrons in either of the directions, parallel and perpendicular to the magnetic field. The approach is unrestrictive with regard to the wave function; it has the distinct advantage over methods that require a basis of functions to describe the wave functions because in numerical solutions one can only have a finite number of such functions. The wave functions determined in the present study came about naturally from the symmetries of the problem and are thus in effect superpositions of a large number of basis functions.  Such an approach however resulted in elliptical partial differential equations for the electrons that were subsequently solved using finite element techniques. It is to be noted that the novel method adopted for determining the direct and exchange interactions between the electrons in the helium atom is also exact in the sense that it does not rely upon any ab initio assumptions to approximate the integrals. These interaction potentials are solved in a natural manner by solving the elliptical partial differential equations, Eqs. (16) and (22). The eigenvalues found in the range of the magnetic field strength parameter \\begin{math}10^{-2} \\leq \\beta, \\beta_{Z} \\leq 10\\end{math} considered in this study were seen to be consistent with previous findings \\cite{Ruder94, Jones1996, Jones1999}. Rational functions were also used to find sufficiently accurate interpolating functions for the binding energies of various states of both the hydrogen and helium atoms in the range of magnetic fields considered herein. These were seen to be accuarate to (an average for all six fits) within 0.8\\%. Potentially such interpolating functions could be used in atmosphere models of neutron stars and white dwarf stars thus obviating the need for involved and laborious calculations of the same.  Thus, the current work describes an unrestricted and computationally less intensive method for calculating the energy levels of atoms in strong magnetic fields. There are in essence three directions in which the current work could be extended. First, the current work can be readily extended to higher magnetic field strengths by changing the domain of calculations appropriately to incorporate the fact that the electrons become more and more bound. Simultaneously, the calculations and the software developed as a part of this study are readily extendable to systems with more than two electrons, i.e., Li, Be, B, C, O etc. Finally, the procedures implemented herein can also be extended towards a multi-configuration framework \\cite{CFF1997}. In essence, the calculations employed herein are for a single configuration of the electrons. Thus having found the energies of the different configurations, it is possible to extend the theory to incorporate a multi-configuration approach which is likely to improve the results already obtained here."
    },
    "0806/0806.4739_arXiv.txt": {
        "abstract": "The muon charge ratio of the lateral muon density distributions in single EAS is studied by simulations, in context of recent proposals to measure this observable in coincidence with EAS observations. While effects of the hadronic interaction do not lead to significant differences of the total $\\mu^+$ and $\\mu^-$ content, the  differences of  the azimuthal variation of the muon densities of opposite charges  and the azimuthal  variation of the muon charge ratio appear to be very much pronounced, dependent on the direction of EAS incidence. This is due to the influence of the geomagnetic field which induces related effects in radio emission from extended air showers. ",
        "introduction": "Primary cosmic rays are dominantly high-energy protons, alpha particles and heavier nuclei  in a relative amount decreasing with the atomic number. When penetrating from the outer space into the Earth's atmosphere they initiate the development of a phenomenon called Extensive Air Showers (EAS) by multiple production of particles in cascading interactions of the primary particles with atmospheric nuclei. The produced secondary radiation establishes an essential feature of our natural environment. It affects material and biological substances and comprises a specific part of the natural radiation background. Its study is of relevance in various scientific problems. Photons, electrons and positrons are the most numerous secondary particles in an EAS event, and the muonic component contributes only to few percent in the single shower. However, at lower primary energies, which are dominating due to the steeply falling primary spectrum, or in case of very inclined showers the electromagnetic shower component gets completely absorbed during the travel through the atmosphere, while the muons (`penetrating component') survive the propagation  through  even larger  slant depths. Hence the inclusive secondary radiation flux in the atmosphere comprises mainly muons (c. 80\\% with c.100 muons s$^{-1}$m$^{-2}$sr$^{-1}$ on sea level). \\\\ The cosmic ray muons originate from decay of hadronic secondaries produced in particle cascades by primary cosmic rays: \\begin{equation*} \\pi^{\\pm} \\rightarrow \\mu^{\\pm} + \\nu_\\mu (\\overline{\\nu}_\\mu ) ~ ~ ~ 100.0 \\% ~ ~ ({\\rm mean~lifetime}~ 2.6 \\cdot 10^{-2} \\mu s) \\end{equation*} \\begin{equation*} K^{\\pm} \\rightarrow \\mu^{\\pm} + \\nu_\\mu (\\overline{\\nu}_\\mu ) ~ ~ ~ ~ 63.5 \\% ~ ~ ({\\rm mean~lifetime}~ 1.2 \\cdot 10^{-2} \\mu s) \\end{equation*} \\noindent The muons decay with a larger mean lifetime ($2.2 \\mu s$): \\begin{equation*} \\mu^+ \\rightarrow e^+ + \\nu_e + \\overline{\\nu}_\\mu \\end{equation*} \\begin{equation*} \\mu^- \\rightarrow e^- + \\overline{\\nu}_e + \\nu_\\mu \\end{equation*} The ratio of the flux of positive to negative muons, the so-called muon charge ratio $R_\\mu (\\mu^+ / \\mu^- )$ is a significant quantity which reflects important features of the hadronic meson  production in cosmic ray collisions~\\cite{ref1,ref2} and can help to discern the primary mass composition~\\cite{ref2,ref3}. \\\\  It is also immediately obvious that the  muon  flux  in  the atmosphere is strongly related to the neutrino flux and that the muon charge ratio \\begin{equation*} R_\\mu (\\mu^+ / \\mu^- ) \\sim R( \\nu_e / \\overline{\\nu}_e ) \\end{equation*} \\noindent provides relevant information for neutrino physics. The atmospheric neutrino measurements with Super-Kamiokande~\\cite{ref4} and other experiments~\\cite{ref5} have revealed that the ratio of muonic to electronic neutrinos is much smaller than the theoretical predictions, $R(\\nu_\\mu/\\nu_e)_{observed} / R(\\nu_\\mu/\\nu_e)_{predicted} \\ll 1$. The deficit was interpreted in terms of neutrino flavour oscillations, confirmed by the observed zenith angular dependence of the measured rates of $\\nu_\\mu$ and $\\nu_e$. \\\\  There are numerous studies of the charge ratio of the atmospheric muon flux, on sea level (see the compilations~\\cite{ref2,ref6,ref7,ref8} and also of the vertical dependence by balloon experiments e.g.~\\cite{ref9,ref10,ref11}. The experimental results provide a highly inclusive information since the atmospheric flux is produced by many different EAS, with primary energies distributed along the steeply falling energy spectrum and the mass composition of the primary flux. When impinging on our Earth's atmosphere they are additionally affected by the geomagnetic field. The influence of the geomagnetic field leads to a dependence of the muon charge ratio from the azimuth of the direction of observation~\\cite{ref12} (East-West effect), in particular for low energy muons, which are dominantly originating from EAS of lower primary energies. \\\\ The measured value of the muon charge ratio which is empirically found to have a value of about $1.25-1.30$, is mainly a result of the positive charge (proton) excess of the primary mass distribution, at $E_\\mu > 10\\,$GeV slightly increasing with  the energy of the observed muons~\\cite{ref8}. The MINOS detector in the Soudan mine has recently published precise results about the charge ratio of atmospheric muons in the TeV energy range and observed significantly higher values of $R_\\mu(\\mu^+/\\mu^-)$~\\cite{ref13}. The rise in the muon charge ratio is expected due to the increasing contribution of kaons to the cosmic ray muon flux and due to an enhanced contribution of the $K^+$ decay~\\cite{ref14,ref15}. \\\\  Modern theoretical approaches of the muon flux and charge ratio start from Monte Carlo simulations of a sufficient number of single EAS events calculated along the primary energy spectrum and chemical distribution of cosmic rays (see ref.~\\cite{ref1} e.g.). The Monte Carlo simulations invoke as generators specific models of the hadronic interactions, which are reflected by the so called energy-weighted moments~\\cite{ref14}. Therefore measurements of the charge ratio are a source of information about the validity of hadronic interaction models. In case of the atmospheric muons it has turned out that the results are rather stable against modifications of the models, except when higher energy muons would be considered. M.~Unger (L3 Collaboration) has put attention~\\cite{ref16} that for muons with $E_\\mu > \\approx 300\\,$GeV the charge ratio would be very discriminative, also due to the increasing influence of kaons. The features of the particular models are expected to be displayed  more distinctly, when the charge ratio of single showers is studied. Apparently there is a small  excess of positive pions already in the single collision process, which leads already in a single shower to a value of $R_\\mu(\\mu^+/\\mu^-) > 1$ (see~\\cite{ref15}). \\\\  In exclusive observation of single EAS which could be specified by a definite primary energy, the direction of incidence and eventually by the mass of the primary, the effects of the different types of the hadronic interactions and of the geomagnetic field are expected to get revealed in a more pronounced way. In particular the lateral distribution of the EAS particles displays an azimuthal variation~\\cite{ref17}, which is influenced by the geomagnetic field, differently for $\\mu^+$ and $\\mu^-$, and leading to an azimuthal variation of the charge ratio of the muon density distribution. \\\\  In this paper we address some aspects of the muon charge ratio of single EAS, especially of the azimuthal variations of the muon density distribution, worked out by simulation studies. Few and far between this topic has been addressed in the past~\\cite{ref18}. The considerations are in context of recent proposals for studies of the muon charge ratio in single EAS, which may be observed with standard detector arrays as used (see ref.~\\cite{ref19}) to study EAS. The effects depend on the direction of incidence and increase with the inclination of shower incidence (zenith angle $\\theta$), i.e. when the distance of muon travel is increased due to the $sec(\\theta)$-dependence. Recently, in view of the possibility to extract charge information of high energy EAS muons the geomagnetic effects on the shower development have been estimated in~\\cite{ref20,ref21} on basis of a modified  Heitler model~\\cite{ref22}, originally known as `toy-model'. The present study yields results on the lateral muon distributions resulting from extensive three-dimensional Monte Carlo simulations with reconstructions of the azimuthal variation of the muon component of inclined showers (with $\\theta=45^\\circ$ as example and with the primary energies of $10^{14}\\,$eV and $10^{15}\\,$eV). ",
        "conclusions": "EAS simulations show that the lateral density distributions of the positive and negative muons are varying not only with the (radial) distance from the shower axis, but also with the azimuth relative to the plane of the incident shower. The reasons are different. In addition to the attenuation effects of charged particles of inclined showers~\\cite{ref15} in the atmosphere by the variation of the traveling distances in the atmosphere, the geomagnetic field affects the travel of positive and negative muons in an opposite way. The geomagnetic effects depend on the direction of the EAS axis relative to the Earth's magnetic vector. This leads to an azimuthal variation of the muon charge ratio of the muon density distribution, which has to be regarded  in context of analyses of experimental data. In the extreme case of very inclined showers (with long slant depths) the Earth magnetic field might be used as magnetic separator~\\cite{ref20,ref21}, at least for muons in the GeV range. Obviously the experimental detection of these features is of great interest for the understanding of the EAS development. Furthermore the quantitative results would also provide some information about the hadronic interaction, in particular when observing higher energy muons. The aspect of the dependence on the hadronic interaction models, currently en vogue, has not been systematically explored in the present paper.  {\\ack We are indebted to R. Engel, T. Huege, D. Heck and A.W. Wolfendale for clarifying discussions about different aspects of the topic. We thank J. Oehlschl\\\"ager and A. Patrascioiu for substantial help in achieving the reported results. These studies have been prompted by the experimental plans of the Cosmic Ray Group of NIPNE, Bucharest, we are collaborating with in this field.}"
    },
    "0806/0806.3605_arXiv.txt": {
        "abstract": "We present an analysis of a previously unpublished radio-continuum observation of \\SNR\\ which at an age of \\ $\\la150$ years is the youngest known in the Galaxy. The observations were made in 1993  using the Australia Telescope Compact Array (ATCA) at two 6-cm frequencies. We note two previously unseen blow-out structures in the north and south direcions. We estimate a flux density of 1.545~Jy, an outer diameter of $\\sim80\\arcsec$ and confirm an expansion  rate of $\\sim0.65\\%$ per year between 1985 and 2008. No polarisation was detected in the radio emission from \\SNR\\ above the 1\\% level. We also present these previously unpublished results as a high resolution reference point from which to study the evolution of SNRs at times for which there is a gap in our knowledge. ",
        "introduction": "It is widely accepted that current catalogues have a distinct deficit of young Galactic Supernova Remnants (SNRs), that is, SNRs $<$2000 years old, with only $\\sim$10 confirmed out of a predicted $\\sim$50 \\citep{1991ARA&A..29..363V, 2003LNP...598...37C}. Of these confirmed SNRs, \\SNR\\ is of particular interest as it is believed to be the youngest in the Milky Way at $\\la150$ years old \\citep{2008ApJ...680L..41R,2008MNRAS.387L..54G}.  Originally identified as a probable SNR by \\citet{1984Natur.312..527G} at 4.9~GHz using the Very Large Array (VLA), \\SNR\\ was described as a shell source with an approximate brightness slightly less than that of the Tycho and Kepler SNRs with a spectral index of $\\sim$--0.7\\footnote{Spectral index defined as $S\\propto\\nu^\\alpha$}. Using the Molonglo Observatory Synthesis Telescope (MOST) Galactic Survey data, \\citet{1994MNRAS.270..835G} confirmed the classification of \\SNR\\ as an SNR: the source was described as featuring a shell-like morphology in the radio with an estimated diameter of 1\\farcm2.  Later \\citet{2000AJ....119..207L} produced a 90-cm image of \\SNR\\ made with the VLA: those authors estimated the 20/90-cm spectral index of the SNR to be $-0.93\\pm0.25$ and the angular  diameter to be 1\\farcm1. \\citet{2004AJ....128.1646N} revisited the data collected by \\citet{2000AJ....119..207L} and -- through the application of superior data reduction techniques -- measured the diameter of \\SNR\\ to be $<$1\\arcmin. Finally, \\citet{2004BASI...32..335G} estimates the diameter of \\SNR\\ based on 1.49~GHz VLA observations made in 1985 to be of 1\\farcm2. In short, multiple radio observations have confirmed that \\SNR\\ has the smallest angular diameter for a known Galactic SNR, indicative of its young age.  Earlier this year \\citet{2008MNRAS.387L..54G} reobserved \\SNR\\ at 4.86~GHz using the VLA after \\citet{2008ApJ...680L..41R} used recent {\\it Chandra} images to show \\SNR\\ had expanded significantly since 1985 and it's X-ray emission appeared  to be purely synchrotron radiation. By comparing these new VLA observations with the 1985 VLA observations made at 1.49~GHz \\citet{2008MNRAS.387L..54G} determined that \\SNR\\ had expanded by $15\\%\\pm2\\%$ over 23 years. Most recently, \\citet{2008arXiv0806.1952M} found that \\SNR's flux density at 843~MHz increased by $1.22\\pm ^{0.24}_{0.16}$ \\% per year over the last two decades.  In this paper we present our results of a previously unpublished 6-cm radio-continuum observation of \\SNR\\ made in 1993 that was found on an examination of the Australia Telescope Compact Array (ATCA) database. It must be emphasised that this radio observation has the highest angular resolution of any other radio image of \\SNR\\ yet presented and does not suffer from the distorted beam shapes that are inherent in observations made using the VLA when observing objects such as this, that are at low Declinations.  We report on the morphology, expansion rate, age, diameter and polarisation properties of this SNR. ",
        "conclusions": "We presented a 6-cm ATCA observation made in 1993 of SNR \\SNR. We used this observation to confirm the estimates of expansion made by  \\citet{2008ApJ...680L..41R} and \\citet{2008MNRAS.387L..54G}. We identified three structures that compose the ring/shell of the SNR as well as two previously unseen possible blow-outs in the north and south directions. We were not able to detect any significant polarisation above the 1\\% level.  The observation presented here (and the accompanying results) are significant in that, unlike most other  previously published radio observations of this SNR, this observation was made from a radio telescope located in the southern hemisphere: this geographic location is better suited for observations of low Declination sources (like \\SNR) than telescopes located in the northern hemisphere. In fact, other published radio observations of this SNR that have been made with the VLA will suffers from distorted beam shapes that are inherent in VLA observations of objects located at low Declinations. In addition, this observation helps to add to our knowledge of this particular SNR (which is now known to be extremely young) which in turn will help us constrain models of the early evolution of SNRs.\"  We would like to emphasise the need for new higher resolution observations with better $uv$ coverage in order to obtain a more complete understanding of \\SNR's morphology (ie. possible blow-outs), as well as to more accurately determine the expansion rate, age and study its magnetic field structure."
    },
    "0806/0806.4163_arXiv.txt": {
        "abstract": "We present high spatial resolution Submillimeter Array observations and supplementary single-dish photometry of the molecular gas and dust around IRAS 04158+2805, a young source with spectral type M5-M6 in the Taurus star-forming region.  A bright, highly elongated dust structure that extends 8\\arcsec\\ ($\\sim$1120\\,AU) in diameter is revealed in a 883\\,$\\mu$m thermal continuum image.  This emission geometry is in good agreement with optical observations that show a similar structure in absorption, aligned perpendicular to bipolar scattered light nebulae.  However, the interferometric data also clearly demonstrate that the submillimeter continuum emission is not centrally concentrated, but rather appears to have a toroidal geometry with substantially lower intensities inside a radius of $\\sim$250-300\\,AU.  Spatially resolved emission from the CO $J$=3$-$2 transition exhibits a velocity gradient along the major axis of the dust structure.  If this kinematic pattern is interpreted as the signature of rotation around a central object, a relatively low mass is inferred ($M_{\\ast} \\sim 0.3$\\,M$_{\\odot}$, with a $\\sim$50\\% uncertainty).  We discuss several possible explanations for the observed gas and dust environment around IRAS 04158+2805, including a flattened envelope with an outflow cavity and a large circumbinary ring.  This source offers unique views of the gas and dust environment surrounding a young low-mass stellar system.  Its properties are generally not commensurate with formation scenarios for such low-mass objects that rely on dynamical ejection, but rather confirms that a single mechanism $-$ molecular cloud core collapse and fragmentation $-$ can produce stars over a wide range of stellar masses (at least an order of magnitude). ",
        "introduction": "Circumstellar material profoundly influences the star formation process.  A large-scale envelope acts as the local mass reservoir during the collapse and growth of a central protostar; a more compact disk regulates how that material is transported onto the star itself.  A great deal of progress has been made in understanding the physical conditions present in the gas and dust surrounding Sun-like stars ($M_{\\ast} \\approx 0.5$-3\\,M$_{\\odot}$).  However, it is unclear if those results can be extrapolated across the entire stellar mass spectrum. The low end of this spectrum ($M_{\\ast} \\le 0.3$\\,M$_{\\odot}$) is of particular interest, as observations of the material around young low-mass objects\\footnote{For simplicity, we will refer to objects with $M_{\\ast} \\lesssim 0.3$\\,M$_{\\odot}$ as low-mass objects, including both very low-mass stars and brown dwarfs.} can help resolve an ongoing debate over their dominant formation mechanism.  The observed abundance of such low-mass objects relative to their more massive counterparts is difficult to account for with the traditional model for isolated star formation.  Two distinct modifications have been proposed: ($a$) the turbulent fragmentation of cloud cores into smaller building units \\citep[e.g.,][]{padoan04}; and ($b$) a dynamical alternative where one component of a multiple system is ejected and prematurely cut off from its accretion reservoir \\citep[e.g.,][]{reipurth01,umbreit05}.  Both scenarios make distinct and conflicting predictions about the gas and dust environments that should be associated with low-mass objects.  In the former, a disk/envelope structure similar to those noted around higher mass T Tauri stars would be expected.  And for the latter, only a truncated disk (with an initial $R_d \\lesssim 10$\\,AU) would survive the tidal stripping and ejection process \\citep{bate03}.  If the remnant material is sufficiently viscous, it could spread to radii up to $\\sim$100\\,AU in $\\sim$1\\,Myr \\citep{armitage97}, but would have shed a vast majority of its initial mass in the ejection process. Therefore, observational constraints on the structure of the gas and dust around low-mass objects offer one avenue to help distinguish how they form.  A variety of observations demonstrate that the signatures of circumstellar gas and dust at small radii (up to a few AU) are indeed common for young low-mass objects \\citep[see the recent review by][]{luhman07b}.  More detailed individual studies confirm that such material has a geometry and composition similar to the disks around more massive T Tauri stars \\citep{natta01,pascucci03,allers06,buoy08}, although typically at lower masses \\citep{klein03,scholz06}.  However, these {\\it unresolved} observations can not unambiguously constrain some key properties of this material, most significantly its spatial structure and extent.  In this article, we present new submillimeter observations of the molecular gas and dust surrounding IRAS 04158+2805.  Located in the $\\sim$1\\,Myr-old Taurus star-forming region, \\iras has a cool central source with an estimated spectral type of M5-M6 \\citep{white04,luhman06,beck07}.  Recently, \\citet{glauser08} provided an initial analysis of this source and the dust structure surrounding it by modeling the broadband spectrum along with optical/infrared scattered light images.  They suggested that their observations are best explained by a large (diameter of 2240\\,AU), massive (dust mass of 1-2$\\times$10$^{-4}$\\,M$_{\\odot}$; gas/dust ratio of $220^{+150}_{-170}$) circumstellar disk around a low-mass star ($M_{\\ast} \\approx 0.1$-0.2\\,M$_{\\odot}$), with no need to include any contribution from a more extended envelope.  Our new data represent the first high spatial resolution view of the gas and dust environment around this and similar cool, young objects at submillimeter wavelengths.  They also provide an opportunity for a rare, albeit crude, estimate of $M_{\\ast}$ from the spatio-kinematics of circumstellar gas that is independent of pre-main-sequence evolution models. These observations are addressed in \\S 2, and the resulting data products are highlighted in \\S 3.  In \\S 4, we discuss the structure of the gas and dust surrounding \\iras and what it can tell us about the central source and the formation of low-mass objects in general.  The results are summarized in \\S 5. ",
        "conclusions": "\\subsection{The Central Object}  Estimating the masses of young stars and brown dwarfs generally requires reference to theoretical models of their structural evolution.  The quantitative reliability of such models for individual stars at ages $\\lesssim$1\\,Myr is highly uncertain, particularly for the low end of the stellar mass spectrum where complicated convection physics, dusty atmospheres, and uncertain molecular opacities and initial conditions can strongly affect the observational diagnostics \\citep[e.g.,][]{dantona94,baraffe02,montalban04}.  Dynamical constraints on $M_{\\ast}$ from the orbital properties of either a companion star \\citep[e.g.,][]{mathieu94} or gas in a circumstellar disk \\citep[e.g.,][]{simon00} provide an extremely valuable independent check that can potentially be used to calibrate these models \\citep[see][]{hillenbrand04}.  \\citet{stassun06} report the only such measurements of young $\\sim$M6 dwarfs to date, for an eclipsing binary system in Orion.  With the first spatially resolved measurements of molecular gas in apparent rotation around a young low-mass object with similar spectral type, the CO observations of \\iras presented here can provide another such estimate of $M_{\\ast}$, although admittedly with a great deal more uncertainty.  Figure \\ref{pv_spec} shows the position-velocity diagram of the CO $J$=3$-$2 line emission for IRAS 04158+2805, where the angular offset is the distance from the observed phase center (such that positions to the east have positive values) and the velocity offset is taken relative to the systemic value ($V_{\\ast} = 7.4$\\,km s$^{-1}$).  Overlaid on the diagram is an inclined ($i = 62$\\degr) Keplerian rotation profile that matches the general kinematic trend in the data rather well (thick curve), with a central point mass of $M_{\\ast} = 0.30$\\,M$_{\\odot}$.  Additional rotation profiles for $M_{\\ast} = 0.15$ and 0.45\\,M$_{\\odot}$ are also shown for reference.  It should be noted that these profiles assume pure Keplerian orbital motion around a central source, which may not be the case in the more complex \\iras environment.  For now, we assume that the dominant kinematic trend is such rotation, but potential complications will be addressed below.  A comparison of this dynamical $M_{\\ast}$ estimate with predictions from stellar evolution models is a challenge because of uncertainties in the luminosity (due to scattered light contamination at short wavelengths), extinction, and effective temperature for the central source.  Recent analysis of the scattered light spectrum from \\iras suggest that this source is cool, with initial spectral type estimates of M6$\\pm1$ \\citep{white04,beck07} refined to M5.25$\\pm$0.25 (optical) and M6$\\pm$0.5 \\citep[infrared;][]{luhman06}.\\footnote{Note that previous spectral type assignments ranged from $\\sim$K7 to M3 \\citep{kenyon98,luhman99}.}  Using the latter classifications and the \\citet{luhman03} empirical effective temperature scale, we estimate $T_{\\ast} = 3050\\pm75$\\,K.  Unfortunately, large uncertainties in the extinction, where estimates range from $A_V \\approx 9$-16 \\citep[e.g.,][]{white04,beck07}, make a luminosity determination difficult. However, we can use pre-main-sequence models for an assumed age in the inferred temperature range to estimate a central mass.  Using the $T_{\\ast}$ range quoted above and an age of $\\sim$1\\,Myr, the \\citet{dantona97} and \\citet{baraffe98} models indicate $M_{\\ast} \\approx 0.09$-0.16\\,M$_{\\odot}$. This mass range does not include the additional uncertainties on the adopted effective temperature scale or the unknown age of the source.  The $M_{\\ast}$ values inferred from the pre-main-sequence models lie roughly a factor of 2 below the nominal value that best describes the kinematics of the CO gas inferred above.  The unquoted and poorly understood uncertainties in both measurements may be the simple explanation for this apparent discrepancy. Without a detailed physical model for the local gas structure around IRAS 04158+2805, it is difficult to estimate a statistically acceptable range of $M_{\\ast}$ from the CO kinematics; there may very well be overlap with the higher end of the $M_{\\ast}$ range from the pre-main-sequence models.  The effective temperature scale adopted above is also uncertain.  A $\\sim$500\\,K increase in the effective temperature (to a value that is typically associated with M2 stars) would reconcile the $M_{\\ast}$ estimates.  Indeed, using an independent method of fitting infrared spectra, \\citet{doppmann05} infer $T_{\\ast} = 3500$\\,K for the \\iras central source, leading to $M_{\\ast} \\approx 0.35$-0.45\\,M$_{\\odot}$ in the aforementioned models.  Moreover, \\citet{hillenbrand04} have shown that pre-main-sequence models generally tend to underpredict $M_{\\ast}$ compared to dynamical measurements in this low-mass range, although typically only by $\\le 20$\\%.  In their recent study that explicitly includes the effects of scattered starlight, \\citet{glauser08} argue for a very luminous central source, $L_{\\ast} \\approx 0.4$\\,L$_{\\odot}$.  While this value must be somewhat degenerate with the dust properties and structure assumed in their model, it is clearly the most robust estimate available.  The pre-main-sequence models imply that such a luminous cool source would be very young, with an age significantly less than 1\\,Myr.  The model mass tracks at such ages are not well understood in light of the uncertain initial conditions.  Clearly, the aforementioned uncertainties can be solely responsible for the apparently different $M_{\\ast}$ estimates from the CO kinematics and pre-main-sequence models.  Alternatively, there is also an appealing and simple physical explanation $-$ the \\iras central source may be a roughly equal-mass binary.  In this scenario, the $M_{\\ast}$ estimate from the CO kinematics refers to the total mass of the stellar system (i.e., the mass interior to the gas that produces the line emission), while the estimate from the pre-main-sequence models assumes only a single star.  Therefore we would expect to see both a high luminosity \\citep[as claimed by][]{glauser08} and $M_{\\ast}$(CO) $\\sim 2 M_{\\ast}$(models).  A variety of surveys suggest that such binaries are fairly common, $\\sim$35-45\\% for M dwarfs \\citep{fischer92,reid97} and $\\sim$10-30\\% for cooler objects \\citep[e.g.,][]{burgasser07}.  While it is difficult to definitively state the properties of the central source(s), we can further assess the nature of \\iras by analyzing the structure of the material that surrounds it.   \\subsection{The Gas and Dust Environment}  Given the elongated structures shown in Figure \\ref{images}, it is reasonable to suggest that the gas and dust around \\iras reside in an exceptionally large circumstellar disk.  \\citet{glauser08} have recently demonstrated that they are able to fit high-quality scattered light images and the SED for this source with a simple disk structure, and without the need to invoke any envelope component.  With the new data presented here, it is worthwhile to re-examine the structure of this material.  \\iras is exceptionally bright at submillimeter wavelengths, with an 850\\,$\\mu$m luminosity larger than $\\sim$85\\% of all other sources detected in the \\citet{aw05} Taurus survey.  Breaking that sample down further, the \\iras submillimeter emission is brighter than $\\sim$90\\% of Class II sources (disk only), but is comparable to the median Class I source (disk + envelope).  If we adopt the standard emissivity for disks (0.034\\,cm$^2$ g$^{-1}$ at 883\\,$\\mu$m, including a gas-to-dust ratio of 100) and assume that the submillimeter continuum emission is optically thin and has a characteristic temperature of $\\sim$20\\,K \\citep[e.g.,][]{beckwith90,aw05,aw07b}, the integrated SMA flux density suggests a large total mass of gas and dust is present around IRAS 04158+2805, $\\sim$0.03\\,M$_{\\odot}$.  This value is in good agreement with the disk mass range inferred by \\citet{glauser08}, $M_d \\approx 0.02$-0.04\\,M$_{\\odot}$.  Estimating masses for circumstellar material in this way is inherently uncertain (perhaps by an order of magnitude), due to the challenge of observationally constraining the optical properties of dust grains and the assumed mass conversion of a trace constituent (dust) to the presumably dominant species (molecular gas).  Perhaps more compelling (and easier to interpret) is the unusual geometry that is observed for the dust structure around IRAS 04158+2805.  The enormous size of the scattered light nebulae noted by \\citet[][2240\\,AU diameter]{glauser08} and the submillimeter continuum emission presented here (1120\\,AU diameter) is exceptionally rare for circumstellar disks \\citep[e.g.,][]{pietu07,watson07,aw07,hughes08}, rivaled only by the circumbinary disk around UY Aur \\citep{duvert98,close98,potter00}.  On the contrary, the observed size is fairly typical (if not on the small end) for larger-scale envelopes \\citep[e.g.,][]{looney00,eisner05}.  Moreover, our resolved observations of the submillimeter continuum reveal that a large central region (perhaps $\\sim$500\\,AU in diameter) has significantly diminished intensity at 883\\,$\\mu$m, suggesting a ring-like or toroidal geometry for the dust structure.  Similar geometries have been inferred for a handful of circumstellar disks, and a variety of underlying causes are possible.  The so-called ``transition\" disks have their inner regions (out to a few tens of AU in radius) largely cleared of observable material, presumably by disk evolution processes like particle growth, photoevaporation, or dynamical interactions with a young planetary system \\citep[e.g.,][]{pietu06,hughes07,brown08}.  However, given the large spatial scale of the central depression and the absence of a corresponding dip in the infrared part of the \\iras SED, as shown in Figure \\ref{sed}, this scenario is unlikely.  Alternatively, extremely dense disks with edge-on orientations can potentially render even the submillimeter emission optically thick, resulting in a central depression in its brightness distribution \\citep[e.g.,][]{wolf08}.  Although the \\iras dust structure is not oriented edge-on, perhaps longer pathlengths through the exceptionally large structure can provide compensating column densities.  If the bulk of the submillimeter emission is optically thick, the spectrum should have a spectral index $\\alpha = 2$, where $F_{\\nu} \\propto \\nu^{\\alpha}$ \\citep{beckwith91}.  Power-law fits to the submillimeter photometry are significantly steeper than the optically thick case, with $\\alpha = 2.7\\pm0.3$ (0.45-1.3\\,mm; $\\alpha = 3.8\\pm0.6$ for 0.85-1.3\\,mm).  At these wavelengths, such a steep spectrum is only noted for a few percent of Class II sources in Taurus and $\\rho$ Oph, but is more common (25-35\\%) in the envelopes around Class I objects \\citep{aw05,aw07b}.  The observed steep submillimeter SED would be difficult to produce unless a large fraction of the dust was optically thin at those wavelengths.  As noted by \\citet{wolf08}, a clear test of this possibility can be made with resolved continuum observations at $\\lambda > 883$\\,$\\mu$m.  In such data, the material should become optically thick only at significantly smaller radii than at 883\\,$\\mu$m, resulting in a more centrally concentrated emission profile (or, equivalently, the null in the visibilities should be detected at a comparatively longer deprojected baseline).  Note that if the material inside this central depression is optically thick at 883\\,$\\mu$m, an enormous reservoir of mass is not accounted for in the above estimate.  A third possibility for explaining the submillimeter continuum depression is that the emitting dust resides in a circumbinary ring.  The material around close binary stars is expected to be dynamically cleared on scales similar to the physical separation of the stellar components \\citep[e.g.,][]{jensen96,guilloteau99}.\\footnote{Recent observations suggest that the circumbinary scenario may be the more appropriate explanation in some cases for the diminished central emission in both transition \\citep{ireland08} and dense, edge-on disks \\citep{guilloteau08}.}  Numerical simulations of this dynamical interaction suggest that a circumbinary disk gap would be cleared out to a radius $\\sim$1.5-3$\\times$ the projected semimajor separation of the stellar components, and the individual circumstellar disks would be truncated at a radius $\\sim$0.2-0.5$\\times$ the projected semimajor separation \\citep[e.g.,][]{artymowicz94}.  If the 883\\,$\\mu$m continuum depression noted here is caused by such clearing, we would expect the projected binary separation to be $\\sim$90-180\\,AU ($\\sim$0\\farcs6-1\\farcs3) for a reasonable range of eccentricities, and the remnant individual circumstellar disks might be truncated at radii of $\\sim$20-90\\,AU.  This configuration is similar to the UY Aur system, where a ring with $\\sim$20\\arcsec\\ diameter is detected in both scattered light and CO around a roughly equal-mass binary with a similar projected separation \\citep{duvert98,close98,potter00}.  The primary difference is in the masses involved; the stellar mass in the \\iras system is $\\sim$4$\\times$ less than for UY Aur, but the circumstellar mass (or rather 850\\,$\\mu$m luminosity) is $\\sim$4$\\times$ higher.  While this possibility is certainly appealing in its reconciliation of both the information on the central source (see \\S 4.1) and the structure of the gas and dust emission, testing the multiplicity of \\iras will be a challenge given the obscuration of the central region by the dust structure.  Decomposing high resolution spectra of scattered starlight may represent the best opportunity in this case.  The \\iras SED, shown in Figure \\ref{sed}, exhibits a strong infrared excess with a flat or slightly rising slope from 2-25\\,$\\mu$m.  A variety of solid-state absorption features are noted in higher resolution infrared spectra, including H$_2$O ice (3\\,$\\mu$m; not shown), silicates (10\\,$\\mu$m), and CO$_2$ ice \\citep[15\\,$\\mu$m;][]{beck07,furlan08}.  The bright and steep submillimeter spectrum has already been discussed above.  While all of these SED properties are not necessarily inconsistent with a large, cold, highly inclined circumstellar disk \\citep[e.g.,][]{menshchikov97,chiang99,pontoppidan05}, taken together they are more common for Class I/Flat-Spectrum objects that harbor a remnant accretion envelope \\citep[e.g.,][]{whitney03,watson04,boogert04,pontoppidan07}. \\citet{furlan08} demonstrate that the \\iras SED features are strikingly similar to those for several other Class I sources in Taurus (see their Fig.~2) and can be successfully reproduced in detail with an envelope model (although the new observations of this source warrant a modification of their geometric assumptions).  So, all of the observational evidence presented here is also consistent with a flattened envelope structure around IRAS 04158+2805.  Such a flattened envelope geometry is predicted by molecular cloud core collapse models that incorporate rotation or magnetic fields \\citep{terebey84,galli93}, and have been clearly detected in other cases, albeit on significantly larger spatial scales than noted here \\citep[e.g.,][]{looney07}.  In this scenario, a partially evacuated outflow cavity at the center of the flattened envelope may explain the central depression in the submillimeter continuum data \\citep[e.g.,][their Models 3 or 4]{whitney03}. \\citet{glauser08} indicate that \\iras drives a substantial H$\\alpha$ jet oriented in the north-south direction, perpendicular to the observed dust structure.  The velocity-integrated CO $J$=3$-$2 morphology does resemble the bases of some molecular outflows \\citep[e.g.,][]{arce06}, but the absence of both a velocity gradient in the proposed flow direction (north-south) and larger scale molecular outflow signatures \\citep{bontemps96,gomez97} are difficult to reconcile with that interpretation.  Nevertheless, some contribution to the CO emission from an outflow and even the envelope (i.e., motions unrelated to Keplerian rotation) would complicate the observed spatio-kinematics.  The dynamical estimate of $M_{\\ast}$ discussed above should be treated with caution with regards to this possibility.  Of course, the envelope and binary ideas are not mutually exclusive, and dynamical clearing of the inner envelope and disks is still a viable alternative.  The images in Figures \\ref{images}c and \\ref{chanmaps} demonstrate that, unlike the submillimeter continuum, the CO emission is centrally concentrated; most of the line emission is from radii inside the toroidal structure traced in the continuum.  While this may suggest that the bulk of the line emission is generated in the disk(s) interior to that structure, it is difficult to provide a reliable origin without a more complete physical model for the \\iras environment.  Future studies of this source should focus on searching for definitive molecular outflow signatures, resolving dense gas tracers that are commonly associated with envelopes \\citep[e.g.,][]{jorgensen07}, observing CO transitions at higher spectral resolution to provide a better tracer of the velocity field, and reconciling the optical/near-infrared scattered light observations with the submillimeter results in an effort to converge on a consistent model of the local gas and dust structure.  Considering the low stellar mass range implied for IRAS 04158+2805, it is difficult to reconcile the large mass reservoir and size of this environment with any formation scenario that relies on it being dynamically ejected from a more massive system \\citep[e.g.,][]{reipurth01,bate03}.  Rather, the observed gas and dust structure is more consistent with the standard scenario, where an accretion disk and perhaps a remnant envelope developed around a single (or binary) source during the gravitational collapse of a molecular cloud core fragment.  Regardless of the precise value of $M_{\\ast}$, \\iras is a remarkable example of the fact that the standard picture for star formation worked out for higher-mass stars is also applicable at the low end of the stellar mass function."
    },
    "0806/0806.3096_arXiv.txt": {
        "abstract": "The cosmic microwave background (CMB) is a rich source of cosmological information. Thanks to the simplicity and linearity of the theory of cosmological perturbations, observations of the CMB's polarization and temperature anisotropy can reveal the parameters which describe the contents, structure, and evolution of the cosmos. Temperature anisotropy is necessary but not sufficient to fully mine the CMB of its cosmological information as it is plagued with various parameter degeneracies. Fortunately, CMB polarization breaks many of these degeneracies and adds new information and increased precision. Of particular interest is the CMB's B-mode polarization which provides a handle on several cosmological parameters most notably the tensor-to-scalar ratio, $r$, and is sensitive to parameters which govern the growth of large scale structure (LSS) and evolution of the gravitational potential. These imprint CMB temperature anisotropy and cause E-to-B-mode polarization conversion via gravitational lensing. However, both primordial gravitational-wave- and secondary lensing-induced B-mode signals are very weak and therefore prone to various foregrounds and systematics. In this work we use Fisher-matrix-based estimations and apply, for the first time, Monte-Carlo Markov Chain (MCMC) simulations to determine the effect of beam systematics on the inferred cosmological parameters from five upcoming experiments: PLANCK, POLARBEAR, SPIDER, QUIET+CLOVER and CMBPOL. We consider beam systematics which couple the beam substructure to the gradient of temperature anisotropy and polarization (differential beamwidth, pointing offsets and ellipticity) and beam systematics due to differential beam normalization (differential gain) and orientation (beam rotation) of the polarization-sensitive axes (the latter two effects are insensitive to the beam substructure). We determine allowable levels of beam systematics for given tolerances on the induced parameter errors and check for possible biases in the inferred parameters concomitant with potential increases in the statistical uncertainty. All our results are scaled to the `worst case scenario'. In this case, and for our tolerance levels the beam rotation should not exceed the few-degree to sub-degree level, typical ellipticity is required to be 1\\%, the differential gain allowed level is few parts in $10^{3}$ to $10^{4}$, differential beam width upper limits are of the sub-percent level and differential pointing should not exceed the few- to sub-arc sec level. ",
        "introduction": "The standard cosmological model accounts for a multitude of phenomena occurring over orders of magnitude of length and angular scales throughout the entire history of cosmological evolution. Remarkably, doing so only requires about a dozen parameters. Perhaps one of the most useful cosmological probes is cosmic microwave background (CMB) temperature anisotropy whose physics is well understood. Complementary cosmological probes can assist in breaking some of the degeneracies inherent in the CMB and further tighten the constraints on the inferred cosmological parameters. Temperature anisotropy alone cannot capture all the cosmological information in the CMB, and its polarization probes new directions in parameter space. B-mode polarization observations are noise-dominated but the robust secondary signal associated with gravitational lensing, which is known up to an uncertainty factor of two on all relevant scales, is at the threshold of detection by upcoming CMB experiments. The lensing signal may have been detected already through its signature on the CMB anisotropy as reported recently by ACBAR (Reichardt et al. [1]). Lensing by the large scale structure (LS) also converts primordial E-mode to secondary B-mode. When high fidelity B-mode data are available a wealth of information from the inflationary era (Zaldarriaga \\& Seljak [2], Kamionkowski, Kosowsky \\& Stebbins [3]), and cosmological parameters that control the evolution of small scale density perturbations (such as the running of the spectral index of primordial density perturbations, neutrino mass and dark energy equation of state), will be extracted from the CMB. At best, B-mode polarization from lensing is a factor of three times smaller than the primordial E-mode polarization, so it is prone to contamination by both astrophysical foregrounds and instrumental systematics. It is mandatory to account for, and remove when possible, all sources of spurious B-mode in analyzing upcoming CMB data, especially those generated by temperature leakage due to beam mismatch, since temperature anisotropy is several orders of magnitude larger than the expected B-mode level produced by lensing.  Beam systematic have been discussed extensively (Hu, Hedman \\& Zaldarriaga [4], Rosset et al. [5], O'Dea, Challinor \\& Johnson [6], Shimon et al. [7]). All the effects are associated with beam imperfections or beam mismatch in dual beam experiments, i.e. where the polarization is obtained by differencing two signals which are measured simultaneously by two beams with two orthogonal polarization axes. Fortunately, several of these effects (e.g. differential gain, differential beam width and the first order pointing error - `dipole'; Hu, Hedman \\& Zaldarriaga [4], O'Dea, Challinor \\& Johnson [6], Shimon et al. [7]) are reducible with an ideal scanning strategy and otherwise can be cleaned from the data set by virtue of their non-quadrupole nature which distinguishes them from genuine CMB polarization signals. Other spurious polarization signals, such as those due to differential ellipticity of the beam, second order pointing errors and differential rotation, persist even in the case of ideal scanning strategy and perfectly mimic CMB polarization. These represent the minimal spurious B-mode signal, residuals which will plague every polarization experiment. We refer to them in the following as `irreducible beam systematics'. We assume throughout that beam parameters are spatially constant. Two recent works (Kamionkowki [8] and Su, Yadav \\& Zaldarriaga [9]) considered the effect of spatially-dependent systematic beam-rotation and differential gain, respectively. This scale-dependence and the associated new angular scale induce non-trivial higher order correlation functions through non-gaussianities which can be both used to optimally remove the space-dependent component of beam rotation [8] and mimic the CMB lensing signal, thereby biasing the quadratic estimator of the lensing potential [9].  To calculate the effect of beam systematics we invoke the Fisher information-matrix formalism as well as Monte Carlo simulations of parameter extraction, the latter for the first time. Our objective is to determine the susceptibility of the above mentioned, and other, cosmological parameters to beam systematics. For the Fisher-matrix-based method and the Monte Carlo simulations we calculate the underlying power spectrum using CAMB (Lewis, Challinor \\& Lasenby [10]). The Monte-Carlo simulations are carried out with COSMOMC (Lewis \\& Bridle [11]). We represent the extra noise due to beam systematics by analytic approximations (Shimon et al. [7]) and include lensing extraction in the parameter inference process, following Kaplinghat, Knox \\& Song [12] and Lesgourgues et al. [13] (see also Perotto et al. [14] for the Monte Carlo simulations) for neutrino mass (and other cosmological parameters) reconstruction from CMB data.  This paper illustrates the effect of beam systematics and its propagation to parameter estimation and error forecasts for upcoming experiments. Our main concern is the effect on the following cosmological parameters: the tensor-to-scalar ratio $r$, the total neutrino mass $M_{\\nu}$ (assuming three degenerate species), tilt of the scalar index $\\alpha$, dark energy equation of state $w$, and the spatial curvature, $\\Omega_{k}$. The lensing-induced B-mode signal is sensitive to all parameters (except the tensor-to-scalar ratio) and peaks at few arcminute scales, while the tensor-to-scalar ratio depends on the energy scale of inflation and the primordial signal peaks at the characteristic horizon size at last scattering, $\\approx 2^{\\circ}$. We note that while the LSS-induced and primordial tensor power B-mode spectra are sub-$\\mu K$ the {\\it shape} of the primordial B-mode spectrum is known (only its amplitude is unknown, Keating [15]) and the secondary LSS-induced B-mode is guaranteed to exist by virtue of the known existence of LSS and E-mode polarization.   The paper is organized as follows. We describe the formalism of beam systematics for general non-gaussian beams and provide a cursory description of a critical tool to mitigate polarization systematics -a half wave plate (HWP), in section 2. The effect of lensing on parameter extraction within the standard quadratic-estimators formalism is discussed in section 3. The essentials of the Fisher matrix formalism are given in section 4 as well as some details on the Monte Carlo simulations invoked here. Our results are described in section 5 and we conclude with a discussion of our main findings in section 6. ",
        "conclusions": "The purpose of this work was to illustrate the effect of beam systematics on parameter extraction from CMB observations. Beam systematics are expected to be significant especially for detecting the B-mode polarization. Ongoing and future experiments must meet very challenging requirements at the experiment design and data analysis phases to assure polarimetric fidelity. Ultimately, a major target of these experiments is the most accurate estimation of cosmological parameters, and for this end it is mandatory to assess, among other issues, the propagation of beam systematics to parameter estimation. The tolerance levels chosen in this work are somewhat arbitrary and may be changed at will, according to the goals of individual experiments, and the numerical values we quote in the tables should be viewed in this perspective.  The only similar work so far to set tolerance levels on beam mismatch in the context of parameter estimation is O'Dea, Challinor \\& Johnson [6] which influenced our present work. However, we expand on this work in several ways. While O'Dea, Challinor \\& Johnson [6] considered only the effect of systematics on the tensor-to-scalar ratio $r$, we consider a family of parameters associated with the B-mode sector: $r$, $M_{\\nu}$, $\\alpha$, $w$ and $\\Omega_{k}$. We set all other cosmological parameters to be consistent with the WMAP values. In order to exhaust the potential of the CMB to constrain these parameters we carried out lensing extraction. In addition, we repeated the analysis for POLARBEAR, CMBPOL-B and QUIET+CLOVER with Monte Carlo simulations and found that the Fisher-Matrix approximation is, in general, inadequate for appraising the biases. We also found that high resolution experiments, such as POLARBEAR are very sensitive to bias from second order gradient effects (i.e. differential ellipticity and differential beamwidth) which is underestimated by the Fisher-matrix-based calculation, but fully accounted for with MCMC simulations. Also, unlike O'Dea, Challinor \\& Johnson [6] our results are presented independently of the scanning strategy details. The only assumption we made was that the scanning strategy is spatially uniform, a condition which can be achieved with or without a HWP which samples the polarization angles in a way which is uniform; both spatially, and in terms of polarization angle. In case that this approximation fails the more general formalism (Shimon et al. [7]) should be used with the added complexity introduced to lensing reconstruction by the scanning-induced non-gaussianity of the systematic B-mode.  We find that parameter bias is the dominant factor and its level actually sets the upper bounds on the beam parameters appearing in Tables IV through IX. Our results show that the most severe constraints are set on the most sensitive experiments for a given tolerance on $\\delta$ and $\\beta$ since these quantities are experiment-dependent (Eq. 26) and since, in general, an experiment with higher resolution and better sensitivity will result in smaller errors $\\sigma_{\\lambda_{i}}$. We expect that the constraints on the systematics should be more demanding so as to realize the potential of experiments. As mentioned above, the most stringent constraints are obtained from the requirement on the bias rather than from increased parameter uncertainty. Again, for the same reason, as shown for specific examples in Fig. 4; the bias {\\it always} exceeds the uncertainty for large enough systematics and this always takes place before the $10\\%$ thresholds in Eqs.(26) are attained. The reason is that for large enough systematics the induced spurious polarization becomes comparable to, or exceeds, the underlying polarization signals, therefore biasing the deduced value. It is easy to visualize configurations in which the bias increases without bound while the `curvature' of the likelihood function (i.e. the statistical error) with respect to specific cosmological parameters does not change. It is also clear from the tables that, in general, the tensor-to-scalar ratio is the most sensitive parameter, and the second most sensitive is $\\alpha$, the running of the scalar index (although there are some exceptions). If the tensor-to-scalar ratio is larger than the case we studied ($r=0.01$), this conclusion may change since $r$ is mainly affected by the overwhelming B-mode systematics on degree scales. $\\alpha$ is predicted to vanish by the simplest models of inflation and was added to parameter space to better fit the WMAP and other cosmological data. As is well-known, information from Ly-$\\alpha$ systems and other LSS probes can, in principle, better constrain $\\alpha$ if their associated systematics can be controlled to a sufficiently accurate level. For these small scales the CMB is not the ideal tool to extract information and the error that beam systematics induce on $\\alpha$ are not significant.  The upper limits we obtained in this work on the allowed range of beam mismatch parameters for given experiments and given arbitrarily-set tolerance levels on the parameter bias and uncertainty, constitute very conservative limits. It can certainly be the case that some of the systematics studied here may be fully or partially removed. This includes, in particular, the first order pointing error which couples to the dipole moment of non-ideal scanning strategies (see Shimon et al. [7]). By removing this dipole during data analysis the effect due to the systematic first order pointing error (dipole) drops dramatically. We made no attempt to remove or minimize these effects in this work. Our results highlight the need for scan mitigation techniques because the coupling of several beam systematics to non-ideal scanning strategies results in systematic errors. This potential solution reduces systematics, which ultimately propagate to parameter estimation, and affect mainly the parameters considered in this work. A brute-force strategy to idealize the data could be to remove data points that contribute to higher-than-the-monopole moments in the scanning strategy. This would effectively make the scanning strategy `ideal' and alleviate the effect of the {\\it a priori} most pernicious beam systematics. This procedure `costs' only a minor increase in the instrumental noise (due to throwing out a fraction of the data) but will greatly reduce the most pernicious reducible beam systematic, i.e. the first order pointing error (`dipole' effect). The lesson is clear: the rich treasures of cosmological parameters deducible from B-mode data  require a combination of high polarimetric fidelity and judicious data mining. Both are eminently feasible upcoming CMB polarization experiments."
    },
    "0806/0806.3269_arXiv.txt": {
        "abstract": "We present stellar population modeling results for 10 newly discovered Lyman alpha emitting galaxies (LAEs), as well as four previously known LAEs at z $\\sim$ 4.5 in the Chandra Deep Field -- South.  We fit stellar population models to these objects in order to learn specifically if there exists more than one class of LAE.  Past observational and theoretical evidence has shown that while many LAEs appear to be young, they may be much older, with \\lya~EWs enhanced due to resonant scattering of \\lya~photons in a clumpy interstellar medium (ISM).  Our results show a large range of stellar population age (3 -- 500 Myr), stellar mass (1.6 $\\times$ 10$^{8}$ -- 5.0 $\\times$ 10$^{10}$ $M$\\sol) and dust extinction (\\dust~= 0.3 -- 4.5 mag), broadly consistent with previous studies.  With such a large number of individually analyzed objects, we have looked at the distribution of stellar population ages in LAEs for the first time, and we find a very interesting bimodality, in that our objects are either very young ($<$ 15 Myr) or old ($>$ 450 Myr).  This bimodality may be caused by dust, and it could explain the \\lya~duty cycle which has been proposed in the literature.  We find that eight of the young objects are best fit with a clumpy ISM.  We find that dust geometry appears to play a large role in shaping the SEDs that we observe, and that it may be a major factor in the observed \\lya~equivalent width distribution in high redshift \\lya~galaxies, although other factors (i.e. outflows) may be in play.  We conclude that 12 out of our 14 LAEs are dusty star-forming galaxies, with the other two LAEs being evolved galaxies. ",
        "introduction": "While high-redshift galaxies can be hard to observe due to dimming with distance, narrowband selection of \\lya~galaxies has proven a very efficient method to select high-redshift galaxies based on a strong emission line (e.g., Rhoads  et al  2000, 2004;  Rhoads \\&  Malhotra 2001; Malhotra \\& Rhoads 2002; Cowie \\&  Hu 1998; Hu et al 1998, 2002, 2004; Kudritzki et al 2000; Fynbo, Moller, \\& Thomsen 2001; Pentericci et al 2000; Ouchi et al 2001, 2003, 2004; Fujita et al 2003; Shimasaku et al 2003, 2006;  Kodaira et  al 2003;  Ajiki et al  2004; Taniguchi  et al 2005; Venemans et al 2002, 2004; Gawiser et al. 2006; Lai et al. 2007, 2008; Nilsson et al. 2007; Finkelstein et al. 2008a).  These objects are of interest, as over 40 years ago Partridge \\& Peebles (1967) proposed that they may be signs of primitive galaxies in formation.  This was easy to understand, as \\lya~photons are copiously produced in star formation regions, and we would expect the first galaxies to be undergoing periods of extreme star formation.  However, not until recently have we had the data to verify this assumption.  The availability of stellar population modeling codes has allowed the derivation of physical parameters of galaxies from photometry alone.  In the past few years (thanks to large surveys such as the Great Observatories Origins Deep Survey (GOODS)), broadband photometry of LAEs has become deep enough to compare objects to these models, learning about such parameters as stellar population age, stellar mass, and dust extinction.  First results from these studies were unsurprising, as stacking analyses showed that an average LAE was young ($\\sim$ 10 -- 100 Myr), low mass (10$^{7 - 8}$ $M$\\sol) and dust free (e.g., Gawiser et al. 2006; Finkelstein et al. 2007; Lai et al. 2008).  Recently, deeper data has allowed the comparison of individual high-redshift LAEs to models for the first time, as chronicled in Chary et al. (2005), Pirzkal et al. (2007), Lai et al. (2007) and Finkelstein et al. (2008a).  These studies have shown a wide range of results, with LAE ages from 1 -- 1000 Myr, masses from 10$^{6 - 10}$ $M$\\sol~and, most surprisingly, dust extinction with A$_{V}$ up to 1.3 mag.  These results show that while some LAEs may be young and dust-free, many are dusty, and some are even evolved (i.e. old and high-mass).  This raises an interesting question, as how can an evolved stellar population produce enough \\lya~photons to be picked up by a narrowband selected survey?  Many scenarios have been proposed (i.e. zero metallicity, top-heavy initial mass function etc.), but these are rather extreme, and given the amount of dust extinction we see, they would be unlikely to produce enough \\lya~photons to explain the observed excesses.  We have thus decided to observationally test a scenario developed theoretically by Neufeld (1991) and Hansen \\& Oh (2006), where the \\lya~equivalent width (EW) actually gets {\\it enhanced} due to a dusty interstellar medium (ISM).  This is counter-intuitive, as dust will strongly suppress any ultraviolet (UV) photons.  The news is even worse for \\lya~photons, as they resonantly scatter off of neutral hydrogen atoms, thus their mean-free-paths are rather small, vastly increasing their chances of encountering a dust grain.  However, as Neufeld, Hansen \\& Oh suggest, an ISM that is very clumpy could actually prevent the \\lya~photons from seeing the dust at all.  This can happen if the dust and neutral hydrogen are evenly mixed together in clumps, with a tenuous ionized medium separating the clumps.  In this geometry, the \\lya~photons stand a very high chance of being resonantly scattered at the surface of these clumps, spending most of their time in the inter-clump medium.  In this manner, the \\lya~photons are effectively screened from ever encountering a dust grain, as the gas and dust are in the same geometry, so the \\lya~gets shielded from encountering dust.  The story is different for continuum photons, as they are not resonantly scattered, thus they will penetrate deeply into a clump, with a strong chance of being scattered or absorbed.  As EW is a measure of line-to-continuum flux, this effectively enhances the observed EW over that due to the stars (or more accurately, the stars interaction with the interstellar hydrogen in their immediate surroundings).  Note that the \\lya~flux itself is not being increased, rather the \\lya-to-continuum ratio is being enhanced (see Finkelstein et al. 2008a, \\S 3.2 for a detailed explanation of this scenario).  In Finkelstein et al. (2008a; hereafter F08a), we analyzed a sample of four LAEs, fitting model spectra to their SEDs, searching for proof that this type of ISM exists.  Three of our objects were best fit by young (5 Myr) dusty (\\dust~$\\sim$ 1 -- 2 mag) stellar populations, similar to those seen by Pirzkal et al. (2007) at z $\\sim$ 5 and Chary et al. (2005) at z $\\sim$ 6.5.  However, the fourth object was best-fit by an old (800 Myr) stellar population, with 0.4 mag of dust arrayed in a clumpy ISM, showing evidence of dust enhancement of the \\lya~EW.  Also interesting was the age distribution, as these galaxies were either very young or very old, although this distribution was hard to quantify with such a small sample.  Using our larger sample, we will now see if the absence of ``teenage'' LAEs is real, and see if most LAEs are primitive, dusty or evolved.  Further detection of this dust-enhancement of the \\lya~EW can help explain the larger than expected EWs seen in many LAEs (e.g., Kudritzki et al. 2000; Malhotra \\& Rhoads 2002; Finkelstein et al. 2007).  This paper is organized as follows.  In \\S 2 we present our observations, including our object selection and redshift information where applicable.  In \\S 3 we present our stellar population models.  Our best-fit results for each object are presented in \\S 4, and we discuss the implications of these results in \\S 5, including suggestions for future improvement.  We present our conclusions in \\S 6.  In this paper we assume Benchmark Model cosmology, where $\\Omega_{m}$ = 0.3, $\\Omega_{\\Lambda}$ = 0.7 and H$_{0}$ = 0.7 (c.f. Spergel et al. 2007).  All magnitudes in this paper are listed in AB magnitudes (Oke \\& Gunn 1983). ",
        "conclusions": "We have presented the results from our analysis of 10 newly discovered, and four previously known, narrowband selected \\lya~galaxies in the CDF--S.  We compared the SEDs of these objects to stellar population models in order to determine their physical properties such as age, mass, and dust extinction.  More specifically, we are interested in finding out whether enhancement of the \\lya~EW due to a clumpy, dusty ISM can be responsible for some of the large \\lya~EWs which we have observed.  We first computed the best-fit stellar population model to each object allowing one SFH.  For the four objects which we previously analyzed in F08a, we confirmed our earlier results, which had three of the objects being fairly young, and one being very old.  The old object (CHa-3) looks to still have a strong \\lya~line due to dust enhancement, shown by its value of the clumpiness parameter (q) of zero.  Although none of the rest of our sample was best-fit by an age over 50 Myr, the majority of our objects appeared to require some amount of clumpy dust enhancement of the \\lya~EW, with 9/14 objects having q $<$ 1, thus dust enhancement is widespread.  As a test, we also allowed objects to be fit by two bursts of star formation: One at a maximally old age of 1.4 Gyr, and one at any time.  Most of our objects showed significantly worse fits with two bursts.  However, we do feel as if it was a necessary exercise to fit all objects to this type of model, as a two-burst population could be a viable alternative explanation to those objects which appear old due to dust enhancement.  To assess the validity of our results, we ran 7000 Monte Carlo simulations, obtaining a best-fit for each object from each simulation.  In the resulting contour plots, we can examine whether the best-fit model truly represents the most likely model by seeing if the best-fit lies in the largest 68\\% confidence region.  If it does not, we assign the object a new age, \\dust~and q based on the largest 68\\% confidence region.  \\fig{fithist2}a and \\fig{fithist2}b plot histograms of the most likely ages and q's for our sample.  The distribution of ages is interesting, in that it implies that LAEs are either very young ( $<$ 15 Myr), or very old ($>$ 450 Myr).  However, as Figure 9 shows, we need a much larger sample before we can see if the bimodality is significant.  There could be many explanations for this, but we propose that this bimodality in LAE stellar population ages may be due to dust.  At the beginning of this work, we asked the question of whether LAEs were primitive, dusty or evolved galaxies.  This work, among others, has shown that while many LAEs are young, they are not primitive.  Using our most likely results, we find a range of \\dust~of 0.3 to 4.5 mag, thus all of our objects have some amount of dust extinction in them.  As dust comes primarily from evolved stars and stellar deaths, the existence of dust is strong evidence that LAEs are not primitive, as the dust has been produced by a previous generation of stars.  We find that out of our 14 candidate LAEs, 12 of them appear to be dusty star-forming galaxies, with ages from 3 - 15 Myr, and \\dust~from 0.4 to 4.5 mag.  The remaining two objects appear to be evolved galaxies, with ages of 450 and 500 Myr, and \\dust~of 0.3 and 3.5 mag respectively, exhibiting \\lya~emission due to dust enhancement of the EW.  The young galaxies in our sample, although they are dusty due to a previous generation of stars, still manage to emit \\lya, mainly due to dust enhancement (8/12 have q $<$ 1).  After a few 10's of Myr, enough massive stars have exploded to further saturate the ISM with dust, and this could explain the drop-off in numbers (\\fig{fithist2}a), and why we don't see any LAEs from 15 - 450 Myr.  After some period of time, the stars have changed the ISM geometry enough so that \\lya~can escape again.  This ISM is now very patchy, which is why both of our old LAEs have dust enhancement of the \\lya~EW in their model spectra.  While this scenario is intriguing, a much larger sample is needed before we can see if this age bimodality is statistically significant.  While we inferred many properties about LAEs in this work, we have also learned that more data is needed before we can truly match our objects to the models.  For most of our objects, we are missing data in a crucial area of the SED, constraining the 4000 \\AA~break.  Without these data points, objects are allowed to be either old or dusty to explain the red colors, a degeneracy which can be fixed with better NIR data.  Future observatories such as the {\\it James Webb Space Telescope} will be sensitive in this regime, and will provide the data needed to better constrain these objects.  Nonetheless, with the data in hand, we can now say that dust enhancement of the \\lya~EW appears to be occurring in the majority of our LAEs, and this effect should be considered in future stellar population studies."
    },
    "0806/0806.2868_arXiv.txt": {
        "abstract": "The mean absolute brightness temperature of the diffuse radio background was measured as a function of frequency in a continuous band between 100 and 200 MHz over an effective solid angle of $\\sim\\pi$~str at high Galactic latitude.  A spectral brightness temperature index of $\\beta=2.5\\pm{0.1}$ ($\\alpha_S=0.5$) was derived from the observations, where the error limits are $3\\sigma$ and include estimates of the instrumental systematics.  Zenith drift scans with central declination $\\delta=-26.5^{\\circ}$ and spanning right ascensions $0 < \\alpha < 10$~h yielded little variation in the mean spectral index. The mean absolute brightness temperature at $\\nu=150$~MHz was found to reach a minimum of $T=237\\pm{10}$~K at $\\alpha=2.5$~h.  Combining these measurements with those of \\citet{1982A&AS...47....1H} yields a spectral index of $\\beta=2.52\\pm{0.04}$ between $150<\\nu<408$~MHz. ",
        "introduction": "The last few years have seen renewed interest in the low-frequency radio sky.  This is due, in part, to the development of new radio arrays designed to study the cosmological epoch of reionization (EOR) through redshifted 21~cm emission from neutral hydrogen in the intergalactic medium (IGM) between $6<z<15$.  At these target redshifts, the 21~cm line is shifted well into meter wavelengths and its characteristic rest frequency of $\\nu=1420$~MHz is reduced to $202>\\nu>89$~MHz.  The detection of the reionization signal is anticipated to be extremely challenging since the low-frequency radio sky is dominated by bright synchrotron radiation and other diffuse emission from the Galaxy, as well as by the integrated contribution of extragalactic continuum sources.  Numerous efforts are underway to investigate strategies for mitigating these foreground contaminants in the planned reionization experiments.  Until new observations are begun, however, these studies must rely either on theoretically motivated arguments for the expected foreground contributions or they must extrapolate measurements from other frequencies, partial sky maps, or limited source catalogs.  Even simple all-sky maps of the Galactic synchrotron contribution must be generated with similar techniques \\citep{2008arXiv0802.1525D}.  This situation produces significant uncertainties in the results of redshifted 21 cm foreground subtraction modeling.  One of the most basic measurements needed for extrapolating foreground properties from other frequencies is the spectral index of the diffuse Galactic and extragalactic emission.  In this regard, upcoming redshifted 21~cm experiments overlap with efforts to study and model foregrounds in the latest generations of cosmic microwave background (CMB) experiments. As those measurements become increasingly sensitivity and target more difficult signatures in the CMB, such as polarization anisotropies, the need to quantify and subtract the contributions of diffuse synchrotron and free-free emission from the Galaxy also increases.  The spectral index, $\\beta$, of the observed sky brightness temperature is typically defined in temperature units as $T\\sim\\nu^{-\\beta}$.  This is related to flux units according to $S\\sim\\nu^{-\\alpha_S}$, where $\\alpha_S=\\beta-2$. Much of our fundamental knowledge of the spectral index properties of the low-frequency radio sky originates from the 1960s and 1970s, when there was considerable interest in constraining spatial variations in $\\beta$ over the frequency range $10<\\nu<1400$~MHz in order to investigate the physical structure of the Galaxy \\citep{1962MNRAS.124..297T, 1966MNRAS.132...79A, 1966MNRAS.133..463P, 1967MNRAS.136..219B, 1974MNRAS.166..355W, 1974MNRAS.166..345S, 1976Ap&SS..44..159S}.  At the time, it was recognized that, along with extragalactic sources, the diffuse Galactic radio emission had structure that was composed of three components, originating from the disk, the spiral arms, and a radio halo, respectively \\citep{1976Ap&SS..44..159S}. Several experiments were performed using sets of identically scaled horns, antennas, or dipole arrays, in addition to receivers on satellites, to measure accurately the spectrum of the Galactic non-thermal radiation.  The individual components were identified and separated by utilizing analysis tools such as temperature-temperature (``T-T'') plots so that their spectral properties could be investigated, both as functions of frequency and Galactic coordinate. These early measurements resulted in several findings, including that the total diffuse spectrum flattens below 10 MHz \\citep{1976Ap&SS..44..159S}, the halo contribution is extremely faint, the spectral index of the disk contribution is dependent on location, and the typical spectral index of the disk contribution steepens rapidly with increasing frequency between 200 and 400~MHz from about $\\beta=2.4$ to $\\beta=2.8$ \\citep{1967MNRAS.136..219B}. More recent investigations \\citep{1987MNRAS.225..307L, 1988A&AS...74....7R, 1988A&A...196..211R, 1999A&AS..137....7R, 1998ApJ...505..473P, 2001MNRAS.327..545J, 2003A&A...410..847P} have confirmed many of these early results and demonstrated that the total spectral index saturates to about $\\beta=2.9$ above 1~GHz. Recently, there has also been renewed interest in absolute sky temperature measurements around 1~GHz in order to investigate predicted deviations in the CMB from a purely black-body (Planckian) frequency distribution.  \\begin{figure*} \\center \\includegraphics[trim=0pc 0pc 0pc 5pc clip]{f1.eps} \\caption[Fourpoint antenna]{ \\label{f_eorantfig} Schematic of the EDGES broadband ``fourpoint'' antenna, which is based on the design of \\citep{fourpoint2}.} \\end{figure*}  Unlike the early experiments, the modern approach to low-frequency radio instrumentation uses broadband antennas and receivers, along with analog-to-digital sampling systems that are capable of directly sampling radio-frequency waveforms without mixing to intermediate frequencies. We have developed such a system \\citep{2008ApJ...676....1B} to cover 100 to 200 MHz for an ``Experiment to Detect the Global EOR Signature'' (EDGES).  This experiment is motivated by theoretical models that predict a weak ($\\lesssim35$~mK), but potentially observable, spectral signature similar to a step function due to the redshifted 21~cm contribution to the mean sky spectrum disappearing as reionization unfolds.  Although the design of the EDGES system is optimized for constraining the smoothness of the low-frequency radio spectrum, a small change to the configuration enables the additional calibration needed to constrain the absolute sky temperature and the spectral index of the diffuse emission.  In this paper, we describe the method of absolute calibration and report the results of the absolute temperature and spectral index measurements.  Because the EDGES antenna is a single dipole with a large field of view, it is incapable of performing the difference measurements employed in the pioneering efforts to isolate the Galactic and extragalactic contributions to the spectrum and, therefore, constrains only the total spectrum.  We begin in $\\S2$ by describing the experimental approach of the EDGES system.  In $\\S3$, we present the procedure for calibrating the sky temperature measurements, along with the values derived from an observing campaign.  We conclude in $\\S4$ and $\\S5$ with a calculation of the spectral index from the measurements and a summary of previous measurements. ",
        "conclusions": "The measurement of the spectral index of the background with a broadband system requires a number of corrections, but we have shown that it is possible to achieve an accurate result competitive with other radio astronomy techniques. As more broadband systems are built, and the systematics are better understood, there is potential for improvements in the accuracy through better modeling and better antenna design for a wider bandwidth of low reflection coefficient.  Based on the trial measurements with the EDGES system presented in this paper, we are confident that the average spectral index at high Galactic latitudes from 100 to 200~MHz lies between $2.4 < \\beta < 2.6$. It should be noted that this is significantly below the the nominal value of $\\beta=2.7$ often assumed in this range of frequencies, and a little less than the high Galactic latitude value of $\\beta\\approx2.6$ that can be extracted from the analysis of \\citet{2008arXiv0802.1525D}. It is consistent with many of the earlier findings summarized in Table~\\ref{t_edges_results}.  Finally, since much of this analysis was motivated by the needs of the redshifted 21 cm experiments, we would also like to bring attention to the utility of including even one simple, well understood antenna in the large, active dipole arrays being developed.  In addition to being used for absolute measurements of the background, one well calibrated antenna can be used for the calibration of aperture arrays with complex beampatterns using the method described by \\citet{1958AuJPh..11...70L}. Such a configuration is sufficient to calibrate the other elements of the array through the redundancy in the correlations on baselines to the single calibrated element and baselines between uncalibrated elements, and can accomplished using only unresolved sources whose flux density need not be known."
    },
    "0806/0806.2435_arXiv.txt": {
        "abstract": "{We explore the relation between black hole mass ($M_{\\rm BH}$) and the motion of the jet components for a sample of blazars. The Very Long Baseline Array (VLBA) 2cm Survey and its continuation: Monitoring of Jets in active galactic nuclei (AGNs) with VLBA Experiments (MOJAVE) have observed 278 radio-loud AGNs, of which 146 blazars have reliable measurements on their apparent velocities of jet components.  We calculate the minimal Lorentz factors for these sources from their measured apparent velocities, and their black hole masses are estimated with their broad-line widths.  A significant intrinsic correlation is found between black hole masses and the minimal Lorentz factors of the jet components, which the Eddington ratio is only weakly correlated with the minimal Lorentz factor, which may imply that the Blandford-Znajek (BZ) mechanism may dominate over the Blandford-Payne (BP) mechanism for the jet acceleration (at least) in blazars. ",
        "introduction": "% \\label{sect:intro} Relativistic jets have been observed in many radio-loud AGNs, which are believed to be formed very close to the black holes. The currently most favored models of the jet formation are BZ and BP mechanisms (Blandford \\& Znajek 1977; Blandford \\& Payne 1982). In these mechanisms, the power of jet is extracted from the disk or black hole rotational energy. The disk-jet connection has been investigated by many authors in different ways (Rawlings \\& Saunders 1991; Falcke \\& Biermann 1995;  Cao \\& Jiang 1999; 2001; 2002, Xie et al. 2007; Xie et al. 2008).   Some different approaches were proposed to estimate the masses of the black holes in AGNs, such as the gas kinematics near a black hole (see Ho \\& Kormendy 2000 for a review and references therein). The central black hole mass derived from the direct measurements on the gases moving near the hole is reliable, but unfortunately, it is available only for very few AGNs. For most AGNs, the velocities of the clouds in broad line regions (BLR) can be inferred from the widths of their broad emission lines. If the radius of the BLR is available, the mass of the central black hole can be derived from the broad-line width on the assumption that the clouds in the BLR are gravitationally bound and orbiting with Keplerian velocities (Dibai 1980). The radius of the BLR can be measured by using the reverberation-mapping method from the time delay between the continuum and line variations (Peterson 1993; Netzer \\& Peterson 1997). Long-term monitoring on the source is necessary for applying this method to derive the radius of the BLR, which leads to a small amount of AGNs with measured black hole masses in this way. Alternatively, a tight correlation was found between the size of the BLR and the optical continuum luminosity, which can be used to estimate the size of the BLR in an AGN from its optical luminosity and then the black hole mass (e.g., Wandel, Peterson \\& Malkan 1999; Kaspi et al. 1996; 2000; Laor 2000).  The kinematic properties of the jet components in blazars were revealed by multi-epoch VLBI observations (e.g., Kellermann et al. 2004; Lister et al. 2005). In this paper, we use a large sample of blazars, of which the proper motions were well measured with VLBA, to explore the relations between the jet speeds and physical properties of blazars, i.e., the black hole masses and Eddington ratios.  The cosmological parameters $\\Omega_{\\rm M}=0.3$, $\\Omega_{\\Lambda}=0.7$, and $H_0=70~ {\\rm km~s^{-1}~Mpc^{-1}}$ have been adopted in this work. ",
        "conclusions": "We find an intrinsic correlation between black hole masses and the minimal Lorentz factors of jet components for a sample of blazars, while no significant correlation between the Eddington rations and the Lorentz factors is present for the same sample. Our main statistical results will not be altered, even if those black holes with masses estimated with continuum luminosities are removed. Our statistical results provide useful clues to the mechanisms of jet formation and acceleration in blazars.  It is believed that the growth of massive black holes in the centers of galaxies is dominantly governed by mass accretion in AGN phases (e.g., Soltan 1982; Yu \\& Tremaine 2002). The massive black holes will be spun up through accretion, as the black holes acquire mass and angular momentum simultaneously though accretion. The spins of massive black holes may also be affected by the mergers of black holes. A rapidly rotating new black hole will be present after the merger of two black holes, only if the binary's larger member already spins quickly and the merger with the smaller hole is consistently near prograde, or if the binary's mass ratio approaches unity (Hughes \\& Blandford 2003). The comoving space density for heavier black holes is much lower than that for lighter black holes (e.g., see the black hole mass function in Yu \\& Tremaine 2002), which means that the probability of the mergers of two black holes with similar masses is lower for heavier black holes. This implies the spins of heavier black holes are mainly regulated by accretion rather than the mergers. Thus, it is natural to expect (in statistical sense) that the heavier black holes have higher spin parameters $a$ than their lower mass counterparts. Volonteri et al. (2007) studied on how the accretion from a warped disc influences the evolution of black hole spins and concluded that within the cosmological framework, one indeed expects most supermassive black holes in elliptical galaxies to have on average higher spin than black holes in spiral galaxies, where random, small accretion episodes (e.g., tidally disrupted stars, accretion of molecular clouds) might have played a more important role. The jets can be accelerated to higher speeds by the heavier black holes, because they are spinning more rapidly (Blandford \\& Znajek 1977). The intrinsic correlation between black hole masses and the minimal Lorentz factors of jet components found in this work is consistent with the Blandford-Znajek mechanism. The properties of accretion disks accretion disk are related with the dimensionless accretion rates $\\dot{m}$ ($\\dot{m}=\\dot{M}/\\dot{M}_{\\rm Edd}\\propto L_{\\rm bol}/L_{\\rm Edd}$). No significant correlation between $L_{\\rm bol}/L_{\\rm Edd}$ and $\\gamma_{\\rm min}$ is found, which implies that the jet acceleration may not be related with the properties of the accretion disk, which may imply that the jet formation is not sensitive to the disk structure. This is, of course, quite puzzling, and to be verified by the future work with a larger blazar sample. Our statistical results implies that the BZ mechanism may dominate over BP mechanism for the jet acceleration in blazars."
    },
    "0806/0806.1020_arXiv.txt": {
        "abstract": "We present a simple, largely empirical but physically motivated model to interpret the mid- and far-infrared spectral energy distributions of galaxies consistently with the emission at ultraviolet, optical and near-infrared wavelengths. Our model relies on an existing angle-averaged prescription to compute the absorption of starlight by dust in stellar birth clouds and in the ambient interstellar medium (ISM) in galaxies. We compute the spectral energy distribution of the power reradiated by dust in stellar birth clouds as the sum of three components: a component of polycyclic aromatic hydrocarbons (PAHs); a mid-infrared continuum characterising the emission from hot grains at temperatures in the range 130--250~K; and a component of grains in thermal equilibrium with adjustable temperature in the range 30--60~K. In the ambient ISM, we fix for simplicity the relative proportions of these three components to reproduce the spectral shape of diffuse cirrus emission in the Milky Way, and we include a component of cold grains in thermal equilibrium with adjustable temperature in the range 15--25~K. Our model is both simple and versatile enough that it can be used to derive statistical constraints on the star formation histories and dust contents of large samples of galaxies using a wide range of ultraviolet, optical and infrared observations. We illustrate this by deriving median-likelihood estimates of the star formation rates, stellar masses, effective dust optical depths, dust masses, and relative strengths of different dust components of 66 well-studied nearby star-forming galaxies from the Spitzer Infrared Nearby Galaxy Survey (SINGS). We explore how the constraints derived in this way depend on the available spectral information. From our analysis of the SINGS sample, we conclude that the mid- and far-infrared colours of galaxies correlate strongly with the specific star formation rate, as well as with other galaxy-wide quantities connected to this parameter, such as the ratio of infrared luminosity between stellar birth clouds and the ambient ISM, the contributions by PAHs and grains in thermal equilibrium to the total infrared emission, and the ratio of dust mass to stellar mass. Our model can be straightforwardly applied to interpret ultraviolet, optical and infrared spectral energy distributions from any galaxy sample. ",
        "introduction": "\\label{intro}  The spectral energy distributions of galaxies contain valuable information about their contents in stars, gas and dust. Direct ultraviolet, optical and near-infrared radiation from stars provides clues on the past star formation history, chemical enrichment and attenuation by dust. Nebular emission lines produced by the gas heated by young stars provide further clues on the current star formation activity and the physical conditions of the star-forming gas. At wavelengths $\\lambda\\ga 3~\\mu$m, the mid- and far-infrared emission reflects the heating of dust in different components of the interstellar medium (ISM) by stars of all ages. Observations at ultraviolet, optical and infrared wavelengths are now becoming available for large samples of galaxies. These include data collected in the ultraviolet by the {\\it Galaxy Evolution Explorer} ({\\it GALEX}, \\citealt{MARTIN05}), in the optical by the Two-degree Field Galaxy Redshift Survey \\citep{2dF} and the  Sloan Digital Sky Survey (SDSS, \\citealt{S02}), in the near-infrared by the Two Micron All Sky Survey (2MASS, \\citealt{2MASS}), in the mid- and far-infrared by the {\\it Infrared Astronomical Satellite} ({\\it IRAS}, \\citealt{IRAS}), the {\\it Infrared Space Observatory} ({\\it ISO}, \\citealt{ISO}) and the {\\it Spitzer Space Telescope} \\citep{SIRTF}, and in the sub-millimetre by the Sub-millimeter Common User Bolometer Array (SCUBA) on the James Clerk Maxwell Telescope \\citep{SCUBA}. Extracting constraints on the stellar populations and ISM of galaxies from these multi-wavelength observations requires the consistent modelling of the emission by stars, gas and dust.  A standard approach to model consistently the emission from stars and dust in galaxies has been to solve the radiative transfer equation for idealised (bulge + disc) spatial distributions of stars and dust (e.g. \\citealt{RR80,ERR90,GORDON01,MIS01,POPESCU00,MISIRIOTIS01}). Early models of this type did not include the evolution of stellar populations. \\citet{GRASIL98} were the first to couple radiative transfer through a dusty ISM and the spectral (and even chemical) evolution of stellar populations. Their model also accounts for the fact that stars are born in dense molecular clouds, which dissipate after some time, and hence, that newly born stars are more attenuated than older stars (see also, e.g., \\citealt{CF00,TUFFS04}). This type of sophisticated model is useful to interpret in detail the emission from individual galaxies in terms of constraints on stellar populations and the spatial distribution and physical properties of the dust. However, because of the complexity of radiative transfer computations, it is not optimised to derive statistical constraints from observations of large samples of galaxies.  A more recent model of starburst galaxies by \\citet[][see also \\citealt{GROVES07}]{DOPITA05} incorporates the consistent treatment of the spectral evolution of stellar populations, the dynamic expansion of \\hii\\ regions and radiative transfer of starlight through gas and dust. The authors of this model provide a simple parameterization of the ultraviolet, optical and infrared spectra of starburst galaxies by adding the spectra of different types of compact \\hii\\ regions and their surrounding photo-dissociation regions. This model provides a fast and flexible tool to interpret starburst galaxy spectra in terms of the physical parameters of star-forming regions. However, it is not designed to be applicable to more quiescent galaxies, in which older stellar populations dominate the emission.  In parallel to these theoretical studies, observations by {\\it IRAS} and {\\it ISO} have motivated the development of simple, empirically calibrated spectral libraries to interpret the infrared emission from galaxies at wavelengths between 3 and 1000~\\mic. For example, \\cite{CE01} and \\cite{DH02} both present single-parameter families of infrared spectra to relate an observed spectral energy distribution to either the total infrared luminosity of a galaxy or the intensity of the interstellar radiation field heating the dust. These libraries can be applied easily to the interpretation of large galaxy samples. They have proved useful to characterise the infrared emission from observed galaxies and to investigate the origin of the cosmic infrared background (e.g. \\citealt{FRANC01, CE01,DE02,LAGACHE03,LAGACHE04, DALE05,MARC06}). A disadvantage of this approach is that it does not relate consistently the infrared emission of the dust to the emission from stellar populations. Another potential limitation is that most existing spectral libraries were calibrated using local galaxy samples, and hence, they may not be applicable to studies of the infrared emission from galaxies at all redshifts (e.g. \\citealt{POPE06, ZHENG07}).  In this paper, we present a simple, largely empirical but physically motivated model to interpret the mid- and far-infrared spectral energy distributions of galaxies consistently with the emission at ultraviolet, optical and near-infrared wavelengths. We compute the spectral evolution of stellar populations using the \\cite{BC03} population synthesis code. To describe the attenuation of starlight by dust, we appeal to the two-component model of \\cite{CF00}. This has been shown to account for the observed relations between the ultraviolet and optical (line and continuum) spectra and the {\\em total} infrared luminosities of galaxies in wide ranges of star-formation activity and dust content \\citep{JB04,KONG04}. We use this model to compute the luminosity absorbed and re-emitted by dust in stellar birth clouds (i.e. giant molecular clouds) and in the ambient (i.e. diffuse) ISM in galaxies. We then distribute this luminosity in wavelength to compute infrared {\\em spectral energy distributions}. We describe the infrared emission from stellar birth clouds as the sum of three components: a component of polycyclic aromatic hydrocarbons (PAHs); a mid-infrared continuum characterising the emission from hot grains at temperatures in the range 130--250~K; and a component of grains in thermal equilibrium with adjustable temperature in the range 30--60~K. In the ambient ISM, we fix for simplicity the relative proportions of these three components to reproduce the spectral shape of diffuse cirrus emission in the Milky Way, and we include a component of cold grains in thermal equilibrium with adjustable temperature in the range 15--25~K.  This simple but versatile model allows us to derive statistical estimates of physical parameters such as star formation rate, stellar mass, dust content and dust properties, from combined ultraviolet, optical and infrared galaxy spectra. To achieve this, we adopt a Bayesian approach similar to that successfully employed by, e.g., \\citet{KAUF03}, \\citet{JB04}, \\citet{AG05} and \\citet{SAL07} to interpret ultraviolet, optical and near-infrared galaxy spectra using only the \\cite{BC03} and \\cite{CF00} models. As an example, we derive median-likelihood estimates of a set of physical parameters describing the stellar and dust contents of 66 star-forming galaxies from the Spitzer Infrared Nearby Galaxy Survey (SINGS, \\citealt{K03}). Our model reproduces well the observed spectral energy distributions of these galaxies across the entire wavelength range from the far-ultraviolet to the far-infrared, and the star formation histories and dust contents of the galaxies are well constrained. We explore how the constraints derived in this way depend on the available spectral information. From our analysis of the SINGS sample, we conclude that the mid- and far-infrared colours of galaxies are tightly related to the specific star formation rate and to other galaxy-wide properties connected to this parameter.  We present our model of the combined ultraviolet, optical and infrared spectral energy distributions of galaxies in Section~\\ref{the_model}. In Section~\\ref{extraction}, we first describe our approach to derive statistical constraints on galaxy physical parameters from multi-wavelength observations. Then, we use this approach to interpret observations of the SINGS galaxy sample taken with {\\it GALEX}, 2MASS, {\\it Spitzer}, {\\it ISO}, {\\it IRAS} and SCUBA. We compare our results to those that would be obtained using previous prescriptions of the infrared emission from galaxies. We also discuss potential sources of systematic errors. Section~\\ref{conclusion} summarises our conclusions. ",
        "conclusions": "\\label{conclusion}  We have developed a simple but versatile model to interpret the mid- and far-infrared spectral energy distributions of galaxies consistently with the emission at ultraviolet, optical and near-infrared wavelengths. Our model relies on the \\cite{BC03} population synthesis code to compute the spectral evolution of stellar populations, and on the two-component model of \\cite{CF00} to compute the total infrared luminosity absorbed and reradiated by dust in stellar birth clouds and in the ambient ISM. We distribute this total infrared energy in wavelength over the range from 3 to 1000~\\mic\\ by considering the contributions by four main dust components: PAH emission, mid-infrared continuum emission from hot dust, warm dust with adjustable equilibrium temperature in the range 30--60~K and cold dust with adjustable equilibrium temperature in the range 15--25~K. We keep as adjustable parameters the relative contributions by PAHs, the hot mid-infrared continuum and warm dust to the infrared luminosity of stellar birth clouds. Cold dust resides (in an adjustabe amount) only in the ambient ISM, where the relative ratios of the other three components are fixed to the values reproducing the observed mid-infrared cirrus emission of the Milky Way. We find that this minimum number of components is required to account for the infrared spectral energy distributions of galaxies in wide ranges of star formation histories.  We have generated a comprehensive library of model galaxy spectra by combining a library of attenuated stellar population spectra (built from stochastic star formation histories and dust contents) with a library of infrared emission spectra. As Fig.~\\ref{seds2} illustrates, these models provide appropriate fits to the observed ultraviolet, optical and infrared spectral energy distributions of nearby galaxies in the SINGS sample, for which data are available from {\\it GALEX}, RC3, 2MASS, {\\it Spitzer}, {\\it ISO}, {\\it IRAS} and SCUBA \\citep{K03}.  We have used this model library to derive median-likelihood estimates of the star formation rate, stellar mass, dust attenuation, dust mass and relative contributions by different dust components to the total infrared luminosity of every SINGS galaxy in our sample. The accuracy of these estimates depends on the available spectral information. We find that, for example, although the total infrared luminosity \\ldust\\ of a galaxy can be roughly estimated using ultraviolet, optical and near-infrared data alone (at least for values in the range from a few $\\times10^{8}$ to a few $\\times 10^{10}L_\\odot$), reliable estimates of this parameter require infrared observations at longer wavelengths.  A main advantage provided by our model is the capacity to study the relation between different physical parameters of observed galaxies in a quantitative and statistically meaningful way. We find that, for example, the specific star formation rate of SINGS galaxies correlates strongly not only with observed infrared colours, but also with several other properties of the galaxies, such as the fraction of total infrared luminosity contributed by dust in the ambient ISM, the contributions by PAHs, warm dust and cold dust to the total infrared luminosity and the ratio of dust mass to stellar mass. These correlations provide important insight into the link between star formation and ISM properties in galaxies. In particular, they allow one to quantify the relations between star formation rate, gas mass, dust mass, stellar mass, dust temperature and distribution of the dust between giant molecular clouds and the diffuse interstellar medium. Studies of these relations at different redshifts will have important implications for our understanding of the physical processes that dominate galaxy evolution.  Our model should be useful for interpreting data from any modern galaxy survey at ultraviolet, optical and infrared wavelengths. It can also be used to design new observations by optimizing the set of observables required to constrain specific physical parameters of galaxies. This model is meant to be used by the astronomical community and we intend to make it publicly available."
    },
    "0806/0806.4119_arXiv.txt": {
        "abstract": "{ We explore the effect of cosmic radiative feedback from the sources of reionization on the thermal evolution of the intergalactic medium. We find that different prescriptions for this feedback predict quite different thermal and reionization histories. In spite of this, current data can not discriminate among different reionization scenarios. We find that future observations both from 21-cm and CMB experiments can be used to break the degeneracy among model parameters provided that we will be able to remove the foreground signal at the percent (or better) level. } ",
        "introduction": "It is well known that the temperature increase of the cosmic gas in ionized regions leads to a dramatic suppression of the formation of low-mass galaxies (see Ciardi \\& Ferrara 2005 for a review). We explore the impact of this effect during cosmic reionization by  considering two different feedback prescriptions: (i) {\\bf suppression model} where galaxies can form stars unimpeded provided that their circular velocity is larger than a critical threshold, which is not fixed to a constant value but evolves according to gas temperature \\cite{CF06}; (ii) {\\bf filtering model}, where, depending on the mass of the galaxy, the fraction of gas available to star formation is reduced with respect of the universal value and it is fully specified by the filtering mass at that redshift \\cite{G00}.  We implement these two different radiative feedback prescriptions into a physically-motivated and observationally tested model of reionization \\cite{CF05,CF06}. Although the two feedback prescriptions predicts quite different reionization and thermal histories (see Fig.~\\ref{history}), in both scenarios it is possible to reproduce a wide range of observational data with a proper choice of few model parameters (the redshift evolution of Lyman-limit absorption systems, the Gunn-Peterson and electron scattering optical depths, the cosmic star formation history, and number counts in high-$z$ sources). Thus, we find that existing data are unable to discriminate among the two reionization histories \\cite{Schneider08}. We then explore alternative methods to break these degeneracies using future 21-cm experiment \\cite{Schneider08} and CMB anisotropy observations \\cite{Burigana08}.  \\begin{figure}[t!] \\resizebox{\\hsize}{!}{\\includegraphics[clip=true]{salvaterra.fig1.ps}} \\caption{\\footnotesize {\\it Top panel:} redshift evolution of the spin (thick lines) and gas kinetic (thin lines) temperatures  predicted by the two models. Solid (dashed) lines refer to suppression (filtering) model. For comparison, we also show the evolution of the CMB temperature $T_\\gamma$ (dotted line). {\\it Bottom panel:} corresponding evolution of the free electron fraction. } \\label{history} \\end{figure} ",
        "conclusions": "We have explored the effect of cosmic radiative feedback from the sources of reionization on the thermal evolution of the intergalactic medium. We implemented two different prescriptions for this feedback into a well-tested, physically-motived model of the early Universe. We found that different prescriptions for this feedback predict quite different thermal and reionization histories. In spite of this differences, current data can not discriminate among different reionization scenarios. Therefore, we explored alternative methods to break this degeneracies using 21-cm experiment and CMB anisotropy observations. We found that future data can distinguish among different reionization histories provided that we would be able to remove the foreground signal at the percent (or better) level."
    },
    "0806/0806.1166_arXiv.txt": {
        "abstract": "In quasi-persistent neutron star transients, long outbursts cause the neutron star crust to be heated out of thermal equilibrium with the rest of the star.  During quiescence, the crust then cools back down.  Such crustal cooling has been observed in two quasi-persistent sources: \\ks{} and \\mxb{}. Here we present an additional \\chand{} observation of \\mxb{} in quiescence, which extends the baseline of monitoring to 6.6~yr after the end of the outburst.  This new observation strongly suggests that the crust has thermally relaxed, with the temperature remaining consistent over 1000~days.  Fitting the temperature cooling curve with an exponential plus constant model we determine an $e$-folding timescale of $465\\pm25$~days,  with the crust cooling to a constant surface temperature of $kT_{\\rm eff}^{\\infty} = 54\\pm2$~eV (assuming $D=10$ kpc).  From this, we infer a core temperature in the range $(3.5-8.3)\\times10^{7}$~K (assuming $D=10$ kpc), with the uncertainty due to the surface composition.  Importantly, we tested two neutron star atmosphere models as well as a blackbody model, and found that the thermal relaxation time of the crust is independent of the chosen model and the assumed distance. ",
        "introduction": "Many low-mass X-ray binaries (LMXBs) are transient, spending the majority of their time in a quiescent state with very low levels of accretion and a small fraction of their time in outburst where the mass accretion rate (and hence X-ray luminosity) increases significantly.  For neutron stars, these repeated outbursts affect the star --  the compression of the crust due to accretion of matter induces electron captures, neutron emissions and pycnonuclear reactions \\citep{haensel90} in the crust, which in turn heat the core.  Over approximately $10^{4}\\textrm{--}10^{5}$~yr a steady-state is reached, in which this deep crustal heating during outburst is balanced by cooling during quiescence \\citep{brown98}.  Typically outbursts in these neutron star X-ray transients only last weeks to months; however, in the so-called quasi-persistent transients, outbursts last several years.  In these objects enough heat is imparted to the neutron star crust that it gets heated significantly out of thermal equilibrium with the rest of the star, which is not the case for the majority of normal X-ray transients.  Therefore, in these quasi-persistent sources, once the source returns to quiescence the crust cools down into thermal equilibrium with the core over an appreciable timescale \\citep[e.g.][]{rutledge02}.  So far, crustal cooling curves have been observed in two such sources: \\mxb{} and \\ks{} \\citep{wijnands01,wijnands03,wijnands04,cackett06}.  For both objects we obtained several \\chand{} and \\xmm{} observations over a period of approximately 4 years after the end of long outbursts. \\citet{cackett06} found that both sources cooled down rapidly, with the cooling well fit by an exponential decay to a constant level.  This was interpreted as the neutron star crust cooling down into thermal equilibrium with the rest of the star.  By comparison to crustal cooling models of \\citet{rutledge02} the cooling curves suggest that the crusts both have a high thermal conductivity, and the cores may require enhanced levels of neutrino emission.  These observations motivated further theoretical study of crustal cooling, and models calculated by \\citet{shternin07} for \\ks{} also have a best fit with a high thermal conductivity crust, but may not require enhanced neutrino emission in the core. In this paper, we present an additional \\chand{} observation of \\mxb{}, extending the quiescent monitoring by 2.8 yr to now cover 6.6 yr after the end of the outburst.  In section \\ref{sec:results} we detail the data analysis and present the results, while in section \\ref{sec:disc} we discuss the cooling curve and implications for the neutron star structure. ",
        "conclusions": "\\label{sec:disc}  We presented a new observation of the quasi-persistent neutron star X-ray transient \\mxb{} in quiescence, extending the quiescent monitoring to 6.6~yr.  Results from the first 6 observations showed that the source had cooled rapidly and indicated that the neutron star crust may have returned to thermal equilibrium with the core.  This new observation shows that the neutron star temperature and flux remained consistent with the previous two \\chand{} observations performed approximately 1000 days before.  The model dependence of the thermal relaxation timescale was investigated with 2 different neutron star atmosphere models as well as a simple blackbody model.  Moreover, we assumed 3 different distances to \\mxb.  The $e$-folding timescales of the cooling curves from all the spectral fits are consistent with each other, demonstrating the robustness of the measurement. The results are consistent with fits to the first 6 observations \\citep{cackett06}.  With the crust thermally relaxed, we can compute the core temperature (here we assume $D=10\\nsp\\kilo\\parsec$).  We integrate the thermal structure equation in the neutron star envelope, following the calculation in \\citet{brown02}.  The inferred core temperature is relatively insensitive to the mass of the neutron star (the proper $T_{\\mathrm{eff}}$ increases with redshift for a fixed \\Teffinfty, but the surface layer becomes thinner with increasing $g$ and reduces the rise in temperature in the  envelope).  There is a significant uncertainty resulting from the depth of the light element layer, however \\citep{brown02}. We find the inferred core temperature to range from $3.5\\ee{7}\\nsp\\K$, for a deep He layer (column of $10^{8}\\nsp\\gram\\usp\\cm^{-2}$) overlaying a pure Fe layer to $8.3\\ee{7}\\nsp\\K$ for an shallow He layer (column of $10^{4}\\nsp\\gram\\usp\\cm^{-2}$) overlaying a layer of heavy  rp-process ashes.  We estimate the time-averaged mass accretion rate for the system to be $7\\ee{-11}\\nsp\\Msun\\usp\\yr^{-1}$, based on the known outburst behavior of 2 outbursts lasting approximately 2.5 yr, with a quiescent period of 21 yr. We estimated the fraction of Eddington luminosity by calculating the ratio of the average persistent outburst flux to the peak type-I X-ray burst flux \\citep[taken from][]{wijnands02}. From this, we estimated the time-averaged crust nuclear heating, assuming the heat deposited in the crust is 1.5 MeV per nucleon \\citep[see][]{brown98,rutledge02b}, to be $\\sim6\\ee{33}\\nsp\\ergs\\usp\\second^{-1}$, for a distance of $10\\nsp\\kilo\\parsec$.  As noted previously \\citep{cackett06,heinke07}, even for the highest core temperature compatible with $T_{\\mathrm{eff}}^{\\infty}$, the neutrino luminosity resulting from modified Urca cooling \\citep[for a review, see][]{yakovlev04} would still be a factor of $\\approx 30$ less than the time-averaged crust nuclear heating. As a check on whether there is a need for enhanced cooling, we also computed the neutrino cooling according to the ``minimal cooling model'' \\citep{page04}, which includes the pair breaking and formation (PBF) neutrino emissivity, but with the $^{1}S_{0}$ channel suppressed following \\citet{steiner08}. In this case the PBF neutrino luminosity from neutrons in the $^{3}P_{2}$ state is still sufficient to balance the time-averaged nuclear heating, if the core temperature is in the upper half of the range given above. Given the uncertainty in the depth of the light element envelope and the superfluid critical temperatures, we cannot exclude that the neutrino emission is solely due to standard cooling processes.  \\begin{figure} \\centering \\includegraphics[angle=270]{f1.eps} \\caption{Flux (top) and temperature (bottom) cooling curves for \\mxb{}. The best-fitting model to the temperatures comprising of an exponential decay to a constant level is shown (solid line), where the constant offset is shown with a dashed line.  The flux cooling curve is then calculated from the fit to the temperatures.  The time t$_0$ = 52159.5 (MJD) is the last day when the source was seen to be active.  Error bars are 1$\\sigma$. Data points are from the nsa fits with the distance fixed at 10 kpc.} \\label{fig:cooling} \\end{figure}  The exponential plus constant cooling curve fits provide a measure of the thermal relaxation time of the crust.  This relaxation time depends on the crust composition and lattice structure \\citep{rutledge02,shternin07}, on the crust thickness and hence surface gravity of the neutron star \\citep{lattimer94}, and on the distribution of heat sources (\\citealt{shternin07,horowitz08a}; Brown \\& Cumming, in preparation). \\citet{shternin07} showed that the cooling timescale in \\ks\\ was best fit by having a high thermal conductivity in the crust, as if it were composed of a locally pure lattice.  This matches molecular dynamics simulations \\citep{horowitz07, horowitz08b}, which find that the dense crust plasma does indeed freeze into an ordered lattice with a high thermal conductivity. As in \\ks, our fits to the cooling of \\mxb\\ are again consistent with such an ordered, low-impurity crust. \\citet{shternin07} noted that the crust may not have completely thermally relaxed; we find a single power-law decay also fits the cooling curve for \\ks{} well, with a power-law index = $-0.12\\pm0.01$ and $\\chi^2_\\nu = 0.2$ when fitting to the temperatures. Further observations are required to determine whether \\ks{} has continued to cool following a power-law decay or if it has reached a constant \\kTeffinfty\\ indicative of a thermally relaxed crust.  \\citet{shternin07} note that the cooling of \\ks\\ can be fit without invoking enhanced neutrino emission in the core.  If \\mxb\\ does have a higher core neutrino emission than \\ks\\ then this may imply that the neutron star in \\mxb\\ is somewhat more massive than the one in \\ks\\ since most equations of state allow higher levels of neutrino emission with increasing central density.  We note, however, that this conclusion assumes that the time-averaged mass accretion rates in both these objects have remained in a steady state."
    },
    "0806/0806.3018.txt": {
        "abstract": "Any calibration of the present value of the Hubble constant ($H_{0}$) requires recession velocities and distances of galaxies. While the conversion of observed velocities into true recession velocities has only a small effect on the result, the derivation of unbiased distances which rest on a solid zero point and cover a useful range of about 4 to 30$\\;$Mpc is crucial. A list of 279 such galaxy distances within $v<2000\\kms$ is given which are derived from the tip of the red-giant branch (TRGB), from Cepheids, and/or from supernovae of type Ia (SNe\\,Ia). Their random errors are not more than $0.15\\mag$ as shown by intercomparison. They trace a {\\em linear\\/} expansion field within narrow margins, supported also by external evidence, from $v=250$ to at least $2000\\kms$. Additional 62 distant SNe\\,Ia confirm the linearity to at least $20,000\\kms$. The dispersion about the Hubble line is dominated by random peculiar velocities, amounting locally to $<100\\kms$ but increasing outwards. Due to the linearity of the expansion field the Hubble constant $H_{0}$ can be found at any distance $>4.5\\;$Mpc. RR\\,Lyr star-calibrated TRGB distances of 78 galaxies above this limit give $H_{0}=63.0\\pm1.6$ at an effective distance of $6\\;$Mpc. They compensate the effect of peculiar motions by their large number. Support for this result comes from 28 {\\em independently\\/} calibrated Cepheids that give $H_{0}=63.4\\pm1.7$ at 15 Mpc. This agrees also with the large-scale value of $H_{0}=61.2\\pm0.5$ from the distant, Cepheid-calibrated SNe\\,Ia. A mean value of $H_{0}=62.3\\pm1.3$ is adopted. Because the value depends on two independent zero points of the distance scale its systematic error is estimated to be 6\\%. -- Other determinations of $H_{0}$ are discussed. They either conform with the quoted value (e.g.\\ line width data of spirals or the $D_{n}\\!-\\!\\sigma$ method of E galaxies) or are judged to be inconclusive. Typical errors of $H_{0}$ come from the use of a universal, yet unjustified P-L relation of Cepheids, the neglect of selection bias in magnitude-limited samples, or they are inherent to the adopted models. % ",
        "introduction": "\\label{sec:1} % It is said sometimes that once in a career, every astronomer is entitled to write a paper on the value of the Hubble constant. To the point, several compilations of the growing literature on $H_{0}$ since 1970 have been made. Those by \\citet{Press:97}, \\citet{Tammann:06} and \\citet{Huchra:07} are examples.  These authors plot histograms of the distribution of $H_{0}$ from about 400 papers since 1970. The sample is so large that the formal error on the average of the histogram is so small that one might infer that the Hubble constant is now known to better than say 1\\%. Of course, what is missing is the fact that most of the values in the literature are not correct. Many suffer from the neglect of the effects of an observational selection bias that {\\em varies with distance}.  We are faced with a problem in writing this review. Do we strive to give a comprehensive history of the distance scale problem beginning with the first determination of the Hubble constant by \\citet{Lemaitre:27,Lemaitre:31,Robertson:28,Hubble:29b}, and \\citet{Hubble:Humason:31,Hubble:Humason:34} to be about $550\\;\\mbox{km}\\;\\mbox{s}^{-1}\\;\\mbox{Mpc}^{-1}$ (units assumed hereafter), coming into modern times with the debates between the principal players? Or do we only write about the situation as it exists today, comparing the ``concordance'' value of $H_{0}=72$ by \\citet{Freedman:etal:01} with the {\\em HST\\/} supernovae calibration value (\\citealt{Hamuy:etal:96}; \\citealt{Tripp:Branch:99}; \\citealt{Suntzeff:etal:99}; \\citealt{STT:06}, hereafter \\citeauthor*{STT:06}; \\citealt{STS:06}, hereafter \\citeauthor*{STS:06}) that gives $H_{0}=62$? We have decided to take the latter course but also to sketch as a skeleton the beginning of the correction to Hubble's 1930-1950 distance scale that started with the commissioning of the 200-inch telescope in 1949. An important comprehensive review of this early period before the Hubble Space Telescope ({\\em HST}) is by \\citet{Rowan-Robinson:85}; the details are not repeated here.    % ******************************************************** % 1.1. Early work on the revision to Hubble's distance scale (1950-1990) % ******************************************************** \\subsection{Early work on the revision to Hubble's distance scale (1950-1990)} \\label{sec:01:1} % Hubble's extragalactic distance scale was generally believed from 1927 to about 1950, beginning with the first determinations of the Hubble constant by the four independent authors cited above. This scale lasted until Hubble's \\citeyearpar{Hubble:29a} distance to M\\,31 was nearly tripled by \\citet{Baade:54} in his report to the 1952 Rome meeting of the IAU. He proposed a revision of the Cepheid P-L relation zero point by about $1.5\\mag$ based on his discovery that RR\\,Lyrae stars could not be detected with the newly commissioned 200-inch Palomar telescope in M\\,31 at $V=22.2$. From this he concluded that M\\,31 must be well beyond the modulus of $(m-M)=22.2$ given earlier by Hubble. The story is well known and is recounted again by \\citet[][Chapter~6]{Osterbrock:01}, in the Introduction to \\citet{TSR:08}, hereafter \\citeauthor*{TSR:08}, and often in histories elsewhere \\citep[e.g.][]{Trimble:96,Sandage:99a} .  Following Baade's discovery, the revision of 1930-1950 scale was begun anew with the Palomar 200-inch telescope, largely following \\citet*{Hubble:51} proposed cosmological program for it. Observational work on the first Cepheid distance beyond the Local Group was completed for NGC\\,2403. Here we made photoelectric calibrations of magnitude scales and used new calibrations of the Cepheid P-L relations \\citep{Kraft:61,Kraft:63,ST:68,ST:69}, and we obtained a revised distance modulus of $(m-M)=27.56$ \\citep{TS:68}. Comparing this with Hubble's modulus of 24.0 showed the large scale difference by a factor of 5.2. Next, the modulus of the more remote galaxy, M\\,101, was determined to be $(m-M) = 29.3$ \\citep{ST:74a} compared with Hubble's modulus of 24.0, giving the large correction factor of 11.5 to Hubble's scale at M\\,101 ($D = 7.2\\;$Mpc). This large stretching was again found in our distance modulus of $(m-M)=31.7$ for the Virgo cluster \\citep{ST:74b,ST:76,ST:90,ST:95}, compared with Hubble's modulus of 26.8. The distance ratio here is a factor of 9.6.  These large factors and their progression with distance came as a major shock in the mid 1970s and were not generally believed \\citep[eg.][etc.]{Madore:76,Hanes:82,deVaucouleurs:82}. However, the new large distances were confirmed for NGC\\,2403 by \\citet{Freedman:Madore:88}, and for M\\,101 by \\citet{Kelson:etal:96} and \\citet{Kennicutt:etal:98}. Although our distance to the Virgo cluster core is still in contention at the 20\\% level, there is no question that the correction factor here is also between 7 and 10 at $20\\;$Mpc.    % ******************************************************** % 1.2. The difficulty of finding H_0 % ******************************************************** \\subsection{The difficulty of finding $H_{0}$} \\label{sec:01:2} % The determination of $H_{0}$, the {\\em present\\/} and hence nearby value of the Hubble parameter, requires -- besides true recession velocities -- distance indicators with known zero point {\\em and\\/} with known intrinsic dispersion. The scatter of the Hubble diagram, $\\log v$ versus $m$ or $(m-M)$, would in principle be a good diagnostic for the goodness of a given distance indicator if it were not also caused by peculiar motions. It is of prime importance to disentangle these two sources of scatter because unacknowledged intrinsic scatter of the available distances introduces a {\\em systematic\\/} increase of $H_{0}$ with distance if {\\em flux-limited samples\\/} are considered, which is normally the case. This is because the mean absolute magnitude of objects in such samples increases with distance due to the increasing discrimination against the less luminous objects. It is important to note that, strictly speaking, this incompleteness bias is {\\em not\\/} the \\citet{Malmquist:20,Malmquist:22} bias which applies only to the {\\em average\\/} effect integrated over the sample being studied; not to individual distances within that sample, each of which must be corrected by a sliding scale.  Neglect of the individual bias values that become progressively larger with increasing distance always gives a Hubble constant that incorrectly appears to increase outward \\citep{deVaucouleurs:58,deVaucouleurs:76,deVaucouleurs:77,Tully:88}.  The widely held view that the increase of $H_{0}$ with distance (up to an unspecified limit) was real deprived the Hubble diagram of its second diagnostic power. The slope of the Hubble line had no longer to be $0.2$, which is the case for {\\em linear\\/} expansion (see hereafter Eq.~\\ref{eq:01}). The apparent increase of $H_{0}$ with distance was not anymore accepted as proof for bias (e.g.\\ \\citealt{Tammann:87} versus \\citealt{Aaronson:87}). It also led to proposals that $H_{0}$ not only varied with distance, but also with direction \\citep{deVaucouleurs:Bollinger:79,deVaucouleurs:Peters:85}. The search for the asymptotic value of $H_{0}$ became self-defeating: one tried to calibrate it at the largest possible distances where, however, the effects of bias are largest.  The bias is always present in a flux limited sample of field galaxies \\citep[][as analyzed using Spaenhauer diagrams]{% Sandage:94a,Sandage:94b,Sandage:95,FST:94}. It is also present in cluster data that are incomplete \\citep{Teerikorpi:87,Teerikorpi:90,Kraan-Korteweg:etal:88, Fouque:etal:90,STF:95,Sandage:08}, and even in field galaxies of any sample that {\\em is\\/} distance limited but if the data are incomplete in the coverage of the distance indicator (apparent magnitude, 21cm line width, etc.) \\citep{Sandage:08}.  However, claims for $H_{0}$ increasing outwards were contradicted by the apparent magnitudes of first-ranked galaxies in clusters and groups. The Hubble diagram of brightest cluster galaxies shows no deviations from linear expansion down to $\\sim\\!2000\\kms$ \\citep[][and references therein]{STH:72,Sandage:Hardy:73,Kristian:etal:78}. This was confirmed down to $\\sim\\!1000\\kms$  in a study of northern and southern groups \\citep{Sandage:75}, which also showed a smooth linear Hubble diagram with no discontinuities over the range of $1000<v<10,000\\kms$. The limit on $\\delta H_{0}/H_{0}$ was $<0.08$, and a proof was given that the Hubble constant does not increase outward. These results were confirmed by \\citet{FST:94} based on the large catalog of 21cm line widths and $I$ magnitudes by \\citet{Mathewson:etal:92a,Mathewson:etal:92b}, and also in the large archive literature cited therein by many others. However, it was so far not possible to tie the local expansion field below $\\la15\\;$Mpc into the large-scale field because of small-number statistics and of large scatter caused by the important effects of peculiar velocities and distance errors. This problem is the subject of Sect.~\\ref{sec:2}.  In parallel to the discussion on distance errors there were many attempts to determine the mean size of the random one-dimensional peculiar velocities $v_{\\rm pec}$ by reading the deviations from the Hubble line vertically as velocity residuals, but this is not easier than to determine the dispersion of the distance indicators because the latter have to be known. In fact the problem is here even deeper. The halted expansion of the Local Group, the retarded expansion by the gravity of the Virgo complex, the large virial velocities in clusters, and the increase of peculiar motions with distance, as manifested by the important velocity of a large volume with respect to the CMB dipole all make it difficult to find the characteristic peculiar velocities of field galaxies.  One of the earliest attempts to determine a cosmological parameter of interest (other than $H_{0}$) was that by Hubble and Humason to measure the mean random velocity of galaxies about an ideal Hubble flow. This, in turn, is related to any systematic streaming, or more complicated systematic motions (a dipole plus even a quadrupole, a shear, or a local rotation) relative to a cosmic frame (\\citealt{Davis:Peebles:83a} for a review; see also \\citealt{Dekel:94}). The discussion by \\citet{Hubble:Humason:31} gave values between $200$ and $300\\kms$ for the mean random motion (they do not quote an rms value) about the ridge line of the redshift-distance relation for local galaxies ($v < 10,000\\kms$).  By 1972 a limit was set of $v_{\\rm pec}<100\\kms$ on local scales \\citep{Sandage:72}. In subsequent papers, too numerous to be cited here, rather lower values were favored \\citep[e.g.][]{ST:75,Giraud:86,Sandage:86,Ekholm:etal:01,Thim:etal:03}. In a representative study \\citet{Karachentsev:Makarov:96} found $v_{\\rm pec}=72\\kms$, supported by later papers of Karachentsev and collaborators. The values of $v_{\\rm pec}$ in function of scale length agree locally (see Sect.~\\ref{sec:2:5}), but clearly increase with distance.  The modest size of the peculiar velocities poses a problem for various hierarchical merging scenarios of galaxy formation which predict mean random motions as high as $500\\kms$ \\citep[cf.][]{Davis:Peebles:83b,Davis:etal:85,Ostriker:93,Governato:etal:97,Leong:Saslaw:04}.     % ****************************************************************** % 2. The local expansion field % ****************************************************************** ",
        "conclusions": "\\label{sec:6} % An intercomparison of RR\\,Lyr, TRGB, and Cepheid distances shows that their dispersion is not more than $0.15\\mag$. The same upper limit holds for SNe\\,Ia as seen from the small scatter in their Hubble diagram at large distances. The four distance indicators stand out because they can provide the most accurate distances within their reach for sizable samples of galaxies and, importantly, their small dispersion makes them highly insensitive to selection bias. Although their reach is drastically different, RR\\,Lyr stars being very short-range, SNe\\,Ia extending to cosmological distances, and the TRGB and Cepheid distances lying in between, there is enough overlap to tie them into a single system of distances.  The combined Hubble diagram of TRGB, Cepheid, and SNe\\,Ia distances shows a well defined Hubble line with slope 0.2, corresponding to linear expansion, over a range of $\\sim\\!250$ to at least $20,000\\kms$. The slope of 0.2, strongly supported also by other evidence (see Sect.~\\ref{sec:2}) implies that the present mean value of the Hubble constant $H_{0}$ {\\em is everywhere the same} (cosmological effects being exempt by definition). Most of the observed dispersion about the Hubble line must be caused by random peculiar motions; allowing for the (small) distance errors they are $70\\kms$ within $7\\;$Mpc and increase outwards to a yet undetermined limit (see Table~\\ref{tab:char}, Col.~10). Lower values are in the literature, but the value here seems well determined from 78 TRGB distances  (Table~\\ref{tab:char}).  The zero point of the Hubble line is set in two {\\em independent\\/} ways. a) The absolute magnitude of the TRGB is determined by 22 RR\\,Lyr star distances and agrees well with other determinations. The adopted magnitude of $M^{I}_{\\rm TRGB}=-4.05$ carries a systematic error of hardly more than $0.1\\mag$. The value holds for [Fe/H]$_{\\rm ZW}=-1.6$ and changes by less than $0.1\\mag$ in the range $-2.0<\\mbox{[Fe/H]}_{\\rm ZW}<-1.3$ typical for old populations (\\citeauthor*{TSR:08}, Fig.~1). The resulting value of $H_{0}=63.0\\pm1.6\\;(\\pm3)$ from 78 distances larger than $4.5\\;$Mpc refers to an effective distance of only $\\sim\\!400\\kms$ where the influence of peculiar velocities is still large, but this is compensated by the large number of TRGB distances. b) Because the P-L relations of the metal-rich Galactic Cepheids and of the metal-poor LMC Cepheids are different they are independently calibrated. The zero point of the Galactic P-L relation rests on Cepheids in Galactic clusters and on physical BBW distances. The zero point of a 10-day Cepheid is confirmed by trigonometric parallaxes to within a few $0.01\\mag$, but the error can increase to $\\sim0.15\\mag$ for Cepheids with the shortest and longest periods depending on the correctness of the adopted P-L slope. The LMC P-L relation with very well determined slope is zero-pointed by an adopted distance of $\\mu^{0}_{\\rm LMC} = 18.54$. This value is based on a multitude of determinations, excluding of course results depending on the P-L relations of Cepheids themselves; the error is again estimated to be $0.1\\mag$. It should be noted that significantly smaller LMC distances come mostly from some a priory assumption on the shape and zero point of the P-L relations of the Galaxy and LMC. -- The Cepheids in other galaxies with metallicities like the Galaxy or LMC are reduced with the corresponding P-L relations; in case of Cepheids with intermediate metallicities the results from the two P-L relations are interpolated. The resulting 29 Cepheid distances larger than $4.5\\;$Mpc give $H_{0}=63.4\\pm1.7$ at an effective distance of $900\\kms$. -- The good agreement of the value of $H_{0}$ from the TRGB and Cepheid distances is highly significant because it is {\\em predicted\\/} from the well supported linearity of the expansion field.  SNe\\,Ia are calibrated through Cepheids and cannot independently contribute to the zero point of the distance scale. But their large number can reduce the statistical error and serve to carry the value of $H_{0}$ to $\\sim\\!20\\,000\\kms$. They give $H_{0}=60.3\\pm2.6$ at an effective distance of $1600\\kms$ and, allowing for a flat Universe with $\\Omega_{\\Lambda}=0.7$, $H_{0}=61.2\\pm0.5$ from 62 SNe\\,Ia at $v>3000\\kms$. The adopted value of \\begin{equation} H_{0} = 62.3\\pm1.3\\quad (\\pm4.0) \\label{eq:H0adopt} \\end{equation} is the unweighted mean from the Cepheids and Cepheid-calibrated SNe\\,Ia averaged with the result from the TRGB. The generous 6\\% systematic error is estimated in Sect.~\\ref{sec:3:4:2}.  The value of $H_{0}$ rests on the two independent zero points set by the TRGB and Cepheid distances. No other zero-pointed distance indicator is available at present, which could carry the distance scale into the expansion field, i.e.\\ to $>4.5\\;$Mpc, for a sufficient number of 20 or more galaxies. But TF distances of a distance-limited sample of spiral galaxies and $D_{n}\\!-\\!\\sigma$ distances out to $2500\\kms$ as well the Hubble diagram of nearby groups and clusters provide at least a consistency check. We are not aware of any serious objection against the adopted value of $H_{0}$.  The literature abounds in larger values of $H_{0}$. Some are based on the untenable view that the LMC P-L relation of Cepheids, whatever its exact shape and zero point, is universal. Others are the result of selection bias, which becomes particularly severe when it is tried to determine $H_{0}$ at the largest distances which can be reached, and from where necessarily only the most luminous objects of their species can enter the catalogs. The importance of selection bias is often underestimated because the quality of the distance indicators is overestimated. The true quality can be determined only if there is broad overlap with high-accuracy distance indicators like RR\\,Lyr stars or TRGB and Cepheid distances, or by consulting the Hubble diagram. The dispersion here, corrected for the reasonably well understood effect of peculiar velocities, gives the random error for a given distance indicator. Also too steep a slope, i.e.\\ $H_{0}$ increasing with distance, is a clear sign of important bias or some other systematic problem of the method. Finally other high values of $H_{0}$ are too model-dependent to be reliable.  Future progress on $H_{0}$ will come from additional near-infrared photometry of Cepheids where they are relatively insensitive to absorption and metallicity. Enormous potential lies still in the TRGB distances. With a somewhat improved understanding of their metallicity dependence, which is in any case small in old populations, they can provide distances to better than $\\pm5\\%$ for well over 1000 galaxies of all types within $\\sim\\!20\\;$Mpc with present techniques and requiring relatively little telescope time. They will thus map the local velocity field in great detail and also yield a high-weight calibration of SNe\\,Ia extending the impact of the method to cosmological distances.   % *********************************************** %  Acknowledgments % ***********************************************"
    },
    "0806/0806.1182.txt": {
        "abstract": "% context heading (optional) % {} leave it empty if necessary {} % aims heading (mandatory) {With this paper we want to investigate the highly variable afterglow light curve and environment of gamma-ray burst (GRB) 060526 at $z=3.221$.} % methods heading (mandatory) {We present one of the largest photometric datasets ever obtained for a GRB afterglow, consisting of multi-color photometric data from the ultraviolet to the near infrared. The data set contains 412 data points in total to which we add additional data from the literature. Furthermore, we present low-resolution high signal-to-noise spectra of the afterglow. The afterglow light curve is modeled with both an analytical model using broken power law fits and with a broad-band numerical model which includes energy injections. The absorption lines detected in the spectra are used to derive column densities using a multi-ion single-component curve-of-growth analysis from which we derive the metallicity of the host of GRB 060526.} % results heading (mandatory) {The temporal behaviour of the afterglow follows a double broken power law with breaks at $t=0.090\\pm0.005$ and $t=2.401\\pm0.061$ days. It shows deviations from the smooth set of power laws that can be modeled by additional energy injections from the central engine, although some significant microvariability remains. The broadband spectral-energy distribution of the afterglow shows no significant extinction along the line of sight. The metallicity derived from \\ion{S}{II} and \\ion{Fe}{II} of [S/H] = --0.57 $\\pm$0.25 and [Fe/H] = --1.09$\\pm$0.24 is relatively high for a galaxy at that redshift but comparable to the metallicity of other GRB hosts at similar redshifts. At the position of the afterglow, no host is detected to F775W(AB) = 28.5 mag with the HST, implying an absolute magnitude of the host M(1500 \\AA{})$>$--18.3 mag which is fainter than most long-duration hosts, although the GRB may be associated with a faint galaxy at a distance of 11 kpc.} % conclusions heading (optional), leave it empty if necessary {} ",
        "introduction": "Gamma-ray bursts (GRBs) and their afterglows offer a powerful tool to probe the high-redshift universe, both through photometry and spectroscopy. The standard fireball model of GRB afterglows \\citep[see][for recent reviews on the topic]{Zhang07, Meszaros06, Gehrels09} predicts a smooth temporal evolution, and the resulting afterglow light curve can be empirically described by a joint smoothly-broken power law, the so-called Beuermann equation \\citep{Beuermann1999}. For many GRBs in the pre-{\\it Swift} era, when  temporally dense afterglow photometry was obtained, the afterglow evolution was found to be smooth as, e.g., in GRB 020813 \\citep{LaursenStanek2003}, GRB 030226 \\citep{Klose2004} and GRB 041006 \\citep{Stanek2005} \\citep[for the complete pre-\\emph{Swift} sample see][]{ZKK}. Out of the total sample of 59 afterglows of that paper, only four of the GRBs analysed showed significant deviations from the expected decay (though we note that about half of the 59 afterglows were not sampled well enough to allow any conclusions). One might be explained by microlensing \\citep[GRB 000301C,][]{GLS2000} and another one by assuming an  inhomogeneously emitting surface \\citep[GRB 011211,][]{JakobssonNA2004}. The other two GRBs which are the only ones showing long-lasting strong deviations are GRB 021004 \\citep[e.g.][and references therein]{deUgarte2005} and GRB 030329 \\citep[e.g.][and references therein]{Lipkin2004}, which incidentally also have the densest optical monitoring.  With the launch of the \\emph{Swift} satellite and its rapid localization capabilities \\citep{Gehrels2004}, the number of highly variable light curves has increased considerably, though there are still examples of very smooth light curves (e.g. that of GRB 080210, which showed a smooth behaviour with a high sampling rate of 1\\,s, A. De Cia, priv. comm.). Some light curves show small, achromatic bumps overlying the smooth power law decay (e.g., GRB 050502A, \\citealt{Guidorzi05}; GRB 061007, \\citealt{Mundell07}; GRBs 090323 and 090328, \\citealt{McBreenLAT}), early bumps with chromatic evolution (e.g., GRB 061126, \\citealt{Perley08a}; GRB 071003, \\citealt{Perley08b}), ``steps'' due to energy injection episodes (e.g., GRB 070125, \\citealt{Updike08}; GRB 071010A, \\citealt{Covino08}; GRB 080913, \\citealt{Greiner080913}, GRB 090926A, \\citealt{Rau090926A, CenkoLAT, Swenson090926A}) or powerful late-time rebrightenings of up to several magnitudes (e.g., GRB 050721, \\citealt{Antonelli06}; GRB 060206, \\citealt{Wozniak060206, Monfardini060206, Stanek060206}; GRB 070311, \\citealt{Guidorzi07}; GRB 071003, \\citealt{Perley08b}). The early time domain, which can now be routinely accessed by rapid follow-up in the \\emph{Swift} era, has yielded more types of variability, like rising afterglows (e.g., GRB 060418, \\citealt{Molinari07}; GRB 060605, \\citealt{Ferrero08}; GRB 060607A, \\citealt{Nysewander07}, \\citealt{Molinari07}; GRB 081008, \\citealt{Yuan10}; see \\citealt{Oates09} and \\citealt{Rykoff09} for further examples) and short-term variability directly linked to the prompt emission (e.g., GRB 041219A, \\citealt{Vestrand041219A, Blake041219A}; GRB 050820A, \\citealt{Vestrand050820A}; GRB 080319B, \\citealt{Racusin08}; GRB 080129, \\citealt{Greiner080129}). In all these cases, dense photometric follow-up during the periods of variability was needed to characterize the phenomena involved.  In addition their use in to studying the GRB phenomenon itself, GRB afterglows can be used to study their galactic environment through absorption line spectroscopy of material in the line-of-sight towards the GRB. Above a redshift of around $z\\sim0.2$, resonant absorption lines from elements present in the interstellar medium (ISM) such as Mg, Zn, Fe, Si, C, and S are shifted into the optical regime and can be studied with ground-based telescopes. For 10 bursts between a redshift of $z=2$ to 6, the metallicity along the line-of-sight in the galaxy could be obtained so far \\citep[e.g.][and references therein]{Savaglio06, Fynbo06a, Price07, Prochaska07, Ledoux09}. The values are usually below solar, but higher than for QSO absorbers at comparable redshifts, some of them even higher than theoretical limits for the formation of collapsars \\citep{Woosley06}. The difference to QSO absorbers can be explained if GRB sightlines probe denser parts of the galaxy, or if GRBs reside in galaxies with higher masses and therefore higher metallicities \\citep{Fynbo08, Pontzen10}. Both for QSO and GRB absorbers, there seems to be a metallicity evolution with redshift \\citep{Savaglio06, Fynbo06a, Price07}, although the slope is different for the two samples. GRB hosts seem to show a low extinction along the line-of-sight \\citep{Prochaska07}, however, relative abundances of heavier elements indicate that some of the ions must be depleted onto dust grains and the depletion pattern resembles the one found in the warm disc and halo of the Milky Way (MW; \\citealt{Savaglio06}). How this can be explained together with the low extinction as also derived from the spectral energy distribution (SED) of the afterglow \\citep{Kann2006, Starling07, Kann07}, is still an open question. One possible solution to this problem is destruction of the dust present in the line-of-sight by the GRB and afterglow radiation \\citep[e.g.,][]{Waxman00, Perna02}, though no strong evidence has ever been found for this. Furthermore, there is clearly a strong observational bias involved, as those GRB afterglows with successful spectroscopy, especially in the case of high-resolution observations, are those which have only low extinctions and thus relatively bright afterglows \\citep{Fynbo09, Kann07}, although rapid observations with large telescopes can achieve detailed spectroscopy of highly extinguished afterglows, as in the case of GRB 080607 \\citep{Prochaska09, Sheffer09, Fynbo09}.  GRB 060526 was detected by the \\emph{Swift} satellite on May 26.686458 (16:28:29.95 UT). The satellite slewed immediately to the burst, detecting both the X-ray and the optical afterglow \\citep{CampanaGCN1}. The BAT instrument on \\emph{Swift} measured two emission episodes. The first one lasted 13.8 s and consisted of two FRED (fast rise exponential decay) peaks, followed by a second symmetric peak between 230 and 270 s \\citep{CampanaGCN2}. The second peak was coincident with a giant X-ray flare followed by a softer flare at 310 s \\citep{CampanaGCN3} also detected in the $v$ band by the UVOT telescope on-board \\emph{Swift} \\citep{BrownGCN}. The gamma-ray fluence was $(4.9\\pm0.6)\\times10^{-7}$ erg cm$^{-2}$ during the first emission episode and $(5.9\\pm0.6)\\times10^{-7}$ erg cm$^{-2}$ during the second, the peak flux of the second episode was however only half of the peak flux of the first one. The photon index of the two epochs changed from $1.66\\pm0.20$ to $2.07\\pm0.18$, thus showing the typical hard-to-soft evolution \\citep{MarkwardtGCN}. The \\emph{Watcher} telescope provided the first ground based detection with $R\\approx15$ mag \\citep{FrenchGCN} 36.2 s after the trigger. \\emph{ROTSE} observations showed a plateau for several thousand seconds after the GRB onset \\citep{RykoffGCN}. A redshift of $z=3.21$ was determined by \\cite{BergerGCN} with the Magellan/Clay telescope. The brightness of the optical afterglow allowed for a dense monitoring which revealed a complex light curve structure including several flares \\citep{HalpernGCN1, HalpernGCN2} and a steepening attributed to a jet break \\citep{ThoeneGCN}.  In this paper, we approach the analysis of GRB 060526 from two directions: Through modelling of the very detailed optical light curve and late optical imaging of the field to detect the host in Sec. 3, and analysis of low to medium resolution spectroscopic observations of absorption lines along the line of sight (Sec. 4). Throughout the paper, we follow the convention $F_\\nu (t)\\propto t^{-\\alpha}\\nu^{-\\beta}$, and use WMAP concordant cosmology \\citep{Spergel2003} with $H_0=71$km s$^{-1}$ Mpc$^{-1}$, $\\Omega_M=0.27$, and $\\Omega_{\\Lambda}=0.73$. Uncertainties are given at 68\\% confidence level for one parameter of interest unless stated otherwise.  %__________________________________________________________________ ",
        "conclusions": "GRB 060526 had a relatively bright afterglow that allowed us to obtain a solid dataset, both photometrically and spectroscopically. We achieved a dense light curve coverage over several days which allowed a detailed study of the afterglow properties, and obtained a series of low resolution but high signal-to-noise spectra to study the host environment.  %\\subsection{A highly variable light curve} The optical light curve can be fitted with a double smoothly-broken power law with a breaks at $t_{b1}=0.090\\pm0.005$ and $t_b=2.216\\pm0.049$ days, and decay slopes of $\\alpha_{plateau}=0.288\\pm0.026$, $\\alpha_1=0.971\\pm0.008$, and $\\alpha_2=2.524\\pm0.052$. The dense sampling of especially the $R_C$-band light curve also reveals additional variability on top of the power laws. These features could be explained either by extended activity of the central engine or through interactions of the shock with the interstellar medium. For the case that the variability arises from external shocks, several mechanisms have been considered. In GRB 021004, both density variations of the external medium into which the GRB jet plows and angular inhomogenities of the jet surface were considered \\citep{NPG03}. However, \\cite{Nakar07} show that density variations would cause much smaller fluctuations than those observed in GRB afterglows and can therefore be ruled out. Another possibility is the injection of additional energy into the shock by slower shells that catch up with the shocked region as it decelerates, this model was used successfully to describe GRB 021004 \\citep{deUgarte2005} and also works better than two other models (double jet and density fluctuations) to describe the highly complex light curve of GRB 030329 \\citep{Huang2006}. Thus, variability can give either information on the medium surrounding the GRB or on the activity of the central engine. A more intriguing possibility is that the flares may be emitted from another region closer to the central engine, resulting from late internal shocks. Powerful X-ray flares that are attributed to late central engine activity have been observed in about 50\\% of all \\emph{Swift} GRBs \\citep[e.g.][]{Burrows2005, Chincarini2007, Krimm2007}, and strong optical/NIR flaring contemporaneous with the GRB prompt emission may also occur \\citep{Vestrand041219A, Blake041219A, Vestrand050820A, Racusin08}, thus making optical flares from late central engine activity an interesting prospect \\citep{KannGCN, Malesani2007}.  Indeed, \\cite{Dai2007} have suggested that the optical variability of the afterglow of GRB 060526 is due to flares from late internal shocks (the very early rapid optical variability we present here is very probably due to central engine activity, as it is seen contemporaneously in gamma and X-rays). \\cite{Khamitov08}, on the other hand, conclude that the short timescale of the variabilities requires the jet to be non-relativistic already at $\\sim1$ day and could then be explained by external density fluctuations. Our analysis lends tentative support to the notion of flares from internal shocks, finding decay slopes for two flares that exceed what should be possible from external shocks. But we caution that the errors of these fits are large due to a low amount of data in the decaying parts. Furthermore, globally, a model using refreshed external shocks is able to account for the light curve variations, although microvariability remains. This creates the intriguing possibility of reverberation effects \\citep[see][for a case of reverberation between gamma-rays and optical emission]{Vestrand050820A}. Short flares in the X-ray or optical bands signal internal shocks from long-term central engine activity, and when these shells catch up with the forward shock front, they re-energise the external forward shock. The detection of such behaviour would probably require dense multi-band observations of a bright afterglow to search for SED changes at high time resolution combined with detailed modelling of the data. This way, one could discern between internal shocks (which are expected to have a different spectral index from the forward shock afterglow) and refreshed, external shocks (which are achromatic). Our data set of the afterglow of GRB 060526 does not allow us such a detailed decomposition.  %\\subsection{A moderate metallicity environment} From the analysis of our low resolution spectra with different resolutions, we detect a LLS and a number of metal absorption lines that all lie at a redshift of $z=3.221$. The low resolution only allows us to derive column densities from measuring the EWs of the absorption lines and adopt a MISC-CoG analysis where we exclude the most saturated as well as blended transitions. We find a best fit for the Doppler parameter of $b=39\\pm3$ km s$^{-1}$ and most of the ions used for the fit lie on the linear part of the CoG which allows a relatively reliable determination of the column densities. The relative abundances of different metals in the spectra indicate some dust extinction, but an intrinsic difference due to enhancement of the production of certain elements cannot be excluded. The very low amount of dust detected in the afterglow SED may indicate that the latter might be the favored possibility.  The column density of neutral hydrogen is rather low compared to other GRBs. We derive a metallicity for the host of $[{\\rm Fe}/{\\rm H}]=-1.09$ which is slightly higher than metallicities determined from other GRB afterglow spectra. According to the definition of QSO absorbers, the host of GRB 060526 is classified as a LLS ($19 < \\log N_{\\rm HI}/{\\rm cm}^{-2} < 20.3$), which seem to have on average higher metallicities than damped Lyman $\\alpha$ systems (DLA; \\citealt{Peroux07}) and a steeper evolution towards lower redshifts. Around redshift 3, however, the metallicities of both samples are within the same range. Also, GRB hosts show a trend towards increasing metallicity with lower redshifts \\citep{Fynbo06a, Savaglio06}. Taking into account that most of the sample used only low-resolution spectra to derive the metallicity (which only gives lower limits for the column densities and the metallicity) this evolution might, however, not be as pronounced as for DLAs and LLS. This might imply that the enrichment of the ISM in the early universe had taken place at earlier times than assumed. The absorbing material along the line-of-sight is mostly in the neutral state, as usually observed for long GRB-DLAs, while sightlines with lower $\\log N_{\\rm HI}$ often contain more ionized material. This might either imply that we have a very small host galaxy or that the GRB is placed somewhere in the outskirts of its host. In general, the low ionization points to a relatively large distance of the absorbing material from the GRB itself.  There is no underlying host galaxy of GRB\\,060526 detected down to a deep limit of 28.5 mag (in F775W AB) in HST/ACS data. At that redshift, this means the host has an absolute magnitude M(1500\\AA{}) $>-18.3$ mag, fainter than an 0.5 L* galaxy at that redshift. Long GRBs have been found to occur in actively star-forming galaxies and star formation is assumed to shift towards smaller and fainter galaxies over time \\citep[e.g.][]{Cowie96} while massive galaxies prove to be rather unchanged throughout the history of the universe \\citep[e.g.][]{Abraham99, Heavens04}. One would therefore expect that GRB hosts should also have higher luminosities towards higher redshifts. \\cite{Fynbo08} concluded that the observed metallicity distribution of GRB hosts (as well as QSO absorbers) at $z\\approx3$ can be explained by the luminosity function of galaxies at that redshift and assuming a luminosity-metallicity relation as derived for other high-redshift samples \\citep{Ledoux06, Erb06}. The non-detection of the host of GRB\\,060526 down to deep limits, however, would not support this suggested evolution. The data neither strongly support nor allow us to rule out that the GRB is associated with the nearby galaxy. While the offset would be very large compared to typical long GRB offsets, it is possible the burst occurred in a locally dense star-forming region which is not detected even in our very deep imaging."
    },
    "0806/0806.0032_arXiv.txt": {
        "abstract": "We investigate the estimation of orbital parameters by least-$\\chi^2$ Keplerian fits to radial velocity (RV) data using synthetic data sets. We find that while the fitted period is fairly accurate, the best-fit eccentricity and $M_p\\sin i$ are systematically biased upward from the true values for low signal-to-noise ratio $K/\\sigma\\lesssim 3$ and moderate number of observations $N_{\\rm obs}\\lesssim 60$, leading to a suppression of the number of nearly circular orbits. Assuming intrinsic distributions of orbital parameters, we generate a large number of mock RV data sets and study the selection effect on the eccentricity distribution. We find the overall detection efficiency only mildly decreases with eccentricity. This is because although high eccentricity orbits are more difficult to sample, they also have larger RV amplitudes for fixed planet mass and orbital semi-major axis. Thus the primary source of uncertainties in the eccentricity distribution comes from biases in Keplerian fits to detections with low-amplitude and/or small $N_{\\rm obs}$, rather than from selection effects. Our results suggest that the abundance of low-eccentricity exoplanets may be underestimated in the current sample and we urge caution in interpreting the eccentricity distributions of low-amplitude detections in future RV samples. ",
        "introduction": "With the rapid increase in the rate of exoplanet detections, it has become feasible to study their statistical properties (for a recent review, see Udry \\& Santos 2007). The distributions of their orbital parameters and correlations with host star properties are crucial for our understanding of planet formation. Up to Feb 2008, over 200 exoplanets have been announced, most of which were detected by the radial velocity (RV) technique. Among the parameters that can be derived from RV data, the orbital eccentricity has a somewhat unexpected distribution, with an extended tail of high eccentricities ($e\\gtrsim 0.1$). Although the exact form of this eccentricity distribution is still uncertain to some extent, especially at the lowest eccentricities (i.e., comparing the distribution in Butler et al. 2006 and in the most recent catalog), it is clear that this distribution is quite different from that of the Solar system. There have already been several theoretical attempts to explain such an eccentricity distribution (e.g., Tremaine \\& Zakamska 2004 and references therein; Juric \\& Tremaine 2007; Ford et al. 2007; Zhou et al. 2007).  However, the largest current exoplanet catalog (e.g., Bulter et al. 2006 and their updates, hereafter Bulter06) is by no means homogeneous. Survey strategies, selection biases in the RV technique and uncertainties in best-fit orbital solutions could all bias the intrinsic distributions of orbital parameters, especially when taken together. Eccentric orbits have larger amplitudes than circular orbits when holding other parameters fixed, while failure to (time) resolve the perihelion approach can lead to non-detection for an eccentric orbit system. These two effects are believed to roughly cancel but detailed simulation is needed (e.g., Endl et al. 2002; Cumming 2004). It is also well known that the errors in least-$\\chi^2$ Keplerian solutions become asymmetric for noisy data (e.g., Ford 2005; Butler et al. 2006), which is especially true for eccentricity. On the other hand, for circular orbits the eccentricity can only be scattered upward by error, so an obvious bias exists for circular orbits. As the RV surveys progress, more and more low amplitude (low signal-to-noise ratio) detections will be reported. Significant uncertainties remain in the distributions of orbital parameters for those exoplanets using best-fit orbital solutions. There are already some cases where the orbit eccentricity is poorly constrained by the RV data (e.g., Jones et al. 2006).  The purpose of the paper is to explore processes that might distort measurements of the shape of the intrinsic eccentricity distribution using synthetic data. First we investigate the reliability of single Keplerian fits to mock RV data sets as function of signal-to-noise ratio, and see if any bias arises from Keplerian fitting to noisy data (\\S\\ref{sec:kep_fit_eff}). Second, we construct statistical models of orbital parameter distributions and pass them through a simulated detection pipeline, in order to see if there is any serious selection effect on eccentricity (\\S\\ref{sec:stat}). We discuss our results and apply them to the current RV planet sample in \\S\\ref{sec:diss}, and we summarize our findings in \\S\\ref{sec:conc}.   \\begin{figure*} \\centering \\includegraphics[width=0.95\\textwidth]{f1.eps} \\caption{Effects of signal-to-noise ratio and number of observations on the best-fit orbital parameters. Filled circles are median values; the true orbital parameters are denoted as horizontal lines (see the text for details), where $\\omega$ and temporal offset $t_0$ are chosen at random. Error bars show standard deviations. From left to right, $N_{\\rm obs}=10,20,60$.} \\label{fig:SN_eff} \\end{figure*} ",
        "conclusions": "\\label{sec:conc}  We have performed simulations of planet detection and orbital parameter fits using mock radial velocity data sets. Two effects that may affect the intrinsic eccentricity distribution are considered: selection bias on eccentricity, and fitting bias in the best-fit orbital solution.  We find that selection bias on eccentricity is negligible, as long as the fitting routine is efficient in finding the global solution. In a realistic survey, this requires a thorough search in the parameter space, and/or with advanced algorithms optimized for the search. Our finding is not in conflict with previous studies (e.g., Endl et al. 2002; Cumming 2004), which claim that the detection efficiency decreases for $e\\gtrsim 0.6$. This is because in previous studies, the detection efficiency as a function of eccentricity is estimated at fixed $K$ while in our study $K$ is larger for more eccentric orbits when other parameters are fixed. On the other hand, we find that for detections with low signal-to-noise ratio and small number of observations, the best-fit eccentricity is biased upward in the median value, which then gives rise to a change in the eccentricity distribution from the intrinsic one.  Inspection of the current exoplanet sample shows only $\\sim 10\\%$ are likely to be significantly affected by the Keplerian fitting bias. However, future radial velocity surveys will contain an increasing number of low amplitude detections if the number of low mass, large semi-major axis exoplanets grows as rapidly as suggested by extrapolation of current results suggest (e.g., Udry \\& Santos 2007). When the sample of low amplitude detections (on average less massive planets and more circular orbits) is large enough for statistical study, the bias in the best-fit eccentricity described here must be taken into account. Also, in individual cases where an accurate estimate of eccentricity is required, i.e., in modeling the habitable zone or tidal heating issues, either high quality RV data or constraints from other observations are required to support reliable conclusions. In the mean time, there is also a need for a more statistically sophisticated understanding of the uncertainties of derived orbital solutions (e.g., Ford 2005; 2006).  Our results suggest that the intrinsic eccentricity distribution may be even more peaked at $e\\approx 0$ than the current observed distribution. Some planet-planet scattering models tend to produce a Rayleigh-distribution of eccentricities (e.g., Juric \\& Tremaine 2007; Ford \\& Rasio 2007) with reduced circular orbits, therefore certain eccentricity damping mechanism such as interactions with a protoplanetary disk may be required to reconcile these models with observations."
    },
    "0806/0806.0497_arXiv.txt": {
        "abstract": "We improve the estimate of the axion CDM energy density by considering the new values of current quark masses, the QCD phase transition effect and a possible anharmonic effect. ",
        "introduction": "The standard model is the greatest triumph of the 20th century particle physics. However, it suffers several naturalness problems. The so-called strong CP problem is one of them. Up to now, the most promising strong CP solution is axion \\cite{PQ77}.  Although its coupling is very small, it could have affected the evolution of the Universe.  On the other hand, the observations like the galaxy rotation curve and the cosmic microwave background radiation tell us that a significant portion of the energy density of the universe cannot be explained by the ordinary baryonic matter. The recent analysis of WMAP data\\cite{Komatsu:2008hk} combining the observation of the large scale structure prefers the standard vacuum energy and cold dark matter ($\\Lambda$CDM) model, with the cold dark matter (CDM) density $\\Omega_{\\rm DM} h^2 \\simeq 0.1143 \\pm 0.0034$.  During the inflation, a sufficiently light scalar field becomes spatially homogeneous and its value is selected stochastically\\cite{Linde82}. After the inflation, the Hubble parameter decreases as the universe expands.  If the Hubble parameter becomes smaller than some scale, usually the scalar mass scale, the scalar field starts to roll down. In the hydrodynamic description, this coherent oscillation works as  cold matter in the evolution of the universe. If the interaction is small enough to keep coherence until the matter-dominant era, it can be a part of the present matter energy density of the universe. The axion is that kind of particle and can be coherent until even now.  Just after the invisible axion was introduced \\cite{Kim79}, people recognized its cosmological implications that the classical axion field oscillation behaves like CDM \\cite{Preskill83}. On the axion CDM energy density, Turner presented a more accurate number including anharmonic effect and numerical instanton calculation in late 1980s \\cite{Turner86}. DeGrand {\\it et. al.} concerned about the possible effects of the chiral phase transition \\cite{DeGrand:1985uq}.   After Turner's pioneering work, many input parameters in axion cosmology, especially the QCD parameters, are refined. We present the improved number for the axion dark matter relic density with these new parameters.  We also discuss the possible effects of the chiral phase transition on the total entropy and the axion number density. This occurs only when supercooling arises. ",
        "conclusions": "In conclusion, we presented a new expression for the cosmic axion energy density  using the new values for the current quark masses, implementing the QCD chiral symmetry breaking and the adiabatic invariant $I$. We included the effect of anharmonic terms also. Using the adiabatic invariant quantity $I$, we presented the correction function $F$ to the anharmonic effect and a factor 1.85 which arose from the overshoot from $\\theta_1$.   Using the bag model for the phase transition in the early universe, the axion number is shown to be conserved in most circumstances. If the axion number change is considered during the phase transition, it is extremely fine-tuned in the $T_{\\rm i}-T_{\\rm f}$ space, in which case the only change is from a small entropy production during the phase transition. In this way, we obtain a factor of few larger bound on $F_a$ compared to the previous ones even for the smaller requirement on $\\Omega_a$. Of course, for a small $\\theta_1$ a strip with large $F_a$s is allowed, which has been used in the anthropic arguments \\cite{anthropic}.  \\vskip 0.5cm \\noindent {\\bf Acknowledgments\\ :} This work is supported in part by the Korea Research Foundation, Grant No. KRF-2005-084-C00001. KJB is supported also by the BK21 Program of the Ministry of Education and Science."
    },
    "0806/0806.3388_arXiv.txt": {
        "abstract": "The atmospheric properties above three sites (Dome C, Dome A and the South Pole) on the Internal Antarctic Plateau are investigated for astronomical applications using the monthly median of the analyses from ECMWF (the European Centre for Medium-Range Weather Forecasts). Radiosoundings extended on a yearly time scale at the South Pole and Dome C are used to quantify the reliability of the ECMWF analyses in the free atmosphere as well as in the boundary and surface layers, and to characterize the median wind speed in the first 100 m above the two sites. Thermodynamic instability properties in the free atmosphere above the three sites are quantified with monthly median values of the Richardson number. We find that the probability to trigger thermodynamic instabilities above 100 m is smaller on the Internal Antarctic Plateau than on mid-latitude sites. In spite of the generally more stable atmospheric conditions of the Antarctic sites compared to mid-latitude sites, Dome C shows worse thermodynamic instability conditions than those predicted above the South Pole and Dome A above 100 m. A rank of the Antarctic sites done with respect to the strength of the wind speed in the free atmosphere (ECMWF analyses) as well as the wind shear in the surface layer (radiosoundings) is presented. ",
        "introduction": "The summits of the Antarctic Plateau are of particular interest for ground-based astronomy because the optical turbulence appears to be confined to a narrow layer close to the ground, \\citep{ma2,la,ag}. Above this layer the atmosphere is exceptionally clear and the turbulence weak (0.27\\arcsec \\cite{la}, 0.36\\arcsec \\cite{ag} at Dome C). Measurements have shown that the height of the turbulent layer above the summits is much lower than above other sites on the Plateau, \\citep{la, maa, ma, ma2}. More precisely, the height of this layer is much larger above the South Pole (220 m \\citep{ma} or 270 m \\citep{tr}), which lies on a slope, than above Dome C (36 $\\pm$10 m \\cite{ag}). Above Dome A the turbulence is expected to be even weaker.  The surface winds of Antarctica are among the principal sources of turbulence near the ground. The dominant source of the surface winds is the sloping terrain and the radiative cooling of the surface \\citep{sch}. The radiative cooling produces a temperature inversion that together with the sloping terrain cause a horizontal temperature gradient. This triggers a surface wind alongside the slope of the terrain. Strong wind shears can therefore occur in the boundary between the surface winds and the winds in the free atmosphere, which in general are geostrophic. This is the main source of instability under conditions of extreme stability, as is the case of the Antarctic atmosphere in winter. Above the summits of the Internal Antarctic Plateau the surface winds should be much weaker than elsewhere on the Plateau due to a lack of the principal cause triggering them: a sloping terrain.  These elements justify the enthusiastic interest of astronomers for this site \\citep{sto,fos}% Studies of the atmospheric properties of these regions are fundamental for applications to ground-based astronomy. However we need a better quantification of these characteristics, using instruments and measurements as well as models and simulations, in order to fill the gaps of uncertainties or doubts that still remain \\citep{forot} , extending our attention to a comparative analysis of different sites of the Internal Antarctic Plateau.  It is important to produce statistical estimate of meteorological parameters at all heights from the ground to verify if atmospheric conditions are {\\it always} advantageous for astronomical applications.  Indeed, it has recently been shown \\citep{gm}, using European Centre for Medium-Range Weather Forecasts (ECMWF) analyses, that in winter the wind speed grows monotonically above $\\sim$10 km a.s.l. (above sea level), achieving median values of the order of $\\sim$30 m s$^{-1}$ at 25 km a.s.l.  At this height a variation of $~\\sim$20 m s$^{-1}$ has been estimated between summer and winter. Such a strong wind might trigger an important decrease of the wavefront coherence time, even in presence of a weak turbulence (see discussion in \\cite{gm}) and, as a consequence, the potentiality of these sites might vanish. It should be therefore interesting to better quantify the median wind speed profile on other sites (with astronomical interest) of the Internal Antarctic Plateau or to retrieve some general indication of the wind speed in the free atmosphere above the Internal Antarctic Plateau.  Besides that, the employment of ECMWF analyses for characterization of the surface layer requires a deeper analysis. \\cite{gm} concentrated their study at heights greater than 30 m thus excluding the surface contribution, assuming that the ECMWF analyses are not optimized for the atmospheric flow near the surface. More recently, studies \\citep{sa} % appeared claiming that the ECMWF analyses can be used to quantify and characterize the atmospheric properties with a good level of accuracy down to the surface level.  In spite of our conviction that this conclusion is the result of a partial analysis (only summer data) we admit that in \\cite{gm} % the authors {\\it assumed} (and they did not proved) the limits of the ECMWF for the surface layer.  We therefore think that is time to provide a dedicated analysis on this subject to know the limit within which we can achieve reliable estimates with General Circulation Models (GCM) i.e. with ECMWF analyses and to identify the domains in which one is forced to employ mesoscale models. The latter are in principle more suitable to better resolve phenomena happening at smaller spatial and temporal scales.  The usefulness of mesoscale models depends on the limitations imposed by the products of the General Circulation Models (that means the ECMWF analyses). It is obvious that, the usefulness of mesoscale models would not be justified for wind speed if ECWMF products could provide answers with a sufficient good accuracy.  Our group is involved in a long term study made with meso-scale models for the simulation of the optical turbulence on the whole troposphere and low stratosphere using the technique described in \\cite{b17} above the Internal Antarctic Plateau\\footnote{Some attempts have been done in the past \\citep{sg} even if with different scientific goals.}. This study is therefore propedeutic to researches done with such a typology of models.  It is therefore fundamental to provide a clear picture of the limitations of the ECMWF products and at the same time, to try to retrieve the maximum number of useful information we can get from such a kind of products.  In this paper we try a first attempt to quantify, above the three sites with some astronomic interests (the South Pole, Dome C and Dome A), the differences of some critical meteorologic parameters that are directly, or indirectly, related to the characteristics of atmospheric turbulence. We expressly select two sites (South Pole and Dome C) for which measurements are available and one site (Dome A) for which no measurements are available.  The reasons for this choice is explained later on (paper's scientific goals).  We use data from the MARS catalogue of the ECMWF and radiosoundings from the South Pole\\footnote{ftp://amrc.ssec.wisc.edu/pub/southpole/radiosonde} and Dome C\\footnote{http://www.climantartide.it}. Analyses data are extracted from the three grid points that are nearest to the sites of interest, i.e. Dome A, Dome C and the South Pole. The coordinates of the sites are given in Table \\ref{ou}. We extracted analyses data from MARS at 00:00 UTC for the whole year of 2005 to assure a complete statistical sample covering all seasons. A more detailed description of the analyses data set is given by \\cite{gm}.\\newline  \\begin{table} \\caption{The geographic coordinates of the sites and the closest grid points from which the ECMWF analyses are extracted.} \\begin{tabular}{llr@{\\degr}l} \\hline Site & Lat. & \\multicolumn{2}{l}{Long.} \\\\ \\hline Dome A$^*$      & 80\\degr22\\arcmin00\\arcsec S &  77 & 21\\arcmin11\\arcsec E \\\\ & 80\\degr30\\arcmin00\\arcsec S &  77 & 30\\arcmin00\\arcsec E \\\\ Dome C $^{**}$  & 75\\degr06\\arcmin04\\arcsec S & 123 & 20\\arcmin48\\arcsec E \\\\ & 75\\degr00\\arcmin00\\arcsec S & 123 & 30\\arcmin00\\arcsec E \\\\ South Pole      & 90\\degr00\\arcmin00\\arcsec S &   0 & 00\\arcmin00\\arcsec E \\\\ & 90\\degr00\\arcmin00\\arcsec S &   0 & 00\\arcmin00\\arcsec E \\\\ \\hline \\end{tabular} \\medskip {\\\\$^*$Measured with GPS by Dr X. Cui (private communication). \\\\$^{**}$Measured with GPS by Prof. J. Storey (private communication).} \\label{ou} \\end{table}   The scientific goals of this paper are: \\begin{enumerate} \\renewcommand{\\theenumi}{(\\arabic{enumi})} \\item To carry out a detailed comparison of radiosoundings/analyses of the wind speed and the potential temperature (the main critical parameters defining the stability of the atmosphere) near the surface (the first 150 m) for winter as well as summer above Dome C and the South Pole. This will permit us to quantify the uncertainty between measurements and ECMWF analyses in this vertical slab.  The idea is therefore to define the conditions in which the ECMWF can be used to characterize, with a good level of reliability and accuracy, some atmospheric parameters and to use this tool to characterize a site for which no measurements are available.  This is of course the interest for a model. Depending on the results of this analysis we will perform comparisons of meteorologic parameters above the South Pole, Dome C and Dome A using ECMWF analyses (Section 2). \\item Using radiosoundings we will estimate the statistic median values of the wind speed in the first tens of meters above the South Pole and Dome C extended from April to November. We will therefore be able to quantify which site shows the better characteristics for nightly astronomical applications (Section 3). \\item We extend the study developed by \\cite{gm} above Dome C for the ECMWF wind speed also to the South Pole and Dome A, both located on the Internal Antarctic Plateau but at different latitude and longitude. In this way, we intend to quantify which site is the best for astronomical applications. Results of this analysis are fundamental to confirm or see in the right perspective the potentialities of Dome C. \\item We extend the analysis of the Richardson number done by \\cite{gm} for Dome C for heights above 30 m to the three sites (South Pole, Dome C and Dome A) in order to quantify the regions and the periods that are less favorable for the triggering of optical turbulence and to identify the site with the best characteristics (Section 5). This result should represent the first estimate of potentialities of Dome A for astronomical applications and this should mean that we are able to provide some reliable results and conclusions even before some measurements are done on that site. This study has therefore a double interest. Firstly the intrinsic result itself. Secondly this analysis should open the path to a different approach for a fast and reliable classifications of potential astronomical sites. \\end{enumerate} ",
        "conclusions": "In this paper we provide a first comparison of the atmospheric characteristics above three different sites on the Internal Antarctic Plateau: Dome C, Dome A and the South Pole. More precisely we try to answer the specific questions defined in the introduction.  (1) The comparison of the ECMWF analyses with the radiosoundings shows that the analyses can accurately describe the atmosphere above the Internal Antarctic Plateau in the whole range from 10 m to 20 km above the surface. During no season does the median difference of the wind speed exceed 1 m s$^{-1}$ above the first 10 m. The median difference of the absolute temperature is within 2 K in the same vertical slab. In the surface layer the wind speed discrepancy between analyses and radiosoundings is slightly larger (2-3 m s$^{-1}$) while ECMWF analyses show a tendency to overestimate the absolute temperature measured by radiosoundings in the lowest level in winter ($\\Delta$T $\\sim$4-5 K). A statistic analysis reveals that most of the radiosoundings explode in winter at about 10-12 km. This does not allow us to estimate the reliability of the ECMWF analyses at these high altitudes.  (2) We proved that the ECMWF analyses do not produce accurate estimates in the surface layers confirming what was only assumed by \\cite{gm}.  This result represents an answer to our original question expressed in the Introduction. The ECMWF intrinsically have limitations for the characterization of the atmospheric flow in this vertical slab. This justify the employment of meso-scale models but, at the same time, also tells us that it will be fundamental to prove that meso-scale models can do better than General Circualtion Models (GCM) in this vertical slab.  (3) We could conclude that above all the three sites the potential temperature in the surface layer is extremely stable even if the ECMWF analyses generally underestimate its gradient when compared to measurements obtained by radiosoundings. Such an effect is particularly evident in winter. Dome A is by far the site with the steepest gradient of potential temperature and wind speed if compared to the South Pole and Dome C.  (4) We proved that the median wind speed in the first meters above the ground is weaker at Dome C than at the South Pole from April to November. However, the wind shear in the surface layer at Dome C is much larger than at the South Pole achieving at 10-20 m a wind speed of 8-9 m s$^{-1}$ in winter. Such a strong wind shear combined with a stable stratification of the air in this layer is most likely to be the cause of the intense optical turbulence that has been measured in the first tens of meters at Dome C \\citep{ag}. Such a strong wind speed at this height might be a source of vibrations produced by the impact of the atmospheric flow on telescope structures and should therefore be taken into account in the conception and design of astronomic facilities.  (5) Median monthly values of the inverse of the Richardson number (1/Ri) indicate that the probability to trigger instabilities is larger above a mid-latitude site (for which we have a reliable characterization of the optical turbulence) than above any of the three sites on the Internal Antarctic Plateau. Above all the three Antarctic sites 1/Ri is visibly smaller than that measured above Mt.Graham (selected as representative of mid-latitude sites). This is the first time that such a conclusion has been achieved and definitely prove that the method presented in \\cite{gm} % is reliable.  (6) Moreover, our analysis permitted us also a more sophisticated discrimination between the quality of the the 1/Ri above the three sites.  Dome A and the South Pole show, indeed, more stable conditions than Dome C above the first 100 m. This is probably due to the polar vortex which, producing an increase of the wind speed in the upper atmosphere, also increases the probability to trigger thermodynamic instabilities.  (7) We showed that it is risky to retrieve estimates of the Richardson number in the surface layer (Fig.\\ref{rs1}) because we did not find an equivalent smoothing effect of the ECMWF analyses for the gradient of the potential temperature nor for the gradient of the wind speed above different sites.  (8) In the free atmosphere, above the first 10 km, the polar vortex induces a monotonic increase of the wind speed in winter that is proportional to the distance of the site to the polar high. Dome C therefore shows the largest wind speed above 10 km in winter. At Dome C the wind speed at 15 km can easily be almost twice that of Dome A and even thrice the wind speed at the South Pole in winter. The wind speed above the South Pole is the weakest among the three sites on the whole 20 km in all seasons. This conclusion put therefore fundamental warnings for astronomical applications.  (9) This study allowed us to draw a first comprehensive picture of the atmospheric properties above the Internal Antarctic Plateau.  In spite of the presence of generally good conditions for astronomical applications, Dome C does not appear to be the best site with respect to the wind speed, in the free atmosphere as well as in the surface layer. Both the South Pole and Dome A show a weaker wind speed in the free atmosphere. Estimates related to the surface layer need to be taken with precaution. ECMWF analyses cannot be used to draw definitive conclusions on comparisons of the three sites in this vertical slab due to their limited reliability in this thin atmospheric slab (see Section 5) and radiosoundings are available only for Dome C and the South Pole. Above Dome A the gradient of the potential temperature is particularly large in the very near surface layer indicating conditions of extreme thermal stability that might be associated to a strong value of the optical turbulence in this vertical range when a thermodynamic instability occurs (possibly even larger than above Dome C).  However our study showed that, to predict the thickness of such a layer we should need measurements or simulations with atmospheric mesoscale model with a higher spatial resolution near the ground that is able to better resolve the evolution of the atmospheric flow. This is a part of our forthcoming activities.  In conclusion, at present, the real solid and unique argument that makes Dome C preferable to the South Pole for astronomical applications is the extreme thinness of the optical turbulence surface layer. We expect at Dome A comparable or even larger optical turbulence values with respect to Dome C in the surface layer. We cannot conclude if the surface layer at Dome A is thinner than that observed above Dome C. However our study clearly indicates that Dome C is not the best site on the Internal Antarctic Plateau with respect to the wind speed in the free atmosphere as well as in the surface layer nor is it the site with the most stable conditions in the free atmosphere. Both Dome A and the South Pole show more stable conditions in the free atmosphere."
    },
    "0806/0806.4082_arXiv.txt": {
        "abstract": "The cosmic acceleration is one of the most significant cosmological discoveries over the last century. The two categories of explanation are exotic component (dark energy) and modified gravity. We constrain the two types of model by a joint analysis with perturbation growth and direct $H(z)$ data. Though the minimal $\\chi^2$ of the $\\Lambda$CDM is almost the same as that of DGP, in the sense of consistency we find that the dark energy ($\\Lambda$CDM) model is more favored through a detailed comparison with the corresponding parameters fitted by expansion data. ",
        "introduction": "The acceleration of the universe is one of the most significant cosmological discoveries over the last century \\cite{acce}. Various explanations of this acceleration have been proposed, see \\cite{review} for recent reviews with fairly complete lists of references of different models. However, although fundamental for our understanding of the universe, its nature remains as a completely open question nowadays.  There are two main categories of proposals. One is that the acceleration is driven by some exotic matter with negative pressure, called dark energy. The other suggests that general relativity fails in the present Hubble scale. The extra geometric effect is responsible for the acceleration. Surely, there are some proposals which mix the two categories. Mathematically, in the dark energy model we present corrections to the right hand side of Einstein equation (matter part), while the correction terms appear in the left hand side of Einstein equation (geometric part).  $\\Lambda$CDM model is the most popular and far simple dark energy model, in which vacuum energy with the equation of state (EOS) $w=-1$ accelerates the universe. From theoretical considerations and by observational implications, people put forward several other candidates for dark energy, such as quintessence ($-1<w<-1/3$), phantom ($w<-1$), etc. Also there are many possible corrections to the geometric part of the theory. One of the leading modified gravity model is Dvali-Gabadadze-Porrati (DGP) model \\cite{dgpmodel}, for a review, see \\cite{dgpre}. In the DGP model, the bulk is a flat Minkowski spacetime, but an induced gravity term appears on a tensionless brane. In this model, gravity appears 4-dimensional at short distances. But, at a distance larger compared to some freely adjustable crossover scale $r_0$ it is altered  through the slow evaporation of the graviton off our 4-dimensional brane world universe into an unseen, yet large, fifth dimension.  We should find the correct, at least exclude the incorrect models in the model sea. The first step is to discriminate between dark energy and modified gravity, whose nature are completely different. To construct a model simulating the accelerated expansion is not very difficult. That is the reason why we have so many different models. Recently, some suggestions are presented that growth function $\\delta(z)\\equiv\\delta\\rho_m/\\rho_m$ of the linear matter density contrast as a function of redshift $z$ can be a probe to discriminate between dark energy and modified gravity \\cite{dgpgrow, growth} models. The growth function can break the degenerations between dark energy and modified gravity models which share the same expansion history.  There is an  approximate relation between the growth function and the partition of dust matter in standard general relativity \\cite{peeb}, \\be \\label{f} f\\equiv\\frac{d\\ln\\delta}{d\\ln a}=\\Omega_m^\\gamma, \\en where $\\Omega_m$ is the density partition of dust matter, $a$ denotes the scale factor, and $\\g$ is the growth index. This relation is a perfect approximation at high redshift region. Also, it  can be used in low redshift region, see for example \\cite{stein}. It is shown that the relation ($\\ref{f}$) is also valid in the case of modified gravity theory \\cite{dgpgrow}. The theoretical value of $\\g$ for $\\Lambda$CDM model is $6/11$ \\cite{growth, stein}, while, for spatially flat DGP model is $11/16$ \\cite{dgpgrow}. The observation data of perturbation{ growth} are listed in Table 1. {\\large \\begin{table} \\begin{center} \\begin{tabular}{ccccc} \\hline\\hline $z$ &  ~~~~~$f_{obs}$ & Reference\\\\ \\hline $0.15$ &~~~~~ $0.51\\pm 0.11$ & \\cite{hawk}\\\\ $0.35$ & ~~~~~$0.70\\pm 0.18$ & \\cite{teg}\\\\ $0.55$ &~~~~~ $0.75\\pm 0.18$ & \\cite{ross}\\\\ $0.77$ &~~~~~ $0.91\\pm 0.36$ & \\cite{guzz}\\\\ $1.4$ &~~~~~ $0.90\\pm 0.24$ & \\cite{angl}\\\\ $3.0$ &~~~~~ $1.46\\pm 0.29$ & \\cite{dona}\\\\ \\hline\\hline \\end{tabular} \\end{center} \\caption{\\label{tab1} Observed perturbation growth as a function of redshift $z$, see also \\cite{weihao} } \\end{table}}  It is shown that $\\Lambda$CDM model is consistent with the current growth data \\cite{favour}. The data seem to weaken a spatially flat DGP model, whose $\\g=11/16$. However, it is found that the growth index $\\g$ is $4/7$ in a non spatially flat DGP model \\cite{weihao}, which is very closed to the index of $\\Lambda$CDM. Thus, the DGP model may be still consistent with current growth data.  In this article we take a different strategy. The previous works were concentrated on the limit of the growth index  and made some approximations on it (often the high $z$ limit was assumed and an approximation was made at linear order), in which only approximate asymptotic value of $\\g$ can be obtained. In fact, the perturbation growth $f$ is a variable with respect to $z$, as displayed in Table 1. By using these growth data we constrain the parameters in $\\Lambda$CDM model and DGP model, respectively.  The other one which is very useful but not widely used in model constraint data is the set of direct $H(z)$. $H(z)$ is derived by a newly developed scheme to obtain the Hubble parameter directly at different redshift \\cite{h(z)}, which is based on a method to estimate the differential ages of the oldest galaxies \\cite{age}. By using the previously released data \\cite{aid}, Simon {\\it et al.} obtained a sample of direct $H(z)$ data in the interval $z\\in (0,1.8)$~\\cite{simon}, just as the same interval of the data of luminosity distances from supernovae. For the present sample of growth data derived with the assumption of the expansion behaviors of the universe, we will present joint fittings to obtain the constraints on the $\\Lambda$CDM and DGP, respectively. Then, through comparing with allowed regions by expansion constraint using different observations, including supernovae (SN), cosmic microwave background (CMB), baryon acoustic oscillations (BAO) etc., we examine which model is more self-consistent.   This article is organized as follow:  In the next section we construct the evolution equation for $f$ in a very general frame. In section III, by using the growth data and $H(z)$ data we present the parameter constraints of $\\Lambda$CDM and DGP, respectively. Our conclusion and some discussions appear in the last section. ",
        "conclusions": "Perturbation growth is a newly developed method to differentiate between dark energy and modified gravity. In the previous works people concentrate on the approximate analytical value of the perturbation growth index of a model, and then compare with the observations. But, the index is not a constant in the history of the universe. We fit dark energy and modified gravity models by using the exact evolution equation of perturbation growth.  The sample of presently available growth data is quite small and the error bars are rather big. Furthermore, we always assumed $\\Lambda$CDM model for deriving the growth data. Thus it seems proper to fit a model by jointing the growth data and the expansion data.  The direct $H(z)$ data are new type of data, which can be used to explore the fine structures of the Hubble expansion history. We put forward a joint fitting by the growth data and $H(z)$ data. The results are summarized as follows: For $\\Lambda$CDM, $\\Omega_{m0}=0.275^{+0.0544}_{-0.0549}$, $\\Omega_{k0}=0.065^{+0.159}_{-0.149}$; For DGP, $\\Omega_{m0}=0.350^{+0.132}_{-0.0974}$, $\\Omega_{r_c}=0.200^{+0.0631 }_{-0.0483}$.  The minimal $\\chi^2$ are 12.26 and 12.11 for $\\Lambda$CDM and DGP, respectively. The permitted parameters of $\\Lambda$CDM by growth+H(z) data show an excellent consistency with the previous results inferred from expansion data. However, for DGP model the discrepancies of the results of growth+H(z) data and expansion data are at least 2$\\sigma$ level. Hence 9 in the sense of consistency, $\\Lambda$CDM is more favored than DGP."
    },
    "0806/0806.4611_arXiv.txt": {
        "abstract": "The NASA Star and Exoplanet Database (NStED) is a general purpose stellar archive with the aim of providing support for NASA's planet finding and characterization goals, stellar astrophysics, and the planning of NASA and other space missions. There are two principal components of NStED: a database of (currently) 140,000 nearby stars and exoplanet-hosting stars, and an archive dedicated to high precision photometric surveys for transiting exoplanets. We present a summary of the NStED stellar database, functionality, tools, and user interface. NStED currently serves the following kinds of data for 140,000 stars (where available): coordinates, multiplicity, proper motion, parallax, spectral type, multiband photometry, radial velocity, metallicity, chromospheric and coronal activity index, and rotation velocity/period.  Furthermore, the following derived quantities are given wherever possible: distance, effective temperature, mass, radius, luminosity, space motions, and physical/angular dimensions of habitable zone. Queries to NStED can be made using constraints on any combination of the above parameters. In addition, NStED provides tools to derive specific inferred quantities for the stars in the database, cross-referenced with available extra-solar planetary data for those host stars.  NStED can be accessed at \\tt{http://nsted.ipac.caltech.edu}. ",
        "introduction": "The NASA Star and Exoplanet Database (NStED) is dedicated to collecting and serving vital published data involved in the search for and study of extrasolar planets and their host stars. The stellar and exoplanet services provide access to stellar parameters of potential exoplanet bearing stars along with exoplanet parameters. The stellar services provided by NStED include the following: \\begin{itemize} \\item Data related to relatively bright nearby stars. \\item The ability to query for individual stars or search by stellar/planetary parameters. \\item Published images, spectra, and time series data related to the stars in the database. \\end{itemize} Complementary to this are the exoplanet services, which includes the following: \\begin{itemize} \\item General data and pubished parameters for the known exoplanets and host stars. \\item Photometric and radial velocity data related to the known exoplanets. \\item A dedicated interface related to exoplanet transit surveys (see companion paper by von Braun et al. (2008) in these proceedings). \\end{itemize} Figure 1 shows an example plot produced using the data served by NStED.  \\begin{figure} \\begin{center} \\includegraphics[width=\\textwidth]{fig1.eps} \\caption{Plot of predicted astrometric wobble for an Earth-sized planet in the habitable zone vs. the apparent V magnitude of the stars, generated using data served by NStED. The stars are sorted by activity level estimates from the R'(HK) index, S index, and X-ray luminosity.} \\end{center} \\end{figure} ",
        "conclusions": "The NStED Stellar and Exoplanet Hosting Star Service provides access to data relevant to exoplanet host stars and bright stars from major catalogues. The search query tools and cross-referencing capabilities make this a powerful engine through which to aid in exoplanet survey programs. NStED provides not only tabular data, but also a wide variety of associated data including images, spectra, and time series (radial velocity and photometric) obervations. NStED is continually updated to reflect the latest results in the literature and to provide published data access to the broader astronomical community. NStED is accessible via {\\tt{http://nsted.ipac.caltech.edu}}.  \\begin{figure} \\begin{center} \\includegraphics[width=\\textwidth]{fig4.eps} \\caption{Aitoff projection of the contents of NStED. Small dots: dwarf stars (for clarity, the giant stars are not plotted); large dots: exoplanet hosting stars; large plus signs: stars with radial velocity curves or photometric lightcurves; open squares/diamonds: stars with images/spectra. For an explanation of the remaining features, see companion paper on the NStED Exoplanet Transit Survey Service by von Braun et al. (2008) in these proceedings.} \\end{center} \\end{figure}"
    },
    "0806/0806.2798_arXiv.txt": {
        "abstract": "We employ chaotic ($\\phi^2$ and $\\phi^4$) inflation to illustrate the important role radiative corrections can play during the inflationary phase. Yukawa interactions of $\\phi$, in particular, lead to corrections of the form $-\\kappa\\phi^4 \\ln (\\phi/\\mu)$, where $\\kappa >0$ and $\\mu$ is a renormalization scale. For instance, $\\phi^4$ chaotic inflation with radiative corrections looks compatible with the most recent WMAP (5 year) analysis, in sharp contrast to the tree level case. We obtain the 95\\% confidence limits $2.4\\times10^{-14}\\lesssim\\kappa\\lesssim5.7\\times10^{-14}$, $0.931\\lesssim n_s\\lesssim0.958$ and $0.038\\lesssim r\\lesssim0.205$, where $n_s$ and $r$ respectively denote the scalar spectral index and scalar to tensor ratio. The limits for $\\phi^2$ inflation are $\\kappa\\lesssim7.7\\times10^{-15}$, $0.929\\lesssim n_s\\lesssim0.966$ and $0.023\\lesssim r\\lesssim0.135$. The next round of precision experiments should provide a more stringent test of realistic chaotic $\\phi^2$ and $\\phi^4$ inflation. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0806/0806.0754_arXiv.txt": {
        "abstract": "{We present the results of the first long-term (2.2 years) spectroscopic monitoring of a gravitationally lensed quasar, namely the Einstein Cross \\obj.  We spatially deconvolve deep VLT/FORS1 spectra to accurately separate the spectrum of the  lensing galaxy from the spectra of the quasar images. Accurate cross-calibration of the observations at 31 epochs from October 2004 to December 2006 is carried out using foreground stars observed simultaneously with the quasar.  The quasar spectra are further decomposed into a continuum component and several broad emission lines.  We find prominent microlensing events in the quasar images A and B, while images C and D are almost quiescent on a timescale of a few months. The strongest variations are observed in the continuum, and their amplitude is larger in the blue than in the red, consistent with microlensing of an accretion disk. Variations in the intensity and profile of the broad emission lines are also reported, most prominently in the wings of the \\ion{C}{III]} and in the center of the \\ion{C}{IV} emission lines. During a strong microlensing episode observed in quasar image A, the broad component of the \\ion{C}{III]} is more magnified than the narrow component. In addition, the emission lines with higher ionization potentials are more magnified than the lines with lower ionization potentials, consistent with the stratification of the  broad line region (BLR) infered from reverberation mapping observations. }  \\FullConference{The Manchester Microlensing Conference: The 12th International Conference and ANGLES Microlensing Workshop\\\\ January 21-25 2008\\\\ Manchester, UK}  \\begin{document} ",
        "introduction": "Most of the quasar microlensing studies so far are based exclusively on broad-band photometric monitoring (e.g. Colley \\& Schild 2003; Udalski et al. 2006). These observations are dominated by variations of the continuum, making it difficult to disentangle between variations of the continuum and of the broad emission lines (BELs). Both regions are affected by microlensing, but in different ways depending on their size.    Several theoretical studies show how multiwavelength lightcurves can constrain the energy profile of the quasar accretion disk and also the absolute sizes of the line-emitting regions (e.g., Schneider \\& Wambsganss 1990; Agol \\& Krolik 1999; Abajas et al. 1999;  Kochanek 2004). In order to investigate the inner structure of quasars, we started the first long-term spectrophotometric monitoring of a lensed quasar. Our target is  the Einstein Cross \\obj, well known for its microlensing induced variability. The spectral variations  of the  four   quasar images  are followed with  the Very Large   Telescope (VLT) from October 2004 to December 2006. In this contribution we summarize the results presented in Eigenbrod et al. (2008). The full analysis of our monitoring data requires detailed  microlensing  simulations in order to constrain the quasar energy profile and  BLR size. These simulations will be the subject of future papers. ",
        "conclusions": "We present the first long-term spectrophotometric monitoring of a gravitationally lensed quasar; the Einstein Cross \\obj. The mean temporal sampling is of one observation every second week. The observations are carried out with the VLT in a novel way, using the spectra of PSF stars, both to deblend the quasar images from the lensing galaxy and to carry out a very accurate flux calibration.   We find that all images of \\obj\\  are affected by microlensing. Image~A  shows an  important brightening episode  at  the end of our monitoring campaign, and   image~B  at the beginning. The continuum of these two images becomes bluer as they get brighter, as expected from microlensing magnification of an accretion disk. We also report microlensing-induced variations of the BELs, both in their integrated line intensities and in their profiles. In   quasar   image   A, the  broad   component   of  the \\ion{C}{III]} line is more magnified than the narrow component. This might indicate that the core of this line is emitted in the NLR. Intensity variations in the BELs are detected mainly  in images A and B. Our  measurements  suggest    that  higher ionization   BELs   like \\ion{C}{IV},  \\ion{C}{III]},  are  more  magnified than lower ionization   lines like \\ion{Mg}{II}. This  is  compatible with reverberation mapping studies and  a stratified structure of  the BLR. There    is  marginal evidence that   regions   of different sizes are responsible for the \\ion{Fe}{II+III} emission.  The very different behaviours of the BELs and the continuum with  res\\-pect to microlensing offer considerable hope to  reconstruct the  two types of regions independently, using inverse ray-shooting simulations."
    },
    "0806/0806.4388_arXiv.txt": {
        "abstract": "We use photometric redshifts to analyse the effect of local environment on galaxy colours at redshifts $ z \\la 0.63$ in the SDSS data release 6. We construct mock SDSS-DR6 catalogues using semi-analytic galaxies to study possible systematic effects on the characterisation of environment and colour statistics due to the uncertainty in the determination of redshifts. We use the projected galaxy density derived from the distance to the nearest neighbours with a suitable radial velocity threshold to take into account the uncertainties in the photometric redshift estimates. Our findings indicate that the use of photometric redshifts greatly improves estimates of projected local galaxy density when galaxy spectra are not available.  We find a tight relationship between spectroscopic and photometric derived densities, both in the SDSS-DR6 data (up to $z=0.3$) and mock catalogues (up to $z=0.63$).  At $z=0$, faint galaxies show a clear increase of the red galaxy fraction as the local density increases.  Bright galaxies, on the other hand, show a constant red galaxy fraction.  We are able to track the evolution of this fraction to $z=0.55$ for galaxies brighter than $M_r=-21.5$ and find that the fraction of blue galaxies with respect to the total population progressively becomes higher as the redshift increases, at a rate of $15\\%/$Gyr.  Also, at any given redshift, bright galaxies show a larger red population, indicating that the star-formation activity shifts towards smaller objects as the redshift decreases. ",
        "introduction": "The study of galaxies in the field and in clusters has revealed the existence of significant correlations between several galaxy properties and their environment. In a  pioneering work by \\citet{dress80}, it was shown for the first time that Galaxy morphologies depend on the local galaxy density. The decrease in the SFR of galaxies in dense environments is a universal phenomenon over a wide range of densities.  For instance, \\citet{gomez03} found that the star formation rate (SFR) of galaxies is strongly correlated with the local galaxy density; \\citet{balogh04} characterise the environment using projected and three-dimensional densities concluding that the present-day correlation between star formation rate and environment is a result of short-timescale mechanisms that take place preferentially at high redshift, such as starbursts induced by galaxy-galaxy interactions (see also \\citealt{baldry04}). However, these studies concern the nearby universe where spectroscopic redshifts allow for estimates of the local density of galaxies. At higher redshifts, spectroscopy is not available for large galaxy samples, and in consequence, the relation between galaxy properties and environment is more poorly understood. An alternative solution to this problem is the use of multi-band photometry to constrain galaxy redshifts. These techniques have been extensively studied (e.g., \\citealt{koo85,con95,gh96,ben00,bolz00,csabai03,annz}), and have proved to be efficient in estimating redshifts for large numbers of galaxies, which opens the possibility of obtaining their intrinsic properties for statistical studies.  One exception to this necessity is that of the VIMOS VLT Deep Survey, for which Cucciati et al. (2006) make use of the available spectroscopic redshifts to estimate local densities; however, their sample is small containing only 6582 galaxies out to $z=1.5$.  Their results show that massive galaxies shift towards redder values for lower redshifts, where only faint galaxies show a blue galaxy population. On the other hand, the use of photometric redshifts to study the evolution of galaxy colours at redshifts higher than $z=0.3$, has only been attempted a few times, in part due to the difficulty in ensuring an accurate statistical measurement of local density. For instance, \\citet{roser} study the colour distribution of galaxies in the CFHTLS-Deep Field Survey \\footnote{http:$//$www.ast.obs-mip.fr/users/roser/CFHTLS\\_T0003} for redshifts in the range $0<z<1.3$, using a well calibrated photometric redshift estimator, New-HyperZ\\footnote{ http:$//$www.ast.obs-mip.fr/users/roser/hyperz }, to find results consistent with those from the VIMOS VLT Deep Survey, indicating the validity of the photometric redshift approach. De Lucia et al.  (2007), find that the population of cluster galaxies in the Las Campanas Distant Cluster Survey (Gonzalez et al., 2001) shows a clear deficit of blue galaxies in low redshift galaxies with respect to clusters at $z\\sim0.8$. The variation of galaxy colours as a function of redshift provides important information on the stellar formation history of galaxies; De Lucia et al. fit their results with models where the star formation is suppressed early on in hostile cluster environments.  In this paper we analyse galaxy properties and their relation to the environment by means of photometric data taken from the Sloan Digital Sky Survey, Data Release 6 (SDSS-DR6, \\citealt{sdss}).  We use galaxy photometric redshifts, $\\zphot$, from the SDSS database, to compute local densities, intrinsic luminosities and colours.  The main advantage of our study is the comparatively large number of galaxies in the SDSS-DR6, at redshifts $0.1<z<0.63$, a range where deep spectroscopic surveys only offer a limited number of objects, and wide spectroscopic surveys are badly affected by selection biases.  \\begin{figure*} \\begin{picture}(450,240) \\put(0,0){\\psfig{file=f1.eps,width=8.cm}} \\put(240,0){\\psfig{file=f2.eps,width=8.cm}} \\end{picture} \\caption{Left panel: $\\zspec$ vs $\\zphot$ relation for the SDSS-DR6 (dashed line);  errorbars correspond to the $10$ and $90$ percentiles. Values corresponding to individual galaxies are shown as grey dots (shown for the $10\\%$ of the sample). The solid line shows the one-to-one relation. The error bars correspond to 2$\\sigma$ deviations.  Right panels: systematic (top right) and stochastic (bottom right) errors in the photometric redshift estimates, for different apparent magnitude ranges (see the figure key).} \\label{f1} \\end{figure*}      This paper is organised as follows. In Section 2 we describe the data used in our analysis, both observational and from models of galaxy formation, Section 3 discusses the photometric redshifts available in the SDSS-DR6 and analyses their advantages and limitations.  Section 4 describes the method used to calculate projected galaxy densities, Section 5 studies the evolution of galaxy colours as a function of their environment with redshift, and finally, Section 6 summarises the results obtained in this work.  Throughout this paper, we adopt the cosmological model characterised by the parameters $\\Omega_m=0.25$, $\\Omega_{\\Lambda}=0.75$ and $H_0=70~h~ {\\rm km~s^{-1}~Mpc^{-1}}$. ",
        "conclusions": "We have studied the evolution of the distribution of galaxy colours in the SDSS-DR6  at intermediate redshifts $0.04<z<0.63$, and its dependence on projected local density, as inferred from nearby neighbours in projection, using photometric redshift information.  A first step consisted in constructing mock catalogues with photometric redshifts and uncertainties derived from the differences between $\\zphot$ and $\\zspec$ in the SDSS-DR6 MGS, in two versions, one with an evolving galaxy population and another one with no evolution.  These mocks are used for two main purposes, i) to determine whether the statistical analyses are affected by systematic errors, and to assess stochastic errors to improve the determination of projected densities for the analysis on real data, and ii) to compare the evolution predicted by the galaxy formation model to the observed evolution.  Using the SDSS-DR6 MGS, we are able to tune our density measurement technique as well as to determine that the estimates of projected densities using photometric redshift information are in agreement with the results using spectroscopic data, out to $z\\simeq0.3$.  We confirm this result in the mock catalogues (both with and without evolution), and extend this test out to much higher redshifts, $z=0.63$, the maximum redshift out to which we can analyse the SDSS-DR6 photometric catalogue avoiding possible confusion problems between galaxies and the PSF. We also test whether the colour distributions of galaxies selected in bins of projected density are comparable when using photometric and spectroscopic distances to select close neighbours for the determination of projected densities.  We find that the distributions are compatible both in the real and mock data.  In this case, the variations in the colour distributions seen in the model are narrower than in the real data, but follow qualitatively similar trends with density and luminosity.  Low luminosity galaxies show mostly a blue population, with signatures of a red population appearing only at the highest density bin; high luminosity galaxies show mostly a red population, with some blue galaxies only at low densities.  Intermediate luminosities show a nice transition from a bimodal colour distribution at low densities to an almost unimodal red distribution at high densities.  These results are in excellent agreement with previous studies (e.g. Balogh et al., 2004).  The evolution of the colour distributions with redshift shows that low luminosity galaxies tend to loose the red population almost completely, while the blue population shifts to even bluer colours; high luminosity galaxies, on the other hand, show a constant shift to bluer colours; this is more easily seen in our study of red galaxy fractions for different redshift intervals.  One important characteristic of this evolution is that  the dependence of the red fraction on local density in high luminosity galaxies, $M_r-5\\log(h_{70})<-20.5$, retains its overall shape and only changes by a constant shift, which for $-23<M_r-5\\log(h_{70})<-21.5$, corresponds to a $\\simeq 60\\%$ increase from median $z=0.545$ to $z=0.05$. Intermediate luminosities, $-21<M_r-5\\log(h_{70})<-19.5$, show a $50\\%$ increase from a median $z=0.19$ to median $z=0.05$ in the red fraction in relatively low projected densities; galaxies in high density environments show a milder change in the red fraction in this redshift range. This evolution can be more easily seen as a function of redshift for fixed values of local density, as is shown in Figure \\ref{f6}, where solid lines and filled symbols correspond to high densities, and dashed lines and open symbols to very low local densities. As can be seen, the slope in the red fraction evolution is approximately constant implying that in the local universe, the blue fraction decreases at a rate of $15\\%/$Gyr. There is an indication that faint and bright galaxies follow a similar trend, but due to the magnitude limit of the SDSS-DR6, the data on the evolution of the fainter subsamples reaches lower redshifts, increasing the uncertainty in the estimate of the slope.  These results are consistent with previous works (Cucciati et al., 2006, De Lucia et al., 2007, \\citealt{martinez}, Ienna \\& Pell\\'o, 2006), where the population of high luminosity galaxies becomes red at higher redshifts than their low luminosity counterparts. This phenomenon has been termed ``downsizing\" in the literature (e.g. \\citealt{marconi04}, \\citealt{shankar04}), in relation to the star formation activity shifting from larger objects towards smaller galaxies, in apparent contradiction to the hierarchical model of galaxy formation.  Several recent works have shown that this behaviour is in fact also natural in a hierarchical Universe (Lagos, Cora \\& Padilla, 2008, Bower et al., 2006, Croton et al., 2006); however, it still remains a challenge for galaxy formation models to reproduce the actual distribution and evolution of galaxy colours which at present is achieved only qualitatively at best, for which this work provides the most accurate measurements carried out to date."
    },
    "0806/0806.4177_arXiv.txt": {
        "abstract": "SNe Ia are good distance indicators because the shape of their light curves, which can be measured independently of distance, varies smoothly with luminosity. This suggests that SNe Ia are a single family of events. Similar correlations are observed between luminosity and spectral properties. In particular, the ratio of the strengths of the \\SiII\\ $\\lambda5972$ and $\\lambda6355$ lines, known as $\\mathcal{R}$(\\SiII), was suggested as a potential luminosity indicator. Here, the physical reasons for the observed correlation are investigated. A Monte-Carlo code is used to construct a sequence of synthetic spectra resembling those of SNe with different luminosities near $B$ maximum. The influence of abundances and of ionisation and excitation conditions on the synthetic spectral features is investigated. The ratio $\\mathcal{R}$(\\SiII) depends essentially on the strength of \\SiII\\ $\\lambda5972$, because \\SiII\\ $\\lambda6355$ is saturated. In less luminous objects, \\SiII\\ $\\lambda5972$ is stronger because of a rapidly increasing \\SiII/\\SiIII\\ ratio. Thus, the correlation between $\\mathcal{R}$(\\SiII) and luminosity is the effect of ionisation balance. The \\SiII\\ $\\lambda5972$ line itself may be the best spectroscopic luminosity indicator for SNe Ia, but all indicators discussed show scatter which may be related to abundance distributions. ",
        "introduction": "Type Ia supernovae (SNe Ia) form a relatively homogeneous supernova subclass. They supposedly originate from accreting C-O white dwarfs that undergo explosive burning as they approach the Chandrasekhar mass limit \\citep{hil00}. Their enormous light output is powered by radioactivity of synthesised \\Nifs\\ and of its daughter nuclide \\Cofs\\ \\citep{arn82}.  Differences in the synthesised \\Nifs\\ mass result in a spread in luminosity of $\\sim\\!\\!3$ mag in the $B$ band \\citep{maz01b,maz07}. \\citet{phi93} found that the decline in $B$-band magnitude within 15 days after $B$ maximum, $\\Delta m_{15}(B)$, correlates inversely with the $B$-band maximum luminosity $M_{B,\\textrm{max}}$ and thus can be used as a luminosity indicator. Thus, SNe Ia can be used as distance indicators in cosmological studies (``standard candles'').  Optical spectra reflect in detail the variations among SNe Ia. In the photospheric phase, which lasts until a few weeks after maximum light, P-Cygni features (blueshifted absorption plus less prominent emission centred at the rest wavelength) are superimposed on a ``pseudo-continuum'' \\citep{pau96}. The line positions and strengths depend on the abundances and on the excitation and ionisation conditions within the debris.  \\citet{nug95} showed that the ratio of the depths of two \\SiII\\ lines ($\\lambda 5972$ and $\\lambda 6355$) at $B$-band maximum, $\\mathcal{R}$(\\SiII), is larger for dimmer SNe, and suggested that this is due to temperature differences. Accordingly, $\\mathcal{R}$(\\SiII) could be used as a distance-independent spectroscopic luminosity indicator \\citep[e.g.][]{bon06}. Other spectral features have subsequently been suggested as possible luminosity indicators (\\citealt{hac06}, hereafter \\citetalias{hac06}; \\citealt{bon06}; \\citealt{fol08}). Still, the detailed physics behind the $\\cal R$(\\SiII)-$L$ correlation have not been explained.  This is the purpose of this work, where we explain the $\\cal R$(\\SiII)-$L$ correlation investigating the details of the connection between the physics of the envelope and line strengths. We present a sequence of spectral models calculated with a Monte-Carlo (MC) radiative transfer code. The synthetic spectra resemble real SNe with different luminosities shortly after $B$ max. In the model atmospheres, we analyse the conditions which influence the line strengths entering $\\cal R$(\\SiII).  In the following, first we describe the concept of the synthetic spectral sequence (Sec. 2). In Sec. 3, we show that our models reproduce the $\\cal R$(\\SiII)-$L$ correlation. In Sec. 4 we explain the correlation analysing the details of the  formation of the lines contributing. After discussing spectral luminosity indicators and other applications of line-strength measurements (Sec. 5), we give conclusions (Sec. 6). ",
        "conclusions": "Based on a sequence of 25 synthetic spectra constructed to fit SNe Ia with different luminosities, we have been able to explain the $\\mathcal{R}($\\SiII$)-L$ relation.  The increasing strength of \\SiII\\ $\\lambda 5972$ is responsible for the increase of $\\mathcal{R}($\\SiII$)$ from luminous to dim objects. This is in turn due to ionisation: in dimmer SNe, \\SiII\\ $\\lambda 5972$ is stronger because the \\SiII\\ fraction increases more strongly than than the level occupation numbers decline. In contrast, \\SiII\\ $\\lambda 6355$ is saturated because of its large optical depths. Thus, the line has a rather constant strength and only a minor influence on $\\mathcal{R}($\\SiII$)$.  An investigation as performed here shows how observed line strengths can be related to SN physics. \\SiII\\ $\\lambda 5972$ probably provides the most reliable information about the Si layer in optical spectra.  If luminosity is to be inferred from spectra, a simultaneous evaluation of several line strengths or ratios is most promising. In order to find additional reliable luminosity indicators, one can apply our methods to further lines caused by single ions."
    },
    "0806/0806.3082_arXiv.txt": {
        "abstract": "We consider the possibility that gravity breaks parity, with left and right handed gravitons coupling to matter with a different Newton's constant and show that this would affect their zero-point vacuum fluctuations during inflation. Should there be a cosmic background of gravity waves,  the effect would translate into anomalous CMB polarization. Non-vanishing $TB$ (and $EB$) polarization components emerge, revealing interesting experimental targets. Indeed if reasonable chirality is present a TB measurement would provide the easiest way  to detect a gravitational wave background. We speculate on the theoretical implications of such an observation. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0806/0806.0155_arXiv.txt": {
        "abstract": "\\noindent We study the formation of monopoles and strings in a model where $SU(3)$ is spontaneously broken to $U(2)=[SU(2)\\times U(1)]/\\ZZ_2$,  and then to $U(1)$. The first symmetry breaking generates monopoles with both $SU(2)$ and $U(1)$ charges since the vacuum manifold is $\\CC P^2$. To study the formation of these monopoles, we explicitly describe an algorithm to detect topologically non-trivial mappings on $\\CC P^2$. The second symmetry breaking creates $\\ZZ_2$ strings linking either monopole-monopole pairs or monopole-antimonopole pairs. When the strings pull the monopoles together they may create stable monopoles of charge 2 or else annihilate. We determine the length distribution of strings and the fraction of monopoles that will survive after the second symmetry breaking. Possible implications for topological defects produced from the spontaneous breaking of even larger symmetry groups, as in Grand Unified models, are discussed. ",
        "introduction": "\\label{model}  Our model contains an $SU(3)$ adjoint field, $\\Phi$, whose vacuum expectation value (VEV) implements the symmetry breaking in Eq.~(\\ref{su3breaking}). Two more $SU(3)$ adjoint fields, $\\Psi_1$ and $\\Psi_2$, acquire VEVs to break the $SU(2)$ subgroup of $U(2)$ to $\\ZZ_2$ as in Eq.~(\\ref{su2breaking}). The Lagrangian for the model is \\begin{eqnarray} L &=& \\frac{1}{4}\\tr[(D_\\mu \\Phi)^2] + \\frac{1}{4}\\sum_{i=1}^2\\tr[(D_\\mu \\Psi_i)^2] \\nonumber\\\\ && - \\frac{1}{8} \\tr (X_{\\mu\\nu}X^{\\mu\\nu}) - V(\\Phi,\\Psi_1,\\Psi_2), \\end{eqnarray} where $D_\\mu\\Phi = \\partial_\\mu\\Phi -i g [X_\\mu,\\Phi]$, $X_{\\mu\\nu}$ is the field strength for the $SU(3)$ gauge field $X_\\mu$, and the potential, $V$, is assumed to have a form that is suitable to give the fields the desired VEVs.  The first stage of symmetry breaking is achieved by the VEV \\begin{equation} \\Phi = \\Phi^{(0)} \\equiv \\eta T^8 \\equiv \\frac{\\eta}{\\sqrt{3}} \\begin{pmatrix} 1 & 0 & 0 \\\\ 0 & 1 & 0 \\\\ 0 & 0 & -2 \\end{pmatrix} , \\end{equation} where $\\eta$ is the energy scale at which the first symmetry breaking occurs and will be set to unity since its value has no effect on the topological structures we are considering. (We could also take $\\Phi = g \\Phi^{(0)} g^\\dag$ for any global $g \\in SU(3)$.) The vacuum manifold at this stage is \\begin{equation} SU(3)/U(2) \\cong \\CC P^2. \\end{equation} Points on $\\CC P^2$ are labeled by three complex numbers $(z_1,z_2,z_3)$, identified under a (complex) rescaling \\begin{equation} Z^T \\equiv (z_1,z_2,z_3) \\cong \\kappa (z_1,z_2,z_3) \\ , \\ \\ \\kappa \\in \\CC, \\ \\ \\kappa\\ne 0. \\label{rescaling} \\end{equation} It will be convenient for us to label the points, following \\cite{Bengtsson:2001yd}, by a point on an octant of a two-sphere given by ${\\bar \\theta}$ and ${\\bar \\phi}$, and two phases, $\\alpha$ and $\\beta$: \\begin{equation} Z^T = (\\sin{\\bar \\theta} \\cos{\\bar \\phi} ~ e^{i\\alpha}, \\sin{\\bar \\theta} \\sin{\\bar \\phi} ~ e^{i\\beta}, \\cos{\\bar \\theta}), \\label{cp2point} \\end{equation} with $0\\le {\\bar \\theta}, {\\bar \\phi}  \\le \\pi /2$ and $0 \\le \\alpha , \\beta \\le 2\\pi$.  The relation between the field $\\Phi$ and a point on $\\CC P^2$ is \\begin{equation} \\Phi =  \\frac{1}{\\sqrt{3}}\\left( {\\bf 1}  - 3\\frac{Z Z^\\dag}{Z^\\dag Z}\\right). \\label{ZtoPhi} \\end{equation}  The second homotopy group of $\\CC P^2$ is known to be the set of integers $\\mathbb{Z}$. A topologically non-trivial configuration can be constructed explicitly by taking ${\\bar \\phi}=0$. The points on the ${\\bar \\phi}=0$ sub-manifold are \\begin{equation} Z^T = (\\sin{\\bar \\theta} ~ e^{i\\alpha}, 0, \\cos{\\bar \\theta}) \\end{equation} and these describe a $\\CC P^1$ subspace of $\\CC P^2$. The points on a two-sphere in physical space, labeled by $(\\theta,\\phi)$, can be mapped onto this $\\CC P^1$ using \\begin{equation} {\\bar \\theta} = \\theta/2 , \\ \\ {\\bar \\phi}=0 , \\ \\ \\alpha = \\phi , \\ \\ \\beta =0. \\label{mmap2} \\end{equation} Equivalently, \\begin{equation} \\Phi = \\frac{1}{2\\sqrt{3}} \\begin{pmatrix} 3\\cos\\theta-1 & 0 & -3\\sin\\theta\\,e^{i\\phi} \\\\ 0             & 2 & 0 \\\\ - 3\\sin\\theta\\,e^{-i\\phi} & 0 & -3\\cos\\theta-1 \\end{pmatrix}. \\label{mmap} \\end{equation} This map represents a simple example of a monopole.  An expression for the topological charge of a monopole can be derived by first constructing the 1-form ``gauge potential'' \\begin{equation} A=\\frac{1}{2i}\\frac{Z^\\dagger dZ-dZ^\\dagger Z}{Z^\\dagger Z}. \\label{gaugepotential} \\end{equation} Note that under the ``gauge transformation'' $Z\\to Ze^{i\\lambda}$, which is a special case of (\\ref{rescaling}), $A$ transforms as $A\\to A+d\\lambda$. The corresponding field strength 2-form is \\begin{equation} F=dA=\\frac{1}{i}\\left(\\frac{dZ^\\dagger\\wedge dZ}{Z^\\dagger Z}- \\frac{dZ^\\dagger Z\\wedge Z^\\dagger dZ}{(Z^\\dagger Z)^2}\\right). \\label{F} \\end{equation} Since this 2-form is exact, its integral over a closed two-surface is a topological invariant --- and moreover is zero unless the surface contains in its interior a point or points where $Z=0$ (so that $A$ is undefined). So the expression for the topological charge in a volume $V$ with closed boundary $\\partial V$ is \\begin{equation} Q = \\frac{1}{2\\pi} \\int_{\\partial V} F = \\frac{1}{4\\pi} \\int_{\\partial V} d^2S^i \\epsilon^{ijk}F_{jk}. \\label{magcharge1} \\end{equation}  There is another way to obtain the expression for the topological charge. We start with the expression known for the 't-Hooft-Polyakov monopole in $SU(2)$ and extend it to $SU(3)$: \\begin{equation} Q = \\frac{1}{8\\pi} \\int_{\\partial V} d^2 S^i f_{abc} \\epsilon^{ijk} n^a \\partial_j n^b \\partial_k n^c, \\label{magcharge2} \\end{equation} where \\begin{equation} n^a = \\frac{Z^\\dag T^a Z}{Z^\\dag Z}, \\end{equation} with $a,b,c=1,\\ldots,8$. Here the $T^a$ are the generators of $SU(3)$, normalized by $\\tr(T^a T^b)=2\\delta^{ab}$, the $f_{abc}$ are structure constants defined by $[T^a,T^b]=2if_{abc}T^c$, and the integration is over the two sphere at infinity. Also note that the vector $n^a$ satisfies $n^a n^a =4/3$. In Appendix~\\ref{topcharge} we show that the two forms for the topolgical charge are equivalent.  It is simple to check that $Q=1$ for the monopole configuration in Eq.~(\\ref{mmap2}) and Eq.~(\\ref{mmap}). The formula in Eq.~(\\ref{magcharge1}) will be useful to locate monopoles in our numerical work described in Sec.~\\ref{algorithm}.  The second stage of symmetry breaking is more involved.  The fields $\\Psi_j$ now also acquire VEVs, which are required to lie in the unbroken $SU(2)$ subgroup, and hence commute with $\\Phi$.  Their magnitudes $\\tr(\\Psi_j^2)$ are fixed by the potential, and they are also required to be mutually orthogonal in the sense that $\\tr(\\Psi_1 \\Psi_2)=0$. Given a value of $\\Phi$ at some spatial point $P$, we need to identify this unbroken subgroup. The standard procedure is to work out commutators of $\\Phi$ with $SU(3)$ generators and to find linear combinations of the generators that commute. In practice, it is easier to first rotate $\\Phi$, say by an $SU(3)$ rotation $R$, to the reference direction, $\\Phi^{(0)}$.  We discuss how to choose $R$ below.  Then the generators of the unbroken $SU(2)$ sit in the $2\\times 2$ upper left corner while the generator $T^8$ of the unbroken $U(1)$ is in the direction of $\\Phi^{(0)}$ itself. With respect to $\\Phi^{(0)}$, the VEVs of $\\Psi_1$ and $\\Psi_2$ can be written in terms of two orthonormal 3-vectors, ${\\bf a}$ and ${\\bf  b}$, as $\\Psi_1^{(0)}={\\bf  a}\\cdot {\\bf T}$ and $\\Psi_2^{(0)}={\\bf  b}\\cdot {\\bf T}$ where \\begin{equation} T^i = \\begin{pmatrix} \\sigma_i &  & {\\bf 0} \\\\ &  & & \\\\ {\\bf 0} &  & 0 \\end{pmatrix},  \\ \\ i=1,2,3, \\label{TSU2} \\end{equation} and $\\sigma_i$ are the Pauli spin matrices. Once $\\Psi_1^{(0)}$ and $\\Psi_2^{(0)}$ are constructed, we can rotate all the fields back to the original point using $R^\\dag$.  The VEVs of $\\Psi_1$ and $\\Psi_2$ break $SU(2)$ down to $\\ZZ_2$, which is the center of $SU(2)$, $\\{ {\\bf 1}, -{\\bf 1}_2 \\}$, i.e.\\  the identity element of $SU(3)$ and $-{\\bf 1}_2 \\equiv {\\rm diag}(- 1 , - 1, 1)$. A string passes through a spatial contour if $\\Psi_1$ and $\\Psi_2$ are such that, on going around the contour, these fields are transformed by the element $-{\\bf 1}_2$ and not by the identity element. The strings are of the $\\ZZ_2$ variety and there is no distinction between a string and an anti-string. Also, there is no known integral formula that can be used to evaluate the winding around the contour. ",
        "conclusions": "\\label{discussion}  We have studied the formation of monopoles connected by strings in an $SU(3)$ model and the results for the distribution of monopoles and strings are summarized in Sec.~\\ref{results}. Here we discuss qualitatively how a similar analysis in realistic grand unified models would proceed. Our experience with $SU(3)$ helps us understand and appreciate the difficulties that are likely to be encountered. As an example, consider the minimal grand unified model based on a $SU(5)$ symmetry group. The symmetry breaking pattern is \\begin{equation} SU(5) \\to [SU(3)\\times SU(2) \\times U(1)]/\\ZZ_3 \\times \\ZZ_2. \\end{equation} and, if the non-Abelian magnetic charges are confined, the relevant symmetry breakings are \\begin{equation} SU(3) \\to \\ZZ_3 \\ , \\ \\ SU(2) \\to \\ZZ_2. \\end{equation} The fundamental magnetic monopoles carry $SU(3)$ and $SU(2)$ charges in addition to the topological $U(1)$ charge. Therefore each monopole will get connected to a $\\ZZ_3$ string and another $\\ZZ_2$ string. Then isolated clusters of monopoles come in two varieties, similar to known baryons and mesons, as shown in Fig.~\\ref{su5singlet}. However, a likely outcome at formation seems to be that, in addition to some isolated baryonic and mesonic clusters, the monopole-string network percolates and we essentially obtain one giant structure, such as depicted in Fig.~\\ref{su5network}.  \\begin{figure} \\includegraphics[width=2.2in,angle=0]{fig6.eps} \\caption{A cluster of 6 monopoles can form a singlet of $SU(3)$ and $SU(2)$, as in ordinary baryons. A bound state of a monopole and antimonopole is also possible, as in ordinary mesons. The $SU(3)$ charge on a monopole is shown in shades of grey (or in color) and the $SU(2)$ charge as a $\\pm$. We have not shown the $U(1)$ charge. $\\ZZ_3$ strings are shown as solid lines; $\\ZZ_2$ strings as dashed lines.} \\label{su5singlet} \\end{figure}  \\begin{figure} \\includegraphics[width=2.2in,angle=0]{fig7.eps} \\caption{Drawing of an infinite monopole-string network that could result from $SU(5)$ grand unified symmetry breaking. The three different shades of circles represent the $SU(3)$ color charge and the plus-minus symbols within the circles the $SU(2)$ charge. The $U(1)$ (hypercharge) charge has not been shown. The isolated clusters of monopoles have to occur in $SU(3)$ and $SU(2)$ singlets.} \\label{su5network} \\end{figure}  It seems hard to explicitly confirm if the network percolates, say by numerical simulation. For example, the vacuum manifold at the first stage of symmetry breaking is 12 dimensional and it also does not fall into a straightforward category like $\\CC P^n$. Determining the distribution of strings is also more complicated since the $SU(3)$ breaking leads to $\\ZZ_3$ strings. These problems do not seem insurmountable but are hard enough that we have not attempted to solve them at the present time.  If very few baryonic clusters form and instead an infinite monopole-string network forms, our experience with string networks \\cite{Aryal:1986cp,Vachaspati:1986cc, Copeland:2005cy,Hindmarsh:2006qn} suggests that the network energy density scales with time and never comes to dominate the universe. Processes such as monopole-antimonopole annihilation and meson formation could dissipate the energy of the network at a rate that is determined by the Hubble expansion. However, this scenario ignores the process of baryon formation from the network. Depending on the rate of this process, we could still have a monopole over-abundance problem coming from the production of baryonic clusters."
    },
    "0806/0806.1962_arXiv.txt": {
        "abstract": "New experiments designed to discover a weakly interacting dark matter (DM) particle via spin dependent scattering can distinguish models of electroweak symmetry breaking.  The plane of spin dependent versus spin independent DM scattering cross sections is a powerful model diagnostic.  We detail representative predictions of mSUGRA, singlet extended SM and MSSM, a new Dirac neutrino, Littlest Higgs with $T$-parity (LHT) and Minimal Universal Extra Dimensions (mUED) models.  Of these models, the nMSSM has the largest spin dependent (SD) cross section.  It has a very light neutralino which would give lower energy nuclear recoils. The Focus Point region of mSUGRA, mUED and the right handed neutrino also predict a very large SD cross section and predict a large signal of high energy neutrinos in the IceCube experiment from annihilations of dark matter in the Sun.  We also describe a model independent treatment of the scattering of DM particles of different intrinsic spins. ",
        "introduction": "The nature of the dark matter (DM) that comprises about 23\\% of the energy density of the Universe is unknown and a large experimental effort is devoted to its explication.  The premise is that dark matter is a stable elementary particle.  The large scale structure of the Universe excludes hot dark matter as the  primary DM component.  There are a number cold and warm dark matter particle candidates, many of which are well motivated by theoretical models of TeV scale physics. For recent reviews see Refs.~\\cite{Baer:2008uu,Hooper:2008sn,Arrenberg:2008wy,dm08:2008aa,ppc08:2008aa}. Some DM candidates (e.g. axino, heavy gravitino, keV sterile neutrino) may have only cosmological consequences.  The axion, which can solve the strong CP problem, may be detectable from their conversion by photons in a strong magnetic field.  The DM particle of models with a conserved discrete symmetry (e.g. the stable particle of supersymmetry with a conserved R-parity) can have rich consequences for collider experiments, direct detection experiments  (the elastic scattering of DM particles on nuclei), and astrophysical experiments (via the detection of gamma rays, positrons, antiprotons, antideuterons, or neutrinos from dark matter annihilations in the galatic halo or the center of the Sun).   A weak scale cross section can naturally lead to a dark matter density of the requisite value to explain the WMAP determination~\\cite{Komatsu:2008hk}.  Our interest here is in the tests of weakly interacting massive particles (WIMP) models via their elastic scattering on nuclei.   In WIMP models a discrete symmetry exists that makes the DM candidate stable.   In the minimal supersymmetric standard model (MSSM), the DM particle is the lightest neutralino, a spin 1/2 Majorana particle.  In a non-SUSY context, a heavy Dirac fermion can be a viable DM candidate~\\cite{Boehm:2003hm,Fayet:2004bw,Adibzadeh:2008pe}.  Models such as Minimal Universal Extra Dimensions~\\cite{Antoniadis:1990ew,Randall:1999ee,Appelquist:2000nn,ArkaniHamed:1998nn,Nath:1999mw,Dicus:2000hm,Servant:2002aq,Macesanu:2005jx,Kong:2005hn,Hooper:2007qk} incorporate a Kaluza-Klein (KK) parity which keeps the lightest KK particle  stable.  Here, the DM particle is spin-1.  In the Littlest Higgs model solution to the gauge hierarchy problem, a discrete $T$-parity ensures that a spin-1 particle is stable~\\cite{Low:2004xc,Hubisz:2004ft}.  Scalar DM, with a spin-0 field can be a consequence of singlet Higgs field extensions of the Standard Model which have been explored in a number of works~\\cite{Silveira:1985rk,McDonald:1993ex,Boehm:2003hm,Burgess:2000yq,Barger:2007im,Davoudiasl:2004be,Fayet:2004bw} and a model with two extra dimensions~\\cite{Dobrescu:2007ec}.  There are a variety of other models with a DM particle that we do not specifically consider here, see e.g.~\\cite{Dolle:2007ce,Chang:2006ra,TuckerSmith:2004jv,Ishiwata:2007ck,Han:1997wn}  SI scattering occurs though Higgs exchange and thus may well be much smaller than SD scattering which occurs through Z-exchange~\\footnote{In SUSY, sfermion exchanges can contribute to both SI and SD scattering.  However, their contributions are typically small, even for light squarks.}.  The calculation of the SI cross section has substantial uncertainties from the estimate of the Higgs coupling to strange quarks~\\cite{Ellis:2005mb,Ellis:2008hf}.   The first WIMP searches focused on spin-independent (SI) scattering on a nucleus and have placed limits approaching $10^{-8}$ pb.  The SI experiments are designed to detect the coherent recoil of the nucleus caused by the DM scattering.  For a heavy nuclear target, the coherent scattering increases the cross-section by the square of the Atomic Number.   The best current limits on SI scattering have been placed by the CDMS 5 tower~\\cite{Ahmed:2008eu} experiment and the XENON10 experiment~\\cite{Angle:2007uj}.  The CDMS upper bound on the WIMP nucleon cross section is $4.6\\times10^{-8}$ pb for a WIMP mass $\\sim 60$ GeV.   The XENON10  bound is $4.5\\times10^{-8}$ pb for a WIMP mass of $\\sim 30$ GeV and $8.8\\times10^{-8}$ pb for a WIMP mass of $\\sim 60$ GeV.  The present SI experiments have detector masses of order 10 kg.  The next generation experiments will have about 100 kg mass.  Detectors with 1000 Kg (1T) size are under design.  The XENON100 experiment is expected to reach a SI cross section sensitivity of $10^{-9}$ pb for a 100 GeV WIMP~\\cite{Aprile:2008dm}.  The LUX experiment, with a 100 kg xenon detector, expects to reach $4\\times10^{-10}$ pb at $M_{DM}\\sim40$ GeV~\\cite{Wang:2008dm}. The ultimate goal of XENON1T is sensitivity to a SI cross section of $10^{-10}$ pb.  One ton Liquid Argon detectors promise competitive sensitivities~\\cite{Nikkei:2008dm,warp,Laffranchi:2007da}.  The SD cross-section occurs through the axial vector coupling to the spin content of the nucleus; there is a $J(J+1)$ enhancement from the nuclear spin $J$. Until recently, the SD cross section limits had been about 6 orders of magnitude weaker than for SI, with the best bound from ZEPLIN-II at 0.07 pb for a DM particle of mass $\\sim50$ GeV scattering on neutrons~\\cite{al.:2007xs}.  The NAIAD, COUPP, and KIMS experiments had placed upper bounds on the cross section for SD scattering of dark matter on protons of 0.5 pb, 0.3 pb and 0.2 pb, respectively, for a DM mass of order 100 GeV~\\cite{Alner:2005kt,Bolte:2006pf,Lee.:2007qn}. The SuperKamiokande (SK) search for neutrinos from DM annihilations in the Sun placed a stronger limit, albeit model dependent, on the SD cross section;  converting the SK neutrino flux limit requires assumptions about the DM mass and the cross section for DM annihilation to neutrinos.  For a DM mass of order 100 GeV the SK limits on the SI and SD cross sections are of order $10^{-5}$ and $0.6\\times10^{-2}$ pb, respectively~\\cite{Habig:2001ei,Desai:2004pq}.   Liquid Xenon has about 50 percent odd isotopes which allows a measurement of SD scattering.  New SD limits  have been reported by the XENON10 experiment of $0.5\\times10^{-2}$ pb for scattering on protons and $0.5$ pb for scattering on neutrons at a WIMP mass of $\\sim 30$ GeV~\\cite{Manalaysay:2008dm}.  Models generally predict the same cross section for scattering on protons and on neutrons, so the most restrictive of the two can be imposed in constraining models.   New techniques with bubble technologies offer great promise for exploring much smaller SD cross sections in future experiments.  The Picasso experiment in SNOLAB uses a superheated liquid in a gel.  Its current proton-WIMP SD cross section bound is of order 1 pb~\\cite{BarnabeHeider:2005pg}.  In a future third stage of this experiment, with a 100 kg detector, sensitivity for SD scattering down to $10^{-4}$ pb is anticipated.  The COUPP experiment~\\cite{Bolte:2006pf,Behnke:2008zz} uses a continuously sensitive bubble chamber with compounds $CF_3I$ and $C_4 F_{10}$ of fluorine, $F^{19}$, and iodine, $I^{127}$.  WIMP scattering from fluorine is determined by $\\sigma_{SD}$ and scattering from iodine is determined by $\\sigma_{SI}$.  %  A strong SD cross section $\\gtrsim 10^{-4}$ pb would also likely yield an observable signal of high energy neutrinos from the Sun due to DM annihilations~\\cite{Barger:2001ur,Cirelli:2005gh,Halzen:2005ar,GonzalezGarcia:2005xw,Barger:2007hj,Barger:2007xf}.  A large signal in neutrino telescopes such as IceCube of high energy neutrinos from DM annihilations in the Sun can occur only if there is a large SD scattering (since the SI cross section is already known to be $\\ll 10^{-6}$ pb.  The capture rate is \\cite{Gould:1992xx} \\be C_{\\odot}= 3.4\\times 10^{20} {\\rm s}^{-1} {\\rho_{local}\\over 0.3 \\text{ GeV/cm}^3} \\left({270\\text{ km/s}\\over v_{local}}\\right)^3\\left({\\sigma_{SD}^H+\\sigma_{SI}^H + 0.07 \\sigma_{SI}^{He}\\over 10^{-6}\\text{ pb}}\\right)\\left({100\\text{ GeV}\\over m_{\\N_1}}\\right)^2, \\label{eq:caprate} \\ee where $\\rho_{local}$ and $v_{local}$ are the local density and velocity of relic dark matter, respectively.  Present limits on the SI and SD scattering cross section on nuclei strongly suggest that an observation of these neutrinos from the Sun would only result from SD capture of WIMPS by hydrogen in the Sun.   To calculate the expected signal rates, we closely follow the approach of Ref.~\\cite{Barger:2007xf}.  Recently, the DAMA/LIBRA experiment and the predecessor DAMA/NaI experiment have reported a combined 8.2$\\sigma$ evidence for an annual modulation signature for DM particles in the galactic halo~\\cite{Bernabei:2008yi}.  The implied SI cross section would be about $10^{-6}$ pb. A possible way to reconcile the tension with other more restrictive bounds from other experiments is that DAMA is not detecting nuclear recoils, but electromagnetic energy deposition from axion-like interactions in the detector.  For now, we set aside this interesting experimental development and focus on nuclear recoil detection of DM.  Our study is devoted to the analysis of what can be learned from the combined measurements of SI and SD processes, which can provide a powerful diagnostic in differentiating models.  We specifically detail representative predictions of mSUGRA~\\footnote{Recently, efforts have been made to distinguish specific areas of mSUGRA parameter space that yield similar collider signatures using measurements from direct detection experiments~\\cite{Altunkaynak:2008ry}.}, singlet extensions of the SM and the MSSM, a new Dirac neutrino, Littlest Higgs with $T$-parity, and Minimal Universal Extra Dimensions models (mUED).  We assume that the DM particle is in thermal equilibrium in the early universe and we impose the requirement that the relic density of DM reproduces the measured WMAP density within $\\pm 10\\%$.  We also present general formulas for SI and SD cross sections for the scattering of WIMPS of intrinsic spin 0, 1/2, 1, and 3/2.   The organization of the paper is as follows.  In section \\ref{sect:scan}, we outline the scan proceedure while in section \\ref{sect:predict} we elucidate the DM scattering predictions of representative models that provide a WIMP candidate and detail how the models could be distinguished by measurement of SI and SD scattering.  Then we proceed, in section \\ref{sect:modindep}, to a model independent analysis of DM scattering for DM particles of arbitrary spin. Finally, we summarize the results of our study in  section \\ref{sect:concl}. ",
        "conclusions": "\\label{sect:concl} Both direct detection experiments for DM scattering and LHC experiments can discover a weakly interacting DM particle and reveal its properties.  In Section \\ref{sect:predict} we showed that the SD versus SI cross section provides an excellent DM diagnostic.  We highlighted this by scanning over the parameter space of six Beyond the SM models which supply viable DM candidates:  mSUGRA, singlet extended SM and MSSM, a stable Dirac neutrino, Littlest Higgs with $T$-parity and Minimal Universal Extra Dimensions.  To thoroughly scan the parameter spaces, we adopted the Bayesian method that is the foundation of the Markov Chain Monte Carlo approach.  The DM models can have distinct phenomenological predictions.  We showed that  the DM model possibilities can be narrowed by measurements of both SI and SD elastic scattering.  The direct signals for DM in recoil and neutrino telescope experiments are complementary to LHC experiments in distinguishing the beyond the standard model physics scenarios~\\cite{Hubisz:2008gg}.    \\begin{figure}[htbp] \\begin{center} \\includegraphics[angle=0,width=0.67\\textwidth]{figs/all-sipsdp.eps} \\caption{The posterior distributions of $\\sigma_{SI}^{DM-p}$ and $\\sigma_{SD}^{DM-p}$ for six models.} \\label{fig:allxsect} \\end{center} \\end{figure}    We  summarize below the model predictions for the DM cross sections; the posterior distributions are summarized in Fig. \\ref{fig:allxsect}. \\bi \\item In mSUGRA, the FP region provides the largest SI and SD cross sections.  This is due to the mixed Higgsino nature of the lightest neutralino; the neutralino couplings to the Higgs and $Z$ bosons are large.  The Bino nature of the lightest neutralino in the CA and AF regions causes these scenarios to have substantially smaller cross sections. \\item The tadpole nMSSM model has large SD scattering, of order $10^{-3}$ pb,  and a wide range of SI cross section.  This is a consequence of the DM annihilation occuring through the $Z$ boson . To counter the small annihilation rate in the early universe (due to the small neutralino mass in the model), the neutralino pair is required to have a larger $Z$ boson coupling, resulting in a large SD cross section. \\item In the singlet extended SM, the DM candidate is spin-0, which gives a vanishing SD cross section.  The SI cross section is generally small, below $\\sim 10^{-8}$ pb, and occurs through Higgs boson exchange.  If SD scattering of DM is observed, the class of models with spin-0 DM would be immediately excluded as being the sole origin of the DM in the universe. \\item For stable Dirac neutrino DM,  the $Z$ boson dominates in the calculation of both the relic density and elastic scattering cross section and makes the SI and SD cross sections tightly correlated and large. \\item In mUED, a 'sweet spot' of $\\sigma_{SD} \\sim {\\cal O}(10^{-6})$ pb exists for which the DM relic density is reproduced.  The relic density is strongly dependent on the curvature parameter and fixes its value.  The KK quarks have approximately the same mass as the inverse curvature and the SD cross section is thus closely tied to the relic density.  On the other hand, the $\\sigma_{SI}$ cross section is more dispersed due to the larger variation of the Higgs boson mass. \\item In the LHT model, the SD interaction occurs through T-odd quarks which have a small hypercharge.  Therefore, the SD cross section in this model is typically very small.  In contrast, the SI scattering proceeds through the Higgs and $T$-odd quarks, giving experimentally accessible SI cross section values.  \\ei  \\begin{figure}[htbp] \\begin{center} \\includegraphics[angle=0,width=0.47\\textwidth]{figs/allbox-sip.eps} \\includegraphics[angle=0,width=0.47\\textwidth]{figs/allbox-sdp.eps} \\caption{Box and whisker plot of the SI and SD cross sections.  The box represents the coverage of the middle 50-percentile and the whisker is $\\sim 100\\%$ coverage.  The xSM is not shown in the SD panel as its cross section vanishes. The large SD cross section of the nMSSM is difficult to probe because of the low ($\\sim 20-30$ GeV) mass of the DM particle in this model.} \\label{fig:boxwhisker} \\end{center} \\end{figure} We provide a summary of the SI and SD cross sections by the box and whisker plots in Fig. \\ref{fig:boxwhisker}.   The boxes represents the coverage of the middle 50-percentile. We summarize the forecast for observing a signal in neutrino telescopes or the scattering cross sections at future recoil experiments in Table \\ref{tab:illust}.  We have also presented model-independent parameterizations of the SI and SD cross sections for DM of arbitrary spins.   \\begin{table}[htdp] \\caption{Survey of models and DM scattering predictions organized by the intrinsic spin of the DM candidate.  The forecast for IceCube detection is also given, wherein we describe in what portion of the parameter space a $5\\sigma$ discovery can be made.  Instances where the likelihood of detection via IceCube are low are marked by $\\times$.  The  $2\\sigma$ ranges for $\\sigma_{SI}^{DM-p}$ and $\\sigma_{SD}^{DM-p}$ are given in pb units.} \\begin{center} \\begin{tabular}{|c|c|cccc|} \\hline $S$ & Model & Candidate & IceCube Prospects & $\\sigma_{SI}^{DM-p}$ &$\\sigma_{SD}^{DM-p} $\\\\ \\hline 0&xSM& $S$ & $\\times$  &  $6\\times10^{-10}-4\\times10^{-9}$ & $\\times$ \\\\%~\\cite{McDonald:1993ex, Boehm:2003hm, Burgess:2000yq, Barger:2007im, Davoudiasl:2004be,Fayet:2004bw} \\hline $\\half$&mSUGRA& $\\widetilde \\chi_1^0$ & Majority of FP region &   $4\\times10^{-10}-8\\times10^{-8}$&   $1\\times10^{-7}-2\\times10^{-4}$ \\\\ % $\\half$&nMSSM& $\\widetilde \\chi_1^0$ & $\\times$ by $E_{\\mu}$ threshold &  $1\\times10^{-9}-5\\times10^{-8}$ &  $10^{-3}-10^{-2}$   \\\\%~ \\cite{Barger:2006kt,Barger:2007nv,Balazs:2007pf} $\\half$&RHN& $\\nu'$ & Portion  & $9\\times10^{-10}-9\\times10^{-8}$  &  $7\\times10^{-8}-3\\times10^{-5}$ \\\\%~ \\cite{Boehm:2003hm,Fayet:2004bw} \\hline $1$&mUED& $B_{1}$ & Portion &  $2\\times10^{-12}-4\\times10^{-11}$&  $1\\times10^{-6}-3\\times10^{-6}$ \\\\%~\\cite{Antoniadis:1990ew,Randall:1999ee,ArkaniHamed:1998nn,Nath:1999mw,Dicus:2000hm,Macesanu:2005jx,Kong:2005hn} $1$&LHT& $A_{H}$ & $\\times$ by $\\sigma_{SD}$ suppression &  $6\\times10^{-11}-4\\times10^{-9}$ &  $2\\times10^{-13}-4\\times10^{-9}$ \\\\%~\\cite{ArkaniHamed:2002qy,Han:2003wu,Schmaltz:2005ky} \\hline \\end{tabular} \\end{center} \\label{tab:illust} \\end{table}%            \\vskip1cm \\hrule \\vskip1cm"
    },
    "0806/0806.4892_arXiv.txt": {
        "abstract": "GD~552 is a high proper motion star with the strong, double-peaked emission lines characteristic of the dwarf nova class of cataclysmic variable star, and yet no outburst has been detected during the past 12 years of monitoring. We present spectroscopy taken with the aim of detecting emission from the mass donor in this system. We fail to do so at a level which allows us to rule out the presence of a near-main-sequence star donor. Given GD~552's orbital period of $103$ minutes, this suggests that it is either a system that has evolved through the $\\sim 80$-minute orbital period minimum of cataclysmic variable stars and now has a brown dwarf mass donor, or that has formed with a brown dwarf donor in the first place. This model explains the low observed orbital velocity of the white dwarf and GD~552's low luminosity. It is also consistent with the absence of outbursts from the system. ",
        "introduction": "\\label{sec:introduction} GD~552 is a blue, high proper motion star ($0.18''/$yr) discovered by Giclas et al. (1970). It was first observed spectroscopically by Greenstein \\& Giclas (1978), who found that it is a cataclysmic variable star (CV) in which a white dwarf accretes matter from a low-mass companion (often termed primary and secondary star, respectively). GD~552's proper motion and position close to the plane of the Galaxy (galactic latitude $4^\\circ$), combined with its blue colour all suggest that it is relatively close by. Greenstein \\& Giclas (1978) suggest a distance of $\\sim 70\\,$pc which gives a transverse motion which is reasonable for  a member of the Galactic disk, and combined with its magnitude $V = 16.5$, suggests $M_V \\sim 12.5$, the equivalent of a $0.6\\,\\msun$ white dwarf with a temperature of only $9000\\,$K. The probable low luminosity of GD~552, which is of central importance to this paper, is backed up by an absence of any observed outbursts, suggesting that it may very rarely or never have outbursts (Cannizzo 1993).  The main spectroscopic study of GD~552 so far has been carried out by Hessman \\& Hopp (1990) (hereafter HH1990) who determined the orbital period of GD~552 of $102.7\\,$min. They observed an extreme Balmer decrement (\\ha:\\hb = 6.2:1.0), indicative of a cool, optically thin disc, another indication of low luminosity (Williams \\& Shipman 1988). HH1990 measured the white dwarf's projected orbital velocity to be $K_1 = 17\\pm4 \\kms$. This is a very low value suggestive of a low inclination system, since for edge-on systems with orbital periods similar to GD~552 $K_1$ is typically $\\sim60-80 \\kms$. However, the emission lines from the accretion disc suggest instead a moderately inclined system as they display clearly separated double-peaks (which come from the outer disc, Horne \\& Marsh 1986) with velocities of $\\pm450 \\kms$ which can be compared to typical peak velocities of $\\sim600 \\kms$ for edge-on systems. To solve this conundrum, HH1990 suggested that the white dwarf in GD~552 is unusually massive -- close to the Chandrasekhar limit in fact -- allowing a low orbital inclination ($\\sim 20^\\circ$) and therefore the small $K_1$ value at the same time as large disc velocities. HH1990 were forced to their model because they assumed that the companion to the white dwarf had to be a main-sequence star.  Since the companion fills its Roche lobe, Roche geometry and the orbital period uniquely specify its density ($36.5\\,{\\rm g}\\,{\\rm cm}^{-3}$), which, if it is a main-sequence star, also fixes its mass, which turns out to be $M_{2} \\sim 0.13\\,\\msun$. This would cause a much larger $K_1$ than observed, unless the inclination is low.  More recent theoretical work has revised HH1990's estimate a little to $M_{2} \\sim 0.15\\,\\msun$ from GD~552's orbital period (Kolb \\& Baraffe 1999), but the problem is qualitatively unchanged. However, this work also suggests a very different scenario. Cataclysmic variables at long orbital periods are thought to evolve towards shorter periods until the companion becomes degenerate at a period near 70 minutes.  After this time, the mass donor increases in size as it loses mass and the orbital period lengthens (e.g. Howell, Nelson \\& Rappaport, 2001; Kolb 1993). Thus although if GD~552 is approaching the period minimum its donor mass must be $M_{2} \\sim 0.15\\,\\msun$, if it has already passed the minimum and is now evolving to longer periods, it would be about a factor of four times less massive, and there would be no need for HH1990's massive white dwarf, face-on model. Moreover, such a system would have an extremely low mass transfer rate, consistent with GD~552's lack of outbursts and probably low intrinsic luminosity. The evolution just described is the standard explanation for the observed minimum orbital periods of CVs, which however fails on two counts.  First, the observed minimum around 80 min is distinctly longer than the theoretical value of 70 min (Kolb \\& Baraffe 1999; although Willems et al. (2005) propose that the period minimum might be longer than expected because of extra angular momentum loss from circumbinary disks), and, second, while we expect most systems to have passed the period minimum there is just one single well-established example of such a system known (SDSS 103533.03+055158.4, Littlefair et al. 2006) although there are a few other good candidates (Mennickent et al. 2001; Littlefair, Dhillon \\& Martin 2003; Patterson, Thorstensen \\& Kemp 2005; Araujo-Betancor et al. 2005; Southworth et al. 2006). Therefore it is of considerable interest to establish whether GD~552 is a pre- or post-period-minimum system.  In this article we carry out a test to distinguish between the two models of GD 552. If the pre-period-bounce model is correct, then, as we show in this paper, we should be able to detect features from the M star in a low mass-transfer rate system such as GD~552. If, conversely, the post-period-bounce model is correct, the donor will be extremely faint and it should not be detectable. In summary, our task is to set the strongest possible limit upon the presence of a hypothetical M dwarf such that we can say that we would have seen it had GD~552 been a pre-bounce system. ",
        "conclusions": "We have used I-band spectroscopy in an attempt to detect the mass donor of the cataclysmic variable GD~552. Failure to detect the donor star puts strong limits on its mass, ruling out a main sequence nature.  We have shown GD~552 data to be consistent with a model in which its components are an ordinary white dwarf and a brown dwarf, at a moderate inclination angle. Population synthesis calculations favour the idea that the mass donor reached the stage of brown dwarf through evolution instead of being born as such. This suggests that GD~552 is likely to be a post-period minimum CV, making it the ``period bouncer'' with the longest securely-determined orbital period known.  We confirm the value obtained by HH1990 for the radial velocity semiamplitude of the white dwarf and determine for it a $T_{\\rm eff}$ of $10900\\pm400$ K by fitting the spectra with a grid of white dwarf models. The temperature determined for the white dwarf in GD552 is consistent with the red edge of the instability strip (11\\,000 $< T_{\\rm eff} <$ 12\\,000 K; Bergeron et al. 1995). From the 5.4\\,h stretch of photometry presented, we find no indication of the presence of pulsations. The lightcurve shows significant varibility but not with the periodicities and amplitudes characteristic of ZZ Cet stars."
    },
    "0806/0806.1686_arXiv.txt": {
        "abstract": "We investigate the dynamical response, in terms of disc size and rotation velocity, to mass loss by supernovae in the evolution of spiral galaxies.  A thin baryonic disc having the Kuzmin density profile embedded in a spherical dark matter halo having a density profile proposed by Navarro, Frenk \\& White is considered.  For a purpose of comparison, we also consider the homogeneous and $r^{-1}$ profiles for dark matter in a truncated spherical halo.  Assuming for simplicity that the dark matter distribution is not affected by mass loss from discs and the change of baryonic disc matter distribution is homologous, we evaluate the effects of dynamical response in the resulting discs.  We found that the dynamical response only for an adiabatic approximation of mass loss can simultaneously account for the rotation velocity and disc size as observed particularly in dwarf spiral galaxies, thus reproducing the Tully-Fisher relation and the size versus magnitude relation over the full range of magnitude.  Furthermore, we found that the mean specific angular momentum in discs after the mass loss becomes larger than that before the mass loss, suggesting that the mass loss would occur most likely from the central disc region where the specific angular momentum is low. ",
        "introduction": "It is well-known that the luminosity $L$ of luminous spiral galaxies correlates tightly with their rotation velocity $V_{\\rm rot}$, as called the Tully-Fisher relation \\citep[hereafter TFR,][]{tf77}. The TFR is usually described as a power-law scaling relation $L\\propto V_{\\rm rot}^{\\gamma}$, and the rms scatter of the TFR is smaller for optical passbands of longer wavelengths. In particular, the I-band TFR is significantly tight and its scatter is approximately as small as 0.4 mag \\citep{w97, kff02}, with a power-law index $\\gamma\\simeq 3-3.5$ \\citep[e.g.][]{pt88, s00}. Light at longer wavelengths from galaxies traces old late-type stars, unaffected by sporadic formation of young stars, and represents the mass of luminous matter.  Thus, the tightness of the I-band TFR suggests the existence of some driving mechanism that depends on mass of galaxies in processes of their formation.   The current standard model of cosmology is the cold dark matter (CDM) model.  Since the initial spectrum of density fluctuations has larger amplitudes for smaller scales in this model, the scenario of structure formation is hierarchical clustering, in the sense that smaller dark haloes cluster to form larger dark haloes hierarchically, thereby creating a scaling relation between mass $M$ and circular velocity $V_{\\rm circ}$ of dark haloes \\citep{f82, bfpr84}. High-resolution $N$-body simulations, based on a power spectrum $P(k)\\propto k^{n}$ with an index $n=-2.5$ appropriate for galaxy scales or reciprocal of wavenumber $k^{-1}$, have provided $M\\propto V_{\\rm cir}^3$, reminiscent of the TFR if a constant mass-to-light $M/L$ ratio is assumed \\citep[e.g.][]{nfw}. Consequently, the CDM model naturally involves a TFR-like scaling relation as observed for luminous spiral galaxies, because a constant $M/L$ is indeed known to be a good approximation to such galaxies \\citep[e.g.][]{fg79}.   When studying the formation of galaxies originated from the growth of density fluctuations in the early universe, however, it is also necessary to take into account the effects of star formation and supernova (SN) explosions in individual galaxies.  In particular, dwarf galaxies have shallow gravitational potential wells, so that energy feedback from SN explosions significantly affects their evolution and hence scaling relations that follow.  In a framework of monolithic collapse scenario of galaxy formation, several authors studied scaling relations for dwarf galaxies in two extreme cases, i.e., the dark matter dominated case in which dark halo dominates the gravitational potential \\citep{ds86}, and the baryon dominated case of self-gravitating galaxies without dark matter \\citep{ya87}.   \\citet{sn99} derived, for the first time by numerical simulations, the TFR for luminous spiral galaxies in the CDM model, using a method of $N$-body and smoothed particle hydrodynamics (SPH) combined with star formation and SN feedback explicitly.  They also showed that the star formation rate is regulated for $M/L$ to be almost constant. However, the zero-point of the TFR turned out to differ from that observed, so that the simulated spiral galaxies were too faint to be consistent with observations at the same circular velocity.  It was therefore pointed out that processes of star formation and SN feedback in current $N$-body/SPH simulations have to be improved \\citep[see also][]{g07, psl07}.   Semi-analytic (SA) models of galaxy formation have been invented as a complementary approach to numerical simulations, in which complex processes such as star formation and SN feedback are simply modeled on galaxy scales.  SA models have indeed succeeded in reproducing many observational results, for example, the luminosity function of galaxies as well as the TFR, by suppressing the formation of dwarf galaxies owing to SN feedback \\citep{sp99, c00, ny04, dlkw04, kjmb05, croton06}.    Recent progress in observational techniques has enabled to see much more properties of distant galaxies and faint dwarf galaxies beyond the limit of magnitude attained so far.  In contrast to the success of SA models in accounting for luminous galaxies, they came to face a serious discrepancy between predicted and observed dynamical properties of dwarf galaxies, while reproducing their photometric properties quite well \\citep[e.g.][]{c00, vdb02}.  Here we would like to stress that the dynamical response to gas removal induced by SN explosions, which has been overlooked in most SA models, is unavoidable in dwarf, less massive galaxies. \\citet{ny03} formulated how the structure of spherical galaxies responds to SN-induced gas removal in a gravitational potential of dark halo, and \\citet{ny04} incorporated this effect into a SA model to show that many properties of elliptical galaxies, including the Faber-Jackson relation \\citep{fj76}, can be reproduced to faint magnitudes of dwarf ellipticals.  A more sophisticated SA model associated with high-resolution $N$-body cosmological simulations also provides a good agreement between predicted and observed properties of galaxies \\citep{nyeyg05}.  Accordingly, as a natural extension of Nagashima \\& Yoshii's trial, it is worth investigating whether such a dynamical effect could also reproduce the observed TFR down to faint magnitudes in spiral galaxies.  In this paper, using the Kuzmin disc \\citep{k52, k56} as a galaxy disc embedded in a spherical dark halo, we present in Section 2 the general formulation for the dynamical response in size and rotation velocity of discs.  In Section 3 we examine such effect for several choices of density distribution of dark halo.  In Section 4 we derive the resulting TFRs and disc size versus magnitude relations.  In Section 5 we summarize the results of this paper.  Detailed derivations of analytic expressions are given in Appendix. ",
        "conclusions": "\\subsection{Losses of mass and angular momentum} The angular momentum distribution in dark haloes is well established theoretically based on the CDM model \\citep[e.g.][]{ct96, ng98}.  It is natural to assume that the baryonic components acquire the same amount of specific angular momentum as that of their host haloes, because the cosmic tidal field provides the angular momentum to dark matter and baryons in the same manner \\citep{w84}. This assumption has been a basis of most theories of galaxy formation.  From observations of dwarf spiral galaxies, \\citet{vdb01a} found that their mean (or {\\it total}) specific angular momenta are very similar to those of their host haloes predicted by those theories. Moreover they found that the the disc to halo mass ratio is smaller than the baryon fraction in the Universe.  This means that only a part of baryons forms discs whereas it keeps an amount of specific angular momentum similar to that of host haloes. They also examined the specific angular momentum distribution in individual haloes, and surprisingly found that it is inconsistent with theoretical prediction by \\citet{b01} based on high-resolution $N$-body simulations.  Secondly, the observed surface density distribution of dwarf spiral galaxies is described as an exponential distribution \\citep[e.g.][]{mfb92}, different from that expected from the CDM model which provides much more centrally concentrated discs when compared with the exponential distribution.  This cannot be solved even if SN feedback is simply considered \\citep{vdb01b}.  One possible solution to the above discrepancies should be obtained by considering the mass loss from the central region of discs, where the gas should have the specific angular momentum lower than that in the outer region.  SN feedback driving the mass loss without the loss of angular momentum leads discs to increasing their specific angular momentum, i.e.,  $(J/M)_{f}/(J/M)_{i}>1$.  This situation is consistent with that considered in this paper, as shown in the previous section.  So far we have assumed rotation-dominated discs throughout gas removal. With this assumption, the specific angular momentum always increases after the mass loss, yielding $(J/M)_{f}/(J/M)_{i}>1$. On the other hand, the transformation of rotation-dominated discs to velocity dispersion-dominated discs, which is the puff-up transformation, leads discs to decreasing the specific angular momentum, yielding $(J/M)_{f}/(J/M)_{i}<1$ \\citep{bs79}. This situation should occur when the gas in the outer region is removed where the specific angular momentum is large.  Such a way of gas removal would be caused by the tidal stripping in clusters of galaxies \\citep[e.g.][]{gg72,fn99,on03}.  \\citet{bs79} proposed that this mechanism should be of lenticular galaxy formation. If so, the fraction of S0 galaxies in clusters will be higher than in fields, and the rotation velocity of dwarf galaxies in clusters will be systematically lower than that in fields.  This should be observationally clarified in future.    \\subsection{Density profiles of dark haloes} We have assumed three types of density profiles of dark haloes, such as the NFW profile with an inner slope of -1 and an outer slope of -3 (equation \\ref{eq:density-nfw}), and two power-law profiles with a slope of 0 (homogeneous) and -1 ($r^{-1}$).  The NFW profile is suggested by $N$-body simulations, and many authors claim that the inner slope will become even steeper after the condensation of baryons owing to the cooling \\citep{bffp86}, which is called as the adiabatic contraction (but see \\citet{sm05} for the effects of random motions). In fact, this implicitly assumes that the cooling timescale is longer than the relaxation timescale for dark haloes to settle into the NFW profile, that is, baryons cool and shrink after virialisation. On galaxy scales, however, this would not be the case. The cooling timescale is much shorter than the dynamical timescale, and the relaxation timescale is similar to or longer than the dynamical timescale. Thus it is not unnatural to assume that a mixture of dark matter and cooled baryons virializes.  In this case, there would be galaxy discs within dark haloes with the NFW profile, without undergoing the adiabatic contraction.  This might be an opposite limiting case to the adiabatic contraction, but recent X-ray observation of the intracluster medium has found that there is no evidence of contraction of the dark halo \\citep{z06}.  Since the central cD galaxy is massive enough to suppress the SN feedback, we do not need to expect the expansion due to dynamical response to SN-induced gas removal.  Thus it is possible to say that a simple assumption of adiabatic contraction does not work in reality.  This should justify our use of the NFW profile for dark haloes that surround galaxy discs.  To clarify this, of course, high-resolution hydrodynamical simulations including the gas cooling processes are required \\citep[e.g.][]{gkkn04}.  \\subsection{Hierarchical formation of galaxies} In Section 4, our results of dynamical response to SN-induced gas removal within dark haloes are applied to scaling relations for galaxies, under the assumption that simple scalings among mass, velocity, and size of baryonic components have been set up {\\it before} the gas removal. It is reasonable to consider that these are inferred from observations of scaling relations for massive galaxies where SN feedback is not effective owing to their deep gravitational potential wells.  This seems to be somewhat in contrast to the approach of, for example, \\citet{vdb01b}, in which the direction of angular momentum vector is assumed to be invariant.  However, it has been shown that the direction can be moved by contiguous accretion of dark matter \\citep{ng98, sw04}. Therefore, it is not assured that the specific angular momentum of discs is similar to that of host haloes averaged over whole regions.  We thus believe that it is reasonable to use the scaling relations for massive galaxies as the initial state for the dynamical response to gas removal,"
    },
    "0806/0806.1223_arXiv.txt": {
        "abstract": "We show that a purely kinetic approach to the excitation of waves by cosmic rays in the vicinity of a shock front leads to predict the appearance of a non-alfv\\'enic fastly growing mode which has the same dispersion relation as that previously found by \\cite{bell04} by treating the plasma in the MHD approximation. The kinetic approach allows us to investigate the dependence of the dispersion relation of these waves on the microphysics of the current which compensates the cosmic ray flow. We also show that a resonant and a non-resonant mode may appear at the same time and one of the two may become dominant on the other depending on the conditions in the acceleration region. We discuss the role of the unstable modes for magnetic field amplification and particle acceleration in supernova remnants at different stages of the remnant evolution. ",
        "introduction": "\\label{sec:intro}  The problem of magnetic field amplification at shocks is central to the investigation of cosmic ray acceleration in supernova remnants. The level of scattering provided by the interstellar medium turbulent magnetic field is insufficient to account for cosmic rays with energy above a few GeV, so that magnetic field amplification and large scattering rates are required if energies around the knee are to be reached. The chief mechanism which may be responsible for such fields is the excitation of streaming instabilities (SI) by the same particles which are being accelerated (\\cite{skilling75c,bell78a,lc83a,lc83b}). The effect of magnetic field amplification on the maximum energy reachable at supernova remnant (SNR) shocks was investigated by \\cite{lc83a,lc83b}, who reached the conclusion that cosmic rays could be accelerated up to energies of order $\\sim 10^4-10^5$ GeV at the beginning of the Sedov phase. This conclusion was primarily based on the assumption of Bohm diffusion and a saturation level for the induced turbulent field $\\delta B/B\\sim 1$. On the other hand, recent observations of the X-ray surface brightness of the rims of SNRs have shown that $\\delta B/B\\sim 100-1000$ (see \\cite{volk} for a review of results), thereby renewing the interest in the mechanism of magnetic field amplification and in establishing its saturation level. It is however worth to recall that the interpretation of the X-ray observations is not yet unique: the narrow rims observed in the X-ray synchrotron emission could be due to the damping of the downstream magnetic field \\cite{pohl} rather than to severe synchrotron losses of very high energy electrons, although this interpretation has some serious shortcomings (see \\cite{morlino} for a discussion).  In this context of excitement, due to the implications of these discoveries for the origin of cosmic rays, \\cite{bell04} discussed the excitation of modes in a plasma treated in the MHD approximation and found that a new, purely growing, non-alfv\\'enic mode appears for high acceleration efficiencies. The author predicted a saturation of this SI at the level $\\delta B/B\\sim M_A (\\eta v_s/c)^{1/2}$ where $\\eta$ is the cosmic ray pressure in units of the kinetic pressure $\\rho v_s^2$, $v_s$ is the shock speed and $M_A=v_s/v_A$ is the Alfv\\'enic Mach number. For comparison, standard SI for resonant wave-particle interactions leads to expect $\\delta B/B\\sim M_A^{1/2} \\eta^{1/2}$. For efficient acceleration $\\eta\\sim 1$, and typically, for shocks in the interstellar medium, $M_A\\sim 10^4$. Therefore Bell's mode leads to $\\delta B/B\\sim 300-1000$ while the standard SI gives $\\delta B/B\\sim 30$. It is also useful to notice that the saturation level predicted by \\cite{bell04} is basically independent of the value of the background field, since $\\delta B^2/8\\pi\\sim (1/2)(v_s/c)P_{CR}$, where $P_{CR}$ is the cosmic ray pressure at the shock surface.  The resonant and non-resonant mode have different properties in other respects as well. A key feature consists in the different wavelengths that are excited. The resonant mode with the maximum growth rate has wavenumber $k$ such that $\\ko=1$, where $r_{L,0}$ is the Larmor radius of the particles that dominate the cosmic ray number density at the shock, namely, for typical spectra of astrophysical interest, the lowest energy cosmic rays at the shock. At the shock location, the minimum momentum is the injection momentum, while at larger distances from the shock, the minimum momentum is determined by the diffusion properties upstream and is higher than the injection momentum, since higher energy particles diffuse farther upstream. When the non-resonant mode exists, its maximum growth is found at $\\ko\\gg 1$. There may potentially be many implications of this difference: the particle-wave interactions which are responsible for magnetic field amplification also result in particle scattering (diffusion). The diffusion properties for resonant and non-resonant interactions are in general different. The case of resonant interactions has been studied in the literature (e.g. \\cite{lc83a}), at least for the situation $\\delta B/B\\ll 1$, but the diffusion coefficient for  non-resonant interactions (in either the linear or non-linear case) has not  been calculated (see however \\cite{zira08b}). The difference in wavelengths between the two modes, in addition to different scattering properties, also suggests that the damping will occur through different mechanisms.  The calculation of \\cite{bell04} has however raised some concerns due to the following three aspects: 1) the background plasma was treated in the MHD approximation; 2) a specific choice was made for the current established in the upstream plasma to compensate for the cosmic ray (positive) current; 3) the calculation was carried out in a reference frame at rest with the upstream plasma, where stationarity is in general not realized (although for small scale perturbations, the approximation of stationarity may be sometimes justified).  In the present paper we derive the dispersion relation of the waves in a purely kinetic approach and investigate different scenarios for the microphysics that determines the compensating current. We show that the fastly growing non resonant mode appears when particle acceleration is very efficient, but whether it dominates over the well known resonant interaction between particles and alfv\\'en modes depends on the parameters that characterize the shock front, its Mach number primarily.  \\cite{bell04} also investigated the developement of the non-resonant modes by using numerical MHD simulations. His results have been recently confirmed by \\cite{zira08a} with a similar approach. \\cite{niem08} made a first attempt to investigate the development of the non-resonant modes by using PIC simulations. In this latter case, the authors find that the non-resonant mode saturates at a much lower level than found by \\cite{bell04}. However, as briefly discussed in \\S~\\ref{sec:concl}, these simulations use a set up that makes them difficult to compare directly with Bell's results.  The paper is organized as follows: in \\S~\\ref{sec:calc} we derive the dispersion relation of the unstable modes within a kinetic approach and adopting two different scenarios for the compensation of the cosmic ray current, namely compensation due to the motion of cold electrons alone (\\S~\\ref{sec:calc1}), and to the relative drift of protons and electrons (\\S~\\ref{sec:calc2}); in \\S~\\ref{sec:modes} we discuss the relative importance of the resonant and non-resonant modes depending on the physical parameters of the system; we also derive analytic approximations for the large (\\S~\\ref{sec:largek}) and small (\\S~\\ref{sec:smallk}) wavenumber limits; finally in \\S~\\ref{sec:res} we study the different modes during the Sedov evolution of a ``typical'' supernova remnant and for different assumptions on the background magnetic field strength; we conclude in \\S~\\ref{sec:concl}.  Throughout the paper we use the expressions {\\it accelerated particles} and {\\it cosmic rays} as referring to the same concept. ",
        "conclusions": "\\label{sec:concl}  We investigated the excitation of streaming instability induced by accelerated particles in the vicinity of a non-relativistic shock wave, typical of supernova shells expanding in the interstellar medium. The calculation is based on kinetic theory, hence we do not require the MHD approximation to hold for the background plasma. We find that the dispersion relation of the waves leads to the appearance of two modes, a resonant and a non-resonant one. The former is the well known unstable mode, discussed by \\cite{zweib79} and \\cite{achter83}, based on a resonant interaction between waves and particles. The latter is similar to that discussed by \\cite{bell04}, who however based his analysis on a set of assumptions that called for further investigation: the calculation of \\cite{bell04} is based on the assumption that the background plasma can be treated in the MHD regime, and makes specific prescriptions on the return current which compensates the cosmic ray current upstream of the shock. Moreover, the whole calculation is carried out in the frame of the upstream plasma, where in principle there is no stationary solution of the problem.  Our kinetic calculations are carried out for two models of the compensating current: in the first model, the return current is established through a population of cold electrons, at rest in the shock frame, which exactly compensate the positive charge of cosmic ray protons. In the second model, the return current is due to a slight drift between ions and electrons in the background plasma upstream of the shock. We have demonstrated that the dispersion relation of the waves is the same in the two cases, to order ${\\cal O}\\left(N_{CR}/n_i\\right)^2$.  The resonant and the non-resonant mode are found at the same time, with growth rates which in the general case are different. The non-resonant mode is almost purely growing and is very apparent when particle acceleration is efficient. The parameter that regulates the appearance of the non-resonant mode is $\\sigma/v_A^2$, where $\\sigma=3\\eta v_S^3/(c R)$. When $\\sigma/v_A^2\\gg 1$, the waves excited in a non-resonant way grow faster than the resonant modes and may lead to a substantial magnetic field amplification.  The strong dependence of $\\sigma$ on the shock velocity implies that the non-resonant mode is likely to be the dominant channel of magnetic field amplification in SNRs in the free expansion phase and at the early stages of the Sedov-Taylor phase of adiabatic expansion. At later times, the non-resonant mode {\\it collapses} on the resonant mode, which keeps providing appreciable growth for longer times, at least if damping mechanisms are neglected. The growth of the fastest non-resonant mode is independent on the strength of the unperturbed initial magnetic field $B_0$.  The non-resonant mode, when present, grows the fastest at wavenumber $k_2$ given by Eq. \\ref{eq:k2}, which in the cases of interest is much larger than $1/r_{L,0}$, where $r_{L,0}$ is the gyroradius of the particles with minimum momentum in the cosmic ray spectrum. These modes are therefore short wavelength waves, which is the main reason why the assumption of stationarity in the upstream frame, as required by \\cite{bell04}, was acceptable, despite the impossibility of reaching actual stationarity in that frame.  The numerical results in this paper were specialized to the case of a power law spectrum $p^{-4}$ of accelerated particles, typical of Fermi acceleration at strong shocks. However, one should keep in mind that the levels of efficiency required for the non-resonant mode to appear are such that the dynamical reaction of the accelerated particles on the shock cannot be neglected (see \\cite{maldrury} for a review). This backreaction leads to several important effects: on one hand the spectra of accelerated particles become concave, and concentrate the bulk of the energy in the form of accelerated particles at the maximum momentum. On the other hand, the efficiently amplified magnetic field also exerts a strong dynamical reaction on the system, provided the magnetic pressure exceeds the gas pressure in the shock region (\\cite{apjlett,long}). This second effect results in an enhanced acceleration efficiency (due to large B-fields) but weaker shock modification (spectra closer to power laws) due to the reduced compressibility of the plasma in the presence of the amplified magnetic field. These non-linear effects cannot be taken into account in the type of calculations presented here, although we do not expect the qualitative character of our conclusions to be affected profoundly by them.  Nevertheless it is useful to go through the possible consequences of the non-linear effects in a somewhat deeper detail: one can expect two types of complications, one of principle and the other in the numerical values of the growth rates. The latter simply derives from the approximations intrinsic in the assumptions we made: for instance the adoption of a power law spectrum of accelerated particles with slope $4$, which clearly fails when a precursor is formed.  There are then complications coming from deeper unknown pieces of Physics or from a not totally satisfactory mathematical approach. For instance, the standard perturbative approach adopted here is based on the assumption of a spatially uniform background. The presence of a precursor invalidates this assumption, although probably not in a dramatic way.  Since the non-resonant mode appears for large values of $k$, the relevant quantities can be assumed to be spatially constant in the precursor on scales $\\sim 1/k$, so that in this respect our calculations are still expected to hold, and probably to a better accuracy for the non-resonant modes ($k \\gg 1/r_{L,0}$) than for the resonant ones ($k \\sim 1/r_{L,0}$). Moreover, as stressed above, the dynamical reaction of the magnetic field leads to weaker modification of the shock, and therefore to spectra with less prominent concavity (closer to $p^{-4}$). Also in this respect, the calculations presented here should serve as a good description of all relevant physical effects related to the growth of the cosmic ray induced instabilities.  More important, the acceleration process is directly affected by the physics of particles' diffusion in the shock region, which in turn is determined by the excited waves. This intrinsic non-linearity cannot be taken into account in perturbative approaches like ours or like Bell's, and one should always be aware of this limitation. Even more, while the diffusion coefficient for resonant modes can at least be derived in quasi-linear theory, at present there is no derivation of the diffusion coefficient associated with scattering on non-resonant modes (see \\cite{zira08b} for a first attempt at discussing this effect).  Another issue that deserves further investigation is that of determining the level of field amplification at which the instability saturates. This cannot be worked out within a linear theory calculation and only numerical simulations can address this issue. Recent efforts in this direction have been made by \\cite{bell04} and \\cite{zira08a} through MHD simulations and by \\cite{niem08} by using PIC simulations. While the first two papers find a saturation level $\\delta B^2/(4 \\pi) \\sim (v_s/c) P_c$, in the third paper a much lower level of field amplification is found. The authors conclude that the existence of large magnetic field amplification through the excitation of non-resonant modes is yet to be established.  Although we agree with this conclusion, we also think that the setup of the PIC simulation by \\cite{niem08} is hardly applicable to investigate the excitation of the Bell instability at shocks, or at least several aspects of it should be studied more carefully. First, they carried out the calculations in a regime in which the condition of strong magnetization, $\\omega\\ll\\Omega_0$, was violated. Second, in order to carry out the calculations, \\cite{niem08} are forced to assume unrealistically large values for the ratio $N_{CR}/n_i$ (of order 0.3 for their most realistic runs). The return current as assumed by \\cite{niem08} corresponds to our second model, which however leads to the same dispersion relation as \\cite{bell04} only at order ${\\cal O}\\left(N_{CR}/n_i\\right)^2$, which is not necessarily the case here. Moreover, the spectrum of accelerated particles is assumed to be a delta-function at Lorentz factor 2, instead of a power law (or more generally a broad) spectrum. It is not obvious that for non-resonant modes this assumption is reasonable. But the most serious limitation of this PIC simulation is in the fact that the authors do not provide a continuous replenishment of the cosmic ray current, which is instead depleted because of the coupling with waves. In the authors' view this seems to be a positive aspect of their calculations, missed by other approaches, but in actuality the cosmic ray current is indeed expected to be stationary upstream, and we think that the PIC simulation would show this too if particles were allowed to be accelerated in the simulation box instead of being only advected and excite waves. Clearly if this were done, the spectrum of accelerated particles would not keep its delta-function shape, but should rather turn into a power-law-like spectrum. The latter issue adds to the absence of a replenishment of the current, which seems to us to be the main shortcoming of these simulations. Overall, it appears that the setup adopted by \\cite{niem08} would apply more easily to the propagation of cosmic rays rather than to particle acceleration in the vicinity of a shock front.  The issue of efficient magnetic field amplification, possibly induced by cosmic rays, has become a subject of very active debate after the recent evidence of large magnetic fields in several shell-type SNRs. The implications of such fields for particle acceleration to the knee region, as well as for the explanation of the multifrequency observations of SNRs, are being investigated. Probably the main source of uncertainty in addressing these issues is the role of damping of the excited waves. For resonant modes, ion-neutral damping and non-linear Landau damping have been studied in some detail: their role depends on the temperature of the upstream plasma and on the shock velocity. For non-resonant modes, being at high $k$, other damping channels could be important (Everett et al., in preparation). Whether SNRs can be the source of galactic cosmic rays depends in a complex way on the interplay between magnetic field amplification, damping, particle scattering and acceleration, together with the evolution of the remnant itself."
    },
    "0806/0806.0319_arXiv.txt": {
        "abstract": "Angle-action coordinates are used to study the relic of an $N$-body simulation of a self-gravitating satellite galaxy that was released on a short-period orbit within the disc of the Galaxy. Satellite stars that lie within $1.5\\kpc$ of the Sun are confined to a grid of patches in action space. As the relic phase-mixes for longer, the patches become smaller and more numerous. These patches can be seen even when the angle-action coordinates of an erroneous Galactic potential are used, but using the wrong potential displaces them. Diagnostic quantities constructed from the angle coordinates both allow the true potential to be identified, and the relic to be dated. Hence when the full phase space coordinates of large numbers of solar-neighbourhood stars are known, it should be possible to identify members of particular relics from the distribution of stars in an approximate action space. This would then open up the possibility of determining the time since the relic was disrupted and gaining better knowledge of the Galactic potential.  The availability of angle-action coordinates for arbitrary potentials is the key to these developments.  The paper includes a brief introduction to the torus technique used to generate them. ",
        "introduction": "\\label{sec:intro}  Within the remarkably successful $\\Lambda$CDM model, galaxy formation is a hierarchical process. Large galaxies, such as the Milky Way, are built up in mergers and the accretion of smaller building blocks \\citep[e.g.][]{WhiteRees1978,SpringelHernquist2003}. The signatures of these processes should still be visible today in the form of substructure such as streams in all components of the Milky Way \\citep[e.g.][]{LyndenBell1995,FreemanBH2002,Helmietal2003,Abadietal2003}.  Evidence of the wealth of substructure in the stellar halo of the Milky Way has increased dramatically over the past 15 years, most notably with observations of the disrupting Sagittarius dwarf galaxy \\citep{IbataGilmoreIrwin1994}, and of the streams visible in the SDSS data \\citep[e.g.][]{FieldofStreams}. Within the disc several substructures are known, as are several mechanisms that might be responsible for them. Stars are born in clusters within the disc, and over a period of several Galactic rotations these clusters evaporate and the stars become phase mixed, spreading in space but retaining closely related orbits. This is commonly referred to as a supercluster -- it is likely the Hyades-Pleiades supercluster formed this way \\citep{Famaeyetal2005}. Substructure can also be created by dynamical interaction with spiral arms, or a rotating bar component; for example the Hercules stream is thought to be associated with the bar of the Milky Way \\citep{Dehnen1999:Bar,Dehnen2000:OLR,Fux2001}.  It has been suggested that the Arcturus group \\citep{Eggen1971} is debris from a merged satellite \\citep{NavarroHelmiFreeman2004}.  Many methods exist for finding this substructure, often based on incomplete phase space information about the stars. In the outer parts of the halo it is possible to find substructure with only knowledge of stellar positions on the sky, sometimes in conjunction with photometric data \\citep[e.g.][]{FieldofStreams}, or radial velocity measurements \\citep[e.g.][]{LyndenBell1995,HelmiWhite1999}. In the solar neighbourhood, approaches that look for common proper motions have been widely used, \\citep[e.g.][]{Chereuletal1999}.  With the availability of full 6D phase space information for an increasing number of stars in the solar neighbourhood, most notably the catalogues resulting from the Geneva-Copenhagen and RAVE surveys \\citep{Nordstrometal2004,Steinmetz2006,Zwitter2008}, and with the prospect of a further increase by several orders of magnitude when Gaia data become available \\citep{GAIA2001}, it is appropriate to consider methods for using these data in full.  As discussed in \\cite{Helmietal1999}, the space of integrals of motion is a very promising one for finding substructure such as superclusters or merger debris. Stars with a single small progenitor (star cluster or small galaxy) will have very similar values for the integrals, which will ensure they are tightly bunched in integral space even after phase mixing has produced a spatial distribution that is effectively featureless. There are additional benefits to using quantities that are not only  integrals but adiabatic invariants, as these are more likely to remain constant as the Galactic potential changes over time.  Previous work has focused on spaces defined by $L_z$ (which is an adiabatic invariant in an axisymmetric potential), and other quantities that can be used as approximate integrals of the motion, such as the total angular momentum \\citep[only an integral of the motion in a spherically symmetric potential,] []{Helmietal1999}, or are not adiabatic invariants, such as the energy \\citep{HelmideZeeuw2000} or the apocentre and pericentre of an orbit \\citep{Helmietal2006}.  In this paper we demonstrate the use of angle-action coordinates to find substructure in the solar neighbourhood.  Actions are adiabatic invariants, and their conjugate variables, the ``angle coordinates'', increase linearly in time.  As \\cite{Tremaine99} has pointed out, these properties make them exceptionally useful for analysing tidal streams. The difficulty of determining the actions of stars in non-spherical potentials has, however, severely restricted their application to date. This is the first in a series of papers in which we show how the concept of orbital tori \\citep{McGB90,KaasB94} makes it possible to exploit the power of angle-action coordinates for practical galactic problems.  In Sections~\\ref{sec:actang}~\\&~\\ref{sec:fit} we briefly introduce angle-action coordinates, and explain how we find them for stars with known phase space positions. In Section~\\ref{sec:sim} we apply them to simulated data of a satellite merger. We show that stars from the satellite that are observed in the solar neighbourhood are confined to a grid of patches in action space, even when rather a poor approximation to the Galactic potential is used. We show further that a diagnostic defined in terms of the angle coordinates enables the true potential to be distinguished from the false one. Section~\\ref{sec:discuss:secular} discusses minor modifications to the analysis that are required to accommodate secular evolution of the Galactic potential over the life of a relic. ",
        "conclusions": "Stars have continued to form at a significant rate throughout the lifetime of the Galaxy's thin disc, and it must be presumed that the disc's mass has increased significantly over the last $5\\Gyr$. Since actions are adiabatic invariants, such secular evolution of the Galactic potential does not affect the distribution of relic stars in the $J_\\phi-J_R$ plane (Fig.~\\ref{fig:randJ}). Secular evolution causes the frequencies at fixed $\\bolJ$ to become explicit functions of time, so we should write $\\bolom(\\bolJ,t)$, and the increment in $\\theta_i$ over time $t$ changes from $\\Delta\\theta_i=\\bolom_it$ to $\\Delta\\theta_i=\\int_0^tdt\\,\\bolom_i$. The conditions for a relic star to be in the solar neighbourhood are given values of $\\Delta\\theta_i\\mod2\\pi$ for $i=R,\\phi$.  Since with secular evolution $\\Delta\\theta_i$ remains a continuous function of $\\bolJ$, these conditions continue to be satisfied only in a grid of patches in action space; secular evolution shifts the patches, but does not blur them. Hence the diagnostic power of plots like Fig.~\\ref{fig:Jsol} is unaffected by secular evolution.  In the presence of secular evolution it becomes necessary in equation (\\ref{eq:thetalapse}) to replace $\\bolom_\\alpha(t-t_0)$ by $\\int_{t_0}^tdt\\,\\bolom_\\alpha$. To evaluate the required integrals, one must adopt a model of the history of the potential, which determines the time dependence of $\\bolom$.  The required model is self-evident if secular evolution is confined to growth in the disc's mass at a known rate. Uncertainties in this rate will make it harder to locate beats in Fig.~\\ref{fig:stat}.  Of course the key to calculating the secular evolution of stellar systems, be they globular clusters or galaxies, is to express their distribution functions in terms of actions \\citep[e.g.][]{SellwoodMcgaugh05}, so the availability of orbital tori for arbitrary potentials opens up new horizons in this area.    A major hope of ``near-field cosmology'' is to identify within the Galaxy groups of stars that were accreted together \\citep{FreemanBH2002}. We have demonstrated the power of angle-action coordinates for doing this by studying the debris of a self-gravitating satellite of mass $3.75\\times10^8\\msun$, released within the plane of a realistic Galactic potential on an orbit with apocentre at $9\\kpc$. On this short-period orbit the satellite's stars become well phase mixed within a couple of gigayears, and they are quite widely distributed in action space.  Nonetheless, the stars that lie within $1.5\\kpc$ of the Sun are concentrated into a grid of patches in action space because only stars with certain frequencies are currently near the Sun.  To see the patchiness of the distribution in action space it is not necessary to use the angle-action coordinates of the true Galactic potential. But the correct potential must be used if statistical measures constructed from the angle coordinates of stars are to show a characteristic pattern of beats from which the time at which  the relic was disrupted can be deduced. Hence our results suggest a two-stage procedure: first a reasonable approximation to the Galactic potential is used to identify relics through the clustering of their stars' points around the nodes of a grid in action space. Then once a relic has been identified, the Galactic potential would be adjusted until the angle-variable diagnostics showed pronounced beats. This second step would not only pin down the Galaxy's potential, but also reveal the time at which  relic was disrupted.  Growth in the mass of the disc since the satellite fell in would have significant effects only on Fig.~\\ref{fig:stat}: to recover this plot it would be necessary to model the time dependence of the Galactic potential, so that the integrals $\\int dt\\,\\bolom$ could be evaluated. We anticipate that with the help of angle-action coordinates this could be done to sufficient accuracy, but defer this refinement to a subsequent publication.  We have neglected the deviations of the Galactic potential from axisymmetry. Could these deviations have a significant impact? \\cite{Juric2008} use star counts in the SDSS survey to show that the Galaxy's thick disc is remarkably axisymmetric near the Sun. This finding suggests that it is legitimate to neglect the bar when searching for relics within the thick disc, such as the Arcturus group. In general, the quadrupole moments of the bar's gravitational potential will decline rapidly outside the end of the bar at $R\\sim3\\kpc$, so stars that are not resonant with the bar will not be strongly affected by it. The observed axisymmetry of the thick disc suggests that few if any of its stars are resonantly trapped by the bar, so their orbits can be safely modelled with an axisymmetric potential. The question of how the -- phase-dependent -- effects of the bar would impact Fig.~\\ref{fig:stat} may prove important.  The exciting possibilities discussed here rest on two foundations. One is the availability of angle-action coordinates for any given potential, and the other is the availability of full phase space coordinates for significant samples of stars.  The torus construction technique developed in a series of paper starting with \\cite{McGB90} can provide angle-action coordinates, and programmes such as the Geneva-Copenhagen, RAVE and Gaia surveys will provide the phase space coordinates.  Currently the torus technique is restricted to either axisymmetric systems or two-dimensional non-rotating bars. However, extension to three-dimensional bars, including bars that are rotating with a constant pattern speed, is in principle straightforward and will be attempted soon.  Clearly when this technique is used to search a real catalogue for relics, and then to analyse them, difficulties will be encountered that we have ignored here. Most obviously one will have to contend with errors in the phase space coordinates of stars (primarily due to errors in distances) and with the difficulty in picking out overdensities in action space against a background of Poisson noise from field stars. We are currently applying the method to $\\sim200\\,000$ stars from the RAVE survey and hope to report the results in the near future."
    },
    "0806/0806.2400_arXiv.txt": {
        "abstract": "Borexino, a large volume detector for low energy neutrino spectroscopy, is currently running underground at the Laboratori Nazionali del Gran Sasso, Italy. The main goal of the experiment is the real-time measurement of sub MeV solar neutrinos, and particularly of the mono energetic (862 keV) $^{7}Be$ electron capture neutrinos, via neutrino-electron scattering in an ultra-pure liquid scintillator. This paper is mostly devoted to the description of the detector structure, the photomultipliers,  the electronics, and the trigger and calibration systems. The real performance of the detector, which always meets, and sometimes exceeds, design expectations, is also shown. Some important aspects of the Borexino project, i.e. the fluid handling plants, the purification techniques and the filling procedures, are not covered in this paper and are, or will be, published elsewhere (see Introduction and Bibliography). ",
        "introduction": "\\label{sec:Intro}  Borexino is a large volume liquid scintillator detector whose primary purpose is the real-time measurement of low energy solar neutrinos. It is located deep underground ($\\simeq$ 3800~ meters of water equivalent, m w.e.) in the Hall C of the Laboratori Nazionali del Gran Sasso (Italy), where the muon flux is suppressed by a factor of $\\approx 10^6$.  The main goal of the experiment is the detection of the monochromatic neutrinos that are emitted in the electron capture decay of $^7Be$ in the Sun \\cite{bib:Borex1}. This measurement is now in progress, and the very first results have been already published in \\cite{bib:be7paper}. However, as shown there, the observed radioactive background is much lower than expected, which results in a potential broadening of the scientific scope of the experiment. Particularly, Borexino now also aims at the spectral study of other solar neutrino components, such as the CNO, pep \\cite{bib:c11}  and, possibly, pp and $^{8}$B neutrinos.  Besides solar physics, the unprecedented characteristics of its apparatus make Borexino very competitive in the detection of anti-neutrinos ($\\bar{\\nu}$),  particularly those of geophysical origin. The physics goals of the experiment also include the detection of a nearby supernova, the measurement of the neutrino magnetic moment by means of a powerful neutrino source, and the search for very rare events like the electron decay \\cite{bib:edecay} or the nucleon decay into invisible channels \\cite{bib:nucleon}.  In Borexino low energy neutrinos ($\\nu$)  of all flavors are detected by means of their elastic scattering of electrons or, in the case of electron anti-neutrinos, by means of their inverse beta decay on protons or carbon nuclei. The electron (positron) recoil energy is converted into scintillation light which is then collected by a set of photomultipliers.  This technique has several advantages over both the water \\che\\ detectors and the radiochemical detectors used so far in solar neutrino experiments. Water \\che\\ detectors, in fact, can not effectively detect solar neutrinos whose energy is below 6 MeV, both because the \\che\\ light yield is low and because the intrinsic radioactive background cannot be pushed down to sufficiently low levels. On the other hand, radiochemical experiments cannot intrinsically perform spectral measurements and do not detect events in real time.  An organic liquid scintillator solves the aforementioned problems: the low energy neutrino detection is possible because of the high light yield, that in principle allows the energy threshold to be set down to a level of a few tens of keV\\footnote{However, the unavoidable contamination of $^{14}$C that is present in any organic liquid practically limits the \"neutrino window\" above $\\approx$ 200 keV}; the organic nature of the scintillator, and its liquid form at ambient temperature, provide very low solubility of ions and metal impurities, and yield the technical possibility to purify the material as required. However, no measurement of the direction of the incoming neutrino is possible and, even more importantly, the neutrino induced events are intrinsically indistinguishable from $\\beta$ and $\\gamma$ radioactivity, posing formidable requirements in terms of radiopurity of the scintillator and of the detector materials.  According to the Standard Solar Model\\footnote{Regardless of neutrino oscillations which are not relevant at this point}, the order of magnitude of sub-MeV solar neutrino interactions is a few tens counts/day for about one hundred tons of target material, corresponding to an equivalent activity of a few $\\cdot ~10^{-9}$ Bq/kg. If one compares this low number with the typical radioactivity of materials (drinking water $\\simeq$ 10 Bq/kg, air $\\simeq$ 10 Bq/kg, rock $\\simeq$ 100-1000 Bq/kg) it is immediately apparent that the core of the Borexino detector must be 9-10 orders of magnitude less radioactive than anything on Earth. Typical radioactive contaminants in solid materials and water are $^{238}$U and $^{232}$Th daughters, and $^{40}$K. Air and therefore normally also commercially available nitrogen are typically contaminated by noble gases like $^{222}$Rn, $^{39}$Ar and $^{85}$Kr.  The necessity to measure such a low neutrino flux with a massive detector poses severe requirements in terms of radiopurity, not only for the scintillator itself, but also for the surrounding materials. Additionally, the neutrino target (100 t of ``fiducial volume`` in Borexino) must be almost completely shielded from external $\\gamma$ radiation and neutrons originating from the rock and from the detector materials.  For almost 20 years the Borexino collaboration has been addressing this problem by developing suitable purification techniques for scintillator, water, and nitrogen, by performing careful material selections, by developing innovative cleaning techniques for metal surfaces, and by building and operating a prototype of the Borexino detector, the Counting Test Facility (CTF). In particular, CTF has played a crucial role in this long R\\&D phase. It is still the only instrument available in the world (except Borexino itself) with the sensitivity to measure the radioactive contamination of a liquid scintillator down to levels as low as $10^{-16}$ g/g in $^{238}$U and $^{232}$Th\\footnote{Here and everywhere in this paper the unit g/g stands for 1 gram of contamination per gram of solution or material}. For more details about the specific requirements in terms of radiopurity of the scintillator and of the detector materials for solar neutrino measurement in Borexino see \\cite{bib:Borex1} and \\cite{bib:Borex2}. For the reader's convenience, we summarize here the main requirements:  \\begin{itemize}  \\item The internal radioactivity of the scintillator must be low enough compared to the expected neutrino signal. Particularly, the design goal was $<10^{-16}$ g/g in $^{238}$U and $^{232}$Th, $<10^{-14}$ g/g in K$_{nat}$\\footnote{K$_{nat}$ is potassium in its natural isotopic abundance.}.  \\item The scintillator must be thoroughly sparged with nitrogen gas in order to remove oxygen (which may deteriorate the optical properties of the scintillator) and air borne contaminants (radioactive). The nitrogen purity requirement is such that the expected background from $^{222}$Rn, $^{39}$Ar and $^{85}$Kr in 100 t of target scintillator must be less than 1 count/day. This corresponds to 0.36 ppm for Ar and 0.16 ppt for Kr.  \\item The total amount of external $\\gamma$ radiation penetrating the central part of the scintillation volume should be below 1 count/day in 100 t. This puts stringent requirements on all materials surrounding the detector, the requirements being more and more stringent for materials closer to the center.  \\end{itemize}  This paper is devoted to the description of the Borexino detector. It is not intended to be a complete reference of the Borexino scientific goals, nor will it provide a comprehensive description of the experiment as a whole. The focus here is the detector, defined as the collection of scintillator volume, containment vessels, light detection devices (photomultipliers and electronics), data acquisition, and calibration systems. We do not cover here the purification plants (a very large fraction of the Borexino equipment) nor the purification techniques adopted to purify scintillator, water and nitrogen. Also, the filling procedures are not covered in this paper. All these very important parts of the experiment are either already published or will be published in the near future.  The paper is structured as follows: section \\ref{sec:Borex} gives a general description of the detector; section \\ref{sec:scintillator} summarizes the main scintillator features; section \\ref{sec:vessels} describes the Inner Nylon Vessels which contain the scintillator and act as ultimate barriers against external contaminations; section \\ref{sec:InnerDet} describes the main detector with its photomultipliers, front end electronics, and data acquisition electronics; section \\ref{sec:OuterDetector} describes the muon detector; sections \\ref{sec:Trigger} and \\ref{sec:Daq} describe the trigger and the data acquisition systems; sections \\ref{sec:CalibHardware} and \\ref{sec:transp} describe the laser based calibration systems for the photomultipliers and for the monitoring of the scintillator transparency; section \\ref{sec:insertion} describes the insertion system for source calibrations.  Finally, the last section provides a brief overview of the detector performance on real data. For more details about detector performance see refs. \\cite{bib:be7paper} and \\cite{bib:neutrinopaper}. ",
        "conclusions": "\\label{sec:Conclu} The construction and the commissioning of the Borexino detector is completed. Data taking has begun on May 15th, 2007 and is going to continue for several years. This paper shows that the detector meets, or in some cases  exceeds, the expected performance. The radioactive background is lower than the design values for several contaminants, particularly for the $^{238}$U and $^{232}$Th daughters. The PMTs of both inner and outer detectors, the electronics, and the trigger system work as expected."
    },
    "0806/0806.0775_arXiv.txt": {
        "abstract": "{} {The disagreement between helioseismology and a recent downward revision of solar abundances has resulted in a controversy about the true neon abundance of the Sun and other stars. We study the coronal Ne/O abundance ratios of nearby stars with modest activity levels and investigate a possible peculiarity of the Sun among the stellar population in the solar neighborhood.} {We used XMM-Newton and Chandra data from a sample of weakly and moderately active stars with log\\,L$_{\\rm X}$/L$_{\\rm bol} \\approx -5\\,...-7 $ to investigate high-resolution X-ray spectra to determine their coronal Ne/O abundance ratio. We applied two linear combinations of strong emission lines from neon and oxygen, as well as  a global-fitting method for each dataset, and crosschecked the derived results.} {The sample stars show a correlation between their Ne/O ratio and stellar activity in the sense that stars with a higher activity level show a higher Ne/O ratio. We find that the Ne/O abundance ratio decreases in our sample from values of Ne/O\\,$\\approx$\\,0.4 down to Ne/O\\,$\\approx$\\,0.2\\,--\\,0.25, suggesting that ratios similar to 'classical' solar values, i.e. Ne/O\\,$\\approx0.2$, are rather common for low activity stars. A significantly enhanced neon abundance as the solution to the solar modeling problem seems unlikely.} {From the coronal Ne/O abundance ratios, we find no indications of a peculiar position of the Sun among other stars. The solar behavior appears to be rather typical of low activity stars.} ",
        "introduction": "The chemical composition of the Sun is used as a reference frame throughout astronomy and therefore the precise determination of the solar abundances is of particular interest. Beside indirect methods utilizing solar wind particles and meteorites, the measurement of photospheric lines is used as the prime method to derive the chemical composition of the Sun. Making use of more and more refined measurements and modeling, the chemical composition of the Sun 'evolved' throughout the decades, however abundance compilations like the ones of \\cite{angr} and more recently by \\cite{grsa} have become widely accepted. \\cite{asplund} proposed the most recent revision of the solar abundance scale based on a new 3\\,D hydrodynamic modeling of the solar atmosphere, resulting in approximately 30\\% lower abundances of many light elements including in particular C, N, O and Ne. While this new solar chemical composition provides a better agreement with e.g. measurements of the local ISM, on the solar side severe discrepancies appeared \\citep{basu04, tur04,bah05}. The light elements are an important source of opacity in the Sun and as a consequence the almost perfect agreement between helioseismology and models of the solar interior based on the abundance values of \\cite{grsa} is seriously disturbed. The disagreement is far beyond measurement errors and estimated modeling uncertainties and therefore a solution of this problem is highly desired.  One proposal to reconcile the new abundances with helioseismology suggests an enhancement of the solar neon abundance \\citep{ant05}. In contrast to C, N, O, Ne has no strong photospheric lines and its abundance is determined from emission lines originating in the transition region and the corona or through composition measurements of solar wind particles. However, both methods do not necessarily trace the true solar photospheric abundances, since chemical fractionation processes in the outer atmospheric layers of stars may be present. For the Sun the FIP (First Ionization Potential) effect is known, leading to an overabundance of low FIP ($\\lesssim$~10\\,eV) elements like Fe or Mg, whereas the high FIP elements Ne and O (that is used as reference element) are thought to remain around photospheric values \\citep[see e.g.][]{gei98}.  While the 'classical' solar Ne/O ratio from \\cite{grsa} is Ne/O=0.18, the overall lower \\cite{asplund} abundances result in a comparable value of Ne/O=0.15. However, the \\cite{asplund} abundances require a significant increase of the solar neon abundance by, depending on the assumed model and abundance uncertainties, a factor of $\\sim$\\,2.5 up to $\\sim$\\,4.0 \\citep{ant05, bah05b}, i.e. Ne/O\\,$\\approx$\\,0.4\\,--\\,0.6, could provide the missing opacity and would reconcile helioseismology with the new abundances derived by \\cite{asplund}. Evidence for an increased neon abundance was presented by \\cite{dra05}, who measured coronal Ne/O ratios in X-ray spectra of various nearby stars, found a value of Ne/O=0.41 and 'solved' the problem by assuming a similar ratio for the Sun, thought noting that solar measurements indicate the lower values, i.e. Ne/O=0.15/0.18 to be correct. Actually their new neon abundance is still slightly below the required enhancement, but assuming other elemental abundances to additionally increase jointly within allowed errors would provide a sufficient opacity increase. Several objections to this solution were raised, especially from solar observers. A reassessment of solar coronal data from the {\\it Solar Maximum Mission} led to an upper limit of Ne/O=0.18$\\pm0.04$ for active regions \\citep{schmelz05} and an analysis of transition region lines observed by the {\\it SOHO} satellite suggested Ne/O=0.17$\\pm0.05$ for the quiet Sun \\citep{you05}, both confirming the 'classical'  solar Ne/O ratio. Although sometimes higher Ne/O values were determined from X-ray spectra of energetic flares or in $\\gamma$-ray production regions \\citep[see e.g. Table 1, sup. data of][]{dra05}, solar measurements are overall consistent with a ratio of Ne/O$\\approx 0.2\\pm$0.05, as also noticed in \\cite{dra05}. Furthermore, the stellar sample used by \\cite{dra05} for their abundance analysis contains mostly active stars and binary systems, which are not comparable to the Sun in terms of activity and are thought to exhibit a different chemical fractionation process. This so-called inverse FIP effect \\citep[see][for overall X-ray properties of stellar coronae]{gue04} is commonly observed in X-ray spectra of active stars. It is possible that the inverse FIP effect modifies photospheric Ne/O ratios, leading to higher coronal values than in the photosphere. Hence it is necessary to determine the Ne/O ratio in relatively inactive stars that are comparable to the Sun. Since low activity stars are intrinsically X-ray faint, only a few nearby stars have been investigated and their Ne/O ratios were mostly found to be below the \\cite{dra05} value but higher than the 'classical' solar one, although measurement errors for individual stars were often quite large.  Two different methods have been used to investigate stellar coronal abundance ratios in general; on one hand, by global spectral modeling that usually includes the determination of key elemental abundances and some kind of emission measure distribution (EMD), on the other hand, by using individual or a combination of emission lines that have a similar emissivity vs. temperature distributions and thus the determined ratio is almost independent of the stellar EMD and the temperature dependence cancels out. This latter method requires adequate and sufficiently strong lines for both elements in question, here Ne and O. Already \\cite{act75} pointed out that the emissivities of the \\ion{O}{viii}~Ly$\\alpha$ and the \\ion{Ne}{ix} resonance line have a similar temperature dependence and are therefore suitable to determine the solar Ne/O ratio, which they determined to Ne/O=0.21$\\pm0.07$. This method has been refined by using linear combinations of suited lines to minimize temperature dependent residuals, specifically \\cite{dra05} included the \\ion{Ne}{x}~Ly$\\alpha$ line and \\cite{lie06} also the \\ion{O}{vii} resonance line. In general, inclusion of more lines, especially from different ionization stages, should minimize the temperature dependence, but adds measurement errors and restricts the method to detectors that spectrally resolve and cover all involved lines with sufficient sensitivity.  In this work we present a study of low to medium activity stars with log\\,L$_{\\rm X}$/L$_ {\\rm bol}\\lesssim -5$ and spectral types mid-F to mid-K, for which high resolution X-ray spectra are available. All stars in our sample are main-sequence stars or only moderately evolved, i.e. they belong to luminosity classes IV\\,--\\,V. We specifically determine coronal Ne/O abundance ratios for these stars by using two linear combinations of Ne and O lines as well as global spectral fits and check the results obtained from the different methods among each other and against  available literature values. We further compare our results to those obtained for more active stars and investigate a possible abundance peculiarity of the Sun.  Our paper is structured as follows. In Sect.\\,\\ref{ana} we describe the applied methods and data used, in Sect.\\,\\ref{ress} we present our results, check the applied methods and discuss implication of our findings in the context of solar and stellar coronal abundances and finally in Sect.\\,\\ref{con} we summarize our findings and conclusions. ",
        "conclusions": "\\label{con}  \\begin{enumerate} \\item We have determined the coronal Ne/O abundance ratio in a sample of low and moderately active stars. For stars in the activity range log\\,L$_{\\rm X}$/L$_{\\rm bol}\\approx -5\\,...-7$ we find a trend of decreasing Ne/O ratio with decreasing activity level. The Ne/O ratio decreases from values of Ne/O\\,$\\approx$\\,0.4 for moderately active stars down to Ne/O\\,$\\approx$\\,0.2\\,--\\,0.25 for low activity stars. The decrease is measured independently of the chosen analysis method, i.e. by global spectral modeling and two linear combinations of strong emission lines.  \\item The Sun as a low activity star fits very well into the picture derived from the stellar X-ray data. A low Ne/O ratio around the  'classical' solar value of Ne/O=0.18 might be rather typical for solar-like stars with comparable activity level. A significant higher neon abundance, i.e. a ratio of Ne/O=0.4\\,--\\,0.6, as proposed to solve the solar interior problem can be virtually ruled out for the coronal composition. It also appears unlikely for the photospheric composition given the present knowledge and observational data on chemical fractionation processes in weakly active stars.  \\item We consider the main findings of this study to be robust, however the presently available data is insufficient to determine the detailed characteristics of the decline. While a linear dependence between Ne/O ratio and activity level describes our data well, some kind of flattening or saturation is expected for very active stars. Possible dependences of abundance fractionation on other stellar parameters beside activity also remain unclear. However, further X-ray observations of nearby low activity stars may successfully address these problems and tighten their re-unification with the Sun.   \\end{enumerate}"
    },
    "0806/0806.0296_arXiv.txt": {
        "abstract": "An approach to the equation of state for the inner crust of neutron stars based on Skyrme-type forces is presented. Working within the Wigner-Seitz picture, the energy is calculated by the TETF (temperature-dependent extended Thomas-Fermi) method, with proton shell corrections added self-consistently by the Strutinsky-integral method. Using a Skyrme force that has been fitted to both neutron matter and to essentially all the nuclear mass data, we find strong proton shell effects: proton numbers $Z$ = 50, 40 and 20 are the only values possible in the inner crust, assuming that nuclear equilibrium is maintained in the cooling neutron star right down to the ambient temperature.  Convergence problems with the TETF expansion for the entropy, and our way of handling them, are discussed. Full TETF expressions for the specific heat of inhomogeneous nuclear matter are presented. Our treatment of the electron gas, including its specific heat, is essentially exact, and is described in detail. ",
        "introduction": "\\renewcommand{\\theequation}{1.\\arabic{equation}} \\setcounter{equation}{0} \\label{intro}  We are concerned with the application of Skyrme-type effective nuclear interactions to the determination of the equation of state (EOS) of the inhomogeneous nuclear matter encountered at nuclear and subnuclear densities in core-collapse supernovas and in the inner crust of neutron stars. In our first paper on this topic \\cite{opp97} we adopted a Wigner-Seitz (WS) model of the inhomogeneous nuclear medium, and used the fourth-order semi-classical temperature-dependent extended Thomas-Fermi (TETF) method to calculate the kinetic energy and entropy. That paper dealt primarily with the conditions prevailing in core-collapse supernovas. The present paper relates rather to the inner crust of neutron stars, describing in particular some modifications to the earlier model, made necessary by two problems that emerge at the much lower temperatures $T$ that are involved: i) the TETF expansion (in powers of $\\hbar^2$) for the entropy converges badly at low $T$; ii) proton shell effects are not negligible at low $T$.  This last point is especially important if one is interested in the neutron-star crust as a possible alternative site for the synthesis of the so-called r-process elements \\cite{latt77,mey89,frei99,gor05}. The usual model of the r-process of nucleosynthesis is associated with the {\\it birth} of a neutron star in a core-collapse supernova, during which ``seed\" nuclei are exposed to an intense flux of neutrons. Rapid (``$r$\") capture of neutrons alternating with beta decay leads to the formation of a string of highly neutron-rich isotopes of a wide range of elements, which, once the source of neutrons is removed, will beta-decay back to the most neutron-rich stable isobar for the given mass number $A$ (see Ref. \\cite{agt07} for a recent review). The alternative picture, of interest here, is associated rather with the {\\it death} of a neutron star, or at least with its partial disruption. Because of the very large densities, the matter in a neutron star is highly neutron rich, and the closer to the center the more neutron-rich it will be. But if for one reason or another matter is ejected from the neutron star it will rapidly decompress, and so will be able to undergo a chain of beta decays, the end-product of which will again be r-process nuclei. Ejection of matter from a neutron star is usually supposed to result from the merger of one neutron star with another, or with a black hole \\cite{frei99,rj01}, but other scenarios have been envisaged, e.g., volcanoes \\cite{dys69}, magnetars \\cite{hl06}, quark stars \\cite{jai07} and explosions resulting from the mass of the neutron star falling below the minimal critical value \\cite{sum98}. However, the precise ejection mechanism is of no concern to us in this paper.  It is convenient at this point to recall that at least three distinct regions can be recognized in a neutron star: a central, locally homogeneous, core, and two concentric shells characterized by different inhomogeneous phases \\cite{pr95}. The outermost of these shells, the ``outer crust\", consists of an electrically neutral lattice of nuclei and electrons. At the surface of the star only nuclei that are stable under natural terrestrial conditions are found (in fact, nuclear equilbrium, discussed below, implies that  only $^{56}$Fe will be found), but on moving towards the interior the increasing density leads to the appearance of nuclei that are more and more neutron rich, until at a mean local density $\\bar{\\rho}$ of around 2.4 $\\times 10^{-4}$ nucleons.fm$^{-3}$ (4.0 $\\times 10^{11}$ g.cm$^{-3})$ neutron drip sets in. This marks the transition to the ``inner crust\", which at least up to a mean density of $\\bar{\\rho}$ = 0.06 nucleons.fm$^{-3}$ consists of neutron-proton clusters, or droplets, immersed in a neutron gas, with the neutralizing electron gas being essentially uniform (we neglect screening effects in this paper). It is equally well established that by the point where the mean density has risen to around $\\bar{\\rho}$ = 0.10 nucleons.fm$^{-3}$, i.e., about 2/3 of the density $\\rho_0$ of symmetric infinite nuclear matter (INM) at equilibrium, the droplet phase no longer exists and has been replaced by the homogeneous phase of the core, which consists primarily of neutrons, with a small admixture of proton-electron pairs, and possibly other particles, including free quarks, closer to the center.  What happens in the transition region over the range 0.06 $\\le \\bar{\\rho} \\le $ 0.10 nucleons.fm$^{-3}$, close to the inside edge of the inner crust, is far less clear. The question cannot be settled by observation at the present time, and theoretical predictions are sensitive to the details of the calculations, in particular to the choice of the effective interaction. For some interactions the transition from the droplet phase to the homogeneous phase is indirect and complex, with a whole sequence of different inhomogeneous phases being formed. At the interface with the homogeneous core these calculations find a ``bubble\" phase, this taking the form of bubbles of neutron gas in a denser liquid of neutrons and protons, the droplet phase having effectively been turned inside out. Furthermore, at slightly lower densities, between the bubble and droplet phases, several so-called ``pasta\" phases are predicted to put in an appearance, these being characterized by exotic, non-spherical shapes \\cite{pr95}. On the other hand, it has been shown that for other effective interactions the situation is much simpler, with no bubble or pasta phases being formed (at least at the assumed zero temperature of a stable neutron star): at a mean density of around $\\bar{\\rho}$ = 0.075 nucleons.fm$^{-3}$ the droplet phase undergoes a transition directly to the homogeneous phase (see Ref. \\cite{dh00}, and references cited therein; also Ref. \\cite{mar05}). In the present paper paper we will avoid these ambiguities by limiting ourselves to values of $\\bar{\\rho}$ less than 0.06 nucleons.fm$^{-3}$, which means that we would not be able to deal with ejection mechanisms that reached even deeper into the star.  Since neutron stars are formed at temperatures of the order of 10 MeV (10$^{11}$ K) and rapidly cool to around 0.1 MeV \\cite{pr95}, it is usually assumed that the final composition of the stable star corresponds to nuclear and beta equilibrium at a temperature of $T$ = 0, the configuration of so-called ``cold catalyzed matter\"; we shall later examine the validity of this assumption. Determining the composition of the outer crust in this picture is straightforward (see, for example, Ref. \\cite{rhs06}): the equilibrating nucleus at each given density (or pressure) is found from the known nuclear masses, as given by experiment or, where mass data are unavailable, a mass model such as the FRDM \\cite{frdm} or HFB-14 \\cite{gsp07} (see also Refs. \\cite{lpt03,pg06} for reviews). We shall therefore not consider the outer crust any further here.  As for the composition of the inner crust of the stable neutron star, the relevant question at a given mean density $\\bar{\\rho}$ is to determine the total number of neutrons $N$, including those of the vapor, and protons $Z$ per cluster. For this one needs the total Helmholtz free energy per nucleon $f$ (including the electronic contribution) at the ambient temperature (usually taken to be zero), as a function of the density and the composition $X \\equiv (Z, A = Z + N)$; one then minimizes $f$ with respect to $N$ and $Z$ at constant $\\bar{\\rho}$. (Alternatively, to determine the composition at a given {\\it pressure} $P$ one minimizes the Gibbs free energy per nucleon $g$ with respect to $N$ and $Z$ at constant  $P$. It follows from the easily proven thermodynamical relation \\beqy\\label{1.1} \\Big(\\frac{\\partial g}{\\partial X}\\Big)_{P,T} = \\Big(\\frac{\\partial f}{\\partial X}\\Big)_{\\bar{\\rho},T} \\eeqy that the two procedures are completely equivalent. We nevertheless find it more convenient to work with the Helmholtz free energy $f$ at given values of $\\bar{\\rho}$: see, for example, Section II of Ref. \\cite{bps71}.)  The pressure in a layer of the crust of density $\\bar{\\rho}$ can then be found by numerical differentiation from the identity \\beqy\\label{1.2} P = \\bar{\\rho}^2\\Big(\\frac{\\partial f}{\\partial \\bar{\\rho}}\\Big)_ {T,X} \\quad ; \\eeqy note particularly that $f$ is a mean quantity, averaged over inhomogeneities, and not a local quantity. With the pressure $P$ determined as a function of the mean density $\\bar{\\rho}$, the values of $P$ and $\\bar{\\rho}$ in any layer of the neutron star, along with the local composition, can be determined through solution of the Tolman-Oppenheimer-Volkoff equation \\cite{tol39,ov39}.  When, for one reason or another, decompression of crustal material begins, the temperature may start to rise. To follow the evolution of this process we shall require the EOS for non-zero values of $T$, and also the specific heat per nucleon at constant volume, $c_v$, given in terms of the entropy per nucleon, $s$, by \\beqy\\label{1.3} c_v = T\\Big(\\frac{\\partial s}{\\partial T}\\Big)_{\\bar{\\rho},X} \\quad . \\eeqy The entropy itself is given in terms of the Helmholtz free energy by \\beqy\\label{1.4} s = -\\Big(\\frac{\\partial f}{\\partial T}\\Big)_{\\bar{\\rho},X} \\quad . \\eeqy Thus all quantities of interest here can be derived from a calculation of $f$ as a function of $\\bar{\\rho}$ and $T$. Note that $s$ and $c_v$, like $f$, are mean quantities.  A popular EOS that has been extensively applied to supernova explosions is that of Lattimer and Swesty~\\cite{ls91}. However, the applicability of this EOS to neutron-star crusts is limited by the fact that it is based on the so-called compressible liquid-drop model without any shell corrections, which at the prevailing low temperatures can be expected to be significant. Actually, both Refs. \\cite{latt77} and \\cite{mey89} attempt to take account of shell effects, although in a rather rudimentary way, by making use of the algebraic bunching technique of Myers and Swiatecki \\cite{ms66}.  In the present paper, as in Ref. \\cite{opp97}, we model the inhomogeneous nuclear medium by a single spherical WS cell, and attempt to incorporate shell effects into this framework microscopically and self-consistently, thereby permitting some measure of continuity of treatment across the interface between the inner and outer crusts. The most obvious way to do this is through the Hartree-Fock (HF) method, as has already been done, for example, by Bonche and Vautherin \\cite{bv81} at finite temperature and by Negele and Vautherin\\cite{nv73} at zero temperature, using the WS approximation. However, we abandoned this approach for the following reason. While protons are strongly bound in the inner crust because of the large neutron excess, and thus show strong shell effects, for neutrons, by the very definition of the inner crust, there will be a continuous spectrum of unbound single-particle (s.p.) neutron states that are occupied. Thus any neutron added to the system must in reality go into this continuum, whence it follows that we should not expect any neutron shell effects. Actually, this conclusion will hold only if the dripped neutrons form a uniform liquid, and in reality scattering of unbound neutrons on the inhomogeneities of the crust may give rise to so-called Casimir or band effects\\cite{bm01,mh02,mbh03}, whose exact evaluation requires the application of the band theory of solids (see Ref. \\cite{cha07} and references quoted therein). Nevertheless, these neutron shell effects are much smaller than the proton ones \\cite{oy94} and have therefore a negligible impact on the EOS and the equilibrium composition of the inner crust, although they are known to be significant for transport properties \\cite{cha06}. However, in practice any HF calculation in the WS approximation involves discretization, giving rise to shell effects for both protons and neutrons. But, as we have argued above, these neutron shell effects must be spurious, and in the HF calculation of Ref. \\cite{nv73} special steps had to be taken to smooth them (see also Refs. \\cite{bst06,cha07}). We conclude that as far as neutrons are concerned the semiclassical extended Thomas-Fermi method is better adapted to a WS approach than is the HF method.  The solution we adopt here to the problem of including the appropriate proton shell corrections without introducing spurious neutron shell corrections is to use the ETFSI (extended Thomas-Fermi plus Strutinsky integral) high-speed approximation to the HF method \\cite{dut,ton,pea,abo1,abo2}. We have already made an exploratory study of the applicability of this method to a WS picture of the EOS, and found that proton shell effects are indeed important \\cite{dop04}, but here, in addition to making much more extensive calculations of the EOS, we improve the $T > 0$ results by taking account of possible shell effects in the entropy, which will manifest themselves in the free energy through the relation \\beqy\\label{1.5} f = e - Ts  \\quad  , \\eeqy where $e$ is the energy per nucleon.  A further development of considerable significance is that the TETFSI method, as we shall refer to this temperature-dependent ETFSI method, is no longer limited to forces whose effective nucleon mass $M^*$ is equal to the real mass $M$. This permits us to use more realistic effective forces with smaller values of the effective mass. Thus in the present calculations the effective interaction that we use is the Skyrme force BSk14, for which the effective mass in symmetric INM at the equilibrium density $\\rho_0$ (0.159 nucleons.fm$^{-3}$) is 0.800$M$, which is to be compared with the value of 0.825$M$ found in extended Brueckner-Hartree-Fock calculations that include three-nucleon forces \\cite{cao06}. This is the force that underlies the Hartree-Fock-Bogoliubov (HFB) mass model HFB-14 \\cite{gsp07}, a force that is eminently suitable for calculating the properties of neutron-star crustal matter, since on the one hand it has been fitted to the properties of neutron matter, as determined by calculations with realistic two- and three-nucleon forces \\cite{fp81}, and on the other hand it gives an excellent fit to essentially all the available mass data ($\\sigma_{rms} = 0.729$ MeV). Given that the neutron-star crust is both inhomogeneous and contains some protons, the high quality of the mass fit is especially relevant, since it means a) that inhomogeneities in nuclear matter (surface effects in droplet-model language) are well modeled, and b) that neutron-proton interactions are well represented. (However, no Skyrme force should be used for the highly supernuclear densities encountered deep within the core of a neutron star.) We stress that in this paper we neglect pairing, as in Refs. \\cite{nv73,bv81}.  In Section II we discuss our parametrization of the WS cell. Section III describes our adaptation of the ETFSI method to the problem of the EOS of the neutron-star inner crust at non-zero temperatures, with particular attention to the convergence properties of the ETF expansion of the entropy. Our formalism is applied in Section IV to the properties of the inner crust of a neutron star (we do not examine in this paper the important question of the rapid decompression of neutron-star matter). The existence of strong proton-shell effects in the inner crust is discussed in this same section. Our conclusions are summarized in Section V. Some important material is to be found in the appendices, notably the TETF expansion of the specific heat (App. A) and a proof of the Strutinsky-integral theorem (App. C). ",
        "conclusions": "We have developed here a high-speed approximation to the HF method for calculating the EOS of the neutron-star inner crust with Skyrme-type forces. Our method, which we refer to as the TETFSI method, models the inner crust in terms of the Wigner-Seitz cell, and consists essentially of a generalization to finite temperatures (and arbitrary effective mass) of the ETFSI method originally developed as a mass model \\cite{dut,ton,pea,abo1,abo2}. An essential difference between our TETFSI method and a full-scale HF calculation of the EOS is that, whereas the latter method inevitably and automatically calculates both neutron and proton shell effects, here we calculate only the latter, since in reality shell effects are much weaker for neutrons than for protons, and will have negligible impact on the composition. In fact, if the HF method is used in a WS picture, as in the classical work of Negele and Vautherin\\cite{nv73} it will lead, because of discretization, to spuriously large neutron shell effects\\cite{cha07}. As in Ref. \\cite{nv73}, we have neglected pairing in this paper, pending the determination of an effective pairing interaction appropriate to the conditions pertaining in neutron-star crusts. Nevertheless, it will be easy to include pairing in the (T)ETFSI framework, as already done in the ETFSI mass models \\cite{dut,ton,pea,abo1,abo2}.  It was found that the TETF expansion of the entropy converges badly at low temperatures for densities typical of inner-crust protons (there was no problem for neutrons). We solved this  difficulty by using the s.p. expression for the proton entropy.  Our exploratory calculations of the EOS were performed with the Skyrme-type force BSk14, a force that was fitted to essentially all the nuclear-mass data, forming thereby the basis of the HFB-14 mass model \\cite{gsp07}. This force is particularly suitable for calculating the properties of neutron-star crustal matter, because it has been fitted to the properties of homogeneous neutron matter while at the same time the good fit to masses ensures that both inhomogeneities and the neutron-proton interaction are well represented.  The calculated composition of the WS cells representing the clustering in the inner crust showed striking shell effects: for $T$ = 0 the proton number $Z$ was limited to the magic values of 50, 40 and 20, the value decreasing with increasing density (at the interface with the outer crust we found continuity with an outer-crust calculation based on the HFB-14 mass model). Although essentially identical results are obtained for $T$ = 0.1 MeV, all our calculated shell effects are wiped out at $T$ = 1 MeV, which means that whether or not shell effects actually exist in the cold crust depends very much on the ``freeze-out\" temperature for nuclear equilibrium. On the other hand, we have shown that even without taking shell effects into account there are considerable differences between our predictions and those of the compressible liquid-drop model on which the EOS of Ref. \\cite{dh01} is based.  We intend to apply the method described here to a study of the synthesis of r-process nuclei in decompressing neutron-star crustal matter. To this end we present here, apparently for the first time, the TETF expressions for the specific heat of an inhomogeneous system of nucleons. In this same context of extensive computations over a wide range of temperature, density and composition, we point out that the (T)ETFSI method lends itself admirably to interpolation, without any loss of precision in the calculated shell effects, essentially because these arise in the sums of quantities that themselves vary smoothly \\cite{pea}."
    },
    "0806/0806.2293_arXiv.txt": {
        "abstract": "Results are presented of a harmonic analysis of the large scale cosmic-ray anisotropy as observed by the Milagro observatory. We show a two-dimensional display of the sidereal anisotropy projections in right ascension generated by the fitting of three harmonics to 18 separate declination bands. The Milagro observatory is a water Cherenkov detector located in the Jemez mountains near Los Alamos, New Mexico. With a high duty cycle and large field-of-view, Milagro is an excellent instrument for measuring this anisotropy with high sensitivity at TeV energies. The analysis is conducted using a seven year data sample consisting of more than 95 billion events, the largest such data set in existence. We observe an anisotropy with a magnitude around 0.1\\% for cosmic rays with a median energy of 6 TeV. The dominant feature is a deficit region of depth (2.49 $\\pm$ 0.02 stat.\\ $\\pm$ 0.09 sys.) $\\times 10^{-3}$ in the direction of the Galactic North Pole centered at 189 degrees right ascension. We observe a steady increase in the magnitude of the signal over seven years. ",
        "introduction": "Observation of the sidereal large scale cosmic-ray (CR) anisotropy at energies of 1 - 100 TeV is a useful tool in probing the magnetic field structure in our interstellar neighborhood as well as the distribution of sources. Cosmic-rays at these energies are almost entirely of Galactic origin and are expected to be nearly isotropic due to interactions with the Galactic magnetic field (GMF) \\citep{AsCR}. The gyro radii of CRs at these energies in a GMF of about $1\\mu$G are from about 100 AU - 0.1 pc which is much smaller than the size of the Galaxy. This has the effect of trapping the CRs in the Galaxy for times on the order of $10^{6}$ years. Inhomogeneities in the GMF are interspersed randomly throughout the galaxy and in effect act as scattering centers for CRs. This randomizes their directions as they propagate leading to a high degree of isotropy on scales of a few hundred parsecs.  Anisotropy can be induced through both large scale and local magnetic field configurations which cause deviations from the isotropic diffusion approximation. At lower energies of around a TeV, the heliosphere may be able to induce a CR excess in the direction of the heliotail and also could modulate the overall CR anisotropy \\citep{NgLC,NgTI}. At higher energies, the contribution of discrete CR sources has been shown to be capable of creating a large scale anisotropy \\citep{PtSNR,SMP}.  Diffusion of CRs out of the Galactic halo can also produce an anisotropy. Since the matter density is higher in the Galactic disk compared to that in the surrounding halo, the diffusion coefficient will generally be much higher in the halo. For this reason, CRs produced in the Galactic disk will tend to diffuse out into the halo creating an anisotropy in the direction perpendicular to the disk. Predictions of the anisotropy have been made using values of the diffusion coefficient inferred using a given propagation model and observational data. This predicted anisotropy can take on values of $10^{-2}$ to $10^{-5}$ depending on the propagation model used \\citep{AsCR}. Given this correlation between the anisotropy and diffusion coefficient, knowledge of the large scale CR anisotropy can be used to constrain diffusion models.  In addition to the above effects, Compton and Getting introduced a theory \\citep{CG} of CR anisotropy which predicts a dipole effect due to the motion of an observer with respect to an isotropic CR plasma rest frame. The anisotropy arising from the Earth's motion around the sun is calculated to be on the order of $10^{-4}$ and would appear in universal time. This effect has been observed by numerous experiments (e.g. see \\cite{ET}, \\cite{TBCG}). There is also a possible sidereal time anisotropy coming from the motion of the Solar System through the Galaxy. This effect is more difficult to predict given that the isotropic CR rest frame is not known. Using the assumption that the CR rest frame does not co-rotate with the Galaxy leads to a predicted anisotropy on the order of $\\sim0.1\\%$. This effect has not been observed \\citep{TB}.  There have been numerous observations of the large-scale sidereal anisotropy in the range of energies $10^{11} - 10^{14}$ eV. In this paper we use large-scale anisotropy to refer to features of the anisotropy in the sky with an extent greater than $\\sim 40^{\\circ}$ in right ascension. Most of these observations are examined by fitting harmonics to the distribution of cosmic-ray events in the right ascension direction. The main features of the anisotropy, the amplitude and phase, are fairly constant over the energy range mentioned above ($3-10 \\times 10^{-4}$ for the amplitude, and $0 - 4$ hr for the phase, see \\cite{SK} for a compilation). At TeV energies, the focus of this paper, the anisotropy is known to have a value on the order of $10^{-3}$ with a deficit region, sometimes called the ``loss-cone\", around $200^{\\circ}$ right ascension, and an excess, or ``tail-in\" region, around $75^{\\circ}$ \\citep{NgLC,NgTI}. In this paper, we present the results of a harmonic analysis of the large scale cosmic-ray anisotropy in the northern hemisphere as observed by the Milagro experiment. ",
        "conclusions": "Previous experiments such as the Tibet Air Shower Array, with a modal energy of 3 TeV, and Super-Kamiokande-I, with a median energy of 10 TeV,  have identified two coincident regions of interest in their sidereal observations: an excess located at $\\sim75^{\\circ}$ r.a.\\ or ``tail-in\" anisotropy, and a deficit at $\\sim200^{\\circ}$ r.a.\\ or ``loss-cone\" anisotropy \\citep{TB,SK}. Both of these regions are consistent with Milagro observations. The ``loss-cone\" coincides with the deep central-deficit region seen in this analysis while the narrow ``tail-in\" excess is clearly observed in another Milagro analysis which is sensitive to features with extent smaller than $\\sim30^{\\circ}$ in r.a. (see \\citet{GW} region A).  The strengthening of the signal in the central-deficit region over time is a result unique to this analysis. Tibet found no evidence of time variation comparing data split into two five-year periods, 1997-2001 and 2001-2005. However, only the second of these time periods overlaps with our data set for which the average value we observe in the deficit region agrees well with their measured deficit.   The anisotropy observed in the galactic cosmic rays could arise from a number of possible effects. The Compton-Getting effect (CG) predicts that due to the motion of the solar system around the galactic center through the rest frame of the cosmic-ray plasma an anisotropy is induced with the form of a dipole with a maximum in the direction of motion. For no co-rotation of the cosmic ray plasma with the Galaxy, the magnitude of the anisotropy is calculated to be $0.35\\%$ given our speed of $\\sim220 km s^{-1}$, while at the other extreme of full co-rotation, the anisotropy would be zero. No evidence of a Galactic CG anisotropy was seen in \\cite{TB}. For no co-rotation, the dipole should have a maximum at r.a.\\ = $315^{\\circ}$ and dec.\\ = $48^{\\circ}$, and a minimum at r.a.\\ = $135^{\\circ}$ and dec.\\ = $-48^{\\circ}$. Since our analysis method yields only a projection of the anisotropy the observed CG effect will be slightly different from the true effect. The CG effect we expect to see in this analysis was determined from Monte Carlo simulation and found to be a dipole with a maximum of $0.14\\%$ at $315^{\\circ}$ r.a.\\ for the declination range between $50^{\\circ}$ to $60^{\\circ}$. This range was considered to try to reduce effects from the central-deficit region. For the actual data, fitting a single harmonic to the projection in r.a.\\ corresponding to declinations $50^{\\circ}$ to $60^{\\circ}$ gives a $\\chi^{2}/d.o.f.\\ = 11505/998$ which is clearly a poor fit. Although this suggests that the observed anisotropy is not dominated by the galactic Compton-Getting effect, its contribution to the anisotropy cannot be ruled out.   In addition to the Compton-Getting effect there is expected to be an anisotropy stemming from the diffusion of cosmic-rays in the interstellar medium. At high energies the main effects are expected to be mainly due to the distribution of discrete CR sources and the structure of the galactic magnetic field in the neighborhood of the solar system. One study conducted consists of a simple diffusion model assuming increased production in the galactic disk due to supernova remnants (SNR) \\citep{PtSNR,SMP}. Also considered was the diffusion of CR out of the galactic halo. This is attractive since we have a deficit region in the general direction of the North Galactic Pole which could be due in part to this diffusion. The main contribution of CR from SNR was considered for sources with distances from Earth of $<1$ kpc and ages $<0.05$ Myr. Calculations performed \\citep{PtSNR,SMP} using these sources and taking into account CR re-acceleration as well as diffusion out of the galaxy gives an anisotropy about 3 times greater than observed, with the main source of the anisotropy due to the Vela SNR located at $128^{\\circ}$ r.a.\\ and $-45.75^{\\circ}$ dec. This model only predicts the magnitude of the expected anisotropy, not its exact phase. It also is remarked that this is a simplified model which assumes an isotropic diffusion tensor which is not explicitly known to be true at these energies.   At energies of $\\sim1$ TeV the heliosphere is believed to have an influence on the distribution of CR \\citep{NgLC,NgTI}. One possible reason for the modulation of the anisotropy on the observed time scale could be due to variations in the heliosphere since we know that it changes in relation to solar output. It is noted that our data begins at the solar maximum and ends near the solar minimum. A recent derivation of the diffusion tensor contains a new component due to perpendicular spatial diffusion which is expected to be an important factor in understanding the anisotropy due to the Galactic disk as well as the modulation of CR in the outer heliosphere \\citep{PerpD}. Finding a consistent explanation of the observed anisotropy and especially its time evolution will be a challenge."
    },
    "0806/0806.2770_arXiv.txt": {
        "abstract": " ",
        "introduction": "\\label{intro}  A characteristic observed feature of cosmological gamma-ray bursts (GRBs) is that they emitted a huge amount of energy in $\\gamma$-rays in a very short time and their isotropic-equivalent $\\gamma$-ray energy (i.e., the total $\\gamma$-ray energy emitted by a GRB if the GRB radiates isotropically) spans a very large range---more than five orders of magnitude. The measured isotropic-equivalent energy of GRBs, $E_\\iso$, appears to have a log-normal distribution with a mean $\\sim 10^{53}$ erg, and a dispersion $\\sim 0.9$ in $\\log E_\\iso$ \\citep{ama06,ama07,li07}.  However, there is evidence that GRBs are beamed \\citep{har99,kul99,sta99}. Assuming that a GRB radiates its energy into two oppositely directed jets, each having a half-opening angle $\\theta_\\jet$. The total solid angle spanned by the jets is then $4\\pi\\omega$, where $\\omega\\equiv 1-\\cos \\theta_\\jet <1$. If the emission of a jet is distributed more or less uniformly on its cross-section, the total $\\gamma$-ray energy emitted by the GRB is approximately $E_\\gamma = \\omega E_\\iso$, smaller than $E_\\iso$ by a beaming factor $\\omega$.  One of the most important discoveries in GRB observations has been that the value of $E_\\gamma$ has a very narrow distribution with a mean $\\sim 10^{51}$ erg comparable to ordinary supernovae, which has led people to claim that GRBs are a standard energy reservoir involving an approximately constant explosion energy (Frail et al. 2001; Piran et al. 2001; Berger, Kulkarni \\& Frail 2003; Bloom, Frail \\& Kulkarni 2003; Friedman \\& Bloom 2005). Theoretical models for interpreting the clustering GRB energy have also been proposed \\citep[see, e.g.,][]{zha02}.  It is well-known that observations of GRBs are strongly fluence or flux-limited, hence GRB samples seriously suffer from Malmquist-type selection biases \\citep{mal20,tee97}. That is, an observer will see an increase in the averaged luminosity or the total energy of GRBs with the distance, caused by the fact that less luminous or sub-energetic bursts at large distances will not be detected. Although this Malmquist bias for a flux-limited sample of astronomical objects looks obvious, sometimes people made serious mistakes in interpreting data by neglecting it. For instance, with a study of nearby galaxies it had been incorrectly claimed that the Hubble constant increases with the distance \\citep{dev72,tee75,san94}.  Unfortunately, as people drew the conclusion on the distribution of the GRB energy and claimed that GRBs are standard energy explosions, they have treated the observed distribution as the intrinsic distribution and have neglected the selection biases that are very important for GRBs at cosmological distances. As a result, the number of faint GRBs has been significantly underestimated, since a GRB will not be detected if its flux or fluence falls below the detection limit. The bright part of the GRB energy distribution suffers from another important selection bias, which arises from the fact that brighter GRBs tend to have smaller jet opening angles and hence smaller probabilities to be detected. This beaming bias is independent of instruments and thus cannot be reduced by improving the sensitivity of detectors.  The aim of this paper is to show that the influence of the selection biases from the fluence limit and the jet beaming is strong enough that the observed distribution of the GRB energy does not represent the intrinsic distribution at all. We present a simple model that explains nicely the observed distribution of the GRB energy, yet the burst energy reservoir in the model is not standard. Hence, the collimation-corrected energy of GRBs can have a very broad intrinsic distribution despite the fact that it is observed to cluster to a narrow distribution. Our results lead to the suggestion that GRBs are not a standard energy reservoir, contrary to the previous claim. ",
        "conclusions": "\\label{conclusion}  The observation that the total energy emitted in $\\gamma$-rays by long-duration GRBs clusters around $10^{51}$ erg \\citep{fra01}, which has been considered as the most intriguing finding in GRB research \\citep{zha04}, is only a superficial result since the strong selection biases from the detector selection effect and the beaming of GRBs have been ignored. The previous claim that the energy output of the central engine of long-duration GRBs has a universal value \\citep{fra01,pir01}, which was derived from the above superficial result, is likely to be incorrect since the observed narrow distribution of $E_\\gamma$ is consistent with a broad intrinsic distribution of $E_\\gamma$.  In fact, our results show that for both $\\log E_\\iso$ (Fig.~\\ref{fig2}) and $\\log E_\\gamma$ (Fig.~\\ref{fig3}), the distribution on the left-hand side to the maximum is well modeled by the cut-off from the fluence limit (the solid curve in Fig.~\\ref{fig1}). It would be surprising that the intrinsic distribution happens to have a low-energy cut-off that is coincident with the fluence limit cut-off.  The influence of the flux or fluence limit of detectors on the observation of GRBs is well-known and has been taken into account either thoroughly or partly in many GRB works, e.g. in deriving the luminosity function of GRBs (Schmidt 1999, 2001; Firmani et al. 2004; Guetta, Piran \\& Waxman 2005; Liang et al. 2007). However, the influence has sometimes been ignored or seriously underestimated. The claim that the collimation-corrected energy of GRBs has a narrow distribution and hence GRBs are a standard energy reservoir is an example where the selection effects have been ignored and wrong physical conclusions have been drawn.  Although it is generally conceived that the jet opening angle is anti-correlated to the GRB energy, in the study on the luminosity function of GRBs the effect of jet beaming has often not been properly taken into account. For example, in \\citet{gue05} and \\citet{lia07}, an isotropic luminosity function was derived by comparing the model prediction with the observed flux or the luminosity distribution without a consideration of beaming, then the derived luminosity function was used to calculate a weighted and averaged beaming factor. As we can see from equation (\\ref{hat_phi}), the isotropic luminosity function derived by them should be the product of the intrinsic isotropic luminosity function and the averaged beaming factor as a function luminosity, not the intrinsic isotropic luminosity itself.  Our results also challenge the proposal that GRBs can be used as standard candles to probe cosmology \\citep[and references therein]{blo03,fri05,sch07}. Although the constancy of the GRB energy is not a necessary condition for GRBs to be standard candles because of the identification of several good correlations among GRB observables, the existence of a large amount of faint bursts that have not been observed might significantly increase the scatter in those correlations. In addition, the recent work of \\citet{but07} indicates that some of those relations arise from partial correlation with the detector threshold and hence are  unrelated to the physical properties of GRBs.  Finally, a prediction of this work that can be tested with future observations is that as the sensitivity of GRB detectors increases the observed distribution of $E_\\gamma$ broadens towards the low-energy end."
    },
    "0806/0806.1290_arXiv.txt": {
        "abstract": "{Lutz-Kelker bias corrected absolute magnitude calibrations for the detached binary systems with main-sequence components are presented. The absolute magnitudes of the calibrator stars were derived at intrinsic colours of Johnson-Cousins and {\\em 2MASS} (Two Micron All Sky Survey) photometric systems. As for the calibrator stars, 44 detached binaries were selected from the {\\em Hipparcos} catalogue, which have relative observed parallax errors smaller than 15\\% ($\\sigma_{\\pi}/\\pi\\leq0.15$). The calibration equations which provide the corrected absolute magnitude for optical and near-infrared pass bands are valid for wide ranges of colours and absolute magnitudes: $-0.18<(B-V)_{0}<0.91$, $-1.6<M_{V}<5.5$ and $-0.15<(J-H)_{0}<0.50$, $-0.02<(H-K_{s})_{0}<0.13$, $0<M_{J}<4$, respectively. The distances computed using the luminosity-colours (LCs) relation with optical ($BV$) and near-infrared ($JHK_{s}$) observations were compared to the distances found from various other methods. The results show that new absolute magnitude calibrations of this study can be used as a convenient statistical tool to estimate the true distances of detached binaries out of {\\em Hipparcos'} distance limit.} ",
        "introduction": "The most reliable physical parameters of stars, e.g. masses, radii, effective temperatures etc., used in testing theoretical predictions of fundamental physical structure of single stars, are obtained from the observations of detached eclipsing binaries with main-sequence components. Masses and radii could be determined within an accuracy better than 1\\% from their light and radial velocity curves \\citep{Andersen1991, Southworthetal2004, Southworthetal2005a}.  There could be several techniques to determine distances for detached eclipsing binaries. The most reliable distances are usually the ones from accurately measured trigonometric parallaxes or interferometric observations of nearby stars. Relatively less reliable distances are the ones based on spectroscopic or photometric parallaxes using absolute magnitudes estimated or computed from the solutions of data obtained by photometric or spectroscopic observations. Such distances, however, require measurements of reliable reddening free apparent magnitudes. The effect of interstellar reddening on the final distance predictions could be larger than expected. Since reddening is less effective at infrared wavelengths, the prediction errors could be minimized by using infrared photometry \\citep{Southworthetal2005b}.  Relation between orbital period, luminosity and colours of the binary stars, called PLC relation, was first proposed by \\cite{Rucinski1974} \\citep[see also,][]{Mochnacki1981} for W UMa-type binaries in order to reduce the scatter in the main-sequence colour-magnitude relation. \\cite{Rucinski1994} established the PLC relation of W UMa-type stars using ($B-V$), ($V-I_{c}$) colours and orbital periods. The PLC relation of W UMa-type binaries was improved later using {\\em Hipparcos} parallaxes \\citep{RucinskiandDuerbeck1997}. In a recent paper, \\cite{Aketal2007} described a similar relation, called PLCs relation for cataclysmic binaries. Although it is possible to find a PLCs relation for detached binaries too, applying this idea to detached binaries makes no sense since the orbital period is not correlated in any physical way to the components. Thus, a relation between the luminosity (absolute magnitude) and colours can be more useful for detached binaries. It should be noted that more accurate calibrations can be obtained by using two colours instead of one colour, since one colour dependent calibrations may include systematic errors \\citep[see, e.g.][]{Biliretal2008}.  In this study, we first estimate the systemic $V$ and $J$ band absolute magnitudes $M_{V}$ and $M_{J}$ of detached binaries using reliable trigonometric parallaxes from {\\em Hipparcos} satellite by assuming that the light in $V$ and $J$ bands comes from the system as a $\\it whole$. Then, we find the dependence of the absolute magnitude on de-reddened colours $(B-V)_{0}$, $(J-H)_{0}$ and $(H-K_{s})_{0}$ to derive absolute magnitude calibrations for detached binaries with observations in Johnson-Cousins and {\\em 2MASS} photometric systems. We should state that our aim is $\\it not$ to measure the absolute magnitudes of each component in the detached binaries in our list.  Unfortunately, measured parallaxes are never free from systematic errors. This problem has been noticed by \\citet[LK,][]{LK73}. Assuming uniform spatial distribution of stars and a uniform distribution of observed parallaxes around the true parallax, LK explained that there must be a systematic error, known as Lutz-Kelker bias, in computed distances which depends only upon relative parallax error ($\\sigma_{\\pi}/\\pi$), where $\\pi$ is observed parallax and $\\sigma_{\\pi}$ is the associated standard error. Because of LK bias, the true parallaxes are always smaller than observed parallaxes \\citep{LK73}. Therefore observed parallaxes are needed to be corrected. This correction procedure is recognized as LK correction.  According to \\cite{J01}, only the studies which are careful to use parallaxes with ($\\sigma_{\\pi}/\\pi)\\leq0.1$ could be excused since the bias could be tolerable. Otherwise, LK bias (if not corrected) would either alter conclusions or invalidate the results. However, the standard LK corrections become significant if ($\\sigma_{\\pi}/\\pi)\\geq0.05$ and increase as relative error gets bigger \\citep{M05}. Thus the parallaxes with such significant errors must be corrected.  Aim of this study is to establish an absolute magnitude calibration for practical usage to estimate distances of detached main-sequence eclipsing binaries which have no tri-gonometric parallax nor known physical structures to imply their absolute brightness, but only have at least an eclipsing light curve indicating they are on the main sequence and detached. Unlike the common practice of predicting binary distances photometrically, which requires estimated absolute magnitudes of each component extracted from the light contributions of each component and/or their physical parameters from a spectroscopic or a light curve solution, this new method does not require extensive work; thus practical and quick, therefore, useful especially when studying large numbers of eclipsing data such as ASAS \\citep{Paczynski06} survey where light curve solutions are not available to most binaries. ",
        "conclusions": "We have suggested absolute magnitude calibrations for detached binaries based on LK corrected {\\em Hipparcos} parallaxes. The calibration equations cover a wide range of colours and absolute magnitudes.  The distances of calibration stars were calculated using the absolute magnitudes obtained from the LCs relations and compared with the {\\em Hipparcos} distances (Fig. 3). Although distances are in agreement in general (concentrated on diagonal), there is a considerable scatter for some systems. Large deviations from the LCs relations are mostly seen for the systems whose distances are higher than $\\sim$100 pc with the exception of RT And, which is closer than 100 pc. We have listed these scattered systems in Table 5. The table shows that distances from literature agree better with our distance estimates from the calibration relation including near-infrared colours. The table also indicates that {\\em Hipparcos} distances are not in good agreement with the distances recorded in literature, in general, although their relative parallax errors are smaller than 0.15. One of the best examples is V1229 Tau. \\cite{Southworthetal2005b} estimated the distance of V1229 Tau using three independent methods and found that {\\em Hipparcos} distance is not in agreement with their results. Our estimates support \\cite{Southworthetal2005b}, but not in agreement with {\\em Hipparcos} distance. Thus, it seems that, in selecting stars for such calibration studies, not only relative parallax errors but also parallaxes higher than 10 mas must be carefully taken into account.  \\begin{table} \\setlength{\\tabcolsep}{4pt} \\center \\caption{The distances of scattered stars in calibration sample obtained from various other sources in literature. Columns 2, 3 and 4 include distances calculated using calibration Eqs. 5, 6 and LK corrected {\\em Hipparcos} parallaxes, respectively. Column 5 shows the distances collected from literature.} \\begin{tabular}{lccccc} \\hline &      Eq. 5   &   Eq. 6   & Hip   &        Lit.  &     \\\\ Star &     $d$ (pc) &     $d$ (pc) &  $d_{LK}$ (pc) &     $d$ (pc) & Refs\\\\ \\hline RT And &    132$\\pm$7 &    129$\\pm$6 &    78$\\pm$6  & 103          & (1) \\\\ V805 Aql &    164$\\pm$8 &    139$\\pm$8 &   192$\\pm$26 & --           & (-) \\\\ AR Aur &    174$\\pm$9 &    157$\\pm$9 &   127$\\pm$12 & 136$\\pm$7    & (2) \\\\ ZZ Boo &     70$\\pm$4 &     82$\\pm$4 &   116$\\pm$10 & 79           & (3) \\\\ AR Cas &    129$\\pm$7 &           -- &   184$\\pm$18 & --           & (-) \\\\ EK Cep &   284$\\pm$15 &    211$\\pm$12&   158$\\pm$14 & 190          & (4) \\\\ RZ Cha &    107$\\pm$5 &    115$\\pm$7 &   195$\\pm$22 & 190$\\pm$10   & (5) \\\\ V624 Her &     94$\\pm$5 &           -- &   152$\\pm$16 & 111          & (6) \\\\ UW LMi &     84$\\pm$5 &     95$\\pm$5 &   141$\\pm$18 &  82          & (7) \\\\ TZ Men &    170$\\pm$8 &    128$\\pm$7 &   108$\\pm$6  & 120$\\pm$10   & (8) \\\\ V1229 Tau &    183$\\pm$9 &    151$\\pm$8 &   116$\\pm$12 & 139$\\pm$4    & (9) \\\\ \\hline \\end{tabular} \\\\ { (1) \\cite{Arevaloetal1995}, (2) \\cite{Semeniuk2000}, (3) \\cite{Popper1998}, (4) \\cite{Popper1987}, (5) \\cite{Andersenetal1975}, (6) \\cite{Lacy1979}, (7) \\cite{Nordstrometal2004}, (8) \\cite{Andersenetal1987b}, (9) \\cite{Southworthetal2005b} } \\end{table}  It has been found that the distances indicated by the LCs relations and various methods are in agreement (see Fig. 4). Although there is some scattering in the figure, almost all systems are located very near to the diagonal line representing equal values. It is interesting to note that the systems in Fig. 4 with third body or apsidal motion do not display larger scatter in comparison to other stars.  Normally, colour-magnitude relation exists for main- sequence single stars which is the main part of HR diagram itself. Similar colour-magnitude relation obviously exists also for main-sequence binaries with a deviation of about 0.50 mag (see Table 3). This could be explained by an argument that the greatest deviation from the relation of single star must occur when two components of the binary are identical. For such cases a binary system would appear 0.75 mag brighter than a single star. So, deviation of binary system cannot be higher than 0.75 mag on the brightness scale. In fact, deviation becomes smaller if one of the stars becomes cooler. When the cooler star's contribution becomes negligible, the system behaves as single star. Assuming that, we have various systems so that expected deviation would be about smaller than 0.40 mag on the brightness scale on an average. This limit is well within the standard deviation predicted in Table 3. Colour axis deviation from single stars are also expected to be negligible. Because at both limits colour of the system is like a single star. When two components were identical there would be no colour deviation. When the temperature difference between components increases system colour approaches to the colour of the dominant component.  The results of this study will help to understand space distributions and evolution of detached binaries in the solar neighbourhood. Indirectly, it will help to improve information about single stars. We suggest that the LCs relations are useful statistical tools to calculate the distances of detached binaries from their optical and near-infrared observations since the LCs relations are based on the most reliable distance estimation method (trigonometric parallax). Distances calculated from the LCs relations can give clues for astrometric observations of these systems, as well. Finally, it should be noted that future astrometric observations of detached binaries such as {\\em GAIA} and {\\em SIM} missions, will refine the LCs and PLCs relations."
    },
    "0806/0806.4232_arXiv.txt": {
        "abstract": "We present the results of the spectroscopic observations of HDE\\,226868, the optical counterpart to the black hole X-ray binary Cyg X-1, from 2001 to 2006. We analyze the variabilities of the two components in the complex H$\\alpha$ line: one P-Cygni shaped component which follows the motion of the supergiant and another emission component moving with an antiphase orbital motion relative to the supergiant, which is attributed to a focused-stellar wind. The results of KOREL disentangling of our spectra indicate that the focused stellar wind is responsible for the major part of the variability of the H$\\alpha$ emission line. The emission of the supergiant component had a small difference between the low/hard and high/soft states, while the focused wind component became strong in the low/hard state and weak in the high/soft state. The wind is nearly undisturbed by the X-ray photoionization during the low/hard state. However, during the high/soft state, the X-rays from the compact object could decelerate the line-driven wind and result in a high mass accretion rate, due to the effect of the X-ray photoionization. The X-ray illuminating could also change the temperature profile of the stellar wind and increase its temperature, and thus decrease the H$\\alpha$ emissivity of the wind, which could explain the H$\\alpha$ variabilities of Cyg X-1 during different X-ray states. ",
        "introduction": "Cygnus X-1 was first discovered during a rocket flight observation in 1964 \\citep{bowyer65} and its optical counterpart was identified with the supergiant star HDE 226868 \\citep{bolton72, webster72}. A 5.6d orbital period was found in the optical spectroscopic observations \\citep{gies82} and the dynamical determination of the binary components provided evidence for the existence of a black hole in Cyg X-1. As the black hole X-ray binary, which was first found in our Galaxy, Cyg X-1 has been extensively studied from radio to $\\gamma$-rays during recent decades. Even though the mass function is precisely known (see \\citet{gies2003}), the masses of the donor star and the black hole are not so well constrained due to the poorly known inclination and the evolutionary status of the supergiant. Using the inclination of $i=35^\\circ$, \\citet{herrero95} estimated the masses of the supergiant and the black hole at 17.8 $M_\\odot$ and 10.1 $M_\\odot$, respectively.   HDE 226868 was classified as an O9.7 Iab supergiant star \\citep{walborn73}, which shows H$\\alpha$ and He {\\small II} $\\lambda$4686 in emission. Although the supergiant nearly fills the Roche lobe, the accretion is mainly via the strong stellar wind from the donor star \\citep{gies86a, gies2003}. The variabilities of the optical lines \\citep{gies86b,ninkov87} on the spectrum of HDE 226868 indicate that the distribution of the stellar wind deviates from a spherical geometry and that an enhanced wind flow exists (``focused stellar wind,\" suggested by \\citet{friend82}) in the direction of the compact object. The focused stellar wind is also revealed by the High-Energy Transmission Grating Spectrometer aboard the Chandra X-Ray Observatory \\citep{miller2005}.   The most recent ephemeris of the 5.6d orbital period has been given by \\citet{lasala98}, \\citet{brocksopp99b}, and \\citet{gies2003} according to their optical spectroscopic observations. In addition, the 5.6d orbital period is also found in the $UBVJHK$ photometry, X-ray and radio data of Cyg X-1 \\citep{brocksopp99a}. Cyg X-1 also shows the superorbital modulation, on a time-scale much longer than the orbital period. A 294d period was first reported by \\citet{priedhorsky83} in X-rays and then by \\citet{kemp83} in the optical. Another $\\sim$ 150d period was found by different authors \\citep{brocksopp99a,lachowicz2006} in X-ray and radio data. This $\\sim$ 150 d period may be caused by the precession of the accretion disk around the compact object \\citep{wijers99}.  A relativistic radio jet with a velocity larger than $\\sim$ 0.6$c$ was detected by \\citet{stirling2001} in their radio observations of Cyg X-1, and therefore Cyg X-1 joins the group of the Galactic microquasar. Microquasar is an X-ray binary with a pair of relativistic radio jets, which is similar to the radio jets found in the active galactic nuclei (see \\citet{mirabel99} for a review). The jet remnant of Cyg X-1 is resolved in radio \\citep{gallo2005} and optical \\citep{russell2007} observations.   Many observational characteristics of the canonical stellar black hole candidate Cyg X-1 are considered evidence that a black hole exisits, similar to other X-ray binary systems. These observational features include the ultra-soft spectra, the high-energy power-law tail above 20 keV, the spectral/temporal transition from the high/soft state to the low/hard state and the X-ray  millisecond variability (see \\citet{tanaka95}; \\citet{mcclintock2006}; \\citet{remillard2006}). Most of the time Cyg X-1 is in a low/hard state and in some years it can transit from the low/hard state to the high/soft state, which can continue from several weeks to several months \\citep{zhang97,brocksopp99a}. Occasionally, Cyg X-1 enters an intermediate state when it fails to make a transition from the low/hard state to the high/soft state \\citep{belloni96,malzac2006}. The radio and X-ray emission has an anticorrelation and when the binary enters the high/soft state, the radio emission will be decreased \\citep{pooley2001}. Most researchers believe the transition state is caused by the physical changes of the gas around the black hole (see \\citet{remillard2006} and references therein). The physical changes in the accretion disk are initially related to the mass loss from the supergiant companion. However, the X-ray radiation of the disk influences the ionization and temperature of the wind and thus also its radiation and dynamics.    Using the method of tomographic separation\\footnote{ This method of decomposition of observed spectra into unknown spectra of two components should be distinguished from the methods of Doppler tomography in which the phase-locked line-profile variations are fitted by projections of smooth distribution of delta-function profiles in velocity space corotating with the orbital motion.} \\citep{bagnuolo91}, \\citet{sowers98} showed that the H$\\alpha$ profiles of Cyg X-1 observed in seasons 1985 and 1986 can be well fitted as a superposition of a P-Cyg profile corresponding to the (approximately spherical) stellar wind of the supergiant and a wide emission peak radiated by the focused stellar wind, which moves in a slightly shifted anti-phase with respect to the orbital Doppler shift of the former component. Using the method of Fourier disentangling \\citep{hadrava95,hadrava97,hadrava04} we have proved the same behavior in the spectra obtained at Ond\\v{r}ejov Observatory in summer 2003 shortly before, during and shortly after one high/soft-state episode of Cyg X-1 \\citep{hadrava07}.    In this article, we present our optical spectroscopic observations of Cyg X-1 from 2001 to 2006. The preliminary results have been introduced in a previous article \\citep{yan2005}. Here we will make a further analysis of the H$\\alpha$ line profiles. Our present data are confined to about one week in each season, meaning they are concentrated into only slightly more than one orbital period. This enables us to compare the solutions of disentangling in different X-ray states, minimizing the possible influence of long-term variations in the structure of the focused stellar wind.      In the following section, the properties of our data are described. The method and results of the disentangling are given in Section 3. Next, the correspondence of our observational results with the standard model of this classical X-ray binary is studied in Section 4. Finally, the conclusions are summarized in Section 5. ",
        "conclusions": "We present and analyze our optical spectroscopic observations of Cyg X-1 from 2001 to 2006. Combined with the RXTE/ASM X-ray observations, we make the following findings about this classical galactic black hole X-ray binary: \\begin{enumerate} \\item We confirm that the H$\\alpha$ line shows two components: a P-Cygni profile moving with the radial velocity curve of the supergiant and a focused stellar wind component that moves with an approximately anti-phase orbital motion relative to the supergiant. The superposition of the two components forms the complex H$\\alpha$ profiles.  \\item The results of KOREL disentangling the H$\\alpha$ spectra during our 2004 and 2006 observations indicate that the focused-stellar wind is responsible for the major part of the H$\\alpha$ variability between different X-ray states. The focused wind component becomes strong during the X-ray low/hard state and weak during the high/soft state. The photoionization and heating of the X-ray photons from the compact object may affect the ionization state and dynamics of the wind from the supergiant. During the high/soft state, the X-ray photoionization could decelerate the gas in the focused stellar wind and result in an increasing mass accretion rate. During the low/hard state, the X-ray is not strong enough to influence the wind ionization state and the compact object has a low-mass accretion rate. The X-ray illumination can form a temperature-inversion layer in the stellar wind. During the high/soft state, the soft X-rays acting onto the focused stellar wind could increase its temperature greatly and thus decrease the H$\\alpha$ emissivity. This could explain the variability of the focused-wind component in H$\\alpha$ during different X-ray states. The strong soft X-ray emission during the high/soft state could also ionize the outer parts of the supergiant winds and render it unable to absorb and re-emit in H$\\alpha$. This scenario is consistent with the small differences between the shapes of disentangled P-Cygni components in the 2004 and 2006 spectra. During our 2001 observations, the wind near the supergiant was also affected by the strong X-ray emission and an extremely weak H$\\alpha$ line was observed.  \\item The H$\\alpha$ lines in our 2006 observations are very strong. We interpret this as the result of the low irradiation of the focused wind by the X-ray photons. The weak incident soft X-rays during the flare could not disturb the temperature profile of the focused stellar wind, while the hard X-rays could heat up the deep layer of the wind and an extra H$\\alpha$ emission could be observed.  \\end{enumerate}"
    },
    "0806/0806.4142_arXiv.txt": {
        "abstract": "We present high precision radial velocity observations of HD\\,17156 during a transit of its eccentric Jovian planet. In these data, we detect the Rossiter-McLaughlin effect, which is an apparent perturbation in the velocity of the star due to the progressive occultation of part of the rotating stellar photosphere by the transiting planet.   This system had previously been reported by \\citet{NaSaOh08} to exhibit a $\\lambda = 62\\arcdeg \\pm 25\\arcdeg$ misalignment of the projected planetary orbital axis and the stellar rotation axis. We model our data, along with the \\citeauthor{NaSaOh08} data, and obtain $\\lambda = 9.4\\arcdeg \\pm 9.3\\arcdeg$ for the combined data set.   We thus conclude that the planetary orbital axis is actually very well aligned with the stellar rotation axis. ",
        "introduction": "Transiting extrasolar planets allow us to perform critical tests of models of planetary system formation and evolution. In our solar system, all of the planets orbit approximately, {\\itshape but not exactly}, in the solar equatorial plane.  \\citet{BeGi05} find that the angle between the plane of the ecliptic and the solar equator is $7.155\\pm0.002^{\\circ}$.   The near co-planarity of the planetary orbits in our solar system has influenced models of planetary system origin from the time of \\citet{Ka1755} and \\citet{La1796} to essentially all modern models \\citep[e.g.][]{Li95,PoHuBo96,Bo00}. Given the level of misalignment in our solar system, and the observation of warps in debris disks around nearby stars such as $\\beta$~Pic \\citep[e.g.][]{BuKrSt95,MoLaPa97,HeLiLa00}, it is obvious that there are common processes which give rise to some small level of spin-orbit misalignment in planetary systems.  The degree of misalignment in real planetary systems must depend on the initial conditions (the initial asymmetries of the collapsing cloud), the physics of disk formation and evolution, and the physics of stellar mass and angular momentum loss during the T~Tauri phase. Most transiting planets have periods of just a few days and orbits of low eccentricity.   Significant exceptions are HD\\,147506b (HAT-P-2b) \\citep{BaKoTo07}, with a 5.6 day orbital period and eccentricity of 0.52, HD\\,17156 \\citep{FiVoMa07,BaAlLa07}, which has the longest orbital period (21 days) and the largest eccentricity (0.67) of all of the transiting exoplanets, and XO-3b, a massive (13.25 $M_J$) planet which has an eccentricity of 0.26 in spite of it short orbital period of 3.192 days \\citep{JKMCBu08}. Scenarios for the formation of short-period planetary systems do not necessarily deliver those planets to their present semi-major axes with low eccentricity.   Type~II migration can result in planets with moderate eccentricities \\citep{SaGo04}, while dynamical interactions among newly formed planets and planetesimals \\citep{JuTr07} or the Kozai mechanism \\citep{HoToTr97,WuMu03} can result in large eccentricities and significant inclination of planetary orbital planes from the plane of the stellar equator (which is presumably close to the plane of the inner portion of the proto-planetary disk). While the shortest period transiting systems have probably been tidally circularized, longer period systems can easily retain large eccentricities over the main-sequence lifetime of the parent star. A planet that has undergone significant gravitational scattering or Kozai excitation would not necessarily retain a low inclination relative to the stellar equator. Thus, the report of a possible spin-orbit misalignment of the HD\\,17156b transiting planet by \\citet{NaSaOh08} is extremely interesting, and calls for further detailed investigation. In this paper, we report our own spectroscopic observations of a transit of HD\\,17156 by its planet, using two telescopes at McDonald Observatory. ",
        "conclusions": "\\label{sec:conclusions} We have reanalyzed the data of \\citet{NaSaOh08} along with our own independent observations of the spectroscopic transit of the eccentric exoplanet HD\\,17156b, and we find that $\\lambda = 9.4\\arcdeg \\pm 9.3\\arcdeg$. We conclude that this exoplanetary system is similar to almost all of the other short-period transiting exoplanetary systems studied using the RM effect so far, in that it shows that the projected planetary orbital axis appears to be aligned with the projected stellar rotation axis, to within our measurement precision. The only remaining notable exception is XO-3b, which was reported by \\citet{HeBoPo08} to show a very significant misalignment of $\\lambda = 70\\arcdeg \\pm 15\\arcdeg$.  The HD\\,17156b system is extremely interesting because it has a very high eccentricity ($e = 0.67$) for such a short orbital period of 21.2\\,days. Most other transiting systems have much shorter orbital periods and eccentricities near zero.   Many of the interesting ways of getting a planet into this type of orbit might well result in the possibility of an orbital plane significantly inclined to the stellar equatorial plane. Dynamical interactions with another nearby planet could result in ejection of the other body and a large orbital inclination of the surviving body.  The presence of a third planet in the system, as suggested by \\citet{ShWeOr08}, would have some effect on the orbital elements of HD\\,17156b.  However, as \\citet{ShWeOr08} discussed, the primary effects are a small oscillation of the eccentricity and a secular advance of the argument of periastron.   As long as the two planets are approximately coplanar, there would be no induced change in the orbital inclination. Thus, it appears that the dynamical process that placed HD\\,17156b into a short-period highly eccentric orbit did not significantly affect the orbital inclination of the system with respect to the stellar equatorial plane."
    },
    "0806/0806.0627_arXiv.txt": {
        "abstract": "Multi-epoch {\\it Chandra} X-ray observations of nearby massive early-type galaxies open up the study of an important regime of low-mass X-ray binary (LMXB) behavior --- long term variability. In a companion paper, we report on the detection of 158 X-ray sources down to a detection/completeness limit of $0.6/1.4\\times 10^{37} {\\rm \\, ergs \\, s}^{-1}$ using five {\\it Chandra} observations of NGC~4697, one of the nearest ($11.3 {\\rm \\, Mpc}$), optically luminous ($M_B < -20$), elliptical (E6) galaxy. In this paper, we report on the variability of LMXB candidates measured on timescales from seconds to years. At timescales of seconds to hours, we detect five sources with significant variability. Approximately $7\\%$ of sources show variability between any two observations, and $16\\pm4\\%$ of sources do not have a constant luminosity over all five observations. Among variable sources, we identify eleven transient candidates, with which we estimate that if all LMXBs in NGC~4697 are long-term transients then they are on for $\\sim 100 {\\rm \\, yr}$ and have a 7\\% duty cycle. These numbers are consistent with those found for brighter LMXBs in M87 and NGC~1399, which suggests that there does not appear to be a measurable difference between the outburst durations of long-term transient neutron star LMXBs and black hole LMXBs. We discuss in detail a transient supersoft source, whose properties are not easily explained by standard explanations for supersoft sources. ",
        "introduction": "\\label{sec:n4697x_intro}  With the angular resolution of the \\textit{Chandra X-Ray Observatory} low-mass X-ray binaries (LMXBs) in nearby early-type galaxies can be detected. For early-type galaxies with a relatively low X-ray--to--optical luminosity ratio, known as X-ray faint galaxies, the emission from LMXBs is the dominant source \\citep[e.g.,][hereafter Papers I and II]{SIB2000,SIB2001}. In typical {\\it Chandra} observations of such galaxies, tens to a few hundred LMXBs are resolved at bright X-ray luminosities, $L_X \\gtrsim 5$--$10\\times 10^{37} {\\rm \\, ergs \\, s}^{-1}$. By stacking multi-epoch {\\it Chandra} observations of early-type galaxies, larger numbers of fainter LMXB candidates can be studied, and brighter LMXB candidates can be studied in greater detail. Such observations open up the study of an important regime of behavior --- long term variability. Low mass X-ray binaries exhibit a wide range of variability, including transient outbursts, type I bursts, and milder fluctuations. Our understanding of LMXB variability in the Galaxy drives models for accretion and LMXB evolution. However, there are several limitations in studying Galactic LMXBs: source distances are known for only a small subset, it is difficult to observe the whole Galaxy at once, and the size of the observed sample is very limited. Observations of nearby early-type galaxies can overcome these limitations of Galactic studies.  Examples of extragalactic LMXB candidates with extreme behaviors, such as luminous supersoft sources (SSs) and flaring sources, have also been examined with {\\it Chandra} in greater detail. In our Galaxy and M31, SSs have very soft X-ray spectra $(\\lesssim 75 {\\rm \\, eV})$ and are generally believed to be accreting white dwarfs (WDs). Since the bolometric luminosities of extragalactic SSs can exceed the Eddington luminosity for a Chandrasekhar mass WD, an alternate hypothesis that these sources contain intermediate-mass ($\\sim 10^2$--$10^3 \\, M_{\\sun}$) accreting BHs has been put forth \\citep{SGS+2002}. Due to the degradation of the soft response of the Advanced CCD Imaging Spectrometer (ACIS) on {\\it Chandra}, the study of SSs has been limited. Some LMXBs exhibit another extreme behavior, namely relatively short timescale (seconds to minutes) flares. Recently, {\\it XMM-Newton} detected type I X-ray bursts in M31 \\citep{PH2005} and we reported a more extreme example of flares in NGC~4697 \\citetext{\\citealp{SSJ2005}, hereafter \\citetalias{SSJ2005}}.  In this paper, we report on the variability of LMXB candidates measured from multi-epoch {\\it Chandra} X-ray observations of NGC~4697, one of the nearest ($11.3 {\\rm \\, Mpc}$), optically luminous ($M_B < -20$), elliptical (E6) galaxy. In a companion paper (Sivakoff et al 2008b, hereafter Paper IV) we report on the properties of LMXB candidates excluding variability, as well as the connection to globular clusters (GCs) revealed by a {\\it Hubble Space Telescope} Advance Camera for Surveys ({\\it HST}-ACS) observation near the galaxy center.  In \\S~\\ref{sec:n4697x_obs}, we recap key aspects of the observations and data reduction of NGC~4697 discussed in Paper IV. Our observations of NGC~4697 provide five approximately equal chances to detect variability within a given $\\sim 40 {\\rm \\, ks}$ observation (\\S~\\ref{sec:n4697x_src_var_intra}), as well as variability between different observations on timescales of $11 {\\rm \\, d}$ to $4.6 {\\rm \\, yr}$ (\\S~\\ref{sec:n4697x_src_var_inter}). We summarize our conclusions in \\S~\\ref{sec:n4697x_conclusion}. Unless otherwise noted, all errors refer to $1 \\sigma $ confidence intervals, count rates are in the $0.3$--$6 {\\rm \\, keV}$ band, and fluxes and luminosities are in the $0.3$--$10 {\\rm \\, keV}$ band. ",
        "conclusions": ""
    },
    "0806/0806.3004_arXiv.txt": {
        "abstract": "The rapid rotation of Be stars may be caused in some cases by past mass and angular momentum accretion in an interacting binary in which the mass donor is currently viewed as a small, hot subdwarf stripped of its outer envelope. Here we report on the spectroscopic detection of such a subdwarf in the Be binary system FY~Canis Majoris from the analysis of data acquired by the {\\it IUE} spacecraft and KPNO Coud\\'{e} Feed Telescope over the course of 16 and 21 years, respectively. We present a double-lined spectroscopic orbit for the binary based upon radial velocities from the {\\it IUE} spectra and use the orbital solutions with a Doppler tomography algorithm to reconstruct the components' UV spectra. The subdwarf is hot ($T_{\\rm eff} = 45\\pm5$~kK) and has a mass of about $1.3 M_\\odot$ and a radius of about $0.6~R_\\odot$. It contributes about $4\\%$ as much flux as the Be star does in the FUV.  We also present observations of the H$\\alpha$ and \\ion{He}{1} $\\lambda 6678$ emission features that are formed in the circumstellar disk of the Be star.  Orbital flux and velocity variations in the \\ion{He}{1} $\\lambda 6678$ profile indicate that much of the emission forms along the disk rim facing the hot subdwarf where the disk is probably heated by the incident radiation from the subdwarf. A study of the FUV infall shell lines discovered in the 1980s confirms their episodic presence but reveals that they tend to be found around both quadrature phases, unlike the pattern in Algol binaries. Phase-dependent variations in the UV \\ion{N}{5} doublet suggest the presence of a N-enhanced wind from the subdwarf and a possible shock-interaction region between the stars where the subdwarf's wind collides with the disk of the Be star. ",
        "introduction": "%  The component stars of massive interacting binaries experience profound transformations over their lifetimes.  Their evolutionary stages are defined by how mass and angular momentum are shared and lost during the mass transfer stages \\citep*{wel01}.  The originally more massive star grows faster and upon reaching a size comparable to its Roche lobe begins mass and angular momentum transfer to the mass gainer companion.  The orbit shrinks until the mass ratio reverses, but mass transfer will continue with an expanding orbit as long as the donor grows in proportion.  The result will be a donor that is stripped of most of its outer envelope and a gainer that now appears rejuvenated by mass transfer \\citep{dra07}.  Mass transfer may also cause the gainer to spin up to a rotational speed close to the critical value where centripetal and gravitational accelerations are equal at the equator \\citep*{pet05}.  \\citet{pol91} and \\citet{van97} suggest that these mass gainer stars may include some of the rapidly rotating, classical Be stars, B-type stars that display emission lines formed in outflowing, circumstellar disks \\citep{por03}.  There is now a substantial body of evidence that the supports the idea that some Be stars were spun up by mass transfer in interacting binaries.  Most of the known massive X-ray binaries consist of a Be star and neutron star companion, the remains of the former donor star \\citep{coe00}. Donor stars that end up with a mass lower than the Chandrasekhar limit will probably avoid a supernova explosion, but will appear first as a hot subdwarf star and then eventually as a massive white dwarf.  Although there are no known examples of Be plus white dwarf binaries, there are a few identified cases of Be plus hot subdwarf binaries.  The first discovery of the spectrum of a subdwarf companion was made from the analysis of UV spectra from the {\\it International Ultraviolet Explorer (IUE)} spacecraft \\citep{tha95} and the {\\it Hubble Space Telescope} \\citep{gie98} of the well-known Be star $\\phi$~Persei.  The hot companion in \\object{$\\phi$~Per} is an object of solar mass and radius but with an effective temperature of $53$~kK.  However, the hot subdwarf is much fainter than the Be star, and its presence is perhaps best observed through its heating effects on the disk gas nearest to the subdwarf \\citep*{ste00,hum01}. \\citet{mai05} show that the binary 59~Cygni shows many of the same characteristics found in $\\phi$~Per and is probably the second example of a Be plus hot subdwarf system.  A third candidate, \\object{HR~2142} \\citep{pet83}, may also belong this category \\citep{wat91} but at the present the evidence is not compelling.  The fact that the nature of the companions in $\\phi$~Per and \\object{59~Cyg} were only discovered after a century of spectroscopic observation, however, highlights the difficulties of their detection \\citep{gie00}.  Here we report on the identification of a hot subdwarf companion in the Be binary \\object{FY~CMa} (HR 2855, HD 58978).  Balmer emission lines in the Be star primary in this system (B0.5~IVnpe, \\citealt*{hil69}; B0.5~IVe, \\citealt*{sle82}) were first discovered by \\citet{pic05}. Although at first glance it appears to be a classical Be star, it has historically displayed much more spectroscopic variability in its optical and UV spectrum than is typical for Be stars \\citep{sle82,pet82,pet88,gra88,cao01}.  Some of the first {\\it IUE} observations of FY~CMa showed evidence of sharp, red-shifted, ``shell'' spectral lines that varied significantly on a timescale of days \\citep{gra88}.  \\citet{pet88} argued that the time scales of variability indicated a binary origin. Spectroscopic variability with a period of 37.26~d was subsequently found by \\citet{riv04} who showed that the strength and profile of the \\ion{N}{5} $\\lambda 1240$ wind line fluctuates with this period.  \\citet{riv04} also presented observations of the H$\\alpha$ and \\ion{He}{1} $\\lambda6678$ emission lines that suggested that they also varied in radial velocity and strength in much the same way as observed in $\\phi$~Per and 59~Cyg, and suggested that FY~CMa might be generically similar to these systems. However, \\citet{riv04} could not confirm binarity as the orbital elements and the nature of the companion remained unknown.  Below we present an investigation of FY~CMa that is based upon a re-examination of the {\\it IUE} spectra and an analysis of a large collection of optical spectra obtained between 1985 and 2006 (\\S2).  We first discuss the use of cross-correlation measurements of the {\\it IUE} spectra to obtain radial velocities for both components and to derive the orbital elements (\\S3).  We then use the orbital velocity curves to reconstruct the individual UV spectra of the components using a Doppler tomography algorithm, and we compare the reconstructed spectra to models to estimate basic stellar properties (\\S4). We then describe the orbital and temporal variations we find in the H$\\alpha$ and \\ion{He}{1} $\\lambda6678$ emission lines (\\S5).  We argue in \\S6 that the emission line variations are probably due to heating of the part of the disk rim that faces the hot subdwarf companion. We also discuss the long-term and phase-dependent variability in the strength and velocity of the \\ion{Fe}{3} shell-type infall lines that were discovered in the mid-1980s and suggest two possible mechanisms for their formation. Finally, we show that phase-dependent profile, strength, and velocity variability in the \\ion{N}{5} wind lines suggest that the subdwarf's wind is N-enhanced and that there is a significant interaction between the subdwarf's wind and the disk of the Be star. ",
        "conclusions": "%  Our spectroscopic results provide the ingredients to estimate the basic elements of the system geometry.  Here we discuss the system dimensions and the evidence that the hot subdwarf is implicated in local heating of the Be star's disk in much the same way as was found for $\\phi$~Per by \\citet{ste00} and \\citet{hum01} and for 59~Cyg by \\citet{mai05}. We present in Figure~8 a sketch of the inferred system geometry as viewed from above the orbital plane.  For the purpose of this illustration, we assume an orbital inclination of $i=70^\\circ$ and a semimajor axis of $112~R_\\odot$ (where the acceptable range is probably $105 - 115 R_\\odot$ for $i=90^\\circ - 66^\\circ$). We also adopt a Be star radius of $R_1=5.3 R_\\odot$ (where the probable range is $4.1 - 7.6 R_\\odot$ for distances of $333 - 625$~pc). The Be star is represented by the solid disk on the left while the hot subdwarf is shown by the small dot on the right.  The double-lobe surrounding them shows the much larger Roche surface boundary.   We can estimate the half light radius of the H$\\alpha$ emitting disk using the method described by \\citet{gru06}.  This estimate is derived from a simple density power law for the disk, the mean equivalent width ($W_\\lambda = -13.6$ \\AA ), the adopted Be star temperature and disk inclination, and an assumed disk outer boundary equal to the volume-equivalent, Roche radius. These parameters lead to an estimated disk to star radius ratio of $R_d/R_1=9.4$, or a half-light emitting radius of $50 R_\\odot$ for the adopted radius of the Be star. We caution that the simplicity of the models may lead to errors in the half-light radius of $\\pm 30\\%$ and that the models assume an azimuthal symmetry for the emission flux, which is probably not the case for the disk emission of FY~CMa (see below).  Nevertheless, this simple estimate of disk radius is indicated by the lighter gray shaded region on the left side of Figure~8.  \\citet{oka01} argue that the disks of Be stars in binaries with nearly circular orbits are truncated at the 3:1 orbital resonance radius, and this occurs at a distance from the Be star of $52 R_\\odot$ that is only slightly larger than our disk half-light radius estimate.  \\placefigure{fig8}     %  We showed in \\S5 that the \\ion{He}{1} $\\lambda 6678$ emission has an orbital modulation that leads to greatest emission strength when the hot subdwarf is in the foreground.  Furthermore we found that the semiamplitude of the orbital Doppler shift of the emission peaks was larger than that of the subdwarf (from Fig.~7, $K_{em}\\approx 173$ km~s$^{-1}$ $> K_2 = 128$ km~s$^{-1}$).  These facts suggest that a significant fraction of the emission forms in the part of the disk near the subdwarf where the Keplerian velocity of the disk gas is larger than the orbital velocity of the subdwarf. Ignoring three-body effects, we can estimate the radius of the emission source by assuming both the disk gas and the subdwarf obey Keplerian motion around the Be star, $R_{em}/a = (M_1/(M_1+M_2)) ((K_1+K_2)/(K_1+K_{em}))^2 = 0.51$ or $R_{em}=58 R_\\odot$. This radius is about $15\\%$ larger than the half-light radius, which suggests that the emission forms mainly along the outer rim of a disk that may be extended towards the hot subdwarf (see Fig.~8). We suspect that this outer disk region is heated by the combined effects of the incident subdwarf radiation and the impact of the subdwarf's wind on the disk.  We imagine that the rest of the disk also contributes to the line emission flux, and, for example, those times when weak emission is observed (see Fig.~7) may correspond to episodes when the gas density declines in the outer rim facing the subdwarf and the flux from the rest of the disk dominates.  Since the heated disk gas nearest the hot subdwarf will have a higher Keplerian velocity, the center of the heated region will move ahead of the axis joining the stars by an amount that depends on the cooling and orbital timescales.  We can estimate the phase offset $\\triangle \\phi$ of this asymmetry in two ways.  First, the central heated region should reach inferior conjunction about midway between the observed velocity extrema at phases $\\phi=0.77$ and 0.17 (see Fig.~7), i.e., at $\\phi=0.96$. Second, the velocity peak separation is presumably caused by the heated region forming in a sector of a ring.  We will observe the largest velocity separations when the mid-point of the ring segment is oriented along our line of sight.  According to the peak separations plotted in Figure~7, greatest peak separations occur around $\\phi=0.42$ and $\\phi=0.92$.   This suggests that the apex of the heated rim is offset by about 0.06 in phase ($22^\\circ$) ahead of the subdwarf inferior junction at $\\phi=0.0$.  From simple geometric arguments, we expect that the maximum half-peak separation $V_m = {\\rm max}(\\triangle V_r /2)$ is related to the emission semiamplitude by $V_m=K_{em} \\sin \\theta$, where $\\theta$ is half the ring sector opening angle.  From the values of peak separation in Table~5, we estimate $V_m=66$ km~s$^{-1}$ and $\\theta = 22^\\circ$.  These values were used to plot a representative heated region in the darker gray shaded area of Figure~8.  The actual extent of the heated ring is probably larger than shown since the separated emission peaks actually sample a range of ring azimuthal angles and not just the termination angles.  The subdwarf may also affect the kinematics of the gas motion in the disk rim.  We show in Figure~9 the orbital phase variations in the appearance of the \\ion{Si}{3} lines near \\ion{Si}{3} $\\lambda 1299$.  The cores of these lines are shell-type features formed in the disk gas seen in projection against the photosphere of the Be star.  We see evidence that the shell feature broadens from $\\phi=0.5$ to $\\phi=0.0$ and apparently shifts from a redshifted to a blueshifted position at phases surrounding conjunction at $\\phi=0.0$. We imagine that the motion of the outer disk gas is influenced by subdwarf as the disk gas passes and overtakes the star. The  gravitational attraction would cause acceleration and outward motion as the disk gas approaches the subdwarf, but the same forces would cause deceleration and inward motion following closest passage. Thus, we would observe infalling (redshifted) disk gas before $\\phi=0.0$ and then outward moving (blueshifted) gas after $\\phi=0.0$. On the other hand, the disk gas projected against the Be star has more nearly Keplerian and tangential motion at $\\phi=0.5$ leading to the appearance of narrow shell lines at that phase. This kind of perturbation in the Keplerian motions may also partially explain the appearance of a redshifted shell feature in the \\ion{He}{1} $\\lambda6678$ profiles near $\\phi=0.9$ (Fig.~6), an orientation where the disk gas projected against the Be star would be infalling and relatively dense.  \\placefigure{fig9}     %    In addition to the shell features discussed above that predominantly are formed in the disk, another circumstellar component is seen in some intermediately-ionized species (e.g. \\ion{S}{3}, \\ion{Fe}{3}). These variable features, which rarely show velocities less than the photospheric value and can display velocities as large as 50~$\\rm km~s^{-1}$ relative to the photosphere of the Be star, were discovered in {\\it IUE} images obtained in the 1980s \\citep*{pet88,gra88} and they clearly show the presence of significant infall of material toward the Be star. Often times these lines are asymmetric with enhanced absorption on the red side of the profile. \\citet*{gra88} identified six infall episodes that are best seen in the \\ion{Fe}{3} (multiplet~34) line at 1895.456~\\AA, since the feature does not have to compete with a strong photospheric component and it is not blended with interstellar or other circumstellar lines.  From the entire set of 97 {\\it IUE} SWP high resolution (HIRES) images that span more than 15 years of observation and our knowledge of the orbit of the system, we now have more insight on the nature of the infall lines. Their behavior in strength and velocity are shown in Figures~10, 11, and 12. Representative examples of the three basic types of observed features, photospheric, photospheric with additional absorption from the Be star's disk, and infall are given in Figure~10. In order to eliminate phase-dependent effects, we have chosen spectra taken at approximately the same quadrature phase on the {\\it trailing} hemisphere of the Be star but at different epochs. Note that the central core of the profile taken during a period of enhanced disk absorption shows the same velocity as the photospheric line, but that the feature formed in infalling material has a significant redshift. In the upper two-part panel of Figure~11 we plot the residual intensity\\footnote{The observed flux divided by the flux in the local continuum outside of the spectral feature.} of the deepest part of the \\ion{Fe}{3} line and its radial velocity versus observation date. All data are included in the intensity plot, but in the velocity plot we use a different symbol for observations that show no, or minimal shell absorption. Points that fall above the {\\it dashed} horizontal line in the upper panel visually appeared to be mostly photospheric. The double horizontal {\\it dashed lines}  in the second panel delineate the domain of photospheric motion. Features that fall above the double line are clearly associated with infall phases. Seven or eight clear infall episodes are observed. The data suggest that the duration of the infall phases varies from days to a year, but significant data gaps from 1979-85 and 1990-94 render it impossible to determine a firm value or its degree of variability. An apparently extended infall phase of about 300 days was seen from HJD~2446919--2447215. In the lower two panels of Figure~11, the same data are plotted versus orbital phase. From the third panel it can be seen that on the average the shell-type absorption is enhanced in the phase interval 0.3--0.7. But from the display in the fourth panel it is apparent that features seen around phase 0.5 tend to display the systemic velocity, implying they are formed in the disk of the Be star. In the fourth panel core velocities from only the spectra showing enhanced shell absorption are compared with the radial velocity curve of the Be star. It is clear that infall features with velocities up to 50~$\\rm km~s^{-1}$ relative to the photosphere are seen {\\it at most phases} with a minimal occurrence around phase 0.5. The infall features tend to cluster around the quadrature points at which their velocities are highest relative to the Be star's photosphere. Many of the phase dependent strength and velocity variations discussed above can also be seen in the gray-scale plot presented in Figure~12.  The episodic infall of material is fundamentally different from that observed in the Algol-type binaries \\citep{pet01}, where mass motion toward a mass gainer is typically seen only between phases 0.75--0.95. In the case of FY~CMa, the infall features predominantly cluster around the quadrature phases from 0.15--0.35 and 0.65--0.90. The velocities of the shell features around phase 0.5 tend to be either photospheric or negative relative to the Be star\\footnote{Evidence for material outflow is often seen in interacting binaries of the Algol type \\citep{pet01} and similar hydrodynamics might be occuring in this Be + subdwarf system.}. \\citet{pet88} discussed various scenarios to explain the observed simultaneous presence of infall and outflow features in this system including a sudden accretion event, the possible presence of magnetic loops, or flow from a polar jet. \\citet{cao01} reported that infall components in the \\ion{N}{5} doublet were contemporaneously observed with those in \\ion{He}{1} $\\lambda5876$ and also suggested that magnetic loops might be involved, but the phase-dependence of the infall features does not support this speculation. The infall is more likely a consequence of an interaction between the wind from the subdwarf and the periphery of the massive disk of the Be star or disk perturbations due to the tidal force produced by the subdwarf. The motion of a shock-interaction region would result from variations in the location of the interface where there is a balance between the dynamic pressure of the subdwarf's wind and the gas pressure at the periphery of the Be star's disk.  It is not clear at this time whether it is the subdwarf or the Be star that might be driving the variability though. Epochs when no enhanced shell absorption is observed would correspond to minimal wind-disk interaction. Alternatively the episodic infall features could simply be a result of gas dynamics in the disk due to a tidal effect from the subdwarf. Hydrodynamical simulations\\footnote{cf. http://harmas.arc.hokkai-s-u.ac.jp/\\%7Eokazaki/BeX/sim/index.html} of the gas motion in a perturbed disk in a Be/X-Ray binary \\citep{oka02} reveal turbulence in the outer disk and even the formation of spiral structure. Similar disk variability could be operating in FY~CMa. The details are probably complex but an interplay of such gravitationally related forces could create time variable and sometimes coherent gas flows. For example, the tidal deformation of the disk leads to an axial elongation that moves ahead of the motion binary.  Disk gas returning from the extended regions would be observed to be moving inwards when viewed at the quadrature phases.     \\placefigure{fig10}     % \\placefigure{fig11}     % \\placefigure{fig12}     %        Finally, we return to the cyclic variation in the \\ion{N}{5} $\\lambda1240$ wind line that first led \\citet{riv04} to find the orbital period of the binary.  The plot of these variations in Figure~13 shows that the greatest extent of blueshifted absorption occurs around $\\phi=0.0$ when the subdwarf is in the foreground.  We note that such a strong \\ion{N}{5} wind feature is quite unusual for Be stars \\citep*{gra87}, and the orbital phasing suggests instead that the feature may form through absorption of the Be star's flux by the wind of the hot subdwarf when it is in the foreground. The profile would mainly be affected by blueshifted wind absorption since any redshifted absorption component would probably by canceled by redshifted emission in the subdwarf's wind.  This kind of orbital wind variation is not seen in the other major wind line, the \\ion{C}{4} $\\lambda 1550$ doublet, and this may suggest that the subdwarf is N-enriched and C-depleted as expected for a star that has been stripped down to near the CNO-burning core.  One problem with this explanation is that the largest wind speeds observed, $\\approx 400$ km~s$^{-1}$, are much smaller than the wind terminal velocity of a few thousand km~s$^{-1}$ expected for a radiatively driven wind from the subdwarf \\citep{lam99}. However, it may be that the wind optical depths are too small for absorption to occur at higher outflow velocities.  \\placefigure{fig13}     %  In 1995 March a sequence of daily {\\it IUE} SWP HIRES spectra of FY~CMa were obtained over the course of 16 days.  Fortuitously these spectra were centered in time around orbital phase 0.0 and extended from phases 0.784 to 0.213.  The temporal variations in the \\ion{N}{5} $\\lambda1240$ doublet are illustrated in a gray-scale diagram in Figure~14. More detail can be seen from the plot of selected observations shown in Figure~15.  The \\ion{N}{5} features initially showed very little structure until a core with the velocity of the center-of-mass of the system appeared around phase 0.97.  The core, still with no velocity shift, reached its greatest strength at conjunction. It broadened by phase 0.08, and had completely vanished by phase 0.21. The core is probably formed in a shock-heated region between the stars where the N-enhanced wind from the subdwarf collides with the disk of the Be star.    \\placefigure{fig14}     % \\placefigure{fig15}     %    Our work demonstrates that FY~CMa represents the third known case of a Be star with a hot subdwarf companion, the remains of an intense binary interaction that stripped down the mass donor to a tiny fraction of its original mass and left the mass gainer spinning rapidly. The subdwarf secondary in FY~CMa has a mass that is close to the Chandrasekhar limit, and lacking a significant H-envelope, it may represent the progenitor of a supernova type SN~Ibc.  The system is now wide enough that it is unlikely that the Be star will be substantially spun down by tidal effects, so its rapid rotation will be important throughout its subsequent evolution. Such rapidly rotating massive stars may be related to the progenitors of collapsars and gamma-ray bursters \\citep{can07}."
    },
    "0806/0806.1830_arXiv.txt": {
        "abstract": "We have performed new laboratory experiments which gave us the possibility to obtain an estimate of the amount of carbon chain oxides (namely C$_3$O$_2$, C$_2$O, and C$_3$O) formed after irradiation (with 200 keV protons) of pure CO ice, at 16 K. The analysis of laboratory data indicates  that in dense molecular clouds, when high CO depletion occurs, an amount of carbon chain oxides as high as 2-3$\\times$10$^{-3}$ with respect to gas phase carbon monoxide can be formed after ion irradiation of icy grain mantles. Then we have searched for gas phase C$_2$O and C$_3$O towards ten low-mass young stellar objects. Among these we have detected the C$_3$O line at 38486.891 MHz towards the low-mass protostar Elias 18. On the basis of the laboratory results we suggest that in dense molecular clouds  gas phase carbon chain oxides are formed in the solid phase after cosmic ion irradiation of CO-rich icy mantles and released to the gas phase after desorption of icy mantles. We expect that the Atacama Large Millimeter Array (ALMA), thanks to its high sensitivity and resolution, will increase the number of carbon chain oxides detected in dense molecular clouds. ",
        "introduction": "One of the main open questions in astrochemistry is the relation between solid phase and gas phase chemistry in dense molecular clouds. In fact in these regions ice mantles form on silicatic and carbonaceous grains after both direct freeze out of gas phase species and grain surface reactions. The presence of icy grain mantles is indirectly deduced from depletion of gas phase species and directly observed in the infrared from absorption features, attributed to vibrational modes of solid phase molecules, superposed to the background stellar spectrum. Icy grain mantles have been detected both in quiescent regions and in star forming regions. In both environments these suffer from energetic processing due to cosmic ions and UV photons. Ion and UV irradiation cause a modification of the structure and of the chemical composition of grain mantles, that is the formation of molecular species not present in the original ice. After desorption of icy mantles molecular species are released to the gas phase which could be enriched by species formed in the solid phase.   Former laboratory experiments have shown that carbon chain oxides (e.g., C$_2$O, C$_3$O, C$_3$O$_2$, C$_4$O, C$_5$O$_2$, C$_7$O$_2$)  are formed after ion irradiation and UV photolysis of CO-rich ice mixtures \\citep[e.g.,][]{gsetal97, gerakines01, trottier04, loeffler05}  along with carbon dioxide (CO$_2$) which is the most abundant molecule formed \\citep[e.g.,][]{gerakines96, loeffler05}.  Carbon chain oxides, namely dicarbon monoxide (C$_2$O) and tricarbon monoxide (C$_3$O) have been detected in the molecular cloud TMC-1 towards the cyanopolyyne peak (hereinafter TMC-1CP) and it has been estimated that fractional abundance of C$_2$O is about 6$\\times$10$^{-11}$ and that of C$_3$O is about 1.4$\\times$10$^{-10}$. These  abundances have been explained by ion-molecule gas phase reactions \\citep{matthews84, brown85, ohishi91, kaifu04}. C$_3$O has also been extensively searched for \\citep{matthews84, brown85} towards other objects, some of which are rich molecular-line sources and which together encompass a wide range of physical conditions. However C$_3$O was not detected in these regions and only upper limits have been estimated. Recently C$_3$O \\citep{tenenbaum06} as well as other O-bearing species, such as H$_2$O \\citep{melnick01, hasegawa06}, OH \\citep{ford03}, and H$_2$CO \\citep{ford04},  have been detected towards the carbon star IRC +10216. This is an asymptotic giant branch star. O-bearing molecules are not expected in carbon stars since the bulk of available oxygen is contained in CO. In order to explain these observations it has been suggested that the increase in the star luminosity is causing the evaporation of orbiting icy bodies \\citep{melnick01}.   Alternatively it has been suggested that gas-phase oxygen-rich chemistry is occurring in the outer shell of the star \\citep{tenenbaum06}.  Here we discuss new laboratory experiments which confirm the formation of carbon chain oxides after ion irradiation of CO ice at low temperature and give us the possibility to obtain a quantitative estimate of the amount of carbon chain oxides formed with respect to initial carbon monoxide (Section 2). We present the new detection of the C$_3$O line at 38486.891 MHz towards the low-mass protostar Elias 18 and we confirm the detection of the same line towards TMC-1CP as already reported by \\citet{kaifu04} (Section 3). On the basis of our laboratory results we suggest that gas phase carbon chain oxides in dense molecular clouds are actually formed in the solid phase after ion irradiation of CO-rich icy grain mantles and released to the gas phase after desorption of icy mantles (Section 4). ",
        "conclusions": "Observations have shown that in dense molecular clouds the fractional abundance of carbon chain oxides (namely C$_2$O and C$_3$O) is of the order of 10$^{-11}$-10$^{-10}$ \\citep[][this work]{matthews84, brown85, ohishi91, kaifu04}. In these regions the fractional abundance of CO is of the order of 10$^{-4}$ then the abundance of carbon chain oxides with respect to CO values about 10$^{-7}$-10$^{-6}$. Laboratory experiments here presented indicate that after ion irradiation of pure CO ice at 16 K the amount of carbon chain oxides formed is of the order of 2-3$\\times$10$^{-3}$ with respect to initial CO (Figure~\\ref{ratios}). In order to estimate the time necessary to obtain in dense molecular clouds the  effects observed in the laboratory  we consider the approximation of effective monoenergetic 1 MeV protons and assume that in dense interstellar regions the 1 MeV proton flux is equal to 1 proton cm$^{-2}$ s$^{-1}$ (see \\citet{mennella03} for a detailed discussion).  However our experimental results were obtained using 200 keV protons. Thus in order to extrapolate the laboratory results to the interstellar medium conditions we assume that they scale with the stopping power (S, energy loss per unit path length) of impinging ions. Using the TRIM code \\citep{ziegler03} we have estimated that for protons S(200keV)/S(1MeV) is 3.8 in the case of pure CO ice. With these hypotheses in mind we have indicated in Fig.~\\ref{ratios}  timescale axis (top x-axis), which gives an estimation of the time (years) necessary to obtain the  effects, observed in laboratory, on interstellar ices. Thus if we assume high CO depletion and that the carbon chain oxides/CO column density ratio obtained in the solid phase after ion irradiation is maintained in the gas phase after desorption of icy grain mantles, then from the exponential equation used to fit the data, we obtain that about 10$^3$ years would be necessary to form the observed column density of carbon chain oxides. As we will discuss below, this time is much shorter than the evolution time scale of dense clouds, thus the observed gas phase abundance of carbon chain oxides could be easily reached even if carbon monoxide is not completely depleted and/or only partial desorption of icy grain mantles takes place.  Towards all the young stellar objects  observed in this work, the solid CO absorption band at 4.67 $\\mu$m has been detected \\citep[e.g.,][]{kerr93, chiar94, chiar95, teresa98} along with the 3 $\\mu$m band due to water ice which is the most abundant solid phase species along these lines of sight. However, recently, a detailed study of the solid CO band profile observed towards a large sample of low mass embedded objects \\citep{ponto03} has shown that typical lines of sight have 60-90\\% of the solid CO in a pure or nearly pure form, suggesting that  interstellar ices are best represented by a layered model rather than a mixed ice \\citep[e.g.,][]{fraseretal04}. The presence of solid CO towards TMC-1CP has never been reported. In Taurus Molecular Cloud a threshold extinction  A$_{th} \\sim 6$ mag is required for CO ice detection \\citep{chiar95}. Gas-phase models use  A$_V$=10 mag in TMC-1 cores \\citep[e.g.,][]{parketal06} while A$_V$=32 mag towards TMC-1A \\citep{teixeira99} thus it seems reasonable to assume that CO ice is also present along the line of sight of TMC-1CP.  Elias 18 resides in a part of the Taurus molecular cloud known as Heiles cloud 2 (HCL2) where low-mass star formation is active. It is a highly obscured object (A$_V$ $\\sim$ 15-19 mag) and it has been suggested \\citep{tegler95} that it is in transition between an embedded young stellar object and an exposed T Tauri star. The IR spectral energy distribution (SED) for this source is typical of a class II or ``flat-spectrum'' YSO with significant optical extinction \\citep{elias78}. Recent observations indicate that Elias 18 has a circumstellar disk oriented close to edge-on and that most of the CO in the disk is incorporated in icy mantles on dust grains, i.e. high depletion is observed \\citep{shuping01}. Mid-IR observations towards Elias 18 show the presence of both the solid CO and CO$_2$ absorption bands \\citep[e.g.,][]{tielens91, chiar95, nummelin01}. The comparison between observations relative to solid CO and laboratory spectra indicates that different ice mixtures can equally well reproduce the observed band profile \\citep{palustraz93, chiar95, chiar98}. However all the fits indicate that a comparable amount of solid CO is in the nonpolar (i.e., CO-rich) and polar (i.e., H$_2$O-rich) components. Among the fits obtained,  it has been shown that the nonpolar component  can be reproduced by the spectrum of ion irradiated pure CO ice (and this is compatible with the hypothesis that CO-rich icy mantles are present along the line of sight in order to form  carbon chain oxides) and the polar component can be reproduced by the spectrum of CO formed after ion irradiation of a H$_2$O:CH$_3$OH ice mixture \\citep{palustraz93}. Detection of the stretching mode band of solid CO$_2$ towards Elias 18 and the comparison of the observed band profile with laboratory spectra have been reported by \\citet{nummelin01}.  However, as discussed by \\citet{ehrenfreund97, gerakines99, ioppolo08}, the profile of the stretching mode band does not strongly depends on the ice mixture and cannot be used to constrain the ice composition along the line of sight.  TMC-1CP is a dense core in the TMC-1 cold dark cloud. Based on gas phase observations and chemical evolution models it has been deduced that its age is about  10$^5$~years  and the density is n$_H$ = 2$\\times$10$^{4}$ cm$^{-3}$ \\citep{parketal06}. Thus the gas would take 10$^9$/n$_H$=5$\\times$10$^{4}$ years to condense on grains \\citep{xander87}. The presence of gas phase species  implies that  desorption processes, such as photodesorption, grain-grain collisions, cosmic ray induced desorption and turbulence, compete with mantle accretion in this region \\citep{boland82, hasegawa93, bringa04}.  Thus detection of C$_3$O towards these lines of sight  is compatible with the hypothesis that this molecular species is formed in the solid phase and released to the gas phase when desorption of icy mantles takes place.  The results here discussed do not exclude that carbon chain oxides are also formed after gas phase reactions in dense molecular clouds. Furthermore we are aware that further observational data are necessary to confirm these results. In fact we plan to search for carbon chain oxides towards other sources in particular hot corinos in Class 0 low mass protostellar objects where evidence of ice mantle evaporation has been reported \\citep{bottinelli07}.  One of the mysteries of interstellar chemistry is the mechanism regulating the balance between gas phase and solid phase species. Carbon chain oxides could be key molecules in this field and thanks to its high sensitivity and resolution the Atacama Large Millimeter Array (ALMA) will give important results increasing the number of detected features in a larger sample of molecular clouds.   Finally, as far as we know, the detection of C$_2$O and C$_3$O  in comets has never been reported. However comets suffer from heavy ion irradiation \\citep{gsetbob91} and CO is abundant in these objects thus we expect  that carbon chain oxides are present in comets too. In fact a tentative detection of carbon suboxide (C$_3$O$_2$) in comet Halley has been reported \\citep{huntress91, crovisier91}. A firm  detection of carbon chain oxides  in comets could also be used to support the hypothesis that the presence of O-bearing species, and in particular C$_3$O, in carbon stars, such as IRC +10216, is due to sublimation of orbiting icy bodies as suggested by \\citet{melnick01}."
    },
    "0806/0806.3697_arXiv.txt": {
        "abstract": "We describe the construction of a database of extremely metal-poor (EMP) stars in the Galaxy. Our database contains detailed elemental abundances, reported equivalent widths, atmospheric parameters, photometry, and binarity status, compiled from papers in the literature that report studies of EMP halo stars with $\\feoh \\le -2.5$. The compilation procedures for this database have been designed to assemble the data effectively from electronic tables available from online journals. We have also developed a data retrieval system that enables data searches by various criteria and illustrations to explore relationships between stored variables. Currently, our sample includes 1212 unique stars (many of which are studied by more than one group) with more than 15000 individual reported elemental abundances, covering the relevant papers published by December 2007. We discuss the global characteristics of the present database, as revealed by the EMP stars observed to date. For stars with $\\feoh \\le -2.5$, the number of giants with reported abundances is larger than that of dwarfs by a factor of two. The fraction of carbon-rich stars (among the sample for which the carbon abundance is reported) amount to $\\sim$ 30 \\% for $\\feoh \\le -2.5$. We find that known binaries exhibit different distributions of orbital period, according to whether they are giants or dwarfs, and also as a function of metallicity, although the total sample of such stars is still quite small. ",
        "introduction": "Extremely metal-poor (hereafter EMP, defined by $\\feoh \\leq -2.5$ in this paper) stars in the Galaxy carry information about the physical conditions in the early epochs when they were born, and are also unique probes of the production of elements by the first generation stars when the Universe emerged from the so-called dark ages. Analysis of their kinematics also provides direct information on the early stages of galaxy formation (e.g., \\cite{Carollo2007}). The chemical compositions of these stars also impose constraints on the nucleosynthesis pathways involved with the formation of the elements throughout the early history of the Galaxy.  It is a long standing problem whether one can identify the low-mass survivors of the first-generation (Population III) stars, those objects born from primordial clouds containing no elements heavier than lithium. If such stars did form it remains possible that they could be found among the EMP stars, since there are processes (such as binary mass-transfer and/or the accretion of interstellar gas polluted by later generation stars) that could effectively ``disguise'' their true nature by making them appear more metal-rich at present.  Thanks to the recent large-scale searches for candidate Very Metal-Poor stars (hereafter VMP, defined by $\\feoh \\leq -2.0$ according to the nomenclature of Beers \\& Christlieb 2005), in particular by the HK survey \\citep{Beers1985, Beers1992} and by the Hamburg/ESO survey \\citep{Christlieb2008}, the number of known VMP stars has increased dramatically since the 1990s. Approximately $\\sim 1200$ and $\\sim 1500$ stars have been identified as VMP to date, on the basis of medium-resolution spectroscopic follow-up of the $\\sim 6000$ and $\\sim 3600$ candidate in HK-I and HES survey, respectively \\citep{Beers2005b}. This number is likely to expand quickly, as additional VMP stars are identified from ongoing efforts such as the Sloan Digital Sky Survey (in particular from SEGUE: Sloan Extension for Galactic Understanding and Exploration, see http://www.sdss.org). Furthermore, high-resolution spectroscopic observations with 8m-class telescopes such as SUBARU, the VLT, and the KECK telescopes are already beginning to elucidate the detailed abundance patterns of VMP stars.  The abundance analyses of EMP stars provide useful information on Galactic chemical evolution by comparison of their abundance patterns with those of more metal-rich stars with $\\feoh \\gtrsim -2.5$, including the globular cluster stars. At present, only three stars are known with metallicities well below $\\feoh = -4$ (all of which have high-resolution abundance analyses available), while more than 100 stars are known with $\\feoh < -3$, roughly half of which have detailed abundance analyses at present. A salient feature of EMP stars is the sharp cut-off below $\\feoh \\sim -3.5$ in the metallicity distribution function. Other important features of EMP stars are the large fraction of carbon-enhanced stars that are known to exist among them, especially below $\\feoh \\sim -2.5$ \\citep{Rossi1999}, as well as the large scatter in the abundances of neutron-capture elements \\citep{Gilroy1988, McWilliam1995b, Ryan1996, Francois2007}. The lighter elements, such as CNO, as well as the {\\it p}- and $\\alpha$-capture elements and {\\it s}-process elements, are used as tools to explore nucleosynthesis from H- and He-burning resulting from binary mass transfer affected by the evolution of low- and intermediate-mass AGB stars \\citep{Suda2004, Lucatello2006,Komiya2007}. For heavier elements, the abundance patterns of individual EMP stars provide crucial information on the r-process elements produced (presumably) by individual supernova events \\citep{Truran1981,Mathews1990}. Such stars are also used as cosmo-chronometers for placing lower limits on the age of the Universe, based on determinations of the abundances of uranium and thorium \\citep{Sneden1996,Wanajo2002}. A handful of stars that exhibit large enhancements of the r-process elements \\citep{Sneden1994,Hill2002, Frebel2007a} have drawn the interest of researchers concerned with nucleosynthesis processes in massive EMP stars and the chemical evolution of the Galaxy. The determination of the isotopic abundances of \\nucm{6}{Li} and \\nucm{7}{Li} by high-resolution spectroscopy \\citep{Smith1993,Hobbs1994,Asplund2006} also impacts observational constraints on Big Bang nucleosynthesis and the astrophysical origins of these elements.  In order to promote studies such as those described above, and to make them more useful in aggregate (e.g., for statistical studies), it is desirable to construct a modern database of the elemental abundances (and other related properties) of metal-poor stars in our Galaxy. Although the available data on the abundances and properties of EMP stars has been greatly increasing in the past decades, thanks to the many high-resolution spectroscopic studies that have been conducted \\citep{Gilroy1988,Ryan1991,McWilliam1995b,Ryan1996,Fulbright2000,Preston2000,Burris2000, Mishenina2001,Aoki2002b,Cohen2002,Carretta2002,Johnson2002,Nissen2002,Cayrel2004,Honda2004b, Cohen2004,Simmerer2004,Spite2005,Barklem2005,Jonsell2005,Aoki2005,GarciaPerez2006a,Cohen2006, Aoki2007b,Francois2007}, there are no present databases that make these data readily available to astronomers and other researchers in order to conduct their own studies. Generally, it is quite difficult (in particular for the non-specialist) to collect the relevant quantities from the widely scattered literature. In part, this is because the data are presented in individual papers using wide varieties of formats, such as text, tables, and figures. Therefore the compilation of pertinent information requires a great deal of human resources for individual investigators, who would benefit greatly from a more automated compilation.  To develop a more effective set of tools for compilation of data for EMP stars, we have adopted a similar set of methodology for data compilation developed by the Japanese nuclear data group\\footnote{The Stellar Abundances for Galactic Archeology database, SAGA. The database will be available at http://saga.sci.hokudai.ac.jp/.} \\citep{Suda2006b}. We adopt the strategy of Hokkaido University Nuclear Reaction Data Center (JCPRG), which has developed tools for compilation via the internet that alleviate much of the human resources required if the data were input manually from the literature \\citep{Otuka2002}. We differ from the JCPRG approach in that we have adopted a relational database management system for data storage ({\\it MySQL}), rather than a text-based master database. We have also adopted their methods of utilizing the database through the internet by developing the tools to retrieve data and draw summary graphs \\citep{Nouri2002,Otuka2005,Pritychenko2006}.  In this paper we describe the structure of the SAGA database, and present some results based on simple queries of the existing system. Our database enables queries of quantities such as the elemental abundances, photometry, atmospheric parameters, binarity, and position in the Galaxy, and the relationships between them. Thereby, we can begin to consider the characteristics of EMP stars in a statistical sense, and better draw global views of the nature of EMP stars in the Galaxy.  The paper is organized as follows. In \\S 2 we describe the compilation and retrieval system for our database. In \\S 3 we elaborate on the global characteristics of EMP stars in our sample. In \\S 4 we present a brief summary. ",
        "conclusions": "We have constructed the SAGA (Stellar Abundances for Galactic Archeology) database of extremely metal-poor stars in our Galaxy. The compiled data are accessible on the web and are opened to all researchers now. Our database includes information on observational details, abundances, atmospheric parameters, photometry, equivalent widths, and binarity status and periods. These data are taken from published papers, with the use of a web-based system of data compilation equipped with useful tools to convert them from various forms of electronic data tables into CSV format. A data retrieval system has been developed which enables the retrieval and plotting of the data selected according to various criteria.  Our sample includes \\nobj\\ stars with distinct object names, roughly half of which are giants. The number of giants becomes twice as large as that of dwarfs if we consider only stars with $\\feoh < -2.5$. The fraction of carbon-enhanced stars ([C/Fe] $\\geq$ +0.5) amounts to $\\sim 30\\%$ among the sample of stars with derived carbon abundance for $\\feoh < -2.5$. The sample stars exhibit a bimodal distribution of V band magnitudes, which is ascribed to the different coverage of effective magnitude range among the large-scale surveys of metal-poor stars. There may exist different distributions of binary periods among the stars with this information available. It is shown for stars with $\\feoh \\lesssim -2.5$ that the binaries with a giant member have typically longer periods than those with a dwarf member, and that there are no dwarfs in binaries having periods of $> 1000$ days yet confirmed. Considering the spatial distribution, our sample may have some biases for the discussion of the properties of metal-poor stars because of the different sampling volumes for dwarfs and giants. In fact, we only have detailed elemental abundances for dwarfs within $\\lesssim$ 5 kpc from the Sun, while giants cover distances extending to more than $\\gtrsim$ 20 kpc in the current sample.  Since the EMP stars in our Galaxy are useful probes for our understanding of the chemical and formation history of our Galaxy, large increases in such data are desired by observers and theoreticians alike. A number of observing projects are planned to increase the stellar sample. For example, the stellar extension program of the Sloan Digital Sky Survey, SDSS/SEGUE, is obtaining medium-resolution spectra from which additional EMP stars may be selected to a depth of up to $\\sim$ 100 kpc \\citep{Beers2004}. LAMOST \\citep{Zhao2006} is a multi-fiber 4m telescope project that will enable up to 4000 stellar spectra to be obtained simultaneously in each exposure. The total survey effort is planned to encompass several million stars. These projects will increase the number of candidate EMP stars by several orders of magnitude in the near future, as compared with the number of known EMP stars known at present. High-resolution spectroscopic follow-up will be required, making use of dedicated programs, such as the proposed WFMOS effort on the Subaru telescope, and with the next generations of Extremely Large Telescopes, with diameters of 30m or more.  It is important for us to understand how large the discrepancies are caused by the independent analyses. In order to check the difference among the derived abundances by different authors, we pick up 17 stars from our sample and compare their derived carbon abundances. The 13 stars among them are retrieved from the sample for which more than 8 papers report the abundance analysis. The 12 stars of them are giants and do not show carbon enhancement. The remaining 4 stars are added to cover the various EMP stars. They are carbon-enhanced stars of giant and dwarf (CS 22948-027 and CS 22898-027, respectively), an extremely iron-poor star having $\\feoh \\sim -4$ (CD-38$^{\\circ}$245), and normal giants for which carbon abundance is reported by more than 5 papers (CS 22169-035). In Figure~\\ref{fig:analyses}, we show the deviations from average values of $\\log g$, $\\teff$, and $\\feoh$ as a function of those of $\\cfe$. It should be noted that all reported carbon abundance is based on 1D LTE model atmosphere and most of the analyses adopt the synthetic spectral technique using CH G band. The majority of the plotted stars are located within the 0.2 dex for $\\cfe$ value, which is well explained in terms of the errors associated with the different values of atmospheric parameters and of the usage of different solar abundance from paper to paper, the latter of which can be important for CNO abundance. Some of the large deviations of the adopted or derived values can also be explained by the analyses based on low-resolution spectra (for example, the case of CS 22898-027), although, the reasons for different results are not necessary obvious for all cases. As can be seen in the left panels of Fig.~\\ref{fig:analyses}, the large differences are highly correlated with the deviations of the adopted values of atmospheric parameters, the latter of which is due to the different way of analyses and corrections for the determinations of surface gravity and effective temperature. In fact, for CS22948-027, the largest discrepancy of atmospheric parameters and metallicity in the figure seems to be caused by both the different method and correction of them.  \\begin{figure} \\begin{center} \\FigureFile(\\textwidth,){figure10.eps} \\end{center} \\caption{Consistency check of the independent analyses for selected 17 stars from SAGA database. Each value denotes the deviation from the average value of adopted or derived quantity in each object for which it is reported. The symbol \"c\" and \"d\" in the legend of object names denotes the ``carbon-rich'' and ``dwarf'', respectively. Othere stars without symbols are giants without carbon enhancement. Note that some of the papers adopt the same set of atmospheric parameters and abundances. Such data are completely overlap with each other in this figure. Note also that the data without values reported by authors do not apper in the figure, which is sometimes the case for $\\cfe$. The majority of the data points are located within the typical errors as enclosed by auxiliary dashed lines in top right panel. }\\label{fig:analyses} \\end{figure}  Accordingly, users should be warned about the possible discrepancy of independent analyses when they use the combined data derived by different authors. At any situations in using our database, users can go back to the original papers at the data retrieval system and check the information on analyses and discussion. The extreme case of discrepancies, if happens, will be discussed in the latest original paper by comparing with the previous works. Therefore, we will not continue to discuss here about the systematic differences between previous works for all objects and elements in the database. For the abundance deternimations with non-LTE scheme or with 3D model atmospheres, their effect for extremely metal-poor stars should also be mentioned in considering the different analyses in more detail. However, it is beyond the scope of this paper and is discussed in the other extensive works for non-LTE abundances (see, e.g, \\cite{Andrievsky2007}) and for 3D model atmosphere \\citep{Asplund2001,Collet2007}.  At present, we are planning to include the information on the analyses adopted by the authors and to implement the option of choosing the LTE or NLTE abundances. On the other hand, It is desirable to improve the quality of data by creating a homogenized dataset that enables us to refine the statistical analysis of abundance trends. For users of interest, we can provide the compiled data related to equivalent width and other necessary data for their re-analysis of the sample.  We plan to continually update the SAGA database with updates as new papers reporting on high-resolution spectroscopic follow-up appear in the literature. We also plan to continue an effort to provide more complete coverage of existing data, by supplementing the SAGA database with stars of higher metallicity, and by extending the temporal coverage to circa 1990. In forthcoming papers, we plan to use the updated SAGA database to discuss more thoroughly the abundance trends of EMP stars, and compare them with theoretical models.\\\\  We are grateful to T. C. Beers for reading the manuscript and for giving helpful suggestions and comments including the denomination of the database. We thank S. Lucatello for kindly providing the abundance data of carbon, nitrogen, and iron in the Hamburg/ESO R-process Enhanced Star (HERES) survey sample. We are also grateful to the anonymous referee for his/her suggestion about the influences of independent analyses. This research has made use of the ADS database, operated at SAO/NASA, USA, mirrored by NAOJ, Japan, and SIMBAD and VizieR database, operated at CDS, France. This work has also made use of the observations with low- to high-dispersion spectroscopy by the optical telescopes all over the world. This work has been partially supported by Grant-in-Aid for Scientific Research (15204010, 18104003, 19740098), from Japan Society of the Promotion of Science."
    },
    "0806/0806.4248_arXiv.txt": {
        "abstract": "We report on the detection of very stable modulations with periods unexpectedly ($\\sim$0.5 \\%) longer than superhump periods during the slowly fading stage of WZ Sge-type superoutbursts in three systems, GW Lib, V455 And and WZ Sge.  These periods are naturally explained by assuming that these modulations are superhumps arising from matter near the tidal truncation radius.  This finding provides an additional support to the hypothetical idea of expansion of the accretion disk well beyond the 3:1 orbital resonance in some low mass-ratio systems.  Combined with the effect of 2:1 resonance, we present an explanation of the origin of positive period derivatives in certain short-period SU UMa-type dwarf novae. ",
        "introduction": "Dwarf novae (DNe) are a class of cataclysmic variables (CVs), which are close binary systems consisting of a white dwarf and a red-dwarf secondary transferring matter via the Roche-lobe overflow. SU UMa-type dwarf novae are a class of DNe, which show superhumps during their long, bright outbursts (superoutbursts) [see e.g. \\citet{vog80suumastars}; \\citet{war85suuma} for basic observational properties].  The origin of superoutbursts and superhumps in SU UMa-type dwarf novae is basically understood as a consequence of thermal and tidal instabilities in the accretion disk (\\cite{osa89suuma}; \\cite{osa96review}), the latter being excited by the 3:1 orbital resonance in the disk (\\cite{whi88tidal}; \\cite{hir90SHexcess}; \\cite{lub91SHa}). The basic observational properties of ordinary SU UMa-type dwarf novae have well been reproduced by this picture.  WZ Sge-type dwarf novae (see e.g. \\cite{bai79wzsge}; \\cite{dow90wxcet}; \\cite{kat01hvvir}) are a subgroup of dwarf novae characterized by large-amplitude (typically $\\sim$ 8 mag) superoutbursts with very long (typically $\\sim$ 10 yr) recurrence times.  Although WZ Sge-type dwarf novae are recognized as a subgroup of SU UMa-type dwarf novae, WZ Sge-type dwarf novae are known to have a number of properties hardly, but not necessarily exclusively, observed in ``textbook'' SU UMa-type dwarf novae.  These properties include: (1) the presence of early superhumps \\citep{kat02wzsgeESH}, (2) (sometimes repetitive) rebrightenings \\citep{kat04egcnc}, (3) nearly constant to positive period derivative ($P_{\\rm dot} = \\dot{P}/P$) of superhumps (\\cite{kat01hvvir}; \\cite{kat03v877arakktelpucma}), and (4) long-lasting fading tails.  The implication of phenomenological relations between some of these properties was first addressed by \\citet{kat98super} [see also \\citet{kat04egcnc} for more discussions].  \\citet{kat98super} presented an idea that the accretion disk can expand beyond the 3:1 resonance during energetic outbursts in low mass ratio ($q = M_2/M_1$) systems exemplified by WZ Sge-type dwarf novae. They argued that the matter beyond the 3:1 resonance can serve as a reservoir supplying matter to the inner disk resulting rebrightenings, and the eccentricity wave propagating outward the 3:1 resonance can explain positive $P_{\\rm dot}$ of superhumps.  This idea thus has a possibility to naturally explain many of peculiar properties of WZ Sge-type dwarf novae.  \\citet{hel01eruma} further introduced an idea of decoupling of thermal and tidal instabilities beyond the 3:1 resonance, and extended the application to ER UMa-type dwarf novae, another subclass of SU UMa-type dwarf novae with similarly low $q$.  In line with these ideas, \\citet{osa02wzsgehump} and \\citet{osa03DNoutburst} presented an overview of an WZ Sge-type outburst based on the disk-instability model of SU UMa-type dwarf novae (\\cite{osa89suuma}; \\cite{osa95wzsge}), and proposed a new conceptual scheme of classifying SU UMa-type dwarf novae based on $q$ and achievable disk radius. \\citet{osa03DNoutburst} also explored the dependence of outburst properties on the matter reaching beyond the 3:1 resonance, and presented a scheme of understanding a variety of superoutbursts. The expanded disk beyond the 3:1 resonance in WZ Sge-type dwarf novae and related objects has thus been favored by theoretical sides. Observational evidence, however, for such an expanded disk had long been rather scarce [see \\citet{kat04egcnc} for description of historical observations], while the recent discovery of different $P_{\\rm dot}$ of superhumps in a variety of superoutburst in the same system \\citep{uem05tvcrv} well matched the scenario by \\citet{osa03DNoutburst}, thus strengthening the expanded disk beyond the 3:1 resonance. The infrared excess during the late stage of WZ Sge-type outbursts (e.g. \\cite{uem08j1021}; \\cite{uem08alcom}) also supports this idea.  In 2007, two spectacular superoutbursts of WZ Sge-type dwarf novae, namely GW Lib and V455 And occurred (\\cite{waa07gwlibiauc}; vsnet-alert 9530).  During the later course of these outbursts, we discovered ``late superhumps'' with unexpectedly long periods and with exceptionally high coherence and stability in their periods.  In this letter, we present an interpretation of these late superhumps originating from the very matter beyond the 3:1 resonance. ",
        "conclusions": ""
    },
    "0806/0806.3973_arXiv.txt": {
        "abstract": "A novel way of looking at the evolution of star clusters is presented. With a \\textit{dynamical temperature}, given by the mean kinetic energy of the cluster stars, and a \\textit{dynamical luminosity}, which is defined as the kinetic energy of the stars leaving the cluster in analogy to the energy of photons emitted by a star, the dissolution of star clusters is studied using a new \\textit{dynamical temperature-luminosity diagram} for star clusters. The investigation contains a parameter-space study of open clusters of up to $N=32768$ single-mass stars with different initial density distributions, half-mass radii, tidal conditions and binary fractions. The clusters show a strong correlation between \\textit{dynamical temperature} and \\textit{dynamical luminosity} and most of the investigated cluster families share a common sequence in such a \\textit{dynamical temperature-luminosity diagram}. Deviations from this sequence are analyzed and discussed. After core collapse, the position of a cluster within this diagram can be defined by three parameters: the mass, the tidal conditions and the binary fraction. Due to core collapse all initial conditions are lost and the remaining stars adjust to the given tidal conditions. Binaries as internal energy sources influence this adjustment. A further finding concerns the Lagrange radii of star clusters: Throughout the investigated parameter space nearly all clusters show a constant half-mass radius for the time after core collapse until dissolution. Furthermore, the ratio of half-mass radius to tidal radius evolves onto a common sequence which only depends on the mass left in the cluster. ",
        "introduction": "\\label{sec:intro} In 1913 Henry Norris Russell presented his work on a relation between the spectral classes of stars and their absolute magnitude at a meeting of the Royal Astronomical Society \\citep{Ru13}. The diagram he showed later became famous as the Hertzsprung-Russel Diagram (Ejnar Hertzsprung was the first to anticipate the existence of a relation between the two quantities) and became one of the most important tools for the study and understanding of stellar evolution.  The temperature-luminosity diagram, which is a derivative of the original Hertzsprung-Russel diagram, shows a tight relation between a star's temperature and its luminosity for the main-sequence phase of stars.  A star cluster is a system which shows some analogies to a star. First of all, the stars within a cluster follow a velocity distribution, which is established through two-body encounters, just like a gas or plasma does through collisions of particles. Therefore it is possible to assign a \\textit{dynamical temperature}, $T$, to the stars in a cluster (Sec.~\\ref{sec:T}). Secondly, a star cluster constantly loses a certain fraction of its stars through escape, like a star constantly emits photons. It thus appears plausible to define a \\textit{dynamical luminosity}, $L$, in terms of the energy carried away by the stars per unit time (Sec.~\\ref{sec:L}).  To find if a dynamical temperature-luminosity relation exists for dynamical systems like star clusters is the motivation of this work, because it is likely that it would prove very useful for describing global cluster properties and evolution. Using numerical simulations, the behaviour of star clusters within a \\textit{dynamical T-L diagram} (Sec.~\\ref{sec:TL}) and the influence of different initial conditions on its development (Sec.~\\ref{sec:var}) is studied.  Throughout the performed parameter-space study the half-mass and tidal radius of the clusters are investigated in detail, as they turned out to not behave as expected. The half-mass radius stays constant for a large fraction of a cluster's life-time and the ratio of half-mass radius to tidal radius evolves along a single sequence, independent of initial conditions.  But first of all a reference model, the ``\\textit{Standard Cluster}'', is defined and investigated in detail (Sec.~\\ref{sec:scl}), which will help establishing a \\textit{dynamical temperature} and a \\textit{dynamical luminosity}. ",
        "conclusions": "The investigations made here show that it is possible to define a \\textit{dynamical temperature} and a \\textit{dynamical luminosity} for star clusters. The \\textit{dynamical temperature-luminosity diagram} established with these two quantities gives a completely new way of looking at cluster evolution and helps to gain valuable insights on the energy evolution of star clusters.  The \\textit{dynamical temperature} is defined through the mean kinetic energy of the stars within a cluster, where binary systems need to be replaced by their centre-of-mass particles. Through this correction the system can still be treated as an ideal gas, which enables to relate the temperature to the velocity dispersion of the stars within the cluster. In this way the \\textit{dynamical temperature} is directly correlated to the velocity distribution of the cluster stars, which is taken to be Maxwellian shaped. The latter implies that there is always a high-velocity tail of stars with velocities above the escape velocity. When these unbound stars leave the cluster, the left-over stars will reestablish a velocity dispersion with a lower \\textit{dynamical temperature} within a relaxation time. A relation between \\textit{dynamical temperature} and the number of escaping stars is therefore expected, since there will always be a certain fraction of stars leaving the cluster.  The \\textit{dynamical luminosity} is defined as the kinetic energy of the stars leaving the cluster, in analogy to the energy of photons emitted by a star. Escaping stars show two different origins, evaporation and ejection, which have to be treated separately. The former is strongly correlated to the temperature, while the latter is due to binary interactions and hence a tracer for structural effects like core collapse or core oscillations. In fact, there is no ejection before core collapse and the occurrence of the first ejected star can be used to define the point in time when the core has reached its densest phase. The limit between \\textit{evaporaters} and \\textit{ejecters} was found to be universal among the investigated clusters at $E_{kin}/(kT) = 5.1 \\pm 0.2$. These two families of escapers are furthermore the reason for splitting up the total luminosity into an \\textit{evaporation luminosity} and an \\textit{ejection luminosity}.  In this way the energy evolution of a star cluster can be understood as follows: A cluster is generating energy in the core through binary burning which causes expansion of the whole cluster. Without a tidal field this leads to infinite expansion, since a single binary can generate large amounts of energy until it is ejected from the cluster, in which case the core has to contract to form a new one. Stars leaving the cluster have to have positive energy, which they can gain through two-body relaxation processes. This resulting evaporation of stars decreases the total energy of the cluster. Ejection mechanisms have to counteract this decrease of energy. Three- or more-body encounters in the core cause not only the ejection of stars with high velocities, but also the ejection of binaries. The latter therefore dissipate the accumulated binding energy and make the total energy of the cluster go to zero (see Fig.~\\ref{E}). A tidal field is setting a limit to the expansion of the cluster, such that the cluster spills over the tidal radius. Evaporation is therefore amplified, since escape is eased.  All computed clusters show a tight relation between \\textit{evaporation luminosity} and \\textit{dynamical temperature}, while the \\textit{ejection luminosity} shows a much larger scatter due to the small number of \\textit{ejecters}, i.e. insufficient statistics. Before core collapse all clusters move towards a \\textit{dynamical main sequence}, where they spend up to half of their total life-time. After core collapse they follow a common cooling sequence, similar to a cooling track of white dwarfs. The duration of the main sequence phase depends on the initial density distribution of the cluster. For initially very extended clusters, i.e. with large initial $R_h/R_{tide}$, core collapse takes a large fraction of the cluster's total life time, during which the temperature stays nearly constant, corresponding to a fixed position in the \\textit{dynamical temperature-luminosity diagram}.  This leads to a large degeneracy in models for equal-mass clusters, since the initial conditions of a cluster on the cooling sequence cannot be traced back. For the \\textit{evaporation luminosity} the only deviations from this cooling sequence are given by models with different tidal conditions, i.e. at different Galactic radii, while the \\textit{ejection luminosity} is additionally affected by the binary fraction. On the other hand this supports the hypothesis that after core collapse, the state of a single-mass system can be fully described by three quantities: the number of stars left in the cluster, the tidal conditions and the binary content. A theory of cluster evolution therefore has to focus on these parameters.  Whether clusters with a stellar mass distribution would also form a \\textit{dynamical main sequence} still has to be investigated. If this is the case, well observed clusters like the Pleiades may be placed within such a diagram and a unique evolutionary track may be assigned to them, which would give direct insights on the former and further evolution of those systems.  Another important finding of this parameter-space study is the phenomenon of a constant half-mass radius for clusters in tidal fields. This holds for the time after core collapse, when a final density distribution is established within the tidal radius, in which the energy production in the core is balanced by the mass loss at the tidal radius. As shown in this work, the tidal field has a significant influence on this value, since clusters at larger Galactic radii show a much larger equilibrium value for the half-mass radius, while isolated clusters do not show this phenomenon at all. This can be expressed in terms of the ratio of half-mass radius to tidal radius, which evolves along a common sequence for all investigated clusters, depending only on mass. The ratio is increasing with time, as the half-mass radius is a constant and the tidal radius is decreasing due to ongoing mass loss. This implies that the energy production efficiency within the half-mass radius is increasing as the cluster loses mass.  The \\textit{Standard Cluster} shows an equilibrium half-mass radius of about 2 pc in the given tidal field, and so do the other models with the same initial mass, tidal conditions and binary fractions (Tab.~\\ref{table1}). An increasing initial mass increases the value of the equilibrium half-mass radius. The models with a higher mass than the \\textit{Standard Cluster} have a slightly declining half-mass radius at the end of their lifetimes, when their ratio of $R_h/R_{tide}$ has reached the common sequence. While the model with 16384 stars still shows a nearly constant $R_h$, the 32k model has a slightly increasing half-mass radius after core collapse until about half its dissolution time, after which it starts decreasing slowly. In a first order approximation it still can be described by a constant but it is expected that this gets more and more inexact the larger the initial mass of the cluster is.  This behaviour contradicts the theory of self-similar evolution of \\citet[p.~59]{Spitzer87}, who found the ratio $R_h/R_{tide}$ to be a constant of about 3, which seems to be a too rough estimate and just holds at the very end of a cluster's life-time. Furthermore, the simple cluster-evolution theory based on the ansatz that evaporation does not change the energy of a cluster and thus the half-mass radius scales with $M^2$ \\citep[p.~525]{Bi87}, can be completely ruled out. This only holds for very extended clusters, for which the initial half-mass radius is larger than the equilibrium value, and for these clusters this ansatz holds only until core collapse because until then no ejection has occurred. In this given set of parameters the constant half-mass radius gives a much better approximation. Nevertheless, this parameter-space study showed how similar single-mass clusters with a wide range of initial conditions evolve, once they have adjusted to the given tidal conditions, which is in most cases right after core collapse.  The discovery of Russell and Hertzsprung was made possible through improvements in distance-determination techniques for stars, just like future developments like GAIA will increase the possibilities of measuring accurate peculiar velocities of stars in the Milky Way. Then, also a large sample of open clusters will be surveyed in detail, which will give the opportunity to reconstruct internal quantities like the velocity dispersion very accurately and also to identify stars which are about to leave a cluster or have left it lately. This means that the actual measurement of the two quantities, \\textit{dynamical temperature} and \\textit{dynamical luminosity}, will become feasible.    \\begin{table*} \\begin{minipage}{158mm} \\centering \\caption{Overview of all computed models. $n$ gives the number of computed models of a particular kind, $\\rho(r)$ is the initial density profile (P: Plummer, K: King) and $R^f_h$ is the fitted value for the half-mass radius after core collapse. Given errors are the standard deviations. The last column gives the fitted values for the limit between \\textit{evaporaters} and \\textit{ejecters} with fitting uncertainties.} \\label{table1} \\begin{tabular}{ccccccccccccc} \\hline $n$ & $N$& $\\rho(r)$  & $R_h$ [pc]  & $f_{bin}$ & $R_{gal}$ [kpc]  & $t_{cc}$ [Myr] & $t_{rel}^0$ [Myr]&$t_{dis}$ [Myr]& $R^f_h$ [pc]& $\\left(\\frac{E_{kin}}{kT}\\right)_{limit}$\\\\ \\hline 27&1000&P&0.8&0&8.5&165$\\pm$20&10.1&2277$\\pm$117&2.01$\\pm$0.09&5.5$^{+1.0}_{-0.9}$\\\\ \\hline 5&1000&K&0.8&0&8.5&162$\\pm$37&10.1&2198$\\pm$120&1.86$\\pm$0.04&5.4$^{+0.8}_{-0.7}$\\\\ \\hline 5&1000&P&0.4&0&8.5&62$\\pm$5&3.6&2039$\\pm$85&1.96$\\pm$0.12&5.7$^{+1.3}_{-1.1}$\\\\ 5&1000&P&1.6&0&8.5&432$\\pm$50&28.7&2467$\\pm$159&2.15$\\pm$0.09&4.8$^{+0.8}_{-0.7}$\\\\ 9&1000&P&2.4&0&8.5&758$\\pm$47&52.6&2573$\\pm$175&2.18$\\pm$0.17& 7.3$^{+8.8}_{-4.0}$ \\\\ 5&1000&P&3.2&0&8.5&1081$\\pm$98&81.0&2592$\\pm$173&2.09$\\pm$0.09&5.9$^{+3.5}_{-2.2}$\\\\ 10&1000&P&4.0&0&8.5&1061$\\pm$116&113.3&2489$\\pm$94&2.03$\\pm$0.13&4.3$^{+1.1}_{-0.9}$\\\\ \\hline 33&500&P&0.8&0&8.5&134$\\pm$32&8.4&1662$\\pm$174&1.88$\\pm$0.12&6.7$^{+1.4}_{-1.1}$\\\\ 8&2000&P&0.8&0&8.5&191$\\pm$27&12.5&3195$\\pm$150&2.13$\\pm$0.05&4.8$^{+0.9}_{-0.8}$\\\\ 5&3000&P&0.8&0&8.5&241$\\pm$11&14.2&3871$\\pm$101&2.29$\\pm$0.03&4.5$^{+1.3}_{-1.0}$\\\\ 4&4000&P&0.8&0&8.5&253$\\pm$19&15.6&4563$\\pm$316&2.38$\\pm$0.14&4.3$^{+1.1}_{-0.9}$\\\\ 3&5000&P&0.8&0&8.5&290$\\pm$9&16.9&4943$\\pm$180&2.48$\\pm$0.13&4.4$^{+1.7}_{-1.2}$\\\\ 1&16384&P&0.8&0&8.5& 442 & 25.7 &  10077 &  2.99 & 3.8$^{+1.6}_{-1.1}$\\\\ 1&32768&P&0.8&0&8.5& 1093 & 33.3 & 15645 &  3.57  & 3.6$^{3.2}_{-1.7}$\\\\ \\hline 5&1000&P&0.8&0&85&193$\\pm$19 &10.1&$>$20000&8.69$\\pm$0.81&5.7$^{+0.8}_{-0.7}$\\\\ 5&1000&P&0.8&0&$\\infty$&146$\\pm$31&10.1&$>$20000&(--)\\footnote[1]{Isolated clusters do not show a constant half-mass radius.}&4.3$^{+0.6}_{-0.5}$\\\\ \\hline 4&1000&P&0.8&0.95&8.5& (--)\\footnote[2]{The presence of primordial binaries prevents the core from collapsing.} &7.8&2012$\\pm$232&2.28$\\pm$0.14&5.7$^{+0.7}_{-0.6}$\\\\ \\hline \\end{tabular} \\end{minipage} \\end{table*}"
    },
    "0806/0806.4554_arXiv.txt": {
        "abstract": "A model which leads to abundant antimatter objects in the Galaxy (anti-clouds, anti-stars, etc) is presented. Observational manifestations are analyzed. In particular, the model allows for all cosmological dark matter to be made out of compact baryonic and antibaryonic objects. ",
        "introduction": "The origin of the observed excess of matter over antimatter in the universe is believed to be pretty well understood now. As formulated by Sakharov~\\cite{ads}:\\\\ 1) nonconservation of baryonic number,\\\\ 2) breaking of C and CP, and\\\\ 3) deviation from thermal equilibrium \\\\ lead to different cosmological abundances of baryons and antibaryons. The cosmological baryon asymmetry is characterized by the dimensionless ratio of the difference between the number densities of baryons and antibaryons to the number density of photons in the cosmic microwave background radiation: \\be \\beta =  \\frac{n_B - n_{\\bar B}}{n_\\gamma} \\approx 6\\cdot 10^{-10} \\label{beta} \\ee There are many theoretical scenarios which allow to ``explain'' this value of the baryon asymmetry, for the review see~\\cite{bs-rev}. Unfortunately ``many'' means that we do not know the single one (or several?) of the suggested mechanisms which was indeed realized. Usually in such cases experiment is the judge which says what is right or wrong. However, it is impossible to distinguish between competing mechanisms having in one's disposal only one number, the same for all the scenarios. We would be in much better situation if $\\beta$ is not a constant over all the universe but is a function of space point, $\\beta = \\beta (x)$. So it is interesting to study the mechanisms which might lead to space varying $\\beta$ and especially, in some regions of space, to $\\beta < 0$, i.e. to possible generation of cosmological antimatter.  There is an increasing experimental activity in search for cosmic antimatter. In addition to the already existing detectors, BESS, Pamella, and AMS, a few more sensitive ones shall be launched in the nearest years, AMS-02 (2009), PEBS (2010), and GAPS (2013), see the review talk~\\cite{picozza} at TAUP 2007. To the present time no positive results indicating an astronomically significant cosmic antimatter have been found but still the bounds are rather loose and as we see in what follows, it is not excluded that the amount of antimatter in the universe may be comparable to that of matter and astronomically large antimatter objects can be in our Galaxy quite close to us.  If this is the case, one should search and may hope to observe cosmic antinuclei starting from $ ^4He $ to much heavier ones, excessive antiprotons and positrons, flux of energetic gamma rays with energies about 100 MeV from $p \\bar p$--annihilation and 0.511 MeV from $e^- e^+$--annihilation, violent phenomena from antistars and anticlouds, and some other more subtle ones.  We cannot say, of course, if there is any reasonable chance to find all that, but at least there is a simple theoretical model according to which galaxies, including the Galaxy, though possibly dominated by matter, may include astronomically significant clumps of antimatter on the verge of possible detection.  This talk consists of the following two main parts: \\\\ I. The mechanism of the antimatter creation leading to considerable amount of antimatter in the Galaxy in the form of compact objects or clouds. \\\\ II. Antimatter phenomenology, observational signatures, and bounds.\\\\   The talk is based on several papers written in collaboration with C. Bambi,  M. Kawasaki, N. Kevlishvili, and J. Silk~\\cite{ad-js,cb-ad,ad-mk-nk}, where a detailed discussion and more complete list of references can be found. ",
        "conclusions": "{1. The Galaxy may possess a noticeable amount of antimatter. Both theory and observations allow for that.}\\\\ 2. Theoretical predictions are vague and strongly model dependent.\\\\ {3. Not only ${ ^4 \\bar{He}}$ is worth to look for but also heavier anti-elements. Their abundances should be similar to those observed in SN explosions.} \\\\ {4. The regions with anomalous abundances of light elements suggest that they consist of antimatter.}\\\\ {5. A search of cosmic antimatter has non-vanishing chance to be successful.}\\\\ 6. Dark matter made of BH, anti-BH, and dead stars is a promising candidate. There is a chance to understand why ${ \\Omega_B =0.05}$ is similar by magnitude to ${ \\Omega_{DM} = 0.25}$."
    },
    "0806/0806.0377_arXiv.txt": {
        "abstract": "Measurements of CMB temperature fluctuations by the Wilkinson Microwave Anisotropy Probe (WMAP) indicate that the fluctuation amplitude in one half of the sky differs from the amplitude in the other half.  We show that such an asymmetry cannot be generated during single-field slow-roll inflation without violating constraints to the homogeneity of the Universe.  In contrast, a multi-field inflationary theory, the curvaton model, can produce this power asymmetry without violating the homogeneity constraint.  The mechanism requires the introduction of a large-amplitude superhorizon perturbation to the curvaton field, possibly a pre-inflationary remnant or a superhorizon curvaton-web structure.  The model makes several predictions, including non-Gaussianity and modifications to the inflationary consistency relation, that will be tested with forthcoming CMB experiments. ",
        "introduction": "Inflation provides a compelling description of the early Universe \\cite{Guth:1980zm}.  The temperature fluctuations in the cosmic microwave background (CMB) \\cite{deBernardis:2000gy,Dunkley:2008ie} and the distribution of galaxies \\cite{Cole:2005sx} agree well with inflationary predictions.  However, there is an anomaly in the CMB: measurements from the Wilkinson Microwave Anisotropy Probe (WMAP) \\cite{Dunkley:2008ie} indicate that the temperature-fluctuation amplitude is larger, by roughly 10\\%, in one hemisphere than in the other \\cite{Eriksen:2003db}.  Fewer than 1\\% of simulated isotropic fluctuation maps exhibit such an asymmetry, and the asymmetry cannot be attributed to any known astrophysical foreground or experimental artifact.  As opposed to the ``axis of evil'' \\cite{deOliveiraCosta:2003pu}, an apparent alignment of only the lowest multipole moments, this asymmetry has gone largely unnoticed (although see \\cite{Gordon:2006ag,Donoghue:2007ze}), and it warrants further theoretical consideration.  \\begin{figure}[b] \\includegraphics[width=8.5cm]{Fig1.eps} \\caption{Measurements of temperature fluctuations in the cosmic microwave background (CMB) show that the rms temperature-fluctuation amplitude is larger in one side of the sky than in the other.  We investigate here whether this may arise as a consequence of a large-scale mode of an inflaton or curvaton.} \\label{fig:supermode} \\end{figure}  In the standard inflation scenario, the Universe undergoes a very long inflationary expansion before the comoving observable Universe exits the horizon during inflation.  Thus, any remnants of a pre-inflationary Universe were inflated away before there could be observable consequences.  This accounts for the smoothness of the primordial Universe as well as its flatness.  It also suggests that primordial density perturbations should show no preferred direction.  The existence of a hemispherical power asymmetry in the CMB challenges this basic prediction of inflation.  A superhorizon perturbation would introduce a preferred direction in the Universe and has been considered as a possible origin of the ``axis of evil\" \\cite{Gordon:2005ai}.    In this article, we investigate how the hemispherical power asymmetry could result from a superhorizon perturbation during inflation,  as depicted in Fig.~\\ref{fig:supermode}.  Since the amplitude of the primordial fluctuations depends on the background value of the fluctuating inflationary field, a large-amplitude superhorizon fluctuation would generate a power asymmetry by varying the background value of the field across the observable Universe.  Of course, the superhorizon fluctuation would make the Universe inhomogeneous, and the near-uniformity of the CMB constrains such departures from homogeneity \\cite{Erickcek:2008jp}.  We begin by showing in section \\ref{sec:singlefield} that the power asymmetry cannot be reconciled with single-field slow-roll inflation without violating constraints to the homogeneity of the Universe.   We then consider an alternative inflationary theory, the curvaton model \\cite{Moroi:2001ct}, which has been suggested as a possible source of a power asymmetry \\cite{Gordon:2006ag}.  In section \\ref{sec:curvaton}, we demonstrate that a superhorizon fluctuation in the curvaton field can generate the observed asymmetry without violating the homogeneity constraints.   The required superhorizon fluctuation in the curvaton field may occur, for example, as a remnant of the pre-inflationary epoch or as a signature of superhorizon curvaton-web structures \\cite{Linde:2005yw}.  The proposed model predicts several signatures, which may soon be tested, in the CMB.  We discuss these signatures and summarize our findings in section \\ref{sec:discussion}. ",
        "conclusions": "\\label{sec:discussion}  The hemispherical power asymmetry in the CMB challenges the assumption that the Universe is isotropic and homogeneous.  A superhorizon perturbation in an inflationary field would introduce a preferred direction in the Universe, and we have investigated this mechanism for generating the observed power asymmetry.  We found that the required superhorizon fluctuation in the inflaton field is inconsistent with measurements of the CMB octupole.  A superhorizon fluctuation in a subdominant scalar field, however, is a viable alternative.  A superhorizon curvaton perturbation can generate the observed power asymmetry without introducing unacceptable anisotropy and non-Gaussianity in the CMB.  We have considered the specific asymmetry $A\\simeq0.2$ reported for WMAP, but our results can be scaled for different values of $A$, should the measured value for the asymmetry change in the future.  In particular, the $f_{\\mathrm{NL}}$ constraint (the lower bound to $R$) in Fig.~\\ref{fig:plot} remains the same, but the upper bound (from the quadrupole) increases as $A$ is decreased.  The lower limit to $\\xi$ also decreases as $A$ is decreased. Here we have also considered a general model in which primordial perturbations come from some combination of the inflaton and curvaton.  Although it may seem unnatural to expect the two field decays to produce comparable fluctuation amplitudes, our mechanism works even if $\\xi=1$ (the fluctuations are due entirely to the curvaton).  Thus, the coincidence is not a requirement of the model.  If the power asymmetry can indeed be attributed to a superhorizon curvaton mode, then the workings of inflation are more subtle than the simplest models would suggest.  Fortunately, the theory makes a number of predictions that can be pursued with future experiments.  To begin, the modulated power should produce signatures in the CMB polarization and temperature-polarization correlations \\cite{Pullen:2007tu}. The curvaton model predicts non-Gaussianity, of amplitude $f_{\\mathrm{NL}} \\gtrsim50$ for $A\\simeq0.2$, which will soon be experimentally accessible.  However, the theory also predicts that the small-scale non-Gaussianity will be modulated across the sky by the variation in $\\bar\\sigma$ (and thus in $\\xi$ and $R$). The presence of curvaton fluctuations also changes other features of the CMB \\cite{Ichikawa:2008iq}.  The ratio of tensor and scalar perturbations ($r$) is reduced by a factor of $(1-\\xi)$ and the scalar spectral index is $n_s = 1-2\\epsilon-(1-\\xi)(4\\epsilon-2\\eta)$.  The tensor spectral index ($n_T$), however, is unaltered by the presence of the curvaton, and so this model alters the inflationary consistency relation between $n_T$ and $r$ and possibly the prospects for testing it \\cite{Smith:2006xf}.  We have here assumed simply that the curvaton decays to the same mixture of baryons, dark matter, and radiation as the inflaton.  However, if the inflaton and curvaton decays products differ, then there may be an isocurvature component \\cite{Lyth:2002my, Lemoine:2006sc}.  Finally, the simplest scenario predicts a scale-invariant power asymmetry; the asymmetry has been found at multipole moments $\\ell \\lesssim 40$, but there are claims that it does not extend to higher $\\ell$ \\cite{Donoghue:2004gu}. If this result holds, it will be interesting to see whether the departure from scale invariance can be obtained by suitably altering the power spectra for the curvaton and inflaton.  For instance, a sudden drop in both $V^\\prime(\\phi)$ and $V(\\phi)$ could enhance the gravitational-potential fluctuations from the inflaton while suppressing the fluctuations from the curvaton \\cite{Gordon:2006ag}; the resulting drop in $\\xi$ would reduce the power asymmetry on smaller scales.  We leave such elaborations for future work."
    },
    "0806/0806.2658_arXiv.txt": {
        "abstract": "Relying on infrared surface brightness fluctuactions to trace AGB properties in a sample of elliptical galaxies in the Virgo and Fornax clusters, we assess the puzzling origin of the ``UV-upturn'' phenomenon, recently traced down to the presence of a hot horizontal branch (HB) stellar component. We find that the UV-upturn actually signals a profound change in the c-m diagram of stellar populations in elliptical galaxies, involving both the hot stellar component and red-giant evolution.  First, we encounter that the strengthening of the UV rising branch is always seen to correspond to a shortening in AGB deployment; this trend can be readily interpreted as an age effect, perhaps mildly modulated by metal abundance.  A comparison between galaxy $\\overline{K}$ magnitudes and population synthesis models confirms that, all the way, brightest stars in ellipticals are genuine AGB members, reaching the thermal-pulsing phase, and with the AGB tip exceeding the RGB tip by some 0.5-1.5~mag. The inferred core mass of these stars is found to be $\\lesssim 0.57$~M$_\\odot$ among giant ellipticals. Coupled with the recognized severe deficiency of planetary nebulae in these galaxies, this result strongly calls for an even more critical blocking effect due to a lengthy transition time needed by the post-AGB stellar core to become a hard UV emitter and eventually ``fire up'' the nebula.  The combined study of galaxy $(1550 - V)_o$ color and integrated H$\\beta$ index points, as an explanation for the UV-upturn phenomenon, to a composite HB with a bimodal temperature distribution, i.e.\\ with both a red clump and an extremely blue component, in a relative proportion of roughly [N(RHB):N(BHB)]~$\\sim$~[80:20]. As far as metallicity of the BHB stellar population is concerned, we find that [Fe/H] values of either $\\simeq -0.7$~dex or $\\gtrsim +0.5$ may provide the optimum ranges to feed the needed low-mass stars (M$_* \\ll 0.58$~M$_\\odot$), that at some stage begin to join the standard red-clump stars. ",
        "introduction": "The so-called ``UV-upturn'' phenomenon \\citep{code79}, i.e., the rising ultraviolet emission shortward of 2000 \\AA, sometimes featuring in the spectral energy distribution (SED) of elliptical galaxies and the bulges of spirals, has been for long a puzzling problem for old galaxy environments dominated by stars of mass comparable to the Sun.  In fact, the implied existence of an important contribution of (long-lived) B stars, hotter than $\\sim 30\\,000$~K and providing up to about 2\\% of the galaxy bolometric luminosity \\citep{rb86}, has been alternately identified with different evolutionary stages. Such stages include binaries \\citep{brown06}, blue stragglers \\citep{bailyn95}, blue horizontal-branch (HB) stars \\citep{dorman95}, asymptotic giant branch (AGB) {\\it manqu\\'e} stars \\citep{greggio}, and post-AGB nuclei of planetary nebulae (PNe) \\citep{rb86} \\citep[see][, for an exhaustive review, and a more recent update by \\citealp{yiyoon04}]{oconnell99}  Resolved color--magnitude (c-m) diagrams of  stellar populations in M32 \\citep{brown98,brown00} have definitely shown that even its relatively poor UV emission almost entirely arises from a fraction of hot HB stars, further complemented by a minority contribution from post-AGB PN nuclei. Still, facing the established interpretative scenario, one is left with at least three important issues that need to be assessed to understand the real nature of the UV-upturn phenomenon.  {\\it (i)} The canonical evolutionary framework experienced in Galactic globular clusters naturally predicts a blue HB morphology only for old, metal-poor stellar populations \\citep{chiosi,rffp}. If this is the case for ellipticals too, then UV stars should represent the $Z \\ll Z_\\odot$ tail of a (supposedly) broad metallicity distribution seen to peak at much higher values, around solar abundance. Clearly, a more composite picture might be envisaged once one admits non-standard models (i.e., including the effects of stellar rotation, helium mixing, differential mass loss, etc.) to account, in particular, for the well known ``second parameter'' dilemma \\citep{sweigart97,buonanno97,catelan01,recioblanco06}. However, this unconventional approach still suffers from a somewhat arbitrary fine-tuning of the key physical assumptions.  {\\it (ii)} Hot HB stars might, nonetheless, also be naturally predicted among super metal-rich stellar models, as far as metal abundance (and the linked helium content) exceeds some critical threshold. Presumably, in this case mass-loss allows stars to reach the HB phase with a conveniently low external envelope, compared to the helium core mass \\citep{dorman,castellani92,yi,buzzoni95,dcruz96}. Such ``extreme HB'' stars (EHB) have actually been observed, for example, in $\\omega$~Cen \\citep{dcruz00}, NGC~6388 and NGC~6441 \\citep{rich97}, and in some old Galactic open clusters as well, like NGC~6791 \\citep{kaluzny92,buson06}; they clearly remain the favorite candidates to explain the evolutionary framework of UV-enhanced elliptical galaxies \\citep{moehler05}.  This hypothesis implies, however, a direct relationship between chemical abundance (modulating the helium core mass at the HB onset) and mass-loss efficiency (to suitably ``peel off'' the stellar envelope along the RGB). As a consequence, one has to expect the UV-to-optical color to be, eventually, one of the most quickly evolving features in the SED of elliptical galaxies \\citep{park97}. In theory, the UV-upturn can fade by several magnitudes as the lookback time increases by a few Gyr, although the effect is still detectable at intermediate redshift ($z \\sim 0.3$) \\citep{brown03,ree06}. Unfortunately, the evolutionary details are extremely model-dependent, and a strong UV excess could be triggered at ages as early as $\\sim 6$ Gyr \\citep{tantalo96} or as late as $\\gtrsim$15 Gyr \\citep{yi}.  {\\it (iii)} An established correlation seems to be in place  between PN luminosity-specific rate and $(B-V)$ color for elliptical galaxies in the Virgo and Fornax clusters, and in the Leo group \\citep{peimbert,hui}. The sense is that reddest metal-rich systems display, at the same time, a stronger UV-upturn \\citep{burstein88} {\\it and} a poorer PN population per unit galaxy luminosity \\citep{buzzoni06}. If the PN event is the final fate for AGB stars at the end of their thermal pulsing phase \\citep{ir83}, then the relative deficiency of nebulae might be evidence of an incomplete (or fully inhibited) AGB evolution of low-mass stars under special environment conditions of the parent galaxy.  As a central issue in this discussion, {\\it it is clear therefore that a preeminent connection should exist between UV excess and AGB distinctive properties of stellar populations in early-type galaxies.}  On account of the Fuel Consumption Theorem \\citep{rb86}, a 1~M$_\\odot$ star of solar metallicity enters its core He-burning phase with, at most, the equivalent of 0.43~M$_\\odot$ of H to be spent as nuclear fuel.\\footnote{Under the most extreme hypothesis of no mass-loss, a 1~M$_\\odot$ star with solar abundance $(Y,Z) = (0.28,0.02)$ starts its HB evolution with a total He amount of roughly $0.62$~M$_\\odot$, of which $\\sim 0.47$~M$_\\odot$ are confined in the core \\citep{sg76} and $Y(1-0.47)\\simeq 0.15$~M$_\\odot$ reside in the envelope. Metals amount to roughly $Z(1-0.47) \\simeq 0.01$M$_\\odot$ and, accordingly, fresh H is 0.37~M$_\\odot$. Taking into account the nuclear rates \\citep[e.g.,][]{cox}, the H+He fuel provides at most the equivalent of $0.37+0.62/10 \\simeq 0.43$~M$_\\odot$ of hydrogen.} This means that, under quite general conditions, post-RGB evolution alone could easily account, in principle, for up to 3/4 of the total bolometric luminosity of a galaxy stellar population \\citep{buzzoni98}. Whether this energy is eventually reduced (if stars loose their fuel before they burn it), or whether it is finally released in the form of ultraviolet or infrared photons, crucially depends on mass-loss and its impact along the entire red-giant evolution. Hence, it is of special pertinence to constrain the relevant physical conditions that affect AGB evolution in favor of an earlier transition of HB stars towards high temperature and enhanced ultraviolet emission.  In this paper we would like to draw the reader's attention to a possibly new and powerful approach to the problem, that can find straightforward applications even to distant galaxies. As explained in Sec.~2, the method relies on surface-brightness fluctuation theory to safely tie infrared effective magnitudes (that {\\it can} be determined for unresolved stellar populations), to stellar luminosity at the AGB tip (that {\\it cannot} be directly observed in distant galaxies). We will show, in Sec.~3, that these results tightly correlate with the ultraviolet properties of ``UV upturn'' elliptical galaxies, allowing a self-consistent physical picture and a quite accurate diagnostic of the post-RGB evolution of their underlying stellar populations, including HB morphology and AGB deployment. Our results will be finally summarized and discussed in Sec.~4. ",
        "conclusions": "In this paper we have carried out a synoptic analysis of the different observational features that have to do with the UV luminosity excess in elliptical galaxies. As far as the canonical picture is assumed, with old stellar populations dominating early-type galaxy luminosity, the appearence of the ``UV upturn''  should readily call for a profound change in the c-m diagram of galaxy stellar populations, not only involving the hot stellar component of the galaxy but also reverberating on red-giant evolution at the low-temperature regime. As we mainly deal with distant, unresolved stellar populations, our analysis has to rely on a combined approach, matching infrared and ultraviolet diagnostic tools in order to probe the main features of the stellar c-m diagram, starting with integrated galaxy photometry.  Theory of surface-brightness fluctuations provides, in this sense, a natural and quite powerful way to go deep inside the problem and, as far as the infrared wavelength interval is considered, we have demonstrated theoretically that a straight and very clean relationship is in place between a macroscopic measure, such as the galaxy fluctuation magnitude, and the corresponding individual magnitude of the brightest stars in turn at the tip of the red-giant (AGB+RGB) phases (Fig.~2).  Played in the $K$ band, this correlation leads, from a measurement of $\\overline{K}$, to a value for $K_{\\rm tip}$: \\begin{equation} K_{\\rm tip}  = 0.75\\ \\overline{K} - 3.1, \\label{eq:calib} \\end{equation} with a $\\pm 0.2$~mag internal uncertainty. As we showed in Sec.~2, our SSP theoretical predictions find full support from the observations, and a direct check on the MC star clusters confirms the $\\overline{K}$ vs.\\ $K_{\\rm tip}$ relationship to be a much more general and deeply intrinsic property of stellar populations, virtually independent from any assumption about age, metallicity, IMF, and mass loss parameters. Given its nature, this relationship cannot, by itself, help disentangle the problem of age/metallicity degeneracy; however, quite fruitfully, it provides us with a very direct probe of AGB properties, in a number of relevant details that directly deal with the mass-loss impact and the mass of dying stars (Fig.~4).  Our effort toward exploring the infrared side of galaxy SEDs has a twofold aim since, as a consequence of the basic principle of energy conservation, any gram of stellar fuel spent to feed ultraviolet luminosity cannot (and will not) be spent at longer wavelengths. This has led to the key issue of this paper, summarized in Fig.~6, that the {\\it strengthening of the UV rising branch is always seen to correspond to a weakening in the AGB luminosity extension}, as traced by galaxy $K$ fluctuation magnitude.  This ``shortening'' in AGB deployment is mainly recognized among giant ellipticals ($\\overline{K}$ becomes fainter with increasing galaxy velocity dispersion, $\\sigma_v$, see Fig.~7), and could mainly be ascribed to an age effect, as the AGB tip naturally fades in luminosity with increasing age of the system (Fig.~3), and high-mass galaxies are recognized to be older than systems of lower mass \\citep[e.g.][]{burstein88,bressan96,liu02,jens03,gonz05a,renziniaraa}. However, the relationship in place likely calls for a more elaborated physical scenario, once the full range of observing evidences is added to our analysis.  \\begin{figure} \\centerline{ \\includegraphics[width=\\hsize]{f12.ps} } \\caption{ $K_{\\rm tip}$ vs.\\ Lick Mg$_2$ index, for the galaxy sample in Table~\\ref{tab_tot}. The Lick index is assumed to trace galaxy metallicity according to the \\citet{buzzoni92} calibration, as reported on the top axis of the plot. The observed decrease in $K_{\\rm tip}$ with increasing $[Fe/H]$ can be mostly explained if more metal-rich galaxies are also older, as in a standard monolythic scenario for galaxy formation. The dashed line on the plot marks the minimum luminosity required for stars to experience the thermal pulsing phase along their AGB evolution, and thus end their evolution as PNe. } \\label{eta} \\end{figure}  {\\it (a)} Besides being old, ``UV upturn'' galaxies are also metal rich (i.e., stronger in Mg$_2$ Lick index). Disregarding any change in mass-loss rate, stellar tracks predict slightly more massive stars to evolve off the MS at a fixed age, with increasing metallicity \\citep[e.g.,][]{bressan94,cantiello03}. This leads to correspondingly more massive AGB stars and a brighter AGB luminosity. Facing the observed trend in galaxy distribution, as summarized in Fig.~12, metallicity effects evidently enter by mitigating the dimming action of age on $K_{\\rm tip}$ with increasing galaxy mass. In any case, the interplay between age and metal abundance actually makes the derived range for the Reimers mass-loss parameter (i.e.\\ $\\Delta \\eta \\simeq 0.4$, as discussed in Sec.~3) a safe upper limit. In fact, it suggests that metal abundance does {\\it not} modulate by orders of magnitude mass-loss efficiency via stellar winds.\\footnote{Observational evidence about the link between metallicity and mass-loss in the Milky Way and the Magellanic Clouds is contradictory. For example, \\citet{groe95} find indications that mass-loss rate (not necessarily mass-loss {\\em efficiency}) in single AGB stars is linearly proportional to $Z$. Conversely, also from data of single stars, \\citet{vanl00} argues that $\\dot M$ is metallicity-independent. From a theoretical point of view, \\citet{cantiello03} models suggest that, if mass-loss is really proportional to metallicity, its effect to dim near-IR effective luminosities on average almost exactly offsets the brightening effect of metallicity itself.}  \\begin{figure} \\centerline{ \\includegraphics[width=\\hsize]{f13.ps} } \\figcaption{Observed ultraviolet color $(1550-V)$ vs. core mass of stars at the AGB luminosity tip, as inferred from eq.~(\\ref{mctip}), for the elliptical galaxy sample of Table~2. The M$_c$ distribution is summarized in the lower histogram, and is the maximum actual mass allowed to luminous stars in the galaxies. One sees that mass of dying stars tends to decrease with increasing UV-upturn strength, being in general $\\lesssim 0.57$~M$_\\odot$ among giant ellipticals. The ``outlier'' objects of Fig.~6 are identified again here, with galaxies labelled according to their NGC number. The arrow for NGC~1316 accounts for the claimed strong internal reddening for this galaxy. Displayed uncertainties for the derived values of M$_c$ take in the full error budget-- each component of $\\sigma(M_c)$ being added in quadrature--, including the contribution of $\\overline{K}$ observations, $K$-band bolometric correction ($\\sigma = \\pm 0.2$~mag), and the $K_{\\rm tip}$ vs. $\\overline{K}$ calibration ($\\sigma = \\pm 0.2$~mag). See text for a discussion. } \\label{mcore} \\end{figure}  {\\it (b)} A match of galaxy $\\overline{K}$ data with the calibration of Fig.~5 confirms that, all the way, brightest stars in ellipticals are genuine AGB members, reaching the thermal-pulsing phase (see also Fig.~12), and with the AGB tip exceeding the RGB tip by some 0.5-1.5~mag. In the temperature range of M giant stars, a major fraction of bolometric luminosity is emitted through the $K$ band, and bolometric correction is a nearly constant quantity that we can estimate from $({\\rm Bol} - K) = +2.75 \\pm 0.2$~mag \\citep{johnson}; we can therefore straightforwardly translate the galaxy fluctuation magnitude into an estimate of the bolometric tip luminosity, $L_*^{\\rm tip}$, and therefrom of the corresponding stellar core mass.\\footnote{A somewhat linear relationship between stellar luminosity and core mass is a general consequence of any evolutionary stage characterized by a (multi) shell-burning regime in the presence of a relatively thin external envelope \\citep{paczynski}. This is actually the case of both pre-He flash evolution along the RGB and the thermal pulsing phase along the AGB \\citep[see][, for a more general discussion]{ir83,boothroyd}.} From our previous calibration (eq.~\\ref{eq:calib}), we can write then \\begin{equation} \\log L_*^{\\rm tip} = -0.4[(0.75\\overline{K}-3.1) +2.75 -4.72] \\end{equation} (where the Sun has magnitude M$_{{\\rm Bol},\\odot} = +4.72$). Following \\citet{boothroyd}, from the assumed core mass-luminosity relation for the solar metallicity range, this leads to \\begin{equation} M_c = {{L_*^{\\rm tip}}\\over 52\\,000}+0.456 \\qquad\\quad [M_\\odot]. \\label{mctip} \\end{equation} Figure~13 reports the inferred core-mass distribution, at the PN onset, for our galaxy sample.  Note that this is the {\\it maximum} actual mass allowed to luminous stars in each galaxy environment, and demonstrates that the {\\it mass of dying stars tends to decrease with increasing UV-upturn strength}, being in general $M_{\\rm dying} \\lesssim 0.57$~M$_\\odot$ among giant ellipticals. For this mass range, PN lifetime is the largest possible, but the timescale for the nebula to be visible is critically constrained by the transition time ($\\tau_{\\rm tt}$) needed by the post-AGB stellar core to be hot enough to ``fire up'' the ejected envelope and become a hard UV emitter \\citep[][]{letizia00,marigo05}. The evident drop of $\\alpha$ among strong UV-upturn galaxies (Fig.~8) might be a direct consequence, therefore, of an increasing blocking effect of $\\tau_{\\rm tt}$ along the inferred $M_c$ range \\citep[i.e., the stellar core takes longer to heat-up than the shell to evaporate;][]{buzzoni06}, to which one has to further add a size cut in the overall PN population, as a result of the EHB progenitors (M$_* \\lesssim 0.52$~M$_\\odot$) evolving as {\\it AGB-manqu\\'e} stars and therefore skipping the nebula event.  {\\it (c)} We remarked, in Sec.~3, the importance of the integrated $H\\beta$ index as a fairly selective tracer of the warm (T$_{\\rm eff} \\simeq 8\\,000-10\\,000$~K) stellar component in the galaxy stellar population. A proper assessment of the photometric contribution from this range of temperatures is of paramount importance in the framework of early-type galaxy evolution, in order to single out any signature of recent (i.e., in the last few Gyr or so) star formation or, conversely, of intervening evolution of the HB morphology among old SSPs, as in a more standard canonical scenario.  As far as UV-upturn galaxies are concerned, the study of $H\\beta$ distribution clearly points to {\\it a substantial lack of A-type stars} in the galaxy mix (see Fig.~9). While, on one hand, this definitely secures the ``quiescent'' nature of these galaxies, it also poses, on the other hand, a stringent constraint on HB morphology in their old-age context. In fact, a bimodal temperature distribution is required for the HB to assure, at a time, both an enhanced UV emission {\\it and} a conveniently low $H\\beta$ feature. Thus, in these systems the expected prevailing bulk of red  HB stars should be accompanied, at some point, by a residual population of blue (metal-poor?) HB objects (coincident with EHB stars in the current empirical classification scheme), peaked at about 20\\,000-40\\,000~K, in a proportion of, roughly, $N_{\\rm RHB}:N_{\\rm BHB}~\\simeq [80:20]$.  On the other hand, to complete the picture, one cannot neglect the masking effects of age distribution, facing a recognized evidence for low-mass ellipticals to display a more silent but also more continuous star formation along their entire galaxy life, that naturally feeds the A-star contribution, thanks to the bluer MS turn-off point (MSTO) exhibited by SSPs in the $~1-3$~Gyr age range, and leads to a younger ``average\" age, compared to high-mass systems. This is what we observe, for instance, among the resolved stellar populations of the Local Group dwarf spheroidals \\citep{mateo98} \\citep[see, in this regard, the illustrative location of M32 in Fig.~9, and also consider the discussion by][]{schiavon04}.  {\\it (d)} Along with our discussion of the $\\overline{K}$ vs.\\ $(1550-V)$ relation, we noticed, in Fig.~6, the presence of a few outliers about 0.7~mag brighter at infrared magnitudes, or alternatively $\\sim 1.5$~mag ``bluer'' in the $(1550-V)$ color, than the main galaxy population. In order to further investigate this issue, we tracked the relevant objects also in other figures, whenever possible.  If galaxy mass ({\\it alias} $\\log \\sigma_v$) is considered as the leading physical parameter to compare outlier location with respect to the bulk of the galaxy distribution (see, for instance Fig.~7), one must conclude that both NGC~4552 and NGC~1389 seem to have a brighter AGB tip rather than a bluer $(1550-V)$ color. With regard to NGC~1389, \\citet{liu02} find that the $\\overline{K_s}$ SBFs of NGC~1389 are also too bright compared to its ($V - I$) color, a fact that would be consistent with either a higher than average metallicity given the age of its most recent burst of star formation or a longer lifetime of its TP-AGB stars \\citep{mouhcine05}. In the case of NGC~4552, however, \\citet{jens96} contribute an interesting piece of information. These authors measure near-IR SBFs for several galaxies in Virgo, and their results are systematically fainter than those obtained by \\citet{pahr94}. Unfortunately, the two groups use slightly different filters (Jensen et al.\\ employ $K^\\prime$, vs.\\ $K_s$~of Pahre \\& Mould), but the discrepancy is larger than can be ascribed to the effect of the filters.\\footnote{\\citet{jens96} attribute the difference to the higher S/N ratio of their data; the S/N ratio of the \\citet{liu02} images is similar to that of Jensen et al.'s.} In particular, Jensen et al.\\ find $\\overline{K^\\prime} = -5.51$ for NGC~4552; assuming $\\overline{K^\\prime} \\equiv \\overline{K_s}$, NGC~4552 would no longer be deviant in the $\\overline{K_s}$ vs.\\ ($1500 - V$) plane. Concerning NGC~1387, the dearth of data for this galaxy in the literature (see Table~2) makes it hard to propose an origin for its departure from the $\\overline{K_s}$ vs.\\ ($1500 - V$) sequence.\\footnote{Note that NGC~1387 complies with the $\\overline{K}_s$ vs.\\ ($V - I$) correlation determined by \\citet{liu02}.} On the other hand, NGC~1387 and NGC~1389 are lenticular galaxies, like the merger remnant NGC~1316, but so are NGC~3384 and NGC~4406, both of which do not deviate from the correlation.\\footnote{The other galaxies in our sample are all genuine ellipticals \\citep{rc3}, except for NGC~224 ({\\it alias} M~31), which is so close, however, that bulge SBFs can be measured without any significant contamination from the disk.} At any rate, NGC~1387 and NGC~1389 constitute privileged candidates for any future ``in-depth'' investigation."
    },
    "0806/0806.2002_arXiv.txt": {
        "abstract": "The zonal winds on the surfaces of giant planets vary with latitude. Jupiter and Saturn, for example, have several bands of alternating eastward (prograde) and westward (retrograde) jets relative to the angular velocity of their global magnetic fields. These surface wind profiles are likely manifestations of the variations in depth and latitude of angular velocity deep within the liquid interiors of these planets.  Two decades ago it was proposed that this differential rotation could be maintained by vortex stretching of convective fluid columns that span the interiors of these planets from the northern hemisphere surface to the southern hemisphere surface. This now classic mechanism explains the differential rotation seen in laboratory experiments and in computer simulations of, at best, weakly turbulent convection in rotating constant-density fluid spheres.  However, these experiments and simulations are poor approximations for the density-stratified strongly-turbulent interiors of giant planets.  The long thin global convective columns predicted by the classic geostrophic theory for these planets would likely not develop.  Here we propose a much more robust mechanism for maintaining differential rotation in radius based on the local generation of vorticity as rising plumes expand and sinking plumes contract.  Our high-resolution two-dimensional computer simulations demonstrate how this mechanism could maintain either prograde or retrograde surface winds in the equatorial region of a giant planet depending on how the density scale height varies with depth. ",
        "introduction": "The complicated flow patterns observed at the cloud tops of giant planets are likely due to a combination of atmospheric phenomena in the shallow surface layer and thermal convection within the deep interior.  The question, which has been debated for several decades, is what differential rotation exists deep below the surface of a giant planet and what dynamics maintains it. Differential rotation is the axisymmetric zonal wind pattern in latitude and radius relative to the deep-seated global magnetic field, which is assumed to be rotating at approximately the mean rotation rate of the planet.  Differential rotation can be maintained by either axisymmetric Coriolis forces arising from a thermally-driven meridional circulation, i.e., as a thermal wind, or by the convergence of Reynolds stress, i.e., the nonlinear transport of longitudinal momentum in latitude and radius.  Magnetic forces, like viscous forces, typically inhibit differential rotation; however in some cases they too can drive differential rotation (e.g. Dormy {\\it et al.} 2002).  Most theoretical studies of this problem have approached it from an Earth-atmosphere context, ignoring the dynamics of the vast interior and using models and approximations only appropriate for shallow atmospheres. See Dowling (1995) for a review of this approach. These models typically assume a hydrostatic balance in the radial direction instead of solving the full momentum equation.  That is, they ignore the radial component of the flow in transporting longitudinal momentum or producing Coriolis forces, two critical elements for maintaining differential rotation in deep convective zones. Instead, these shallow-atmosphere models of giant planets usually rely on a thermal wind scenario (e.g. Allison 2000) or the convergence of latitude-longitude Reynolds stress to drive zonal winds. The latter produces a retrograde (i.e., westward directed) zonal wind in the equatorial region (e.g. Williams 1978, Cho and Polvani 1996) because the local vertical component of the planetary rotation rate (the only component these models consider) increases with latitude.   Retrograde zonal flow in the equatorial region is observed on surfaces of Uranus and Neptune (Hammel {\\it et al.} 2005) but a prograde (i.e., eastward directed) equatorial jet is observed on our gas giants, Jupiter (Porco {\\it et al.} 2003) and Saturn (Sanchez-Lavega {\\it et al.} 2000).  The zonal wind at higher latitudes on the gas giants has a latitudinally banded pattern of alternating retrograde and prograde flows. A banded pattern of zonal flows with a prograde equatorial jet can be obtained by tuning a heating function distributed in latitude and radius that continually nudges the temperature toward a prescribed profile that drives the desired zonal wind pattern (e.g. Williams 2003). However, in such a model the zonal wind profile in latitude and radius is the result of the shallow layer assumption and the mathematical fit to the surface observations via the ad hoc heating function; therefore it does not provide a dynamically-consistent prediction or explanation.  It may indeed be the case that the zonal winds on giant planets are confined to the shallow surface layers and driven by heat sources and instabilities there without any influence from the deep convection below.  However, hydrostatic shallow-atmosphere models assume this from the start and so are not capable of predicting this to be the case. The Galileo probe measured a doubling of the zonal wind speed with depth in the surface layer of Jupiter (Atkinson {\\it et al.} 1998). A scale analysis (Ingersoll and Pollard 1982) suggests that the zonal winds on Jupiter and Saturn extend well below their surfaces. Probably the strongest indication of the existence of deep convection is that spherically symmetric evolutionary models with detailed equations of state and opacities predict convection throughout the deep fluid interiors of our gas giants (e.g. Hubbard {\\it et al.} 1999, Guillot 2005, and T. Guillot private communication). Therefore, to predict what differential rotation exists below the surface and to understand how it is maintained one needs a global dynamically-consistent model that extends deep below the surface, with sufficient physics and the full set of three-dimensional (3D) equations of motion. The model may also need good representations of radiative transfer and moist convection near the surface (Ingersoll {\\it et al.} 2000), a hydrogen phase transition well below the surface separating the semi-conducting and metallic regions (Nellis 2000) and magnetic field generation with its feedback on the flow.  An alternative to a shallow-atmosphere model is a dynamically-consistent 3D global model from the geodynamo community.  Most giant planet studies using this approach apply the Boussinesq approximation to the set of equations that describes deep rotating convection.  That is, a constant background density is assumed. These 3D global models do produce banded zonal winds and a prograde equatorial jet without neglecting the dynamics of the deep interior and without prescribing ad hoc heating distribution functions (e.g. Sun {\\it et al.} 1993, Christensen 2002, Stanley and Bloxham 2004, Heimpel {\\it et al.} 2005). The studies find that the convergence of Reynolds stress, both latitude-longitude and radius-longitude, dominates over the thermal wind mechanism in maintaining differential rotation and that the kinetic energy in the meridional circulation is typically several orders of magnitude less than that in the differential rotation (e.g. Christensen 2002). That is, these studies suggest that meridional circulation is maintained by the Coriolis forces arising from the differential rotation, not the other way around (i.e., the differential rotation is not a thermal wind).  However, because these Boussinesq models assume a constant background density and relatively large viscosity, the differential rotation in these simulations is maintained by the vortex stretching of large convective columns due to the sloping impermeable boundary.  Unlike these Boussinesq models, the interior of a giant planet has a significant density stratification, especially near the surface, and much smaller viscosity, which produces strong convective turbulence characterized by a broad spectrum of scales (e.g. Glatzmaier 2005). This turbulence is dominated by small-scale fluid parcels, which initially develop as thermal boundary layer instabilities, then detach and either rise from the inner boundary or sink from the outer boundary. Their dynamical evolution in the bulk of the convection zone is determined by nonlinear vortex-vortex interactions and the local density stratification, not by the spherical curvature of distant boundaries. The deep-convection constant-density models from the geodynamo community produce flow amplitudes relatively independent of depth, opposite of that assumed in the hydrostatic shallow-atmosphere models from the climate community. Both, however, completely ignore the maintenance of differential rotation by radial flow through a density stratification.  Glatzmaier and Gilman (1981, 1982), using quasi-linear solutions to the anelastic equations of motion, examined the role of density stratification in generating vorticity and maintaining differential rotation by compressional torque without requiring the vortex stretching of columnar convection. Ingersoll and Pollard (1982) and Busse (1986) addressed the role of density stratification using scale and asymptotic analyses.  However, they retained Busse's original assumption of vortex-stretching by geostrophic columnar convection spanning the entire interior, replacing the column length by the integrated column mass.  Here we examine the deficiency of the classic vortex-stretching mechanism for maintaining differential rotation in giant planets. Then we describe a more robust mechanism that accounts for their large density stratifications and strong turbulence without requiring geostrophic columns that span the interior.  We study this mechanism with nonlinear simulations of turbulent convection that focus on the convergence of radial-longitudinal Reynolds stress for maintaining differential rotation in radius. ",
        "conclusions": "We have argued that the extremely thin convective columns required to span from the northern to southern boundaries in the classical vortex-stretching mechanism would likely not develop in the turbulent interiors of giant planets.  As an alternative, we have investigated the maintenance of differential rotation due to the effects of rising fluid expanding and sinking fluid contracting within a density-stratified equatorial plane of a giant planet. Our simple 2D computer simulations illustrate how this density-stratification mechanism can maintain differential rotation by compressional torques acting locally on small convective plumes. We have discussed differences in the flow structures produced with constant-density (Boussinesq) models and those produced by density-stratified (anelastic) models, which are most significant in the outer region of a giant planet where the density scale heights are smallest.  The opposite differential rotation profiles illustrated in figures 1 and 2 suggest that the oppositely directed equatorial winds observed on our ice giants compared to those on our gas giants might be due to their different radial profiles of density (Hubbard {\\it et al.} 1991, Guillot 2005).  Our models also demonstrate how this mechanism can maintain differential rotation in radius when the interior is fully convective (i.e., no solid core) and within a convection zone above a convectively stable interior.  The 2D models we have presented here, however, are meant to simply demonstrate this fundamental mechanism, which has been neglected in most previous models of giant planets; they are not meant to predict the pattern or extent of differential rotation below the surface of a giant planet.  Our Cases 1 and 2 span only two density scale heights, not the much larger number which exist within a giant planet, most of which occur in the shallow layers where the density-stratification mechanism is most effective. We have also neglected the hydrogen phase transition, which is predicted to exist at roughly 90\\% of Jupiter's radius and 50\\% of Saturn's radius (Nellis 2000), and the magnetic field that is generated in the outer semi-conducting region where the electrical conductivity rapidly increases with depth.  In addition, we have not addressed the maintenance of differential rotation in latitude, which would require a 3D global simulation with a density stratification in spherical radius.    \\bigskip \\noindent{\\bf Acknowledgments} \\smallskip  We thank F. Busse, U. Christensen, C. Jones, P. Olson, P. Roberts and D. Stevenson for discussions.  T.R. is supported by an NSF Astronomy and Astrophysics Postdoctoral Fellowship under award 0602023. Support for this research was provided by grants from the NASA {\\it Planetary Atmospheres Program} (NAG5-11220), the NASA {\\it Outer Planets Research Program} (NNG05GG69G), the NASA {\\it Solar and Heliospheric Physics Program} (NNG06GD44G) and from the {\\it Institute of Geophysics and Planetary Physics} at Los Alamos National Laboratory and the University of California Santa Cruz. Computing resources were provided by NSF at the {\\it Pittsburgh Supercomputing Center} and by an MRI funded Beowulf cluster at UCSC (AST-0521566), by the NASA {\\it Advanced Supercomputing Division} and by DOE at the {\\it National Energy Research Scientific Computing Center}.     \\bigskip \\noindent{\\bf References} \\smallskip  \\refpar Allison, M., A similarity model for the windy jovian thermocline. {\\it Planet. 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Sci.}, 1978, {\\bf 35}, 1399-1424.  \\refpar Williams, G.P., Jovian dynamics. Part III: Multiple, migrating, and equatorial jets. {\\it J. Atmos. Sci.}, 2003, {\\bf 60}, 1270-1296.   \\appendices"
    },
    "0806/0806.2833_arXiv.txt": {
        "abstract": "We consider the statistical relationship between the growth rate of activity in the early phase of a solar cycle with its subsequent amplitude on the basis of four datasets of global activity indices (Wolf sunspot number, group sunspot number, sunspot area, and 10.7-cm radio flux). In all cases, a significant correlation is found: stronger cycles tend to rise faster. Owing to the overlapping of sunspot cycles, this correlation leads to an amplitude-dependent shift of the solar minimum epoch. We show that this effect explains the correlations underlying various so-called precursor methods for the prediction of solar cycle amplitudes and also affects the prediction tool of Dikpati et al. (2006) based upon a dynamo model.  Inferences as to the nature of the solar dynamo mechanism resulting from predictive schemes which (directly or indirectly) use the timing of solar minima should therefore be treated with caution. ",
        "introduction": "Solar activity is the driver of space weather, which has practical consequences for human activities in space. This is one reason why the search for methods to predict its (short-term and long-term) future levels has found much interest in the literature. Further motivation arises from the potential implications for understanding the {\\em origin} of solar activity: a reliable method of predicting the amplitude of future solar cycles could provide a constraint on dynamo models. On the other hand, the converse proposition is not necessarily true, as has been pointed out by \\citet{Bushby:Tobias:2007}: the nonlinear dynamics of the dynamo might be such as to make mid- to long-term prediction impossible even if an almost perfect physical understanding of the dynamo mechanism is achieved.  In many cases, recipes for prediction are inferred from correlations found in historical records of measured quantities, which are (directly or indirectly) related to solar activity. As illustrated by \\citet[][cf. his Fig. 14.2]{Wilson:1994} and \\citet[][cf. their Fig.~6]{Lantos:Richard:1998}, the success of most methods in actually {\\em predicting} the unknown amplitude of a future cycle is rather disappointing. Nevertheless, the correlations between several `precursors', i.e., quantities measured during the descending or minimum phase of a cycle and the amplitude of the subsequent cycle \\citep[e.g.,][]{Hathaway:etal:1999, Schatten:2003} might have a non-random origin and thus call for a physical explanation. This could have implications for dynamo models.  In this paper, we consider the effect of the overlapping of solar cycles in combination with their asymmetric shape on the correlations between precursors and following cycle amplitudes. In this connection, the important aspect of the asymmetry is the difference of the amplitude-dependent ascent rate in the early cycle phase (related to the so-called Waldmeier effect) compared to the decay rate near the end of a cycle. Since sunspot cycles overlap for typically 2 to 3 years \\citep{Harvey:1992a}, this asymmetry affects the timing of the activity minima, which are pivotal epochs for most precursor methods. We show that these effects can explain the correlations upon which such methods are based, without necessarily implying a direct physical connection between the precursor quantity and the following cycle. We also show that the essence of the Waldmeier effect, i.e., that stronger cycles tend to show a faster rise of activity levels during their ascending phase than weaker cycles, is a robust property present in all activity indices.  On this basis, we explain how cycle asymmetry and cycle overlapping may also affect the dynamo-based prediction method of \\citet{Dikpati:Gilman:2006}, in spite of recent claims to the contrary \\citep{Dikpati:etal:2008}. ",
        "conclusions": "We have confirmed a highly significant correlation between the growth rate of activity during the early phase of a solar cycle and its maximum amplitude for all global activity indices, i.e., Wolf and group sunspot numbers, total sunspot area, and 10.7-cm radio flux. On the other hand, there is no significant correlation between the decay rate in the late cycle phase and the cycle amplitude.  Owing to the overlapping of individual cycles, this asymmetry leads to an amplitude-dependent shift of the minimum epoch, thus explaining (fully or partly) the predictive power of precursor methods which (directly or indirectly) use the timing of the activity minimum as a pivotal point. The resulting correlation in the sunspot area data probably also affects the predictions with the dynamo-based model of \\citet{Dikpati:Gilman:2006}.  For our understanding of the origin of the solar magnetic field, it is important to disentangle the effects of `real' physical precursors, i.e., properties of the old cycle directly affecting the flux generation for the next cycle or early high-latitude manifestations of the new cycle, from apparent precursors, which derive their predictive power from the the amplitude-dependent shift of the minimum epoch."
    },
    "0806/0806.2438_arXiv.txt": {
        "abstract": "We report the first results from the GammeV search for chameleon particles, which may be created via photon-photon interactions within a strong magnetic field. Chameleons are hypothesized scalar fields that could explain the dark energy problem. We implement a novel technique to create and trap the reflective particles within a jar and to detect them later via their afterglow as they slowly convert back into photons. These measurements provide the first experimental constraints on the couplings of chameleons to photons. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0806/0806.2054_arXiv.txt": {
        "abstract": "We demonstrate that, for the case of quasi-equipartition between the velocity and the magnetic field, the Lagrangian-averaged magnetohydrodynamics $\\alpha-$model (LAMHD) reproduces well both the large-scale and small-scale properties of turbulent flows; in particular, it displays no increased (super-filter) bottleneck effect with its ensuing enhanced energy spectrum at the onset of the sub-filter-scales.  This is in contrast to the case of the neutral fluid in which the Lagrangian-averaged Navier-Stokes $\\alpha-$model is somewhat limited in its applications because of the formation of spatial regions with no internal degrees of freedom and subsequent contamination of super-filter-scale spectral properties.  \\resp{}{We argue that, as the Lorentz force breaks the conservation of circulation and enables spectrally non-local energy transfer (associated to Alfv\\'en waves), it is responsible for the absence of a viscous bottleneck in MHD, as compared to the fluid case.  As LAMHD preserves Alfv\\'en waves and the circulation properties of MHD, there is also no (super-filter) bottleneck} found in LAMHD, making this method capable of large reductions in required numerical degrees of freedom; specifically, we find a reduction factor of $\\approx 200$ when compared to a direct numerical simulation on a large grid of $1536^3$ points at the same Reynolds number. ",
        "introduction": "When large-scale numerical simulations of astrophysical or geophysical magnetohydrodynamics (MHD) are desired, all dynamical scales of the physical system are rarely, if ever, resolved.  For this reason, sub-grid-scale (SGS) modeling of MHD dynamics in the context of computations in the geophysical and astrophysical context is required. This modeling can be achieved implicitly, in the simplest example by employing a dissipative numerical scheme, or it can be done explicitly by creating a Large Eddy Simulation (LES--see \\cite{MK00} for a recent review). Explicit methods for MHD are not as pervasive as they are in engineering, or for geophysical and atmospheric flows.  In fact, modeling for MHD is a relatively new field (see \\cite{PFL76,Y87}).  One problem with extending the LES methodology for hydrodynamic turbulence to MHD is that most LES are based upon eddy-viscosity concepts \\cite{MK00}, which can be related to a known power law of the energy spectrum \\cite{ChLe1981} (although generalizations can be devised, see e.g. \\cite{BaPoPo+2008}), or upon self-similarity.  For MHD, the underlying assumption of locality of interactions in Fourier space is not necessarily valid \\cite{AMP05a,MAP05} (a contradiction of self-similarity) and spectral eddy-viscosity concepts \\cite{ZSG02} cannot be applied in a straightforward manner as neither kinetic nor magnetic energy is a conserved quantity and the general expression of the energy spectrum is not known at this time \\cite{I64,K65,GoSr1995,GaNaNe+2000,MiPo2007a,PoMiMo+2008,MaCaBo2008}. Purely dissipative models \\cite{TFS94,AMK+01} are inadequate as they ignore the exchange of energy at sub-filter scales between the velocity and magnetic fields and such models have been shown to suppress small-scale dynamo action \\cite{HB06} and any inverse cascade from the sub-filter scales \\cite{MC02}.  \\add{}{A satisfactory LES for MHD has been proposed} for the case starting with some degree of alignment between the velocity and magnetic fields \\cite{LS91,MC02}.  Other restricted-case MHD-LES are applicable to low magnetic Reynolds number \\cite{PPP04,KM04,PMM+05}. Extensions of spectral models to MHD based on two-point closure formulations of the dynamical equations proposed recently look promising in the analysis of turbulent flows and of the dynamo mechanism \\cite{BaPoPo+2008}. Finally, though technically not an LES, there are also hyper-resistive models for MHD which require rescaling of the length (wavenumber) scales to a known direct numerical simulation (DNS) \\cite{HB06}.  One model which can be written as an LES is the Lagrangian-averaged MHD (\\lamhd) equations \\cite{H02a, H02b, MP02}.  It has been shown to reproduce a number of features of DNS. \\resp{}{In two dimensions (2D) for Taylor Reynolds numbers ($R_\\lambda$) up to $\\approx5000$ it has been shown to reproduce selective decay, the inverse cascade of mean-square vector potential, and dynamic alignment between the velocity and magnetic fields \\cite{MMP05a} as well as the statistics \\vier{of} small-scale cancellation \\cite{PGMP05} and intermittency \\cite{PGHM+06}.  In three dimensions (3D) at Reynolds numbers ($Re$) of $\\approx500$, \\lamhda reproduced the inverse cascade of magnetic helicity (associated with the development of force-free magnetic field) and the helical dynamo effect \\cite{MMP05b}.} \\neu{It has also been tested (up to kinetic $Re\\approx3000$, magnetic $Re\\approx300$) for its ability to predict the critical magnetic Reynolds number for a non-helical dynamo at low magnetic Prandtl number \\cite{Mi2006a}. \\lamhda performed well in all these tests. Its equivalent hydrodynamic model, the Lagrangian-averaged Navier-Stokes (LANS) equations, also performed well in tests at $R_\\lambda\\lessapprox300$ (see \\cite{CHO+05} and references in \\cite{PGHM+07a}).  However, above $Re\\approx3000$ ($R_\\lambda\\approx800$),} it was shown that placing the filter width in the inertial range leads to contamination of the super-filter-scale properties (such as the spectra) \\vier{for \\lans.} We refer here to this effect as the super-filter-scale bottleneck, which \\resp{}{may be} different in nature from the viscous bottleneck observed in some DNS of the Navier-Stokes equations. The contamination may be linked to the formation of spatial regions in the flow with no internal degrees of freedom (so-called ``rigid bodies'') \\cite{PGHM+07a}, which also correspond to the development of a secondary inertial range of the LANS equations at sub-filter scales. This \\resp{}{super-filter-scale contamination} provides an effective constraint on the filter size and, hence, on the available reduction of the total number of the (numerical) degrees of freedom (\\dof) needed to reproduce the large-scale dynamics of the flow at a given Reynolds number; a factor of $\\approx 10$ \\resp{}{can be achieved.  This limitation is not apparent} in low and moderate \\resp{}{Reynolds number} (resolution) simulations (e.g., $64^3$ \\lansa compared with $256^3$ DNS) as the scale separation is not enough for the above-mentioned phenomenon of contamination of small-scale spectra because of rigid body regions in the flow to appear. \\resp{}{The bottleneck (and super-filter-scale contamination) was not studied as such but neither was it observed in 2D \\lamhda for high Reynolds number \\cite{MMP05a,PGMP05,PGHM+06}. 3D \\lamhda has only been tested} at \\add{}{more} moderate Reynolds number \\cite{MMP05b} (see also \\cite{MiPoSu2008} for a recent review). The aim of the present work is, thus, to \\resp{}{determine if} \\lamhda in three space dimensions, \\add{}{for higher Reynolds number} \\resp{}{develops problems similar to that of \\lans.  Specifically, we test for the existence of} spatial regions with no available internal degrees of freedom. We show in the following that \\lamhda behaves better in this respect than \\lans, and, thus, continues to appear as a promising model for MHD flows. ",
        "conclusions": "In this paper, we have tested the \\lamhda model against high Reynolds number direct numerical simulations (up to \\resp{}{Reynolds numbers of $\\approx 9200$)} and in particular we have focused our attention on the dynamics of small scales near the $\\alpha$ cut-off. We find that the small-scale spectrum presents no particular defect; specifically, we find that, unlike in the hydrodynamical case, the Lagrangian-averaged modeling for MHD exhibits, even at large Reynolds numbers, neither a positive-power-law spectrum nor any contamination of the super-filter-scale spectral properties.  \\resp{}{This difference between \\lansa and \\lamhda} is not due to the inclusion of a hyper-diffusive term in \\lamhda that stems from the derivation of the model; rather, it stems from fundamental differences between hydrodynamics and MHD. Indeed, neither the (non-consistent) removal of hyperdiffusion from \\lamhda nor the examination of scales much smaller than $\\alpha$ gave any indication of problems similar to those caused by the zero-flux regions found in computations using \\lans.  These regions limited the computational gains of using \\lansa as a LES in hydrodynamics to a factor of only $10$ in computational degrees of freedom or $30$ in computation time.  \\lamhda is not subject to the same limitations and, as we demonstrated, a gain of a factor of $200$ in the number of degrees of freedom, or a factor of $1300$ in computation time, obtains when comparing to the highest Reynolds number in turbulent MHD available today in a DNS.  There are two obvious candidates to explain the lack of a (super-filter-scale) bottleneck effect in \\lamhd: the enhanced (hyper-)diffusion in \\lamhda compared with \\lans, and \\resp{}{physical differences between fluids and magneto-fluids, specifically, spectrally nonlocal transfer via Alfv\\'en waves and its associated} breaking of the circulation conservation.  The first candidate would eliminate the super-filter-scale bottleneck by removing energy from the system and precluding the formation of a secondary range below the filtering scale $\\alpha$ (note that this term becomes of the same order as the ordinary diffusion when $l\\sim\\alpha$).  Simulations of \\lamhda performed without the hyper-diffusion term \\resp{}{ruled out} this scenario, as no super-filter bottleneck was found.   The second candidate is the \\resp{}{presence of the Lorentz force in MHD (and \\lamhd) which} breaks down the circulation conservation \\resp{}{and provides the restoring force for Alfv\\'en waves.  Both properties were shown to be preserved by \\lamhd.  In Navier-Stokes, the development of helical filaments could quench local interactions \\cite{MT92,T01} depleting the energy transfer and leading to the viscous bottleneck.  However, in MHD, the conservation of the circulation ($d\\Gamma/dt=0$ in the absence of dissipation) is broken by the Lorentz force, which modifies Kelvin's theorem (see Eq. (\\ref{eq:kelvin})).  The forcing term is associated with the Alfv\\'en waves, and represents the removal of circulation (and of kinetic energy) that is transfered to the magnetic field. Note that in Fourier space, the term scales as $k E_M(k)$ and is dominant compared to the dissipation in the inertial range. This term precludes the formation of rigid bodies, giving as a result a larger net flux towards smaller scales and a resulting larger dissipation in MHD/\\lamhd. This is illustrated in Fig. \\ref{FIG:PDFS}. This sink of circulation may also be the cause of the lack of a viscous-scale bottleneck in MHD. In LANS it was shown \\cite{PGHM+07a,PiGrHoMi+2008} that conservation of the circulation (except for viscosity) leads to the formation of rigid bodies that fill a substantial volume of the fluid, and that in turn substantially decrease the energy flux to small scales, reduce dissipation, and create the super-filter scale bottleneck.} In \\lamhd, the destruction of sub-filter-scale rigid bodies by large scale magnetic field and shear \\resp{}{results} as the presence of a magnetic field permits the development of long-range interactions in spectral space \\cite{MAP05,AMP05a,Al2007a}.  This can also explain why $\\alpha-$models for other non-local equations, or for problems that do not preserve the circulation provide good SGS models.  As an example, the use of \\lansa in primitive equations ocean modeling gives satisfactory results, e.g. in its reproducing the Antarctic circumpolar current baroclinic instability that can be seen only at substantially higher resolutions when using direct numerical simulations \\cite{HeHoPe+2008}.  \\resp{}{Energy is dissipated in MHD flows through two different processes. Viscosity is responsible for the dissipation of mechanical energy, while Ohmic losses are responsible for dissipation of magnetic energy. Mechanical and magnetic energy are not conserved separately, but rather coupled as illustrated by the existence of Alfv\\'en waves, which correspond to oscillations of the magnetofluid with the velocity field parallel or anti-parallel to the magnetic field, and associated to the interchange of magnetic and kinetic energy.  In MHD, it is believed that most of the total energy in the flow is finally dissipated (mediated by this interchange) through Ohmic losses, in a process that involves reconnection of magnetic field lines. This is supported by several simulations of MHD turbulence \\cite{HaBrDo2003,Mi2007a} and is consistent with phenomenology. While in hydrodynamics small scales are permeated by a myriad of vortex filaments, in MHD the dominant dissipative structures are current sheets, where strong gradients of the magnetic field and their associated strong currents lead to rapid Ohmic dissipation.  Sub-grid models attempt to replace the physical processes of small-scale dissipation by processes that mimic the non-linear transfer of energy to smaller scales (where energy is in reality dissipated, but now in scales that are not resolved by the model). In traditional LES, this is done with enhanced turbulent viscosities. Note that the eddy viscosity is not obtained from the linear dissipative term (the term that describes the actual physical process responsible for the dissipation) but from the non-linear terms in the equations (the terms that describe the coupling between fields at different scales). The final goal is not to capture the dissipation processes, but to be able to preserve (with computational gains) the large scale dynamics.}  \\resp{}{Lagrangian averaged models take a different (although related, see e.g., \\cite{PGHM+06}) approach. Besides adding (in some cases, as in the case of MHD) an enhanced viscosity, the non-linear terms are modified at small scales. This modification changes the time-scale of the energy cascade, and as a result changes the scaling law of the energy spectrum $E(k)$ at sub-filter scales. This change leads to changes in the dissipation, as the dissipation is in the original equations proportional to $k^2E(k)$. The end result (an enhanced dissipation that is intended to mimic the transfer of energy to smaller scales in the unresolved scales) should be the same as in a traditional LES: gains in computing costs preserving as much information of the large scale flow as possible.  As in the case of LES, the actual dissipation process is not as important as the fact that large-scale dynamics should be reproduced with minimal contamination \\vier{by} the sub-grid model. We believe the results presented here (and in earlier work \\cite{MMP05a,MMP05b,PGMP05,PGHM+06,Mi2006a}) show this is the case, and allow the use of the LAMHD equations as a subgrid model of MHD turbulence. However, considering the differences observed between LANS and LAMHD, we discuss the dissipation processes in LAMHD.  Two mechanisms for dissipation can be identified in LAMHD: dissipation of mechanical energy through the viscosity, and dissipation of magnetic energy through (enhanced) Ohmic losses. From the equations, the total variation of energy goes as \\cite{MMP05a}: $dE/dt = -\\nu\\left<\\boldsymbol{\\omega}\\cdot\\bar{\\boldsymbol{\\omega}}\\right> -\\eta\\left<j^2\\right>$ and as a result the mechanical energy dissipation scales as $k^2E_V(k)$ while the magnetic energy dissipation scales as $(1+ \\alpha^2k^2)k^2E_M(k)$. The extra $k^2$ factor in the latter gives more dissipation than in the LANS case. This excess of magnetic dissipation in LAMHD mimics, as previously mentioned, the dominant contribution to dissipation by Ohmic losses in MHD. This hyperdiffusion is required in the sub-filter scales to accurately model the total energy dissipated at the unresolved scales.  This was demonstrated by our experiments with a modified \\lamhd, where we (non-consistently) removed the hyperdiffusive term and found the resulting model to fail as a LES.}  \\resp{}{Yet} another way to understand the differences between \\lansa (for incompressible isotropic and homogeneous flows) and \\lamhda is to consider the derivation of these models \\cite{H02a} using the generalized Lagrangian-mean (GLM) formalism \\cite{AM78}.  This form of Lagrangian averaging describes wave, mean-flow interactions.  For the case of weak turbulence, where the nonlinear transfer is dominated by waves, GLM requires in principle no closure. As a result, GLM gives an exact closed theory for the evolution of the wave activity. On the other hand, when there are no waves (as in incompressible Navier-Stokes) or when eddies dominate the transfer, a closure is required.  One possible closure assumes that fast fluctuations are just advected by the mean flow (basically, Taylor's frozen-in hypothesis for the small scale turbulent fluctuations) and leads to the several \"$\\alpha$-models\" that include \\lansa and \\lamhd.  In this context, it is not surprising for subgrid models based on GLM to perform better in the presence of Alfv\\'en waves (for \\lamhd) or Rossby and gravity waves (for the Lagrangian-averaged primitive equations \\cite{HeHoPe+2008}).  The more relevant the waves are to the dynamics, and to the non-linear coupling of modes in the system, the less relevant is the hypothesis behind the closure.  Furthermore, the $\\alpha$-model equations can then be expected to be a better approximation to the problem at hand, that is, to be closer to an exact closure of the original system of equations.  \\vier{In the fluid case, the application of the ``Taylor'' closure that smaller-than-$\\alpha$ scale fluctuations are swept along by the large-scale flow results in the fluctuations having greatly reduced interactions.  This allows for a reduction in computational expense and leads to the super-filter-scale bottleneck by quenching spectrally non-local interactions.  In the \\lamhda case, the small-scale $\\vec{z}^+$ ($\\vec{z}^-$) fluctuations are swept along by the large-scale $\\bar{\\vec{z}}^-$ ($\\bar{\\vec{z}}^+$) flow. Small-scale fluctuations advected by two different fields may now collide and nonlinearly interact.  The second part of the model is the preferential hyperdiffusion of Alfv\\'en waves with wavelengths shorter than $\\alpha$.  This damps rather than quenches nonlinear interactions among the small scales.  This more gentle suppression of the transfer of energy to smaller scales reduces the numerical resolution requirements without forming a bottleneck.}  It was noted in \\cite{MMP05b} when assessing the properties of \\lamhda in the dynamo context that the overall temporal evolution was satisfactory, e.g. with a correct growth rate, although the growth of the magnetic seed field started slightly earlier in the \\lamhda run than in the DNS. One can speculate as to whether this delay is linked to the super-bottleneck effect of \\lansa (which prevails when the magnetic field is negligible compared to the velocity, the two modeling approaches, \\lamhda and \\lans, being dynamically consistent).  This point is left for future work; one could determine as well at what ratio of magnetic to kinetic energy the overshooting of spectra in \\lansa disappears for \\lamhd.  Also deserving of a separate study is to investigate the behavior of \\lamhda when anisotropies that appear at small scales \\cite{MiPo2007a} are present; this would be essential when a uniform magnetic field is imposed to the overall flow. The evaluation of the behavior of the model when computing spectra in the perpendicular and parallel directions (with respect to a quasi-uniform magnetic field, computed by locally averaging the field in a sphere of radius comparable to the integral scale) remains to be done but is somewhat time consuming. An analysis of the structures that develop in the highly turbulent \\lamhda flow studied in the preceding section is also left for future work; of particular interest is the occurrence of Kelvin-Helmholtz like roll-up of current sheets as observed at high resolution \\cite{MiPo2007a}; however, the choice of the parameter $\\alpha$ in the present paper was made on the basis of questioning the existence or lack thereof of a rigid-body high-wavenumber $k^{+1}$ spectrum and, thus, was not optimized for the study of the inertial range properties of the flow for which a much smaller value of the length $\\alpha$ could be used.   Finally, how far resolution can be reduced when using \\lamhda as a LES for various statistics of interest will also require further detailed study. The present study shows that, to reproduce the super-filter-scale energy spectrum in three dimensions, gains by a factor of 1300 in computing time can be achieved. The need to reproduce higher order statistics can decrease these gains. As an example, in two-dimensional MHD, it was shown that gains when using \\lamhda as a subgrid model depend for high order moments on the order that one wants to see to be accurately reproduced \\cite{PGHM+06}."
    },
    "0806/0806.4358_arXiv.txt": {
        "abstract": "We consider an ensemble of closed strings in a compact space with stable one cycles and compute the speed of sound  resulting from string thermodynamics. Possible applications to the issue of Jeans instability in string gas cosmology are mentioned. ",
        "introduction": "\\label{sec:1} String gas cosmology \\cite{BV} (see \\cite{SGrevs} for some recent reviews)  is a scenario of the very early universe based on taking into account new degrees of freedom and new symmetries which characterize string theory but are absent in point particle field theories. In string gas cosmology (SGC), matter is treated as a gas of closed strings. According to SGC, the radiation phase of standard cosmology was preceded not by a period of inflation, but by a quasi-static Hagedorn phase during which the temperature of the string gas hovers close to the Hagedorn temperature \\cite{Hagedorn}, the maximal temperature of a gas of closed strings in thermal equilibrium. Thus, it is hoped that the scenario will not have a singularity in its past. SGC provides a possible dynamical explanation \\cite{BV} of why there are only three large spatial dimensions (see, however, \\cite{Columbia,Danos} for some concerns), provides a simple and physical mechanism of stabilizing all of the size \\cite{Patil2} (see also \\cite{Watson,Patil1}) and shape \\cite{Edna} moduli of the extra dimensions, leaving only the dilaton multiplet un-fixed. The dilaton, in turn, can be fixed \\cite{Danos2} by making use of non-perturbative effects like gaugino condensation.  It has recently been suggested \\cite{NBV,Ali,BNPV1} that thermal fluctuations of a gas of closed strings on a toroidal space will, in the context of the background cosmology of the string gas scenario, generate an almost scale-invariant spectrum of adiabatic, coherent cosmological fluctuations, fluctuations which have all of the correct properties to explain the recent high precision observations of cosmic microwave background anisotropies . Thermal fluctuations are given by the specific heat capacity. For a string gas living in a toroidal space the heat capacity depends on the area of torus. The deviation from extensivity is, in fact, a result of the enormous number of winding modes that become excited close to Hagedorn temperature. This result suggests that the heat capacity $c_V(R)$ scales holographically as a function of the radius $R$ of a volume embedded inside the total box, i.e. $c_V(R) \\sim R^2$. This leads to the scale-invariance of the spectrum of cosmological fluctuations \\cite{NBV,Ali,BNPV1}, and implements a concrete realization of the argument that holography almost always leads to scale invariance of perturbations \\cite{Banks}. SGC, in fact, produces a scale invariant spectrum of scalar metric perturbations with a slight red tilt, again like what is obtained in most inflationary models. The spectrum of gravitational waves, on the other hand, is characterized \\cite{BNPV2} by a spectrum with a slight blue tilt, unlike the slight red tilt which is predicted in inflationary models. This yields a way to observationally distinguish between the predictions of SGC and inflation \\cite{Stewart}.  In order to develop string gas cosmology into a viable alternative to inflationary cosmology, further issues need to be addressed. How are the horizon, flatness, size and entropy problems of Standard Big Bang cosmology addressed in string gas cosmology? What explains the overall isotropy of the cosmic microwave background (horizon problem)? why is the universe so large (size problem) and contains so much entropy (entropy problem) compared to what would be expected if the universe emerges from the Big Bang with a scale commensurate with its initial temperature? Recall that it was these questions which motivated the development of inflationary cosmology \\cite{Guth}.  If SGC is embedded in a bouncing universe cosmology, as can be realized \\cite{Biswas} in the ghost-free higher derivative gravity theory discussed in \\cite{Siegel}, then the horizon, size and entropy problems do not arise \\footnote{See \\cite{Natalia} for an approach to resolving the size problem without assuming a bouncing cosmology.}. The flatness problem, however, persists. In this Note we focus on the flatness problem. The flatness problem has two aspects: firstly, what explains the overall nearly spatially flat geometry of the current universe, and, secondly, why there are no large amplitude small-scale fluctuations which will collapse into black holes and prevent the SGC scenario from working. It is this second aspect of the flatness problem which will be addressed in this Note (the first one is again not present if the Hagedorn phase of SGC occurs around the time of a cosmological bounce and we assume that at some point in the contracting phase the universe was as large as it is today and the spatial curvature is comparable to the current spatial curvature).  In cosmology, the Jeans length determines whether small scale structures can collapse. The Jeans length, in turn, is determined by the speed of sound. If the speed of sound is of the order unity (in units in which the speed of light is unity), then the Jeans length, the length below which perturbations are supported against collapse by pressure, is given by the Hubble length $H^{-1}$, where $H$ is the expansion rate of space (see e.g. \\cite{MFB} for a review of the theory of cosmological fluctuations and \\cite{RHBrev} for an overview). In a radiation-dominated universe the speed of sound is $c_s^2 = 1/3$ and hence small-scale instabilities do not occur. In contrast, in a matter-dominated universe $c_s = 0$ and Jeans instability on small scales occurs. It is this gravitational instability which leads to the formation of nonlinear structures in the present universe.  In this Note we compute the speed of sound in the Hagedorn phase of SGC. Matter in the Hagedorn phase is dominated by relativistic strings containing both momentum and winding modes. On this basis, we might expect that locally flat space is protected against the Jeans instability. On the other hand, the net pressure is small since the positive pressure of the momentum modes is cancelled by the negative pressure of the winding modes. This leads to the expectation that the speed of sound might be small.  In this Note we show that the speed of sound calculated using the microcanonical ensemble is very small and positive. As soon as the radius of torus becomes two orders of magnitude or so larger than the string length, the speed of sound goes to zero exponentially with the radius. On this basis, it appears that SGC may suffer from a small-scale Jeans instability problem, since, in order to make contact with the current universe, the radius $R$ needs to be at least 1 mm (if the string scale is comparable to the Grand Unification scale). However, since the microcanonical ensemble has limited applicability in describing the physics of subspaces of the entire space, we need other techniques in order to unambiguously resolve the Jeans instability puzzle of SGC.  There has been some previous work on inhomogeneities in the early phases of SGC. Specifically, in the context of taking the background of string gas cosmology to be described by dilaton gravity, as in \\cite{TV,BEK,Borunda}, it has been shown \\cite{Watson0} that there is no growth of cosmological inhomogeneities. Dilaton gravity cannot, as is now realized \\cite{Betal,KKLM}, provide a consistent background for the Hagedorn phase of string gas cosmology because of the rapid time variation of the dilaton. Since the dilaton velocity was playing an important role in the considerations of \\cite{Watson0}, we have to revisit the issue of stability towards growth of fluctuations.  In the following section, we will review the microcanonical approach to string thermodynamics. We discuss the physics of an ideal gas of strings, the equilibrium conditions and the energy distribution in the gas. In Section 3 we will focus on the microcanonical ensemble to find the speed of sound for string gas fluctuations. In the absence of a well-defined canonical ensemble, the interpretation of the speed of sound becomes subtle. The final section contains a discussion of the interpretation of results and conclusions. ",
        "conclusions": "In this note we have computed the speed of sound in a gas of closed strings on a three dimensional torus, assuming that the toroidal radii are equal. Our main result is that the speed of sound is positive but highly suppressed when $R \\gg l_s$.  The positivity of the speed of sound has implications for the stability of cosmological perturbations. In general relativity, for adiabatic fluctuations, the scalar metric perturbation variable $\\Phi$ (the fluctuation of the metric component $g_{00}$ in longitudinal gauge - the coordinate system in which the metric is diagonal - whose physical meaning is that of the relativistic generalization of the Newtonian gravitational potential) satisfies the following equation of motion (see e.g. \\cite{MFB} for a review of the theory of cosmological perturbations) \\be \\label{Phieq} \\Phi^{\\prime \\prime} + 3 {\\cal H} \\Phi^{\\prime} - c_s^2 \\nabla^2 \\Phi + \\bigl[ 2 {\\cal H}^{\\prime} + (1 + 3 c_s^2) {\\cal H}^2 \\bigr] \\Phi \\, = \\, 0. \\ee In the above, we have used conformal time $\\eta$ which is related to the physical time variable $t$ via the cosmological scale factor $a(t)$. \\be dt \\, = \\, a(t) d\\eta \\, . \\ee The derivative with respect to conformal time is denoted by a prime, and ${\\cal H}$ is the scale factor in conformal time. Neglecting the expansion of the universe, the solutions of (\\ref{Phieq}) are oscillatory if the speed of sound is positive: \\begin{equation} \\Phi_k(\\eta) \\, = \\, Ae^{ikc_s\\eta} + Be^{-ikc_s\\eta}  \\,  . \\end{equation} In particular, there is no instability of the metric fluctuations.  On small scales, however, we should be concerned with the instability of density fluctuations. On sub-Hubble scales and in a matter-dominated universe, the density fluctuation $\\delta \\rho$ obeys the following equation of motion \\be \\label{Newtdens} \\ddot{ \\delta \\rho} - c_s^2 \\nabla^2 \\delta \\rho - 4 \\pi G \\rho_0 \\delta \\rho \\, = \\, 0 \\, , \\ee once again neglecting both the expansion of space and the possible presence of entropy fluctuations. This equation shows that fluctuations on scales larger than the Jeans length (given by the wave number for which the second and the third term in the above are equal) grow. If the speed of sound is comparable to the speed of light, the Jeans length is of the order of the Hubble length $H^{-1}$ and there is no Jeans instability problem. However, if the speed of sound is negligible (as in our case) then there is a potential Jeans instability problem \\footnote{Note that the expression for the speed of sound (\\ref{speed}) is independent of the string coupling $g_s$. The string thermodynamics is defined in the limit of small $g_s$ the same way as in \\cite{AW}. The Jeans length can be found using $R\\geq \\sqrt{\\frac{\\pi c_s}{g_s^2\\rho}}$.}.  Since in the Hagedorn phase of string gas cosmology $p/\\rho$ is vanishingly small for large values of $R$, the equation of state is like that of a matter-dominated universe. However, since the basic objects which make up the gas are not point particles but extended relativistically moving strings, it is unlikely that equation (\\ref{Newtdens}) applies to describe matter fluctuations in SGC. Putting these considerations together, we conclude that our study has so far not resolved the concern that SGC might suffer from a Jeans instability problem.  The interpretation of speed of sound in SGC is somewhat subtle. The reason is that in the thermodynamics of a string gas the equivalence of the microcanonical and canonical ensembles is lost due to the exponentially growing density of states close to the Hagedorn temperature. Thus, results concerning the change in thermodynamical quantities as the size of the entire sample box is changed cannot immediately be applied to questions related to sub-boxes, like the question we are addressing here.  In standard thermodynamics the analogy between the canonical ensemble and the microcanonical ensemble  in the limit of large box sizes comes about because we can approximate the partition function -the Laplace transform of the density of states - by a saddle-point approximation, finding the average value of $E$ for any sub-system of interest with small fluctuations about the mean. However, if the density of states grows exponentially with energy, like in the case of strings, the saddle point approximation breaks down. If one tries to push the canonical ensemble further, one obtains divergent fluctuations about the mean value.  Although non-extensive thermodynamics might seem very counter-intuitive, even classically in the presence of gravity it is somewhat inevitable. According to the ergodicity theorem, a system that evolves for a long time will be able to reach any small neighborhoud of a point in phase space. Therefore, there are trajectories which run into regions in phase space that correspond to black holes. Once these black holes are nucleated out of thermal fluctuations, they grow and render the canonical ensemble ill-defined."
    },
    "0806/0806.3863_arXiv.txt": {
        "abstract": "{} {X-ray emission is an important diagnostics to study magnetic activity in very low mass stars that are presumably fully convective and have an effectively neutral photosphere.} {We investigate an archival XMM-Newton observation of LHS 2065, an ultracool dwarf with spectral type M9.} {We clearly detect LHS 2065 at soft X-ray energies in less than 1\\,h effective exposure time above the 3\\,$\\sigma$ level with the PN and MOS1 detector. No flare signatures are present and we attribute the X-ray detection to quasi-quiescent activity. From the PN data we derived an X-ray luminosity of $L_{\\rm X} = 2.2 \\pm 0.7 \\times 10^{26}$~erg/s in the 0.3\\,--0.8\\,keV band, the corresponding activity level of log~$L_{\\rm X}$/$L_{\\rm bol}\\approx -3.7$ points to a rather active star. Indications for minor variability and possible accompanying spectral changes are present, however the short exposure time and poor data quality prevents a more detailed analysis.} {LHS 2065 is one of the coolest and least massive stars that emits X-rays at detectable levels in quasi-quiescence, implying the existence of a corona. } ",
        "introduction": "The ultracool dwarf star LHS~2065 (GJ~3517) with a spectral type of M9V is a nearby star located at a distance of 8.6\\,pc. These late-type, very low mass stars are generally assumed to be fully convective and hence a solar-type dynamo is not expected to work. Further, their cool photospheres should be effectively neutral, leading to a high electric resistivity. Therefore the presence of magnetic activity phenomena in the outer atmospheric layers of these stars, especially outside transient events like flares, is remarkable. X-ray emission, among other diagnostics, can put important constraints on the possible dynamo and activity models.  LHS~2065 is among the latest main-sequence stars detected in X-rays. While \\cite{fle93} derived an 2\\,$\\sigma$ upper limit on its X-ray luminosity of $L_{\\rm X}< 3.7 \\times 10^{26}$~erg/s from the RASS (ROSAT All Sky Survey) data, \\cite{schmitt02} report a detection of an X-ray flare in 68\\,ks pointed observations performed in April/May and Oct./Nov. 1997 with the ROSAT HRI. During this event with a decay time of 1.4\\,h, a peak luminosity of $L_{\\rm X}= 4 \\times 10^{27}$~erg/s was derived using a nominal count to flux conversion factor since the ROSAT HRI had no intrinsic energy resolution. In addition to this large flare event, a smaller flare is present in the observations and some indications for 'quiescent' X-ray emission of LHS~2065 were deduced from the absence of clear flaring signal in the light curve. Specifically, \\cite{schmitt02} report non-flaring X-ray emission from LHS~2065 for a period of a few days in autumn 1997 at the level of $L_{\\rm X}= 2.6 \\times 10^{26}$~erg/s, while for the data taken half a year earlier an upper limit of  $L_{\\rm X} \\le 1.8 \\times 10^{26}$~erg/s was derived. With $L_{\\rm bol}=1.2\\times 10^{30}$\\,erg/s \\citep{goli04}, one arrives at an activity level of  log~$L_{\\rm X}$/$L_{\\rm bol} = -3.7$ in quasi-quiescence, i.e. a level not too far from the saturation level around  log~$L_{\\rm X}$/$L_{\\rm bol} = -3$, while inactive stars like the Sun have log~$L_{\\rm X}$/$L_{\\rm bol} \\approx -7$. Consequently, while being X-ray faint in absolute values, LHS~2065 appears to be a rather active star, despite its low temperature of T$_{eff}\\lesssim 2400$\\,K. Further evidence for significant magnetic activity on LHS~2065 is its  H${\\alpha}$ emission \\citep{moh03}, suggesting the existence of a chromosphere, and in its strong magnetic field of B$f >3.9$~kG \\citep{rei07}. We note that it was not possible to disentangle the parameters {\\it B} and {\\it f} with the applied method that utilizes the profiles of FeH lines. In view of the rarity of X-ray observations of ultracool dwarfs we examined the XMM-Newton observations of LHS~2065 and report in this note a clear X-ray detection of LHS~2065 with the EPIC detector at soft X-ray energies below 1\\,keV, despite short effective exposure times and unfavorable high background conditions. In Sect.\\,\\ref{ana} we describe the observation and data analysis, in Sect.\\,\\ref{res} we present our results and summarize our findings in Sect.\\ref{sum}. ",
        "conclusions": "\\label{sum}  \\begin{enumerate} \\item We detected quasi-quiescent X-ray emission from the ultracool dwarf LHS~2065 with spectral type M9 at soft X-ray energies below 1.0\\,keV with a significance above 3\\,$\\sigma$ in the PN and MOS1 detector. The derived X-ray luminosity of $L_{\\rm X} = 2.2 \\pm 0.7  \\times 10^{26}$~erg/s leads to an activity level of log~$L_{\\rm X}$/$L_{\\rm bol}\\approx -3.7$, pointing to a rather active star. To our knowledge LHS~2065 is the coolest main sequence star with securely detected quasi-quiescent X-ray emission.  \\item XMM-Newton is well suited to study faint, soft X-ray sources, and a deeper X-ray observation of LHS~2065 has the potential to put stronger constraints on the X-ray properties of active stars at the very cool end of the main sequence.   \\end{enumerate}"
    },
    "0806/0806.2990_arXiv.txt": {
        "abstract": "{In isothermal disks the migration of protoplanets is directed inward. For small planetary masses the standard type~I migration rates are so fast that this may result in an unrealistic loss of planets into the stars.} {We investigate the planet-disk interaction in non-isothermal disks and analyze the magnitude and direction of migration for an extended range of planet masses. } { We have performed detailed two-dimensional numerical simulations of embedded planets including heating/cooling effects as well as radiative diffusion for realistic opacities. } {In radiative disks, small planets with $M_{\\rm planet} < 50 M_{\\rm Earth}$ do migrate outward with a rate comparable to absolute magnitude of standard type~I migration. For larger masses the migration is inward and approaches the isothermal, type~II migration rate. } {Our findings are particularly important for the first growth phase of planets and ease the problem of too rapid inward type-I migration. } ",
        "introduction": "\\label{sec:introduction} Planets form in disks surrounding young stars. The growing protoplanets undergo an embedded phase where the gravitational interaction with the ambient gaseous disk results in a change of its orbital elements. For protoplanets with masses below about 30 Earth masses the disk is not disturbed too strongly and the interaction can be treated in the linear approximation. Calculations of the total disk torques acting on the planet lead generally for these small masses to a reduction of the semi-major axis, i.e. to an inward migration \\citep{1979ApJ...233..857G, 1997Icar..126..261W, 2002ApJ...565.1257T, 2004ApJ...602..388T}. It soon turned out that the inward drift of this type-I migration is very fast and the planets might be lost before they can grow to larger objects \\citep{1993Icar..102..150K}. This problem has become more visible after comparing population synthesis models with the characteristics of observed planetary systems \\citep{2004A&A...417L..25A, 2008ApJ...673..487I}. To avoid this rapid phase of inward migration, alternative scenarios have been sought. In a turbulent disk, migration occurs stochastically with inward and outward phases which slows down the migration \\citep{2005A&A...443.1067N}. Departures from the linear regime at around 10-20 $M_{\\rm Earth}$ can also lead to reduced inward migration \\citep{2006ApJ...652..730M}. However, both processes are not sufficient to solve the problem. The planet trap scenario to halt planetary migration \\citep{2006ApJ...642..478M} requires a positive density gradient which may not be given in general.  To simplify the calculations, nearly all of the analytical and numerical studies devoted to study the planet-disk interaction process have focussed on isothermal disks, where the temperature is a given function of the position in the disk. Early work on non-isothermal disks focussed on high mass embedded planets and did not notice a strong effect on migration \\citep{2003ApJ...599..548D, 2006A&A...445..747K}. Using a fully three-dimensional radiative calculations of an embedded small mass planet, \\citet{2006A&A...459L..17P} have shown in a very important work that migration can be significantly slowed down or even reversed when thermal effects are included. Subsequent analysis indicate that this behaviour is related to a radial entropy gradient in the flow \\citep{2008ApJ...672.1054B, 2008A&A...478..245P}. Recently, \\citet{2008arXiv0804.4547P} have shown that for small planet masses, the combination of radiative and viscous diffusion may allow for long-term unsaturated positive torques and possible outward migration.  In this letter we investigate this possibility in more detail for a whole range of planetary masses, which will allow us to estimate its effect on the long term evolution of the planet. For that purpose we perform two-dimensional numerical hydrodynamical simulations of embedded planets in radiative disks. A method to treat the three-dimensional radiative transfer approximately in these 2D simulations will be outlined in the next section. Our results on the migration rate for various masses (in Section 3) indicate that for masses smaller than about 50~$M_{\\rm Earth}$ the torques remain unsaturated in the long run and migration is indeed directed outward, while larger planets drift inward. The consequence for the migration process and the overall evolution of planets in disks is discussed. ",
        "conclusions": "\\label{sec:summary} We have performed an investigation of the migration of planets in disks using two-dimensional numerical simulation including heating/cooling effects as well as radiative diffusion.  Using different formulations of the energy equation, we first show that for a planet mass of $20 M_{\\rm Earth}$, migration is directed inwards in the isothermal and adiabatic situation, while inclusion of radiative effects leads to an outward migration of the planet. This finding supports the torque reversal mechanism in radiative disks due to corotation effects as suggested by \\citet{2008ApJ...672.1054B}.  A detailed parameter study for planetary masses in the range between $10^{-5}$ and $10^{-3} M_\\odot$ shows that the effect is limited to planets in the low mass regime, $M_p \\leq 50 M_{\\rm Earth}$ where corotation effects are indeed important.  Larger mass planets open up gaps in the disk and the migration rate becomes similar to the isothermal case.  Our findings are particularly important for the first growth phase of planets and may ease the problem to too rapid inward type~I migration. Depending on the mass accretion rate onto the planet the growing planetary embryos can spend an extended time span in an outward migration phase and avoid loss into the star. However, close-in planets exist for a range of planetary masses, according to the observations. Thus, a significant and long outward migration phase may create new difficulties. Whether this problem really exists can only be answered by following the actual long-term migration of planets through the disk including its mass growth.  Constructing the necessary migration histories of planetary cores, to be used in population synthesis models, requires suitable scaling laws for the migration process as a function of disk parameter ($\\Sigma(r)$, $T(r)$) for realistic accretion disks with net mass flow.  The present study can be used as a starting point for these larger parameter studies. The inclusion of three-dimensional effects and additional physics (MHD, self-gravity, mass accretion) will make the models even more realistic in the future."
    },
    "0806/0806.0392_arXiv.txt": {
        "abstract": "We examine the constraints on high-redshift star formation, ultraviolet and X-ray pre-ionization, and the epoch of reionization at redshift $z_r$, inferred from the recent WMAP-5 measurement, $\\tau_e = 0.084 \\pm 0.016$, of the electron-scattering optical depth of the cosmic microwave background (CMB).  Half of this scattering can be accounted for by the optical depth, $\\tau_e =$ 0.04--0.05, of a fully ionized intergalactic medium (IGM) at $z \\leq z_{\\rm GP} \\approx$ 6--7, consistent with Gunn-Peterson absorption in neutral hydrogen. The required additional optical depth, $\\Delta \\tau_e = 0.03 \\pm 0.02$ at $z >$ \\zGP, constrains the ionizing contributions of ``first light\" sources.  WMAP-5 also measured a significant increase in small-scale power, which lowers the required efficiency of star formation and ionization from mini-halos. Early massive stars (UV radiation) and black holes (X-rays) can produce a partially ionized IGM, adding to the residual electrons left from incomplete recombination.  Inaccuracies in computing the ionization history, $x_e(z)$, and degeneracies in cosmological parameters ($\\Omega_m$, $\\Omega_b$, $\\sigma_8$, $n_s$) add systematic uncertainty to the measurement and modeling of $\\tau_e$.  From the additional optical depth from sources at $z >$ \\zGP, we limit the star-formation efficiency, the rate of ionizing photon production for Pop~III and Pop~II stars, and the photon escape fraction, using standard histories of baryon collapse, minihalo star formation, and black-hole X-ray preionization. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0806/0806.1195.txt": {
        "abstract": "{}{}{}{}{} % 5 {} token are mandatory  \\abstract % context heading (optional) %{High-mass protostellar objects (HMPOs) coincide with massive clumps which are good candidates to correspond to the early phases of the formation of high-mass stars.} {The study of physical  and chemical properties of massive protostars is critical to better understand the evolutionary sequence which leads to the formation of high-mass stars.} % Aims {IRAS~18151$-$1208 is a nearby massive region ($d =3$~kpc, $L \\sim 2$\\ttp{4}~\\lsol) which splits into three cores: MM1, MM2 and MM3 (separated by 1\\arcmin--2\\arcmin). We aim at (1) studying the physical and chemical properties of the individual MM1, MM2 and MM3 cores; (2) deriving their evolutionary stages; (3) using these results to improve our view of the evolutionary sequence of massive cores.} % Methods {The region was observed in the %CS (2--1), (3--2), (5--4), C$^{34}$ (2--1), %(3--2), H$_2$CO (3$_{22}$--2$_{21}$), (3$_{12}$--2$_{11}$), %HCO$^+$ and H$^{13}$CO$^+$ (1--0), and N$_2$H$^+$ (1--0). CS, C$^{34}$S, H$_2$CO, HCO$^+$, H$^{13}$CO$^+$, and N$_2$H$^+$ lines at mm wavelengths with the IRAM 30m and Mopra telescopes. We use 1D and 2D modeling of the dust continuum to derive the density and temperature distributions, which are then used in the RATRAN code to model the lines and constrain the abundances of the observed species.} % Results heading (mandatory) {All the lines were detected in MM1 and MM2. MM3 shows weaker emission, or even is undetected in HCO$^+$ and all isotopic species. MM2 is driving a newly discovered CO outflow and hosts a mid-IR-quiet massive protostar. The abundance of CS is significantly larger in MM1 than in MM2, but smaller than in a reference massive protostar such as AFGL~2591. In contrast the N$_2$H$^+$ abundance decreases from MM2 to MM1, and is larger than in AFGL~2591.} % conclusions heading (optional), leave it empty if necessary {Both MM1 and MM2 host an early phase massive protostar, but MM2 (and mid-IR-quiet sources in general) is younger and more dominated by the host protostar than MM1 (mid-IR-bright). The MM3 core is probably in a pre-stellar phase. We find that the N$_2$H$^+$/C$^{34}$S ratio varies systematically with age in the massive protostars for which the data are available. It can be used to identify young massive protostars.  %A good indication of a decrease with time of the line ratio N$_2$H$^+$/C$^{34}$S is observed and proposed to be systematically used to search for young massive protostars. } ",
        "introduction": "How high-mass stars form is still an open issue \\citep[\\textit{e.g.}][]{zinnecker2007}. It is particularly not clear whether the formation process for OB/high-mass stars is different from the way solar-type/low-mass stars form. Stars more massive than $\\sim$10~\\msol\\ may form like a scaled-up version (high accretion rates) of the single (or monolithic) collapse observed for the low-mass stars, or require a more complex process in which competitive accretion inside the central regions of a forming cluster may play a decisive role \\citep[\\textit{e.g.}][]{bonnell2004}. In the first scenario, the observed high-mass clumps (100 to 1000~M$_\\odot$; 0.5~pc in size) are expected to fragment to form self-gravitating high-mass cores (10 to 100~M$_\\odot$; 0.01$-$0.1~pc in size; \\textit{e.g.} \\citealt{krumholz2007}) which would collapse individually to form a massive single or binary stars. In the second scenario, the competitive accretion is expected to occur inside the high-mass clumps. The study of the properties of high-mass clumps is therefore a central observational issue to progress in our understanding of the earliest phases of high-mass star formation.  From a selection of IRAS sources not associated with any bright radio source but having the IRAS colors of \\uchii\\ regions \\citep[as defined by][]{wood1989}, \\citet{sridharan2002} have built a sample of so-called High Mass Protostellar Objects (hereafter HMPOs) which would correspond to the pre-\\uchii\\ phase of the formation of high-mass stars. \\citet{beuther2002a} found that the HMPOs were systematically associated with massive clumps as detected in the dust continuum and in CS line emission. These clumps are good candidates to correspond to the earliest phases of high-mass star formation. They have not yet formed any bright H{\\sc{ii}} regions and contain a large amount of gas at high densities. The precise evolutionary stages of the individual HMPOs might however be very diverse and they require to be derived individually through dedicated, detailed studies.  Recently, \\citet{motte2007} investigated the whole Cygnus~X complex and obtained a first unbiased view of the evolutionary scheme for high-mass clumps and cores. Like for the HMPOs, the massive clumps ($\\sim$0.5~pc in size) in Cygnus~X could be resolved into massive cores ($\\sim$0.1~pc). Roughly half of these cores were found to be very bright in the (mid)-infrared \\citep[such as AFGL~2591; \\textit{e.g.}][]{vandertak1999}, and therefore very luminous: the mid-IR high-luminosity massive cores. The other half are weak or not detected in the mid-IR, hereafter the mid-IR-quiet massive cores. Surprisingly all the mid-IR-quiet massive cores were however found to drive powerful SiO outflows. They could therefore be safely understood as the precursors of the infrared high-luminosity massive cores.   IRAS~18151$-$1208 is a rather typical ($L \\sim 10^4$~\\lsol) and relatively nearby (3~kpc) HMPO \\citep{sridharan2002}. \\citet{beuther2002a} show that the clump actually splits into four individual cores MM1, MM2, MM3 and MM4 (see Fig.~\\ref{fig:msx-map} for an overview of the region). In addition MM1 seems to further split into two cores, separated by 16\\arcsec.  We hereafter refer to MM1-SW for the secondary peak in the South-West of MM1 which possibly hosts an embedded protostar \\citep{davis2004}.  The weakest core MM4 is clearly outside the main region, hence it will not be further considered in this paper.  A CO outflow toward MM1 was discovered by \\citet{beuther2002b}.  The MM1 and MM2 cores have sizes roughly two times larger, and similar masses (using the same dust emissivity and temperature) than the most massive cores in \\citet{motte2007}. The IRAS source coincides with MM1 and no significant IRAS contribution can be safely attributed to MM2 or MM3. MM3 is the least massive and compact core and could actually be a still quiescent or pre-stellar core. The IRAS~18151$-$1208 region is therefore particularly interesting to study since it hosts three individual cores which could be interpreted as high-mass star formation sites in three different evolutionary stages. Millimeter and sub-millimeter wave observations are reported in \\citet{mccutcheon1995} and \\citet{beuther2002a}. The near-IR counterparts (H$_2$ jets and HH objects) of the CO outflow driven by MM1 have been imaged by \\cite{davis2004}.  Molecular line observations compared with results of line modeling based on a physical model of the source is a classical technique to constrain physical and chemical properties of protostellar objects \\citep[see][for a review]{ceccarelli1996,vandertak2005}. While different approaches can be adopted, the most reliable and most often used method consists in constraining first the physical (density and temperature distributions) model from the dust continuum emission (mid-IR to millimeter wavelengths), and then use this physical model to derive the fractional abundances of molecular species from line emission modeling \\citep{vandertak1999, vandertak2000, schoier2002, belloche2002, hatchell2003}. Due to the dramatic changes in physical conditions inside the protostellar envelopes (increase of density and temperature, radiation, shocks) chemical evolution is observed and can be modeled thanks to chemical network codes. It is even expected that chemistry could provide a reliable clock to date protostellar objects \\citep[\\textit{e.g.}][]{vandishoeck1998,doty2002,wakelam2004}.  After the description of the molecular observations and of the continuum data from the literature (Sect.~2), the results are presented in Sect.~3. Section~4 details the modeling procedure from the fits of the spectral energy distributions (hereafter SEDs) for MM1 and MM2 using MC3D\\footnote{See \\citet{wolf1999} for details.} code (Sect.~4.1.1 and 4.2.1) to the non-LTE calculations of the line profiles and intensities for all observed molecules using RATRAN\\footnote{See \\citet{vandertak1999,vandertak2000b} for details.} code (Sect.~4.1.2, 4.1.3, 4.2.2, and 4.2.3). In Section~5, we discuss the results of this detailed analysis in order to improve our observational view of the evolution of high-mass cores from pre-stellar to high luminosity massive protostars.   \\begin{figure}[t!] \\centering \\includegraphics[width=\\columnwidth]{overlay.jpg}\\\\ \\caption{Image of the MSX 8\\micron\\ emission toward IRAS~18151$-$1208 (color scale) overlaid with the 1.2~mm map (white polygon) by \\citet{beuther2002a} (white contours at 5 and 10~\\%, and grey contours from 20~\\% to 90~\\% of the maximum). The large and small crosses indicate the positions of the IRAS source and of MM1-SW (see text) respectively. The black and white star symbols show the positions of the methanol and water masers respectively. The green polygon displays the region mapped in the present study (see Fig.~\\ref{fig:maps}).} \\label{fig:msx-map} \\end{figure}  \\begin{table}[t!] \\caption{IRAS~18151$-$1208 sources characteristics. Offsets from IRAS position, J2000 coordinates and velocity in the local standard of rest \\citep{beuther2002a} are reported.}             % title of Table \\label{tab:sources}      % is used to refer this table in the text \\begin{tabular}{lccccc} \\hline \\hline source & $\\Delta\\alpha$[\\arcsec]& $\\Delta\\delta$[\\arcsec] & $\\alpha$[J2000] & $\\delta$[J2000]  & $\\varv$ [km/s] \\\\ \\hline MM1 & 13.2 & -4.9 & 18$^h$17$^m$58.0$^s$ & -12$^\\circ$07'27\" & 33.4  \\\\ MM2 & -98.9 & -32.8 & 18$^h$17$^m$50.4$^s$ & -12$^\\circ$07'55\" & 29.7  \\\\ MM3 & -72.3 & 26.5 & 18$^h$17$^m$52.2$^s$ & -12$^\\circ$06'56\" & 30.7  \\\\ \\hline \\end{tabular} \\end{table}  %__________________________________________________________________ ",
        "conclusions": "Here we summarize our conclusions on the massive protostellar objects MM1 and MM2 of the IRAS~18151$-$1208 region.  \\begin{enumerate} \\item The three cores of the region, MM1, MM2 and MM3 are physically linked and have probably been formed from a single parental cloud or clump. \\item We detected for the first time a CO outflow driven by MM2. It clearly establishes the protostellar nature of MM2. In contrast MM3 does not show any outflow activity and is therefore most probably a pre-stellar core. \\item Following the evolutionary scheme discussed in \\citet{motte2007}, MM1 is a high luminosity IR (or IR-bright) massive protostar while MM2 is an IR-quiet massive protostar. \\item We have established that while an inner flattening of the matter distribution is required to reproduce the SED of MM1, a simple 1D spherical geometry is enough to well model molecular line observations. In contrast, MM2 does not even require a inner flattening since it is not detected in the mid-infrared range. % \\item We suppose that MM1 interacted with its environment for a longer time %   than MM2 which is still embedded and shows envelope interaction clues %   through its higher turbulent velocity, its higher H$_2$CO abundance and %   its para-to-ortho ratio greater than 1 in its inner part. \\item A significant depletion of CS in the inner parts of the MM1 and MM2 cores is required to fully reproduce the observed CS line emission. % \\item We suggest that, extrapolating our results, mid-IR-quiet cores are %   younger than mid-IR-bright cores. \\item We find that the abundance ratio between CS and N$_2$H$^+$ could be a very good evolution tracer for high-mass protostellar cores hosting high-mass protostars. \\end{enumerate}"
    },
    "0806/0806.3358_arXiv.txt": {
        "abstract": " ",
        "introduction": "In this article the role of the conformal symmetries in the Jackiw-Pi model is studied at quantum level in $1$-loop approximation. The Jackiw-Pi model, first described in \\cite{JP1,JP2,JP3}, can be seen as an extension of the non-linear Schr\\\"odinger model with $U(1)$ Chern-Simons gauge fields. A similar study for the non-linear Schr\\\"odinger model was performed in \\cite{DeKokVanHolten}.  Like the non-linear Schr\\\"odinger model, a variation on the Jackiw-Pi model has been used to study the physics of gasses of bosonic particles \\cite{Barashenkov}. The Jackiw-Pi model has also been considered in the context of Aharo\\-nov-Bohm scattering \\cite{BergmanLozano}, which essentially is the scattering of charged particles in the plane by a magnetic flux tube. In reference \\cite{BergmanLozano} the renormalization of coupling constant $g^2$ is studied. In this article the full renormalization at $1$-loop level is presented, and at this order of perturbation theory only $g^2$ is seen to renormalize: the matter and gauge fields, mass and electric charge $e$ are found not to renormalize.  The Jackiw-Pi model is a variation on the Abelian Higgs-model, with the Maxwell kinetic term replaced by a Chern-Simons term and with the matter Lagrangian taken to be non-relativ\\-is\\-tic. In the Abelian Higgs-model vortices arise as topological defects. For a specific choice of the coupling constants $e$ and $g^2$, i.e.\\ $g^2 = -e^2$, stationary self-dual vortices arise in the Jackiw-Pi model as well \\cite{JP1, JP2, JP3, Dunne, HorvCSLect}.  Like the non-linear Schr\\\"odinger model, the classical Jackiw-Pi model possesses, next to translation, rotation, Galilei and gauge invariance, also scale invariance and special conformal invariance. Due to the $1$-loop renormalization effects the scale and special conformal symmetries of the Jackiw-Pi model are seen to be broken precisely as in the non-linear Schr\\\"odinger model \\cite{DeKokVanHolten}. However, this breaking does not occur when the renormalized coupling constants $e$ and $g^2$ are fixed to the value $g^2 = \\pm e^2$. In particular, in the case of the minus sign both the conformal symmetries and the vortex solutions are seen to survive quantization.  This paper is structured as follows. Sect.\\ 2 contains a short introduction to the Jackiw-Pi model and its vortex solutions. Sect.\\ 3 reviews the symmetries of the model. Sect.\\ 4 describes the quantization, whilst the 1-loop corrections are calculated in sect.\\ 5. In sect.\\ 6 the scale dependence of the coupling constants is computed, and sect.\\ 7 contains a summary and our conclusions. ",
        "conclusions": "\\label{JPconclusie}  It has been shown that to first order in perturbation theory only the coupling constant $g^2$ is renormalized: neither the fields, nor the electric charge are scale dependent. Moreover, if the coupling constants in the classical model are chosen such that $g^2 = \\pm e^2$, this relation is preserved and $g^2 = g_R^2$ is constant as well.  In all other cases the renormalized coupling constant $g_R^2$ becomes a function of the scale $\\Lambda$ and the scale and special conformal symmetries are anomalous. Since only $g^2$ is renormalized, the argument given in \\cite{DeKokVanHolten} makes it clear that the time-dependence of the corresponding charges $D$ and $K$ is given by \\be \\frac{dK}{dt} & =& -t\\frac{dD}{dt} ; \\\\\\nonumber \\frac{dD}{dt} &=& \\frac{1}{2}\\, \\beta(g_R^2)\\, \\int d^2x\\, \\Fg^{\\dag\\,2} \\Fg^2 = \\frac{1}{4\\pi}\\left(g_R^4-e_R^4\\right) \\int d^2x\\, \\Fg^{\\dag\\,2} \\Fg^2. \\ee From this expression it follows that the conformal symmetries indeed survive quantization in $1$-loop approximation if and only if the condition $g^2 = g_R^2 = \\pm e_R^2$ is satisfied.  It so happens that for the minus sign this condition is exactly the one the coupling constants need to satisfy such that the classical self-dual vortex solutions exist (\\ref{jw.1}). In this special case not only the conformal symmetries survive quantization, but also the classical self-dual solutions. The condition $g^2 = - e^2$ is thus essential for the existence of self-dual vortex solutions in the classical theory and in the quantum theory as well.  Assuming $\\alpha = e^2$ to be less than unity, and choosing $g_R^2$ less than $\\alpha$ at some reference energy scale $\\Lambda_*$, the Jackiw-Pi model is seen to have the unusual property that the model is in the perturbative regime for all values of $\\Lambda$, as shown in figure  \\ref{grplot2b}. For small values of $\\Lambda$ $g_R^2$ is seen to be positive, signifying a repulsive interaction between the particles, and for large values of $\\Lambda$ $g_R^2$ is negative, signifying an attractive interaction between the particles.  In the case $\\alpha$ less than unity, when $g_R^2$ is larger than $\\alpha$ at the reference energy scale, the model has a perturbative regime either for small values of $\\Lambda$ or for very large values, as shown in figure  \\ref{grplot2a}. When $\\alpha$ is larger than one, there only is a perturbative regime for intermediate values of $\\Lambda$ when $g_R^2$ at the reference scale is chosen less than $\\alpha$, as shown in figure \\ref{grplot2b}.      {\\small"
    },
    "0806/0806.4114_arXiv.txt": {
        "abstract": "We present an initial matching of the source positions of the \\chandra\\ \\nb\\ \\xrss\\ to the new UKIDSS-GPS \\nir\\ survey of the \\nb. This task is made difficult by the extremely crowded nature of the region, despite this, we find candidate counterparts to $\\sim 50\\%$ of the \\xrss. We show that detection in the $J$-band for a candidate counterpart to an \\xrs\\ preferentially selects those candidate counterparts in the foreground whereas candidate counterparts with only detections in the $H$ and $K$-bands are more likely to be \\nb\\ sources. We discuss the planned follow-up for these candidate counterparts. ",
        "introduction": "The \\citet{wang02} and \\citet{muno03} \\chandra\\ surveys of the \\nb\\ revealed a large population of weak, hard point-like \\xrss\\ that could account for up to 10\\% of the previously observed ``diffuse'' X-ray emission from the \\nb.  Numerous studies have tried to characterise these newly discovered \\xrss\\ of the \\nb\\ surveys based on their X-ray properties. \\citet{pfah02} suggest that a large fraction may be wind-accreting neutron stars with high mass companions. \\citet{muno03,muno06} and \\citet{ruit06} propose that they could be white dwarfs accreting from main sequence counterparts (cataclysmic variables, polars and intermediate polars). \\citet{will03} and \\citet{liu06} believe them to be neutron stars with low mass companions and \\citet{wu07} have speculated that they could be isolated neutron stars and black holes accreting from the interstellar medium. However, the weak nature of the sources means that positive identification of the majority of the sources is impossible based solely on the X-ray data. Identification of the stellar counterparts to these \\xrss\\ we will allow the differentiation between these possibilities.  As shown in \\citet{band05}, to identify candidate counterparts to \\xrss\\ in the Nuclear Bulge, very high resolutions imaging is required to avoid issues of confusion due to the high stellar density of the region. In addition to that, the extremely high, variable extinction towards the Nuclear Bulge requires these observations to be performed in the \\nir\\ or longer wavelengths. Previously, only the DENIS and \\tmass\\ surveys had observed the \\nb\\ entirely in the \\nir, however as will be shown, their depth and resolution is insufficient for the purposes of identifying the counterparts to the majority of the \\nb\\ \\xrss. ",
        "conclusions": "We have presented an initial matching of the source positions of the \\chandra\\ \\nb\\ surveys of \\citet{wang02} and \\citet{muno03} to the new UKIDSS-GPS \\nir\\ survey of the \\nb. In doing so, we have identified candidate counterparts to $\\sim50\\%$ of the \\xrss\\ in this extremely crowded and heavily extincted region. We show that these candidate counterparts are consistent with the overall stellar distribution of observations towards the \\nb\\ although the relative proportions of the populations are different.  Candidate counterparts that are observed in the $J$-band have a higher tendency to be foreground sources and those observed only in the $H$ and $K$-bands are more likely to be \\nb\\ sources. This is most likely to be an effect of the high levels of extremely spatially variable extinction (see additional poster contribution by Gosling et al. for further details) towards the \\nb.  Further observations including narrow-band \\brg\\ imaging and \\nir\\ spectroscopy using the FLAMINGOS-2 instrument as well as comparison to other datasets such as the \\spitzer\\ observations of the \\nb\\ will be used to better identify the true companions to these \\xrss\\ and so gain an understanding as to the nature of the \\nb\\ \\xrs\\ population."
    },
    "0806/0806.1513_arXiv.txt": {
        "abstract": "Using the Galaxy as an example, we study the effect of $\\vec j \\times \\vec B$ force on the rotational curves of gas and plasma in galaxies. Acceptable model for the galactic magnetic field and plausible physical parameters are used to fit the flat rotational curve for gas and plasma based on the observed baryonic (visible) matter distribution and $\\vec j \\times \\vec B$ force term in the static MHD equation of motion. We also study the effects of varied strength of the magnetic field, its pitch angle and length scale on the rotational curves. We show that $\\vec j \\times \\vec B$ force does not play an important role on the plasma dynamics in the intermediate range of distances $6-12$ kpc from the centre, whilst the effect is sizable for larger $r$ ($r \\geq 15$ kpc). ",
        "introduction": "Observed flat rotational curves of many galaxies have been subject of long-term controversy. The observational fact that the azimuthal velocity of gas and stars in the galactic plane is constant over a large range of the distances from the centre of a galaxy has yielded two main explanations. In an attempt to save the assertion that the Newtonian gravitational theory holds over the cosmological distances, one such theory assumes the presence of non-baryonic massive dark halo surrounding a spiral disk. In this scenario, gravitational acceleration $GM_<(r)/r^2$ which balances the centrifugal acceleration $V^2(r)/r$, is assumed to vary as $1/r$. This means that the mass enclosed within a certain radius $r$, $M_<(r)$, scales as $\\propto r$. However, this is not what is observed at large radii of the Galaxy. The second possible explanation of the flat rotational curves is that the Newtonian gravity does not apply on cosmological scales and further modifications are due (Milgrom 1983). Historically the latter explanation was not favoured due to absence of the general relativistic extension of the theory. However, this drawback was alleviated by the formulation of the generalisation of Einstein's general relativity based on a pseudo-Riemannian metric tensor and a skew-symmetric rank three tensor field, called metric-skew-tensor gravity (MSTG). The latter leads to a modified acceleration law that can explain the flat rotation curves of galaxies and cluster lensing without postulating exotic dark matter (Moffat 2005). Recently, Brownstein \\& Moffat (2006) have shown that MSTG can provide a good explanation to the flat rotational curves of a large sample of low and high surface brightness galaxies and an elliptical galaxy. Their MSTG fits were compared to those obtained using Milgrom's phenomenological MOND model and to the predictions of the Newtonian / Kepler acceleration law.  In this work, using the Galaxy as an example, we study the effect of $\\vec j \\times \\vec B$ force on the rotational curves of gas and plasma in galaxies. We use an acceptable model for the galactic magnetic field and plausible physical parameters to fit the flat rotational curve for gas and plasma based on the observed baryonic (visible) matter distribution and $\\vec j \\times \\vec B$ force term in the static MHD equation of motion. It should be mentioned that the present work was complete when author became aware of a similar earlier work by Nelson (1988). The latter studied the dynamical effect of magnetic stress on the tenuous outer gaseous discs of galaxies. Nelson (1988) used an earlier, less observationally constrained model (Sofue et al. 1986) for the magnetic field. No fit to an observational rotational galactic curve was presented. Here we advance the earlier hypothesis by choosing a more realistic model for the galactic magnetic field, as well as perform fit to the Milky Way rotational curve.  Other significant previous developments include:  Battaner et al.~(1992) who argue that an azimuthal magnetic field can carry slightly ionised gas with the general galactic rotation, rendering dark matter unnecessary. It was shown for the illustrative case of M31, that a magnetic field of 6 $\\mu$G is required, and the synchrotron emission of relativistic electrons in this field is compatible with the observations. However this was not without a subsequent debate (Katz 1994). More recent arguments in favour of magnetic fields, in this context, were also presented in (Battaner \\& Florido 2000; Battaner et al. 2002; Battaner \\& Florido 2007). However, based on virial constraints Sanchez-Salcedo \\& Reyes-Ruiz (2004), on contrary, show that azimuthal magnetic fields hardly speed up H I disks of galaxies as a whole. This demonstrates that the role of magnetic fields in the rotational curves of galaxies is an active area of research  and there is no general agreement to date.  Our analysis shows that $\\vec j \\times \\vec B$ force does not play an important role on the plasma and gas dynamics in the intermediate range of distances $6-12$ kpc from the centre, whilst the effect of the force is significant for larger $r$ ($r>15$ kpc). ",
        "conclusions": "Earlier work of Nelson (1988) conjectured about importance of the dynamical effect of magnetic stress on the tenuous outer gaseous discs of galaxies. One of the conclusions of Nelson (1988) was that an increase in the pitch angle of the magnetic field yields higher rotational velocities (this result is also corroborated in our model -- see our Fig.2). The model presented here is an improvement on Nelson (1988) work in that we used a more realistic magnetic field model of Han \\& Qiao (1994), and we perform an actual fit to the Milky Way rotational curve. The latter is possible because our model is simpler and presents an analytical expression for the rotational velocity Eq.(10) as opposed to the need for solving an ordinary differential equation (Eq.(7) from Nelson (1988)). Other previous developments include (Katz 1994; Battaner et al.~1992; Battaner \\& Florido 2000; Battaner et al. 2002; Battaner \\& Florido 2007; Sanchez-Salcedo \\& Reyes-Ruiz 2004).  In the present study we investigate the effect of $\\vec j \\times \\vec B$ force on the dynamics of gas and plasma. We show that as far as rotational curve of gas and plasma of the Milky Way is concerned, inclusion of $\\vec j \\times \\vec B$ force only provides a tolerable fit to the rotational curve of the Galaxy  for $r>15$ kpc from the centre, but fails in the intermediate range  $6-12$ kpc. In principle, a tolerable fit can be obtained for all radii with the stronger magnetic field of $B_0 \\geq 11 $ $\\mu$G, but such high values are not observed. Further study is needed to clarify whether the model formulated in this work can be used to fit rotational curves of other known galaxies where the magnetic fields are stronger.  Other weaknesses of this model include:  (i) How well galactic plasma couples to the magnetic field (for $\\vec j \\times \\vec B$ to be effective). Naturally this coupling is prescribed by the degree of ionisation of the medium, which in turn, is prescribed by the Saha equation and is sensitive to the temperature. In general, initial temperatures of galaxies are expected to be high because so called virial temperature (page 557 from (Gilmore et al. 1989)) $T_{\\rm virial} \\simeq G M m_p /(k R)$, where symbols have usual meaning, for a typical size galaxy is of the order of $10^6$ K. However, after cooling phase galactic discs are much cooler at about $\\simeq 10^4$ K. Quireza et al. (2006) quote electron temperatures in the disc of galaxy of the order of $10^4$ K which means that degree of ionisation of the galactic disc is sufficient to couple plasma to the magnetic field and $\\vec j \\times \\vec B$ force. After all, solar photosphere which is at temperature of only 6000K is commonly described by MHD approximation, despite low degree of ionisation and the presence of large concentration of neutrals. Also, in addition to thermal collisions some significant ionisation may be provided by the cosmic rays (mostly protons) that are accelerated at the bow and termination shocks. A substantial flux of cosmic rays is produced in a shock at Galactic north, a direction toward which our Galaxy has long been known to be moving in the Local Supercluster with the velocity of 200 km s$^{-1}$ (Medvedev \\& Melott 2007).   (ii) The origin of the magnetic field in the galaxy itself is deeply coupled with the Galaxy's dynamics and MHD via the dynamo mechanism. The field strength and morphology are dependent on the dynamics of the plasma, which is a function of density, temperature, turbulent velocity, and galactic rotation. Therefore, the centrifugal force due to galactic rotation acts both on the plasma  and magnetic field, and not on the plasma alone. Such advanced topics are naturally beyond the scope of the simple model presented here.  The overall conclusion of this work is that $\\vec j \\times \\vec B$ does not play an important role in the plasma dynamics in the intermediate range of distances $6-12$ kpc from the centre, whilst the effect is considerable for larger $r$ ($r \\geq 15$ kpc)."
    },
    "0806/0806.4270_arXiv.txt": {
        "abstract": "The afterglow of GRB 050401 presents several novel and interesting features : \\begin{enumerate} \\item An initially faster decay in optical band than in X-rays. \\item A break in the X-ray light curve after $\\sim$ 0.06 day with an unusual slope after the break. \\item The X-ray afterglow does not show any spectral evolution across the break while the R band light curve does not show any break. \\end{enumerate}  We have modeled the observed multi-band evolution of the afterglow of GRB 050401 as originating in a two component jet, interpreting the break in X-ray light curve as due to lateral expansion of a narrow collimated outflow which dominates the X-ray emission. The optical emission is attributed to a wider jet component. Our model reproduces all the observed features of multi-band afterglow of GRB 050401.  We present optical observations of GRB 050401 using the 104-cm Sampurnanand Telescope at ARIES, Nainital. Results of the analysis of multi-band data are presented and compared with GRB 030329, the first reported case of double jet.\\\\ ",
        "introduction": "\\label{sec:Introduction} The optical and X-ray light curves of Gamma Ray Burst (GRB) afterglows, in the simplest cases, show a power law decay with an index $\\alpha \\sim 1.0$. Deceleration of the relativistic shock wave generated by the explosion which results in GRB can explain the power law decay of the GRB afterglows. The most common deviation from the power law decay behaviour of the afterglow light curves is an achromatic break seen in the light curve. This break has been seen in a significant number of GRB afterglows and has been successfully explained as being due to the sideways expansion of the collimated ejecta from the explosion. In the post \\emph{Swift} era, many more deviations from this simple behaviour of the afterglow light curve have been detected. \\emph{Swift} with its capabilities of quick slewing towards the source has been able to observe GRB afterglows as early as a few tens of seconds after the burst. In this early part of the evolution the GRB afterglows commonly exhibit a steep decay with $\\alpha \\sim$ 3 to 5 with the usual definition $F_{\\nu}(t) \\propto t^{-\\alpha}\\nu^{-\\beta}$ where $F_{\\nu}(t)$ is the observed afterglow flux at frequency `$\\nu$' and time $t$. The phase of steep decay lasts for about a few hundred seconds after which a slower decay, with $\\alpha \\sim 0.5$, of the afterglow starts. About a few thousands of seconds after the burst the afterglow starts decaying steeply again with $\\alpha > 1.0$.  Many GRB afterglows observed by \\emph{Swift} show puzzling features in the light curves like (1.) early steep decay [$\\alpha \\sim 3 ~to~ 5$] and (2.) Chromatic breaks (breaks seen in some wavebands but not others) with $\\Delta \\alpha \\sim 1.0$ which are difficult to explain using the standard fireball model \\citep{Rees1992,Meszaros1993}. It has been shown by \\citet{PTO2006,RW2006} that the puzzling features of the X-ray afterglow light curves can be fitted using one or two components with exactly the same empirical functional form, viz. an exponential fall followed by a powerlaw decay of flux with time, although it has not yet resulted into any physical understanding of the behaviour of the X-ray afterglow. While there is no clear understanding of the early steep decays of GRB afterglows, a few plausible explanations have been put forward : see e.g. \\citet{Steep01,Steep02}. The flat decay of X-ray afterglow light curves which follows the steep decay have been, in some cases, explained as being due to energy injection from the central engine, probably a magnetar \\citep{Zhang_01,Zhang_02}. From the study of chromatic breaks seen in six well sampled afterglow light curves \\citet{Pan2006} concludes that if both, the optical and the X-ray afterglows, were to arise from the same outflow then the chromaticity of light curve breaks can rule out energy injection or the structure of the jet as the possible reasons of it.  One such GRB afterglow with puzzling features in optical and X-ray light curves is GRB 050401. GRB 050401 triggered \\emph{Swift}-BAT at 14:20:15 UT on 2005 April 01 \\citep{GCN3162}. The X-ray afterglow was detected by \\emph{Swift}-XRT \\citep{GCN3161} about 130 seconds after the trigger and the optical afterglow candidate was confirmed by ground based observations by \\citet{GCN3164}. The burst duration $\\rm T_{90}$ is estimated to be $\\sim$ 33 seconds \\citep{GCN3173}. Using the measured spectral redshift of the afterglow ($z = 2.9$) \\citep{GCN3176} and the fluence \\citep{GCN3173,050401_MDP,GCN3179} the isotropic equivalent energy released during the explosion turns out to be  $1.4 \\times 10^{54}$ for a flat universe with $\\rm \\Omega_{m} = 0.3, \\Omega_{\\Lambda} = 0.7 ~and~ H_{0} = 70 ~km~s^{-1} ~Mpc^{-1}$.\\\\  Multiband afterglow of GRB 050401 also presents some puzzling features which can be summarized as follows : \\begin{enumerate} \\item A break in the X-ray light curve after $\\sim$ 0.06 day with an unusual slope after the break \\citep{050401_MDP,050401_Watson}. \\item The X-ray afterglow does not show any spectral evolution across the break while the R band light curve does not show any break \\citep{050401_MDP,050401_Watson}. \\item A large extinction inferred from X-ray afterglow which is not consistent with the observed optical afterglow \\citep{050401_Watson}. \\end{enumerate}  The optical observations are presented in \\S~\\ref{sec:Observations}. We have done some preliminary analysis of the light curves which is discussed in \\S~\\ref{sec:lightcurve}. We have tried to explain the multi-band behaviour of the GRB afterglow using a double jet model which is described in \\S~\\ref{sec:model} along with the previous attempts by others using a different model. In the Discussion section (\\S~\\ref{sec:Discussion}) molecular clouds as a plausible explanation for the large extinction is presented (\\S~\\ref{sec:disc_extinction}). The only other GRB afterglow which has been explained using a similar double jet model is the GRB 030329 \\citep{030329_38,030329_05}. We compare the physical features of GRB 030329 and GRB 050401 in \\S~\\ref{sec:disc_compare}. Our conclusions are summarized in \\S~\\ref{sec:Summary}. ",
        "conclusions": "\\label{sec:Discussion}  \\subsection{A plausible explanation for the large extinction inferred from X-ray absorption} \\label{sec:disc_extinction}  It is now well established from the observations of GRB hosts that long GRBs preferentially occur in massive star forming regions e.g. \\citet{WB2006}. The massive star forming regions host large molecular clouds. Typical column densities of cold molecular clouds are $\\rm > 10^{22} ~cm^{-2}$, densities $100 - 10^{4} cm^{-3}$ and sizes $\\rm \\sim 20 ~pc$. Giant molecular clouds are even denser ($\\rm 10^{4} - 10^{7} cm^{-3}$) and larger ($\\rm \\sim 100 ~pc$) \\citep{Shore}. It is possible that one such cloud in the host galaxy of GRB 050401 happens to fall along our line of sight which can explain the large extinction inferred from the X-ray spectrum. We consider the possibility of radiation from the double jet of GRB 050401 being obscured by a molecular cloud so aligned that it covers the narrow jet of GRB 050401 completely while the wide jet is partially covered. By changing the fractional coverage of wide jet by the cloud we measured change in the value of reduced $\\chi^{2}$ of the fit. In effect, this amounts to adjusting the intrinsic luminosity of the wide jet upwards with increasing covering factor to match the observed optical flux. This results in the relative contribution of the wide jet to the X-ray afterglow to increase, affecting the fit quality. Keeping all other parameters fixed at their best fit values obtained for zero coverage, we find that a covering fraction of 60\\% can be accomodated within a range of $\\Delta \\chi^{2}/dof = 1$. Beyond this the reduced $\\chi^{2}$ rises sharply and reaches $\\Delta \\chi^{2}/dof > 15$ for a covering factor of $\\sim 90\\%$. For the observed column density of $1.7 \\times 10^{22}$ cm$^{-2}$ \\citep{050401_MDP}, and assuming typical densities ($100-1000$ cm$^{-3}$) of the molecular clouds, the size of the molecular cloud could be estimated to be around $5-55$ parsecs. It is therefore probable that one such molecular cloud partially obscures our view of GRB 050401. This situation is illustrated in Figure~\\ref{fig:doublejet}. \\begin{figure*} \\begin{center} \\includegraphics[width=6cm,angle=-90]{DoubleJet_side_4.ps} \\caption{A side view of the double jet (not to scale). The observer is on the axis of the jets and at a distance of 24 Gpc (which can be considered as at infinity for geometric purposes in this figure). The arrows indicate the afterglow light rays emanating from the jets. The intervening molecular cloud, of size larger than 5 to 55 parsecs, responsible for the observed large extinction is sitting at a distance of about 100-1000 parsecs from the GRB. The estimated diameters of the jets around 0.05 days turn out to be about $2 \\times 10^{-3} ~pc$ and $> 2 \\times 10^{-2}~pc$ respectively for the narrow and the wide jet. The large cloud covers a significant portion of the central narrow jet and partially covers the wide jet when seen from the observer's point of view. As a result, the optical radiation from the narrow jet is completely extincted. Most of the optical radiation from the wide jet does not suffer from this extinction.} \\label{fig:doublejet} \\end{center} \\end{figure*} At this point, we would like to point out two possible caveats in the double jet model proposed here :  The separation of the optical and X-ray emitting regions, as proposed in the present model, is motivated by the large discrepancy of  about 8 magnitudes between the amount of optical extinction inferred from soft X-ray absorption and that from observed optical-IR spectrum of the GRB050401 afterglow.  It should, however, be kept in mind that the \\citet{PS1995} relation used to predict $A_V$ from X-ray absorbing column $N_H$ is an empirical one, and cannot be considered fully reliable in all circumstances. For example, a metallicity higher than solar by a factor of ~10, or a dust-to-gas ratio lower by a similar factor, can reconcile the X-ray absorption with observed optical extinction.  Such explanations in this case cannot be ruled out, and have been already discussed by \\citet{050401_Watson}.  The second caveat is that the model presented here requires a rather special geometrical alignment - the two jets of the GRB should shine through the outer edge of a molecular cloud, much larger in size than the transverse extent of the jet working surface, in such a manner as to provide large extinction to the inner jet but much less to at least half the outer component. This requires that the outer edge of the cloud be dense, and have a strong density gradient to differentially affect the two jet components.  An elongated, cigar-shaped cloud with its axis nearly parallel to the line of sight, would also help such a scenario. We also note that the size of the cloud required, as estimated by us using an average density, is prone to large uncertainties if its shape is unusual or if large density gradients are present. \\subsection{GRB 050401 and GRB 030329 : A comparison} \\label{sec:disc_compare}  The only other GRB whose afterglow has been explained as being due to double jet is the GRB 030329 \\citep{030329_38, 030329_05}. optical and X-ray light curves of GRB 030329 afterglow showed a near simultaneous break at 0.55 day whereas the radio light curves had a break at about 10 days after the burst. \\citet{030329_38} have explained the two breaks as being due to lateral expansion of the two co-axial jets of different opening angles ($\\sim 5^{\\circ}$ and $\\sim 17^{\\circ}$).  In the case of GRB 050401, afterglow light curves do not show the presence of two different breaks. Instead, absence of a break at optical frequencies till late times ($\\sim 13$ days after the burst) leads us to infer the presence of a wider jet with opening angle larger than $29^{\\circ}$ while a steep break ($\\Delta \\alpha \\sim 0.8$) at 0.06 day after the burst in X-ray light curve can be explained as a jet break due to lateral expansion of a narrow jet of opening angle $1.15^{\\circ}$.  The wider jet of GRB 030329 was estimated to be marginally more energetic than the narrower jet \\citep{030329_38, 030329_05}. Similarly, in the case of GRB 050401, we find, that the wider jet is marginally more energetic than the narrower jet.  \\subsection{GRB 050401 and the Ghirlanda Relation :} \\label{sec:Ghirlanda_relation}  It has been found that the collimation corrected energies ($E_{\\gamma}$) of the GRBs are correlated with the peak energy of the GRB spectrum as measured in the frame of reference of the source ($E_{peak}^{src}$). This correlation is also called as the Ghirlanda relation \\citep{Ghirlanda2004}. Unfortunately, the $E_{peak}^{src}$ for GRB 050401 is not available as it falls outside the energy range of BAT. However, \\citet{050401_Sato} have used the Konus-Wind spectral data to find $E_{peak}^{obs}$. From their analysis \\citet{050401_Sato} finds that in order to satisfy the Ghirlanda relation the afterglow light curve of GRB 050401 should exhibit a jet break $\\sim 10^{4} ~s$ after the burst. This lower limit of the allowed range for jet break time is close to the break seen at 0.06 day in the X-ray light curve of GRB 050401, which we interpret as a jet break corresponding to the narrow jet in our model.  \\citet{050401_Sato} quantifies the Ghirlanda relation as $E_{peak}^{src} = A ~E_{\\gamma,52}^{0.706}$ where $E_{\\gamma,52}$ is the collimation corrected energy released in $\\gamma$ rays during the burst, in units of $10^{52}$ ergs. Using a sample of a large number of GRBs \\citet{050401_Sato} constrains the value of the proportionality constant $A$ : $1950 < A < 4380 $. Using the estimated value of $E_{\\gamma}^{iso} \\sim 10^{54}$ ergs and the $1.15^\\circ$ as the opening angle of the narrow jet in our double jet model, the $E_{\\gamma}$ turns out to be $2 \\times 10^{50}$ ergs. Using $E_{peak}^{src} = 447^{+75}_{-64} keV$ for GRB 050401 as reported by \\citet{050401_Sato} along with $E_{\\gamma} = 2 \\times 10^{50}$ ergs we estimate $A = 7076^{+2597}_{-1897}$. This value of $A$ is within $2~\\sigma$ of $A = 4380$, the higher limit on $A$ obtained considering the sample of GRBs satisfying the Ghirlanda relation. Having discussed this, we would also like to point out that the Ghirlanda relation has sometimes been critisized as being due to selection effects rather than being an intrinsic correlation \\citep{Butler2008}.    We have reported VRI band observations of GRB 050401 afterglow on 1st Apr. 2005. Also, we have modeled the afterglow of GRB 050401 as due to two physically distinct collimated outflows, using our own VRI band photometry along with the observations available in the literature, and compared with GRB 030329. Our main conclusions about GRB 050401 are as follows : \\begin{itemize} \\item [1.] We showed that the light curves of GRB 050401 afterglow can not be explained under the assumption of continuous energy injection. The flatter decay, which appealed for the continuous energy injection model, can instead be explained by low values of electron energy distribution index $p$. \\item [2.] The afterglow of GRB 050401 can be well fit by the double jet model with the interpretation that the break in the X-ray light curve at $\\sim$ 0.06 day after the burst is due to a narrow collimated jet expanding sideways. The obscured optical emission is attributed to a wider which did not undergo significant sideways expansion until at least $\\sim 13$ days after the burst. \\item [3.] Kinematically, we find that the wider jet is slightly more energetic, than the narrow jet. This result is similar to what was found in the double jet of GRB 030329. \\item [4.] Our interpretation of the break in the X-ray light curve at 0.06 days after the burst as a jet break is consistent with the Ghirlanda relation. \\end{itemize}"
    },
    "0806/0806.3666_arXiv.txt": {
        "abstract": "In this paper it has been shown that the neutrino bulk viscous stresses can give rise to the late time acceleration of the universe. It is found that a number of spatially flat FRW models with a negative deceleration parameter can be constructed using neutrino viscosity and one of them mimics a $\\Lambda$CDM model. This does not require any exotic dark energy component or any modification of gravity. ",
        "introduction": "That the present universe is undergoing an accelerated expansion has now been firmly established. The initial indications came from the supernovae data \\cite{1} and were soon confirmed by many high precision observations including the WMAP \\cite{2}. Theoretical physics has thus been thrust with the challenge of finding the agent, dubbed ``{\\em dark energy}\", which can drive this acceleration. Naturally a host of candidates appeared in the literature which can provide this antigravity effect. However, the dark energy should become dominant only during the later stages of matter era so that nucleosynthesis in the early universe and the large scale structure formation in the matter dominated regime could proceed unhindered and make the universe look the place where we live in now. \\par Amongst the various dark energy candidates, the cosmological constant $\\Lambda$ is certainly the most talked about one. It matches different observational requirements quite efficiently and has been known in cosmology for quite a long time for the various roles it could play. The insurmountable problem is of course that of the huge discrepancy between the required value of $\\Lambda$ and that predicted theoretically \\cite{3}. \\par The quintessence models, where an effective negative pressure generated by a scalar field potential drives the acceleration, work extremely well to fit into various observational constraints. For a very brief review, we refer to Martin's recent work \\cite{4}. But none of the potentials employed for the purpose can boast of any sound theoretical motivation. Non minimally coupled scalar field theories like Brans - Dicke theory can also be used where even a dark energy is not required \\cite{5}, but the value of the Brans - Dicke parameter $\\omega$ needs to be given a very small value contrary to the observational requirements. Scalar - tensor theories with a dark energy, particularly where an interaction between the dark energy and the geometrical scalar field like the Brans - Dicke scalar field is allowed so as to alleviate the coincidence problem, appear to do well \\cite{6}. But the nature of the interaction is hardly well-motivated. Also, a recent work shows that energy has to be pumped in to the dark matter from the dark energy sector as demanded by the second law of thermodynamics \\cite{7}. This is indeed counter-intuitive in view of the fact that the dark energy dominates over the dark matter only during the later stages of evolution. Chaplygin gas models \\cite{8}, modified gravity theories \\cite{9} and many other models are proposed. They all have their success stories as well as failures in some way or the other. In the absence of a clear verdict in favour of a particular dark energy candidate, all of them have to be discussed seriously. There are excellent reviews regarding different models and their relative merits \\cite{10}. \\par One important general feature is that all these models use either some kind of an exotic field or some modification of the firmly established general relativity - the effect of the modification being hardly required by other branches of well established physical theories, and the possibility of an actual physical detection of them appears to be quite a far-fetched one. \\par Recently a well known sector of matter, whose existence in abundance has been firmly established, namely the neutrino distribution, has been proposed as a candidate for the dark energy \\cite{11}. The motivation comes from particle physics, and for a brief but comprehensive review we refer to \\cite{12}. However, in these models, neutrinoes are normally the sector which `{\\em feels}' the existence of dark energy and cannot really solve the problem by itself, i.e, without a quintessence potential. For a review of the neutrino properties in a cosmological context, we refer to \\cite{dolgov}. \\par In the present work, we treat the neutrinoes completely classically and show that bulk viscous stresses in the distribution can indeed do the trick. \\par The advantage of the neutrinoes is that they are real objects, the method of detection being quite well conceived. The neutrinoes were decoupled from the background radiation quite early in the evolution when the temperature was as high as $3 \\times 10^{10} K$ \\cite{13}. Thus the interaction of the neutrinoes with other forms of matter can be ignored and hence the problem of the ``direction of the flow of energy\" \\cite{7} does not arise and the model becomes much more tractable. \\par Neutrino viscosity, both in the form of shear \\cite{14} or in the form of bulk viscosity \\cite{15} had been investigated quite a long time back for various purpose. There has been a renewed interest in the neutrino viscosity quite recently as well \\cite{16}. The problem of dissipative effects like viscosity or heat conduction had been that of a parabolic transport equation, which could allow the signals to travel with super-luminal speed resulting in a violation of causality. But the extended irreversible thermodynamics, which modifies the transport equation by including a relaxation time and a further divergence term to avoid this problem, is now quite well understood \\cite{17} and the modified transport equation had already been quite extensively used in cosmology \\cite{18}. \\par In what follows, we employ a two - component non-interacting matter sector, one is the normal cold dark matter and the other being a neutrino distribution. The latter is endowed with bulk viscosity which produces a negative stress. It is shown that a very simple accelerated model can be constructed from this. The model looks simple as the shear viscosity is neglected in order to be consistent with the isotropic nature of the universe. ",
        "conclusions": "As a conclusion, one can say that a sufficient bulk viscous stress in the neutrino distribution can potentially serve the purpose of a dark energy. This does not require any ill-motivated scalar potential or an otherwise unwarranted modification of general relativity. However, the model is far from being complete. The amount of the bulk viscous stress has to be sufficient to drive the acceleration, and the correct phenomenological connection between the quantities like $\\eta$, $\\tau$ and $\\rho$ has to be found out so that temperature of neutrinoes has a realistic profile. The neutrino viscosity can in fact give a wide range of accelerating models. One of them, the effective $\\Lambda$CDM model is discussed here. But the other example mentioned also works, and leaves a possibility of finding many others within the scope of it. It has to be searched which solution is favoured from the consideration of stability as well as that of observational bounds."
    },
    "0806/0806.2660_arXiv.txt": {
        "abstract": "Accretion disks in which angular momentum transport is dominated by the magnetorotational instability (MRI) can also possess additional, purely hydrodynamic, drivers of turbulence. Even when the hydrodynamic processes, on their own, generate negligible levels of transport, they may still affect the evolution of the disk via their influence on the MRI. Here, we study the interaction between the MRI and hydrodynamic turbulence using local MRI simulations that include hydrodynamic forcing. As expected, we find that hydrodynamic forcing is generally negligible if it yields a saturated kinetic energy density that is small compared to the value generated by the MRI. For stronger hydrodynamic forcing levels, we find that hydrodynamic turbulence modifies transport, with the effect varying depending upon the spatial scale of hydrodynamic driving. Large scale forcing boosts transport by an amount that is approximately linear in the forcing strength, and leaves the character of the MRI (for example the ratio between Maxwell and Reynolds stresses) unchanged, up to the point at which the forced turbulence is an order of magnitude stronger than that generated by the MRI. Low amplitude small scale forcing may modestly suppress the MRI. We conclude that the impact of hydrodynamic turbulence on the MRI is generically ignorable in cases, such as convection, where the additional turbulence arises due to the accretion energy liberated by the MRI itself. Hydrodynamic turbulence may affect (and either enhance or suppress) the MRI if it is both strong, and driven by independent mechanisms such as self-gravity, supernovae, or solid-gas interactions in multiphase protoplanetary disks. ",
        "introduction": "The magnetorotational (MRI) or Balbus-Hawley instability \\citep{vel59,chandra61,balbus91,balbus98} underpins the most important --- and possibly only --- source of outward angular momentum transport in a wide class of well-ionized accretion disks. The MRI destabilizes disk flows in which ${\\rm d}\\Omega^2 / {\\rm d}r < 0$, and leads to a state of sustained magnetohydrodynamic (MHD) turbulence that transports angular momentum outward \\citep{hawley95,brandenburg95,armitage98,hawley00,papaloizou03,hirose06}. The majority of the transport is mediated by Maxwell rather than Reynolds stresses. Following convention, the efficiency of angular momentum transport within disks is measured via an equivalent Shakura-Sunyaev (1973) $\\alpha$ parameter, which can be expressed in terms of the fluctuating velocity and magnetic fields as, \\begin{equation} \\alpha P_0 = \\rho v_r \\delta v_\\phi + \\frac{B_r B_\\phi}{\\mu_0}, \\label{alpha} \\end{equation} where $P_0$ is the thermal pressure. The first term on the right-hand side of this equation represents the Reynolds (or fluid) stress and the second term represents the Maxwell (or magnetic) stress. There is no strict equivalence between the evolution of MHD turbulent disks and models that assume an $\\alpha$ shear viscosity \\citep{balbus98,pessah07}, but for our purposes $\\alpha$ is a convenient measure of the efficiency of angular momentum transport.  The fact that the MRI dominates the transport of angular momentum within accretion flows does not, of course, imply that other sources of turbulence do not exist within disks. The most striking example occurs in galactic disks, in which turbulence can be driven by thermal instabilities and supernova explosions in regions that are unstable to the MRI \\citep{piontek05}. However, even in ``normal\" accretion disks around stars or compact objects there are numerous possibilities, including self-gravity in sufficiently massive disks \\citep{toomre64}, convection \\citep{stone96}, and Kelvin-Helmholtz instabilities excited by the interaction between gaseous and solid components of protoplanetary disks \\citep{cuzzi93}. Additional fluid motions can also be driven within disks due to the gravitational influence of embedded planets \\citep{bate02} or binary companions \\citep{spruit87}, though the wave-like fluid motions induced by these sources are different from those initiated by turbulence. Some of these processes may drive fluid motions within the disk whose kinetic energy density is comparable to that resulting from the MRI. In particular, a self-gravitating disk can remain stable against fragmentation while generating an effective $\\alpha \\simeq 10^{-1}$ \\citep{gammie01,rice05} --- as large or larger than that produced by the MRI. Since the fluid motions resulting from hydrodynamic processes are uncorrelated with the MRI, it is reasonable to suspect that there could be non-trivial interactions between the MRI and other turbulent processes in disks within which both are operating simultaneously. Indeed, \\cite{fromang04} found that in a self-gravitating, MRI-turbulent disk, the strength of the angular momentum transport from the self-gravity was both weaker, and had a different time-dependence, when compared to a disk in which self-gravity alone was at work. This result is striking, since the disk simulated by \\cite{fromang04} was dominated by low-order (azimuthal wavenumber $m=2$) self-gravitating structure whose scale was much larger than the most unstable MRI scales.  In detail, different drivers of hydrodynamic turbulence may influence the MRI in a unique manner, requiring a case-by-case study (self-gravity, for example, is a special case since it {\\em does} yield outward transport of angular momentum, whereas most other hydrodynamic mechanisms yield negligible or even inward transport). We do not address such subtleties here, but rather study how generic driven hydrodynamic turbulence of specified strength interacts with the MRI. One's expectation is obviously that hydrodynamic turbulence that is weak (say, in terms of the saturated value of the kinetic energy density) compared to that driven by the MRI ought to leave the MRI unscathed, with significant interaction developing when the two sources of turbulence are of comparable strength. Very strong turbulence can amplify magnetic fields in the disk {\\em independent} of the MRI, though this may not necessarily be accompanied by angular momentum transport. Our goal in this paper is to test such order of magnitude intuition.  The focus of this paper is on small-scales, which can be captured most effectively using local shearing-box simulations. Such simulations have significant limitations that must be borne in mind. In particular, the strength of angular momentum transport in zero-net flux simulations with purely numerical viscosity and magnetic diffusivity has a marked dependence on numerical resolution \\citep{gardiner05,fromang07a}. Moreover, when physical values for the transport coefficients {\\em are} used (as in this paper) the strength of turbulence depends upon the magnetic Prandtl number $Pm \\equiv \\nu / \\eta$ as well as on the Reynolds number \\citep{lesur07,fromang07b}. What this means for real disks --- which when highly ionized have values of the diffusivity $\\eta$ and viscosity $\\nu$ that are much smaller than can currently be simulated --- is unclear, but an obvious implication for numerical experiments is that the absolute value of $\\alpha$ derived from shearing-box simulations must be viewed with caution. For the time being, constraints derived from modeling of observed systems may be more reliable \\citep{kpr07}. For this reason, our focus here is on how the strength of angular momentum transport varies in the presence of additional hydrodynamic turbulence, rather than its absolute value.  The plan of this paper is as follows. In \\S2 we describe the set-up for our runs, which make use of the {\\sc PENCIL} MHD code previously employed for both disk and turbulence calculations \\citep{haugen04}. In \\S3 we present results, which concentrate on the influence of different levels of hydrodynamic turbulence on the saturation level, structure, and angular momentum transport efficiency of the MRI. These results are summarized and discussed in \\S4. ",
        "conclusions": "In this paper we have used local shearing-box simulations of accretion disks to study the interaction between the magnetorotational instability and hydrodynamic turbulence in disks where both coexist. We have studied both large scale forcing, in which energy is injected at scales $\\lambda \\sim H$, the disk scale height, and to a more limited extent small scale forcing where $\\lambda \\ll H$. We find that the effect of the additional energy input from the hydrodynamic turbulence on the MRI varies depending upon both the strength of the forcing and, to some extent, on its spatial structure.  For large scale forcing the results are straightforward. When hydrodynamic forcing is ``weak\" -- in the sense that the hydrodynamic forcing, on its own, yields a saturated kinetic energy density in the disk that is less than or equal to that generated by the MRI alone -- the MRI is essentially unaffected at the level of precision accessible to our simulations. This result is no surprise. The MRI is a robust instability, which is present within differentially rotating flows containing almost arbitrary magnetic field geometries \\citep{balbus98}. Low amplitude random perturbations do not affect it. When stronger forcing is imposed, we find that both the saturation value of the magnetic field and the strength of angular momentum transport can be substantially boosted. The flow retains many of the characteristics of the MRI --- such as a similar ratio between magnetic field strength and $\\alpha$, and a similar ratio of Maxwell to Reynolds stresses, even in a regime when the hydrodynamic forcing (which, by itself, yields no transport at all) is formally dominant. The application of heuristic dynamo arguments to the MRI is suspect \\citep{balbus98}, but we tentatively attribute the numerical behavior to the more efficient regeneration of vertical field in the presence of hydrodynamic turbulence.  The small scale forcing results are more nuanced, and given the limited number of simulations we have performed should be regarded as preliminary. In this regime we find that only the strongest level of forcing boosts the strength of angular momentum transport, whereas lower forcing levels actually suppress transport. We attribute this different behavior to the fact that there are two important considerations that affect the saturation amplitude of the MRI. One is the strength of the vertical magnetic field on relatively large scales (typically a fraction of $H$), similar to the most unstable linear MRI wavelengths. Enhancement of the vertical field on this scale -- which is readily accomplished with large scale forcing but which requires an inverse cascade in the small scale case -- boosts the strength of the MRI. The second is the dissipation scale. Turbulent driving that increases the amplitude of kinetic energy at or near this scale may enhance magnetic field dissipation, ultimately reducing the saturation level of magnetic fields in the disk as a whole. This is the inverse of the physical process invoked to explain the dependence of $\\alpha$ on the magnetic Prandtl number \\citep{lesur07, fromang07b}. Small scale forcing may also directly destroy the correlations between $B_r$ and $B_\\phi$ that result in Maxwell stress.  Our results allow us to infer which additional physical effects are likely to be able to affect the MRI in accretion disks. We first observe that if the additional turbulence is ``powered\" by the MRI itself (i.e. if the turbulence derives energy from the gravitational potential well as a result of MRI-driven angular momentum transport), then generically it is unlikely to be as powerful as the MRI. Such turbulence will have at most a small effect on the magnitude and character of angular momentum transport. As an explicit example, we would not expect an MRI-active disk that was additionally unstable to convection to differ much from one in which the vertical structure was stable against convective motions.  We can also consider sources of turbulence that are independent of the MRI, in the sense that they would exist even in a (hypothetical) disk that was absent magnetic fields entirely. Physical examples include self-gravity, thermal instability, and turbulence stirred up by the interaction between the gaseous and solid components of protoplanetary disks. There is no reason why these sources of turbulence should not generate fluid motions of greater amplitude than those produced by the MRI (this is likely to be the case in self-gravitating disks near the fragmentation threshold, and locally in some regions of protoplanetary disks). In this regime, our results suggest that the MRI will have a non-trivial interaction with the hydrodynamic turbulence. If the forcing occurs at scales comparable to $H$ or larger, we find that the interaction will likely boost the strength of angular momentum transport. Even for quite violent forcing -- up to an order of magnitude in excess of that required to produce parity with the MRI -- we find that it is most accurate to think of the coupled system as one with boosted MHD turbulence, rather than as a hydrodynamic system passively advecting magnetic fields. Conversely, small scale forcing, unless it is very strong, may actually suppress the saturation level of magnetic fields in the disk and their associated angular momentum transport."
    },
    "0806/0806.0179_arXiv.txt": {
        "abstract": "The relic cosmic background neutrinos accompanying the cosmic microwave background (CMB) photons may hide a universal lepton asymmetry orders of magnitude larger than the universal baryon asymmetry.  At present, the only direct way to probe such an asymmetry is through its effect on the abundances of the light elements produced during primordial nucleosynthesis. The relic light element abundances also depend on the baryon asymmetry, parameterized by the baryon density parameter ($\\eta_{\\rm B} \\equiv n_{\\rm B}/n_{\\gamma}$), and on the early-universe expansion rate, parameterized by the expansion rate factor ($S \\equiv H'/H$) or, equivalently by the effective number of neutrinos (N$_{\\nu} \\equiv 3 + 43(S^2 - 1)/7$). We use data from the CMB (and Large Scale Structure: LSS) along with the observationally-inferred relic abundances of deuterium and helium-4 to provide new bounds on the universal lepton asymmetry, finding for $\\eta_{\\rm L}$, the analog of $\\eta_{\\rm B}$, 0.072$\\pm$0.053 if it is assumed that \\Nnu = 3 and, 0.115$\\pm$0.095 along with \\Nnu = 3.3$^{+0.7}_{-0.6}$, if \\Nnu is free to vary. ",
        "introduction": "The standard models of particle physics and cosmology assume that in the early, radiation-dominated universe only the known, massless or light ($mc^2 \\ll kT$) particles, including three flavors of light, active neutrinos (N$_{\\nu} = 3$), contribute to energy density driving the universal expansion ($H^2 = 8\\pi G\\rho/3$; $\\rho = \\rho_{\\rm R}$). It is also generally assumed that any universal lepton asymmetry is very small, of order the baryon asymmetry.\\footnote{Charge conservation ensures that the very small proton-antiproton asymmetry, of order $\\eta_{\\rm B}$, is balanced by a correspondingly small asymmetry between electrons and positrons but, there are no such constraints on the size of any asymmetry between neutrinos and antineutrinos.} In analogy with the measure of the baryon asymmetry provided by $\\eta_{\\rm B}$, an asymmetry between neutrinos and antineutrinos of flavor $\\alpha$ ($\\alpha$ = e, $\\mu$, $\\tau$) can be described in terms of the neutrino chemical potential $\\mu_{\\alpha}$ or, in terms of the dimensionless degeneracy parameter $\\xi_{\\alpha} \\equiv \\mu_{\\alpha}/kT$ by, \\be \\eta_{L} \\equiv \\Sigma_{\\alpha} {n_{\\nu_{\\alpha}} - n_{\\bar{\\nu}_{\\alpha}} \\over n_{\\gamma}} = {\\pi^{3} \\over 12\\zeta(3)}\\Sigma_{\\alpha} \\left[\\left({\\xi_{\\alpha} \\over \\pi}\\right) + \\left({\\xi_{\\alpha} \\over \\pi}\\right)^{3}\\right]. \\ee  For Big Bang Nucleosynthesis (BBN) the electron neutrinos play a key role through their charged-current weak interactions, which help to regulate the neutron-to-proton ratio ($p + e^{-} \\leftrightarrow n + \\nu_{e}$, $n + e^{+} \\leftrightarrow p + \\bar\\nu_{e}$, $n \\leftrightarrow p + e^{-} + \\bar\\nu_{e}$).  Since the BBN-predicted \\4he abundance is, to a very good approximation, determined by the neutron-to proton ratio at BBN, changes from the standard model value of this ratio will be reflected in its primordial abundance and, to a lesser extent, in the abundances of the other light elements produced during BBN.  These relic abundances therefore provide probes of a universal lepton asymmetry. For further discussion and previous analyses, see \\eg \\citealt{wfh, reeves, yb, bg, by,sch79, dr, scherrer, freese, terasawa, boes85, kang, kohri, esposito, ichi02, barger, kneller04, cuoco, serpico, steigman07, popa07, popa08}.  There is another way in which a significant neutrino degeneracy may influence the early evolution of the Universe.  In the standard models of particle physics and cosmology, the energy density at, or just prior to, BBN is contributed by the cosmic background radiation photons, \\epm pairs, and three flavors of extremely relativistic neutrinos.  In the standard cosmology, the neutrinos constitute ~40\\% of the total energy density.  Any modification of the early-Universe energy density (or expansion rate; see \\eg \\citealt{vs08}) can be parameterized in terms of the ``effective\" number of neutrinos by $\\rho \\rightarrow \\rho' \\equiv \\rho + \\Delta$N$_{\\nu}\\rho_{\\nu}$.  For the standard models, \\Nnu = 3 at BBN, while in the post \\epm annihilation epoch probed by the CMB, \\Nnu = 3.046 \\citep{mangano05}.  A secondary effect of neutrino degeneracy is to increase the energy density in relativistic neutrinos predicted by the standard model, so that \\Nnu $\\rightarrow$ N$_{\\nu} + \\Sigma_{\\alpha} \\Delta$N$_{\\nu}(\\xi_{\\alpha})$, where \\be \\Delta{\\rm N}_{\\nu}(\\xi_{\\alpha}) \\equiv {30 \\over 7}\\left({\\xi_{\\alpha} \\over \\pi}\\right)^{2} + {15 \\over 7}\\left({\\xi_{\\alpha} \\over \\pi}\\right)^{4}. \\ee  In general, if the three active neutrino flavors mix only with each other, neutrino oscillations ensure that their degeneracies equilibrate prior to BBN \\citep{LS01,dolgov02,wong02,abb02}.  In the following we will assume this is the case and use \\xie = $\\xi_{\\mu} = \\xi_{\\tau}$.  As a result, \\be \\eta_{\\rm L} = {\\pi^{3} \\over 4\\zeta(3)}\\left[\\left({\\xi_{e} \\over \\pi}\\right) + \\left({\\xi_{e} \\over \\pi}\\right)^{3}\\right], \\ee and \\be \\Sigma_{\\alpha} \\Delta{\\rm N}_{\\nu}\\left(\\xi_{\\alpha}\\right) = {90 \\over 7}\\left({\\xi_{e} \\over \\pi}\\right)^{2} + {45 \\over 7}\\left({\\xi_{e} \\over \\pi}\\right)^{4}. \\ee Notice that, if $\\xi_{e} = \\xi_{\\mu} = \\xi_{\\tau}$, then for $|\\xi_{e}| \\la 0.2$, $\\Sigma_{\\alpha} \\Delta$N$_{\\nu}(\\xi_{\\alpha}) \\la 0.05$. Later, in \\S6, we will relax this assumption so that $\\xi_{e} \\neq \\xi_{\\mu}$ and/or $\\xi_{\\tau}$)].  The only effect of non-zero values of $\\xi_{\\mu}$ and/or $\\xi_{\\tau}$ is to change the effective value of N$_{\\nu}$, while non-zero values of \\xie also contribute to \\Nnu and, more importantly, they modify the neutron-to-proton ratio at BBN. ",
        "conclusions": "\\begin{figure} \\centerline{\\epsfxsize=5truein\\epsffile{figure3.eps}} \\caption{ The left hand panel shows the probability distribution for \\xie, inferred from BBN and the adopted relic abundances of D and \\4he, for the case where \\Nnu = 3.  In the right hand panel, the dashed curve is the BBN constraint convolved with the CMB/LSS constraint on \\eten alone, and the solid curve is the BBN constraint convolved with the CMB/LSS constraint on \\eten and \\Nnu. } \\label{fig:xipdf} \\end{figure}  \\begin{figure} \\centerline{\\epsfxsize=2.5truein\\epsffile{figure4.eps}} \\caption{ The dashed curve shows the probability distribution for \\Nnu for \\xie = 0 inferred from BBN and the relic abundances of D and \\4he convolved with the CMB/LSS constraints on \\eten and \\Nnu. The solid curve shows the same for \\xie $\\neq$ 0. } \\label{fig:nnupdf} \\end{figure}  The left hand panel of Figure \\ref{fig:xipdf} shows the probability distribution of \\xie for the standard expansion rate S = 0 (\\Nnu = 3), derived after marginalizing over \\eten, from BBN and the adopted primordial abundances of D and \\4he.  The right hand panel of Figure \\ref{fig:xipdf} shows the probability distribution of \\xie for the more general case where a non-standard expansion rate S $\\neq$ 0 (\\Nnu $\\neq$ 3) is allowed.  The constraints are based on combining BBN and the adopted primordial abundances of D and \\4he with the independent constraints from the CMB/LSS.  Additional information from the CMB/LSS constraints on \\eten are used before marginalizing over \\eten and \\Nnu to produce the dotted curve, while additional information from the CMB/LSS constraints on both \\eten and \\Nnu are used before marginalizing over \\eten and \\Nnu to produce the solid curve.  The dashed curve in Figure \\ref{fig:nnupdf} shows the probability distribution of \\Nnu for \\xie = 0, derived from BBN and the adopted primordial abundances of D and \\4he with the independent constraints from the CMB/LSS after marginalizing over \\eten. The solid curve in Figure \\ref{fig:nnupdf} shows the probability distribution of \\Nnu for the more general case where a non-zero neutrino degeneracy \\xie $\\neq$ 0 is allowed. The constraints are based on combining BBN and the adopted primordial abundances of D and \\4he with the independent constraints from the CMB/LSS.  Additional information from the CMB/LSS constraints on \\eten and \\Nnu are used before marginalizing over \\eten and \\xie. Both of these constraints on \\Nnu are consistent with each other and with the standard model prediction of \\Nnu $=$ 3.  Our results can be used to constrain any deviation in the universal expansion rate from its standard model value due to the increase in radiation energy density from neutrino degeneracy.  Our constraint on the neutrino degeneracy leads to $\\Sigma_{\\alpha} \\Delta$N$_{\\nu}(\\xi_{\\alpha}) \\leq 0.03$ at 95\\% confidence.  In addition, the constraint on the neutrino degeneracy yields a corresponding constraint on any lepton asymmetry, $\\eta_{\\rm L}$: $\\eta_{\\rm L} = 0.072\\pm0.053$ for \\Nnu = 3 and $\\eta_{\\rm L} = 0.115\\pm0.095$ when \\Nnu $\\neq$ 3.  Using the constraints on $\\eta_{10}$, N$_{\\nu}$, and \\xie, the BBN-predicted primordial abundance of \\7li may be inferred.  For \\Nnu = 3, [Li] = $2.63^{+0.04}_{-0.05}$ and for \\Nnu $\\neq 3$, [Li] = $2.62^{+0.05}_{-0.06}$. Both of these are considerably higher than the value ([Li]$_{\\rm P} = 2.1\\pm 0.1$) determined from observations of metal-poor halo stars (\\citet{ryan}, \\citet{asplund06}) without any correction for depletion, destruction, or gravitational settling.  If, however, the correction proposed by \\cite{korn06} is applied, the predicted and observed \\7li abundances may, perhaps, be reconciled.  It remains an open question whether this lithium problem is best resolved by a better understanding of stellar physics or, if it is providing a hint of new physics beyond the standard model.  In our analysis we have assumed that $\\xi_{e} = \\xi_{\\mu} = \\xi_{\\tau}$. Suppose instead that $\\xi_{\\mu} = \\xi_{\\tau} \\neq \\xi_{e}$ or, that $\\xi_{e} = \\xi_{\\mu} \\neq \\xi_{\\tau}$ \\citep{dolgov04}. Since our constraints on \\xie come from BBN, and they constrain \\xie to be sufficiently small that the allowed degeneracy has minimal effect on the universal expansion rate ($\\Delta$N$_{\\nu}(\\xi_{e}) \\la 0.01$), the only way to probe non-zero values of $\\xi_{\\alpha} \\equiv \\xi_{\\mu} \\equiv \\xi_{\\tau}$ or $\\xi_{\\alpha} = \\xi_{\\tau}$, is through their effect on the expansion rate ($S$ or, $\\Delta$N$_{\\nu}$).  In these cases, it is possible that $\\xi_{\\alpha}$ may be $\\gg \\xi_{e}$.  For \\Nnu = 3.3$^{+0.7}_{-0.6}$, $\\Sigma_{\\alpha} \\Delta$N$_{\\nu}(\\xi_{\\alpha}) \\la 1.7$ at $\\sim 2\\sigma$.  If it is assumed that $\\xi_{\\alpha} = \\xi_{\\mu} = \\xi_{\\tau}$, then $|\\xi_{\\alpha}| \\la 2.34$ and $|\\eta_{\\rm L}| \\la 5.0$. If, instead, it is assumed that $\\xi_{\\mu} = \\xi_{e} \\ll \\xi_{\\tau}$ (or, vice-versa for $\\xi_{\\mu}$), then $|\\xi_{\\tau}| \\la 4.12$ and $|\\eta_{\\rm L}| \\la 7.6$.  Of course, our results are sensitive to the relic abundances we have adopted.  The simple but accurate fitting formulae \\citep{kneller04} we have used make it easy to reevaluate our constraints for any adopted abundances.  The constraint on \\xie is sensitive to the \\4he abundance and is relatively insensitive to small changes in the D abundance.  For example, we repeated our analysis for a different primordial \\4he mass fraction, ${\\rm Y}_{\\rm P} = 0.247\\pm0.004$. This alternate abundance, in combination with the D abundance used in this paper, and the independent constraints on \\eten and \\Nnu from the CMB/LSS from \\cite{vs08} yields \\xie = 0.023$\\pm$0.041.  Our results here are similar to, but considerably more restrictive than those of \\citet{barger} and of \\citet{serpico}, due to the improved constraints on \\Nnu and \\eten from the WMAP 5-year and other CMB and LSS data.  The analysis described here seems to be related to that in recent papers by \\citet{popa07,popa08}.  However, we fail to understand how they derive their constraints and why they find so much tighter bounds on \\Nnu and so much weaker bounds on $\\xi_{e}$.  Except from its contribution to the radiation energy density and the early Universe expansion rate, a lepton asymmetry in the neutrino sector is invisible to the CMB.  Future CMB experiments will improve the constraint on N$_{\\nu}$ by measuring the neutrino anisotropic stress more accurately. According to \\cite{bash04}, PLANCK should determine N$_{\\nu}$ to an accuracy of $\\sigma(N_{\\nu}) \\sim 0.24$ and CMBPOL, a satellite based polarization experiment, might improve it further to $\\sigma(N_{\\nu}) \\sim 0.09$, independent of the BBN constraints.  Although still too large to provide a measure of the neutrino degeneracy, the tighter constraint on \\Nnu can be used to further narrow the allowed ranges of \\xie and \\Nnu shown in Figure \\ref{fig:nnuxi}.  \\begin{center} ACKNOWLEDGMENTS \\end{center} This research is supported at The Ohio State University by a grant (DE-FG02-91ER40690) from the US Department of Energy.  We thank J. Beacom and G. Gelmini for useful discussions.   \\singlespace \\small"
    },
    "0806/0806.3270_arXiv.txt": {
        "abstract": "{Ribas and collaborators have recently proposed that an additional, $\\sim$\\,5\\,$M_{\\oplus}$ planet orbits the transiting planet host star GJ\\,436. Long-term dynamical interactions between the two planets leading to eccentricity excitation might provide an explanation for the transiting planet's unexpectedly large orbital eccentricity. In this paper we examine whether the existence of such a second planet is supported by the available observational data when the short-term interactions that would result from its presence are accounted for. We find that the model for the system suggested by Ribas and collaborators lead to predictions that are strongly inconsistent with the measured host star radial velocities, transiting planet primary and secondary eclipse times, and transiting planet orbital inclinations. A search for an alternative two planet model that is consistent with the data yields a number of plausible solutions, although no single one stands out as particularly unique by giving a significantly better fit to the data than the nominal single planet model. We conclude from this study that Ribas and collaborator's general hypothesis of an additional short-period planet in the GJ\\,436 system is still plausible, but that there is not sufficient evidence to support their claim of a planet detection.} ",
        "introduction": "The GJ\\,436 system is unique among the nearly 250  extrasolar planetary systems identified so far\\footnote{A regularly updated list of reported exoplanets can be found at http://exoplanet.eu.}. It contains a ``Hot Neptune'' planet (planet ``b'') that was originally discovered with high precision Doppler spectroscopy by \\citet{butler04} and that was later found to transit by \\citet{gillon07b}. This planet is the only known member of its class that transits its host star. Therefore, it is an interesting target for the particular investigations that are feasible for transiting exoplanets \\citep[for a recent review of the observational techniques applicable to transiting exoplanets see][]{charbonneau07}. Since the discovery of the planet's transiting nature, follow-up studies have been carried out with the \\textit{Spitzer Space Telescope} (hereafter referred to as \\textit{Spitzer} for brevity) by \\citet{gillon07a}, \\citet{deming07}, \\citet{demory07}, and \\citet{southworth08}; and the \\textit{Hubble Space Telescope} (\\textit{HST}) by \\citet{bean08}.  The GJ\\,436 system is also unusual because planet b has a significantly non-circular orbit despite its proximity to its host star. The planet has an orbital eccentricity $e$\\,=\\,0.15\\,$\\pm$\\,0.01 as determined from analyzing the radial velocities of the host star with the constraint provided by the observed time of the planet's secondary eclipse that was observed with \\textit{Spitzer} \\citep{deming07,demory07}. An elliptical orbit for such a short-period planet ($P$\\,=\\,2.6\\,d) is potentially at odds with the predictions of tidal theory, which suggests that the planet's orbit should become circularized on a timescale of $\\lesssim$\\,$10^{8}$\\,yr \\citep{maness07, deming07}.  \\citet{maness07} found a significant linear trend in the radial velocities measured for GJ\\,436 over 6 years superimposed on the signal from planet b. This discovery was interpreted to mean that GJ\\,436 likely has an additional, but not necessarily planetary-mass, companion in a long-period orbit. \\citet{maness07} investigated whether a long-period planet consistent with the radial velocity trend could be the perturber leading to excitation of planet b's orbital eccentricity. They found that it was possible, but far from certain owing to the unconstrained nature of the object causing the slow acceleration of GJ\\,436. For example, \\citet{maness07} suggested that a roughly Saturn mass planet in a 25\\,yr orbit with $e$\\,$\\sim$\\,0.2 would be consistent with the radial velocities and provide the necessary regular perturbations.  Recently, \\citet[][hereafter RFB]{ribas08} have proposed another explanation for GJ\\,436b's high eccentricity. They suggest that an additional, short-period planet in the system would provide the necessary regular dynamical impulse to the transiting planet so that it would maintain its high orbital eccentricity in the face of tidal circularization over long timescales \\citep[although this result has recently been questioned by][]{mardling08}. Such a planet also met their requirement to cause the transiting planet's orbital inclination to change by 0.1\\degr\\,yr$^{-1}$. They saw this change in inclination as the reason why \\citet{butler04} didn't discover that GJ\\,436b transited despite ostensibly having achieved the necessary photometric precision and sampling in their search. Their hypothesis is that the planet simply wasn't transiting at that epoch, while it was when observed at a later epoch by \\citet{gillon07b} and the subsequent investigators mentioned above.  With the possible existence of another short-period planet in mind, RFB studied the radial velocities of GJ\\,436 provided by \\citet{maness07}. They identified a low-significance peak (20\\% false alarm probability) in a periodogram of the single planet residuals and used that as the starting point for a two planet fit. Assuming Keplerian orbits, they were able to obtain a two planet model that fit the radial velocities significantly better than a single planet model.  The second planet in the RFB model has an orbital period, $P$\\,=\\,5.1859\\,d, which puts it close to a 2:1 orbital resonance with the transiting planet, and minimum mass, $M\\,sin\\,i$\\,=\\,4.7\\,$M_{\\oplus}$. RFB proposed that this second planet exists based on the success of their model for providing a source for the transiting planet's eccentricity, a reason for the non-detection of transits by \\citet{butler04}, and a fit to the radial velocities. If confirmed this would be the lowest mass planet yet found around a nearby, main sequence star. Therefore, the claim deserves further scrutiny.  One particular aspect of the RFB study that merits further investigation is the possible sensitivity of observables to gravitational interactions between the two planets in their model. These two planets would be in close, moderately eccentric, and possibly non-coplanar orbits and so their mutual perturbations might be significant on short timescales in addition to the long timescales that RFB only considered. If they are, then RFB's claimed detection might be spurious because their model did not account for them. The critical issue is that Keplerian orbits, which RFB used for their modeling of the radial velocities, are only strictly valid for the two-body problem (i.e. a single planet orbiting a single star and also in the absence of significant General Relativity effects). Such orbits are only a sufficient approximation for modeling the data for multi-planet systems when the planet-planet interaction timescales are much longer than the length of the observations. For systems where short-term interactions are occurring, or are even possible, a model based on direct integrations of the equations of motion (i.e. Newtonian orbits) should be calculated to check the validity of the Keplerian approximation. If the observables of interest would be significantly different when accounting for the interactions, then the Newtonian orbit model must be used.  In this paper we assess the consistency of RFB's model of the GJ\\,436 planetary system with the observed host star radial velocities, transiting planet primary (transit) and secondary eclipse times, and transiting planet orbital inclinations when using Newtonian rather than Keplerian orbits. Pioneering work by \\citet{laughlin01} and \\citet{rivera01} have demonstrated that radial velocities with precisions on the order of a few m\\,s$^{-1}$ are sensitive to short-term dynamical interactions for certain exoplanet systems. \\citet{agol05} and \\citet{holman05} have shown that transit timings measured with precisions of a few seconds up to a few minutes are quite sensitive to additional planets with masses down to even the terrestrial level. Transit timings for a planet near to a low-order resonance, as RFB propose for GJ\\,436b, are particularly sensitive to very low-mass planets \\citep{steffen05}. ",
        "conclusions": "We have shown that a Super-Earth planet like that one proposed by RFB could still exist in the GJ\\,436 system, although their specific claim of a detection is erroneous. Recently, \\citet{alonso08} carried out another investigation into the plausibility of RFB's proposed planet using a novel constraint on the allowable inclination change of the transiting planet and reached a conclusion similar to ours. Ultimately, more observational data are needed to further constrain the architecture of the GJ\\,436 system and its evolutionary history. RFB's proposed planet would have been the lowest-mass one yet discovered around a nearby star so the results of this work reemphasize the importance of considering planet-planet interactions when interpreting observations of multi-planet systems. Keplerian orbits are still a useful approximation for modeling many such systems, but their appropriateness for a particular case should always be tested.\\\\[-15pt]"
    },
    "0806/0806.0843.txt": {
        "abstract": "~~}We report 20 and 6 cm VLA deep observations of the CDF-S including the Extended CDF-S. We discuss the radio properties of 266 cataloged radio sources, of which 198 are above a 20 cm completeness level reaching down to $43~\\mu$Jy at the center of the field. Survey observations made at 6 cm over a more limited region covers the original CDF-S to a comparable level of sensitivity as the 20 cm observations.  Of 266 cataloged radio sources, 52 have X-ray counterparts in the CDF-S and a further 37 in the E-CDF-S area not covered by the 1 Megasecond exposure. Using a wide range of material, we have found optical or infrared counterparts for 254 radio sources, of which 186 have either spectroscopic or photometric redshifts (Paper II). Three radio sources have no apparent counterpart at any other wavelength. Measurements of the 20 cm radio flux density at the position of each CDF-S X-ray source detected a further 30 radio sources above a conservative 3-sigma detection limit.   X-ray and sub-mm observations have been traditionally used as a measure of AGN and star formation activity, respectively.  These new observations probe the faint end of both the star formation and radio galaxy/AGN population, as well as the connection between the formation and evolution of stars and SMBHs. Both of the corresponding gravitational and nuclear fusion driven energy sources can lead to radio synchrotron emission.  AGN and radio galaxies dominate at high flux densities. Although emission from star formation becomes more prominent at the microjansky levels reached by deep radio surveys, even for the weakest sources, we still find an apparent significant contribution from low luminosity AGN as well as from star formation.  % ",
        "introduction": "}  This is the first of a series of papers based on Very Large Array (VLA) 1.4 GHz (21 cm) and 5 GHz (6 cm) observations of the Chandra Deep Field South (CDF-S) which also includes the Hubble Ultra Deep Field (UDF) and the Extended Chandra Deep Field South (E-CDF-S).  Paper II \\citep{M08a} presents the optical and near IR counterparts to the observed radio sources and Paper III (P. Tozzi et al, in preparation) their X-ray spectral properties.  Paper IV (P. Padovani et al, in preparation) discusses the source populations. Other papers will deal with a 1.4 GHz survey for low surface brightness sources in the CDF-S, a deeper 5 GHz survey, and a 1.4 GHz survey of the E-CDF-S \\cite{M08}.  Among the most fundamental issues in astrophysics is when and how stars, galaxies, and black holes formed and how they evolved with cosmic time.  Star and galaxy formation is a complex process in which mergers, shocks and accretion all appear to play important roles. Active galactic nuclei (AGN) which occur as a result of accretion onto a massive black hole, is also an important ingredient in galaxy evolution. Since these processes have complicated signatures throughout the electromagnetic spectrum, multiwavelength observations are needed to determine their evolution, dynamics and properties as a function of cosmic time.  Deep X-ray studies  \\citep[e.g.,][]{2005ARA&A..43..827B,2002AJ....124.2351B, 2004AJ....128.2048B, 2003A&A...399...39R} as well as the deep radio surveys indicate an increasing contribution to the X-ray emission at fainter flux densities from star forming activity. Various lines of evidence suggests the simultaneous presence of AGN and star formation and a causal relation between super massive black holes (SMBHs) and star formation. Previous deep radio surveys, \\citep[e.g.,][]{mux05,2000A&A...361L..41G} also appear to detect radio emission caused by both AGN and star formation in the same galaxy.  From investigations of the relationship between black hole mass and bulge luminosity \\citep{1995ARA&A..33..581K, 1998AJ....115.2285M}, or between black hole mass and bulge velocity dispersion \\citep{2000ApJ...539L...9F}, it appears that AGN processes  occur in star-forming galaxies with both processes being driven by, or connected by, an unknown phenomenon \\citep{2004ApJ...600..580G, 1999MNRAS.308L..39F}. Understanding the interaction of stars, galaxies and AGN phenomenon and their evolution with time are fundamental to an understanding of galaxy formation in the early universe.  Such studies may also shed new light on the formation and variety of the nearby, quiescent galaxies.  Studies of cosmic evolution for radio galaxies \\citep{1990MNRAS.247...19D}, optically selected quasars \\citep{2001AJ....121...54F, 2000MNRAS.317.1014B}, radio loud quasars \\citep{1996Natur.384..439S, 2005A&A...434..133W}, X-ray sources \\citep{2005ARA&A..43..827B}, and star formation \\citep{1998ApJ...498..106M} all indicate a similar pattern of an increase in luminosity and/or density up to epochs corresponding to redshifts of 1 to 2 followed by a rapid decline to current epochs. The near simultaneity (on cosmic time scales) of these apparently very different phenomena is curious.  What caused the cutoff in the formation of stars, galaxies, and quasars all around the same time?  Does the presence of an AGN stimulate star formation, or do high rates of star formation lead to AGN?  Or, is there some underlying phenomenon which is responsible for both?    \\subsection{Radio Studies}  The extragalactic radio source population ranges from normal galaxies with luminosities near $10^{19} \\rm ~Watts~Hz^{-1}$ to galaxies whose radio emission is as much as $10^7$ times greater owing to regions of massive star formation and/or to an AGN.  The population of radio sources in the sky with flux densities greater than 1 mJy is dominated by AGN driven emission in which virtually all of the energy is generated from the gravitational potential associated with a SMBH in the nucleus.  For these sources, the observed radio emission includes the classical extended jet and double lobe radio source which may extend up to Megaparsecs from the parent galaxy, as well as compact radio components that are more directly associated with the energy generation and collimation near the central engine of the AGN.  We will use the term ``radio galaxy'' to refer to those sources where the radio emission is primarily from extended lobes or jets with inverted power-law spectra, and which have optical counterparts that are of galactic dimensions and are typically associated with bright Elliptical galaxies.  We use ``AGN'' for those radio sources, including quasars, which are less than an arcsecond in size, have flat or inverted radio spectra characteristic of an opaque synchrotron source and whose optical counterparts appear stellar. For both the radio galaxies and AGN the power source is thought to be due to the heating of material from infall to a SMBH ranging from $10^6$ solar masses for low luminosity AGN to $10^8$ or more solar masses in the case of radio galaxies and quasars.  In all cases the radio emission is believed to be due to synchrotron emission from relativistic electrons in magnetic fields of the order of $10^{-5}$ Gauss.  Below 1 mJy there is an increasing contribution to the radio source population from synchrotron emission resulting from relativistic plasma ejected from supernovae of young stars associated with massive star formation in galaxies or groups of galaxies where mergers or interactions appear to be important \\citep{win95,ric98,fom02,2005A&A...441..879C,fom06}.  However, the mix of star-formation and AGN related radio emission and the dependence on  epoch is not well determined. Massive star-formation phenomena is a strong function of cosmic time and produces a variety of radiation at all wavebands: synchrotron emission at radio wavelengths; warm dust emission at infra-red and sub-mm wavelengths; shock and hot gas emission at optical and ultra-violet wavelengths; and thermal X-ray emission from the hot gas where massive stars are formed and from X-ray binary stars.  At least for low redshifts, the radio emission is tightly correlated with the FIR emission \\citep{1992ARA&A..30..575C, con91}. Because of the lack of absorption at radio wavelengths, the measured radio flux density can give a good estimate of star formation activity over $ \\sim 10^8$ years at fainter levels than are probed with FIR observations. However, the radio-FIR relation has not yet been tested out to large redshifts.  Also, \\citet{2003A&A...399...39R} report a radio- X--ray correlation for star forming galaxies.  Radio observations are particularly valuable in distinguishing between AGN and star-formation, as they see through the gas and dust which often surrounds both star forming regions and AGN. At radio wavelengths, routine observations even for the faintest sources can have the sub-arcsecond resolution needed to distinguish between radio galaxies, AGN and starburst morphology. In the later case, the radio source has dimensions about that of a galactic disk or of single regions of active star formation typically in the range 0.1 to 1 arcsecond at z $\\sim$ 1.  The AGN, by contrast have dimensions of the order of $10^{-4}$ to $10^{-2}$ arcsec, characteristic of an opaque synchrotron source of 1 to 100 parsecs, while the more powerful radio galaxies show the characteristic jet-multiple lobe morphology, often along with a milliarcsec AGN component.  High angular resolution of radio observations are also critical to unambiguously identify the host galaxy needed to establish the SED from observations over a wide range of wavelengths, to determine redshifts and luminosities, as well as the morphology of the host galaxy and to reveal associations in groups or clusters.  Previous deep radio observations have been made with the VLA in e.g., the Hubble Deep Field North \\citep{ric98,ric00}, in SSA13 \\citep{fom02,fom06} and other fields \\citep{1993ApJ...405..498W, mit85, cil00, 2003A&A...403..857B, 2004MNRAS.352..131S} and with the WSRT \\citep{oor85, 2000A&A...361L..41G, 2002AJ....123.1784D}.  In the southern hemisphere, the Australia Telescope Compact Array, (ATCA) has been used to study the Phoenix field \\citep{1998MNRAS.296..839H}, the Hubble Deep Field South \\citep{2003AJ....125..465H}, and the Chandra Deep Field South and ELAIS SWIRE fields\\citep{nor2006, A06}. Combined VLA and MERLIN observations of the HDFN \\citep{mux05} with a resolution of $0.2''$ show the complexity of the microjansky radio emission suggesting a mixture of AGN and star formation which contributes to the sub-millijansky radio population. In Paper IV, we show that the well known increase in the number of sub-millijansky radio sources is apparently not entirely due to a population of star forming galaxies, and that the fraction of AGN contribution to the sub-millijansky population is greater than previously thought.  Most of the microjansky radio sources found in these deep radio surveys are identified with galaxies brighter than about R=26.  About 15\\% are fainter; many are red in color and may be detected in the z and K band.  These extremely-red objects (EROs) clearly show the presence of dust which does not affect the centimeter radio emission, but produces the large IR and sub-mm emission and obscures the optical emission.  The $\\sim 0.2''$ position accuracy available with the VLA is crucial in locating their X-ray and optical counterparts and in determining the nature of the emission process.  This is particularly important for sub-mm sources with radio counterparts, as the sub-mm observations alone have poor positional accuracies and often ambiguous identifications.  Comparison of the radio and the X-ray position may help distinguish whether the radio emission is dominated by extra-nuclear star formation or by an AGN.  \\subsection{X--ray Studies}  X-ray luminosities greater than about $10^{35}$ Watts are usually, but not always, associated with AGN, QSOs, Seyferts, and other active galaxies containing broad emission lines \\citep{2003A&A...412..689P}. However, not all AGN are observed as strong X--ray sources in the Chandra or XMM bands.  Some may be heavily obscured \\citep{DRRP05} or Compton thick, a phenomenon which occurs for absorbing column densities of equivalent hydrogen atoms assuming solar metalicity are larger than $\\sim 10^{24}$~cm$^{-2}$, when the reflected component is expected to dominate the X-ray spectrum.  Intrinsic absorption of less than about $10^{22}$ to $10^{24}$~cm$^{-2}$, corresponds to Compton thin sources where the X-ray emission is usually easily observed.  At least 80 percent of these sources in the CDF-S are associated with narrow emission line type II objects.  Starting nominally at $1.5 \\times 10^{24}$~cm$^{-2}$,  the Comptonization effects start to be relevant, direct emission is strongly suppressed, and therefore we expect that most of the X--ray emission will be due to reflection from a slab of cold material.  In this case the observed X-ray emission may be only of the order of a few percent of the intrinsic power, and such objects are therefore expected to be only weak X-ray sources.  According to \\citet{2005MNRAS.357.1281W} half of the hard X-ray background (XRB) above 5 kev remains unresolved and is shown to be consistent with the XRB being due to Compton-thick sources at $z \\sim 1$.  This means that XMM and Chandra observations are missing a large part of this population of sources. Radio observations may be sensitive to these Compton-thick sources with suppressed X-ray emission; for example, \\cite{2005Natur.436..666M} use radio and FIR data to identify highly absorbed X-ray weak AGN.  X-ray sources with luminosities in the range $\\sim 10^{34}$ Watts to $10^{35}$ Watts include a mixture of low luminosity AGN and star forming galaxies which cannot be distinguished from X-ray observations alone.  Although the weaker X-ray sources appear to be increasingly dominated by star formation, unlike the weak radio surveys, the deep X-ray surveys do not show any evidence for mergers or other interactions \\citep{2005ARA&A..43..827B}.  \\subsection{The Chandra Deep Field South}  The CDF-S is probably the most intensely studied region of sky with extensive multiwavelegth observations using the most powerful space and ground facilities, and is uniquely suited for studies of the co-evolution of star formation and AGN. High sensitivity X-ray observations are available from Chandra \\citep{2002ApJS..139..369G,2005ApJS..161...21L}.  The GOODS \\citep{2004ApJ...600L..93G} and GEMS \\citep{RBB04} multiband imaging programs using the HST Advanced Camera for Surveys (ACS) give sensitive high resolution optical images over the CDF-S field \\citep{B06}. Ground based imaging and spectroscopy are available from the ESO 2.2m and 8m telescopes, and IR observations from the Spitzer Space Telescope. Also, located near the center of the CDF-S, is the Hubble Ultra Deep Field (UDF) with its unprecedented sensitivity in four optical bands reaching a limiting magnitude as faint as 29 at 775 nm.   In this paper we report new radio observations of the CDF-S, including the Hubble UDF, and the E-CDF-S, made with the NRAO Very Large Array (VLA) at 1.4 GHz (20 cm) and 5 GHz (6 cm).  The effective angular resolution was $3.5''$ and  minimum rms noise as low as $8.5~\\mu$Jy per beam at 20 cm and $7~\\mu$Jy per beam at 6 cm.  These deep radio observations complement the larger area, but less sensitive lower resolution observations of the CDF-S discussed by \\cite{nor2006} and \\cite{A06}. ",
        "conclusions": "}    The VLA survey of the CDF-S and E-CDF-S has cataloged 266 radio sources above a limiting flux density of $43~\\mu$Jy in the most sensitive part of the field.  Typical radio position accuracy of better than 1 arcsecond and multiwavelength imaging allowed optical or NIR counterparts to be identified for more than 95 percent of the cataloged source with spectroscopic or photometric redshifts available for about 70 percent of the radio sources.  Although for most of the unidentified sources the lack of a unique OIR counterpart is due to ambiguities resulting from multiple faint galaxies within the radio position uncertainty, there are a few empty fields with no apparent counterpart down to limiting magnitudes as faint as mag 25.9.  Most of the radio sources are unresolved or only barely resolved with the 3.5 arcsecond VLA beam at 20 cm, although about ten percent are well resolved with angular sizes greater than 5 arcseconds.  Most of the resolved sources have weak low surface brightness extensions.  Ten sources, all of which have a total flux density greater than 1 mJy, have multiple components characteristic of classical radio galaxies, while the submillijansky population is thought to be due to synchrotron radiation from regions of active star formation and low luminosity AGN.  Eighty-nine radio sources in our complete radio catalog were found to have X-ray counterparts in either the 1 Megasecond Chandra catalog or in the E-CDF-S. In addition to the cataloged radio sources, we also give the measured radio emission from X-ray sources found in the CDF-S catalog.  An additional 30 CDF-S X-ray sources were detected at 20 cm were found with SNR between 3-sigma and 5-sigma.  Below a few hundred microjanskys, the differential radio source count for the CDF-S approaches the ''Euclidian\" slope of -2.5 corresponding to the evolving population of starforming regions and low luminosity AGN.  The scatter among different surveys is large, and may be due, at least in part, to cosmic variance.  Although, the uncertainties in the instrumental corrections for the weaker sources, especially for resolution, are also large and may account for some of the discrepencies among observers, it is hard to explain the observed large scatter at the few hundred microjansky level as the result of instrumental corrections.  However, we note the significant difference in the flux density scales of the VLA and ATCA observations of the same CDF-S.  Padovani et al. (in preparation) use these data to discuss the radio/X-ray/OIR relationships and the evidence for contribution to the microjansky radio emission from AGN as well as from star formation  \\cite{M08} has used the VLA at 1.4 GHz to cover the full E-CDF-S with a sensitivity $\\sim$~$8~\\mu$Jy rms, which will be complemented by deeper new Chandra X-ray and Spitzer IR observations as well as new VLA 6 cm observations.  Later, ALMA with its unprecedented sensitivity and resolution at sub-millimeter wavelengths will help to unravel the nature of the sub-millijansky radio source population and the relation between AGN and star forming activity."
    },
    "0806/0806.2040_arXiv.txt": {
        "abstract": "We compare the results from a state-of-the-art semi-analytic model of galaxy formation with spectro-photometric observations of distant galaxy clusters observed in the range $0.8\\leq z\\leq 1.3$. We investigate the properties of their red sequence (RS) galaxies and compare them with those of the field at the same redshift. In our model we find that i) a well-defined, narrow RS is obtained already by $z\\approx 1.2$; this is found to be more populated than the field RS, analogously to what observed and predicted at $z=0$; ii) the predicted U-V rest-frame colors and scatter of the cluster RS at $z=1.2$ have average values of 1 and 0.15 respectively, with a cluster-to-cluster variance of $\\approx 0.2$ and $\\approx 0.06$, respectively. The scatter of the RS of cluster galaxies is $\\approx $ 5 times smaller than the corresponding field value; iii) when the RS galaxies are considered, the mass growth histories of field and cluster galaxies at $z\\approx1.2$ are similar, with 90 \\perc of the stellar mass of RS galaxies at $z=1.2$ already formed at cosmic times $t=2.5$ Gyr, and 50 \\perc at $t=1$ Gyr; v) the predicted distribution of stellar ages of RS galaxies at $z=1.2$ peaks at $3.7$ Gyr for both cluster and field populations; however, for the latter the distribution is significantly skewed toward lower ages. When compared with observations, the above findings show an overall consistency, although the average value $\\Delta_{U-V}\\approx 0.07$ of the observed cluster RS scatter at $z\\approx1.2$ is smaller than the corresponding model central value. We discuss the physical origin and the significance of the above results in the framework of cosmological galaxy formation. ",
        "introduction": "Environment-dependent properties of galaxies constitute a natural test-ground for cosmological theories of galaxy formation. These envisage the properties of the galaxy populations to originate from the primordial density field dominated by the Dark Matter (DM); the collapse and growth of small-scale perturbations (leading to galaxy progenitors) is thus modulated by the underlying large-scale density field whose overdense regions will collapse to form groups and clusters of galaxies.  Among the galaxy properties, the color distribution constitutes a major issue, as colors measure the ratio of the present star formation rate to the overall mass of formed stars in a given galaxy. Three fundamental observational results deal with the star formation history and the color distribution of field galaxies: a) The global star formation decrease from $z=1.5$ to the present by almost 2 dex (see, e.g., the compilation by Hopkins 2004 and references therein); b) The downsizing (Cowie et al. 1996), i.e., a stronger ($\\sim 3$ dex) and earlier (starting at least at $z\\approx 2$) decline of star formation in massive galaxies compared  to a slower ($\\approx$ 1 dex) decline of the star formation rate of the present blue/late-type galaxies from $z\\sim 1.5$; correspondingly, an inverse correlation of the age of stellar population with the galaxy stellar mass is observed; c) The bimodality (Strateva et al. 2001; Baldry et al. 2004; Bell et al. 2004), i.e., a marked segregation in color between the blue galaxy population (dominant at faint magnitudes and for low-mass galaxies) and the red population (dominant for bright, massive galaxies) populating the ''red sequence'' (RS).  In hierarchical clustering models, the first property naturally results from the high star formation rate at high redshifts $z\\gtrsim 2$ sustained by the high cooling efficiency in the dense environments of galaxy progenitors and by effective starbursts; the decline at $z\\lesssim 2$ follows from the exhaustion of the gas reservoir converted into stars at higher $z$. The interpretation of downsizing in hierarchical galaxy formation is also straightforward: bright, massive galaxies form from the coalescence of protogalaxies which collapsed in a biased, overdense region of the primordial density field; thus such progenitors collapsed and began to form stars at earlier cosmic epochs, and their cold gas reservoir is exhausted earlier thus yielding older (and hence redder) stellar population in the final galaxy. The origin of the bimodality is more controversial; recent works (see discussion in Neistein, van der Bosch, Dekel 2006; Dekel \\& Birnobim 2006; Menci et al. 2006) suggest that self-regulation of star formation (due to Supernovae feedback) can be effective below a critical (DM) mass scale $M\\approx 5\\times 10^{11}\\,M_{\\odot}$; the inclusion of feedback from AGNs then quenches further star formation in massive haloes.  The above properties are strongly modulated by the galaxy environment. Not only the morphology has long been recognized to strongly depend on the environment (with the fraction of early-type galaxies increasing in groups and clusters of galaxies, Spitzer \\& Baade 1951), but also the fraction of galaxies belonging to the blue population has been recently recognized to decrease in dense environments (Baldry et al. 2004, 2006; Gerke et al. 2007; Cooper et al. 2006). The environmental dependence of the blue fraction $f_b$ is stronger than the above mentioned luminosity dependence observed in the field. This means that the environmental density enhances the probability for an early star formation followed by quenching even for low or intermediate mass galaxies.  The scenario emerging from the above observational framework is the following: galaxies flow from the blue population to the red sequence (corresponding to the decline of star formation at $z<2$) at rates and cosmic epochs that depend very strongly on their mass and environment. Massive galaxies have converted most of their gas into stars and moved to the red sequence very early at $z>2$ in both cluster and field, since at high redshift their star formation was not effectively self-limited by feedback. On the other hand, lower mass galaxies $M<10^{12}\\,M_{\\odot}$ move to the RS later in the field compared to clusters. The speed-up of the transition to the red sequence of intermediate/low- mass galaxies in dense environments (groups and clusters) can be contributed by a number of physical processes: i) Biased galaxy formation; in principle, galaxies later included in groups and clusters form from clumps which collapsed in biased, overdense regions of the primordial field, hence characterized by an earlier star formation. However, the initial overdensity corresponding to the final group or cluster may be so tiny that such an effect might be too small.  ii) Starbursts resulting from galaxy merging and fly-by (sometimes referred to as ''harassment''); while merging is actually enhanced in groups compared to clusters (see Cavaliere, Colafrancesco \\& Menci 1992), the minor - but more frequent - starbursts originated from galaxy grazing encounters may significantly accelerate the star formation at high redshift in groups and clusters, and hence anticipate the transition to the red sequence. Note however that this process is effective mainly for mid-size galaxies which strike the best compromise between cross section and abundance.  iii) Strangulation; in small-mass galaxies the gas may be stripped off when they are included in the cluster (Larson, Tinsley \\& Caldwell 1980); in most semi-analytic models this is captured by dispersing their hot gas mass through the whole host halo (see, e.g., Balogh, Navarro \\& Morris 2000). Note that our model, similarly to canonical semi-analytic approaches,  does not deal with the ram pressure stripping of the dense, cold gas in satellite galaxies. iv) AGN feedback; to some extent, such a mechanism must be at work, and there is a general agreement (Ciotti \\& Ostriker 1997; Silk \\& Rees 1998, Haehnelt, Natarajan \\& Rees 1998, Fabian 1999) that it may contribute to suppress any residual star formation in massive galaxies; in models where the AGN phase is triggered by galaxy interactions, the AGN feedback depends on the environment which modulates the encounter probability. The effectiveness of the AGN feedback in quenching the black hole growth and the subsequent star formation has been confirmed by recent aimed simulations of galaxy collisions triggering AGN activity (Di Matteo, Springel \\& Hernquist 2005).  We have investigated the role of such a process in a previous paper (Menci et al. 2006) showing that it can affect the color distribution of most massive galaxies by removing the residual fraction of such galaxies that otherwise would populate the blue branch of the color distribution.  The relative role of all these processes is still matter of investigation. Their dependence on the mass-scale of the host structure (group or galaxy cluster) would provide an important clue to disentangle their contribution to the observed environment-dependent properties of galaxies. Several recent papers (Cooper et al. 2007, Gerke et al. 2007) point out that the difference in the properties of the RS between field and overdense environments remains unchanged up to $z\\approx 1$ but tends to vanish at redshifts $z\\gtrsim 1.3$. Since these results are based on samples which probe the environments up to the group scale (rich clusters are not included), the epoch $1.2\\lesssim z\\lesssim 1.5$ would represent a measure of the effectiveness of poor/medium environments in modulating the star formation of member galaxies; at such redshifts, the environment-induced decline of the star formation would become stronger than the average cosmological decline. Probing the transition to the RS of galaxies in richer, extreme environment at $1\\lesssim z\\lesssim 1.5$ is then crucial.  In this context, recent HST/ACS observation of a sample of clusters at $0.8<z<1.3$ have fueled much progress. E.g., Blakeslee et al. (2003, 2006) and Mei at al. (2006a; 2006b) have found a well defined and extremely narrow red sequence with an intrinsic scatter of $\\sim 0.05$ mag (in the color $i_{775}-z_{850}$ as observed with HST/ACS), much smaller than the corresponding field RS. This points toward the existence of a strong environment-dependent effect which, for galaxies in extremely dense environments, has already been effective at $z\\approx 1.5$, thus supporting the view that the denser the environment, the earlier the transition to the RS.  How do such results compare with hierarchical models ? While such models essentially capture the transition of faint galaxies to the RS (at least in the redshift range $z\\leq 0.8$, see De Lucia et al. 2007; for results of  hydro/N-body simulations see Romeo et al. 2008 ), the tightness of the RS constitutes an impressive constraint for them, since in the hierarchical scenario the growth of DM haloes constitutes - on average - a gradual process, driven by the merging of sub-clumps, the progressive cooling of gas at the center followed by a quiescent star formation. Since the tightness of the observed RS indicates a fast transition of the galaxies to the same region of the color distribution acting on very short time scales in the early epochs, this is particularly challenging for hierarchical models which predict clusters of galaxies to be the last structures to virialize (typically at $z<2 $). On the other hand, in such models, high-redshift clusters also constitute the most biased environments for galaxies to form, hence enhancing the effectiveness of the processes i)-iv) described above in accelerating the transition to the RS. Thus, comparing the observed properties of the RS in high-z clusters with model predictions constitutes an extremely sensitive probe of hierarchical models and of the physical processes there included.  Here we compare the results of a state-of-the-art semi-analytic model of galaxy formation with the properties of the RS of galaxy clusters at high redshift, the most extreme (biased) galaxy environments observed so far. We also emphasize that the observations used in this work are currently the most suited for such an analysis, as with ACS observations in the $i$ and $z$ bands one can achieve the photometric accuracy to reliably estimate color scatters out to $z\\sim\\! 1.3$.  The model we adopt is well suited for such an investigation since it includes all the processes i)-iv) discussed above, which are though to contribute to the environmental dependence of the star formation histories. In particular, it includes both the burst of star formation due to grazing encounters (important in the early building up of massive galaxies) and a description of the AGN evolution which has been extensively tested against observations both in the optical (Menci et al. 2003) and in X-rays (Menci et al. 2004b).  In previous papers (Menci et al. 2005; 2006) we have successfully compared the model with observations concerning the galaxy color distribution at low redshifts as a function of both luminosity and environment, and at high-redshift we showed how the inclusion of AGN feedback allows to suppress residual star formation in massive galaxies yielding the observed abundance of EROs.  The paper is organized as follows: in Sects. 2 and 3 we briefly recall the main properties of the semi-analytic model we use, and the observational data set we compare with. In Sect. 4.1 we show the model results concerning the color-magnitude relation both at low and at high redshift in the cluster and in the field, and we briefly discuss how it compares with existing observations. In sect. 4.2 we focus on the properties of the RS in high-redshift clusters, and compare the model results concerning the scatter and the normalization of the RS with the ACS observations of distant clusters. We also provide predictions concerning the distribution of the scatter and the normalization of the RS among the cluster population, which could be verified with existing or future observations. In Sect. 4.3 we also compare the mass assembly history and the age distribution of $z\\approx 1.2$ galaxies in the field and in clusters resulting from the model with estimates derived from fitting observed galaxy colors and spectra with spectral synthesis models. In Sect. 4.4 we interpret our results on the basis of the physical processes which determine the scatter of the RS in hierarchical models. In Sect. 5 we draw our conclusions. ",
        "conclusions": "In this paper we have compared the results from a state-of-the-art semi-analytic model of galaxy formation to spectrophotometric observations of the most distant galaxy clusters observed in the range $0.8\\lesssim z\\lesssim 1.3$, to investigate the properties of the red sequence (RS) galaxies in high-redshift clusters and to compare them with those of the field in the same stellar mass range.  A long standing prediction of cosmological galaxy formation models (see Kauffmann 1996; Diaferio et al. 2001 ) is that the {\\it average} stellar populations of galaxies in clusters are older, since the star formation processes are accelerated in structures originated from overdense regions of the primordial density field. Recent observational progress (Baldry et al. 2004) has revealed in detail such an average reddening of galaxies for increasing environmental density, due to an increased population of the red branch of the bimodal color distribution. Thus, the redder {\\it average} colors of the cluster galaxy populations we observe today are not caused by a global shift of the ages of the galaxy populations toward lower values, but instead by a different balance between the populations of the two branches of the color distribution.  In particular our results can be summarized as follows  \\begin{itemize}  \\item The speed up of the transition of galaxies from the blue cloud to the red sequence which takes place in dense environments is qualitatively reproduced by the hierarchical model for a wide range of galaxy masses (being faster for more massive galaxies) and cosmic times. In the case of distant galaxy clusters, the most extreme environments which provide the most challenging test for hierarchical models, such a fast transition in the model is able to reproduce a well defined red sequence in clusters at $z\\approx 1.3$, which is narrower than in the field, as observed. The fraction of galaxies which belong to the RS is larger in clusters than in the field from $z=0$ to $z=1.5$, and stays nearly constant in such an interval.  \\item In agreement with observations, we find that at $z\\approx 1.2$ the RS is not universal, i.e. is subject to a cluster-to-cluster variance over the cluster population. The predicted color normalization is close to the observed average value and is distributed in a narrow range $\\langle c\\rangle \\approx 0.9-1.1$, when different clusters are considered. Observations (e.g. Mei et al. 2007) suggest a cluster-to-cluster variation of $\\sim\\! 0.05$ mag, a measurement which is limited by sample size and uncertainties in the K-correction and passive luminosity evolution in the redshift interval $0.8\\leq z\\leq 1.3$.  \\item The scatter of the model RS of cluster galaxies is found to be much tighter than that of the field by a factor ranging from 5 to 10. However, on average, they are still larger than the observed tight scatters ($\\approx 0.05-0.08$ mag) which are consistent only with the lower envelope of model distribution enclosing 75 \\perc of the clusters selected in the Monte Carlo simulation.  \\item When {\\it only} the RS galaxies are considered, the mass growth histories of field and cluster galaxies at $z\\approx1.2$ are similar. This is consistent with the mass growth derived from observations through SED and spectral fitting (Gobat et al. 2008) of massive early-type galaxies with adjustable star formation laws. Both observations and the model converge in indicating that at a cosmic time $t\\approx 2.5$ Gyr more than 90 \\perc of the stellar mass finally assembled in the RS galaxies at $z=1.2$ has already been formed in both field and clusters. The model also suggests that at cosmic times $t\\approx 1$ Gyr the stellar mass fraction was about 50 \\perc in both field and clusters, whereas spectral-synthesis modeling of the spectrophotometric data suggests a faster time scale in assembling stellar mass in clusters. Even though our model predicts the scatter of the growth histories to be larger in the field than in clusters, a slightly faster assembly time scale seems to be needed to account for the very tight observed RSs.  \\item The corresponding predicted ages of the stellar populations in RS galaxies at $z\\approx 1.2$ reflect the above behavior, with distributions peaking at values $\\tau\\approx 3.5$ Gyr in both field and clusters. These are consistent with the values derived from observations through SED+spectral fitting, although the latter indicate a slightly larger (by $\\lesssim 0.5$ Gyr) age for cluster RS galaxies as compared with the model (Gobat et al. 2008; Rettura et al. 2008).  The model prediction of a larger dispersion of ages for field RS galaxies compared to clusters is in good agreement with observations. Indeed, a negligible fraction of RS galaxies are predicted and observed to have ages shorter than 2 Gyr, while an appreciable - though minor- fraction ($\\approx 0.2$) of field RS galaxies have ages shorter than that.  \\item The reduced scatter in color, ages, stellar mass growth, of cluster RS galaxies at $z\\approx 1.2$ compared to field RS galaxies is originated from intrinsic properties of the merging histories of cluster galaxies. These are characterized by an earlier (compared to the field) assembly of galaxy progenitors into a main progenitor at times $t\\lesssim 1.5$ Gyr. After such a time, the bulk of the stellar mass formation of cluster RS galaxies takes place in their main progenitor; thus the scatter in mass growth histories, ages and colors is considerably suppressed, since the variance introduced by the different merging histories is considerably reduced compared to the field, where the star formation is distributed among a larger number of progenitors.  \\end{itemize}  The picture emerging from the above results is the following. For increasing cosmic time, the exhaustion of cold gas in galaxies and the subsequent decrease of the star formation activity drives the galaxies to flow from the ''blue cloud'' region of the color-magnitude space to the RS; the more massive the galaxies (and the earlier their gas conversion into stars), the earlier is the epoch of drift to the RS.  In dense environments, and in particular for high redshift clusters, the drift of faint galaxies from the blue cloud to the RS is accelerated by various processes (namely, processes i)-iv) recalled in the introduction) so that also fainter galaxies rapidly drift to the RS. As a result, the blue fraction in clusters is significantly smaller than in the field, and at high-redshift a well defined, narrow RS is already in place.  When only the galaxies that have drifted to the RS are considered, their {\\it average star formation history} does not depend strongly on the environment and is not responsible for the reduced scatter of the RS in clusters compared to the field; the effect of the environment is to accelerate the drift to the RS (thus enhancing the relative population of the red branch compared to the blue branch in the color distribution) rather than changing the properties of the RS itself. This is supported by the detailed measurements of the color distribution of galaxies as a function of the environment in Baldry et al. (2004).  On the other hand, the {\\it mass assembly history} of RS galaxies in clusters is considerably different from that occurring in the field, and has a significant impact on the scatter of the RS.  Since cluster galaxies are formed from clumps that collapsed within high-density regions of the primordial density field, their assembly into a dominant progenitor is faster compared to the field, the formation of the bulk of the stellar mass of cluster RS galaxies takes place in their main progenitor, and the variance introduced to the star formation history by the different merging histories is considerably suppressed. The more biased the galaxy environment, the larger the above effect, so that for RS galaxies in high-redshift clusters the dispersion of mass growth histories, ages and colors is appreciably reduced compared to the field.  Despite the overall good agreement between our models and observations on the color, luminosity and mass distributions of galaxies in high and low density environments, our analysis has shown that even faster mass assembly or shorter star formation time scales are needed to reproduce the surprisingly tight RS in the massive clusters at $z\\approx 1.2$ (see fig. 5). This is extremely difficult to achieve in the model, since the dynamical history of DM haloes does not constitute a tunable component for a given initial power spectrum. Indeed, we have also verified that changing the normalization of the power spectrum, while affecting the abundance of clusters at a given redshift, is not effective in further narrowing down the scatter of the color-magnitude relation. Thus, we propose that the acceleration of mass assembly in cluster galaxies (making the bulk of the effect) must be complemented with a further acceleration of star formation in RS cluster galaxies. This is indicated also by the slightly mismatch between model and observations in the stellar mass growth history (fig. 6) and average ages (fig. 7), when the RS cluster galaxies are considered. The modeling of such an additional enhancement of early star formation in the progenitors of cluster RS galaxies will probably require a finer treatment of the environmental dependence of the interaction-driven starbursts. An additional issue is the slope of the cluster RS obtained from the model. Current observations find no significant evolution with redshift of the RS slope which remains close to $| \\delta (U-B)/\\delta B| \\approx 0.04$ (Mei et al. 2008, submitted), whereas the model predicts a (non-evolving) flat slope. This can be due to both an over-quenching of star formation at low masses, but also to a poor treatment of chemical evolution (e.g., the contribution of SNI is not properly accounted for in semi-analytic models) which is known  to affect the RS slope through metallicity effects (Kodama \\& Arimoto 1997).  Finally, we mention two additional consequences of our interpretation of the model results, both consistent with preliminary indications from observations. First, at redshift $z\\gtrsim 1.5$ our model predicts the scatter of the cluster RS to extend to a wider range of values compared to the low-redshift interval $z\\lesssim 1.5$ (see fig. 4) and to easily reach values larger than 0.2 (see fig. 5), consistent with preliminary results by Gobat et al. (2008). Second, the model predicts the speed-up of star formation in dense environments to be larger for small mass galaxies (with $M\\lesssim 10^{10.5}\\,M_{\\odot}$) (as inferred, e.g., by fig. 3) compared to massive RS galaxies (which indeed are predicted to have similar mass assembly histories in the field and in clusters, see fig.6). This is also qualitatively consistent with preliminary results by Gobat et. a. (2008), which indicate a lower average age of stellar populations in clusters only for galaxies with $M\\lesssim 10^{11.2}\\,M_{\\odot}$. We shall address in detail both points in a next paper."
    },
    "0806/0806.3701_arXiv.txt": {
        "abstract": "It is widely assumed that the observed universe is accelerating due to the existence of a new fluid component called dark energy. In this article, the thermodynamics consequences of a nonzero chemical potential on the dark energy component is discussed with special emphasis to the phantom fluid case.  It is found that if the dark energy fluid is endowed with a negative chemical potential, the phantom field hypothesis becomes thermodynamically consistent with no need of negative temperatures as recently assumed in the literature. ",
        "introduction": "The  current idea of an accelerating Universe driven by dark energy is based on a  large convergence of independent observational results, and its explanation constitutes one of the greatest challenges for our current understanding of fundamental physics \\cite{SN,CMB}. The origin and the nature of dark energy is still a mystery, however, there is no doubt that its existence is beyond the domain of the standard model of particle physics \\cite{review}.  Among a number of possibilities to describe this dark energy component, the simplest and most theoretically appealing way is by means of a cosmological constant $\\Lambda$, which acts on the Einstein field equations as an isotropic and homogeneous source with a constant equation of state parameter $p/\\rho=-1$. On the other hand, although cosmological scenarios with a $\\Lambda$ term might explain most of the current astronomical observations, from the theoretical viewpoint they are plagued with some fundamental problems thereby stimulating the search for alternative dark energy models driven by different candidates \\cite{model,list1}.  In 1995, Lima and collaborators \\cite{limamaia,LA04} analyzed several thermodynamic and statistical properties of a dark energy fluid. They  assumed a dark energy fluid component phenomenologically described by an equation of state $p=\\omega \\rho$ with null chemical potential. Later on, Lima and Alcaniz (and independently Brevik et al. \\cite{B}) stressed that their theoretical  thermodynamic treatment ruled out the case of phantom energy because the comoving entropy of a dark component with $\\omega < -1$ is negative (see Reference \\cite{REF} for other interesting discussions of phantom fluids).  However,  thermodynamic arguments in favor of the phantom hypothesis were put forward by Gonz\\'alez-D\\'{i}az and Sig\\\"uenza \\cite{gonzalez}. They claimed that the temperature of a phantomlike fluid is always negative in order to keep its  entropy positive definite (as statistically required).  In this work we reanalyze the thermodynamics properties of an expanding universe  filled with a dark energy fluid endowed with a non-zero chemical potential. As we shall see, the main effect of a chemical potential is that the phantom scenario becomes thermodynamically consistent with no need to assume negative temperatures. ",
        "conclusions": "In this paper we have investigated the thermodynamics properties of a dark energy fluid with equation of state, $p=\\omega \\rho$, by assuming that its chemical potential is different from zero.  In Figure 1, we summarize the main results of the present analysis including a chemical potential to the dark energy fluid.  As discussed in previous section, the regions with $S < 0$ are  thermodynamically forbidden. Note also that many dark energy fluids satisfy the combined constraints regardless of the $\\mu$ sign, that is, a large interval of  negative $\\omega$ values is allowed from thermodynamic considerations.  However, a phantom like behavior ($\\omega < -1$) is permitted only for $\\mu < 0$. It should be stressed that for $\\mu=0$ one finds  $\\omega_{min}=-1$ (see Eq. (\\ref{accord})) in accordance to the results previously derived by Lima and Alcaniz \\cite{LA04}. The present analysis with $\\mu \\neq 0$  also opens  the possibility for an equation of state parameter $\\omega < -1$, thereby recovering the idea of a phantom dark energy with no need of negative temperatures.  Therefore, as far as we known, the inclusion of a negative chemical potential to the dark energy fluid is the only way to save the phantom hypothesis without to violate basic thermodynamic properties."
    },
    "0806/0806.0228_arXiv.txt": {
        "abstract": "We analytically and numerically study the motion of electrons along a magnetic loop, to compare with the observation of the propagating feature of the non-thermal microwave source in the 1999 August 28 solar flare reported by Yokoyama et al. (2002). We model the electron motion with the Fokker-Planck equation and calculate the spatial distribution of the gyrosynchrotron radiation. We find that the microwave propagating feature does not correspond to the motion of electrons with a specific initial pitch angle. This apparent propagating feature is a consequence of the motion of an ensemble of electrons with different initial pitch angles, which have different time and position to produce strong radiation in the loop. We conclude that the non-thermal electrons in the 1999 August 28 flare were isotropically accelerated and then were injected into the loop. ",
        "introduction": "\\label{sec:introduction} In solar flares, a large amount of released energy is expended for production of non-thermal particles. They propagate in the solar corona and produce non-thermal emissions such as hard X-rays (HXRs), $\\gamma$-rays, and microwaves. The observed non-thermal emissions have key information on acceleration and transport mechanisms of particles.  Flare non-thermal emissions are signatures of propagation of accelerated particles. For HXRs, \\cite{1994PhDT.......335S} analyzed flares with double-footpoint HXR sources observed with {\\Yohkoh} \\citep{1991SoPh..136....1O}, and showed that the time variation of the emissions from double sources coincides with each other within 0.1 s. This indicates that electrons, not ions, propagate along the magnetic loop, reach the footpoints, and then emit HXRs there. \\cite{2003ApJ...595L..69L} observed intense HXR and $\\gamma$-ray emissions in the 2003 July 23 flare, by using {\\it Reuven Ramaty High Energy Solar Spectroscopic Imager} \\citep[{\\RHESSI};][]{2002SoPh..210....3L}. They found that the location of the $\\gamma$-ray line emission source, which corresponds to a precipitation site of accelerated ions, does not coincide with the HXR sources. This implies that the acceleration site and/or propagation path of ions differs from that of electrons.  For meter-wavelength radio emissions, the so-called type III radio burst \\citep[e.g.,][]{1998ARA&A..36..131B} is thought to be due to electromagnetic conversion of Langmuir waves, which are excited by electrons propagating from a flare site toward the interplanetary space along open magnetic field lines. For centimeter-wavelength radio emissions, the microwave source located at a remote distance from a flare main site \\citep[microwave remote source;][]{1996SoPh..165..275H,1997SoPh..173..319H,1997ApJ...489..976N} is interpreted as the precipitation of electrons propagating from the main site where the energy release and particle acceleration occur. \\cite{1999PASJ...51..483H} showed that the light curve at the microwave remote source has a similar profile to the HXR light curve at the main site, but is slightly delayed by roughly the loop transit time of electrons. {On the other hand, \\cite{2008ApJ...677.1367A} reported a flare in which the microwave light curve slightly leads the HXR. They argued that this time delay, opposite to that found by \\citeauthor{1999PASJ...51..483H}, could be realized if electrons are injected into an asymmetric magnetic loop toward a footpoint with a stronger magnetic field.} \\cite{2008ApJ...673..598M} analyzed the spatial and spectral variations of HXR and microwave emissions in the 2003 May 29 flare, and showed that the observed characteristics can be interpreted as a consequence of the electron transport called trap-plus-precipitation \\citep{1976MNRAS.176...15M}.  There are a few microwave observations of the direct detection of propagating electrons from a time series of images. \\cite{1994kofu.symp..199B} reported the motion of the 8.4 GHz and 15 GHz microwave sources with a speed of 3,000 ${\\rm km \\; s^{-1}}$. \\cite{2000ApJ...533L.167W} reported the rapid motion of the 0.33 GHz microwave source with a speed of $2.6 \\times 10^{4} \\; {\\rm km \\; s^{-1}}$. The most striking observation was reported by \\cite{2002ApJ...576L..87Y} (hereafter Y2002). They observed a flare that occurred on 1999 August 28, using Nobeyama Radioheliograph \\cite[NoRH;][]{1994PROCIEEE...82..705} with 0.1 s cadence. From a time series of NoRH 17 GHz images, they found two different classes of propagating features of non-thermal microwave sources along the loop (see Fig. 3 in Y2002). One is a slower propagation at a speed of $6 \\times 10^{3} \\; {\\rm km \\; s^{-1}}$, and the other is a faster propagation at a speed of $9 \\times 10^{4} \\; {\\rm km \\; s^{-1}}$.  The origin of the slower propagating feature is discussed by \\cite{2007A&A...465..613S}. They interpreted this feature as propagation of low-frequency whistler waves that can trap the parent electrons by wave-particle interactions at their turbulent front. The phase velocity of low-frequency whistler waves is roughly in agreement with the observation. The origin of the faster propagating feature is discussed by Y2002. They interpreted this propagating feature as a free stream of relativistic electrons along the loop. Assuming that the 17 GHz microwave-emitting electrons are almost relativistic ($v \\sim c > 9 \\times 10^{4} \\; {\\rm km \\; s^{-1}}$), they concluded that the parent electrons are injected with large pitch angle, $\\sim 70^{\\circ}$.  The interpretation of Y2002 on the microwave propagating feature is subject to the following assumptions. First, they assumed that electrons freely stream along the loop. However, since the flare loop must be a converging magnetic loop, electrons suffer the magnetic mirroring force. Second, they assumed that the apparent motion of the microwave source, which is generated via {\\gyros} radiation \\citep{1969ApJ...158..753R,1972SoPh...26..151T,1981ApJ...251..727P,1985ARA&A..23..169D,1999spro.proc..211B}, corresponds to the trajectory of the parent electrons. However, since the {\\gyros} radiation mechanism intricately depends on many physical parameters, it is not necessarily evident that the apparent motion of the radiation source is identical to the trajectory of the electrons with a specific pitch angle.  In this paper we reconsider the rapidly propagating feature of the microwave source in the 1999 August 28 flare and address the pitch-angle distribution of the electrons, by refining the treatments of the emission mechanism as well as the electron motion. In {\\S}~\\ref{sec:analytic-treatment} we present an analytic solution of the electron motion in a converging magnetic loop. {\\S}~\\ref{sec:numerical-model} shows our numerical model of the electron motion along the loop, in which we use the gyro-averaged Fokker-Planck equation to determine the electron phase space distribution. We calculate the {\\gyros} intensity distribution along the loop from the calculated electron distribution for comparison with the observation. In {\\S}~\\ref{sec:result-discussion} we present our calculation results. The results do not support the interpretation of Y2002 but suggest that the electrons in the 1999 August 28 flare were isotropically accelerated and then were injected into the loop. In {\\S}~\\ref{sec:summary} we summarize this paper. ",
        "conclusions": "We present our model calculation results and address whether the apparent propagating motion of the microwave source reported by Y2002 actually corresponds to the motion of electrons injected with a specific initial pitch angle. To discuss this, we perform calculations for two cases of the pitch-angle distribution of the initial condition ($\\phi(\\mu)$ in eq. (\\ref{eq:ini_con})): narrow-angle injection ({\\S}~\\ref{sec:narrow-band-case}) and broad-angle injection ({\\S}~\\ref{sec:broad-band-case}) cases.  \\subsection{Narrow-Angle Injection}\\label{sec:narrow-band-case} We first consider that initially injected electrons have an almost unique pitch angle, following the interpretation made by Y2002. To simulate this case, we give an initial pitch-angle distribution peaking at $\\mu = \\pm \\mu_{\\rm p}$: \\begin{eqnarray} \\phi(\\mu) = \\exp \\left[ -\\left(\\frac{|\\mu|-\\mu_{\\rm p}}{0.1}\\right)^{2}\\right]. \\label{eq:ini_mu_narrow} \\end{eqnarray}  The top left panel of Figure \\ref{fig:result} shows the time variation of the electron number distribution along the loop. We show the calculation result with $\\mu_{\\rm p} = \\mart{1-M^{-1}} = 0.612$, which is equal to the loss cone angle cosine $\\mu_{\\rm c}$. The vertical axis corresponds to the spatial coordinates from the loop top to the footpoint. We use the 1 MeV electron distribution integrated over pitch angle for this illustration. The solid and dashed lines show the trajectories of electrons with initial pitch-angle cosine of respective $\\pm 0.61$ and $\\pm 0.4$, obtained by solving equation (\\ref{eq:drdt}) with equations (\\ref{eq:mag_mom}) and (\\ref{eq:pbr}) (not eq. (\\ref{eq:quadb})) by using the 4th order Runge-Kutta method.  Initially injected electrons move toward the footpoint and about a half of them reach there after 0.2 - 0.4 s. They are lost from the calculation box because they are initially in the loss cone. Remaining electrons bounce back by the magnetic mirror and move toward the opposite direction. Subsequent diffusion by Coulomb collisions is ignorable because the collision time  ($\\sim 200$ s) is much longer than the loop transit time ($< 1$ s).  { We calculate the electron distribution not only in space but also in energy (see eqs. (\\ref{eq:fp}) and (\\ref{eq:ini_con})). Since the electron propagation is velocity-dispersive, the time variation of the electron number distribution differs with energy. However, relativistic electrons show almost the same profile as the top left panel of Figure \\ref{fig:result} because their velocity is nearly $c$. }  The middle and bottom left panels of Figure \\ref{fig:result} show the time variation of the 17 GHz {\\gyros} intensity distribution along the loop in the top- and side-view cases, respectively. Note that the vertical scale of the middle left panel is slightly different from other panels because of the projection effect of the loop.   The strong radiation source is localized in space along the loop. In the side view (bottom left), radiation primarily comes from the footpoint. In the top view (middle left), strong radiation comes from the intermediate position between the loop top and footpoint. These results are attributed to the dependence of the {\\gyros} emissivity on a viewing angle as well as the parent electron distribution along the loop. We interpret these results by using the analytic solution in {\\S} \\ref{sec:analytic-treatment}.  The {\\gyros} radiation is primarily produced by mildly relativistic ($\\sim 1$ MeV) electrons with the pitch-angle cosine of $\\mu \\sim \\beta \\cos \\theta$ \\citep{1981ApJ...251..727P,1990ApJ...354..735L}. Given all electrons with the initial pitch-angle cosine $\\mu_{0}$ at the loop top, their pitch-angle cosine is given by equation (\\ref{eq:mag_mom}) with (\\ref{eq:quadb}) as a function of $r$. The viewing angle $\\theta$ is also given as a function of $r$, which is different between the top- (eq.~(\\ref{eq:theta_disk})) and side-view ($\\theta = 90^{\\circ}$) cases. By solving $\\mu = \\beta \\cos \\theta$ with respect to $r$, we obtain the position $r_{\\rm g}$ at which strong radiation are produced. In the top view, the equation is written as \\begin{eqnarray} \\mart{\\mu_{0}^{2} - (1-\\mu_{0}^2)(M-1)\\left(\\frac{r_{\\rm g}}{r_{\\rm f}}\\right)^{2}} = \\beta \\tanh (k r_{\\rm g}). \\label{eq:rg_disk} \\end{eqnarray} The solution $r_{\\rm g}$ of equation (\\ref{eq:rg_disk}) is smaller than $r_{\\rm f}$, as long as $\\mu_0$ is smaller than the critical value (obtained by taking $r_{\\rm g} = r_{\\rm f}$ in this equation). In the narrow-angle injection case, we assume that injected electrons have an almost unique initial pitch angle, $\\mu_{0} \\sim \\mu_{\\rm p}$. When $\\mu_{\\rm p} = \\mart{1-M^{-1}}$, equation (\\ref{eq:rg_disk}) is \\begin{eqnarray} \\beta \\tanh (k r_{\\rm g}) \\sim \\mart{\\left(1-M^{-1}\\right)\\left\\{1-\\left(\\frac{r_{\\rm g}}{r_{\\rm f}}\\right)^{2}\\right\\}},\\label{eq:rg_disk_rewrite} \\end{eqnarray} which gives $r_{\\rm g}$ of an intermediate value between 0 and $r_{\\rm f}$. Therefore the strong radiation comes from the intermediate position between the loop top and footpoint. This explains the intensity distribution in the middle left panel of Figure \\ref{fig:result}.  In the side view, $\\theta = 90^{\\circ}$ is independent of $r$. Then the equation $\\mu = \\beta \\cos \\theta = 0$ results in \\begin{eqnarray} r_{\\rm g} = \\frac{\\mu_{0}}{\\mart{(1-\\mu_{0}^{2})(M-1)}} r_{\\rm f},\\label{eq:rg_limb} \\end{eqnarray} which again gives $r_{\\rm g}$ of an intermediate value between 0 and $r_{\\rm f}$ as long as $\\mu_{0} < \\mu_{\\rm c} = \\mart{1-M^{-1}}$. When $\\mu_{0} \\sim \\mu_{\\rm p} = \\mart{1-M^{-1}}$, this results in $r_{\\rm g} \\sim r_{\\rm f}$. Therefore the strong radiation comes from the footpoint. This explains the intensity distribution in the bottom left panel of Figure \\ref{fig:result}.        Based on these discussions, we conclude that electrons injected into the loop with an almost unique pitch angle do not yield the propagating feature of the radiation source along the loop.    \\subsection{Broad-Angle Injection}\\label{sec:broad-band-case} Next, we consider that initially injected electrons have an isotropic pitch-angle distribution: $\\phi(\\mu) = {\\rm constant}$. Electrons with a small initial pitch angle can reach the footpoint while those with a large initial pitch angle are confined to a narrow region around the loop top.  The top right panel of Figure \\ref{fig:result} shows the time variation of the electron number distribution along the loop. The solid, dashed, and dash-dotted lines show the trajectories of electrons with initial pitch-angle cosine of respective $\\pm 0.61$, $\\pm 0.4$, and $\\pm 0.2$. As expected, electrons are broadly distributed in the loop compared with the narrow-angle injection case (top left panel).  The middle and bottom right panels of Figure \\ref{fig:result} show the time variation of the 17 GHz {\\gyros} intensity distribution along the loop in the top- and side-view cases, respectively. The intensity distribution is quite different from that in the narrow-angle injection case (middle and bottom left panels). The strong radiation source is broadly distributed along the loop. We can see the propagating feature of the strong radiation source from the loop top to the footpoint in both the middle and bottom right panels, similar to the observation of Y2002. For example, rapidly propagating features are found during 0.3 - 0.5 s and 0.9 - 1.2 s in the bottom right panel. These features do not result from the motion of electrons with a specific initial pitch angle but from the motion of an ensemble of electrons with different initial pitch angles.  We discuss the apparent motion of the radiation source by using the analytic solution in {\\S} \\ref{sec:analytic-treatment}. For mathematical simplicity, we consider the side-view case. As mentioned in {\\S}~\\ref{sec:narrow-band-case}, the {\\gyros} radiation is primarily produced by the electrons with $\\mu \\sim 0$ in the side view. It is found from equations (\\ref{eq:sol_mu}) and (\\ref{eq:omega}) that the pitch-angle cosine of an electron with the initial pitch-angle cosine $\\mu_{0}$ becomes zero at \\begin{eqnarray} t_{\\rm g} = \\frac{(m/2) \\pi}{\\omega} = \\frac{(m/2)\\pi}{\\mart{(1-\\mu_{0}^{2})(M-1)}} \\frac{r_{\\rm f}}{v}, \\label{eq:t_g} \\end{eqnarray} where $m$ is an odd integer. This equation shows that when $\\mu_{0}$ is smaller (i.e., an initial pitch angle is larger) $t_{\\rm g}$ is smaller. At $t = t_{\\rm g}$, the electron position is given by equation (\\ref{eq:sol_r}) with $\\sin(\\omega t_{\\rm g}) = 1$. Using this, $\\mu_{0}$ in equation (\\ref{eq:t_g}) can be removed and it is rewritten as \\begin{eqnarray} t_{\\rm g} = \\frac{m \\pi}{2} \\mart{1 + \\frac{1}{M-1}\\left(\\frac{r_{\\rm f}}{r_{\\rm g}}\\right)^2} \\frac{r_{\\rm g}}{v}. \\label{eq:ana_solution} \\end{eqnarray} This equation shows that when $r_{\\rm g}$ is smaller $t_{\\rm g}$ is smaller. Therefore, $r_{\\rm g}$ is smaller when $\\mu_{0}$ is smaller. Electrons produce strong radiation at different time and position, depending on their initial pitch angle. By this equation, one can trace the position of the strong {\\gyros}-emitting electron in $(r,t)$ space. Figure \\ref{fig:result_zoom} is the zoomed image of the bottom right panel of Figure \\ref{fig:result}, to compare the analytic solution (eq. (\\ref{eq:ana_solution})) with the calculation result. The thick black line shows the analytic solution with $m=1$ and $v=c$. As can be seen, the analytic solution well explains the propagating feature of the radiation source.  Based on these discussions, we interpret the microwave propagating feature as follows. Electrons with a larger (smaller) initial pitch angle emit microwaves earlier (later) at the position closer to (farther from) the loop top, as seen in the peak of the dash-dotted (solid) line in Figure \\ref{fig:result_zoom}. This difference, due to the difference of the initial pitch angle of the parent electrons, appears as the propagating motion of the strong microwave source along the loop from the loop top to the footpoint.  The propagation of the radiation source seen in Figure \\ref{fig:result} is periodic. This is because the initially injected electrons remain to oscillate in the loop $(m=1,3,5,\\dots)$. Actual flares, on the other hand, show the light curve with a duration (tens of second) longer than the electron loop transit time scale ($\\lsim 1$ s), indicating that electrons are continuously injected into the loop \\cite[e.g.,][]{2008ApJ...673..598M}. The periodic propagation seen in our calculation is probably occulted by the emissions from subsequently injected electrons, except the first one. The observable propagating motion is the first line $(m=1)$. The time of electrons reaching the footpoint and the propagation speed of the radiation source are calculated from equation (\\ref{eq:ana_solution}) with $r_{\\rm g} = r_{\\rm f}$, \\begin{eqnarray} t_{\\rm end} = \\frac{\\pi}{2}\\mart{\\frac{M}{M-1}} \\frac{r_{\\rm f}}{v},\\; v_{\\rm prop} \\simeq \\frac{r_{\\rm f}}{t_{\\rm end}} = \\frac{2}{\\pi}\\mart{\\frac{M-1}{M}}v. \\label{eq:vprop} \\end{eqnarray} { The propagation speed of the radiation source is determined from the electron velocity and the magnetic mirror ratio. When the magnetic field strength at the emission site is on the order of 100 Gauss, effective energy of electrons producing $\\gsim 10 \\;{\\rm GHz}$ {\\gyros} radiation is $\\gsim 1 $MeV \\citep{1999spro.proc..211B}, hence their velocity is nearly $c$. The propagation speed, $v_{\\rm prop} \\simeq 0.4c$ with $M=1.6$, is almost independent of the observation frequency, as long as parent electrons are relativistic. This frequency-independent microwave propagating feature is found not only at 17 GHz but also at higher frequencies up to 34 GHz in our calculation. }   We performed the calculation for two ideal cases of the view (top and side) because the {\\gyros} radiation depends on the viewing angle. The apparent propagating motion seen in the top-view case (middle right panel of Fig. \\ref{fig:result}) looks different from that in the side-view case (bottom right panel), but can qualitatively be understood in the same manner as the side view. Difference of the view, which depends on the magnetic field configuration and location of the loop, does not influence our interpretation on the microwave propagating feature.   It is supposed that a specific electron accounts for the propagating feature of the microwave source. To produce the strong radiation, the pitch angle $\\alpha$ of the electron should be close to $\\theta$ at any time and position throughout its propagation along the loop. This is satisfied when (1) both $\\alpha$ and $\\theta$ do not change, or (2) $\\theta$ changes in the same manner as $\\alpha$. The former case is realized if the field strength and the viewing angle do not change along the loop. The latter case may be realized if the loop is on the limb with east-west orientation, because $\\theta$ as well as $\\alpha$ increase toward the footpoint in such configuration. These cases are, however, limited ones. Our interpretation holds in any probable field configuration and loop location.       We presented analytic and numerical treatments of the electron motion along a magnetic loop, and consider the rapidly propagating feature of the non-thermal microwave source reported by Y2002. We studied this issue by describing the electron distribution with the {\\FP} equation and by calculating the {\\gyros} radiation.  We first assumed that electrons injected into the loop have an almost unique initial pitch angle, following the interpretation made by Y2002. These electrons do not yield the propagating feature of the radiation source along the loop. This does not support the interpretation of Y2002.  We next assumed that electrons injected into the loop have an isotropic pitch-angle distribution. In this case, the intensity distribution shows the apparent motion of the strong radiation source from the loop top to the footpoint, similar to the observation. This feature is interpreted as the motion of an ensemble of electrons (not a specific electron), which have different time and position to produce strong radiation due to the difference of their initial pitch angle. To show the microwave propagating feature, the injected electrons should be broadly distributed in pitch-angle space.  We discuss a probable physical process of electrons in the 1999 August 28 flare. It is thought that this flare was triggered by the interaction between the microwave-propagating loop and a compact loop. Such configuration is suggested by \\cite{1999PASJ...51..483H}, called ``double-loop flare''. \\cite{1999PASJ...51..483H} concluded that in the double-loop flare electrons are accelerated at the region where two loops interact. Based on this model, Y2002 interpreted that the acceleration site in the 1999 August 28 flare is where the propagation of the non-thermal microwave source starts, that is, the injection region. We conclude that the electrons in the 1999 August 28 flare started to propagate along the loop just after being isotropically accelerated at the site.  In {\\S} \\ref{sec:narrow-band-case} we discussed the calculation for the narrow-angle injection case with $\\mu_{\\rm p} = 0.612$ only. The discussion can be applied to electrons with different pitch-angle distributions. When the electron pitch-angle distribution is concentrated perpendicular to the magnetic field, $\\mu_{\\rm p} = 0$. In this case, it is found from equations (\\ref{eq:rg_disk}) or (\\ref{eq:rg_limb}) with $\\mu_0 \\sim 0$ that the strong radiation source is confined to the loop top. When the electron pitch-angle distribution is concentrated parallel to the field, on the other hand, $\\mu_{\\rm p} = 1$. In this case the solution of equations (\\ref{eq:rg_disk}) or (\\ref{eq:rg_limb}) with $\\mu_0 \\sim 1$ is $r_{\\rm g} = \\infty$, meaning that there is no position within the loop to satisfy the condition for strong radiation. Neither of these cases results in the microwave propagation.  {We calculated the {\\gyros} intensity from the approximation of \\cite{1981ApJ...251..727P}, which is valid when ambient plasma density at the emission site is so low that the absorptions and the Razin suppression are not important. We have to use a more exact formula of the {\\gyros} intensity \\citep[e.g.,][]{1969ApJ...158..753R,2003ApJ...587..823F} in case that microwaves are produced at the high-density site in which these effects are important.}  {As well as the intensity, one can utilize the degree of circular polarization of the {\\gyros} radiation for diagnostics of the pitch-angle distribution of parent electrons \\citep[e.g.,][]{2008ApJ...677.1367A}. \\cite{2003ApJ...587..823F} found from their numerical calculations that the degree of polarization increases when the electron pitch-angle distribution is more anisotropic and/or the line of sight is more parallel to rather than perpendicular to the magnetic field line. Propagating electrons in a converging magnetic loop have a more anisotropic distribution at the footpoint than the loop top. Furthermore, it is expected for the disk flares (top view) that the magnetic field line at the footpoint is close to parallel to the line of sight. Therefore, {\\gyros} emission is more likely to be polarized at the footpoint than the top of the electron-propagating loop. The 1999 August 28 flare certainly showed that the degree of polarization increases toward the footpoints from the loop top (see Fig. 2 in Y2002), agrees with the above statement.}  The propagating feature of the microwave source gives us great opportunities to study the {\\gyros} radiation mechanism and the electron transport, and to constrain the pitch-angle distribution of the injected electrons which is crucially important to understand the electron acceleration mechanism. As is evident from many hard X-ray observations showing footpoint sources \\citep[e.g.,][]{1994PhDT.......335S,1999spro.proc..321S}, propagation of electrons and the microwave source along the loop commonly occurs in the flare. However, detection of such phenomena is quite difficult. The 1999 August 28 flare is the unique event in that NoRH detected the microwave propagating feature during its observational period until end of 2004 since its operation start in 1992 June \\citep{Shimojo_priv}. This event was an extremely well resolved one in both space and time by NoRH.  For further study, improvement of radio observatories is important. Observations with high temporal ($\\lsim 0.1 \\; {\\rm s}$) and spatial ($\\sim 1''$) resolutions more clearly resolve the microwave source because it propagates with a speed close to the speed of light along a loop with typical length $\\lsim 100'' \\sim 7 \\times 10^{4} \\; {\\rm km}$. Such observations should be implemented at frequencies greater than $\\sim 10 \\; {\\rm GHz}$ which correspond to the optically-thin regime of the microwave emission in typical solar flares, to obtain the spectral property of non-thermal electrons. Since the {\\gyros} radiation depends on the viewing angle with respect to the loop, the center-to-limb variation of the microwave distribution is studied to address the electron distribution in the loop. Such study has been carried out statistically by e.g., \\cite{1985PASJ...37..575K} and \\cite{2002SoPh..206..177S} with spatially-unresolved data. Statistical study on the center-to-limb variation of the spatial distribution of the microwave emission further gives constraints on the pitch-angle distribution of the injection electrons. Our numerical study will be of help to the future observational study for understanding the electron dynamics in solar flares."
    },
    "0806/0806.0702_arXiv.txt": {
        "abstract": " ",
        "introduction": "Despite their discovery now dating back almost a century, a full explanation for the R-stars eludes us. The division into early-R and later-R now seems to be a division into true-R and N (or J), respectively. We will assume this dichotomy in what follows and pursue an explanation for the early-R stars. The main features to be explained are 1) they are Carbon stars. I.e. they have atmospheric \\textit{n}(C) $>$ \\textit{n}(O). 2) Their spectral type is otherwise K. 3) They are enhanced in $^{12}$C, $^{13}$C, $^{14}$N, but seem to have solar [Fe/H], oxygen, and s-process abundances. 4) Their luminosity (about $100\\,\\Lo$) identifies them as clump giants, that is, low mass stars burning He in their cores; and most peculiarly of all 5) long term studies by McClure~(2007) failed to find any early R-stars in binary systems. Both the luminosity and the solar s-process abundances imply the R~stars have not reached the thermally pulsing AGB phase. In contrast are the N-stars, rich in s-process elements and with luminosities in excess of $2000\\,\\Lo$, leading to their interpretation as AGB stars having undergone third dredge up. The most thorough investigation of the composition of the R-stars was that of Dominy (1984).   The fact that R-stars are observed only as single stars leads to the notion, initially counter-intuitive, that they must all have originated as binaries. The argument is that a single star would not evolve any differently to a widely separated binary, so if R-stars are not found in binaries at all then they cannot exist as single stars. Hence they must be exclusively binary in origin and their current singularity is assumed to be due to a merger event. Given that approximately 20\\% of normal late-type giants are binaries, and that none of these stars are observed in binaries, we assume that every R-star is the product of a coalescence.   In normal low-mass single-star evolution, neutrino losses at high density cause a temperature inversion in the degenerate core of stars ascending the red-giant branch. Eventually the triple-$\\alpha$ reaction ignites at the point where the temperature peaks, which is no longer at the centre. A strong flash occurs, perhaps leading to $10^9\\,\\Lo$ from He-burning. A convective region develops and extends from the off-centre temperature maximum almost all the way to the H-rich envelope. It seems that contact is not made between these two convective regions (except for the case of very low [Fe/H]: see Fujimoto et al. 1990, 2000; Hollowell et al. 1990; Schlattl et al. 2001, 2002; Picardi et al. 2004; Komiya et al. 2007). After the flash dies down, there is a second flash, somewhat closer to the centre but substantially less energetic. This repeats a few times until the flash moves to the centre, and then central He burning is initiated. The energy released from the explosive He-burning has effectively lifted the degeneracy of the core and enables it to now burn He quiescently.   The first attempt at an explanation for the R-stars was made by Paczy\\'{n}ski and Tremaine (1977). They showed that, if the core-flash could be ignited sufficiently far from the centre of the star, that is, at a much larger core mass than normal, then a dredge-up episode follows the flash and carbon is dredged to the stellar surface.  This would explain the observed $n({\\rm C}) > n({\\rm O})$ in the R-stars: which are thought to be core He burning stars, and thus would be the progeny of this unusual core-flash. It remained to explain why only a small fraction of core flashes produced such dredge-up or, alternatively, why only a small fraction of core flashes begin at much larger core mass than normal.   This model was the preferred explanation for R-stars until the discovery that they are all single stars. A merged binary model was the basis for a recent study by Izzard et al. (2007) to explore merger scenarios using binary star population synthesis. They identified possible formation channels that lead to an R-star outcome. The most promising scenario was the merging of  a He white dwarf and a first-ascent red-giant. Typical He white dwarf masses are about $0.15 - 0.2\\,M_{\\odot}$, and the red-giant mass is around $1 - 2\\,\\Mo$. The merger is hypothesised to lead to a more rapidly rotating core than normal which then supports the core more than in the normal case. The core flash is hence delayed and ignites at a larger core mass, generating dredge-up in the manner found by Paczy\\'{n}\u0301ski and Tremaine (1977).   In this paper we try to take the next step in investigating this model, by looking at some basic nucleosynthetic constraints. We assume that the merger event has already occurred. Furthermore we assume that the star has returned to hydrostatic equilibrium, which allows us to use a hydrostatic stellar evolution code to model the evolution and nucleosynthesis. We try to force a core-flash event that is followed by dredge-up of carbon, and see if the resulting abundances are consistent with those observed in R-stars. ",
        "conclusions": "The R-stars continue to resist theorists' attempts to determine their origin. The binary merger hypothesis seems to be the best candidate at present, but direct calculations of this stage are unavailable and we are forced to make small steps toward validating, or otherwise, this qualitative model. In this paper we have simulated the events that would follow a late ignition of the core flash. If this ignition occurs further from the centre than is normal, then we confirm that dredge-up of carbon may result. Our calculations show that a substantial fraction of this carbon must then be exposed to burning via the CN cycles (and possibly ON). The observed low C isotope ratio remains a problem for the calculations shown here: the observed value indicates that essentially all of the added material has been burned to equilibrium via CN (and possibly ON) cycling. Only in that case can we match the observed $^{12}$C/$^{13}$C ratio. But then we burn too much C into N, overproducing N and destroying so much C that the star is no longer a carbon star.  Further advances in understanding the R-stars may require fully 3D hydrodynamical calculations of the merger event. Such work may be possible soon using the {\\it Djehuty\\/} code (e.g. Dearborn, Lattanzio and Eggleton, 2006).         \\bigskip\\bigskip\\bigskip\\bigskip"
    },
    "0806/0806.1638_arXiv.txt": {
        "abstract": "We discuss penumbral fine structure in a small part of a pore, observed with the CRISP imaging spectropolarimeter at the Swedish 1-m Solar Telescope (SST), close to its diffraction limit of 0\\farcs16. Milne--Eddington inversions applied to these Stokes data reveal large variations of field strength and inclination angle over dark-cored penumbral intrusions and a dark-cored light bridge. The mid-outer part of this penumbra structure shows $\\sim$0\\farcs3 wide spines, separated by $\\sim$1\\farcs6 (1200~km) and associated with 30\\degr{} inclination variations.  Between these spines, there are no small-scale magnetic structures that easily can be be identified with individual flux tubes. A structure with nearly 10\\degr{} more vertical and \\emph{weaker} magnetic field is seen midways between two spines. This structure is co-spatial with the brightest penumbral filament, possibly indicating the location of a convective upflow from below. ",
        "introduction": "The discovery of dark cores in sunspot penumbral filaments (Scharmer et al.\\@ 2002) suggests that the basic elements of penumbral fine structures are observable. The nature of individual dark cores was investigated with high-resolution multi-line spectra, but without polarization information, by Bellot Rubio et al.\\@ (2005). They concluded that the cores are associated with weaker magnetic field strength, by 100--300~G\\@. Langhans et al.\\@ (2007) measured the azimuthal variation of circular polarization signal in the 6302~{\\AA} \\ion{Fe}{1} line of regular spots at different heliocentric distances. They found that dark cores are associated with a \\emph{strongly} reduced field strength and a magnetic field that is more horizontal than for their lateral brightenings by about 10\\degr{}--15\\degr{}. Analysis of spectropolarimetric data from the Japanese satellite Hinode also shows lower field strength in the dark cores, but only by 100--150~G, and small inclination changes of about 4\\degr{} (Bellot Rubio et al.\\@ 2007).  We describe the first spectropolarimetric observations with CRISP, an imaging spectropolarimeter built for the Swedish 1-m Solar Telescope (SST). The large, unobscured aperture of the SST corresponds to a diffraction-limited resolution that is twice as high as that of Hinode. We present CRISP observations of penumbral fine structure made at a spatial resolution close to the SST diffraction limit of 0\\farcs16. Using Milne--Eddington (ME) inversions applied to these data, we discuss spatial variations of the magnetic field and line-of-sight (LOS) velocities for penumbral structure seen over parts of a large pore. ",
        "conclusions": "We have presented ME inversions based on data obtained with the CRISP imaging spectropolarimeter, used with the 1-m SST. The spatial resolution of this Stokes data represents a break-through in ground-based spectropolarimetry and a major improvement also as compared to recent Hinode data. The large variations in field strength and inclination angle inferred from the present data can to a large extent can be explained with the high spatial resolution of the SST/CRISP data.  The penumbra observed is partial, covering only a small part of the pore observed. We have analyzed dark-cored filamentary structures intruding into this pore and a lightbridge-like structure, detached from the surrounding photosphere. The three dark-cored structures analyzed are associated with strongly reduced field strength (around 50\\% relative to their surroundings), and a significantly more horizontal magnetic field (by 15\\degr--50\\degr) than outside the dark cores. The variations in field strength and inclination across the dark-cored filament are consistent with analysis of earlier SST magnetogram (Stokes V) data (Langhans et al. 2007). Even with the high spatial resolution of that and the present SST data, the magnetic field is found to be far from horizontal above the penumbral dark cores.  The inclination map shows a pronounced `spine' structure (Lites et al.\\@ 1993) with a magnetic field that is locally more vertical by $\\sim 30\\degr$ and that, except in the outer penumbra, is locally stronger by about 150~G\\@. Within these spines, the LOS velocities are strongly reduced. The spines seen in Fig.\\@~\\ref{context} are separated by about 1200~km. Such widely separated spines were evident also in earlier SST magnetogram data, discussed by Scharmer et al. (2007). We do not find any evidence for horizontal flux tubes with diameters in the range 100-250~km modeled in numerous papers (e.g., Borrero et al. 2007; Tritschler et al. 2007; Ruiz Cobo \\& Bellot Rubio 2008), in the outer penumbra. Midways between two of the spines, a faint spine-like structure is seen in the inclination map. The locally weaker and more \\emph{vertical} magnetic field of this structure is in contradiction with an interpretation in terms of a horizontal flux tube. The region with weaker field strength coincides with the brightest penumbral filament seen in the continuum. A possible interpretation is that the faint spine-like structure is related to a weak convective upflow, making the magnetic field overlying that upflow locally more vertical. This interpretation must be regarded as speculative in view of the small LOS velocities measured. Future CRISP observations of larger sunspots are likely to clarify this and also to provide more critical constraints on models."
    },
    "0806/0806.4537_arXiv.txt": {
        "abstract": "{We derive the star formation histories of early-type galaxies at $z\\simeq1.2$ in both low and high density environments. To this purpose, we compare the co-added spectroscopic and 8-9 band photometric data of 43 mass selected early-type galaxies in the massive cluster RDCS J1252.9-2927 and the GOODS/CDF-S field with a large grid of composite stellar population models based on the Bruzual \\& Charlot templates. We find that the cluster early-type galaxies formed the bulk of their stars approximately 0.5 Gyr earlier than early-types in the field, whereas field early-types presumably finish forming their stellar content on a longer time scale. Such a difference is particularly evident at masses $\\lesssim 10^{11} M_\\odot$, whereas it becomes negligible for the most massive galaxies. While our differential analysis of the stellar population parameters of cluster and field galaxies in the same mass range convincingly shows distinct star formation histories, the absolute age difference remains model dependent. Using the star formation histories that best fit the SEDs of the red sequence galaxies in RDCS 1252.9-2927, we reconstruct the evolution of the cluster red sequence and find that it was established over 1 Gyr and is expected to dissolve by $z\\approx 2$.} ",
        "introduction": "Massive early-type galaxies are a good tracer of the early mass assembly in the Universe. The study of their spectrophotometric and morphological properties, especially at high redshift, over a range of environmental densities, can significantly constrain the different models of structure formation: the monolithic collapse (e.g. Eggen et al. \\cite{Eggen62}), in which early-type galaxies result from a single burst of star formation at high redshift, and the hierarchical model (e.g. Toomre \\cite{Toomre77}), where they form by the merger of low mass progenitors. This latter process is naturally expected in a $\\Lambda$CDM cosmology and predicts different formation histories whether a galaxy is in a low-density environment or the member of a cluster (e.g. De Lucia et al. \\cite{DeLucia06}). Indeed, the analysis of the fossil record via line-strength indices shows that massive early-type galaxies in local high-density environments are at least 1.5 Gyr older than their counterparts in low-density regions (Thomas et al. \\cite{Thomas05}, S\\'{a}nchez-Bl\\'{a}zquez et al. \\cite{Sanchez-Blazquez06}, Clemens et al. \\cite{Clemens06}), while from the mass-to-light ratio of cluster and field galaxies up to $z\\sim1$, van Dokkum \\& van der Marel (\\cite{vanDokkum07}) find a lower value of $\\sim$0.4 Gyr. On the other hand, early-type galaxies appear to have formed at an early time and been in place at $z \\simeq 2$ (Bernardi et al. \\cite{Bernardi98}, van Dokkum et al. \\cite{vanDokkum01}), with little star formation happening ever since.  This suggests that studying the star formation history of early-type galaxies at $z > 1$ allows one to place stronger constraints on structure formation models than at low redshift, where any difference has been smoothed out by billions of years of mostly passive evolution and occasional mergers. In such a case, much sparser data are available, as few massive galaxy clusters have been observed so far at $z>1$. One of these, RDCS J1252.9-2927 (Rosati et al. \\cite{Rosati04}) at $z=1.24$, has had an extensive multi-wavelength spectroscopic coverage. In this paper, we use the spectrophotometric data of early-type galaxies in RDCS J1252.9-2927 to reconstruct their general star formation history and compare it with early-type galaxies from the GOODS/CDF-S at similar redshift. An independent analysis of the same data sets, also including morphological properties and far-UV rest-frame photometry is presented in Rettura et al. (\\cite{Rettura08}). This paper is structured as follows. In Sect. 2, we describe our data and sample selection; in Sect. 3 we present our methodology and in Sect. 4 we describe and discuss the results of our analysis.  We assume a $\\Lambda$CDM cosmology with $\\Omega_m=0.3$, $\\Omega_{\\Lambda}=0.7$ and $H_0=70$ km s$^{-1}$ Mpc$^{-1}$. All magnitudes in this paper are given in the AB system (Oke \\cite{Oke74}), unless stated otherwise. ",
        "conclusions": "We compared the underlying stellar population properties of $z \\simeq 1.2$ early-type galaxies in the very high density environment of the massive cluster RDCS 1252 with those in the GOODS field.  We derived star formation histories by fitting both the spectroscopic and the broad band photometric data with a large grid of stellar population synthesis models.  To this purpose, we select 43 early-type galaxies (21 in the field and 22 in the cluster) all with stellar masses greater than $M_{lim}=5\\times10^{10}M_{\\odot}$, $i_{775}-z_{850} \\geq 0.8$ and the absence of the [OII]$\\lambda$3727 line in the spectra. These criteria naturally select all galaxies on the red sequence of RDCS 1252. For each sample, we use the co-added spectrophotometric data of the galaxies and compare them with BC03 models of exponentially declining star formation rates with an additional burst.  We find a small although significant difference in the star formation histories of the cluster and field populations, suggesting that the cluster galaxies form the bulk of their stars $\\sim\\!0.5$ Gyr earlier than their counterparts in the field, with massive early-type galaxies having already finished forming stars at $z>1.5$ in both environments. This difference is particularly evident at masses $\\lesssim 10^{11} M_\\odot$, which are characterized by a longer star formation time scale, resulting in a final formation time delayed by $\\sim$1 Gyr, whereas it becomes negligible for the most massive galaxies. While our differential analysis of the stellar population parameters of cluster and field galaxies in the same mass range convincingly shows distinct star formation histories, the absolute age difference remains model dependent. In an accompanying paper, Rettura et al. (\\cite{Rettura08}) have analyzed the rest-frame far-UV flux of the same sample of early-type galaxies and found that field galaxies are at least 0.5 mag brighter than the RDCS 1252 galaxies in the same mass range. Our best fit models consistently predict this magnitude difference, which is indicative of the longer star formation time scale of the field galaxies. We have verified that such a difference in derived star formation histories in the two environments cannot be ascribed to incompleteness of the mass selected samples, which would tend to rather increase such an effect.  We also used extensive Monte-Carlo simulations to identify possible biases in the model fit, and discussed the effect of inherent degeneracies such as metallicity and dust. We note that independent studies of massive early-type galaxies based on the measurement of the mass-to-light ratios of massive early-type galaxies in high- and low-density environments (Treu et al. \\cite{Treu05}, van der Wel et al. \\cite{vanderWel05}, van Dokkum \\& van der Marel \\cite{vanDokkum07}) have reached similar conclusions.  We also used the best fit star formation histories from the 9-band SEDs of the red sequence galaxies in RDCS 1252 to predict that a tight ($\\Delta(U-B)_z=0.05$ mag) red sequence at $z=1.2$ is established over approximately 1 Gyr and dissolves by $z\\approx 1.9$. This suggests that for massive clusters, which have long reached virialization by redshift 1.2, we do not expect a significant red sequence at $z>2$, i.e. a $(U-B)_z$ color scatter well above 0.1 mag.  These observations and analysis can be used to provide significant constraints on galaxy evolution models in a hierarchical scenario (e.g. Menci et al. \\cite{Menci08}), which predicts the evolution in high-density environments to be accelerated compared to the field."
    },
    "0806/0806.3856_arXiv.txt": {
        "abstract": "We present three dimensional hydrodynamical simulations aimed at studying the dynamical and chemical evolution of the interstellar medium (ISM) in isolated dwarf spheroidal galaxies (dSphs). This evolution is driven by the explosion of Type II and Type Ia supernovae, whose different contribution on both the dynamics and chemical enrichment is taken into account.  Radiative losses are effective in radiating away the huge amount of energy released by SNe explosions, and the dSph is able to retain most of the gas allowing a long period ($\\ge 2-3$ Gyr) of star formation, as usually observed in this kind of galaxies.  We are able to reproduce the stellar metallicity distribution function (MDF) as well as the peculiar chemical properties of strongly O-depleted stars observed in several dSphs. The model also naturally predicts two different stellar populations, with an anti-correlation between [Fe/H] and velocity dispersion, similarly to what observed in the Sculptor and Fornax dSphs. These results derive from the inhomogeneous pollution of the SNe Ia, a distinctive characteristic of our model. We also applied the model to the peculiar globular cluster (GC) $\\omega$ Cen in the hypothesis that it is the remnant of a formerly larger stellar system, possibly a dSph. ",
        "introduction": "Due to their proximity, the galaxies of the Local Group (see Mateo 1995 and Geisler et al. 2007 for a review) offer an unique opportunity to study in details their structural, dynamical and chemical properties and to test different theories of galaxy formation.  Owing to their low metallicity and lack of neutral hydrogen, it was initially believed that dSphs are relatively simple objects whose ISM is completely removed by SN II explosions after a very short intense star formation period (e.g. Dekel \\& Silk 1986).  Doubts about this picture come from the high resolution spectroscopy of several dSphs showing a wide range in metallicity. For istance, Shetrone et al. (2001) have observed stars in Draco and Ursa Minor with values of [Fe/H] in the range $-3\\leq$[Fe/H]$\\leq -1.5$. The same authors also found that their observed dSphs have [$\\alpha$/Fe] abundances that are $0.2$ dex lower than those of Galactic halo field stars at the same metallicity. This suggests that the bulk of the stars in these systems formed in gas self polluted by SNe II as well as SNe Ia and that the star formation (SF) must continue over a relatively long timescale in order to allow a sufficient production of iron by SNe Ia (Ikuta \\& Arimoto 2002, Lanfranchi \\& Matteucci 2004, Marcolini et al. 2006; but see also Recchi et al. 2007, Salvadori et al. 2008).  A complex star formation history (SFH) is further suggested by several facts: $i$) isolated low mass dSphs such as Phoenix (Young et al. 2007) and Leo T (de Jong et al. 2008) were able to form stars up to 100 Myr ago; $ii$) the SFHs of dwarf galaxies are strongly dependent on their local environment, the fraction of passively evolving galaxies dropping from $\\sim$70\\% in dense environments, to zero in the rarefied field (Haines et al. 2007 ); $iii$) dwarf ellipticals and dSphs cluster around the dominant spirals galaxies, while gas rich star forming dwarf Irregulars are found at larger distances (van den Bergh 1994). These points highlight the role of the environment (tidal interaction/ram pressure stripping), and disfavors a scenario in which the evolution is due uniquely to internal processes.  \\begin{figure}[] \\begin{center} \\includegraphics[width=2.6in]{draco.ps} \\includegraphics[width=2.6in]{zstellenam.ps} \\caption{Left: distribution of the logarithm of the density in three orthogonal planes at $\\sim$ 400 Myr after the beginning of the simulation. Note the inner bubble carved by the SNe explosions and the shocks propagating outward. Right: [Fe/H] distribution of the models resembling the Draco dSph (left panels) and the $\\omega$ Cen globular cluster (right panels) together with the corresponding [$\\alpha$/Fe]-[Fe/H] diagrams; the blue line and blue dots represent stars with [$\\alpha$/Fe]$\\le$0.2 (i.e. affected by SNe Ia inhomogenous pollution, see text). The lower panels represents the assumed SFH for the two models.} \\label{fig1} \\end{center} \\end{figure}  Here we briefly report some results obtained by Marcolini et al. (2006, 2007, 2008) with a model which turns to be consistent with many properties of the Draco dwarf and, with minimal assumptions, to the chemical properties of the peculiar system $\\omega$ Centauri (Marcolini et al. 2007), which is believed to be the remnant of an ancient dSph. ",
        "conclusions": ""
    },
    "0806/0806.1548_arXiv.txt": {
        "abstract": "We report the results of our third survey for formaldehyde (H$_2$CO) 6$\\,$cm maser emission in the Galaxy. Using the Very Large Array, we detected two new H$_2$CO maser sources (G23.01$-$0.41 and G25.83$-$0.18), thus increasing the sample of known H$_2$CO maser regions in the Galaxy to seven. We review the characteristics of the G23.01$-$0.41 and G25.83$-$0.18 star forming regions. The H$_2$CO masers in G23.01$-$0.41 and G25.83$-$0.18 share several properties with the other known H$_2$CO masers, in particular, emission from rich maser environments and close proximity to very young massive stellar objects. ",
        "introduction": "The first formaldehyde (H$_2$CO) 6$\\,$cm maser was discovered in 1974 toward NGC$\\,$7538 (Downes \\& Wilson 1974; Forster et al. 1980); in the following $\\sim 28\\,$years, H$_2$CO masers were detected only toward two other regions in the Galaxy: Sgr B2 (Whiteoak \\& Gardner 1983) and G29.96$-$0.02 (Pratap et al. 1994). This low number of H$_2$CO maser regions was unexpected (e.g., Forster et al. 1985; Gardner et al. 1986) given the detection of many H$_2$CO maser spots in a single source (Sgr B2, Whiteoak \\& Gardner 1983; Mehringer et al. 1994), the widespread distribution of formaldehyde molecules as exemplified by Galactic H$_2$CO 6$\\,$cm absorption (e.g., Watson et al. 2003; Araya et al. 2002; Downes et al. 1980; Dieter 1973), and the apparently common astrophysical conditions needed for the maser excitation if the masers were pumped by radio continuum radiation (see Pratap et al. 1992).  The idea that H$_2$CO masers may be pumped by background radio continuum was initially proposed by Boland \\& de Jong (1981) to explain the maser in NGC$\\,$7538 IRS1. However, the low detection rate of new H$_2$CO masers in dedicated surveys (Mehringer et al. 1995; Forster et al. 1985), the non-detection of maser emission from the H$_2$CO 2$\\,$cm transition (e.g., Hoffman et al. 2003) and low emission measure (or non-detection) of radio continuum sources near several of the known H$_2$CO masers, indicate that the pumping mechanism in most cases cannot be due to radio continuum excitation (e.g., see Araya et al. 2007c for the case of IRAS$\\,$18566+0408). The low detection rate of H$_2$CO masers led several authors to speculate that formaldehyde masers are rare because specific and/or short-lived physical conditions may be needed for the excitation (e.g., Forster et al. 1985; Mehringer et al. 1995; Araya et al. 2007a).  In an effort to understand the H$_2$CO maser phenomenon and its place during the formation of massive stars, we conducted two surveys for H$_2$CO masers using Arecibo, the Green Bank Telescope, and the VLA\\footnote{The Very Large Array (VLA) is operated by the National Radio Astronomy Observatory (NRAO), a facility of the National Science Foundation operated under cooperative agreement by Associated Universities, Inc.} that resulted in the detection of two new Galactic H$_2$CO maser regions (IRAS$\\,$18566+0408, with emission first detected by Araya et al. 2004 and confirmed to be a maser by Araya et al 2005;  and G23.71-0.20 with details in Araya et al. 2006, from a survey later reported by Araya et al. 2007b). In this article we present the results of our third survey for H$_2$CO masers. ",
        "conclusions": "\\label{discussion}   \\subsection{G23.01$-$0.41}  \\subsubsection{H$_2$CO 6$\\,$cm Maser Emission}  H$_2$CO 6$\\,$cm maser emission was detected toward G23.01$-$0.41 with the VLA in the BnA $\\rightarrow$ B reconfiguration (Figure~1), and confirmed with VLA-D observations conducted approximately 8 months later (\\S2.2, Figure~2). The peak position, LSR velocity, and line width measurements are consistent in both runs (Table~3). If the VLA-D spectrum is smoothed to the channel width of the lower spectral resolution observations (VLA BnA$\\rightarrow$B, Table~3), the peak intensities are consistent within 2$\\sigma$, revealing no apparent variability of the line in a $\\sim 8$ month period.    As mentioned in $\\S3$, the H$_2$CO 6$\\,$cm maser in G23.01$-$0.41 was not spectrally resolved in the VLA BnA$\\rightarrow$B observations (line width $<$ 1.0\\kms, Figure~1). In the high spectral resolution observations (Figure~2) the line was barely resolved, and appears to be the superposition of two components. The separation of the two possible overlapping lines is $\\la$0.4\\kms.   \\subsubsection{Infrared Environment}  Far-IR emission is found at the location of the H$_2$CO maser as shown by IRAS data\\footnote{IRSA-IRAS Sky Survey Atlas, http://irsa.ipac.caltech.edu/Missions/iras.html}. However, the far-IR emission at the H$_2$CO maser location is overlapped with several nearby sources. MSX 21.4$\\mu$m data toward G23.01--0.41 is also affected by confusion due to blended emission with a source $\\sim 20$\\arcsec~east of the H$_2$CO maser position. At 70$\\mu$m, MIPSGAL Spitzer data show that the maser is coincident with a compact (FWHM $<$ 30\\arcsec) and strong (saturated) FIR source. Since the emission is saturated, a reliable measurement of the 70$\\,\\mu$m size and position cannot be obtained; nevertheless, the peak position appears to be within $\\sim 3$\\arcsec~of the H$_2$CO maser, i.e., certainly within the telescope angular resolution at 70$\\,\\mu$m ($\\sim 20$\\arcsec; Rieke et al. 2004). At the 24$\\,\\mu$m Spitzer MIPS band, the IR source coincident with the H$_2$CO maser is substantially less affected by saturation and the center position can be more reliably determined; the 24$\\,\\mu$m source and the H$_2$CO maser are less that 0.3\\arcsec~apart.  The IR source coincident with the H$_2$CO maser is detected in mid/near-IR observations from the Spitzer/IRAC GLIMPSE survey (Benjamin et al. 2003). Figure~3 (upper panel) shows a three color image (3.6$\\mu$m blue, 4.5$\\mu$m green, 8.0$\\mu$m red) of the G23.01$-$0.41 region. The Spitzer source at the H$_2$CO maser location shows strong 4.5$\\mu$m excess indicative of shocked gas (e.g., Smith et al. 2006). There is also extended 4.5$\\mu$m excess emission west of the H$_2$CO maser position (Figure~3, inset upper panel).  Testi et al. (1998) detected a near infrared (NIR) source toward G23.01$-$0.41, that was also detected by 2MASS. The 2MASS source is $\\sim 4$\\arcsec~displaced from the H$_2$CO maser position. Given this large angular offset, it is unlikely that the NIR source is responsible for the excitation of the masers in G23.01$-$0.41 ($\\S 4.1.4$).  \\vspace{-0.5cm}  \\subsubsection{Radio Continuum and Thermal Molecular Lines}   We detected no 6$\\,$cm radio continuum toward the maser position to a 5$\\sigma$ level of 2.4\\mjyb~($\\theta_{syn} = 1$\\asec$6\\times1$\\asec1). Using the VLA in the C configuration ($\\theta_{sys} \\sim 1$\\arcsec), Codella et al. (1997) report an upper limit of the 1.3$\\,$cm radio continuum of 0.4\\mjyb~(3$\\sigma$).  Based on our VLA-D observations, the nearest 6$\\,$cm radio continuum source is approximately 50\\arcsec~north of the maser position. The radio continuum source (a compact H{~\\small II} region that will be discussed in a future paper) has associated H$_2$CO absorption at approximately the same velocity as the H$_2$CO maser, indicating that the H$_2$CO maser belongs to an extended molecular cloud that shows clear evidence for the presence of more evolved massive stars. The LSR velocity of the H$_2$CO absorption (75.0\\kms) implies two possible kinematic distances: 4.8 or 10.8$\\,$kpc. Codella et al. (1997) as well as Caswell \\& Haynes (1983) preferred the far kinematic distance\\footnote{Caswell \\& Haynes (1983; see also Caswell et al. 1995b, Anglada et al. 1996, Testi et al. 1998, Forster \\& Caswell 1999) report a far kinematic distance of 12.8$\\,$kpc instead of our value of 10.8$\\,$kpc because they assumed a Sun -- Galactic Center distance of 10.0$\\,$kpc, whereas we assume 8.5$\\,$kpc (Brand \\& Blitz 1993).}; however the distance ambiguity has not been resolved (Forster \\& Caswell 1999). For example, Harju et al. (1998) list the near kinematic distance for the source (see also Scoville et al. 1987, Pestalozzi et al. 2005).  A number of thermal molecular lines have been detected toward G23.01$-$0.41. Caswell et al. (2000) detected (quasi-)thermal emission of the CH$_3$OH lines at 107.0$\\,$GHz\\footnote{A maser component was also found superimposed on the thermal profile of the CH$_3$OH 107.0$\\,$GHz line.} and 156.6$\\,$GHz (see also Slysh et al. 1999). Harju et al. (1998) detected broad (line width $>$ 20\\kms) SiO J=2--1 and J=3--2 emission with the SEST telescope. CO and $^{13}$CO emission was also detected in single dish surveys (Scoville et al. 1987; Jackson et al. 2006).  Interferometric (Nobeyama Millimeter Array and Plateau de Bure Interferometer) observations of $^{12}$CO, $^{13}$CO, and C$^{18}$O have been recently reported by Furuya et al. (2008) with an angular resolution $\\sim$6\\arcsec. Their $^{12}$CO spectrum shows high velocity gas, in particular prominent red-wing emission that traces a massive ($>50$\\Mo) molecular outflow, possibly located close to the plane of the sky.  High density molecular gas in the region is evident from detection of CS ($J=1-0$) and NH$_3$ (1,1) lines with the MIT Haystack 37$\\,$m telescope (Anglada et al. 1996). Single dish NH$_3$ (1,1), (2,2), and (3,3) observations were also conducted by Codella et al. (1997) with the Medicina 32$\\,$m telescope. Codella et al. (1997) derived a rotational temperature of 13$\\,$K.  High angular resolution NH$_3$ observations were conducted by Codella et al. (1997) with the VLA ($\\theta_{syn} \\sim 1$\\arcsec). The NH$_3$ (3,3) data show a compact ($<10$\\arcsec) molecular core that has a SE--NW elongation (Figure~3, lower panel). Codella et al. (1997) derived the following parameters for the ammonia core: deconvolved angular diameter = 3\\asec3, T$_k = 58\\,$K, M$_{\\mathrm{vir}} = 886$\\Mo, and $n_{\\mathrm{H_2}} = 6.9 \\times 10^6\\,$cm$^{-3}$. The H$_2$CO maser is coincident with the NH$_3$ (3,3) peak emission (Figure~3, lower panel).  Furuya et al. (2008) recently conducted a detailed multi-wavelength study of G23.01$-$0.41 at high angular resolution ($<$ 10\\arcsec) using CO (see above), HNCO, and CH$_3$CN. Their observations confirm the presence of a hot molecular core characterized by a CH$_3$CN rotation temperature of $\\sim 120\\,$K and a core mass of $\\sim 380$\\Mo. The velocity distribution of the CH$_3$CN gas is consistent with rotation of a molecular core oriented almost perpendicular to the outflow direction. The elongation and velocity gradient of the CH$_3$CN core is parallel to the elongation of the NH$_3$ emission shown in Figure~3.    \\subsubsection{Other Astrophysical Masers}     A variety of masers have been detected toward G23.01$-$0.41 in single dish surveys. Methanol masers have been found at 6.7$\\,$GHz (Menten 1991; Caswell et al. 1995b; Szymczak et al. 2002; see also catalogs by Pestalozzi et al. 2005 and Xu et al. 2003), 12$\\,$GHz (MacLeod et al. 1993; Caswell et al. 1995a; B{\\l}aszkiewicz \\& Kus 2004), 107.0$\\,$GHz (Caswell et al. 2000), 44$\\,$GHz (Slysh et al. 1994), and 95$\\,$GHz (Val'tts et al. 2000). H$_2$O 22$\\,$GHz and OH ($^2\\Pi_{3/2}~J=3/2$ ground state) masers have also been reported (Szymczak et al. 2005; Szymczak \\& G\\'erard 2004; Caswell \\& Haynes 1983).  High angular resolution observations of the CH$_3$OH 6.7$\\,$GHz masers have been conducted (Caswell, unpublished data; see Caswell et al. 2000); the position of the CH$_3$OH maser is coincident with that of the H$_2$CO maser within 1\\arcsec~rms. Forster \\& Caswell (1989, see also Forster \\& Caswell 1999) conducted VLA observations of H$_2$O 22$\\,$GHz ($\\theta_{syn} = 3$\\asec$9\\times1$\\asec8) and OH 1665$\\,$MHz ($\\theta_{syn} = 6$\\asec$5\\times1$\\asec0) masers in G23.01$-$0.41. The H$_2$CO maser is located within $2$\\arcsec~of the OH and H$_2$O maser positions (see Figure~3 lower panel).   \\subsubsection{The Nature of the G23.01$-$0.41 Star Forming Region}   The abundant available multi-wavelength data shows that G23.01$-$0.41 is an active site of massive star formation. The H$_2$CO maser is located at the peak of a molecular core traced by NH$_3$ (3,3) and CH$_3$CN. The hot molecular core also harbors a variety of molecular masers. Given the highly confused far IR environment, it is not possible to reliably measure the bolometric luminosity of the massive stellar object pinpointed by the H$_2$CO maser (Figure~3); however, based on the available mid and far IR data, the upper limit of the luminosity is $\\sim 10^6$\\Lo. As in the case of the infrared source associated with the H$_2$CO maser in IRAS$\\,18566+0408$ (Araya et al. 2007c), the H$_2$CO maser in G23.01$-$0.41 is found toward a source with 4.5$\\mu$m infrared excess, which likely indicates shocked gas in an outflow. Considering the presence of a hot molecular core at the location of the H$_2$CO maser, the detection of several other maser species, evidence for outflow and shocked gas based on CO, SiO and 4.5$\\mu$m infrared excess emission, absence of radio continuum, and evidence for rotation of the molecular core, the H$_2$CO maser appears to pinpoint the location of a very young massive stellar object in a evolutionary stage prior to the formation of a radio bright ultra-compact H{~\\small II} region. The central object may still be undergoing accretion; further observations should be made to clarify this point.    \\subsection{G25.83$-$0.18}   \\subsubsection{H$_2$CO 6$\\,$cm Maser Emission}   An the case of G23.01$-$0.41 ($\\S4.1$), maser emission in G25.83$-$0.18 was first detected in the VLA survey ($\\S2.1$) and then confirmed with higher spectral resolution observations (VLA-D, $\\S2.2$) $\\sim 8\\,$months later. The H$_2$CO maser shows a double peaked line profile (Figures~1 and 2). The separation between the two peak components is 1.5\\kms. No significant difference in the maser flux density was found between the two epochs.   As reported in Table~3, the line profile of the high spectral resolution observations is well fit by the superposition of three Gaussian profiles; two broad components (0.7 and 0.9\\kms~FWHM) and a narrow component (FWHM = 0.29\\kms). Whether the line profile is composed of only two non-Gaussian maser lines or three (or more) components is unclear. For example, the profile could be due to an asymmetric line (the blue-shifted component) overlapped with a Gaussian line (the red-shifted component, see Figure~2).   \\subsubsection{Infrared Environment}   The H$_2$CO maser is at the center of an infrared dark cloud as revealed by 8.0$\\mu$m \\linebreak Spitzer/GLIMPSE observations (see Figure~4). The H$_2$CO maser is located between a source with strong 4.5$\\mu$m excess emission (green in Figure~4, upper panel inset) and a source brighter at 8$\\mu$m (red in Figure~4). The peak of the 4.5$\\mu$m source is offset ($\\sim 3$\\arcsec) from the location of the H$_2$CO maser. No infrared emission was detected with MSX toward the position of the H$_2$CO maser; no 2MASS source was found coincident with the H$_2$CO maser either.  Data from Spitzer MIPSGAL at 70$\\mu$m reveal a strong ($>100\\,$Jy) far IR source whose peak is within $\\sim$1\\arcsec~from the position of the H$_2$CO maser. At low level, the 70$\\mu$m emission is extended and comprises a neighboring infrared source; however, the core emission is compact (FWHM$\\,<\\,$25\\arcsec), i.e., close to the theoretical telescope diffraction limit. The IR source coincident with the H$_2$CO maser is also detected in the 24$\\mu$m MIPS band.    \\subsubsection{Radio Continuum and Thermal Molecular Lines}  We detected no radio continuum toward the position of the H$_2$CO maser to a level of 2.8\\mjyb~(5$\\sigma$) in the VLA BnA $\\rightarrow$ B observations ($\\theta_{syn} = 1$\\asec$6\\times 1$\\asec$1$); the 5$\\sigma$ upper limit set by the VLA-D data is 5\\mjyb~(see Figure~4, upper panel). The nearest radio continuum source to the H$_2$CO maser is $\\sim$2\\arcmin~to the west (G25.80$-$0.16; Figure~4, upper panel) which was also detected in the continuum at 8.64 and 6.67$\\,$GHz by Walsh et al. (1998).   The H$_2$CO maser is coincident with a mm and sub-mm core. Walsh et al. (2003) detected compact 450$\\mu$m and 850$\\mu$m emission; the H$_2$CO maser is located within 4\\arcsec~of the peak position of the sub-mm source (the JCMT beam is approximately 8\\arcsec~and 15\\arcsec~at 450$\\mu$m and 850$\\mu$m, respectively). Hill et al. (2005) conducted SEST/SIMBA observations of G25.83$-$0.18 and detected 1.2$\\,$mm emission ($S_\\nu = 5.4\\,$Jy); the FWHM of the mm source is 60\\arcsec. Based on the mm detection, Hill et al. (2005) report a mass of $2.8 \\times 10^3$\\Mo~or $8.8 \\times 10^3$\\Mo~depending on the kinematic distance (see below). The H$_2$CO maser is coincident with the 1.2$\\,$mm peak within $\\sim 1$\\arcsec.   We detected H$_2$CO 6$\\,$cm absorption in the region (Figure~4, middle and lower panels, VLA-D observations). The main H$_2$CO absorption clump is coincident with the infrared dark cloud, and shows a shell-like brightness distribution (Figure~4, middle panel)\\footnote{The shell-like distribution of H$_2$CO absorption (Figure~4, middle panel) is not due to overlapping H$_2$CO maser emission and absorption at the central position; the measured velocity of the H$_2$CO maser is not in the velocity range used to obtain the H$_2$CO absorption image.}. The H$_2$CO maser is located at the projected spatial center of the shell and near the radial velocity edge of the H$_2$CO absorption line (Figure~4, lower panel). The shell brightness distribution could be due to: 1.) H$_2$CO in gas phase is less abundant in the inner regions of the shell, for example, due to chemical gradients, smaller total molecular density, or depletion (see Young et al. 2004 for the case of pre-protostellar cores); and/or 2.) the excitation conditions for H$_2$CO absorption are less favorable in the inner regions of the clump (e.g., Zhou et al. 1990 explained the detection of an H$_2$CO shell-like structure in B335 as a consequence of higher molecular density at the core that quenches the anomalous absorption). Observations of other H$_2$CO transitions as well as other mid and high density tracers are needed to fully investigate the nature of the shell-like H$_2$CO absorption source. Nevertheless, the recent high angular resolution detection of NH$_3$ at the position of the H$_2$CO depression (Longmore et al. 2007, see below) suggests that the H$_2$CO shell-like structure is due to high density ($>10^5\\,$cm$^{-3}$) molecular gas in the center of the molecular core that quenches H$_2$CO anomalous absorption.  We also detected H$_2$CO absorption close to the radio continuum source (compare Figure~4 upper and middle panels). The similar velocity of the H$_2$CO absorption gas associated with the continuum source and with the infrared dark cloud indicates that both are part of the same star forming complex. Assuming that the LSR velocity of the H$_2$CO absorption line traces the systemic velocity of the cloud, the two possible kinematic distances to this massive star forming region are 5.6 and 9.7$\\,$kpc. Since the H$_2$CO maser is associated with an infrared cloud seen in absorption against the mid-IR galactic background (Figure~4), it is likely that the region is located at the near kinematic position.  G25.83$-$0.18 has been detected in CH$_3$CN, HCO$^+$, and H$^{13}$CO$^+$ with the Mopra telescope by Purcell et al. (2006). Based on the CH$_3$CN data, they derived a rotation temperature of $\\sim 50\\,$K, and found evidence for infalling motion based on HCO$^+$ data. In addition, Longmore et al. (2007) detected an optically thick NH$_3$ core at the position of the H$_2$CO maser from ATCA observations. The line width of the NH$_3$ (1,1), (2,2), (4,4), and (5,5) transitions range between 6 and 30\\kms, indicating the possible presence of a molecular outflow partially traced by NH$_3$.    \\subsubsection{Other Astrophysical Masers}   Walsh et al. (1998) conducted ATCA observations of 6.7$\\,$GHz CH$_3$OH masers in \\linebreak G25.83$-$0.18; five maser components were found with LSR velocities between of 90.7 and 99.0\\kms. The H$_2$CO maser is located $\\sim 2$\\arcsec~south of the CH$_3$OH maser clump\\footnote{The absolute position accuracy of the CH$_3$OH masers is $\\sim 1$\\arcsec, Walsh et al. (1998).} (Figure~4, upper panel, inset).   Single dish observations of the Class II CH$_3$OH 6.7 and 12$\\,$GHz masers were conducted by B{\\l}aszkiewicz \\& Kus (2004) with the Toru\\'n telescope. The peak LSR velocity of the 6.7 and 12$\\,$GHz masers were 91.3 and 90.7\\kms, i.e., coincident in velocity with the H$_2$CO maser.  Ellingsen (2005) conducted observations with the Mopra Telescope of the 95.1 GHz Class I CH$_3$OH maser transition, and detected several maser lines (peak maser emission at 90.2\\kms) overlapped with a broad component. The broad line (4.3\\kms~line width) could be due to thermal emission given that its peak velocity (94.2\\kms) coincides with the H$_2$CO absorption peak velocity (see $\\S4.2.3$).   Other maser species have also been detected in single dish surveys. Szymczak \\& G\\'erard (2004) detected a OH 1667$\\,$MHz maser line with a peak LSR velocity of 92.9\\kms; their figure A.1 also shows detection of a possible high velocity ($\\sim 120$\\kms) OH 1612$\\,$MHz line. Szymczak et al. (2005) conducted H$_2$O 22$\\,$GHz observations with the 100$\\,$m Effelsberg telescope and detected several H$_2$O maser features within a velocity range of $\\sim 50$\\kms, centered at 94.7\\kms.    \\subsubsection{The Nature of the G25.83$-$0.18 Star Forming Region}  The coincidence of the H$_2$CO maser with other maser species (in particular with Class II CH$_3$OH masers), with an infrared dark cloud and molecular core, the absence of radio continuum emission, the detection of a mm, sub-mm and far-IR source, and the presence of 4.5$\\mu$m excess emission toward the center of the infrared dark cloud, imply that G25.83$-$0.18 is a very young region of massive star formation in an evolutionary stage prior to the ultra-compact H{~\\small II} region phase. The precise location and luminosity of the protostar responsible for the excitation of the different maser species is unclear; higher sensitivity continuum observations are needed to reveal the position of the exciting source. G25.83$-$0.18, as well as G23.01$-$0.41, may be classified as Group 2 cores following the Longmore et al. (2007) nomenclature, i.e., warm NH$_3$ cores associated with CH$_3$OH masers but no detectable radio continuum.   \\subsection{G23.01$-$0.41 and G25.83$-$0.18 with respect to the other known H$_2$CO Maser Regions}   Including the two new masers reported in this work, H$_2$CO 6$\\,$cm masers have been detected toward seven regions in the Galaxy, and in a total of 15 maser `spots' (at 1\\arcsec~resolution). The H$_2$CO masers in G23.01$-$0.41 and G25.83$-$0.18 share similar characteristics with most of the other known H$_2$CO maser regions, in particular: (1.) H$_2$CO masers are found in regions that harbor a variety of other molecular masers (e.g., Hoffman et al. 2003; Mehringer et al. 1994), (2.) the velocity difference of the H$_2$CO masers with respect to the systemic velocity of the clouds is typically less than 6\\kms (e.g., Pratap et al. 1994; Araya et al. 2004), which suggests that H$_2$CO masers do not originate in high velocity outflows, (3.) the flux density of the known masers is less than $\\sim$2$\\,$Jy (e.g., Hoffman et al. 2007), (4.) all known H$_2$CO masers have been detected in regions of massive star formation (Araya et al. in prep.), (5.) excluding Sgr B2, most of the H$_2$CO masers show double peaked profiles with separations smaller than 3\\kms~(e.g., NGC$\\,$7538 IRS1, Forster et al. 1985; IRAS$\\,$18566+0408, Araya et al. 2007d; see also review by Araya et al. 2007a), (6.) even though bright radio continuum sources may be found in the same star forming complexes, most H$_2$CO masers are located toward sources characterized by weak (or no) compact radio continuum emission, typically undetected at a few mJy sensitivity levels (e.g., Araya et al. 2005; this work), (7) excluding some of the masers in Sgr B2, the known H$_2$CO masers appear to be associated with very young massive stellar objects (in an evolutionary phase prior to the formation of radio-bright ultra-compact H{~\\small II} regions) that have strong far IR emission and molecular core counterparts (e.g., $\\S 4.1$ and 4.2).  Many H$_2$CO maser regions show evidence of outflows/jets and shocked gas based on 4.5$\\mu$m excess emission from Spitzer/IRAC data (see, for example, Araya et al. 2007c in the case of IRAS$\\,$18566+0408). Further evidence comes from molecular data such as SiO, H$_2$, and H$_2$S (e.g., Zhang et al. 2007; Beuther et al. 2007b; Kraus et al. 2006; Gibb et al. 2004; Maxia et al. 2001; Harju et al. 1998), and radio continuum observations (e.g., Araya et al. 2007c). However, the H$_2$CO masers may not be directly associated with the shocked material (e.g., note the offset between the 4.5$\\mu$m excess source and the H$_2$CO maser in G25.83$-$0.18; Figure~4, upper panel).   \\vspace*{-0.5cm}"
    },
    "0806/0806.2899_arXiv.txt": {
        "abstract": "We study the incidence rate of damped $\\Lya$ systems associated with the host galaxies of gamma-ray bursts (\\GDLAs) as functions of neutral hydrogen column density ($\\NHI$) and projected star formation rate (SFR) using cosmological SPH simulations.  Assuming that the occurrence of GRBs is correlated with the local SFR, we find that the median $\\NHI$ of \\GDLAs\\ progressively shifts to lower $\\NHI$ values with increasing redshift, and the incidence rate of \\GDLAs\\ with $\\log \\NHI > 21.0$ decreases rapidly at $z\\ge 6$. Our results suggest that the likelihood of observing the signature of IGM attenuation in GRB afterglows increases towards higher redshift, because it will not be blocked by the red damping wing of DLAs in the GRB host galaxies.  This enhances the prospects of using high-redshift GRBs to probe the reionization history of the Universe.  The overall incidence rate of \\GDLAs\\ decreases monotonically with increasing redshift, whereas that of \\QDLAs\\ increases up to $z=6$. A measurement of the difference between the two incidence rates would enable an estimation of the value of $\\egrb$, which is the mass fraction of stars that become GRBs for a given amount of star formation. Our predictions can be tested by upcoming \\highz\\ GRB missions, including {\\it JANUS (Joint Astrophysics Nascent Universe Scout)} and {\\it SVOM (Space multi-band Variable Object Monitor)}. ",
        "introduction": "\\label{section:intro}  A number of authors have proposed using GRBs to probe the history of cosmic star formation and the reionization of the Universe \\citep[e.g.,][]{Totani97a, Miralda98a, Lamb00, Barkana04}, neither of which is well-understood \\citep[see, e.g.,][]{Holder03, Nag06c}. To date, observations of high-redshift (hereafter \\highz) quasars and galaxies have been able to constrain reionization only up to $z\\sim 7$ \\citep{Fan06a}. However, if GRBs are associated with the deaths of massive stars \\citep[e.g.,][]{Woosley93, Paczynski98}, then theoretical studies imply that GRBs may be detectable out to $z\\simeq 10-20$ through their prompt $\\gamma$-ray emission and afterglows \\citep{Lamb00, Ciardi00, Gou04, Inoue07}. This raises the possibility of using GRBs to investigate the reionization history of the Universe, and the {\\it Swift} satellite has indeed detected \\highz\\ GRBs with bright afterglows \\citep{Cusumano06, Haislip06, Kawai06}.  However, GRB lines-of-sight (LOSs) tend to probe more the inner parts of galaxies than the random QSO LOSs do, and are therefore often associated with neutral hydrogen (\\HI) absorption \\citep{Pro07a}. Indeed, analyses of the afterglow spectra reveal the presence of DLAs in the red damping wing \\citep{Vreeswijk04, Berger06, Watson06, Ruiz07}, and the DLAs may hide the absorption signatures of the neutral IGM \\citep{Totani06}. If such cases dominate the \\highz\\ GRB afterglow spectra, then it may be difficult to use GRBs to probe the detailed reionization history of the Universe \\citep{McQuinn08a}. Thus, it is important to understand the redshift evolution of the incidence rate of \\GDLAs\\ at $z\\ge 6$ as a function of $\\NHI$.  In this {\\it Letter}, we use cosmological SPH simulations based on the concordance $\\Lam$ cold dark matter (CDM) model to study the $\\NHI$ distribution and the incidence rate of \\GDLAs\\ as a function of redshift between $z=1-10$.  The DLAs associated with quasar LOSs are often referred to as \\QDLAs.  Since quasars serve as randomly distributed background beacons in the Universe, \\QDLAs\\ can be more broadly interpreted as all the \\HI\\ gas clouds that satisfy the DLA criterion ($\\NHI > 2\\times 10^{20}$\\,cm$^{-2}$), regardless of whether or not they have been intersected by quasar LOSs.  We adopt the latter broad interpretation of \\QDLAs\\ in this paper, and by this definition \\GDLAs\\ are a subset of \\QDLAs.  \\begin{figure*} \\begin{center} \\includegraphics[angle=0, scale=0.4]{dist_norm.eps} \\includegraphics[angle=0, scale=0.4]{sfrdist.eps} \\caption{Distribution of all DLAs (including both \\QDLAs\\ and \\GDLAs) as a function of $\\NHI$ (panel [{\\it a}]) and $\\Ssfr$ (panel [{\\it b}]) at $z=1-10$. The top axis in each panel indicates corresponding values of $\\NHI$ and $\\Ssfr$ based on the empirical \\citet{Kennicutt98} law. } \\label{fig:nh} \\end{center} \\end{figure*} ",
        "conclusions": "\\label{sec:discussion}  Using cosmological SPH simulations, we have examined the redshift evolution of incidence rates of \\GDLAs, assuming that long GRBs are correlated with local SFR. The distribution of \\GDLAs\\ is intrinsically different from that of \\QDLAs, and the incidence rate of \\GDLAs\\ decreases monotonically towards \\highz, whereas that of \\QDLAs\\ increases from $z=1$ to $z=6$. Quasars are assumed to be randomly distributed background sources in the sky, which illuminate the DLA gas in foreground galaxies. GRBs can also serve as randomly distributed beacons with respect to the DLA gas in foreground galaxies, but for \\GDLAs, GRBs are not random background sources because they are in the same host galaxy.  We find that the incidence rate of \\GDLAs\\ with $\\log \\NHI > 21.0$ decreases rapidly at $z\\ge 6$, suggesting that the likelihood of observing the IGM attenuation signature in GRB afterglows increases toward higher redshifts, without being blocked by the red damping of DLAs in the GRB host galaxies.  This enhances the prospects for using \\highz\\ GRBs to probe the reionization history of the Universe. Our predictions can be tested by upcoming \\highz\\ GRB missions, including {\\it JANUS (Joint Astrophysics Nascent Universe Scout)} and {\\it SVOM (Space multi-band Variable Object Monitor)}.  It might be hoped that it would be possible to estimate the incidence rate of \\GDLAs\\ by accumulating a large GRB sample.  However, because GRBs are not random background sources for \\GDLAs, this would require a prohibitively large GRB sample to estimate the area covering fraction of \\GDLAs\\ from GRB observations alone.  If long GRBs do indeed trace star-forming regions, and if all the LOSs to star-forming regions are coincident with DLAs, then one could estimate the total \\GDLA\\ incidence rate simply by measuring the area covering fraction of star-forming regions from deep imaging surveys of galaxies. When the number of GRBs becomes comparable to that of QSOs, one should expect non-negligible new intervening DLAs in GRB LOSs that are not associated with GRB hosts, since GRB afterglows now act as random beacons.  An alternative possibility would be to search for quasars in the proximity of GRBs, or vice versa.  Such a search of QSO-GRB pair-LOSs would yield a coincidence probability between \\QDLAs\\ and \\GDLAs, which roughly corresponds to the ratio of the two incidence rates. For this purpose, only those systems, for which the QSOs are in the background and the GRBs in the foreground, can be used. So far, no such cases have been identified, but a combination of all-sky GRB surveys (e.g., BATSE, Swift, GLAST) and optical-IR imaging surveys of galaxies (e.g., SDSS, Pan-STARRS, LSST) may prove successful in the future. A constraint on the ratio of the two incidence rates would make it possible to estimate the ratio $\\egrb / \\Ssfr$. In addition, deep observations of GRB host galaxies could constrain $\\Ssfr$ independently.  Then combining the above two constraints would allow us to estimate the value of $\\egrb$."
    },
    "0806/0806.2850_arXiv.txt": {
        "abstract": "Statistical modeling of nuclear data provides a novel approach to nuclear systematics  complementary to established theoretical and phenomenological approaches based on quantum theory.  Continuing previous studies in which global statistical modeling is pursued within the general framework of machine learning theory, we implement advances in training algorithms designed to improved generalization, in application to the problem of reproducing and predicting the halflives of nuclear ground states that decay 100\\% by the $\\beta^-$ mode.  More specifically, fully-connected, multilayer feedforward artificial neural network models are developed  using the Levenberg-Marquardt optimization algorithm together with Bayesian regularization and cross-validation.  The predictive performance of models emerging from extensive computer experiments is compared with that of traditional microscopic and phenomenological models as well as with the performance of other learning systems, including earlier neural network models as well as the support vector machines recently applied to the same problem. In discussing the results,  emphasis is placed on predictions for nuclei that are far from the stability line, and especially those involved in the r-process nucleosynthesis. It is found that the new statistical models can match or even surpass the predictive performance of conventional models for beta-decay systematics and accordingly should provide a valuable additional tool for exploring the expanding nuclear landscape. ",
        "introduction": "INTRODUCTION} \\begin{quotation} \\small \\it ``Numbers are the within of all things.''  Pythagoras of Samos \\end{quotation}  This work is devoted to the development of artificial neural network models which, after being trained with a subset of the available experimental data on beta decay from nuclear ground states, demonstrate significant reliability in the prediction of $\\beta^-$ halflives for nuclides absent from the training set.  The work represents an exploratory study of the degree to which the existing data determines the mapping from proton and neutron numbers to the corresponding $\\beta^-$ halflife.   There is an urgent need among nuclear physicists and astrophysicists for reliable estimates of  $\\beta^-$-decay halflives of nuclei far from stability~\\cite{A1,eirhnh2}.   Among nuclear physicists this need is driven both by the experimental programs of existing and future radioactive ion beam facilities and by the stresses placed on established nuclear structure theory  as totally new areas of the nuclear landscape are opened for exploration. For nuclear astrophysicists, such information is intrinsic to an understanding of supernova explosions -- the initialization of the explosion, the subsequent neutronization of the core material, and the strength and fate of the shock wave formed -- and the nucleosynthesis of heavy elements above \\rm{Fe}, notably the r-process~\\cite{eirhnh3,eirhnh4,eirhnh5}. Both the element distribution on the r-path and the time scale of the r-process are highly sensitive to the $\\beta$-decay properties of the neutron-rich nuclei involved.  In the nuclear chart there are spaces for some 6000 nuclides between the $\\beta$-stability line and the neutron-drip line. Except for a few key nuclei, $\\beta$ decay of r-process nuclei cannot be studied in terrestrial laboratories, so the required information must come from nuclear models. Over the years, a number of approaches for modeling of $\\beta^-$-decay halflives have been proposed and applied.  These include the more phenomenological treatments, such as the Gross Theory (GT), as well as microscopic approaches based on the shell model and the proton-neutron Quasiparticle Random-Phase Approximation ($pn$QRPA) in various versions. More recently, hybrid macroscopic-microscopic and relativistic models have come on the scene.  Some of these approaches emphasize only global applicability, while others seek self-consistency or comprehensive inclusion of nuclear correlations.  Table 1 of Ref.~\\onlinecite{111} provides a convenient summary of a number of the competing models of beta-decay systematics.      In Gross Theory, developed by Takahashi, Yamada and Kondoh~\\cite{22}, gross properties of $\\beta^-$ decay over a wide nuclidic region are predicted by averaging over the final states of the daughter nucleus.  Subsequently, various refinements and modifications of this treatment have been introduced.   The most current of these is the so-called Semi-Gross Theory (SGT), in which the shell effects of only the parent nucleus are taken into account~\\cite{12}.  On the other hand, in the calculations of $\\beta^-$-decay halflives within the shell model, the detailed structure of  $\\beta$ strength function is considered.  Results exist for lighter nuclei and nuclei at  $N= 50, 82$, and $126$.  (See Refs.~\\onlinecite{eirhnh9, eirhnh10} for recent calculations.)  Due to the limits set by the size of the configuration space, calculations are not possible for heavy nuclei.  Several groups have carried out extensive $pn$QRPA studies including pairing.  Efforts along this line by Klapdor and co-workers~\\cite{eirhnh11} began in the framework of the Nilsson single-particle model, including the Gamow-Teller residual interaction in Tamm-Dancoff approximation (TDA), with pairing treated at the BCS level \\cite{2}.  This approach has been complemented and refined by Staudt et ~al.~\\cite{5} and Hirsch et al.~\\cite{eirhnh14}, using $pn$QRPA with the Gamow-Teller residual interaction. The later study by Homma et al.~\\cite{6}, denoted NBCS + $pn$QRPA, includes a schematic interaction also for the first-forbidden (\\textit{ff}) decay.  The Klapdor group has extended the $pn$QRPA theory to calculate $\\beta$-decay halflives in stellar environments using configurations beyond 1$p-$1$h$ \\cite{eirhnh16}.  The starting point of the $\\beta$-decay calculations of M\\\"{o}ller and co-workers is the study of nuclear-ground-state masses and deformations based on the finite-range droplet model (FRDM) and a folded-Yukawa single-particle potential~\\cite{24}. The  $\\beta$-decay halflives for the allowed Gamow-Teller transitions have been obtained from a $pn$QRPA calculation after the addition of pairing and Gamow-Teller residual interactions, in a procedure denoted FRDM + $pn$QRPA~\\cite{eirhnh18, 7}. In the latest calculations the effect of the \\textit{ff} decay has been added by using the Gross Theory ($pn$QRPA +\\textit{ff}GT)~\\cite{8}.  Non-relativistic $pn$QRPA calculations that aim at self-consistency include the Hartree-Fock-Bogoliubov + continuum QRPA (HFB + QRPA) calculations performed with a Skyrme energy-density functional for some spherical even-even semi-magic nuclides with   $N= 50, 82, 126$~\\cite{eirhnh21}. The extended Thomas-Fermi plus Strutinski integral method (ETFSI) (an approximation to HF method based on a Skyrme-type force plus a  $\\delta-$function pairing force) has been elaborated and applied to large-scale predictions of $\\beta^-$  halflives~\\cite{56}. Recently, the density functional + continuum QRPA (DF + CQRPA) approximation, with the spin-isospin effective \\rm{NN} interaction of the finite Fermi system theory operating in the \\textit{ph} channel, has been developed for ground-state properties and Gamow-Teller and \\textit{ff} transitions of nuclei far from the stability line, and applied near closed neutron shells at $N= 50, 82, 126$  and in the region ``east'' of $^{208}{\\rm Pb}$~\\cite{eirhnh23,111}. In the relativistic framework, a $pn$QRPA calculation ($pn$RQPRA) based on a relativistic Hartree-Bogoliubov description of nuclear ground states with the density-dependent effective interaction DD-MEI* has been employed to obtain Gamow-Teller $\\beta^-$-decay halflives of neutron-rich nuclei in the $N\\simeq 50$   and $N\\simeq 82$ regions relevant to the r-process~\\cite{57}. Recently, an extension of the above framework to include momentum-dependent nucleon self-energies was applied in the calculation of  $\\beta$-decay halflives of neutron-rich nuclei in the  $Z\\simeq 28$ and  $Z\\simeq 50$ regions~\\cite{eirhnh77}.  Despite continuing methodological improvements,  the predictive power of these conventional, ``theory-thick'' models is rather limited for  $\\beta^-$-decay halflives of nuclei that are mainly far from stability.  The predictions often deviate from experiment by one or more orders of magnitude and show considerable sensitivity to quantities that are poorly known. In this environment, statistical modeling based on advanced techniques of statistical learning theory or ``machine-learning,'' notably artificial  neural networks (ANNs)~\\cite{19,15} and support vector machines (SVMs)~\\cite{15,27,eirhnh28}, offers an interesting and potentially effective alternative for global modeling of  $\\beta^-$-decay lifetimes.  Such approaches have proven their value for a variety of scientific problems in astronomy, high-energy physics, and biochemistry that involve function approximation and pattern classification~\\cite{59,60}. Statistical modeling implementing machine-learning algorithms is ``theory-thin,'' since it is driven by data with minimal guidance from mechanistic concepts; thus it is very different from the ``theory-thick'' approaches summarized above. Any nuclear observable $X$ can be viewed as a mapping from the atomic and neutron numbers $Z$ and $N$ identifying an arbitrary nuclide, to the corresponding value of the observable (the $\\beta$ halflife, in the present study).  In machine learning, one attempts to approximate the mapping $(Z,N)\\to X$ based only on an available subset of the data for X, i.e., a body of {\\it training data} consisting of known examples of the mapping.  One attempts to {\\it infer} the mapping, in the sense of Bayesian probability theory as expounded by Jaynes~\\cite{eirhnh31}. Thus, one is asking the question: ``To what extent does the data, and only the data, determine the mapping $(Z,N)\\to X$?''  The answer (or answers) to this question should surely be of fundamental interest, when confronted with databases as large, complex, and refined as those existing in nuclear physics.  A learning machine consists of (i) an input interface where, for example, input variables $Z$ and $N$ are fed to the device in coded form, (ii) a system of intermediate elements or units that process the input, and (iii) an output interface where an estimate of the corresponding observable of interest, say the beta halflife $T_{\\beta}$ appears for decoding. Given an adequate body of training data (consisting of input ``patterns'' or vectors and their appropriate outputs), a suitable learning algorithm is used to adjust the parameters of the machine, e.g., the weights of the connections between the processing elements in the case of a neural network. These parameters are adjusted in such a way that the learning machine (a) generates responses at the output interface that closely fit the halflives of the training examples and (b) serves as a reliable predictor of the halflives of the test nuclei absent from the training set. In the more mundane language of function approximation, the learning-machine model provides a means for \\textit{interpolation} or \\textit{extrapolation}.  Neural-network models have already been constructed for a range of nuclear properties including atomic masses, neutron separation energies, ground state spins and parities, and branching probabilities  for different decay channels, as well as  $\\beta^-$-decay halflives~\\cite{59,60,13,14,70,29}. Very recently, global statistical models of some of these properties have also been developed based on support vector machines~\\cite{25,55,eirhnh38}.  In time, there has been steady improvement of the quality of these models, such that the documented performance of the best examples approaches or even surpasses that of the traditional ``theory-thick'' models in predictive reliability.  By their nature, they should not be expected to compete with traditional phenomenological or microscopic models in generating new physical insights. However, their prospects for revealing new regularities are by no means sterile, since the explicit formula created by the learning algorithm for the physical observable being modeled is available for analysis.  We present here a new global model for the halflives of nuclear ground states that decay 100\\% by the $\\beta^-$ mode, developed by implementing the most recent advances in machine-learning algorithms.   Sec.~II describes the elements of the model, the training algorithm employed, steps taken to improve generalization, the data sets adopted, and the coding schemes used at input and output interfaces. Performance measures for assessing the quality of global models of beta lifetimes are reviewed in Sec. III. The results of our large-scale modeling studies are reported and evaluated in Sect. IV. Detailed comparisons are made with experiment, with a selection of the theory-driven GT and $pn$QRPA global models, and with previous ANN and SVM models. This assessment is followed by the presentation of specific predictions for nuclei that are situated far from the line of stability, focusing in particular at those involved in r-process nucleosynthesis. Finally, Sect. V summarizes the conclusions of the present study and considers the prospects for further improvements in statistical prediction of halflives. ",
        "conclusions": ""
    },
    "0806/0806.2133_arXiv.txt": {
        "abstract": "{Semi-analytic treatments of the evolution of orbits of weakly interacting massive particles (WIMPs) in the solar system suggest that the WIMPs bound to the solar system may enhance the direct detection rate relative to that of the unbound population by up to a factor of order unity, and boost the flux of neutrinos from WIMP annihilation in the Earth by up to two orders of magnitude. To test these important but uncertain results, we perform a suite of numerical orbit integrations to explore the properties of the bound WIMP population as a function of the WIMP mass and the scattering cross section with baryonic matter.  For regions of WIMP parameter space presently allowed by experiments, we find that (i) the bound WIMP population enhances the direct detection rate by at most $\\sim 1\\%$ relative to the rate from unbound halo WIMPs; (ii) it is unlikely that planned km$^3$-scale neutrino telescopes will detect neutrinos from WIMP annihilation in the Earth; (iii) the event rate from neutrinos produced by WIMP annihilation in the Sun may be much smaller than implied by the usual calculations, which assume that WIMPs scattered onto bound orbits are rapidly thermalized in the Sun.}  \\FullConference{Identification of dark matter 2008\\\\ August 18-22, 2008\\\\ Stockholm, Sweden}  \\begin{document}  \\newcommand{\\sigmapsi}{$\\sigma_p^{SI}$} \\newcommand{\\sigmansi}{$\\sigma_n^{SI}$} \\newcommand{\\sigmapsd}{$\\sigma_p^{SD}$} \\newcommand{\\sigmansd}{$\\sigma_n^{SD}$} \\newcommand{\\kms}{km~s$^{-1}$} ",
        "introduction": "WIMPs in the solar system may be detected in (i) ``direct detection'' experiments, which measure the recoil of nuclei during interactions with astrophysical WIMPs \\cite{akerib2006b} and have placed interesting constraints on the WIMP mass $m_\\chi$ and spin-independent (or -dependent) elastic scattering cross section with protons \\sigmapsi\\ and neutrons \\sigmansi\\ (\\sigmapsd\\ and \\sigmansd); and (ii) ``indirect detection'' experiments,  including neutrino telescopes.  The next generation of neutrino telescopes may detect the high-energy neutrinos from annihilation of WIMPs captured in the Sun and Earth \\cite{amram1999}.  For a given WIMP model, event rates in direct and indirect detection experiments are determined by the phase-space distribution function (DF) of WIMPs.  The fiducial assumption is that these event rates are dominated by WIMPs from the Galactic halo, passing through the solar system on unbound orbits \\cite{jungman1996}.  However, the following two processes may create additional long-lived populations of WIMPs bound to the solar system. (i) Damour \\& Krauss \\cite{damour1999} argued that secular gravitational interactions with the planets could increase the perihelia of Galactic WIMPs scattered in the Sun to larger radii, quenching further scattering by solar nuclei and dramatically increasing the WIMP lifetimes.  Their semi-analytic estimates indicated that the bound population could increase direct detection rates by up to a factor of $\\mathcal{O}(1)$, and enhance the neutrino flux from the Earth by two orders of magnitude \\cite{damour1999}. (ii) Gould and Lundberg \\& Edsj{\\\" o} \\cite{gould1991} estimated that the density of bound WIMPs at Earth due to gravitational capture and scattering by the planets was $\\sim 1\\%$ of the density of unbound WIMPs if additional elastic scattering in the Sun were unimportant.  Because of their low geocentric speeds the bound WIMPs could enhance the capture rate in the Earth by up to an order of magnitude, hence boosting the neutrino flux.  Semi-analytic techniques cannot, however, describe the full spectrum of behavior in few-body gravitational systems like the solar system.  Numerical orbit integrations are the only reliable tool for this task.  In this paper, we summarize the results of a set of such integrations, which are described further in \\cite{peter2008b}. ",
        "conclusions": ""
    },
    "0806/0806.2419_arXiv.txt": {
        "abstract": "We constrain the uncertainty in waiting times for detecting the first double-neutron-star (DNS) mergers by gravitational wave observatories. By accounting for the Poisson fluctuations in the rate density of DNS mergers and galaxy space density inhomogeneity in the local Universe, we define a detection `zone' as a region in a parameter space constrained by the double neutron star merger rate and two LIGO operations parameters: an observation horizon distance and science run duration. Assuming a mean rate of about 80 DNS mergers per Milky Way galaxy Myr$^{-1}$, we find a 1/20 chance of observing a merger by Enhanced LIGO in only 1 yr of observation. The minimum waiting time and temporal zone width for an Advanced LIGO sensitivity are much shorter and imply that there is a 95\\% probability of detecting a DNS merger in less than 60 days and a 1/20 chance of a first detection in about 1 day. At the 5\\% probability threshold for a first detection, we find that the effect of galaxy clusters on detection is smoothed out and may only influence detection rates after 5--10 years observation time. ",
        "introduction": "Double neutron star (DNS) binary mergers are potentially strong sources of detectable gravitational wave (GW) emission. Nine DNS binary systems have been discovered, eight of them are in the Galactic disk and one is in a globular cluster (Stairs 2004, Lorimer et al. 2006). Energy loss from GW emission \\citep{Tayl89} causes an orbital in-spiral until the binary system merges, resulting in a burst of GWs usually described as a `chirp' signal.  The US LIGO (Laser Interferometer Gravitational-wave Observatory) is searching for, among other potential sources, DNS mergers in the local Universe and has extended its sensitivity horizon to these events to some tens of Mpc (Abbott et al. 2007). LIGO sensitivity to a DNS merger depends on its distance and relative orientation to the event. To aid searches a catalogue has recently been compiled of  distances, sky positions and blue luminosities of galaxies in the local Universe: the Compact Binary Coalescence Galaxy (CBCG) catalogue \\citep{kopp08}.  Theoretical calculations and simulations predict a wide range of delay times between DNS formation and merger (see Belczynski et al. 2006). Their results support the argument that a significant fraction of DNS mergers follow massive star formation. Blue luminosity is a known indicator of massive star formation and may also provide a means of tracking the DNS merger rate, for binaries with merger times that are short compared to the Hubble time. The relatively frequent occurrence of such systems is supported by the discovery of the double pulsar system J0737--3039A,B. With an orbital period of only 2.45 hr, it will coalesce in only 87 Myr. In the local Universe, out to some hundreds of Mpc, we assume, like most other studies, e.g. \\cite{Phin91} and others, that the DNS merger rate scales with blue luminosity.  Interestingly, some short gamma ray bursts (GRBs), observed by the Swift satellite, may also be the high-energy EM emissions from DNS mergers. This is evidenced by the optical localization of several short hard bursts with their host galaxies. The identification of short GRBs with different types of galaxies and distances from galaxy centers is related, in a non-trivial way, to the DNS velocity and the delay times between formation and merger (see e.g. Belczynski et al. 2006). In contrast to the binary pulsar systems, the offset position for GRB050509b at 40 $\\pm 13$ kpc from its elliptical galaxy G17 implies large kick velocities \\citep{grind06}. This discrepancy between the observed radio pulsar binaries and short GRB locations could indicate that DNS mergers may occur in both old stellar populations, such as globular clusters, and relatively younger environments.  Even though the exact  distribution of DNS mergers with galaxy types is not well constrained, it is reasonable to assume that there is a strong correlation between massive star formation rates, DNS merger rate and the blue luminosity of the progenitors.   With the assumption that galactic blue luminosity is a direct measure of massive star formation in the local Universe, the CBCG catalogue can be used to infer the DNS merger rate at extra-galactic distances by scaling from the blue luminosity of our Galaxy. This requires knowledge of the Galactic DNS merger rate. In order to estimate the DNS coalescence rate, Kalogera et al. (2001) used a semi-empirical approach, based on the observed properties of known DNS systems and pulsar survey selection effects, to obtain scale factors that correct for the unobserved fraction of existing systems. With this model and more recent pulsar survey results that include J0737--3039, \\cite{Kalog04} presented bounds for the merger rate for Galactic disk DNS systems, yielding $\\mathcal R_{\\mathrm{DNS}} = 83^{+209.1}_{-66.1}$ Myr$^{-1}$. \\cite{Reg05}, using numerical simulations, find that mergers are possible between $2 \\times 10^5$ yr and the age of the Universe. Using both evolutionary and statistical models, \\cite{Reg05} find a Galactic merger rate of 17 Myr$^{-1}$,  similar to the lower bound calculated by \\cite{Kalog04}. It is clear from these studies that the DNS merger rate is highly uncertain by several orders of magnitude. To account for this uncertainty, we employ a lower range of values for the DNS merger rate of $1-100$ Myr$^{-1}$ per Milky Way Galaxy. ",
        "conclusions": "Figure 2 plots the detection zone defined by equation (\\ref{ezone}). The $D_{5\\%}^{\\mathrm{PEH}}$ curve shows that the inhomogeneity of the local Universe affects the PEH at horizon distances less than 50 Mpc. Waiting times corresponding to $D_{\\mathrm{h}}<50$ Mpc are reduced compared with a uniform galaxy distribution (dot-dashed line) but the deviation manifests after relatively long observation times i.e. $5-10$ yr. For a horizon distance corresponding to a sensitivity of Enhanced LIGO, $\\sim 60$ Mpc, the waiting time is about 1 yr and the detection zone spans about 60 yr. Assuming a horizon distance of Advanced LIGO, the waiting time is about 1 day and the detection zone spans 60 days. This implies that there is a 90\\% probability of observing at least one DNS merger in $1-60$ days of a science run at Advanced LIGO sensitivities.  Figure 3 plots the fractional difference between the PEH$_{5\\%}$ contours for a model that assumes a uniform galaxy distribution and one that includes inhomogeneities in galaxy density in the local Universe. There is a maximum factor of about 1.4 decrease in the horizon distance from clustering in the local Universe, corresponding to an observation time of 7--10 years. Considering that Enhanced and Advanced LIGO should have horizon distances out to the nearly homogenous galaxy space density regime, it is clear that this effect will not be important for science runs of practical duration.   \\cite{kopp08} show that the relatively small horizon distances of the two initial LIGO detectors can vary because of the relative orientation of the detector to the sky positions of galaxies in the local Universe. This is because the detector antenna pattern must coincide with the sky position of the DNS host galaxies for a detection. The galaxies become sparser and the sky separations become larger as the horizon distance becomes smaller. In this analysis, the effect of different detector locations has not been included, as this study focusses more on detector sensitivities comparable to those of Enhanced and Advanced LIGO. The associated horizon distances of these detectors extend from the border of inhomogenous to mostly uniform galaxy space densities.  Figure 4 plots the detection zone at a fixed horizon distance corresponding to that of Advanced LIGO, but with a varying mean DNS merger rate of $1-100$ Myr$^{-1}$ per Milky Way Galaxy to account for the uncertainty in this rate. The corresponding minimum waiting times at the 5\\% probability level for the two rate extremes are 0.5 and 90 days.  In summary, inclusion of the Poisson uncertainty for the minimum waiting times of DNS mergers detectable by GW observatories provides a more realistic means for defining a detection zone in terms of the duration of a science run. Even if the rate density of DNS mergers is well constrained, the study demonstrates how the waiting times can vary widely over an observation period. To highlight this, the model calculates a 1/20 chance of observing a DNS merger by Enhanced LIGO in only 1 yr of observation assuming an event rate of about 80 Myr$^{-1}$ in the Milky Way Galaxy. The horizon distance of Enhanced LIGO is also in the distance regime where inhomogeneities in galaxy space density will only marginally influence the probability of a detection.  The waiting times and zone width at an Advanced LIGO sensitivity are much shorter in duration and imply a first detection at the 95\\% probability level will occur in less than 60 days and there is a 1/20 chance of a first detection in about 1 day. If the merger rate is of order 1 Myr$^{-1}$ in the Milky Way Galaxy, the waiting times at a 5\\% probability level assuming an Advanced LIGO horizon distance are extended to about 90 days."
    },
    "0806/0806.2243_arXiv.txt": {
        "abstract": "In this paper we briefly present our works on the relic gravitational waves (RGW) and the CMB polarization in the accelerating universe. The spectrum of RGW has been obtained, showing the influence of the dark energy. Compared with those  from non-accelerating  models, the shape of the spectrum is approximately similar, nevertheless, the amplitude of RGW now acquires a suppressing factor of the ratio of matter over dark energy $\\propto \\Omega_m/\\Omega_{\\Lambda}\\sim 0.4$ over almost the whole range of frequencies. The RGW spectrum is then used as the source to calculate the spectra of CMB polarization. By a two half Gaussian function as an approximation to the visibility function during the photon decoupling, both the ``electric\" and ``magnetic\" spectra have been analytically derived, which are quite close to the numerical ones. Several physical elements that affect the spectra have been examined, such as the decoupling process, the inflation, the dark energy, the baryons, etc. ",
        "introduction": "The existence of gravitational waves is a major prediction of General Relativity that has not yet been directly detected. On other hand, inflationary models  predict, among other things, a stochastic  background of relic gravitational waves (RGW) generated during the very early stage of expanding universe. \\cite{starobinsky}$^-$\\cite{Grishchuk77} Therefore, the detection of RGW plays a double role in relativity and cosmology. For a number of gravitational  detections, ongoing or under development, the spectrum of RGW represents one of their major scientific goals. However, the current expansion of the universe has been found to be an accelerating one, probably driven by dark energy. This will have important implications on RGW and its detections. As is known, the cosmic background radiation has certain degree  of polarization generated via Thompson scattering during the decoupling in the early universe. \\cite{Basko}$^-$\\cite{BondEfstathiou87} In particular, if the tensorial perturbations (RGW) are present at the photon decoupling in the universe, then magnetic type of polarization will be produced.\\cite{Polnarev80}$^-$\\cite{HuWhite} This would be a characteristic feature of RGW on very large scales, since the density perturbations will not generate this magnetic type of polarization. Besides the generation of linear polarization, the rotation of linearly polarized EM propagation by RGW has also been first studied in Refs.\\cite{Ni,Ni2}, and WMAP polarization data has already been used to constrain the effect in cosmological distance to 0.1 rad, which is important for fundamental physics. For both theoretical and observational studies, it is necessary to examine the effects of the dark energy on RGW and on CMB anisotropies and polarization. In this talk I shall present our calculational results on these issues.   First I will present briefly our result of the spectrum of RGW, both analytical and numerical, in the accelerating Universe $\\Omega_{\\Lambda}+\\Omega_m = 1$. As a double check, we have also derived an approximation of the spectrum analytically. The results from both calculations are consistent with each other. Discussions are given on the possible detections.     Then I will mention sketchily our analytic calculation  of the CMB polarization produced by Thompson scattering in the presence of the RGW. The resulting spectra  are quite close to the numerical one computed from the CMBFAST code, and have several improvements over the previous analytic results. Moreover, the formulae bear the explicit dependence on such important processes, as the decoupling, the inflation, the dark energy, the baryons. ",
        "conclusions": "Our calculations of RGW have shown that in the low frequency range  the peak of spectrum is now located at a frequency $\\nu_E \\simeq (\\frac{\\Omega_m}{\\Omega_{\\Lambda}})^{1/3} \\nu_H$, where $\\nu_H$ is the Hubble frequency, and there appears a new segment of spectrum between  $\\nu_E$ and $\\nu_H$. In all other intervals of frequencies $\\geq  \\nu_H$, the spectral amplitude acquires an extra factor $\\frac{\\Omega_m}{\\Omega_{\\Lambda}}$, due to the current acceleration, otherwise the shape of spectrum is similar to that in the decelerating models. The amplitude  for the model $\\Omega_{\\Lambda}=0.65$ is $\\sim 50 \\%$ greater than that of the model $\\Omega_{\\Lambda}=0.7$. The spectrum sensitively depends on the inflationary models, and a larger $\\beta$ yields  a flatter spectrum, producing  more power. Both the LIGO bound  and the nucleosynthesis bound point out that the inflationary model $\\beta=-1.8$ is ruled out, but the model $\\beta=-2.0$ is still alive.   Our  analytic polarization spectra of CMB has the following improvements.  (i)  The analytic  result of CMB polarization is quite close to the numerical result from the CMBFAST code. The dependence of polarization on the dark energy and the baryons are analyzed. A smaller $\\Omega_{\\Lambda}$ yields a higher amplitude and shifts the peaks to large scales. A larger $\\Omega_b$ yields a lower amplitude and shifts the peaks to large scales.  (ii) Our half-Gaussian approximation of the visibility function fits analytically better than the simple Gaussian fitting, and its time integration yields a parameter-dependent damping factor. This improves the spectrum $ \\sim 30\\%  $ around the second and third peaks.  (iii) The second order of tight coupling limit reduces the amplitude of spectra by $\\sim 58\\%$, comparing with the first order.  (iv)  A larger value of the spectrum index $n_T$  of RGW and a larger ratio $r$ yield higher polarization spectra.   {\\flushleft ACKNOWLEDGMENT: Y.  Zhang would like to thank the organizers of Third International ASTROD Symposium on Laser Astrodynamics, Space Test of Relativity and Gravitational-Wave Astronomy. He also thanks Dr. L. Q. Wen for interesting discussions. The work has been supported by the CNSF No. 10773009, SRFDP,  and CAS. W. Zhao has been supported by Graduate Student Research Funding from USTC.}"
    },
    "0806/0806.0888_arXiv.txt": {
        "abstract": "We investigate the relationship between two massive star-forming galaxy populations at redshift $z\\sim 2$; i.e. submillimetre galaxies (SMGs) and BzK-selected galaxies (BzKs). Out of 60 SMGs  found in the Subaru/XMM-Newton deep field, we collect optical--NIR photometry of 28 radio counterparts for 24 SMGs, based on refined sky positions with a radio map for 35 SMGs (Ivison et al. 2007). We find a correlation between their $K$-band magnitudes and $BzK$ [$\\equiv (z-K)-(B-z)$] colours: almost all of the $K$-faint ($K_\\mathrm{AB} > 21.3$) radio-detected SMGs have $BzK>-0.2$, and therefore BzKs. This result gives strong support to perform direct optical identification of SMGs by searching for BzKs around SMGs. We calculate the formal significance ($P'$ value) for each of the BzK associations around radio-undetected SMGs, and find 6 new robust identifications, including one double identification. From this analysis, we obtain the current best estimate on the surface density of BzK-selected SMGs, which indicates that only $\\sim1$ per cent of BzKs are SMGs. If BzKs are normal disk-like galaxies at $z\\sim 2$ as indicated by the correlation between their star formation rate (SFR) and stellar mass and also by dynamical properties, SMGs are likely to be merging BzKs. In this case, a typical enhancement of SFR due to merging is only a factor of $\\sim 3$, which is an order of magnitude lower than that of local ULIRGs. This may indicate that most of the merging BzKs could be observed as SMGs. Considering a possible high fraction of mergers at $z\\sim 2$ (at least it would be higher than the fraction at $z\\la 1$ of $\\sim 10$ per cent), it is rather puzzling to find such a low fraction of SMGs in the progenitor population, i.e. BzKs. ",
        "introduction": "\\label{sec:introduction}  Among many cosmological surveys, the submillimetre (submm) survey is very unique, in the sense that the expected flux density of sources is almost insensitive to redshift for $z\\approx $ 1 -- 8, owing to the strong negative $K$-correction \\citep[e.g.][]{2002PhR...369..111B}. Although the current sensitivity allows us to detect only the brightest infrared galaxies in the universe, it is possible to detect massive starbursts and gas rich QSOs at extreme redshifts $z\\gg 6$ if exist \\citep[e.g.][]{2008MNRAS.383..289P}. However, most of the submm galaxies (SMGs) currently identified lie at $z\\la 3$. This is not because of the detection limit as noted above, but because of the `identification limit', owing to a large beam size of current (sub)mm telescopes used for surveys. The radio emission provides a high-resolution substitute for the infrared emission observed in the submm \\citep[e.g.][]{2001ApJ...548L.147C,2003ApJ...585...57C, 2002MNRAS.337....1I,2004MNRAS.355..485B,2004ApJ...606..664D,2007MNRAS.380..199I}, and hence the detection limit in the radio has set the upper limits on redshifts of identified SMGs \\citep{2005ApJ...622..772C}. The best way to identify optical counterparts of SMGs is to have high resolution submm images with interferometers \\citep{2006ApJ...640L...1I,2007ApJ...671.1531Y,2007ApJ...670L..89W,2008ApJ...673L.127D}. Although time consuming, it allows us to perform direct optical/near-infrared (NIR) identifications from the submm position.  Here we adopt another method of direct optical identification using optical/NIR colours of galaxies around the submm source. The expected number of random optical/NIR associations within the error circle of the submm position is not negligible; e.g. there would be $\\sim 4$ $K$-band sources ($K_\\mathrm{AB} < 23)$ on average within a $r<8''$ radius (see below). However, if SMGs have rather confined optical-NIR colours and the surface density of objects with similar colours is reasonably low, it will be possible to directly identify SMGs without using costly high-resolution submm images. As in the radio identification method \\citep{1986MNRAS.218...31D}, we can calculate the formal significance of each optical/NIR association, given the number counts of colour-selected galaxies. To do this we need a firm basis for the colours of SMGs.  Since the discovery of SMGs, extremely red objects (EROs) usually defined with $(I-K)_\\mathrm{Vega}\\ga 4$ or $(R-K)_\\mathrm{Vega}\\ga 5$ and distant red galaxies with $(J-K)_\\mathrm{Vega}>2.3$ have been paid special attention as candidates of SMG counterparts \\citep[e.g.][]{1999MNRAS.308.1061S,2002MNRAS.331..495S,2003ApJ...597..680W,2004MNRAS.354..193C,2004AJ....127..728F,2004MNRAS.351..447C,2005MNRAS.358..149P}. It turns out that the fraction of EROs or DRGs in SMGs is not very high and their surface densities are not low enough to reject random associations. Among the NIR-selected galaxy populations in the literature, BzK-selected star-forming galaxies \\citep[BzKs --][]{2004ApJ...617..746D} may be the most promising counterparts of SMGs \\citep{2005ApJ...633..748R,2007ApJS..172..132B,2007MNRAS.381.1154T}, which lie at $1.4\\la z \\la 2.5$ and include heavily obscured galaxies, such as SMGs, as an extreme subset.  In this paper, we indeed find that almost all of the $K$-faint (indicating higher redshifts) radio-detected SMGs are BzKs. We then use BzKs as a key galaxy population to identify radio-undetected SMGs at $1.4\\la z \\la 2.5$. Since the BzK-selection is based on observed wavelengths covering a spectral break at 4000\\,\\AA, we can naturally extend this selection technique to the higher redshift range , e.g. $z\\gg 3$, in the future by using a combination of wavebands at longer wavelengths. We also emphasize that the understanding of the physical relation between SMGs and BzKs would be important to reveal the evolution of galaxies at $z\\sim 2$, given that they have similar stellar masses of $\\sim 10^{11}$\\,M$_\\odot$, and spatial correlation lengths of $\\sim 10$\\,$h^{-1}$\\,Mpc \\citep{2004ApJ...611..725B,2006ApJ...638...72K,2007ApJ...660...72H}. We discuss the hypothesis of SMGs being merging BzKs.  We adopt the sample of SMGs from the SCUBA HAlf Degree Extragalactic Survey \\citep[SHADES --][]{2005MNRAS.363..563M,2006MNRAS.372.1621C} as described in Section 2. In Section 3, we investigate typical optical/NIR colours of the radio-detected SMGs. We apply our direct identification method for radio-undetected SMGs in Section 4. We then discuss a possible evolutionary link between SMGs and BzKs in Section 5. Finally we give our summary in Section 6. Throughout this paper, we assume a cosmology with $\\Omega_\\mathrm{m} =0.3$, $\\Omega_\\Lambda =0.7$ and $H_0 =70$$\\,$km$\\,$sec$^{-1}$$\\,$Mpc$^{-1}$. All magnitudes in this paper use the AB system unless otherwise noted.    \\begin{table*} \\begin{minipage}{120mm} \\tabcolsep=5pt \\caption{Photometry of radio-detected SMGs} \\begin{tabular}{lllcccccccccc}% \\hline \\multicolumn{1}{c}{Name} & \\multicolumn{1}{c}{R.A.} & \\multicolumn{1}{c|}{Dec.} & \\multicolumn{1}{c|}{$K_s$} & \\multicolumn{1}{c|}{$z'-K_s$} & \\multicolumn{1}{c|}{$B-z'$} & \\multicolumn{1}{c|}{$R-K_s$} & \\multicolumn{1}{c|}{$J-K_s$} \\\\ \\multicolumn{1}{c}{} & \\multicolumn{2}{c}{[J2000]} & \\multicolumn{1}{c}{[AB]$^a$} & \\multicolumn{4}{c}{[AB]$^b$} \\\\ \\hline SXDF850.01&  34.37758&  -4.99355&   22.01$\\pm$0.07&    2.59&    1.60&    3.46&    2.54\\\\ SXDF850.02&  34.51490&  -4.92432&   21.50$\\pm$0.04&    2.61&    1.11&    2.95&    2.13\\\\ SXDF850.03&  34.42561&  -4.94124&   18.77$\\pm$0.01&    1.02&    2.92&    1.71&    0.55\\\\ SXDF850.04&  34.41116&  -5.06093&   21.04$\\pm$0.03&    2.83&    1.43&    3.54&    1.37\\\\ SXDF850.05&  34.51192&  -5.00856&   20.34$\\pm$0.01&    2.26&    2.55&    4.05&    1.50\\\\ SXDF850.07&  34.41192&  -5.09108&   21.22$\\pm$0.03&    2.67&    1.30&    3.12&    1.75\\\\ SXDF850.08&  34.43402&  -4.93019&   22.95$\\pm$0.10&    3.59&    0.44&    3.64&  $\\ge 0.41^c$\\\\ SXDF850.10&  34.60375&  -4.93423&   21.93$\\pm$0.05&    2.14&    1.41&    2.69&    1.60\\\\ SXDF850.12&  34.49710&  -5.08446&   22.42$\\pm$0.05&    1.19&    2.06&    1.64&    0.90\\\\ SXDF850.16&  34.55784&  -4.96193&   21.84$\\pm$0.03&    2.39&    1.10&    3.17&    1.07\\\\ SXDF850.19&  34.61584&  -4.97691&   21.53$\\pm$0.04&    2.25&    1.45&    3.17&    0.96\\\\ SXDF850.21$^d$ &  34.42711&  -5.07356&   15.51$\\pm$0.01&    0.33&    0.85&    0.53&   -0.18\\\\ SXDF850.23&  34.42679&  -5.09602&   23.54$\\pm$0.12&    2.67&    0.49&    3.15&    1.35\\\\ SXDF850.24b&  34.39473&  -5.07499&   20.74$\\pm$0.02&    2.24&    2.53&    3.51&    1.16\\\\ SXDF850.27&  34.53304&  -5.02927&   21.85$\\pm$0.03&    3.01&    1.61&    3.91&    1.33\\\\ SXDF850.28a&  34.52867&  -4.98692&   19.20$\\pm$0.01&    2.13&    3.26&    3.33&    1.13\\\\ SXDF850.28b&  34.52839&  -4.98821&   20.73$\\pm$0.01&    2.25&    2.95&    3.99&    1.23\\\\ SXDF850.30&  34.41660&  -5.02102&   21.37$\\pm$0.03&    3.95&    2.70&    5.29&    2.07\\\\ SXDF850.31&  34.39938&  -4.93208&   19.71$\\pm$0.01&    1.92&    2.85&    3.00&    1.09\\\\ SXDF850.35&  34.50349&  -4.88503&   19.68$\\pm$0.01&    1.91&    2.22&    3.04&    1.04\\\\ SXDF850.37&  34.35240&  -4.97823&   22.09$\\pm$0.07&    3.46&    1.32&    4.35& $\\ge 0.73^c$\\\\ SXDF850.47a&  34.39322&  -4.98260&   21.00$\\pm$0.02&    2.02&    3.52&    4.03&    0.97\\\\ SXDF850.47b&  34.39337&  -4.98329&   21.18$\\pm$0.02&    1.69&    2.04&    2.89&    0.97\\\\ SXDF850.47c&  34.39039&  -4.98269&   19.76$\\pm$0.01&    2.13&    1.79&    3.27&    0.93\\\\ SXDF850.52a&  34.52133&  -5.08157&   19.12$\\pm$0.01&    1.02&    1.70&    1.47&    0.61\\\\ SXDF850.52b&  34.52065&  -5.08368&   21.49$\\pm$0.02&    2.88&    3.18&    4.83&    1.55\\\\ SXDF850.77&  34.40069&  -5.07595&   20.35$\\pm$0.01&    1.95&    3.12&    3.36&    1.12\\\\ SXDF850.96&  34.50083&  -5.03784&   22.07$\\pm$0.04&    1.73&    0.99&    2.41&    1.03\\\\ \\hline\\end{tabular}\\label{phot_table} \\medskip Note --- $^a$ Total (Petrosian) magnitudes from the UKIDSS/UDS catalogue. $^b$ Aperture-matched ($2''$) colours, corrected for galactic extinction. $^c$ 5\\,$\\sigma$ upper limit in $J$. $^d$ A nearby galaxy with large angular size. Colours are derived from total magnitudes. \\end{minipage} \\end{table*} ",
        "conclusions": "% We discuss a possible evolutionary link between SMGs and BzKs and its implications for physical properties of SMGs, by using the star formation rate (SFR)-stellar mass ($M_*$) relation of BzKs. In the following, we discuss only SMGs satisfying the BzK selection criterion and $K<20$ (Vega), in order to securely derive stellar masses. We find seven such SMGs (singly associated ones) from our sample. In order to increase the sample, we also use the sample of radio-detected SMGs in \\cite{2005ApJ...622..772C}, which have spectroscopic redshifts and $K$-band photometry \\citep{2004ApJ...616...71S}. We find seven SMGs in their sample, which have a starburst-dominated optical spectrum, a redshift at $1.4 \\le z \\le 2.5$ (i.e. the redshift range of BzKs), and $K<20$ (Vega). We refer to this sample and the sample from the SXDF as the Chapman sample and the SXDF sample, respectively.    \\subsection{Are SMGs merging BzKs?} The starburst activity of SMGs seems to be induced by galaxy interactions/mergers as indicated by their morphology and dynamical properties \\citep{2004ApJ...611..732C,2006MNRAS.371..465S,2006ApJ...640..228T, 2008arXiv0801.3650T,2007ApJ...671..303B}. Immediate progenitors of SMGs would be gaseous star-forming galaxies as massive as SMGs. Most plausible candidates for this parent population would be star-forming BzKs, given a similar stellar mass and coevality with SMGs. Although the SFRs of BzKs are sometimes as high as those of local ULIRGs \\citep{2005ApJ...631L..13D}, the kinematical properties of BzKs are more similar to quiescent disc galaxies \\citep{2006Natur.442..786G,2007ApJ...671..303B}. This quiescent nature of star formation in BzKs is also supported by the measurement of the star formation efficiency from the CO luminosity, which is an order of magnitude lower than that of SMGs and similar to local spirals \\citep{2008ApJ...673L..21D}. In the following discussion, we assume that some dynamical perturbations occurring in BzKs induce the vigorous star formation of SMGs.   \\subsection{SFR and stellar mass of SMGs} We derive the stellar masses of SMGs, using the empirical formulae by Daddi et al. (2004, their Eq. 6 and 7 for the SED-fitting--derived mass): \\begin{eqnarray} \\log (M_* / 10^{11} M_\\odot) &= &-0.4(K^{\\mathrm{tot}} - K^{11}) + \\Delta \\log M_*, \\\\ \\Delta \\log M_* &=& 0.218 [(z-K)-2.29] \\end{eqnarray} where $K^{11}=19.51$ (Vega) is the $K$-band magnitude corresponding on average to a mass of $10^{11}$\\,M$_\\odot$. Since these formulae are calibrated for $K<20$ (Vega), we restrict our sample of SMGs to those in this magnitude range as noted above. The correction term with $(z-K)$ colour is applied only for the SXDF sample, since $z$-band photometry is not available for the Chapman sample in the literature. The stellar masses derived from these formulae are based on the Salpeter initial mass function extending between 0.1 and 100\\,M$_\\odot$. In deriving the stellar masses, we adopt Petrosian magnitudes from the UKIDSS catalogue for estimation of total $K$-band magnitudes. For the Chapman sample, we adopt the total $K$-band magnitudes in Smail et al. (2004), which are estimated with the aperture photometry of a large (4$''$) diameter. We find that the mean stellar mass of 14 SMGs (7 from SXDF and the other 7 from the Chapman sample) is 1.2$\\times 10^{11}$\\,M$_\\odot$. If we included SMGs with $K>20$ (Vega) as well, the mean stellar mass would reduce about 20\\,\\%. The uncertainty in the stellar mass on single objects is about 60\\,\\% (Daddi et al. 2004). However, we caution that the uncertainty may be much larger than this, since SMGs would have a large fraction of mass in substantially obscured stellar populations. For example, \\cite{2005ApJ...635..853B} derive on average 5 times higher stellar masses from rest-frame $K$-band magnitudes than those reported from modeling the UV/optical SEDs of SMGs in Smail et al. (2004).   The SFRs are derived from radio 1.4\\,GHz fluxes and redshifts. We estimate the FIR luminosities from the radio-infrared luminosity relation, $q_L= \\log (L_\\mathrm{FIR}/ \\nu_0 L_\\mathrm{1.4\\,GHz})$ found for SMGs \\citep{2006ApJ...650..592K}, where $q_L=2.14$ and $\\nu_0 = 4.52$\\,THz. The radio luminosities at the rest-frame 1.4\\,GHz are calculated with $L_\\mathrm{1.4\\,GHz} = 4\\pi D_L^2 S_\\mathrm{1.4} (1+z)^{\\alpha-1}$ where $D_L$ and $S_{1.4}$ are the luminosity distance and the 1.4\\,GHz flux density, and $\\alpha$ is the spectral index ($S\\propto \\nu^{-\\alpha}$). We adopt $\\alpha = 0.7$ \\citep{1992ARA&A..30..575C}. For the SXDF sample, we assume a redshift of $z=1.9$, an average redshift of BzKs \\citep{2004ApJ...617..746D}. For the Chapman sample, we use the spectroscopic redshifts from Chapman et al. (2005). We convert the derived FIR luminosity to the SFR using a relation SFR = $L_\\mathrm{FIR} / 5.8\\times 10^9$\\,L$_\\odot$ \\citep{1998ApJ...498..541K}. For the combined sample of SMGs, we derive an average SFR of 680\\,M$_\\odot$\\,yr$^{-1}$. This becomes 600\\,M$_\\odot$\\,yr$^{-1}$ if we include the SMGs with $K>20$ (Vega). If we adopt the widely-used local radio-FIR relation by \\cite{1992ARA&A..30..575C}, instead of that for SMGs by \\cite{2006ApJ...650..592K}, the derived SFRs increases by a factor of $\\sim 3$.   Local ULIRGs are also considered for comparison. We adopt the average of the sample in \\cite{2002ApJ...580...73T}. We use the dynamical mass, instead of the stellar mass, which gives a solid upper limit on the stellar mass. The averages of the dynamical masses and the SFRs are 1.3$\\times 10^{11}$\\,M$_\\odot$ and 340\\,M$_\\odot$\\,yr$^{-1}$, respectively.    \\begin{figure}% \\resizebox{8cm}{!}{\\includegraphics{sfr_mass.eps}} \\caption{SFR vs.\\ stellar mass of SMGs. For local ULIRGs we adopt the dynamical mass. Solid circles indicate SMGs. Solid circles inside diamonds indicate SMGs with spectroscopic redshifts. Open circles are for BzKs with $K<22$ mag (Vega), taken from Daddi et al. (2007) with MIPS 24$\\mu$m-based SFRs. The large solid circle indicates the average of SMGs, whose error bars represent the standard deviation of the combined SMG sample. The large open square corresponds to the average of local ULIRGs. Solid and dashed lines indicate the observed SFR-stellar mass relation of star-forming galaxies at $z\\sim 2$ and 0, respectively. Solid and dashed arrows show a typical enhancement of SFR experienced by SMGs and local ULIRGs, respectively. } \\label{mass} \\end{figure}   \\subsection{SFR enhancement of SMGs} We now consider how much SFRs are enhanced in SMGs, owing to dynamical perturbations. Such an enhancement may be estimated from a comparison of the average SFR of SMGs with that of BzKs, a suspected progenitor population of SMGs, with a similar stellar mass. We note that the increase in the stellar mass would be less than a factor of $\\sim2$ in merging of gas-righ galaxies to produce still gaseous SMGs, in which we could neglect the increase in stellar mass after merging. The average SFR of progenitor galaxies may be estimated from the correlation between SFRs and stellar masses found for BzKs or local star-forming galaxies, once we derive the stellar mass of SMGs or local ULIRGs.  In Figure \\ref{mass}, we plot the SFRs and stellar masses of SMGs and local ULIRGs. \\cite{2007ApJ...670..156D} present the correlation between the SFR and the stellar mass of BzKs, SFR = 200$\\times M_{11}^{0.9}$ where $M_{11} = 10^{11}$\\,M$_\\odot$, which is depicted as the solid line in the figure. Also, the same correlation but for local blue galaxies found in the Sloan Digital Sky Survey -- SFR = 8.7$\\times M_{11}^{0.77}$ \\citep{2004MNRAS.351.1151B,2007A&A...468...33E} -- is shown as the dashed line.  Comparing the average SFR and $M_*$ of SMGs to the SFR-$M_*$ relation of BzKs, we estimate that a typical enhancement of SFR experienced by SMGs is around a factor of 3 indicated as the solid arrow in Figure \\ref{mass}. This enhancement is rather moderate as if SMGs follow the same SFR-$M_*$ relation. In the local universe, such a moderate enhancement is, for example, found for M82 (Figure 18 in Elbaz et al. 2007), whose starburst activity is likely to be induced by interactions with M81. Compared to local ULIRGs, this enhancement is less significant by an order of magnitude. We estimate the SFR enhancement of local ULIRGs to be a factor of $\\sim 30$ (the dotted arrow), comparing the mean SFR and $M_*$ of local ULIRGs with the SFR-$M_*$ relation of local star-forming galaxies. This is a lower limit, since we use the dynamical mass for local ULIRGs. If progenitors of SMGs are BzKs with a stellar mass comparable to SMGs, then the SFR enhancement could {\\it not} be as high as local ULIRGs. As \\cite{2006ApJ...640..228T} indicated, SMGs are probably experiencing maximum starbursts \\citep{1999ApJ...517..103E}, determined roughly by the dynamical time scale of galaxies and negative feedback due to super novae. The maximum SFR expected for SMGs would be around 600\\,M$_\\odot$\\,yr$^{-1}$ \\citep {2006ApJ...640..228T}, which is already comparable to the observed SFRs.  What are implications of a moderate ($3\\times$) SFR enhancement of SMGs? In the local universe, such a moderate enhancement is found for typical interacting pairs of galaxies (not necessarily in the most luminous merging phase) with a projected distance of $r_p < 20$ kpc \\citep{2007arXiv0711.3792L}. The numerical simulation by \\cite{2007A&A...468...61D} also indicates that a typical enhancement of SFR during the major merger phase is only a factor of $<5$. If this moderate enhancement is a typical case at $z\\sim 2$ as well, the fraction of SMGs in star-forming BzKs could be as high as the merger fraction. Although the merger fraction at $z>1$ is quite uncertain, it would be at least similar to or higher than the value at $z\\la 1$ of $\\sim 10$\\,\\% \\citep[e.g.][]{2003AJ....126.1183C,2008ApJ...672..177L}, because of the hierarchical clustering properties of the galaxy formation process. This argument leaves us an open question why the fraction of SMGs in BzKs is not $>10$ per cent, but only $\\sim 1$ per cent. A bias of the submm surveys, which tend to miss galaxies as luminous as SMGs but with higher dust temperatures, should play a role here, but the effect would be only a factor of $\\sim 2$ \\citep{2004ApJ...614..671C}. Also we are missing some fraction of radio-undetected, and hence cool-dust SMGs, which are not extremely red (see Section 4.2). Even if we consider luminous dusty galaxies at $z\\sim 2$ with both higher and lower dust temperatures, a correction for the incompleteness would be less than an order of magnitude, and therefore not enough to reconcile the apparent inconsistency.   We could explain this apparent inconsistency between the rareness of SMGs and its moderate SFR enhancement as follows. At $z\\sim 2$, galaxy interactions may be too common, which means that high specific SFRs (mass-normalized SFRs) of BzKs are a result of galaxy interactions to some extent. In this case, there would be almost no net enhancement of SFR in typical interacting galaxies over the SFR-$M_*$ relation. The SFR enhancement of SMGs looks moderate, but might be indeed significant as galaxies at $z\\sim 2$, possibly caused only by rare merging events, such as those with some particular orbit parameters, very equal mass ratios, or multiple mergers. For example, according to the numerical simulation by \\cite{2007A&A...468...61D}, high gas concentrations and hence high SFR enhancement can be realized not in direct mergers, but in retrograde mergers.   We should bare in mind that SMGs are not the only galaxy population experiencing major mergers. SMGs will merely be a tip of iceberg. \\cite{2007ApJ...656....1L} suggest that UV morphologies of BzKs and UV-selected (BX/BM) galaxies are not distinguishable from that of SMGs. This indicates either that not only SMGs but also BzKs and BX/BM galaxies are experiencing merger-induced starbursts,  or simply that the UV morphology is not a reliable indicator of a major merger. On the other hand, \\cite{2007ApJ...671..303B} show that SMGs are a dynamically distinctive galaxy population with large mass concentration, i.e. dynamically hotter than BzKs and BX/BM galaxies. Therefore, the UV morphology may not be very useful to identify dynamically hot galaxies at $z\\ga 2$, as suggested by no obvious correlations between the UV morphology and other galactic properties, such as SFR, outflow and stellar mass \\citep{2007ApJ...656....1L}. We need more direct measures on dynamical properties of galaxies at $z\\sim 2$, in order to identify kin of SMGs. For this purpose, Atacama Large Millimeter Array \\citep[ALMA --][]{2006SPIE.6267E...2B} will play an important role, which is capable of measuring gas dynamics of $\\sim1$-mJy sources, such as BzKs \\citep{2005ApJ...631L..13D}."
    },
    "0806/0806.0841_arXiv.txt": {
        "abstract": "It is well known that the leptogenesis mechanism offers an attractive possibility to explain the baryon asymmetry of the universe. Its particular robustness however comes with one major difficulty: it will be very hard if not impossible to \\emph{test} experimentally in a foreseeable future, as most of the mechanics \\emph{typically} takes place at high energy or \\emph{results} from suppressed interactions, without unavoidable low-energy implications. An alternate approach is taken by asking: can it be at least falsified? We show that \\emph{possible} discoveries at current and future colliders, most notably that of right-handed gauge interactions, would indeed forbid at least the \"canonical\" leptogenesis mechanisms, namely those based on right-handed neutrino decay. General lower bounds for successful leptogenesis on the mass of the right-handed gauge boson $W_R$ are given. Other possibilities to falsify leptogenesis, including from the observation of a $Z'$, are also considered. ",
        "introduction": "The recent evidence for neutrino masses has brought forward leptogenesis \\cite{fy} as a very attractive mechanism to explain the baryon asymmetry of the universe. Along this mechanism, the baryon asymmetry of the universe is explained by the same interactions as the ones which can explain the neutrino masses. In the most straightforward seesaw model, which assumes right-handed neutrinos in addition to the standard model particles, both neutrino masses and leptogenesis originate from the Yukawa interactions and lepton number violating  Majorana masses of the right-handed neutrinos \\begin{equation} {\\cal L} \\owns - \\overline{L} \\,{\\widetilde H} \\, {Y_\\nu^\\dagger}  \\, N -\\frac{1}{2}\\,\\overline{N} \\, {m_N} \\,{{N}^c} +\\text{h.c.} \\end{equation} where $L$ stands for the lepton weak doublets and $\\tilde H$ is related to the standard Brout-Englert-Higgs (hereafter simply Higgs) doublet $H \\equiv (H^+, H^0)$ by $\\tilde H = i \\tau_{2} H^*$.  However, testing this mechanism will be a very difficult task for several reasons. If the right-handed neutrinos have a hierarchical mass spectrum, due to neutrino mass constraints, leptogenesis through $N$ decay can lead to the observed amount of baryon asymmetry e.g.~only if it involves right-handed neutrinos with masses above $\\sim 10^8$~GeV \\cite{predi,di}. As a result they cannot be produced at colliders. Moreover there are many more parameters in the Yukawa coupling matrices which can play an important role for leptogenesis, than there are (not too suppressed) low energy observables which could constrain these parameters.\\footnote{A possible exception to that arises in the supersymmetric case from the effects of Yukawa couplings on the running of the slepton masses \\cite{raidal}. This nevertheless assumes that universality of lepton soft mass terms must be present (an assumption which requires to be tested) and, for any real test of leptogenesis, would require to observe a long series of rare leptonic decays not necessarily expected to be all close to the present corresponding experimental bounds.}  If the right-handed neutrinos have instead a quasi-degenerate spectrum (for at least 2 of them), leptogenesis can be efficient at lower scales \\cite{lowscale} but generically in this case the neutrino mass constraints require suppressed values of Yukawa couplings, which hampers their production at colliders.  For leptogenesis to be both efficient and tested at low energy, not only is a  quasi-degeneracy between 2 right-handed neutrinos required, but also a special flavour structure which allows for larger Yukawa couplings while preserving  the light neutrino mass constraints,\\footnote{This case can be realized if the Yukawa induced dimension 6 operator coefficients are unsuppressed (decoupling from the suppressed neutrino mass dimension 5 ones). This does not necessarily require cancellations of the various entries. It requires that some of entries are smaller than others, as in the inverse seesaw, see e.g.~\\cite{ValleConcha,Smirnov,ABBGH}. But it e.g.~leads only to lepton conserving channels with rather large background at LHC \\cite{delAguila}.} and/or a right-handed neutrino production mechanisms other than through the Yukawas and associated neutrino mixings.  In this paper we consider the problem of testing leptogenesis mechanisms  the other way around. While they cannot confirm leptogenesis, could low energy observations at least  exclude it? We propose one particularly clear possibility, namely the observation of  a right-handed charged gauge boson $W_R$. It is known that for high mass right-handed neutrinos and $W_R$, around $10^{10}$ GeV or higher, the $W_R$ can have suppression effects on leptogenesis through dilution and scattering, but, in the specific case of reheating after inflation, they can also boost the $N$ abundances \\cite{Carlier,Cosme,2sarkar} and hence relax the constraints on Yukawa couplings. Not surprisingly, with a low scale $W_R$ the suppression effects are dramatically enhanced. Actually, see section 2, they turn out to be so strongly enhanced that, even with a maximal CP asymmetry of order unity, leptogenesis cannot be a sufficient cause of the matter excess anymore.   Right-handed gauge interactions lead in particular to much larger suppression effects at low scale than left-handed interactions do in other contexts (i.e.~than in leptogenesis from scalar \\cite{typeIIlepto,typeIIleptoeffic} or fermion \\cite{typeIIIlepto} triplet decays, whose efficiency have been calculated in Refs.~\\cite{typeIIleptoeffic,typeIIIlepto}). This is due to the fact that at the difference of triplets,  a single $N$ can interact through $W_R$ exchange with fermions which are all in thermal equilibrium, which induces more efficient, and hence dangerous, scatterings and decays. In particular, some of the scatterings involving the $W_R$ turn out to induce a very large suppression due to the fact that they do not decouple through a Boltzmann suppression. The production of $N$'s through a light $W_R$, often presented as the easiest way to produce $N$'s, is therefore incompatible with successful leptogenesis, and even enhanced $N$ production from reheating cannot compensate for the large suppression. The lower bounds on the mass of the $W_R$, required for successful leptogenesis, are given in section 3.  The possible discovery of a low-energy $W_R$ has recently been the object of several analysis by LHC collaborations \\cite{Wlhcferrari,Wlhckras,LHCstudies}. It should be feasible up to $m_{W_R} \\sim 3$-5~TeV (see more details, and additional possible searches, in section 7).   The observation of a $W_R$ is not the only possibility to exclude canonical neutrino decay leptogenesis from current energy data. We give a list of other possibilities in section 5, considering in particular the implications of the observation of a $Z'$ at LHC. The case of other leptogenesis seesaw models with not only or without right-handed neutrinos is briefly considered in section 6. ",
        "conclusions": "We have shown that the discovery at LHC or future accelerators, of a $W_R$ coupling to a right-handed neutrino and a right-handed charged lepton, would rule out the possibility to create any relevant lepton asymmetry from the decay of right-handed neutrinos, see Fig.~2. A $W_R$ induces extra $N$ decay channels inducing large dilution and washout effects, as well as very fast gauge scatterings (whose decoupling doesn't occur through Boltzmann suppression). We determined bounds on $m_{W_R}$ and $m_N$ for successful leptogenesis, given in Fig.~\\ref{Wbounds} and Eqs.~(\\ref{boundsreso}) and (\\ref{boundshier}). Similarly we discussed how the discovery of other particles generally expected in presence of right-handed gauge interactions, or of a $Z'$, could also affect leptogenesis, ruling it out too in some cases. Leptogenesis from the decay of scalar or fermion triplet would be also basically ruled out in presence of a $N$ or both a $N$ and a $W_R$ around the TeV scale, unless there is a flavour symmetry to protect one flavour combination from the washout due to these states."
    },
    "0806/0806.4465_arXiv.txt": {
        "abstract": "We present stellar metallicities derived from Ca II triplet spectroscopy in over 350 red giant branch stars in 13 fields distributed in different positions in the SMC, ranging from $\\sim$1\\arcdeg\\@ to $\\sim$4\\arcdeg\\@ from its center. In the innermost fields the average metallicity is [Fe/H] $\\sim -1$. This value decreases when we move away towards outermost regions. This is the first detection of a metallicity gradient in this galaxy.  We show that the metallicity gradient is related to an age gradient, in the sense that more metal-rich stars, which are also younger, are concentrated in the central regions of the galaxy. ",
        "introduction": "The chemical enrichment history of a galaxy is related to the origin and distribution of nuclear species in its stars and gas. The chemical elements are mainly produced by stars which drive the enrichment of the interstellar medium by ejecting material containing the product of the stellar nucleosynthesis from which the new generations of stars are formed. In addition, gas flows also play an important role in chemical enrichment, diluting the products of the stellar nucleosynthesis with unenriched material from outside the galaxy, and mixing metals from one part of the system to another (e.g. bringing metal-rich gas into metal-poor regions). Thus, the study of chemical evolution of galaxies involves understanding the spatial distribution and temporal evolution of various elements by taking into account the processes of star formation, the distribution of stars according to their masses and chemical compositions, and the final yields of various elements and detectable remnants of parent stars. Until recently, only the chemical enrichment history of the solar vicinity could be studied in detail. However, the modern multiobject spectrographers attached to the 8-10 m class telescopes allow us to study the chemical enrichment history of the nearest Local Group galaxies.  Because of their proximity, and the fact that they present a wide range of ages and metallicities, the Magellanic Clouds are attractive objects to study chemical enrichment histories. In a previous paper \\citep[][hereafter paper III]{carrera07b}, we investigated the chemical enrichment history of the LMC. In this paper, we will focus in the study of the SMC. There are considerable less studies of the SMC as compared with the LMC, probably due to: (i) its irregular appearance, with complex kinematics; (ii) its distance, located further away than the LMC; and (iii) its depth in the line of sight, which remains a subject of controversy.  Most of the information about the stellar populations of the SMC has been obtained from its cluster system \\citep[e.g.][]{costahatz98,piatti01,piatti05a}. The cluster age distribution does not show the age-gap observed in the LMC \\citep{costahatz98,mighell98}. From a sample of seven clusters older than 1 Gyr, \\citet{rich00} suggest that star formation was stronger in two main episodes, one $\\sim$8$\\pm$2 Gyr ago and another $\\sim$2$\\pm$0.5 Gyr ago. However, there does not seem to be any age interval completely lacking of objects \\citep{rafelski05}. There is only one old cluster, NGC 121, that is younger than most of the LMC globular clusters \\citep{piatti05a}. However, the SMC shows a significant old field population \\citep[hereafter Paper I]{noel07}.  To our knowledge, there are only three studies in which the star formation history (SFH) of the SMC field population was derived \\citep{dolphin01,harriszaritsky04}. From a deep color--magnitude diagram (CMD) \\citet{dolphin01} found than the star formation in a small field in the periphery of this galaxy was relatively constant until about 2 Gyr ago, with no star formation occurring since then. However, this could be biased by the fact that their field was specifically chosen because it did not have young stars. \\citet{harriszaritsky04} studied a 4\\arcdeg$\\times$4$\\fdg$5 field in the central region of the galaxy. From shallower CMDs they found that the star formation in the SMC has had two main episodes: one which formed the oldest populations and lasted until 8.5 Gyr ago, and a recent one that started around 3 Gyr ago. No\\\"el et al (2008, in preparation) have also obtained accurate SFHs for the fields presented in this paper, using CMDs reaching the oldest main sequence turnoffs with good photometric accuracy. They find that two main episodes of star formation in all fields, one at old ages ($\\simeq 10$ Gyr ago) and another one at intermediate ages ($\\simeq 5$ Gyr ago), in addition to young star formation in the wing fields.  Detailed determinations of chemical abundances exist only for the youngest population of the SMC \\citep[i.e.][]{hill97,venn99,hunter07}, and its chemical enrichment history has been mainly determined from studies from its cluster system. The cluster age--metallicity relationship (AMR) has been obtained by \\citet{piatti05a} and \\citet{mighell98}, mainly from photometric indicators. An initial chemical enrichment has been found, followed by a period of relatively slow increase in the metal abundance. Clusters more metal-rich than [Fe/H]$\\geq$-1 are younger than 5 Gyr. Since then, the metallicity has again increased until now. On average, the SMC is more metal-poor than the LMC. The very recent work by \\cite{idiart07}, which have obtained chemical abundances in a sample of SMC planetary nebulae, has found a similar result, with the exception that the chemical enrichment episode at a very early epoch is not observed.  In the present work we  focus on obtaining stellar metallicities of individual red giant branch (RGB) stars in the field population of the SMC from Ca II triplet (hereafter CaT) spectroscopy. These stars have been selected in 13 fields distributed at different positions in the SMC ranging from $\\sim$1\\arcdeg\\@ to $\\sim$4\\arcdeg\\@ from its center. Deep photometry of these fields has been presented in Paper I . The procedure followed to select the targets is explained in Section \\ref{targetselection}. The data reduction is discussed in Section \\ref{datareduction}. The radial velocities of the stars in our sample are obtained in Section \\ref{radialvelocities}. In Section \\ref{cat} we discuss the calculation of the CaT equivalent widths and the determination of metallicities. Section \\ref{agedetermination} presents the method used to derive ages for each star by combining the information on their metallicity and position on the CMD. The analysis of the data is presented in Section \\ref{analysis}. The metallicity distribution of each field and the possible presence of a metallicity gradient is discussed in Section \\ref{analisismetallicitydistribution}. The derived AMRs for each field are presented in Section \\ref{agemetallicity}. The main results of this paper are discussed in Section \\ref{discursion}. ",
        "conclusions": "\\label{discursion}  Using CaT spectroscopy, we have derived stellar metallicities for a large sample of RGB field stars in 13 regions of the SMC situated at different galactocentric distances and positions angles. We have found a radial metallicity gradient, which is most evident for those fields situated toward the South, where we covered a large galactocentric radius. For a given galactocentric distance, the mean metallicities for fields situated at different position angles are very similar. The inner fields have a mean metallicity of [Fe/H] $\\sim-1$, which is similar to that of the cluster metallicity distribution.   We have obtained the AMR of each field from the combination of metallicities, derived from CaT spectroscopy, and the position of stars in the CMD. All fields have similar AMRs, which are also similar to the cluster\ufffds one \\citep{piatti05a}. All show a rapid initial increase of metallicity, followed by a very slow chemical enrichment period. A new relatively fast chemical enrichment episode is observed in the last few Gyrs in the fields within $\\sim$2\\arcdeg\\@ of the center with enough young stars to sample it. From the information on the AMRs, we conclude that coeval stars have the same metallicity everywhere in the SMC. The observed metallicity gradient is therefore related to an age gradient, because the youngest stars, which are also the most metal-rich, are concentrated in the central regions of the galaxy.  In a forthcoming paper we will try to reproduce the observed AMR with chemical evolution models using accurate SFRs, as a function of time, which are being derived by our group in each field (No\\\"el et al. 2008, in preparation)."
    },
    "0806/0806.4186_arXiv.txt": {
        "abstract": "The presence of black holes (BHs) at the centers of dwarf elliptical galaxies (dEs) has been argued both theoretically and observationally. Using archival HST/WFPC2 data, we found the Virgo cluster dwarf elliptical galaxy VCC128 to harbor a binary nucleus, a feature that is usually interpreted as the observable signature of a stellar disk orbiting a central massive black hole. \\citet{db06} estimated its mass M$_\\bullet \\sim 6 \\times 10^6-5\\times 10^7$~M$_\\odot$. One of the most robust means of verifying the existence of a BH is radio continuum and/or X-ray emission, however because of the deficiency of gas in dEs, radio continuum emission forms the best option here. We have tried to detect the X-band radio emission coming from the putative black hole in VCC128 when it accretes gas from the surrounding ISM. While we made a positive 4$\\sigma$ detection of a point source 4.63$''$ south-west of the binary nucleus, no statistically significant evidence for emission associated with the nuclei themselves was detected.  This implies either that VCC128 has no massive central black hole, which makes the nature of the binary nucleus hard to explain, or, if it has a central black hole, that the physical conditions of the ISM (predominantly its density and temperature) and/or of the surrounding accretion disk do not allow for efficient gas accretion onto the black hole, making the quiescent black hole very hard to detect at radio wavelengths. ",
        "introduction": "The masses of the massive central black holes observed in many galaxies exhibit a variety of scaling relations, such as the M$_\\bullet - \\sigma$ relation between black-hole mass, M$_\\bullet$, and central velocity dispersion, $\\sigma$ \\citep{g00,fm00} and the $v_{\\rm circ} - \\sigma$ relation between circular velocity, $v_{\\rm circ}$, and central velocity dispersion \\citep{f02,b03,pi05,bfg06}. The latter can be interpreted, via the M$_\\bullet - \\sigma$ relation, as a relation between the mass of the central black hole and the total mass of the host galaxy. These correlations suggest a strong coupling between the formation of massive central black holes and the formation of galaxies.  Literature abounds with scenarios for producing massive central black holes. Models designed to explain the properties of QSOs in massive galaxies typically produce central seed black holes with a minimum mass of M$_\\bullet \\gtrsim 10^6$M$_\\odot$, e.g. through the collapse of a massive central gaseous disk or a supermassive stellar object \\citep{lr94,h98,sr98}. This seed can grow even more massive by feeding from a surrounding gas disk. These models are motivated by the necessity of producing supermassive black holes very rapidly after the Big Bang and therefore may be biased to large seed masses.  The gravitational collapse of a relativistic star cluster, supposedly born during a starburst, might produce black hole seeds with masses up to $10^4$~M$_\\odot$ \\citep{s04}. Scenarios for growing less massive black holes, aptly called intermediate-mass black holes, also exist in the literature \\citep{mr01,mh02,pm02}. Their applicability seems to be restricted to very dense systems, such as, for instance, globular clusters or galactic nuclear star clusters, because they rely on dynamical friction to sink stellar-mass seed black holes, possibly originating from Pop{\\sc iii} objects as suggested by \\citet{mr01}, to the center of the system on short enough timescales.  The picture becomes even more complicated if the fact that galaxies grow via a process of hierarchical merging is taken into account. Recently, the impact of the co-evolution of central black holes and their host galaxies has been incorporated in semi-analytical models (SAMs) of galaxy evolution. E.g., \\citet{kh00}, \\citet{m07} and \\citet{ka05} allow massive central black holes to grow by two processes:~{\\em (i)} the influx of gas into the galaxy center due to the starburst and the strongly non-axisymmetric forces that accompany a galaxy-galaxy merger, and {\\em (ii)} merging of the black holes after the host galaxies have merged. These SAMs can reproduce the aforementioned scaling relations between the black-hole mass and the properties of the host galaxy. Using a high-resolution small-volume cosmological N-body/SPH simulation, \\citet{mhbs07} show they can grow M$_\\bullet \\sim 10^6$~M$_\\odot$ central black holes in Milky-Way size halos by a redshift of $z \\sim 6$ by allowing central black holes to merge if their host galaxies merge and by gas accretion during mergers.  However, when the members of a binary black-hole system merge, there can be a considerable beaming of the gravitational radiation emitted during the final stages of the plunge. This results in a recoil of the newly formed black hole. Fully relativistic calculations of black-hole binary mergers have shown that this recoil velocity can be as high as a few hundred~km~s$^{-1}$, depending on the holes' spin, orientation, and orbital eccentricity \\citep{p07}. This could potentially hamper the retention of black holes in merging low-mass protogalaxies and by extension the build-up of supermassive black holes.  Clearly, the search for massive central black holes in low-mass galaxies is of considerable interest for the study of the symbiosis of host galaxies and their central black holes and for constraining the minimum mass of seed black holes. Nuclear activity in dwarf galaxies has turned up a number of central black holes with estimated masses in the range M$_\\bullet \\sim 10^4-10^6$~M$_\\odot$ \\citep{fh03,b04,gh04}. Dynamical modeling of ground-based stellar kinematics of dwarf elliptical galaxies ruled out the presence of central black holes with masses M$_\\bullet \\gtrsim 10^7$~M$_\\odot$ \\citep{g02}.  \\begin{figure*} \\includegraphics[width=18cm]{f1.eps} \\caption{Contourplot of the VLA radio observations on top of a V-band image of the center of VCC128. The figure shows the $12''\\times 12''$ region centered on the binary nucleus, which is just resolved in this image. The white contours correspond in increasing order to the contourlevels 3$\\sigma$, 4$\\sigma$, 5$\\sigma$, \\ldots with $1\\sigma$ = $30\\mu$Jy, while the gray contours correspond to $-3\\sigma$. The detected source, at $4\\sigma$ significance, can be seen SW of the center (encircled). The synthesized beam is shown in the bottomleft corner. A V-band image of entire VCC128 is shown in the topright corner. \\label{ima_128}} \\end{figure*}  Obtaining and modeling high-quality stellar kinematics of faint dwarf galaxies is, unfortunately, a very time-consuming way of searching for inactive central black holes. This prompted us to look for a more efficient, photometry-based method. Using archival HST images, we found the Virgo Cluster dwarf elliptical galaxy (or dE) VCC128 to harbour a binary nucleus \\citep{db06}. At the time, only two other galaxies with binary nuclei were known: M31 and NGC4486B \\citep{l93,l96,l05}. In both cases, the binary nucleus was interpreted as the observable signature of a stellar disk orbiting a central massive black hole \\citep{t95,l96}. This suggested that the binary nucleus in VCC128 might also be the hallmark of a central massive black hole.  However, the presence of a central massive black hole in VCC128 still needs to be established by independent means. An observationally rather undemanding method for detecting massive central black holes is through the radio emission they produce when they accrete gas. In Section \\ref{rce}, we give a short discussion of this method, followed by an account of our radio continuum observations of VCC128 with the Very Large Array (VLA) in Section \\ref{obs}. We end by presenting our results and conclusions in Section \\ref{results}. ",
        "conclusions": "\\label{results}  We created a radio continuum map of the central $12''\\times 12''$ region of VCC128 with an rms noise of $30~\\mu$Jy/beam. This map is presented in Fig. \\ref{ima_128}, overplotted on a V-band image of VCC128 obtained with FORS1 mounted on the VLT (ESO program 079.B-0632(B)). The data reduction and analysis of the optical imaging of VCC128 will be reported in detail elsewhere. A 4$\\sigma$ detection of a point source was made at R.A. $12h~14m~59.66s$, DEC. $+9^\\circ~ 33'~50.94''$ (J2000). This is 4.63$''$ SW of the binary nucleus. The source has a flux density of $1.66 \\pm 0.05$~mJy. Unfortunately, no emission was found at the optical position of the nucleus.  Since we do not detect any radio emission coming from the center of VCC128, we can only place a $S_\\nu \\approx 90~\\mu$Jy $3\\sigma$ upper limit on the radio continuum emission of the putative central black hole. Assuming a flat radio spectrum, this corresponds to a $3\\sigma$ upper limit on the 5~GHz radio luminosity of $L_{\\rm cont} = \\nu L_\\nu \\approx 1.5 \\times10^{35}$~erg~s$^{-1}$. This is lower than the expected radio luminosity of an actively accreting BH with a mass M$_\\bullet \\gtrsim 10^6$~M$_\\odot$.  We conclude that either {\\em (i)} VCC128 does not contain a central massive black hole or, alternatively, {\\em (ii)} VCC128 might contain a central massive black hole with a mass of the order of a few $10^6$~M$_\\odot$ but the gas near the galaxy center is either too hot or too rarefied or conditions near the black hole are unsuitable for efficient accretion. In the former case, one needs to rethink the nature of the binary nucleus. As detailed in \\citet{db06}, all other explanations for the binary nucleus other than it being the observational signature of a stellar disk rely on rather contrived chance configurations."
    },
    "0806/0806.4523_arXiv.txt": {
        "abstract": "%  The experimental technique of fluorescence light observation is used in current and planned air shower experiments that aim at understanding the origin of ultra-high energy cosmic rays. In the fluorescence technique, the geometry of the shower is reconstructed from the correlation between arrival time and incident angle of the signals detected by the telescope. The calculation of the expected light arrival time used so far in shower reconstruction codes is based on several assumptions. Particularly, it is assumed that fluorescence photons are produced instantaneously during the passage of the shower front and that the fluorescence photons propagate on a straight line with vacuum speed of light towards the telescope. We investigate the validity of these assumptions, how to correct them, and the impact on reconstruction parameters when adopting realistic conditions. Depending on the relative orientation of the shower to the telescope, corrections can reach 100~ns in expected light arrival time, 0.1$^\\circ$ in arrival direction and 5~g~cm$^{-2}$ in depth of shower maximum. The findings are relevant also for the case of ``hybrid'' observations where the shower is registered simultaneously by fluorescence and surface detectors. ",
        "introduction": "\\label{sec-intro}   Understanding the origin and nature of ultra-high energy (UHE) cosmic rays above $10^{19}$~eV is a major challenge of astroparticle physics~\\cite{reviews}. These cosmic rays are studied by detecting the atmospheric showers they initiate. Current and planned air shower experiments~\\cite{auger,hires-gzk,telarray,euso,owl} use the technique of fluorescence light observation: shower particles deposit energy in the atmosphere through ionisational energy loss. Part of this energy (of order $10^{-4}$) is emitted isotropically at near-UV wavelengths in de-excitation processes. These fluorescence photons can be detected by appropriate telescope systems operating in clear nights. Typically, pixel cameras with 25$-$100~ns timing resolution are used, where an individual pixel covers a field of view of about 1$-$1.5$^\\circ$ in diameter (see e.g.\\ Ref.\\ \\cite{auger}). The signal (light flux per time) is registered as a function of the viewing direction of the pixels.  The first step to reconstruct the primary parameters of an observed air shower is given by the determination of the shower geometry. An accurate geometry reconstruction is, for instance, decisive for directional source searches; but it is also a prerequisite for reconstructing other important shower parameters such as the primary energy or the depth of shower maximum. We note that also the shower energies obtained from Auger ground array data are calibrated by the fluorescence telescopes~\\cite{roth_icrc07}.  The determination of the shower geometry is commonly performed in two steps in the fluorescence technique~\\cite{flyseye}. First, the ``shower-detector-plane'' (SDP) is determined as the plane spanned by the (signal-weighted) viewing directions of the triggered camera pixels (Fig.~\\ref{fig-sdp}). Next, the geometry of the shower within this SDP is reconstructed based on the correlation between arrival time of the signals and viewing angle of the pixels projected into the SDP. The measured time-angle correlation is compared to the one expected for different shower geometries, and the best-fit geometry is determined. For the calculation of the expected time-angle correlation, the following function is in use (following e.g.\\ Ref.\\ \\cite{bunner,flyseye,sokolsky}):  \\begin{equation} t_{i}  =  t_{0}+\\frac{R_{p}}{c_{\\rm vac}}\\tan \\left( \\frac{\\chi _{0}-\\chi _{i}}{2} \\right) \\label{arrival_old} \\end{equation}  where $t_i$ is the arrival time of the photons at camera pixel $i$ (usually, a signal-weighted average arrival time is taken from the time sequence observed in a pixel), $t_0$ is the time at which the shower axis vector passes the closest point to the telescope at a distance $R_p$, $c_{\\rm vac}$ is the vacuum speed of light, $\\chi_0$ is the angle of incidence of the shower axis within the SDP, and $\\chi_i$ is the viewing angle of pixel $i$ within the SDP (see also Fig.~\\ref{fig-sdp}). Comparing the expected $t_i$$-$$\\chi_i$ correlation to the observed one ($i=1...n$ for $n$ triggered pixels), the best-fit parameters $R_p$, $t_0$ and $\\chi_0$ in Eq.~(\\ref{arrival_old}) are found by a $\\chi^2$-minimization. Together with the SDP derived previously, the shower geometry is then fully determined and can also be expressed in terms of shower impact point, arrival direction, and ground impact time.  \\begin{figure}[t] \\begin{center} \\includegraphics[width=10.0cm]{fig1.eps} \\caption{ Sketch of the shower geometry and quantities used in the derivations. } \\label{fig-sdp} \\end{center} \\end{figure}  Eq.\\ (\\ref{arrival_old}) is derived as follows. Assuming the fluorescence light to be emitted by a point-like object moving at $c_{\\rm vac}$  along the shower axis vector, the shower propagation time $\\tau_{{\\rm shower},i}$ from point $S_i$ to the point at reference time $t_0$ on the shower axis (Fig.~\\ref{fig-sdp}) can be expressed as  \\begin{equation} \\tau_{{\\rm shower},i} = \\frac{R_p}{c_{\\rm vac}\\cdot \\tan(\\chi_0 - \\chi_i)}~. \\label{4_shower} \\end{equation}  Next, assuming the fluorescence photons to propagate on straight lines with $c_{\\rm vac}$, the light propagation time $\\tau_{{\\rm light},i}$ from $S_i$ to the telescope is  \\begin{equation} \\tau_{{\\rm light},i} = \\frac{R_p}{c_{\\rm vac}\\cdot \\sin(\\chi_0 - \\chi_i)}~. \\label{4_prop} \\end{equation}  With Eqs.\\ (\\ref{4_shower}) and (\\ref{4_prop}), and assuming an instantaneous emission of the fluorescence photons at $S_i$, the expected arrival time $t_i$ (relative to the time $t_0$ of closest approach of the shower to the telescope) of fluorescence photons at a pixel viewing at an angle $\\chi_i$ becomes  \\begin{eqnarray} t_{i} & = & t_{0} - \\tau_{{\\rm shower},i} + \\tau_{{\\rm light},i} \\nonumber \\\\ & = & t_{0} + \\frac{R_{p}}{c_{\\rm vac}} \\left( \\frac{1}{\\sin (\\chi_0 - \\chi_i)}-\\frac{1}{\\tan (\\chi_0 - \\chi_i)}\\right) \\nonumber  \\\\ & = & t_{0}+\\frac{R_{p}}{c_{\\rm vac}}\\tan \\left( \\frac{\\chi _{0}-\\chi _{i}}{2} \\right) \\label{4_arrival_old} \\end{eqnarray}  which equals Eq.\\ (\\ref{arrival_old}).  Thus, the derivation of Eq.\\ (\\ref{arrival_old}) for calculating the expected time-angle correlation is based on the following assumptions:  \\begin{itemize} \\item the spatial structure and the propagation of the shower disk can be approximated by a point-like object moving at $c_{\\rm vac}$, \\item the fluorescence light is produced instantaneously, \\item the fluorescence light propagates with $c_{\\rm vac}$, \\item the fluorescence light propagates on a straight line. \\end{itemize}  In this article, we investigate the validity of these assumptions. The impact of the corrections on reconstruction parameters is studied. The results are relevant both for observations with fluorescence telescopes alone and for ``hybrid'' observations where the shower is registered by fluorescence and surface detectors.  The plan of the paper is as follows. In Section~\\ref{sec-effects}, we discuss the various assumptions and their corrections individually. In Section~\\ref{sec-impact}, the impact on shower reconstruction is studied. Conclusions are given in Section~\\ref{sec-concl}. ",
        "conclusions": "\\label{sec-concl}  The assumptions used in the ``classical'' function of Eq.~(\\ref{arrival_old}) for reconstructing the shower geometry from fluorescence light observations were investigated. The finite shower thickness leads to an energy deposition in air by secondaries which is delayed, compared to the shower front, by about 5$-$6~ns (with some dependence on the specific light collection algorithm employed). The emission of fluorescence light is further delayed due to the finite lifetime of the transitions which, due to quenching, is altitude dependent. Typical values are a few nanoseconds up to 20~km height, and $>$15~ns for heights above 30~km. The propagation speed of light is reduced by the index of refraction of air. The delay, compared to a propagation with vacuum speed of light, depends on the locations of emission point and telescope, and can exceed 20$-$25~ns. Finally, another effect of refraction is the bending of light, which also depends on the locations of emission point and telescope. Angular differences between the apparent and actual emission point of 0.02$^{\\circ}$ can occur, which may correspond to time shifts of several 10~ns. This effect can also lead to a slight tilt of the SDP.  All these corrections can be considered as geometrical ones, i.e.\\ they are independent of specific properties of the individual showers other than their geometry. The corrected function for geometry reconstruction is given in Eq.~(\\ref{arrival_new}). Compared to the previous approach, which assumed maximum propagation speed of both light and particles as well as no other delays, the effects of delayed timing (including the effect of bending of light) accumulate. In total, differences of up to $\\sim$100~ns in estimated light arrival time are possible. Air shower experiments with comparable, or better, time resolution should take these effects into account. This refers both to data reconstruction and to implementing these effects in the shower$-$detector simulation. In terms of overall shower reconstruction parameters, corrections are typically 0.03$-$0.05$^\\circ$ in arrival direction (with a systematic trend of overestimating the zenith angle when neglecting the effect), $\\simeq$ 0.5$-$1\\% in energy and 2$-$3~g~cm$^{-2}$ in $X_{\\rm max}$, but may in some cases exceed 0.1$^\\circ$ and 5~g~cm$^{-2}$. This is to be compared to typical reconstruction accuracies of $\\sim$0.6$^\\circ$ (directional resolution) \\cite{bonifazi05} and $\\sim$11~g~cm$^{-2}$ (systematic $X_{\\rm max}$ uncertainty) \\cite{unger07} in case of Auger hybrid events.  The increase in computing time for event reconstruction is modest, particularly when applying the corresponding corrections only when approaching convergence in the minimization process (increase of $\\sim$20\\% or less, depending on implementation). Some of the effects investigated in this work might be relevant also for shower detection techniques other than fluorescence telescope observations at ultra-high energy, e.g.\\ Cherenkov light observations of air showers.    {\\it Acknowledgements:}  We would like to thank our Colleagues from the Pierre Auger Collaboration for many fruitful discussions, in particular Fernando Arqueros, Jose Bellido, Bruce Dawson, Philip Wahrlich and the members of the Auger group at the University of Wuppertal. Figs.\\ \\ref{sim_data} and \\ref{sim_dtheta} were produced using Auger software packages~\\cite{offline,offline2}. This work was partially supported by the German Ministry for Research and Education (Grant 05 CU5PX1/6)."
    },
    "0806/0806.3303_arXiv.txt": {
        "abstract": "{The recently launched  satellite, \\emph{Fermi Gamma-ray Space Telescope\\/}, is expected to find out if cosmic-ray (CR) protons are generated from supernova remnants (SNRs), especially  RX~J1713.7$-$3946, by observing the GeV-to-TeV $\\gamma$-rays. The GeV emission is thought to be bright if the TeV emission is hadronic, i.e., of proton origin, while dim if leptonic. } {We reexamine the above view using a simple theoretical model of nonlinear acceleration of particles to calculate the gamma-ray spectrum of Galactic young SNRs. } { } {If the nonlinear effects of CR acceleration are considered, it may be impossible to distinguish  the evidence of proton acceleration from leptonic  in the $\\gamma$-ray spectrum of Galactic young SNRs like RX~J1713.7$-$3946. On the other hand, future  km$^3$-class neutrino observations will likely find a clear evidence of the proton acceleration there. } {} ",
        "introduction": "Recently, {\\it Fermi Gamma-ray Space Telescope (Fermi)}\\footnote {http://fermi.gsfc.nasa.gov/}, observing GeV $\\gamma$-ray photons, has been launched. The GeV $\\gamma$-ray observations with {\\it Fermi} are expected to identify the accelerators of Galactic cosmic-ray (CR) protons whose energy extends up to the ``knee'' energy ($\\approx10^{15.5}$~eV). At present, the most probable candidate for the CR accelerator is a young supernova remnant (SNR). Since the detections of synchrotron X-rays  in some SNRs show evidence of electron acceleration \\citep{koyama95}, the current unsolved issue is whether the SNRs produce high-energy protons or not. TeV $\\gamma$-ray observations are important to address this problem. So far, TeV $\\gamma$-rays  have been detected from several young SNRs \\citep{enomoto02,aharonian04,aharonian05c,katagiri05}. They arise from either leptonic (CMB photons up-scattered by high energy electrons) or hadronic ($\\pi^0$-decay photons generated via accelerated protons) processes, and it is generally difficult to separate these processes using  only  the TeV energy band; the study of wide-band, GeV-to-TeV spectra is necessary.    RX~J1713.7$-$3946 (hereafter RXJ1713) is a representative SNR from which bright TeV $\\gamma$-rays have been detected. The H.E.S.S. experiment measured the TeV spectrum and claimed that its shape was better explained by the hadronic model \\citep{hess06,hess07}. So far, compared with other young SNRs, the TeV $\\gamma$-ray spectrum of RXJ1713 is the most precisely measured and the energy coverage is wide, from 0.3 to 100 TeV, so that we can obtain the best constraints on theoretical models.   Recently, time variation of synchrotron X-rays was discovered in RXJ1713 \\citep{uchiyama07}. If the variation timescale is determined from the synchrotron cooling of X-ray emitting electrons, the magnetic field is estimated to be $B\\sim$~mG. If so,  the leptonic, one-zone emission model \\citep[e.g.,][]{aharonian99} cannot explain the TeV-to-X-ray flux ratio, supporting the hadronic origin of TeV $\\gamma$-rays. It should be noted that the amplified magnetic field is theoretically expected \\citep[e.g.,][]{lucek00,giacalone07}. In this case, according to the standard diffusive shock acceleration theory, the maximum energy of accelerated protons is estimated as \\citep{aharonian99} \\begin{equation} E_{{\\rm max},p}=8\\times10^3 \\frac{B_{\\rm mG}t_3}{\\eta_g} \\left(\\frac{v_s}{4000~{\\rm km}~{\\rm s}^{-1}}\\right)^2 {\\rm TeV}~~, \\end{equation} which can be comparable to the knee energy. Here, $B_{\\rm mG}$, $v_s$, $t_3$, and $\\eta_g$ are the magnetic field strength in units of mG, the shock velocity, the age of the SNR in units of $10^3$~yr, and the gyrofactor, respectively.   However, at present, there are several issues to be addressed, as the above picture on RXJ1713 is not yet proved. First, if $B\\sim{\\rm mG}$ and TeV emission is hadronic, then in order to explain the measured flux of radio synchrotron emitted by primary electrons, the electron-to-proton ratio at the SNR should be anomalously small, $K_{ep}\\sim10^{-6}$ \\citep{uchiyama03,butt08}, which is far below the observed value at the earth and estimated values in the nearby galaxy \\citep{katz08}. This might be resolved if the electrons are accelerated in the later stages of SNR evolution, when the value of $K_{ep}$ is different from the present value \\citep{tanaka08}, although further discussions are necessary. Second, the hadronic scenario may be inconsistent with the molecular cloud (MC) observations \\citep{fukui03}. RXJ1713 is surrounded by MCs, which might suggest collision with them and high target number density. If the TeV $\\gamma$-rays are hadronic, such a region should be brighter than observed \\citep{plaga07}.   In addition, if the measured  width of the synchrotron X-ray  filaments at the shock front of SNRs  is determined by the synchrotron cooling effect \\citep{uchiyama03,vink03,bamba03b,bamba05a,bamba05b}, the magnetic field is independently estimated as $B\\approx0.1$~mG \\citep{parizot06}, which is an order of magnitude smaller than that estimated by \\citet{uchiyama07}. Also, the cutoff energy of TeV $\\gamma$-ray spectrum is low, so that in the one-zone hadronic scenario $E_{{\\rm max},p}$ is estimated as 30--100~TeV \\citep{villante07}, which is approximately two orders of magnitude lower than the knee energy. If $E_{{\\rm max},p}<100$~TeV and $B\\approx1$~mG, then Eq.~(1) tells us $\\eta_g\\ga80$, implying far from the ``Bohm limit'' ($\\eta_g\\approx1$) which is inferred from the X-ray observation \\citep{parizot06,yamazaki04} or expected theoretically \\citep{lucek00,giacalone07}. This statement is recast if we involve recent results of X-ray observations. The precise X-ray spectrum of RXJ1713 is revealed, which gives $v_s=3.3\\times10^8\\eta_g^{1/2}$~cm~s$^{-1}$ \\citep{tanaka08}. Then, Eq.~(1) can be rewritten as $E_{{\\rm max},p}=5\\times10^3B_{\\rm mG}t_3$~TeV. Hence, in order to obtain $E_{{\\rm max},p}<100$~TeV, we need $B\\la20~\\mu$G  in the context of the hadronic scenario of TeV $\\gamma$-rays. One might think that the volume filling factor of the region with $B\\approx1$~mG is small and that the average field strength is smaller, e.g., $B\\approx0.1$~mG. However, even in this case, $E_{{\\rm max},p}$ is more than 100~TeV, which contradicts the observed $\\gamma$-ray spectrum beyond 10~TeV.  In these circumstances, {\\it Fermi} will give us important information on the $\\gamma$-ray emission mechanism. So far, the GeV emission has been thought to be bright if the TeV emission is hadronic, while dim if leptonic. However, this argument is not so straightforward if the nonlinear model of CR acceleration is considered. In the next section, we calculate the photon spectrum using a simple semi-analytic model taking into account nonlinear effects. Indeed, we show that in a certain case, the hadronic emission spectrum in the GeV-to-TeV band is  similar to the leptonic one. ",
        "conclusions": "\\label{sec:conclusion}  Some  models predict relatively bright GeV $\\gamma$-rays compared with those considered above. If the CR back-reaction effect on the particle spectrum is small (inefficient-acceleration case), then the energy spectral index of protons is $s_p\\approx2$, and the $\\pi^0$-decay $\\gamma$-ray emission shows a roughly flat $\\nu F_\\nu$-spectrum, $\\nu F_\\nu\\propto\\nu^0$, in the GeV--TeV band (model~IIa/IIIa in Table~\\ref{table:1}). The predicted flux is marginally consistent with the EGRET upper limit, $\\nu F_\\nu\\approx5\\times10^{-11}$~erg~s$^{-1}$cm$^{-2}$ at 1--10~GeV \\citep{hess06,hartman99}. On the other hand, if the magnetic field is moderately strong, the synchrotron cooling effect causes steepening of the electron spectrum over a wide energy range --- typically $s_e\\approx3$ \\citep[e.g., see \\S~19.3, eq.~(19.16) of][]{longair94}. In this case, leptonic IC emission in the GeV--TeV band again shows a nearly flat $\\nu F_\\nu$-spectrum (model~IIb in Table~\\ref{table:1}). Therefore, these $\\gamma$-ray emission models (IIa/IIIa and IIb) cannot be distinguished. This has been  discussed in \\citet{ellison07}, where $B\\approx60~\\mu$G.   In summary,  it may be difficult to differentiate between hadronic and leptonic  emission by the spectral shape of the GeV-to-TeV $\\gamma$-ray emission of Galactic young SNRs like RXJ1713 (Table~\\ref{table:1}). As shown in this paper, when the GeV $\\gamma$-ray flux is  relatively low (e.g., $\\nu F_\\nu\\propto\\nu^{0.5}$), both an efficient acceleration model with hadronic $\\gamma$-ray emission (model~Ia) and a leptonic, weak magnetic-field model with inefficient acceleration (model~IIIb) may give similar spectral shapes. On the other hand, as already pointed out in \\citet{ellison07}, when the GeV emission is relatively bright (e.g., $\\nu F_\\nu\\propto\\nu^0$), one may not be able to distinguish the hadronic model in the inefficient case (models~IIa/IIIa) from the leptonic one with a moderately strong magnetic field  (model~IIb). This conclusion may, at least qualitatively,  be applicable to other young SNRs emitting TeV gamma-rays, such as RX~J0852.0$-$4622 \\citep{katagiri05,aharonian05c}. {\\it Fermi} will likely provide us with rich information on the emission mechanism of RXJ1713 and other young SNRs. However, one should only draw  conclusions with great care, even in the {\\it Fermi} era. Probably,   neutrino observation with km$^3$-class detectors such as IceCube~\\citep{Achterberg:2007qp} or KM3NeT~\\citep{Kappes:2007ci} will finally resolve  the problem \\citep{crocker02,halzen02,kistler06,vissani08,huang08,halzen08}. As shown in Fig.~\\ref{fig:neutrino}, if the observed TeV $\\gamma$-ray emission is hadronic, then  the expected neutrino spectrum at the source is above the atmospheric neutrino background at around 5 -- 10~TeV, which may become `the smoking gun' of proton acceleration in Galactic young SNRs.      \\begin{figure} \\includegraphics[width=20pc]{0845fig5.eps} \\caption{ $\\nu F_\\nu$-spectra of TeV $\\gamma$-rays (solid line) and $\\mu$-neutrinos (dashed line) calculated within the hadronic two-zone model in the efficient acceleration case. The solid line is the same as in Fig.~2. We have averaged the vacuum oscillation effects among neutrinos. The dotted line shows the daily averaged atmospheric  neutrino flux expected in KM3NeT~\\citep{Kappes:2007ci,kappes07}.  } \\label{fig:neutrino} \\end{figure}"
    },
    "0806/0806.0520_arXiv.txt": {
        "abstract": "{} {The separation of foreground contamination from cosmic microwave background (CMB) observations is one of the most challenging and important problem of digital signal processing in Cosmology. In literature, various techniques have been presented, but no general consensus about their real performances and properties has been reached. This is due to the characteristics of these  techniques that have been studied essentially through numerical simulations based on semi-empirical models of the CMB and the Galactic foregrounds. Such models often have different level of sophistication and/or are based on different physical assumptions (e.g., the number of the Galactic components and the level of the noise). Hence, a reliable comparison is difficult. What actually is missing is a statistical analysis of the properties of the proposed methodologies. Here, we consider the {{\\it Internal Linear Combination}} method (ILC) which, among the separation techniques, requires the smallest number of {{\\it a priori}} assumptions. This feature is of particular interest in the context of the CMB polarization measurements at small angular scales where the lack of knowledge of the polarized backgrounds represents a serious limit.} {The statistical characteristics of ILC are examined through an analytical approach and the basic conditions are fixed in a way to work satisfactorily. A comparison with the FastICA implementation of the {{\\it Independent Component Analysis}} (ICA) method, one of the most celebrated techniques for blind signal separation, is made.} {ILC provides satisfactory results only under rather restrictive conditions. This is a critical fact to take into consideration in planning the future ground-based observations (e.g., with ALMA) where, contrary to the satellite experiments, there is the possibility to have a certain control of the experimental conditions.} {} ",
        "introduction": "The experimental progresses in the detection of cosmological and astrophysical emissions require a parallel development of data analysis techniques in order to extract the maximum physical information from data. An example of very interest is represented by signals that are the mixture of the emission of distinct physical mechanisms. The study of the underlying physical processes needs the separation of the different components that contribute to the observed signals. This is the case of the Cosmic Microwave Background (CMB) observations where there is the necessity to separate the CMB from diffuse foregrounds originated by our own Galaxy. In this context, an extensive literature is available and many approaches have been proposed \\citep[for a review, see][]{del07}. However, no general consensus about their real performances and properties has been reached.  This is due to the approach followed to determine the characteristics of these  techniques that is based on numerical simulations that make use of semi-empirical models of the CMB and the Galactic foregrounds. The point is that such models often have different level of sophistication and/or are based on different physical assumptions (e.g., the number of the Galactic components and the level of the noise). Hence, a reliable comparison is difficult. A trustworthy assessment of the real capabilities of such methodologies should require a rigorous analysis of their theoretical statistical characteristics independently of the specific context where they are applied. However, at least in our knowledge, in literature nothing has never been presented in this sense \\citep[however, see][~for a discussion concerning the power spectrum]{sah07}. Things become even more serious in the case of CMB polarization measurements, where the available {\\it a priori} information is quite limited and the use of {\\it blind} separation techniques obligatory \\citep{sti06}.  A situation as this one is clearly unsatisfactory. This also because in a near future some innovative ground-based experiments are planned for polarization observations at extraordinary high spatial resolution as with the Atacama Large submmillimetre/millimetre array (ALMA). One important advantage of these experiments is that, contrary to the satellite observations, they will allow a certain control of the experimental conditions. Hence, a fully exploitation of the capabilities of the instruments implies a careful preparation of the observations in such a way the obtained data are of sufficient quality to permit an effective application of the chosen separation methodology. For this reason, in this work we start exploring the capabilities of algorithms aimed at a careful subtraction of the foreground sources when the amount of available {\\it a priori} information is limited (an expected situation for polarization observations). In particular, we consider one of the most used approaches to the separation of different emissions, the {\\it internal linear combination} method (ILC), which was adopted for instance in the reduction of the data from the Wilkinson Microwave Anisotropy Probe (WMAP) satellite for CMB observations \\citep{ben03}. Among the separation techniques, ILC requires the smallest number of {\\it a priori} assumptions. For the reason presented above, we do not perform an application to any astrophysical dataset here, but rather we study the general properties of this technique in order to fix the conditions under which it can be expected to produce reliable results. We also make a critical comparison with the architecture and capabilities of the FastIca implementation of the ICA approach, which is the other main handle to the separation when little {\\it a priori} information is available \\citep{sti06}.  In the following, the available data are assumed in form of $N_o$ maps, taken at different frequencies, containing $N_p$ pixels each. More precisely, if $\\Si$ provides the value of the $p$th pixel for a map obtained at channel ``$~i~$'' \\footnote{In the present work, $p$ indexes pixels in the classic spatial domain. However, the same formalism applies if other domains are considered as, for example, the Fourier one.}, our starting model is: \\begin{equation} \\label{eq:image} \\Si = \\Scmbi + \\Sgal + \\Nmthi(p) \\end{equation} where $\\Scmbi$, $\\Sgal$ and $\\Nmthi(p)$ are the contributions due to the CMB, the diffuse Galactic foreground and the experimental noise, respectively.  Although not necessary for later arguments, it is assumed that all of these contributions are representable by means of stationary random fields. At least locally, in many experimental situations this is an acceptable approximation. If not, in any case it is often made since it permits a statistical treatment of the problem of interest and the results can be used as benchmark in the analysis of more complex scenarios. In the present context, this assumption holds on small patches of the sky. Finally, without loss of generality, for easiness of notations the random fields are supposed the realization of zero-mean spatial processes. In the present work the contribution of non-diffuse components (e.g., due to SZ cluster, point-sources, \\ldots) is not considered and it is supposed to be already removed through other methodologies.  The paper is organized as follows: in Sec.~\\ref{sec:formalization} the relevant analytical formulas are introduced. In Sec.~\\ref{sec:ILC} the statistical framework of ILC in noiseless observations is presented and compared with that derived for FastICA. Simulations are also presented. In Sec.~\\ref{sec:ILCnoise} the case of noisy observations is considered. Finally, the conclusions are drawn in Sec.~\\ref{sec:conclusions}. ",
        "conclusions": "\\label{sec:conclusions}  The arguments presented in the previous sections show that for a safe use of the ILC estimator some conditions have to be satisfied. In particular: \\begin{enumerate} \\item The observations have to cover a sky area much wider than the spatial scale (correlation length) of the observed maps in such a way to allow an accurate estimate of the cross-covariance matrix $\\Cb_{\\Sb}$; \\item Model~(\\ref{eq:model})-(\\ref{eq:Amatrix}) has to hold. In particular, ILC cannot be expected to produce satisfactory results if some of the Galactic templates depend on the observing channels. This point can be realized if a Galactic channel-dependent template $\\Sbb_j$ is thought as the sum of $N_j$ channel-independent templates. In this case, model~(\\ref{eq:model})-(\\ref{eq:Amatrix}) still holds but with an effective number of Galactic components that now is $N_c^* = \\sum_{j=1}^{N_c} N_j$. Since the number of observing channel is typically rather limited, in practical applications it has to be expected that $N_o < N_c^* + 1$ and then the point below applies; \\item The condition $N_o \\ge N_c^* + 1$ has to hold, i.e. the number of the observing channels has to be equal or larger than the number of the physical components (CMB included) that contribute to form the observed maps. In the contrary case, the solution can suffer a severe distortion; \\item The level of the noise has to be rather low otherwise a severe bias can be introduced in the ILC estimator. This can be removed but at the price of a possible remarkable amplification of the noise influence. However, especially in situations of high ${\\rm SNR}$, the use of the unbiased ILC estimator can offer some advantages. \\end{enumerate} From these points it is not difficult to realize that the ILC estimator is not trivial to use. In particular, the first two points conflict each other. In fact, in the case of wide maps, model~(\\ref{eq:model})-(\\ref{eq:Amatrix}) is not applicable since on large spatial scales the Galactic templates are expected to depend on the observing channel. In this respect, a simple tests can be of help that is based on the analysis of the eigenvalues of $\\widehat \\Cb_{\\Sb}$. In fact, if $N_o \\geq N_c^* + 1$, matrix $\\widehat \\Cb_{\\Sb}$ has to be (almost) singular with a number of (almost) zero elements in the diagonal matrix $\\widehat \\Db$ of Eq.(\\ref{eq:svd}) equal to $N_o - N_c^* - 1$. Unfortunately, this test can be thwarted by the noise; because of the statistical fluctuations, the entries of $\\widehat \\Db$ that should be close to zero can take larger value.  The obvious conclusion is that, in order to obtain reliable results, the use of ILC requires a careful planning of the observations. The area of the sky to observe as well as the tolerable level of the noise are factors that have to be fixed in advance. To try an ``{\\it a posteriori}'' correction of the distortions introduced in the ILC solution by the violation of the above conditions is a quite risky operation. The question is that all of these distortions critically depend on the true solution that one tries to estimate and thereof they cannot be obtained from the data only. The alternative represented by the numerical simulations that make use of semi-empirical templates is quite risky. In particular, there is the concrete possibility to force spurious features in the final results. This holds also in the case these results are used as ``{\\it prior}'' in more sophisticated separation techniques.  A last comment regards the use ILC in the Fourier domain. Since this method provides the same result independently of the domain, it could seem that there is no particular benefit to work in the Fourier one. Actually, this can be not true if the maps have different spatial resolutions (i.e. the observing channels have different point spread functions) and/or whether a frequency-dependent separation is desired \\citep{teg03}. Although in this way it is possible to improve some properties of the separated maps (e.g., the spatial resolution), it is necessary to stress that this has a cost in the amplification of the noise level."
    },
    "0806/0806.2655_arXiv.txt": {
        "abstract": "We present optical spectra of 14 emission-line stars in M33's giant HII regions NGC 592, NGC 595 and NGC 604: five of them are known WR stars, for which we present a better quality spectrogram, eight were WR candidates based on narrow-band imagery and one is a serendipitous discovery. Spectroscopy confirms the power of interference filter imagery to detect emission-line stars down to an equivalent width of about 5 \\AA\\ in crowded fields. We have also used archival HST/WFPC2 images to correctly identify emission-line stars in NGC 592 and NGC 588. ",
        "introduction": "The overwhelming majority of stars in the Universe display absorption lines in their visible range spectrum. Emission lines in stellar spectra are therefore notable exceptions which betray fascinating properties such as the presence of unusually strong chromospheric activity, rapid rotation or in the case of interest to this paper, strong mass loss and high luminosity. Nearby giant H~{\\sc ii} regions are hosts to an interesting zoo of emission-line objects (Walborn \\& Fitzpatrick 2000), most of them being post-main sequence massive stars of notable interest for our understanding of stellar evolution at the top end of the initial mass function. They also offer important testbeds to understand the more distant, unresolved starbursts: individual stars can be counted and spectroscopically classified, allowing a direct comparison with the modelisation of the spatially integrated properties of their ionizing cluster (see Vacca et al. 1995 for such a comparison in 30 Doradus, or Bruhweiler, Miskey \\& Smith Neubig 2003 in NGC 604). Among massive stars with emission lines, those of Wolf-Rayet (WR) type are the easiest to detect and classify because of their strong and broad emission lines in the visible part of the spectrum. Thanks to surveys in Local Group galaxies (Massey \\& Johnson 1998) and improvements in theoretical models (Meynet \\& Maeder 2005), the evolutionary status of WR stars is now understood well enough to use them as diagnostics to infer the properties of starburst regions. For instance, in unresolved clusters or starburst knots of distant galaxies, the equivalent width of the `WR-bumps' are good indicators of the age and upper mass limit of the stellar population (Pindao et al. 2002).  The small spiral galaxy M33 is host to four giant H{\\sc ii} regions bright enough to have their own NGC number: NGC 604, the second most luminous starburst cluster in the Local Group; then, in decrasing order of H$\\alpha$ luminosity, NGC 595, NGC 592 and NGC 588. Despite their different galactocentric distances, these four regions have very similar metallicities, with $12 +$ log O/H $= 8.4 - 8.5$ (Magrini et al. 2007). Two papers published in the fall of 1981 presented the spectroscopic discovery of WR stars in these clusters: D'Odorico \\& Rosa (1981) derived a surprisingly large (50) number of WR stars in NGC 604, while Conti \\& Massey (1981) noticed that some WR stars in the four regions were excedingly luminous. Both studies however suffered from a lack of spatial resolution. More WR candidates were identified by interference filter imagery (taking advantage of their strong HeII emission) and spectroscopically confirmed by Massey \\& Conti (1983), Armandroff \\& Massey (1985, 1991) and Massey, Conti \\& Armandroff (1987). The most detailed catalogue of WR stars (with spectral classification) and WR candidates in M33 is presented in Massey \\& Johnson (1998). Drissen, Moffat \\& Shara (1990, 1993) identified more WR candidates based on high resolution CCD images with interference filters; but until now however, none of these were spectroscopically confirmed.  In the first paper (Abbott et al. 2004) of this series dedicated to WR stars in M33, we presented new spectra of one Of, 14 WN, one transition-type WN/WC and 26 WC stars in the field of M33. In this second paper, we present spectra of known WR stars and most of the WR candidates in the giant HII regions NGC 604, NGC 595 and NGC 592. ",
        "conclusions": "We have spectroscopically confirmed the presence of emission lines in a sample of WR candidates selected by interference filter imagery. Most of them are genuine WR stars, but we have also detected transition-type evolved O stars, with an He II equivalent width as low as 5 \\AA\\ . As demontrated here and in previous publications, this technique is very efficient, especially in crowded fields (see also Hadfield \\& Crowther 2007), with a very high success rate and relatively few false detections. The census of the WR population in NGC 588, NGC 592, NGC 595 and NGC 604 is now essentially complete: only three stars (NGC 604-WR2, NGC 604-WR4 and NGC 595-WR2), which are members of very dense and barely resolved groups, lack a clear identification although their WR nature is not in doubt. We have also obtained good quality spectrograms of previously known WR stars in these regions, allowing a better spectral type identification. The four giant HII region studied here harbour about 20\\% of the entire known Wolf-Rayet population of M33, with a WC/WN ratio significantly lower (0.25) than that of the field (0.4; Massey \\& Johnson 1998)."
    },
    "0806/0806.0650_arXiv.txt": {
        "abstract": "We construct a catalog of radio sources detected by GB6 (6\\,cm), FIRST and NVSS (20\\,cm), and WENSS (92\\,cm) radio surveys, and the SDSS optical survey.  The 2.7\\,million entries in the publicly-available master catalog are comprised of the closest three FIRST to NVSS matches (within 30\\arcsec) and vice-versa, and unmatched sources from each survey.  Entries are supplemented by data from the other radio and optical surveys, where available.  All objects with even a small probability of physical association are included, such that catalog users can easily implement their own selection criteria for data analysis.  We perform data analysis in the $\\sim3000\\deg^2$ region of sky where the surveys overlap, which contains 140,000 NVSS-FIRST sources, of which 64,000 are detected by WENSS and 12,000 by GB6.  About one third of each sample is detected by SDSS.  An automated classification method based on 20\\,cm fluxes defines three radio morphology classes: complex, resolved, and compact. Radio color-magnitude-morphology diagrams for these classes show structure suggestive of strong underlying physical correlations.  Complex and resolved sources tend to have a steep spectral slope ($\\alpha\\sim-0.8$) that is nearly constant from 6 to 92\\,cm, while the compact class (unresolved on $\\sim5\\arcsec$ scale by FIRST) contains a significant number of flat-spectrum ($\\alpha\\sim0$) sources.  In the optically-detected sample, quasars dominate the flat-spectrum compact sources while steep-spectrum and resolved objects contain substantial numbers of both quasars and galaxies.  Differential radio counts of quasars and galaxies are similar at bright flux levels ($>100$\\,mJy at 20\\,cm), while at fainter levels the quasar counts are significantly reduced below galaxy counts.  The optically-undetected sample is strongly biased toward steep-spectrum sources.  In samples of quasars and galaxies with SDSS spectra (2,885 and 1,288 respectively), we find that radio properties such as spectral slope, morphology, and radio loudness are correlated with optical color and luminosity. ",
        "introduction": "\\label{sec:introduction}  Quasars and powerful radio galaxies dominate the observed counts of continuum radio sources above milliJansky flux levels, and display spectacular morphological variety that is correlated with other properties such as spectral slope and luminosity.  The unification paradigm for radio galaxies and quasars \\citep{urry_review,jackson99} attempts to explain much of this rich variety of observational data as arising from essentially the same anisotropic processes which appear very different to us because of varying viewing angles to the radio jets.  This conjecture has fundamental implications for our understanding of quasars and galaxies, but for a strong test of the unification paradigm one essentially needs a large statistical sample with well-controlled selection criteria and robust estimates of the source morphology, as well as appropriate models to interpret the data.  Statistical studies of radio emission from extragalactic sources are entering a new era, resulting from the availability of large sky radio surveys that are sensitive to milliJansky flux levels \\citep[e.g.,][]{first,nvss,debreuck}.  The catalogs based on these surveys contain large numbers of sources, have high completeness and low contamination, and are available in digital form.  The wide wavelength region spanned by these surveys, from 6\\,cm for GB6 to 92\\,cm for WENSS, and detailed morphological information at 20\\,cm provided by FIRST and NVSS, allow significant quantitative and qualitative advances in studies of radio sources.  In addition, the optical catalog obtained by the Sloan Digital Sky Survey \\citep[SDSS;][]{york} can be used to separate quasars from galaxies, and the redshifts measured by SDSS allow a comprehensive study of the optical-radio correlation for quasars and galaxies.  FIRST and NVSS, conducted separately at the Very Large Array (VLA), were the first radio surveys with sufficiently high angular resolution to allow unambiguous matching with deep optical surveys, providing identifications for a large number of radio sources. The two surveys have the same radio frequency, but FIRST goes slightly deeper with higher resolution and smaller sky coverage (\\S\\ref{sec:surveys}). \\citet{machalski} and \\citet{sadler99,sadler02} measured the radio luminosity function of radio-loud active galactic nuclei (AGN) and star-forming galaxies by  cross-correlating the NVSS with spectroscopic galaxy surveys. \\citet{magliocchetti} matched FIRST to the 2 degree Field Galaxy Redshift Survey \\citep{2dfgrs}, using spectra to classify galaxies as ``classical'' radio galaxies, starburst and late-type galaxies, and Seyfert galaxies.  \\citet{i02} cross-correlated the FIRST survey with the SDSS photometric survey, resulting in a much larger sample of optical identifications (about one third of FIRST sources are matched to an SDSS source, and 0.16\\% of SDSS sources are matched to a FIRST source, with a matching radius of $2\\arcsec$), with galaxies outnumbering quasars 5:1, and a negligible fraction of radio stars in the sample.  \\citet{best} cross-correlated the SDSS spectroscopic galaxies sample with both FIRST and NVSS, and found that AGN dominate radio counts down to 5\\,mJy at 20\\,cm.  The more recent work of \\citet{mauch07} includes the largest radio-selected galaxy sample available from a single radio survey, combining the NVSS with the 6 degree Field Galaxy Survey \\citep{6dfgs}; they confirm that radio-loud AGN and star-forming galaxies have quite different distributions in the plane of radio power versus absolute K band (infrared) magnitude. The study of radio and optical properties of quasars was extended by \\citet{devriesW} who found that 10\\% of SDSS quasars have detectable radio cores at 1.4\\,GHz ($>0.75$\\,mJy), and 1.7\\% have double-lobed morphology, i.e., are associated with multiple FIRST components.  In this paper, we describe the construction of a unified catalog of radio sources detected in the 6\\,cm GB6, 20\\,cm FIRST and NVSS, and 92\\,cm WENSS radio surveys, and the SDSS optical survey (Data Release 6, hereafter DR6).  We began by merging the FIRST and NVSS surveys into a single catalog containing over 2\\,million sources detected by at least one survey; about 500,000 sources are detected by both FIRST and NVSS.  Where available, we supplement this survey with data obtained by the GB6 and WENSS surveys that enable the computation of radio spectral slopes; about 30,000 sources are detected by all four radio surveys.  The radio sources were also cross-correlated with the optical SDSS catalogs; nearly 92,000 FIRST sources have an SDSS counterpart within $2\\arcsec$.  The radio and optical surveys chosen for inclusion in the unified catalog primarily cover the northern celestial hemisphere.  We have opted to use these surveys in order to take advantage of the high astrometric accuracy of FIRST, which, designed to cover the same region of sky as the SDSS, is limited to the northern Galactic cap.  Several recent large-area radio surveys covering a wide range of frequencies are also available in the south.  These include the Sydney University Molonglo Sky Survey at 36\\,cm \\citep[SUMSS;][]{mauch03}, the Parkes-MIT-NRAO survey at 6\\,cm \\citep[PMN;][]{gregory}, and the Australia Telescope 20-GHz (1.5\\,cm) survey \\citep[AT20G;][]{massardi}.  The main advantages of the unified catalog presented in this paper are the multi-wavelength radio data (92\\,cm and 6\\,cm in addition to 20\\,cm), and the increased number of optical identifications from the SDSS DR6. By combining these five extensive surveys, we have assembled precise astrometric measurements, flux at multiple wavelengths, spectral index, morphological information (size, shape, and orientation of resolved objects), and optical identifications into a single comprehensive catalog of radio objects.  At optical wavelengths, the SDSS provides colors and classification (i.e., into quasars and galaxies) for the $\\approx30\\%$ of sources detected optically.  The unified catalog is primarily a resource of low-redshift ($z\\lesssim2$) quasars and radio galaxies with AGN, but may also prove to be an important source of rarer objects such as radio stars, high-redshift quasars, and high-redshift galaxies.  The area observed by all five surveys is nearly $3,000\\deg^2$.  The unified catalog provides comprehensive multi-wavelength observations at greater depth and for a larger number of sources than any previously available catalog of radio objects.  The limits in sky coverage of the catalog are defined by the FIRST and NVSS sky coverage (Fig.~\\ref{fig:position plots}).  The unified catalog thus includes sources in the Galactic plane which are imaged by the NVSS. Galactic sources must be studied with care: due to the nature of interferometry, radio surveys (depending on their angular resolution) generally do a poor job of imaging highly-extended sources such as Galactic HII regions and supernova remnants, whose angular sizes often reach several arcmin or more.  The analysis presented in this paper is limited to the sky covered by FIRST, which is greater than $30\\degr$ from the Galactic plane.  In this paper, we discuss scientific applications of the unified radio catalog and present a preliminary data analysis.  In a companion paper (A. Kimball et al. in preparation), we will expand upon and refine this analysis by comparing  the source distribution in radio morphology, radio color, and optical classification space to the models of \\citet{barai06,barai07}.  The remainder of the paper is laid out as follows. In \\S\\ref{sec:surveys}, we describe the surveys used to create the unified radio catalog.  In \\S\\ref{sec:matching} we discuss the creation of the catalog, including completeness and efficiency of the matching algorithms.  In \\S\\ref{sec:analysis} we present a preliminary analysis of the radio source distribution according to radio morphology, radio color, optical identification, and, for sources with spectra, redshift and luminosity.  We also discuss some catalog applications.  We summarize our results, and discuss the suitability of the catalog for comparison with radio evolution models, in~\\S\\ref{sec:discussion}. ",
        "conclusions": "The information provided by this catalog will enable many diverse studies.  In this section, we briefly remark on applications toward the investigation of the radio quasar/galaxy unification theory, the search for radio stars, and the selection of possible high-redshift galaxies.  We then describe the usefulness of the catalog for studies of the evolution of the radio universe by comparison with models, which we discuss in detail in a companion paper (A. Kimball et al. in preparation).  \\subsection{Unification paradigm for radio-loud active galactic nuclei}  The orientation-based unification scenario \\citep{urry_review} was originally motivated by the similarity of radio emission from sources that appear very different optically, i.e. quasars and galaxies.  The theory assumes that sources whose emission is dominated by a radio core or lobes are members of the same ``parent'' population, but differ in appearance because their highly anisotropic emission is viewed from different angles.  This anisotropy is predominantly due to relativistic motion of the plasma in the inner jets and Doppler boosting (``beaming'') when the angle between the line of sight and the plasma velocity vector is small \\citep{krolik}.  Hence, the same object could appear as a core source when viewed at small angles (with the ``boosted'' core outshining the lobes), and as an extended double-lobe or a core-lobe source otherwise.  Understanding the radiation anisotropies in AGNs is required to unify the different types; that is, to identify each single, underlying AGN type that gives rise to different observed classes through different orientations \\citep{antonucci,urry_review}.  The unified catalog allows users to classify many quasars and galaxies according to their radio morphology (e.g., core sources, lobe sources) and spectral behavior.  Statistical studies of such objects will lend evidence for or against the unification paradigm by the investigation of number statistics, the size versus orientation distribution, and environment.  The unification scenario has successfully explained many observed properties of bright radio sources in previous research \\citep{barthel,urry,padovani,lister,urry_review}, but recent studies suggest there may be some intrinsic differences between radio quasars and galaxies \\citep{willott02}.  \\subsection{A method for selecting high-redshift ($z>1$) galaxy candidates} \\label{subsec:highz gals}  Galaxies at high redshift are important for studies of large scale structure and galaxy evolution, and in fact the higher the redshift of a galaxy the more useful it is for these studies.  Candidate high-redshift galaxies can be selected from the unified catalog based on their radio morphology, radio spectral slope, and lack of an optical counterpart.  The lack of optical counterpart requirement is intended to select sources so distant that their observed optical emission lies blueward of the rest-frame $4000$\\AA ~break \\citep[e.g., ][]{madau}.  As the optical sources used herein were selected in the $r$ and $i$ bands at 6165\\AA ~and 7481\\AA ~respectively, this requirement will tend to select galaxies at redshifts of $z\\sim1$ and higher.  A steep radio spectrum has often been used as a criterion to find high-redshift sources, owing to the fact that higher redshift objects are observed at higher rest-frame frequencies, and because radio galaxy spectra tend to flatten below $\\sim300$\\,MHz in their rest frame \\citep[e.g.,][and references therein]{cruz}. Additionally, as discussed in \\S\\ref{subsec:optical ids I}, in the compact morphology subclass a steep radio spectrum is a likely indicator of an optically-resolved source.  To find a sample of candidate high-redshift galaxies, we select sources which are unresolved by FIRST, undetected by the SDSS, and have steep radio spectra.  Specifically, we begin with the sample of objects identified by NVSS, FIRST, and WENSS which have compact radio morphology.  We require $\\alpha^{92}_{20}<-0.5$ and no SDSS match within $3\\arcsec$.  Note that heavily dust-obscured galaxies are also likely to be found in such a sample. In the catalog overlap region, the above criteria select 9,953 sources. Visual examination of SDSS images of all the candidates revealed 334 suspect sources, i.e. radio positions coinciding with a diffraction spike, scattered light, a nearby star or galaxy, etc.  Removing the suspicious sources results in a sample of 9,619 candidate $z>1$ galaxies.  It will be interesting to compare this sample with deeper optical surveys and observations at other wavelengths, as the data become available.  This sample is available for download on the catalog website (sample~K).  \\subsection{The search for radio stars} \\label{sec:radiostars}  With the great increase in the sensitivity of radio surveys in the last several decades, along with the more accurate source positions afforded by radio interferometry, comes the ability to search for fainter and rarer radio sources. A small fraction of stars have significant non-thermal radio emission, such as dMe flare stars \\citep{white} and cataclysmic variables \\citep{chanmugam,mason}, but most stars have only weak thermal radio emission. Because quasars and stars lie in different locations in SDSS color-color diagrams \\citep{i02,richards01}, it may be possible to select a clean sample of radio stars based on their photometric colors.  To investigate this possibility, we select a sample of radio star candidates from the unified radio catalog.  We begin by selecting point sources from the catalog overlap region which coincide with a FIRST radio detection.  We use a conservative matching distance of $1\\arcsec$ to limit the contamination due to random matches.  To ensure a sample for which the majority have spectra, we require $i<19$.\\footnote{The SDSS quasar target selection algorithm flags all $i<19.1$ sources within $2\\arcsec$ of a FIRST detection \\citep{dr5quasars}.}  We eliminate the region of color-color space where quasars are commonly found: $u-g<0.8$ and $-0.2<g-r<0.6$.  The number of candidates remaining after each step of the selection algorithm is shown in Table~\\ref{table:radiostars}.  The resulting sample of 532 candidate radio stars contains 406 objects with SDSS spectra.\\footnote{The remaining 126 stellar candidates do not have spectra for one or more of the following reasons: they are not in the region of sky included in the SDSS DR6 spectroscopic sample, they are too close on the sky to another spectral target, they failed the quasar target selection algorithm because of the bright limit or bad photometry, or they were not targeted because the photometry in a previous SDSS data release led to an extended rather than point source optical morphology identification.} From the SDSS spectral classifications,\\footnote{Spectroscopic identifications determined using the SDSS ``SpecBS'' Princeton reductions.  For details, see http://spectro.princeton.edu/.} we determine that 78 are stars, 313 are quasars, and 15 are galaxies.  The results of applying the selection to the SDSS DR5 quasar catalog \\citep{dr5quasars} are also shown for comparison.  The completeness of the quasar catalog is $\\sim95\\%$ \\citep{richards02,vandenberk}.  Seven of the known quasars were identified as galaxies by the SDSS spectroscopic pipeline, and two as stars.  The latter are actually the super-position of a star and a quasar: the spectra are primarily stellar, but also contain at least one broad emission line.  Note that 33 of the candidate radio sources classified as ``quasar'' are absent from the quasar catalog.  These are sources that were spectrocopically-observed for the sixth SDSS data release, but are not in DR5.  The sky density of unresolved $i<19$ SDSS sources is $\\sim1000\\deg^{-2}$, while the density of quasars is $\\sim10\\deg^{-2}$.  However, out of nearly 6 million point sources with $i<19$, we found only 78 confirmed radio stars, implying that the likelihood of a star being a radio source is approximately $10^{-5}$, while the likelihood for a quasar is about $8\\times10^{-2}$.  Consequently, despite using a conservative color cut to eliminate probable quasars, the sample is dominated by radio quasars rather than radio stars.  Using this color-color selection is therefore not an efficient way to find radio stars.  One must have spectral classification in order to select a clean radio star sample; photometric observations alone are not sufficient.  \\subsection{Radio galaxy evolution models}  In the companion paper, we attempt to understand and quantitatively explain the observed source counts for the nine subclasses defined by radio morphology and optical identification (Fig.~\\ref{fig:radio colorcolor optical}), with the aid of models.  \\citet{barai06,barai07} investigated a large number of input parameter variations for several radio galaxy evolution models and compared the resulting distributions of radio power, spectral index, linear size, and redshift to the low-frequency ($\\sim$2\\,m) Cambridge catalogs \\citep{laing,eales,mcgilchrist}.  They found that none of the models was able to provide good simultaneous fits to all observables for all three catalogs, but noted that the observational datasets were too small to adequately constrain the model parameters.  The catalog presented in this paper represents a significant increase in the size of radio samples for which luminosity, size, spectral slope, and redshift can be measured.  We will compare the observed distributions from the unified catalog to distributions of sources from mock catalogs (obtained by applying observational selection criteria to the computer-generated radio skies). The distributions will help to constrain the increasingly sophisticated models of radio galaxy evolution.  The models will thus help explain the physics and evolution leading to the distributions seen in the unified catalog.  \\subsubsection{Defining a long-wavelength flux-limited sample} \\label{subsubsec:wenss sample}  The radio galaxy evolution models currently perform best at long wavelengths, where emission is dominated by the radio lobes.  The most robust approach for comparing models and observations is therefore to define a flux-limited sample using long-wavelength measurements, such as the 92\\,cm flux from the WENSS survey.  Fortunately, because most radio sources are brighter at longer radio wavelengths, a source observed in WENSS (faint limit of 18\\,mJy) is almost guaranteed to be detected at 20\\,cm (faint limits of 1\\,mJy for FIRST and 2.5\\,mJy for NVSS).  However, the 92\\,cm flux limit must be chosen carefully in order to avoid biasing against sources with radio spectra so steep that they are too faint at 20\\,cm to be detected by NVSS.  The color-magnitude diagrams and spectral slope histograms presented in the data analysis of this paper indicate that very few radio sources have a spectral slope index less than -1.5 between 20 and 92\\,cm.  To ensure an NVSS detection for all sources with $\\alpha^{92}_{20}<-1.5$ in a flux-limited sample requires $t_\\mathrm{WENSS}\\leq12.9$ (see Eq.~\\ref{eq:alpha9220}).  We positionally matched all $t_\\mathrm{WENSS}<12.9$ sources in the WENSS survey to NVSS and FIRST, and found that only 1\\% did not have a match in NVSS.  We interpret this value as an upper limit on the fraction of sources with $\\alpha^{92}_{20}<-1.5$.  Ninety-five per cent of the WENSS sources have both a FIRST detection and an NVSS detection within $30\\arcsec$.  The sample of NVSS-FIRST-WENSS matches with $t_\\mathrm{WENSS}<12.9$ in the unified catalog contains 47,567 sources.  Over 12,000 of those have optical counterparts, from which a cosmological distance can be determined via photometric redshift techniques.  Photometric redshifts from the SDSS are accurate to $\\sim0.06$ for luminous red galaxies with redshifts $z<0.7$ \\citep{nikhil}.  A further 3,000 of the radio-optical sample additionally have SDSS spectroscopic redshifts.  This dataset therefore represents a major size increase compared to the sample of 327 sources used by Barai \\& Wiita."
    },
    "0806/0806.2149_arXiv.txt": {
        "abstract": "The cold dark matter scenario predicts that a large number of dark subhalos should be located within the halo of each Milky-way sized galaxy. One tell-tale signature of such dark subhalos could be additional milliarcsecond-scale image splitting of quasars previously known to be multiply-imaged on arcsecond scales. Here, we estimate the image separations for the subhalo density profiles favoured by recent N-body simulations, and compare these to the angular resolution of both existing and upcoming observational facilities. We find, that the image separations produced are very sensitive to the exact subhalo density profile assumed, but in all cases considerably smaller than previous estimates based on the premise that subhalos can be approximated by singular isothermal spheres. Only the most optimistic subhalo models produce image separations that would be detectable with current technology, and many models produce image separations that will remain unresolved with all telescopes expected to become available in the foreseeable future. Detections of dark subhalos through image-splitting effects will therefore be far more challenging than currently believed, albeit not necessarily impossible. ",
        "introduction": "While the cold dark matter (CDM) scenario has been very successful in explaining the formation of large-scale structures in the Universe \\citep[see e.g.][for a review]{Primack}, its predictions on the scales of individual galaxies have not yet been confirmed in any convincing way. One particularly interesting feature of the CDM model is the high level of halo substructure generated. According to current simulations, each galaxy-mass CDM halo should contain large numbers of subhalos \\citep[typically accounting for $\\approx 5$--10\\% of its total mass; e.g.][]{Nurmi et al.,Diemand et al. a} in the dwarf-galaxy mass range. However, these subhalos do not appear to correspond to luminous structures, as the dark halo of the Milky Way would then contain a factor of 10--100 more satellite galaxies than observed, provided that each subhalo corresponds to a luminous dwarf galaxy \\citep[e.g.][]{Klypin et al.}. A similar lack of dwarf galaxies compared to the number of dark halos predicted is also evident on scales of galaxy groups \\citep{Tully et al.}. One way out of this problem is to assume that most of these low-mass halos correspond to so-called dark galaxies \\citep[][]{Verde et al.}, i.e. objects of dwarf-galaxy mass which either do not contain baryons or in which the baryons have not formed many stars. While a number of very faint satellite galaxies have recently been detected \\citep[e.g.][]{Zucker et al.,Simon et al.}, it is still far from clear that these exist in sufficent numbers to account for the subhalos predicted by CDM \\citep[][]{Simon & Geha}.  Gravitational lensing may in principle offer a route to detecting completely dark galaxies, and the magnification associated with dark halo substructures has long been suspected to be the cause of the flux ratio anomalies observed in multiply-imaged quasars \\citep[e.g.][]{Mao & Schneider,Kochanek & Dalal}. However, the subhalos predicted by current CDM simulations appear too few to explain the flux ratios of many of these systems \\citep[e.g.][]{Mao et al.,Metcalf a,Maccio & Miranda}. This indicates that other mechanisms, like lensing by low-mass field halos outside the main lens \\citep{Metcalf b,Miranda & Maccio}, stellar microlensing in the lens galaxy \\citep{Schechter & Wambsganss} and absorption or scintillation in the interstellar medium \\citep{Mittal et al.} may also be at work.  A similar but more direct way of putting the CDM subhalo predictions to the test could be to search for small-scale {\\it image separations} produced by strong lensing from subhalos. \\citet{Yonehara et al.} argue that by targeting quasars that are already known to be multiply-imaged on arcsecond scales, there should be a significant probability of detecting image splitting by subhalos at scales of milliarcseconds (so-called millilensing or sometimes mesolensing). If the sources are extended and can be resolved at scales smaller than the Einstein radii of the subhalos, then not only the number densities of the subhalos but also their internal density profiles can be constrained \\citep{Inoue & Chiba b}. It has been argued that this technique may become observationally feasible with ALMA \\citep{Inoue & Chiba a} or future space-VLBI missions like VSOP-2 \\citep{Inoue & Chiba c}. If correct, this would constitute a major step forward in the study of dark halo substructures.  These proposals do however rely on the assumption that the image separations (or Einstein radii) of subhalo lenses are similar to those produced by singular isothermal spheres. This is an assumption that is difficult to justify, since neither simulations, theoretical arguments nor observations favour a density profile of this form for dark matter halos in the relevant mass range.  N-body simulations based on the CDM paradigm typically predict dark matter halos to have inner density profiles of the form $\\rho(r)\\propto r^{-\\alpha}$ with central density slopes $\\alpha\\approx 1$ \\citep[e.g.][hereafter NFW]{NFW}, whereas the singular isothermal sphere (SIS) assumes $\\alpha=2$. While the SIS has proved to be a successful model for observed galaxy-mass lenses \\citep[e.g.][]{Rusin et al.,Koopmans et al.}, this is believed to be due to the luminous baryons residing in the inner regions of these objects. This component contributes substantially to the overall mass density in the centre, and its formation over cosmological time scales may also have caused the CDM halo itself to contract, thereby steepening the inner slope of its density profile \\citep[e.g.][]{Gnedin et al.,Maccio et al.,Gustafsson et al.,Kampakoglou}.  Dwarf galaxies are well-known to be more dark matter dominated than their more massive counterparts, and some of the mechanisms proposed for keeping the missing satellites of the Local Group dark assume that most of their baryons have been lost early in the history of the Universe, or are prevented from collapsing into their central regions \\citep[e.g.][]{Barkana & Loeb,Read et al.}. Hence, it seems reasonable that the inner slope $\\alpha$ of their density profiles should be shallower than that of an SIS. Indeed, \\citet{Ma} finds that to explain the distribution of angular separations among strong lenses, the lens population must shift from an SIS-type profile ($\\alpha = 2$) to something more resembling the NFW profile ($\\alpha \\approx 1$) at halo masses of $\\lesssim 10^{12}\\ M_\\odot$. Observational constraints on the slope of the innermost density profile of both field and satellite galaxies in the dwarf galaxy mass-range typically favour a slope even shallower than this \\citep[$\\alpha \\approx 0$--1; e.g.][]{de Blok & Bosma,Spekkens et al.,Zackrisson et al.,Gilmore et al.}. This is potentially bad news for the detectability of low-mass subhalos through millilensing effects, since the image separation expected from strong lensing by a halo of fixed virial mass is extremely sensitive to the slope of its inner density profile. In fact, as pointed out by \\citet{Trentham et al.}, the image separation expected from strong lensing by dark galaxies would be unobservably small if their density profiles are of NFW type. An inner slope of the density profile even shallower than this would just make the problem worse.  As low-mass halos in the field fall into the potential wells of larger halos, tidal interactions with their host halos and with other subhalos will cause these objects to lose mass. Since this mass loss preferentially takes place in the outer regions of the subhalos, their density profiles change over time \\citep[e.g.][]{Hayashi et al., Kazantzidis et al.}, thereby altering their lensing properties compared to more isolated field halos of the same mass. Here, we take a critical look at the prospects for strong-lensing detections of dark subhalos in the dwarf-galaxy mass range, by deriving the image separations expected for subhalo density profiles favoured by current simulations. ",
        "conclusions": "Our results indicate that detection of subhalos through gravitational image-splitting is likely to be considerably more challenging than suggested in previous studies, due to the smaller image separations predicted for subhalo density profiles more realistic than the SIS models often adopted. In fact, no currently planned telescope will be able to resolve the image separations produced by subhalos with density profiles of the type suggested by the most realistic simulations currently available (H03 \\& K04). We stress, however, that these simulations do not necessarily represent the final word on this issue. NFW density profiles were adopted for the subhalo progenitors in both the H03 and K04 simulations, which means that the slope of the central density profile prior to stripping was {\\it assumed} rather than derived from the simulations themselves. If the subhalos would have steeper central density slopes (e.g. of M99 type), these would give rise to image separations that could be resolved even with existing telescopes.  Despite the somewhat bleak detection prospects presented here, there are at least two effects that can potentially improve the detectability of image splitting by subhalos: baryon cooling and the presence of intermediate mass black holes.  Baryon cooling would cause the subhalo (or its progenitor halo) to contract \\citep[e.g.][]{Gnedin et al.,Maccio et al.,Gustafsson et al.,Kampakoglou}, thereby increasing the central density and boosting the image separation. It is, however, not clear how strong this effect is likely to be, since this depends on the details of the mechanism that prevents the dark subhalo from forming stars. Baryon cooling is usually assumed to be associated with star formation, and most attempts to explain the lack of bright subhalos therefore propose that the baryons have either been lost early in the history of the Universe, or are prevented from cooling by the ultraviolet background provided by reionization \\citep[e.g.][]{Barkana & Loeb,Read et al.}. Neither mechanism is likely to result in any significant halo contraction. Nonetheless, claims of dark baryons in the form of cold gas in galactic disks have been made \\citep{Bournaud et al.}, implying that there may be some route for gas to cool without forming stars or being detected by the usual H$_2$ tracers \\citep[e.g.][]{Pfenniger et al.,Pfenniger & Combes}. One may therefore speculate that there could be alternative solutions to the missing-satellite problem, in which subhalos are kept dark even though baryon cooling has taken place.  Intermediate mass black holes \\citep[IMBHs, see e.g.][]{van der Marel,Zhao & Silk,Noyola et al.} with masses of $\\sim 10^2$--$10^4\\ M_\\odot$ would also boost the image separations, if present in the centres of dark subhalos. This could for instance be the case if the empirical relations between the mass of a supermassive black hole and its dark matter halo \\citep{Ferrarese} would extend into the dwarf-galaxy mass range. Even in the case of $\\kappa=\\gamma=0$, an IMBH of mass $10^4 \\ M_\\odot$, would give an image separation of $\\approx 4\\times 10^{-4}$ arcsec for $z_\\mathrm{l}=0.5$ and $z_\\mathrm{s}=2.0$. This is sufficient to allow VSOP-2 to resolve the image splitting, regardless of the density profile of the subhalo hosting the IMBH.  Of course, even if the Ferrarese relations would hold for luminous dwarf galaxies, the do not necessarily do so for dark ones, since this depends on the formation details of IMBHs. Moreover, a population of halo IMBHs {\\it not} associated with subhalos could form an undesired background of millilensing events that would obfuscate attempts to study subhalos through image-splitting effects.  Finally, there is an important issue related to the limitation of current CDM halo simulations. Current N-body simulations only resolve scales down to $R\\sim 0.001 r_\\mathrm{vir}$, whereas for the steeper density profiles -- i.e. those producing detectable image separations for dwarf-galaxy masses -- a non-negligible fraction of $\\bar{\\Sigma}(<r)$ at the  Einstein radius comes from radii smaller than this. This makes the exercise of deriving dark halo lensing properties from fitting functions based on current N-body data somewhat hazardous. Unless theoretical arguments \\citep[e.g.][]{Hansen & Stadel} can be used to determine the slope of the density profile at even smaller radii, robust predictions for the image separations produced by low-mass halos may be very difficult to derive. Observational limits on the lensing optical depth as a function of $\\Delta\\theta$, based on searches for subhalo image splitting, may eventually turn out to be a useful way of setting constraints on the internal structure of dark matter halos at scales below the resolution of numerical simulations. Of course, the optical depth does not depend only on the density profile of the subhalos, but also on their mass function and the spatial distribution within the macrolens. Detailed estimates of the optical depth for strong lensing by subhalos in the dwarf-galaxy mass range is outside the scope of this study, but will be covered in paper II of this series (Riehm et al. 2008, in preparation)."
    },
    "0806/0806.0466_arXiv.txt": {
        "abstract": "{The solar mid-infrared metallic emission lines have already been observed and analyzed well, and the formation scenario of the Mg\\,I 12 $\\rm\\mu$m lines has been known for more than a decade. Detections of \\textit{stellar} emission at 12 $\\rm\\mu$m have, however, been limited to Mg\\,I in very few objects. Previous modeling attempts have been made only for Procyon and two cool evolved stars, with unsatisfactory results for the latter. This prevents the lines' long predicted usage as probes of stellar magnetic fields.} {We want to explain our observed Mg\\,I emission lines at 12 $\\rm\\mu$m in the K giants Pollux, Arcturus, and Aldebaran and at 18 $\\rm\\mu$m in Pollux and Arcturus. We discuss our modeling of these lines and particularly how various aspects of the model atom affect the emergent line profiles.} {High-resolution observational spectra were obtained using TEXES at Gemini North and the IRTF. To produce synthetic line spectra, we employed standard one-dimensional, plane-parallel, non-LTE modeling for trace elements in cool stellar atmospheres. We computed model atmospheres with the MARCS code, applied a comprehensive magnesium model atom, and used the radiative transfer code MULTI to solve for the magnesium occupation numbers in statistical equilibrium.} {The Mg\\,I emission lines at 12 $\\rm \\mu$m in the K giants are stronger than in the dwarfs observed so far. We present the first observed stellar emission lines from Mg\\,I at 18 $\\rm \\mu$m and from Al\\,I, Si\\,I, and presumably Ca\\,I at 12 $\\rm \\mu$m. We successfully reproduce the observed Mg\\,I emission lines simultaneously in the giants and in the Sun, but show how the computed line profiles depend critically on atomic data input and how the inclusion of energy levels with $n \\ge 10$ and collisions with neutral hydrogen are necessary to obtain reasonable fits.} {} ",
        "introduction": "\\label{Introduction} Metallic solar emission lines around 12\\,$\\rm\\mu$m were first identified by \\citet{Chang83}, with the most prominent lines originating from transitions\\footnote{Quantum state $nl$, where $n$ denotes the principal quantum number and $l$ the orbital} $7i~\\rightarrow~6h$ (12.32\\,$\\rm\\mu$m) and $7h~\\rightarrow~6g$ (12.22\\,$\\rm\\mu$m) between Rydberg states of neutral magnesium. Additional Rydberg emission lines from Al\\,I, Si\\,I, and tentatively Ca\\,I were identified as well \\citep{Chang83,Chang84}. The Mg\\,I line formation scenario remained unclear until \\citet{Chang91} and \\citet{Carlsson92}, hereafter C92, in two independent studies reproduced the emission features by employing standard plane-parallel numerical radiative transfer with a detailed atomic model and a reliable solar atmosphere. They confirmed an origin below the atmospheric temperature minimum, refuted a chromospheric line contribution, and established a non-LTE (LTE: Local Thermodynamical Equilibrium) formation scenario. The solar lines have subsequently been used in, e.g., Mg\\,I statistical equilibrium analysis by Zhao et al.\\ (1998). C92 proposed a general Rydberg line formation mechanism for the highly excited metal lines, which implied that all visible metal emission lines in the solar spectrum around 12\\,$\\rm\\mu$m originated in the photosphere. A detailed non-LTE modeling of the Al\\,I emission has been carried out by \\citet{Baumueller96}, where they confirmed this mechanism.\\\\ The lack of suitable spectrometers and the low stellar flux in the mid-infrared have in the past made high-resolution spectroscopy in this wavelength region possible only for the Sun and a few luminous nearby stars. \\citet{Ryde04} observed the 12\\,$\\rm\\mu$m Mg\\,I emission features in Procyon, and successfully reproduced the line profiles by employing the same modeling technique as C92. \\citet{Uitenbroek96} observed and modeled the evolved stars Arcturus ($\\alpha$ Boo) and Betelguese ($\\alpha$ Ori). Using the same model atom as C92, they were unable to fit the line profiles of the $7i \\rightarrow 6h$ Mg\\,I transition, which appeared both in emission (Arcturus) and absorption (Betelguese). Their observational sample also included five M giants and supergiants, in which the line appeared in absorption. However, \\citet{Ryde06} investigated water vapor lines for Betelguese in the same spectral region, and found a water line that coincided with the wavelength of the Mg\\,I 12.32\\,$\\rm\\mu$m line. The group successfully modeled the water line, without considering the Mg\\,I blend (which we predict to be very weak, see Sect.\\ \\ref{Observations_disc}). This may explain the sample of observed M star absorption at 12.32 $\\rm \\mu$m, since water vapor is expected in these stars, whereas the Mg\\,I emission line contribution should be minor.\\\\ A well known potential use for the Mg\\,I lines is as probes of magnetic fields, which play a fundamental role in the underlying physics of a cool stellar atmosphere. Zeeman line-splitting from an external magnetic field increases quadratically with wavelength, while the Doppler broadening only has a linear dependence. Thus a line's sensitivity to magnetic fields becomes higher at longer wavelengths. The splitting of the solar emission lines was pointed out early and has been extensively analyzed. We have performed observations of the magnetically active dwarf $\\epsilon$ Eridani. These will be reported on in a forthcoming paper (Richter et al., in preparation), hence we defer further discussions about stellar disk-averaged magnetic fields until then.\\\\ Prior to (stellar) diagnostic applications, however, we should make sure that we are able to model and understand these lines in a range of stars. So far, as mentioned above, modeling attempts for evolved stars have been unsuccessful. We address this issue here by analyzing high-resolution observational spectra, which show strong Mg\\,I emission lines in the three giants Pollux (K0\\,III), Arcturus (K1.5\\,III), and Aldebaran (K5\\,III). We model and analyze simultaneously the three K giants and the Sun, with particular emphasis on influences from atomic data, and discuss why previous modeling attempts have not succeeded. The organization of the paper is as follows; in Sect.\\ \\ref{Observations}, we describe the observations. In Sect.\\ \\ref{Departure} we review some concepts about the formation of the infrared Mg\\,I emission lines and in Sect.\\ \\ref{Modeling} we describe our modeling procedure. Results are presented in Sect.\\ \\ref{Results} and we discuss them and give our conclusions in Sect.\\ \\ref{Discussion}. ",
        "conclusions": "\\label{Discussion}  The model atom extension and the introduction of collisions with neutral hydrogen remove the discrepancy between observed and modeled emission in our study. \\textit{Before} the model atom was extended, we undertook a number of tests to investigate the large discrepancies between observations and models. We therefore start our discussion by a short summary of these, with the purpose to simplify future work.\\\\ The lines are sensitive to the photoionization rates as these affect the recombination to the Rydberg levels. We thus made an \\textit{ad-hoc} increase in the photospheric MARCS radiation field until the modeled Mg\\,I emission lines matched the observed, but found that the required additional mean intensities gave rise to surface fluxes that by far exceeded observed ones. We included a chromospheric temperature-rise in Arcturus and found it to have a negligible impact, with the lines forming in atmospheric layers below the temperature minimum. We computed MARCS models in spherical geometry and employed the spherical version of MULTI, S-MULTI \\citep{Harper94}, but differences from plane-parallel models were small. No attempt to analyze influences from atmospheric inhomogeneities has been made in this work. A discussion about how granulation affects the solar lines can be found in \\citet{Rutten94}.  \\subsection{The model atom extension} \\label{Extension}  All Mg\\,I departure coefficients are shifted upwards (increased) by the extension of the model atom, but the upward shift is more pronounced in the higher energy levels. It is this change in the departure coefficient \\textit{ratio} in the line-forming regions that is sufficient to cause a significant change in results, i.e., higher emission peaks for the larger atom. The enhanced collisional coupling in the uppermost Mg\\,I levels and to the Mg\\,II ground state strengthens the cascading process in rather the same manner as the model atom with top-levels $n=9$ did when compared to one reaching only $n=7$ (see Sect. \\ref{Enlargement}). Qualitatively, more recombinations enter at the top-levels, channel down through transitions that take part of the Rydberg ladder (see Fig.\\ \\ref{Fig:terms}) and set up larger departure coefficient differences between adjacent levels. The extension thus has a significant impact on the mid-infrared emission lines, whereas the overall character of the Mg\\,I statistical equilibrium remains.\\\\ To include the energy levels with $n \\ge 10$ seems especially pertinent when applying the model atom on diluted stellar atmospheres with low surface gravities (as shown by the large differences in the modeled line profiles of the K giants). These are more influenced by radiative transitions, and thus the extension ensures that the Rydberg level replenishment from the ion state is properly accounted for by including top-levels that are fully dominated by collisions.  \\subsection{Effects from extra collisions} \\label{Hcoll_disc}  Higher rates of collisional excitation and ionization affect the 12 $\\rm\\mu$m lines in the implemented stellar model atmospheres differently. Figs.\\ \\ref{Fig:12.22um_ArAl} and \\ref{Fig:12.22um_PoSu} show how (by the introduction of collisions with neutral hydrogen) the Mg\\,I 12 $\\rm\\mu$m emission is reduced in the Sun, increased in Arcturus and Aldebaran and almost unchanged in Pollux. Apparently, a homogeneous increase of collisional rates actually results in stronger emission in the low surface gravity atmospheres of Arcturus and Aldebaran (which one would perhaps not expect since, generally, collisions act to thermalize lines toward LTE).\\\\ We have analyzed this result by computing additional models for the Sun and Arcturus, where we included radiatively allowed bound-bound collisions with neutral hydrogen for 1) only the three transitions $7i \\rightarrow 6h$, $7h \\rightarrow 6g$, and $7g \\rightarrow 6f$ and 2) only transitions with $\\Delta n,\\Delta l = -1$ and $n \\geq 8$ and $l \\geq 4$. The first test imposes stronger collisional coupling only in the $n = 7 \\rightarrow 6$ transitions themselves, whereas the second serves to strengthen the collisional bound-bound coupling in transitions in the uppermost levels that take part of the Rydberg ladder (see Fig.\\ \\ref{Fig:terms}), while maintaining all other rates.\\\\ We quantify the results of this experiment by the normalized flux maximum in the 12.22 $\\rm\\mu$m line. In the first case, both models give lower emission, as expected from the direct thermalizing effect on the $n = 7 \\rightarrow 6$ transitions. The modeled flux maximum above the continuum decreased by $\\sim$\\,30\\,\\% and $\\sim$\\,25\\,\\% for the Sun and Arcturus respectively, as compared with the model excluding collisions with neutral hydrogen. For the second case, however, the emission increased with a similar percentage in the solar model, whereas in Arcturus, the modeled flux maximum doubled its value. Thus, in a star like Arcturus the enhancement effect from highly excited lines (see Sect.\\ \\ref{Departure}) dominates the reduction effect from the higher collisional rates in the line transitions themselves, so that when introducing collisions with neutral hydrogen homogeneously throughout the model atom, the outcome is an emission increase (as seen in Fig.\\ \\ref{Fig:12.22um_ArAl}). We can understand this by noting that, e.g., the ratio between collisional and radiative deexcitation rates for the $8i \\rightarrow 7h$ transition (supplying the 12.22 $\\rm\\mu$m line) is $C_{ji}/R_{ji}$\\,$\\sim$\\,15.0 (Sun) and $\\sim$\\,0.2 (Arcturus) in typical line-forming layers when collisions with neutral hydrogen are excluded. In principle, this means that the contribution from hydrogen is needed in the giants to ensure an efficient Rydberg ladder.  \\subsection{Observations of Rydberg emission lines around 12 $\\rm\\mu$m} \\label{Observations_disc}  The observed emission-line flux spectra for Pollux and Arcturus in the 12 $\\rm\\mu$m region closely resembles the solar limb intensity spectrum, whereas the solar disk-center spectrum lacks emission features from other elements than magnesium (see Fig.\\ \\ref{Fig:12um_obs}). It is evident that strong emission features from the K giants, as compared with solar-type dwarfs, appear for more metallic Rydberg lines than magnesium, and the observability of different elements so far follows the same pattern as in the Sun. Future observations will tell if this observed trend remains for a larger sample. Model tests with K dwarfs indeed predict lower Mg\\,I emission for dwarfs than for giants also within the same spectral class (supported as well by observations of the magnetically active K dwarf $\\epsilon$ Eridani, Richter et al., in preparation).\\\\ For cooler K giants, absorption from water vapor starts to influence the 12 $\\rm\\mu$m spectrum and we have detected a blend at 12.32 $\\rm\\mu$m in Arcturus and Aldebaran. Observations from this spectral region in the yet cooler M supergiant Betelguese \\citep{Ryde06} reveal no emission lines above the noise level, which is consistent with our modeling of the Mg\\,I lines in this star (using the same stellar parameters as in \\citeauthor{Ryde06}).  \\subsection{Comparison with other studies}  \\citet{Uitenbroek96} modeled the Mg\\,I 12.32 $\\rm\\mu$m line in Arcturus and concluded that the computed line was, roughly, half as strong as the observed. When using the 'C92' model, we find an even larger discrepancy for the 12.22 $\\rm\\mu$m line (see Fig.\\ \\ref{Fig:12.22um_ArAl}). As already discussed, these longstanding discrepancies between observations and models for K giants are removed when using our new model atom. We note also that \\citeauthor{Uitenbroek96} did not detect the water vapor absorption line, which is blended with the Mg\\,I 12.32 $\\rm\\mu$m line in our Fig. \\ref{Fig:12um_obs}. The extension of the model atom changes the results in the solar case as well (as compared with C92 and \\citeauthor{Ryde04} 2004), placing the Mg\\,I lines in higher emission, but the former results are recovered by the introduction of collisions with neutral hydrogen.\\\\ We have settled here with the rather questioned, albeit standard, Drawin recipe for collisions with neutral hydrogen while we await results from more rigorous quantum mechanical calculations. We have shown the influence from these collisions on the formation of the mid-infrared Mg\\,I emission lines, however we stress that it has not been an aim of this investigation to put detailed empirical constraints on their efficiency. Such a task would require a larger set of lines, including also other wavelength regions. Nevertheless, we may still compare our adopted scaling factor to the Drawin formula for Mg\\,I collisions with neutral hydrogen, $S_{\\rm H}=10^{-3}$, with other values from the literature. Mentioned in Sect.\\ \\ref{Atom} was the exponential decrease resulting in $S_{\\rm H}=3 \\times 10^{-10}$ for the 12 $\\rm\\mu$m transitions \\citep{Zhao98}, a model which was later abandoned by the same group in favor of a constant $S_{\\rm H}=0.05$ \\citep{Gehren04}, inferred only from optical lines. In a recent non-LTE abundance study of magnesium in metal-poor stars \\citep{Mashonkina08}, $S_{\\rm H}=0.1$ is used. The value we find based on the mid-infrared lines is more than one order of magnitude lower than the values obtained from these two optical studies. Our model atom has not been applied to optical lines, however such a combined study should be given high priority in future work. We thus conclude that the mid-infrared emission lines from near-by giant stars may be suitable diagnostics for testing atomic input data in future non-LTE analyses.\\\\ Finally, as we are now able to model and explain the observed emission lines for both dwarfs and giants, diagnostic applications regarding stellar disk-averaged magnetic fields are possible."
    },
    "0806/0806.4245_arXiv.txt": {
        "abstract": "{Different cosmological data are consistent with an accelerated expansion produced by an exotic matter-energy component, dubbed ``dark-energy''. A cosmological constant is a possibility since it satisfies most of the observational constraints.% In this work, the consequences of such a component in the dynamics of groups of galaxies is investigated, aiming to detect possible effects in scales of the order of few Mpc. % The Lema\\^itre-Tolman model was modified by the inclusion of the cosmological constant term and, from the numerical solution of the equations of motion, a velocity-distance relation was obtained. This relation depends on two parameters: the central core mass and the Hubble parameter. The non-linear fit of such a relation to available data permitted to obtain masses for five nearby groups of galaxies and for the Virgo cluster as well as estimates of the Hubble constant.% The analysis of the present results indicates that the velocity-distance relation derived from the modified Lema\\^itre-Tolman model as well as that derived from the ``canonical''  model give equally acceptable fits to the existent data. As a consequence, any robust conclusion on the effects of the cosmological constant in the dynamics of groups could be established. The mean value of the Hubble parameter derived from the present study of local flows is $H_0 = 65\\pm 7$ km/s/Mpc.%  \\bigskip  {\\bf key words.} Galaxies: kinematics and dynamics; Cosmological parameters;  } {}{}{}{} ",
        "introduction": "Small groups of galaxies are very common structures in the universe and may contribute up to $\\sim$ 50\\% of its matter content (Huchra \\& Geller 1982; Geller \\& Huchra 1983; Nolthenius \\& White 1987). Early estimates of mass-to-light ($M/L$) ratios for groups based on the virial relation lead to values typically of the order of 170~$M_{\\odot}/L_{B,\\odot}$ (Huchra \\& Geller 1982). However, new and high quality data on galaxies situated in nearby groups (Karachentsev 2005 and references therein) yield values around 10-30~$M_{\\odot}/L_{B,\\odot}$, considerably smaller than past estimates. If these new estimates are correct, then the local matter density would be only a fraction of the global matter density. The virial relation gives trustful masses only if groups are in dynamical equilibrium. Observers assume, in general, that this is the case if the crossing time, ratio between a suitable defined radius, characterizing the dimension of the group and the typical velocity of galaxies belonging to the system, is less than the Hubble time. Using cosmological simulations, Niemi et al. (2007) identified groups with the same algorithm used by observers. Their analysis indicates that about 20\\% of them are not gravitationally bound and that such a fraction increases when the apparent magnitude limit of the survey increases. Moreover, Niemi et al. (2007) do not have found any correlation between the virial ratio $2T/\\mid W\\mid$ and the crossing time, a result independent of the magnitude limit of the mock catalogue and that weakens the criterium usually adopted to characterize the dynamical equilibrium.  Besides the question related to the mass estimates of groups, a second and long-standing problem concerns the fact that dispersion of the peculiar velocities over the Hubble flow is quite small, usually referred as the ``coldness'' of the local velocity flow (Sandage \\& Tammann 1975; Giraud 1986; Schlegel et al. 1994; Ekholm et al. 2001). The presence of the dark energy has been invoked as a possible explanation for the smoothness of the local Hubble flow (Chernin 2001; Teerikorpi et al. 2005). Dark matter simulations by Governato et al. (1997) for cosmological models with $\\Omega_m = 1$ or $\\Omega_m = 0.3$ are, according to these authors, unable to produce systems embedded in regions having ``cold'' flows, i.e., with 1-D dispersion velocities of the order of 40-50 km/s. From simulations based on a $\\Lambda CDM$ cosmology, Macci\\`o et al. (2005) and Peirani \\& de Freitas Pacheco (2006, hereafter PP06) estimated values for the 1-D dispersion velocity, averaged within a sphere of $\\sim$ 3 Mpc radius, of 80 and 73 km/s respectively.  New simulations were recent reported by Hoffman et al. (2007), who have compared results issued from $CDM$ and $\\Lambda CDM$ cosmologies with identical parameters, except for the presence or not of the cosmological constant term. They claim that no significant differences are noticed in the velocity flow around galaxies having properties similar to those observed in the neighborhood of the Milky Way (MW).  In the case of groups constituted by a central and dominant mass, corresponding to a single or to a pair of massive galaxies and an external ``cloud'' of low mass satellites, an alternative approach to the virial relation was already proposed in the eighties by Lynden-Bell (1981) and Sandage (1986), based on the Lema\\^itre-Tolman (LT) model (Lema\\^itre 1933; Tolman 1934). The LT model describes quite well the dynamics of a bound central core embedded by an extended halo, which approaches asymptotically a homogeneous Friedmann background. In this situation, three main distinct regions can be distinguished: i) the central core, in which the shell crossing has already occurred, leading to energy exchanges which transform radial into random motions; ii) the zero-velocity surface, boundary which separates infalling and expanding bound shells and iii) the ``marginally'' bound surface (zero total energy), segregating bound and unbound shells. Density profiles resulting from the LT model were discussed by Olson \\& Silk (1979) and an application of this model to the velocity field close to the Virgo cluster was made by Hoffman et al. (1980), Tully \\& Shaya (1984) an Teerikorpi et al. (1992), among others. The LT model, modified by the inclusion of a cosmological constant, was revisited by PP06, who applied their results for a sample of galaxies in the outskirt of the Virgo cluster and for dwarf galaxies in the vicinity of the M31-MW pair.  The velocity-distance ($v$-$R$) relation is a photograph of the kinematic status of the different shells at a given instant. Such a relation, derived either from the ``canonical'' or the modified LT model, depends only on two parameters: the core mass and the Hubble parameter. Thus, if pairs of values (v, R) are known for satellite galaxies of a given group, a non-linear fit can be performed between data and the theoretical $v$-$R$ relation, permitting a simultaneous determination of the core/central galaxy mass and of the Hubble parameter (PP06). In this paper, the $v$-$R$ relation modified by the inclusion of a cosmological constant, was calculated for the present age of the universe by using an improved integration algorithm with respect to that employed originally by PP06. The fit of the resulting numerical values gives a slight different solution, which was applied to data on the groups M81, Sculptor, CenA/M83 and IC342/Maffei-I, allowing estimates of their masses as well as of the Hubble parameter $H_0$. The previous analysis by PP06 of the Local Group and of the Virgo cluster was also revisited. As we shall see, the inclusion or not of the cosmological constant in the LT model does not affect considerably the resulting masses. However, the resulting values of the Hubble parameter are systematically higher for the LT model when compared to the modified LT model, but mean values resulting from both approaches are consistent, within the uncertainties, with other independent determinations. This paper is organized as follows: in Section 2 a short review of the model is presented; in Section 3 an application to nearby groups is made and, finally, in Section 4 the main conclusions are given. ",
        "conclusions": "The velocity profile for the LT model, modified by the inclusion of a cosmological constant, was calculated. The inclusion of a such a term in the equation of motion modifies some characteristics of the ``canonical'' LT model. Shells inside the zero-velocity surface collapse earlier and, as a consequence, for a given distance negative velocities higher than those derived from the ``canonical'' model are obtained. Moreover, shells whose maximum expansion radius is beyond the critical value $R\\sim 1.39R_0$ will never collapse, since their acceleration becomes positive.  Data on dwarf galaxies belonging to nearby groups and galaxies in the outskirt of the Virgo cluster are well represented by such a model, indicating that these objects are either collapsing or expanding, to fallback in the future, if their distances are smaller than $\\sim 1.39R_0$. Moreover, such a good agreement between the theoretical $v$-$R$ relation and data implies also that the use of the virial to estimate the core masses is questionable.  However, the $v$-$R$ relation derived from the ``canonical'' LT model gives an equally acceptable fit and the present results cannot be used as an argument in favor of the detection of effects due to the cosmological constant in scales of the order of few Mpc. Core masses derived from both models agree to within 10\\% but the Hubble parameter estimated from the ``canonical'' LT model is systematically higher than values resulting from the modified model. Nevertheless, mean values are compatible with other independent estimates.  The mean value of the 1-D dispersion velocity derived from the five groups (Virgo excluded) investigated in this work is $43\\pm 7 km/s$, smaller than  values derived from cosmological simulations within scales of 1-3 Mpc, i.e., 73 km/s (PP06) and 80 km/s (Macci\\`o et al. 2005). It is worth mentioning that Axenides \\& Perivolaropoulos (2002) studied the dark energy effects in the growth of matter fluctuations in a flat universe. They concluded that the dark energy can indeed cool the local Hubble flow but the required parameters to make the predicted dispersion velocity of the order of 40 km/s are ruled out by observations which constrain either the present dark energy density or the equation of state parameter w(=$P_x/\\varepsilon_x$). Thus, the dark energy with a time independent equation of state cannot explain the observed quietness of the local Hubble flow, which remains an enigma."
    },
    "0806/0806.4904_arXiv.txt": {
        "abstract": "Polaris, the nearest and brightest classical Cepheid, is a single-lined spectroscopic binary with an orbital period of 30 years. Using the High Resolution Channel of the Advanced Camera for Surveys onboard the {\\it Hubble Space Telescope\\/} ({\\it HST\\/}) at a wavelength of $\\sim$2255~\\AA, we have directly detected the faint companion at a separation of $0\\farcs17$. A second {\\it HST\\/} observation 1.04~yr later confirms orbital motion in a retrograde direction. By combining our two measures with the spectroscopic orbit of Kamper and an analysis of the {\\it Hipparcos\\/} and FK5 proper motions by Wielen et al., we find a mass for Polaris~Aa of $4.5^{+2.2}_{-1.4} \\, M_\\odot$---the first purely dynamical mass determined for any Cepheid. For the faint companion Polaris~Ab we find a dynamical mass of $1.26^{+0.14}_{-0.07} \\, M_\\odot$, consistent with an inferred spectral type of F6~V and with the flux difference of 5.4~mag observed at 2255~\\AA\\null. The magnitude difference at the $V$ band is estimated to be 7.2~mag. Continued {\\it HST\\/} observations will significantly reduce the mass errors, which are presently still too large to provide critical constraints on the roles of convective overshoot, mass loss, rotation, and opacities in the evolution of intermediate-mass stars.  Our astrometry, combined with two centuries of archival measurements, also confirms that the well-known, more distant ($18''$) visual companion, Polaris~B, has a nearly common proper motion with that of the Aa,Ab pair. This is consistent with orbital motion in a long-period bound system. The ultraviolet brightness of Polaris~B is in accordance with its known F3~V spectral type if it has the same distance as Polaris~Aa,Ab. ",
        "introduction": "Cepheid variable stars are of central importance in galactic and extragalactic astronomy.  They are the primary standard candles for measuring extragalactic distances, and they provide critical tests of stellar-evolution theory. Surprisingly, however, until now there has not been a single Cepheid with a purely dynamical measurement of its mass.  Polaris ($\\alpha$ UMi) is the nearest and, at 2nd magnitude, the brightest classical Cepheid, albeit one with a small light amplitude in its 3.97-day pulsation period (Turner et al.\\ 2005 and references therein). The amplitude, which had been decreasing for several decades, now appears to have stabilized and may be increasing again (Bruntt et al.\\ 2008 and references therein). The {\\it Hipparcos\\/} parallax of Polaris indicates a luminosity consistent with pulsation in the first overtone (Feast \\& Catchpole 1997).  Polaris is the brightest member of a triple system (see Kamper 1996 and references therein). The well-known visual companion, Polaris~B, is an 8th-magnitude F3~V star at a separation of $18''$.  The Cepheid itself is a member of a single-lined spectroscopic binary with a period of 30~yr. In this paper we report the first direct detection of the close companion, from which we derive the first entirely dynamical mass measurement for a Cepheid.  Cepheid masses are a key parameter for testing stellar evolutionary calculations.  Beginning in the 1960's, discrepancies were found in the sense that Cepheid masses derived from pulsation modeling were lower than those derived from evolutionary tracks.  A revision in envelope opacities brought pulsation and evolutionary masses closer together, partially alleviating this ``Cepheid mass problem.''  However, recent evolutionary and pulsation constraints for Galactic (Bono et al.\\ 2001b; Caputo et al.\\ 2005; Keller 2008; Natale, Marconi, \\& Bono 2008) and Magellanic (Bono, Castellani, \\& Marconi 2002; Keller \\& Wood 2006) Cepheids still imply a discrepancy in masses at the $\\sim$15-20\\% level.  The luminosities of Cepheids depend on the mass of the helium-burning core; physical mechanisms affecting the helium core mass include mixing due to convective core overshoot during the main-sequence phase, mass loss, stellar rotation, and radiative opacity. A directly measured mass for Polaris would thus provide an important constraint on this theoretical framework. ",
        "conclusions": "The results of this study are as follows:  \\begin{enumerate}  \\item We have used UV imaging with the ACS/HRC onboard the {\\it HST\\/} to make the first direct detection of the close companion of the classical Cepheid Polaris.  \\item We confirm orbital motion in a retrograde sense, based on two observations a year apart.  \\item By combining our {\\it HST\\/} measurements with the single-lined spectroscopic orbit (Kamper 1996) and the FK5 and {\\it Hipparcos\\/} proper motions (Wielen et al.\\ 2000), we derive a dynamical mass for the Cepheid Polaris~Aa of $4.5^{+2.2}_{-1.4}\\,M_\\odot$---the first purely dynamical mass for any Cepheid.  \\item The dynamical mass is smaller than values estimated either from pulsational properties or evolutionary tracks, but the error bars are still large enough that the discrepancies have not achieved statistical significance.  \\item The close companion Polaris Ab has a dynamical mass of $1.26^{+0.14}_{-0.07}\\,M_\\odot$. This is consistent with a spectral type of about F6~V, inferred from the UV brightness of Ab.  \\item The more distant and well-known companion Polaris~B has a UV flux consistent with its known spectral type of F3~V, lying at the same distance as the Cepheid. The proper motion of Polaris~B is shown to be very similar to that of Aa,Ab, consistent with motion in a wide but bound orbit around the close pair.  \\item Continued {\\it HST\\/} imaging, including two more observations that have been approved for our own program, will decrease the errors on the dynamical mass of Polaris, allowing a critical test of stellar-evolution theory and the influence of such effects as convective overshoot, mass loss, rotation, and opacities.  \\end{enumerate}"
    },
    "0806/0806.3715_arXiv.txt": {
        "abstract": "The mid- to far-infrared emission of the outbursting FU Orionis objects has been attributed either to a flared outer disk or to an infalling envelope.  We revisit this issue using detailed radiative transfer calculations to model the recent, high signal-to-noise data from the IRS instrument on the {\\em Spitzer Space Telescope}.  In the case of FU Ori, we find that a physically-plausible flared disk irradiated by the central accretion disk matches the observations. Building on our previous work, our accretion disk model with outer disk irradiation by the inner disk reproduces the spectral energy distribution between $\\sim$~4000~\\AA~ to $\\sim 40 $~$\\mu$m. Our model is consistent with near-infrared interferometry but there are some inconsistencies with mid-infared interferometric results. Including the outer disk allows us to refine our estimate of the outer radius of the outbursting, high mass accretion rate disk in FU Ori as $\\sim$ 0.5 AU, which is a crucial parameter in assessing theories of the FU Orionis phenomenon.  We are able to place an upper limit on the mass infall rate of any remnant envelope infall rate to $\\sim 7 \\times$10$^{-7}$ M$_{\\odot}$yr$^{-1}$ assuming a centrifugal radius of 200 AU.  The FUor BBW 76 is also well modelled by a 0.6 AU inner disk and a flared outer disk. However, V1515 Cyg requires an envelope with an outflow cavity to adequately reproduce the IRS spectrum.  In contrast with the suggestion by Green et al., we do not require a flattened envelope to match the observations; the inferred cavity shape is qualitatively consistent with typical protostellar envelopes.  This variety of dusty structures suggests that the FU Orionis phase can be present at either early or late stages of protostellar evolution. ",
        "introduction": "The FU Orionis systems are a small but remarkable class of variable young stellar objects (YSOs) which undergo outbursts in optical light of 5 magnitudes or more \\citep{herbig77}, with a F-G supergiant optical spectra and K-M supergiant near-infrared (near-IR) spectra dominated by deep CO overtone absorption. FU Orionis objects (FUors) have been modelled as a high mass accretion disk around pre-main-sequence stars (Hartmann $\\&$ Kenyon 1985, 1987a, 1987b) to explain the infrared color excesses and the variations of both the spectral type and rotational line width with wavelength (Hartmann $\\&$ Kenyon 1996; Paper I).  \\cite{herbig77} argued that at least some FU Orionis outbursts must be repetitive, and \\cite{Lee96} suggested that this could be explained by infall from an envelope to the disk, replenishing the disk mass for further outbursts. The infall picture is also suggested by the presence of scattered light envelopes around FUors, suggesting that they are objects in early stages of star formation \\citep{herbig77,goodrich87}. In the evolution sequence, FUors may play significant role in transfering a large amount of mass ($\\geq$10$\\%$) to the central star, which is even higher than the mass accumulated in the T Tauri phase \\citep{leebook}.  However, not all FUors show the large mid-infrared (mid-IR) excesses that clearly demand dense infalling envelopes. In particular, \\cite{ALS1987} suggested that FU Ori itself had only a depleted or low-density envelope. \\cite{kenyon91} suggested that a pure flared disk model could explain FU Ori while an infalling envelope was needed for V1057 Cyg, but \\cite{turner97} proposed that both FU Ori and V1057 Cyg required flattened envelopes. Finally, taking advantage of IRS spectra obtained with the Spitzer Space Telescope, \\cite{jgreen06} concluded that the SEDs of V1057 and V1515 Cyg required envelopes and derived crude models for these objects, while FU Ori and BBW 76 might be explained with flared disks only. Some recent studies about the silicate features also show FUors can be classified as two categories and some objects are evolved (FU Ori and BBW 76) with only disks left \\citep{quanz2007}.  The IRS spectra provide us with the opportunity to perform a much more detailed SED analysis for FUors. In Paper I, we developed detailed accretion disk models to study the inner disk of FU Ori and we derived an inner disk size $\\sim$ 1 AU. In this paper, we re-examine the interpretation of mid-IR excesses, taking advantage of the IRS spectra, with more detailed radiative transfer models, to resolve the disk/envelope problem. We will describe the observational data in \\S 2. In \\S 3, the method to calculate the temperature structure of the irradiated surface and the resulting spectrum is described.  Model results for three FU Orionis objects (FU Ori, BBW 76, and V1515 Cyg) are presented in \\S 5.  Finally, in \\S 6 we discuss some implications of our results. ",
        "conclusions": "With the latest IRS spectra and a detailed radiative transfer model including irradiation, we revisited the envelope/flared disk problem of FUors. We confirmed the results of \\cite{jgreen06} that FU Ori and BBW 76 have flared outer disks while V1515 Cyg has a circumstellar envelope. However, instead of a flattened envelope and a large opening angle \\citep{jgreen06}, we derived a more reasonable cavity structure.  For FU Ori, near-infrared interferometry also supports our model, while mid-infrared interferometry does not fit perfectly and needs further modelling. Including the contribution from the outer flared disk, we gave a tighter constraint on the outer radius of the inner disk $\\sim$0.5 AU, which is essential to test outburst theories (Paper I).  We also suggest using a spectral index $n$ from 14 $\\mu$m to 40 $\\mu$m to classify FUors into early (Category I) and late (Category II), which is less sensitive to the viewing angle. The variety of FU Orionis envelopes suggests that FU Orionis phenomenon is observed at different stages of the protostars, and every YSO may experience several FU Orionis outbursts from very early to late stages of the infalling envelope.  This research was supported in part by the University of Michigan. In addition, this work was 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. We gratefully acknowledges John D. Monnier who drew our attention to the interferometry and gave us many helpful advices."
    },
    "0806/0806.0897_arXiv.txt": {
        "abstract": "Using simultaneous high spatial (1.3 arc sec) and temporal (5 and 10 s) resolution H-alpha observations from the 15 cm Solar Tower Telescope at ARIES, we study the oscillations in the relative intensity to explore the possibility of sausage oscillations in the chromospheric cool postflare loop. We use standard wavelet tool, and find the oscillation period of $\\approx$ 587 s near the loop apex, and $\\approx$ 349 s  near the footpoint. We suggest that the oscillations represent the fundamental and the first harmonics of fast sausage waves in the cool postflare loop. Based on the period ratio $P_1/P_2 \\sim$ 1.68, we estimate the density scale height in the loop as $\\sim$ 17 Mm. This value is much higher than the equilibrium scale height corresponding to H-alpha temperature, which probably indicates that the cool postflare loop is not in hydrostatic equilibrium. Seismologically estimated Alfv\\'en speed outside the loop is $\\sim$ 300-330 km/s. The observation of multiple oscillations may play a crucial role in understanding the dynamics of lower solar atmosphere, complementing such oscillations already reported in the upper solar atmosphere (e.g., hot flaring loops). ",
        "introduction": "The coupling of complex magnetic field and plasma generates variety of magnetohydrodynamic (MHD) waves and oscillations in various solar structures. These MHD waves and oscillations are one of the important candidates for coronal heating and solar wind acceleration. The idea of exploiting observed oscillations as a diagnostic tool for determining the physical conditions of the coronal plasma was first suggested by Roberts et al. (1984). Until recently, the application of this idea has been marked by the lack of high-quality observations of coronal oscillations. However, this situation has changed dramatically, especially due to space-based observations by the Solar and Heliospheric Observatory (SOHO), the Transition Region and Coronal Explorer (TRACE), and most recently with the high resolution spectra from the Hinode spacecraft. The fast kink wave is most frequently observed mode as it can be directly detected by periodic spatial displacement of coronal loop axis (Nakariakov et al. 1999; Aschwanden et al. 1999; Wang \\& Solanki 2004). Using temporal series image data from CDS/SOHO, recently O'Shea et al. (2007) have found the first evidence of fast-kink standing oscillations in the cool transition region loops as well. On the other hand, the fast MHD sausage wave causes the variation of pressure and magnetic field in a coronal loop and therefore can be observed as intensity oscillations or periodic modulation of coronal radio emission (Nakariakov et al. 2003). Recently, Srivastava  et al. (2008) have reported the signature of the leakage of chromospheric magnetoacoustic oscillations into the corona near the south pole. These observations provide a basis for the estimation of coronal plasma properties (Nakariakov \\& Ofman 2001).  Recently, Verwichte et al. (2004) have detected interesting phenomenon of simultaneous existence of fundamental and first harmonics of fast kink oscillations (see also De Moortel \\& Brady 2007; Van Doorsselaere et al. 2007). However, the ratio between the periods of fundamental and first harmonics $P_1/P_2$ was significantly shifted from 2, which later was explained as a result of longitudinal density stratification in the loop (Andries et al. 2005; McEwan et al. 2006). The rate of the shift allows to estimate the density scale height in coronal loops, which can be few times larger compared to its hydrostatic value (Aschwanden et al. 2000).  In this paper, we report the first evidence of fundamental and first harmonics of fast sausage oscillations in cool postflare loop. We use high spatial (1.3 arc sec) and temporal (5 and 10 s) resolution H-alpha observations from the 15 cm Solar Tower Telescope at ARIES, Nainital, India, to study the oscillations in the relative intensity. Using the standard wavelet tool, we find multiple oscillation periods at different parts of the loop, which are interpreted as the fundamental and first harmonics of fast sausage mode. In section 2, we describe the observations and data reduction. We describe the wavelet analysis in section 3. In section 4, we present our theoretical model and its link to coronal seismology. We present results and discussion in the last section. ",
        "conclusions": "}  Using high spatial (1.3 arc sec) and temporal (5 and 10 s) resolution H-alpha observations, we found intensity oscillations with different periodicity in different parts of cool postflare loop. The intensity shows the oscillation with the period of $\\approx$ 587 s at the loop apex and with the period of $\\approx$ 349 s near loop footpoint. We interpret the oscillations as signature of the fundamental and first harmonics of fast sausage mode. It is difficult to say whether the oscillations are due to the trapped or leaky modes. However, intensity plot likely shows the decrement of the oscillation amplitude, therefore we suggest the presence of leaky nature of these modes. Seismologically estimated Alfv\\'en speed outside the loop is $\\sim$ 300-330 km/s. Using the period ratio $P_1/P_2$, we also estimate the density scale height in the loop as $\\sim$ 17 Mm. This value is much higher than the equilibrium scale height at low H-alpha temperature, therefore we suggest that cool postflare loops are not in hydrostatic equilibrium. In the movie of the observations, we clearly see the unidirectional mass motion from western foot point to eastern foot point in the selected loop which indicates its departure from hydrostatic equilibrium.  In conclusions, we report the first observational evidence of multiple oscillations of fast sausage modes in cool chromospheric postflare loop. Future detailed observational search should be investigated, especially using recent space-based observations."
    },
    "0806/0806.4150_arXiv.txt": {
        "abstract": "We apply magnetohydrodynamic (MHD) modeling to the radio galaxy Hercules A for investigating the jet-driven shock, jet/lobe transition, wiggling, and magnetic field distribution  associated with this source.  The model consists of magnetic  tower jets in a galaxy  cluster environment, which has been discussed in a series of our papers.  The profile of underlying ambient gas plays an important role in jet-lobe morphology.  The balance between the magnetic pressure generated by axial current and the ambient gas pressure  can determine the lobe  radius.  The jet  body is confined jointly by  the external  pressure and gravity  inside the  cluster core radius $R_{\\rm  c}$, while  outside $R_{\\rm c}$  it expands  radially to form  fat  lobes  in  a  steeply  decreasing  ambient  thermal  pressure gradient.  The  current-carrying jets  are responsible for  generating a strong,   tightly   wound  helical   magnetic   field.   This   magnetic configuration will be unstable  against the current-driven kink mode and it visibly grows  beyond $R_{\\rm c}$ where a  separation between the jet forward  and return  currents  occurs.  The  reversed  pinch profile  of global  magnetic  field  associated  with  the jet  and  lobes  produces projected $\\bolB$-vector distributions aligned with the jet flow and the lobe edge. AGN-driven shock powered  by the expanding magnetic tower jet surrounds the jet/lobe  structure and heats the ambient  ICM.  The lobes expand subsonically; no  obvious hot spots are produced  at the heads of lobes.  Several  key features in  our MHD modeling may  be qualitatively supported by the observations of Hercules A. ",
        "introduction": "\\label{sec:INT} In the present paper,  we continue our discussion of the dynamics  of  extragalactic  jets  in cluster  environments  within  the framework of the ``magnetic  tower'' model \\citep[]{LB94, L96}, which we have  analyzed  in  a   series  of  papers.   Magnetohydrodynamic  (MHD) mechanisms are  frequently invoked to model  the launching, acceleration and  collimation  of  astrophysical  jets \\citep[see,  {\\it  e.g.},][and references  therein]{F98,  M01}.   An underlying  large-scale  (coronal) poloidal  field for  producing the  magnetically driven  jets  is almost universally assumed in  many theoretical/numerical models.  However, the origin and existence of such  a galactic magnetic field are still poorly understood.   In   contrast   with   the   large-scale   field   models,   Lynden-Bell \\citep[]{LB94,  L96}  examined the  expansion  of  the local  force-free magnetic loops anchored to the star  and the accretion disk by using the semi-analytic  approach.   Twisted  magnetic  fluxes  due  to  the  disk rotation make the magnetic loops  unstable and splay out at a semi-angle 60\\degr \\  from the rotational  axis of the disk.   Global magnetostatic solutions of  magnetic towers with  external thermal pressure  were also computed  by   \\citet[]{L01}  using  the   Grad-Shafranov  equation  in axisymmetry  \\citep[see also,][]{L02, LR03,  UM06}.  Full  MHD numerical simulations  of magnetic  towers  have been  performed in  two-dimension (axisymmetric)  \\citep[]{R98, T99, U00,  K02, K04a}  and three-dimension \\citep[]{K04b}.   Magnetic towers  are also  observed in  laboratory experiments \\citep[]{HB02, L05}.   The first in our series, \\citet[][hereafter Paper I]{L06}, described the basic assumptions  and approaches in the numerical  modeling of magnetic tower  jets.    The  evolution   of  the  tower   jets  in   a  constant density/pressure  background  was examined  there.   The  second in  our series,  \\citet[][hereafter  Paper   II]{N06}  investigated  the  global structure  of  magnetic   tower  jets  in  a  gravitationally-stratified atmosphere,  in   terms  of  the   MHD  wave  structure,   radial  force equilibrium, and  collimation.  A large current flowing  parallel to the jet  bulk flow  plays  an essential  role  in the  determining the  lobe radius,  as does  the background  pressure  profile.  The  third in  our series,  \\citet[][hereafter Paper  III]{N07} investigated  the stability properties  of magnetic  tower jets  on cluster  scales.  Current-driven instabilities are  responsible for the  non-axisymmetric structures; the external  ``kink'' ($m=1$)  mode grows  outside $R_{\\rm  c}$,  while the internal  ``double  helix''  ($m=2$)  mode  grows  predominantly  inside $R_{\\rm c}$.  This  is the forth in our series of  papers to examine the nonlinear magnetic tower jets. Here we examine our model's applicability to  an individual  source, Hercules  (Herc)  A (3C  348), by  performing three-dimensional MHD simulations.   Herc A  is one of  the most powerful  double-lobed radio sources  in the sky.  Its total  energy content is approximately $3  \\times 10^{60}$ erg and it is  identified with the central compact diffuse  (cD) galaxy of a cluster  \\citep[]{SH93} at  a low  redshift of  $z=0.154$ \\citep[]{S99}. The X-ray luminosity of $4.8 \\times 10^{44}$ erg s$^{-1}$ \\citep[]{GL04} is  thought to be  due to  bremsstrahlung of  the very  high temperature intracluster medium (ICM).   Herc A has a peculiar  radio structure; the morphology looks  like a  classical double-lobed Fanaroff-Riley  type II (FR II)  radio source  \\citep[]{FR74}, but it  has no  compact hotspots, instead   showing  an   unusual  jet-dominated   morphology  as   in  an Fanaroff-Riley type I  (FR I) radio source, which  was first revealed by \\citet[]{DF84}.  Furthermore, the two jets in Herc A are quite different in  appearance.  The  eastern jet  appears to  have  continuous twisting flows, while  the western  jet leads to  a unique sequence  of ``rings''. The jets in Herc A are well-collimated initially but they flare suddenly at a certain distance from the nucleus.  An explanation for the presence of such different  structures in the same source  is still unresolved in the literature.  Herc  A is not a typical FR I  source, either; its jets are well-collimated and has knot components.  Therefore, Herc A might be classified an intermediate case: FR I/II.   Various models  have been proposed  for explaining the formation  of the jets and rings  in Herc A, including both  kinematic \\citep[]{M88, MS96, SM02}  and dynamic  \\citep[]{M91, S02}  points of  view.   By performing axisymmetric,  two-dimensional  hydrodynamic simulations,  \\citet[]{M91} proposed a model  in which an {\\em over}-dense  and {\\em over}-pressured jet undergoes  a sudden expansion when it  becomes highly over-pressured compared  to the background  gas at  a certain  distance from  the core. This causes  a subsequent  conical expansion, giving  rise to  the ``ice cream cone'' shape of the lobe. However, at present, there is nothing in the  X-ray  observations  to  suggest  that the  jet  becomes  thermally over-pressured  as  required  in  this  model.   In  \\citet[]{S02},  the formation of  ring structure  in the western  jet is interpreted  as the result  of  ribbon-like,  annular  shocks propagating  through  the  jet cocoon.  Also, these  authors conclude that the absence  of hot spots in Herc A results  from the dynamics of turbulent  and entraining backflows in the radio lobe.   Recently,  \\citet[]{GL03}  have revealed  detailed  structures of  total intensity, spectral index, polarization, and projected magnetic field in Herc A by  using multifrequency VLA imaging.  The  jet disruption in the eastern  lobe  is clearly  visible,  implying  flow instabilities.   The projected magnetic field, after correcting for Faraday rotation, closely follows the edges  of the lobes, the jets, and the  rings; the two lobes have generally  similar field patterns.  This  observational picture may disagree  with the  turbulent structure  inside the  lobes  suggested by \\citet[]{S02}.   Magnetic field distributions associated  with the extragalactic jets and their  comparisons with  the observations  have been  examined  by using Stokes  parameters, Faraday  rotation  measure (RM),  and the  projected magnetic  field ($\\bolB$-vector)  \\citep[e.g.,][]{L81,  C89, C92,  HR99, HR02,  LCB06}. In  general, helical  magnetic fields  produce asymmetric transverse  brightness and polarization  profiles; they  are symmetrical only if  the field is purely toroidal  or the jet is  at 90$^{\\circ}$ to the line of sight  \\citep[]{L81}.  \\citet[]{LCB06} conclude that, to the first approximation,  the field  configuration in FR  I jet over  10 kpc scales is a mixture of toroidal and longitudinal components.   \\citet[]{GL04}'s {\\it ROSAT} X-ray  observations of the intracluster gas in the Herc A cluster  have revealed extended X-ray emission coming from a  compact source  at the  center of  cluster, out  to a  radius  of 2.2 Mpc. The  azimuthally-averaged X-ray surface brightness  profile is well fitted by a modified King ($\\beta$) model \\citep[]{K72, CF76}, with core radius  $R_{\\rm c}=121  \\pm 10$  kpc and  $\\beta=0.74 \\pm  0.03$.  Their X-ray analysis supports the conclusion  that the inner jets are confined thermally by the cluster core  gas: thermal pressure in the cluster core is almost 10 times larger than the equipartition pressure in the eastern jet  with the standard  assumptions (including  no protons).   The radio lobes  are  largely  positioned  {\\em  beyond} the  X-ray  core  radius, allowing  for  projection, so  they  are  expanding  essentially into  a power-law atmosphere  with density falling  as $R^{- 3 \\beta}  \\sim R^{- 2.22}$, quite close to the $R^{\\-2}$  profile needed to give the lobes a self-similar structure \\citep[]{F91}.   These observational evidence motivate us to apply our dynamical magnetic tower  jet model to  Herc A.  To our  knowledge, no  MHD model  has been applied to Herc A before though there are several hydrodynamic models as we have  already referred. Radio observations suggest  that the magnetic fields need to be taken  into account to discuss about jet/lobe dynamics in Herc  A. We believe that  a key clue to  understanding the transition between the jet and the lobe may be whether or not the inner jets become over-pressured  during their  propagation in  the X-ray  thermal cluster gas.   However, the jet  internal pressure  may be  magnetically, rather than  thermally, dominated.   This  is in  contrast  to previous  models \\citep[]{M91}.  Of particular interest here are the jet/lobe transition, wiggling,  and  the  magnetic  field distribution  associated  with  the eastern jet of Herc A, which  are not yet fully discussed among the past hydrodynamic models of Herc A. We  do not discuss the formation of rings in the western jet in this paper.   This paper is  organized as follows.  In \\S 2,  we outline our numerical methods.  In \\S  3, we describe our numerical  results.  Discussions and conclusions are given in \\S 4 and \\S 5. ",
        "conclusions": "By  performing 3D MHD  simulations, we  have investigated  the nonlinear dynamics of  magnetic tower  jets in a  galaxy cluster  environment.  In this fourth paper of the  series, we have applied our numerical modeling of the magnetic tower to a specific source, Hercules A.   To our knowledge,  this is the first trial of  discussing Hercules A jet by  using   a  magnetohydrodynamic  model.    Projected  magnetic  field distributions associated with the jets and lobes of Hercules A have been revealed  in  recent  radio  observation  and  therefore,  the  dynamics including magnetic fields plays a role in understanding this source.  In the present  paper, we have investigated the  jet-driven shock, jet/lobe transition,  wiggling,  and  projected  magnetic field  distribution  of Hercules  A and, in  general, these  feasutres can  be explained  by our magnetic tower model.   Our conclusions are summarized as follows: \\begin{enumerate} \\item A transition from the narrow jet  to the fat lobe can occur at the cluster core radius.  This interesting  feature can also be confirmed in recent observations  \\citep[]{GL03, GL04}.   In our models,  the tightly collimated helical  field configuration of a magnetic  tower jet expands abruptly  beyond the  core radius  due  to the  decreasing external  gas pressure.  Thus,  the pressure profile  of the surrounding ICM  plays an important role in determining the morphology of a magnetic tower jet.   \\item  The   expanding  magnetic   tower  jet  produces   the  preceding hydrodynamic shock wave  in the ambient ICM. This  can be interpreted as AGN-driven shock in X-ray observation  and the derived shock Mach number is identical to \\citet[]{N05}'s  observation result.  The magnetic tower lobes expand subsonically and thus no  hot spots are produced at the end parts of lobes. This may be a general understanding of FR I or low power FR II radio galaxies \\citep[]{K06}.   \\item The  size of the magnetic  tower lobe can  be accurately estimated from the  jet axial current  and the external  gas pressure at  the lobe edge. The  estimated lobe  size derived by  our numerical  simulation is qualitatively   comparable    with   the   physical    scale   seen   in \\citet[]{GL03}. In  our model, if the  background gas has  a steep slope ($\\beta$) in the King profile,  the magnetic jet can potentially lead to a huge magnetic tower lobe.   \\item Magnetic tower jets, which have a reversed magnetic pinch profile, are subject to current-driven instabilities.  Specifically, the apparent distortions by  the non-axisymmetric  kink mode are  visible at  the jet body inside  the lobe rather than  the central core.  By expanding the lobe the  edge of axial currents could be  free against the exciting external kink mode; this situation  may be suppressed inside the cluster core radius.   \\item The synthetic polarimetry of our magnetic-tower jets is consistent with the gross polarization features observed in the jets of Hercules A. The magnetic tower model produces projected $\\bolB$-vector distributions that are similar to those  observed in the global extragalactic jet-lobe system, especially along the jets and the lobe edges.   \\item Magnetic tower  jets may produce a self-consistent  picture of the field  configuration discussed in  several non-dynamical  model of  FR I jets.  It consists  of toroidally  dominated spine  inside the  lobe and poloidally dominated sheath at the lobe edge (a closed poloidal magnetic flux with a  toroidal magnetic component, rather than  one side directed helical configuration). \\end{enumerate}"
    },
    "0806/0806.1891_arXiv.txt": {
        "abstract": "{Cygnus~OB2 is one of the richest OB associations in the local Galaxy, and is located in a vast complex containing several other associations, clusters, molecular clouds, and HII regions. However, the stellar content of Cygnus~OB2 and its surroundings remains rather poorly known largely due to the considerable reddening in its direction at visible wavelength.} {We investigate the possible existence of an extended halo of early-type stars around Cygnus~OB2, which is hinted at by near-infrared color-color diagrams, and its relationship to Cygnus~OB2 itself, as well as to the nearby association Cygnus~OB9 and to the star forming regions in the Cygnus~X North complex.} {Candidate selection is made with photometry in the 2MASS all-sky point source catalog. The early-type nature of the selected candidates is confirmed or discarded through our infrared spectroscopy at low resolution. In addition, spectral classifications in the visible are presented for many lightly-reddened stars.} {A total of 96 early-type stars are identified in the targeted region, which amounts to nearly half of the observed sample. Most of them have featureless near-infrared spectra as expected from OB stars at the available resolution. Another 18 stars that display Brackett emission lines can be divided between evolved massive stars (most likely Be stars) and Herbig Ae/Be stars based on their infrared excesses. A component associated with Cygnus~OB9/NGC~6910 is clearly identified, as well as an enhancement in the surface density of early-type stars at Cygnus~X North. We also find a field population, consisting largely of early B giants and supergiants, which is probably the same as identified in recent studies of the inner $1^\\circ$ circle around Cygnus~OB2. The age and large extension of this population discards a direct relationship with Cygnus~OB2 or any other particular association.} {Earlier claims of the possible large extent of Cygnus~OB2 beyond its central, very massive aggregate seem to be dismissed by our findings. The existence of a nearly ubiquitous population of evolved stars with massive precursors suggests a massive star formation history in Cygnus having started long before the formation of the currently observed OB associations in the region.} ",
        "introduction": "\\object{Cygnus~OB2} is one of the richest OB associations known in our Galaxy, and it is the most nearby one of its kind at less than 2~kpc from the Sun (Kn\\\"odlseder~\\cite{knoedlseder03}). It is also a unique resource for the observational study of the upper end of the stellar mass function (e.g., Herrero et al.~\\cite{herrero99}): O stars, Wolf-Rayet (WR) stars, Luminous Blue Variables (LBVs), B[e] stars... are all classes with representatives in Cygnus~OB2. This association thus provides, in a single complex, the most complete and accessible showcase of the variety found among the hottest and most massive stars.  $UBV$ photometry reveals the presence of hundreds of early-type, heavily-obscured stars in Cygnus~OB2 (Massey \\& Thompson~\\cite{massey91}; Kiminki et al.~\\cite{kiminki07}). However, a more real measure of its richness is given by star counts in the near infrared, where the large, patchy foreground extinction in the direction of the association is greatly reduced. Using this technique, Kn\\\"odlseder~(\\cite{knoedlseder00}) has inferred a content of $\\sim 2600$~OB stars, including over 100 O-type stars and evolved stars having massive progenitors. Support for this estimate has been provided by Comer\\'on et al.~(\\cite{comeron02}) using near-infrared low resolution spectroscopy. High signal-to-noise ratio spectra in the visible of the least obscured objects of this sample carried out by Hanson~(\\cite{hanson03}) has confirmed the classification of the vast majority of the early-type star candidates listed by Comer\\'on et al.~(\\cite{comeron02}). However, it has also shown that a large fraction of them are giant and supergiant B stars rather than O stars. This suggests a picture more complex than that of a single, recent starburst less than 3~Myr ago giving rise to the association (Massey \\& Thompson~\\cite{massey91}). It also casts doubts on the actual membership of these stars in Cygnus~OB2, given their older ages and their extended spatial distribution when compared to the strongly-clustered, earliest-type stars at the center of the association. Since the precursors of the extended population dominated by early B-type giants and supergiants must have been O-type stars, the region must have been producing massive stars long before the formation of the major young clusters and associations currently observed there.  We investigate the nature of this population and its possible relationship with Cygnus~OB2 by studying the early-type stellar content in a region lying between $1^\\circ$ an $2^\\circ$ from the center of the association. Earlier results of this work have already revealed the existence of interesting objects possibly related to Cygnus~OB2 well beyond the boundaries of the central cluster. Examples are \\object{WR~142a}, a WC8 Wolf-Rayet star about $1.3^\\circ$ from the center of Cygnus~OB2 (Pasquali et al.~\\cite{pasquali02}), and the very massive O4If runaway star \\object{BD$+43^\\circ 3654$} more than two degrees away (Comer\\'on \\& Pasquali~\\cite{comeron07}). The study of the outskirts of Cygnus~OB2 is also motivated by its overlap with other structures tracing recent star formation, most notably the Cygnus~OB9 association and several compact HII regions and embedded clusters belonging to the Cygnus~X molecular cloud complex (Schneider et al.~\\cite{schneider06}).  We use the same technique based on 2MASS $JHK$ colors used in Comer\\'on et al.~(\\cite{comeron02}) to identify OB stars in the targeted region. Infrared spectroscopy allows us to identify new young intermediate-mass stars and evolved massive stars, both characterized by their emission in Br$\\gamma$, as well as other likely early-type stars without noticeable emission lines. We also provide spectral types of many of these objects whose relatively light obscuration makes them accessible to spectral classification in the blue. This paper can thus be regarded as a continuation of the previous work of Comer\\'on et al.~(\\cite{comeron02}) and Hanson~(\\cite{hanson03}), now probing the largely uncharted region beyond the limits of Cygnus~OB2 explored thus far. The reader is referred to the recent review by Reipurth \\& Schneider~(\\cite{reipurth08}) for a detailed description of the current state of knowledge of the vast extent of Cygnus star forming complexes, and, in particular, of the Cygnus~OB2 association and the Cygnus~X region. ",
        "conclusions": "}  We have extended previous work carried out mainly by Comer\\'on et al.~(\\cite{comeron02}) and Hanson~(\\cite{hanson03}) in characterizing the population of early-type stars in the vicinity of Cygnus~OB2. We have done this by identifying a magnitude-limited sample of candidate early-type stars through near-infrared imaging in the area between $1^\\circ$ and $2^\\circ$ from the center of Cygnus~OB2, confirming them through near-infrared spectroscopy, and providing an accurate spectral classification of many of them through spectroscopy in the visible. Our main results can be summarized as follows:  \\begin{itemize}  \\item The studied area is rich in early-type stars, including evolved massive emission-line stars, likely Herbig Ae/Be stars, and some deeply-embedded massive young stellar objects. As in previous studies, a color-based selection criterion is found to be efficient in selecting candidate early-type stars, although with a considerable degree of contamination by other types of objects, mainly cool stars and foreground late-B, A, and F stars. Near-infrared spectroscopy at modest resolution and signal-to-noise ratio suffices to identify the actual emission-line stars.  \\item Our early-type selection criteria recovers many stars belonging to the Cygnus~OB9 association and its cluster NGC~6910, which are partly included in the studied area. We also find a significant number of reddened early-type stars and Herbig Ae/Be star candidates in the area occupied by the Cygnus~X North complex, which includes the compact HII regions DR~17, DR~21, DR~22, and DR~23. The massive star distribution is, however, not confined to the HII regions.  \\item Many early-type stars are found not to be associated with any known structure. Visible spectroscopy is available for most of this field population, and shows that it is strongly dominated by giant and supergiant early-B stars. The characteristics of this population appear to be the same as discussed by Hanson~(\\cite{hanson03}) for the additional early-type stars identified by Comer\\'on et al.~(\\cite{comeron02}) in the vicinity of Cygnus~OB2. It most probably represents a population significantly older than the massive stellar cluster at the center of Cygnus~OB2.  \\item We have extended the color-based selection criterion to estimate the surface density of early-type candidates in an area of 200~square degrees centered on Cygnus~OB2, statistically correcting for the unrelated background and foreground populations. We find that the field population is not particularly related to Cygnus~OB2 nor to any other structure of the region, being all-pervasive. The need for massive precursors to produce this extended field population suggests that massive star formation in the whole Cygnus region has proceeded for a long time, well before the currently identified OB associations were formed. Such interpretation rests on the assumption underlying this work that all the early-type stars are located at similar distances, and it would be challenged if the field population were either foreground or background to the OB associations, or if it spread over a wide range of distances. Although this possibility needs to be clarified by future studies, we do not find evidence in our temperature-absolute magnitude diagrams for a systematic difference between the distance to the members of this population and the adopted distance of Cygnus~OB2.  \\end{itemize}  We believe that the present study conclusively dismisses the case for a large extent of Cygnus~OB2 much beyond the boundaries of its central concentration, which was already suggested by Garmany \\& Stencel~(\\cite{garmany92}) and supported by Comer\\'on et al.~(\\cite{comeron02}), but then questioned by Hanson~(\\cite{hanson03}). Previous studies such as those by Massey \\& Thompson~(\\cite{massey91}) and Kn\\\"odlseder~(\\cite{knoedlseder00}) clearly show that most of the massive stellar content of the association still awaits identification and classification (see also Kiminki et al.~\\cite{kiminki07}, for recent work on the confirmation of new OB members in the inner regions of Cygnus~OB2). Our results indicate that future studies of Cygnus~OB2 aiming to characterize its upper main sequence should concentrate on heavily-reddened members near its core, rather than on identifying new members at large distances from it.  We note that over half of the early-type stars for which we have obtained confirming near-infrared spectroscopy do not yet have a detailed spectral classification available. This is particularly important for the new candidates identified in Cygnus~X North, whose proximity to compact HII regions, some of which contain embedded clusters, leads us to suspect a young age for most of them. If confirmed, the extent and composition of this group may warrant its status as an OB association in its own right, although we feel that a better characterization is still needed before taking that step. On the other hand, the as yet unclassified stars located elsewhere offer a potential for the discovery of new interesting members deserving detailed study, such as WR~142a or BD$+43^\\circ 3654$. The fact that the vast majority of stars that form the basis for the present study are recognized here as early-type members of the region for the first time is a reminder of the large amount of observational work still needed to characterize the stellar population in the Cygnus region. The importance of this task is stressed by the rapid progress being made at different wavelengths in the characterization of its interstellar content."
    },
    "0806/0806.1069_arXiv.txt": {
        "abstract": "This paper describes a conceptual framework for understanding kinetic plasma turbulence as a generalized form of energy cascade in phase space. It is emphasized that conversion of turbulent energy into thermodynamic heat is only achievable in the presence of some (however small) degree of collisionality. The smallness of the collision rate is compensated by the emergence of small-scale structure in the velocity space. For gyrokinetic turbulence, a nonlinear perpendicular phase mixing mechanism is identified and described as a turbulent cascade of entropy fluctuations simultaneously occurring at spatial scales smaller than the ion gyroscale and in velocity space. Scaling relations for the resulting fluctuation spectra are derived. An estimate for the collisional cutoff is provided. The importance of adequately modeling and resolving collisions in gyrokinetic simulations is biefly discussed, as well as the relevance of these results to understanding the dissipation-range turbulence in the solar wind and the electrostatic microturbulence in fusion plasmas. ",
        "introduction": "\\label{sec_intro}  What is turbulence? Modulo many definitional and interpretational subtleties \\cite{Frisch,Tsinober}, {\\em turbulence is multiscale disorder}: we tend to say that we are dealing with a turbulent system if we have detected (measured, observed, simulated, intuited) chaotic fluctuations of some field(s) over a broad range of scales. In plasmas, these fluctuating fields are the electric and magnetic fields and the distribution function of the particles (either measured directly or accessible partially via its moments: density, flow velocity, temperature). So turbulence is defined as a {\\em syndrome} \\cite{Stewart_movie}: it is identifed by its symptoms.\\footnote{We thank T A Yousef for bringing to our attention this analogy, which is particularly apt in fusion contexts, where turbulence is indeed a disease that gives rise to anomalous transport, prevents plasma confinement and thus hampers humanity's progress toward the hydrogen-powered future.} The next logical step is to ask what causes the development of the problem in the first place. The short answer is energy injection: in physical systems, turbulence is stirred up by some source of energy, which is system-specific and can be in the form of direct mechanical forcing (spoon in a tea cup, supernovae in the interstellar medium), boundary conditions (airplane wing), or various instabilities feeding on background equilibrium gradients (tokamak microturbulence, solar convection, magnetorotational turbulence in accretion discs). The fluctuation energy injected into the system is nearly always dissipated into heat. Because the dissipation mechanisms available to the system have to do with its material properties (microphysics) and are usually unrelated to the energy-injection mechanism (macrophysics), there is more often than not a scale separation between the energy-injection, or energy-containing, scale (the {\\em outer scale}) and the much smaller dissipation scale (the {\\em inner scale}). In order to dissipate energy, the system has to bridge this gap and one way for this to happen is for the nonlinear interactions to fill the intermediate scale range with fluctuations --- giving rise to multiscale disorder, or turbulence (there are, of course, other ways, e.g., shock or current-sheet formation, but we will not consider them here).  The simplest illustration of the argument made above is the case of a Navier-Stokes neutral fluid, whose velocity field $\\vu$ satisfies \\bea \\label{NS} \\dd_t\\vu + \\vu\\cdot\\vdel\\vu = -\\vdel p + \\nu\\nabla^2\\vu + \\vf, \\quad \\vdel\\cdot\\vu = 0, \\eea where $p$ is pressure, $\\nu$ the molecular viscosity of the fluid, and the body force $\\vf$ stands in for the outer-scale energy injection. The kinetic energy of the fluid then satisfies \\bea {\\rmd\\over \\rmd t}\\intrV {u^2\\over2} = \\eps - \\nu\\intrV |\\vdel\\vu|^2, \\label{NS_energy} \\eea where $V$ is the system volume and $\\eps = (1/V)\\intr \\vu\\cdot\\vf$ is the injected power per unit volume. In a stationary state, the injection and dissipation terms on the right-hand side of this equation must balance, even though $\\eps$ is finite and viscosity is small, or, more precisely, the viscous term in \\eqref{NS} is negligible at the outer scale. The balance is accomplished by transferring kinetic energy to small scales, where the velocity gradients are large, compensating for the viscosity's smallness. The viscous (inner) scale to which the energy has to travel in order to be dissipated is, on dimensional grounds, $\\lvisc\\sim(\\nu^3/\\eps)^{1/4}\\sim\\lf\\Re^{-3/4}$, where $\\lf$ is the outer scale and $\\Re=\\urms\\lf/\\nu$ is the Reynolds number. The system becomes turbulent when $\\Re\\gg1$, i.e., $\\lvisc\\ll\\lf$, so fluctuations arise over a broad band of scales.\\footnote{{\\em A one-paragraph review of the Kolmogorov--Obukhov 1941 turbulence theory} \\cite{K41,Obukhov}: If it can be assumed (by no means an automatic certainty!) that the energy is transported {\\em locally} from scale to scale \\cite{Richardson}, the energy flux through the intermediate scales $\\lf\\gg\\lambda\\gg\\lvisc$ (the {\\em inertial range}) must be constant and equal to $\\eps$. Assuming that fluctuations are isotropic in the inertial range, we have $\\ul^2/\\taul \\sim \\eps$, where $\\ul$ is the characteristic relative velocity of fluid elements separated by a distance $\\lambda$ and $\\taul$ is the characteristic nonlinear interaction time ({\\em energy-cascade time}) at this scale. For a local cascade, dimensionally, $\\taul\\sim\\lambda/\\ul$ and the Kolmogorov scaling law immediately follows: $\\ul\\sim(\\eps\\lambda)^{1/3}$. \\label{fn_K41}} ",
        "conclusions": "\\label{sec_conc}  Let us now come back to the original motivation for the above developments: the necessity to understand how the distribution function is brought to collisional scales in the velocity space. We have seen that this is done by transferring the energy injected at the outer scale down to the ion and electron gyroscales via a multichannel cascade of generalized energy through phase space. Below the ion gyroscale, the phase-space nature of the cascade becomes particularly manifest as the ion distribution function simultaneously develops small scales in the gyrocentre and velocity space via a nonlinear perpendicular phase mixing process. We have described this process as a Kolmogorov-like turbulent cascade enabled by a constant flux of ion entropy and derived scaling relations for the fluctuations of the distribution function and the electric potential.  Using these scalings \\eqref{scalings}, let us now estimate the collisional cutoff in phase space. As we explained in \\secref{sec_gen}, the collisional scale is reached if the velocity-space correlation scale is $\\delta v/\\vthi\\sim(\\nui/\\omega)^{1/2}$. Using \\eqref{vperp_scale} and estimating $\\omega\\sim\\taul^{-1}$, we~get \\bea {\\dvc\\over\\vthi}\\sim {1\\over\\kc\\rho_i} \\sim \\lt(\\nui\\taur\\rt)^{3/5}\\sim {\\leps^{1/5}\\rho_i^{2/5}\\over\\mfp^{3/5}}, \\label{cutoff} \\eea where $\\leps = m_i\\ni\\vthi^3/\\eps$ and $\\taur\\sim(m_i\\ni\\rho_i^2/\\eps)^{1/3}$ is the fluctuation time scale at $\\kperp\\rho_i\\sim1$. This formula is perhaps our most consequential result for numerical applications: it tells us what it means to have a well-resolved gyrokinetic simulation of plasma turbulence and shows that the resolution requirements in the gyrocentre and velocity spaces are fundamentally linked. In this context, it is clear that adequate modeling of collisions \\cite{Abel_etal,Barnes_etal} and controlled velocity-space resolution \\cite{Barnes_vres} are imperative for gyrokinetic simulations.  No matter how small the collisional cutoff \\eqref{cutoff} is, all of the energy channelled into the sub-gyroscale entropy cascade will reach this cutoff in finite time --- roughly the nonlinear interaction scale $\\taur$ evaluated at the ion gyroscale. Since the process is nonlinear, this time is amplitude dependent. If the principle of critical balance (\\secref{sec_GK}) holds at the ion gyroscale, $\\taur$ should be roughly equal to the linear parallel propagation time scale at $\\kperp\\rho_i\\sim1$. Importantly, the time to reach the collisional cutoff does not depend on the collision rate --- just like in hydrodynamic turbulence (\\secref{sec_intro}), the time to reach the viscous scale is the turnover time at the outer scale, independent of viscosity.  Another important conclusion is that the dissipation range ($\\kperp\\rho_i>1$), even in the absence of kinetic Alfv\\'en waves, is filled with electrostatic fluctuations due to the ion entropy cascade. This is a purely kinetic effect invisible in any fluid models. In fusion plasmas, this may be relevant for identifying the nature of electrostatic fluctuations found between the ion and electron gyroscales.\\footnote{Another possible source of electrostatic fluctuations in the dissipation range is an inverse cascade of another conserved quantity of electrostatic gyrokinetics, $\\intr\\ephi^2$. Its conservation is a particular consequence of a larger set of more general gyrokinetic conservation laws valid in 2D (and only under some additional assumptions in 3D) \\cite{Tome,Plunk_PhD}. However, since its cascade is inverse, it is only relevant if there is a source of energy below the ion gyroscale --- as, e.g., in the case of ETG turbulence~\\cite{Dorland_ETG}.} In space physics, the great variability of the observed spectra in the dissipation range \\cite{Smith_etal} might be speculatively attributed to varying proportions of energy contained in the entropy and kinetic-Alfv\\'en-wave cascades~\\cite{Tome}.  These results are only the first glimpse of what one finds if one adopts the view of plasma turbulence as a kinetic cascade in phase space. We believe that further studies conducted in this vein, both numerical \\cite{Howes_etal3,Tatsuno_etal} and analytical \\cite{Tome,Plunk_PhD}, will unveil much new physics and many new and tantalizing questions.  \\ack We gratefully acknowledge continued interactions with I~G~Abel, M~A~Barnes, C~H~K~Chen, T~S~Horbury, R~Numata and T~A~Yousef, who are involved in ongoing collaborations with us on some of the topics discussed in this paper. AAS was supported by an STFC Advanced Fellowship and by the STFC Grant ST/F002505/1. GGH and TT were supported by the US DOE Center for Multiscale Plasma Dynamics. WD, GWH, GGH, GGP and TT thank the Leverhulme Trust International Network for Magnetised Plasma Turbulence (Grant F/07 058/AP) for travel support."
    },
    "0806/0806.1296_arXiv.txt": {
        "abstract": "{In the last few years we have developed stellar model atmospheres which included effects of anomalous abundances and strong magnetic field. In particular, the full treatment of anomalous Zeeman splitting and polarized radiative transfer were introduced in the model atmosphere calculations for the first time. The influence of the magnetic field on the model atmosphere structure and various observables were investigated for stars of different fundamental parameters and metallicities. However, these studies were purely theoretical and did not attempt to model real objects.} {In this investigation we present results of modelling the atmosphere of one of the most extreme magnetic chemically peculiar stars, HD\\,137509. This Bp SiCrFe star has the mean surface magnetic field modulus of about $29$\\,kG. Such a strong field, as well as clearly observed abundance peculiarities, make this star an interesting target for application of our newly developed model atmosphere code.} {We use the recent version of the line-by-line opacity sampling stellar model atmosphere code \\llm, which incorporates the full treatment of Zeeman splitting of spectral lines, detailed polarized radiative transfer and arbitrary abundances. We compare model predictions with photometric and spectroscopic observations of HD\\,137509, aiming to reach a self-consistency between the abundance pattern derived from high-resolution spectra and abundances used for model atmosphere calculation. } {Based on magnetic model atmospheres, we redetermined abundances and fundamental parameters of HD\\,137509 using spectroscopic and photometric observations. This allowed us to obtain a better agreement between observed and theoretical parameters compared to non-magnetic models with individual or scaled-solar abundances.} {We confirm that the magnetic field effects lead to noticeable changes in the model atmosphere structure and should be taken into account in the stellar parameter determination and abundance analysis.} ",
        "introduction": "\\label{intro} The atmospheric structure of magnetic chemically peculiar (CP) stars deviates from that of normal stars with similar fundamental parameters due to unusual chemistry, abundance inhomogeneities and the presence of strong magnetic field. These effects are not considered in the standard model atmosphere calculations, possibly leading to errors in the stellar parameter determination and abundance analysis. To circumvent this long-standing problem of stellar astrophysics, we have developed a new line-by-line opacity sampling model atmosphere code \\llm\\ \\citep{llm}. Using this tool in the series of recent papers, we investigated in detail the effects of anomalous Zeeman splitting \\citep{zeeman_paper1}, polarized radiative transfer \\citep{zeeman_paper2} and inclination of the magnetic field vector \\citep{zeeman_paper3} on the model structure, energy distribution, hydrogen line profiles, photometric colors and the magnitude of bolometric corrections for a grid of model atmospheres with different effective temperatures and metallicities. For the first time we were able to obtain new results applying direct and self-consistent modeling of all effects mentioned above and to answer the question how does the magnetic field act at different temperatures and what one could expect if the magnetic field is ignored in calculations of model atmosphere of magnetic CP stars. It was shown that the strength of the magnetic field is the key characteristic controlling the magnitude of the magnetic field effects and the polarized radiative transfer should be taken into account. In contrast, the orientation of the magnetic field vector does not have much influence on any of the observed stellar characteristics and, thus, can be safety ignored in the analysis routines.  So far our models with the magnetic field effects included have been developed and applied only in the context of purely theoretical studies. Here we make the first attempt to model the atmosphere of a real star. In this work we use the \\llm\\, stellar model atmosphere code to investigate the atmospheric structure of the star, taking into account individual chemical composition, anomalous Zeeman splitting and polarized radiative transfer.  HD\\,137509 (HIP\\,76011, NN~Aps) is a B9p SiCrFe chemically peculiar star with a strong reversing longitudinal field and variable lines of \\ion{Si}{ii} and iron-peak elements \\citep{mathys91,mathys_lanz}. \\citet{paper1} (hereafter \\paper1) has detected resolved Zeeman split lines in the spectrum of HD\\,137509, showing that this star is characterized by a non-dipolar magnetic field geometry with a mean surface field strength of  about $29$\\,kG. This is the second-largest magnetic field ever found in a CP star (the first place is occupied by the well-known Babcock's star, \\citet{preston}). The atmospheric parameters, $\\teff=12\\,750\\pm500\\,K$ and $\\log g=3.8\\pm0.1$, were derived in \\paper1\\ using theoretical fit to the observed \\hbeta\\, and \\hgamma\\, line profiles based on the \\atl\\, \\citep{kurucz13} model with enhanced metallicity, $[M/H]=+1.0$. The appearance of such a strong magnetic field and the presence of outstanding abundance anomalies inferred in \\paper1\\ allow us to use HD\\,137509 as a test ground for the application of the new generation magnetic model atmospheres.  In the next section we briefly describe the techniques employed to construct magnetic model atmospheres. Results of the calculations for HD\\,137509 are presented in Sect.~\\ref{results}. Main conclusions of our study are summarized in Sect.~\\ref{concl}. ",
        "conclusions": "\\label{concl}  In the present paper we have constructed advanced theoretical stellar model atmospheres incorporating accurate treatment of the individual abundances pattern, Zeeman splitting, polarized radiative transfer and compared the results with the observations of extreme magnetic CP star HD\\,137509. With the mean surface field of $\\langle B \\rangle \\approx 29$\\,kG, this object has the second largest magnetic field among CP stars. Strong overabundance of iron-peak elements and extreme underabundance of helium in the atmosphere of this star opens a possibility to investigate the importance of taking individual abundances into account when constructing model atmospheres of magnetic chemically peculiar stars. Theoretical model atmosphere calculations were compared with the hydrogen line profiles and metallic spectrum of HD\\,137509. The Str\\\"omgren and Geneva photometric parameters were also investigated. The main conclusions of our study can be summarized as follows: \\begin{itemize} \\item We found that the effect of individual abundances dominates the change of Balmer \\hbeta\\, and \\hgamma\\, line profiles compared to the model with solar-scaled abundances. This implies that once the abundances of the star have been determined using an approximate model, it is necessary to recalculate the model atmosphere in order to ensure the consistency between abundance pattern and the model structure. The magnetic field has less influence on the hydrogen lines, however it should be taken into account for stars with very strong magnetic fields. \\item Modification of the atmospheric temperature-pressure structure due to presence of the magnetic field and peculiar abundances generally has little impact of the metal line profiles comparing to that of non-magnetic one. \\item For HD\\,137509, the effect of magnetic field on photometric colors is very important for some photometric parameters. Occasionally the combined impact of the magnetic field and the realistic chemistry is more important than the effect of using only individual abundances. Generally, magnetic model atmospheres allowed us to obtain a better agreement between almost all observed and theoretical color indices. Thus, we can conclude that magnetic field should be taken into account in the analysis of stars with strong magnetic fields. \\item We showed that the analysis of the spectra of such extreme Bp stars with strong magnetic fields and unusual chemistry as HD\\,137509 requires a self-consistent approach. Once the abundances of the most important elements are derived using an approximate model atmosphere, it is necessary to recompute the model with new abundances trying to fit various observables such as hydrogen line profiles, photometric colors, and energy distribution (if available). For HD\\,137509 we found that the simultaneous fit to the hydrogen line profiles and photometrical indices employing the model with both magnetic field and individual abundances included requires as much as $1000$\\,K correction to the effective temperature and $0.4$~dex correction to the surface gravity compared to the results obtained using simple scaled-solar models. \\item We expect that the overall energy distribution of the star is strongly modified by magnetic line blanketing. However, the lack of accurate energy distribution covering meaningful wavelength region in the spectrum of HD\\,137509 precludes us from reaching robust quantitative conclusions. We emphasize that availability of the flux distributions for this and other magnetic CP stars is of great importance for the stellar parameter determination and for verification of the new generation model atmospheres of magnetic CP stars. \\end{itemize}"
    },
    "0806/0806.0962_arXiv.txt": {
        "abstract": "We have studied the rapid X-ray time variability in 149 pointed observations with the \\textit{Rossi X-ray Timing Explorer} (RXTE)'s Proportional Counter Array of the atoll source 4U~1636--53 in the banana state and, for the first time with RXTE, in the island state. We compare the frequencies of the variability components of 4U~1636--53 with those in other atoll and Z-sources and find that 4U~1636--53 follows the universal scheme of correlations previously found for other atoll sources at (sometimes much) lower luminosities. Our results on the hectohertz QPO suggest  that the mechanism that sets its frequency differs from that for the other components, while the amplitude setting mechanism is common. A previously proposed interpretation of the narrow low-frequency QPO frequencies in different sources in terms of harmonic mode switching is not supported by our data, nor by some previous data on other sources and the frequency range that this QPO covers is found not to be related to spin, angular momentum or luminosity. ",
        "introduction": "\\label{sec:intro}  Low-mass X-ray binaries (LMXBs) can be divided into systems containing a black hole candidate (BHC) and those containing a neutron star (NS). The accretion process onto these compact objects can be studied through the timing properties of the associated X-ray emission \\citep[see, e.g., ][for a review]{Vanderklis06}.  \\citet{Hasinger89} classified the NS LMXBs based on the correlated variations of the X-ray spectral and rapid X-ray variability properties. They distinguished two sub-types of NS LMXBs, the Z sources and the atoll sources, whose names were inspired by the shapes of the tracks that they trace out in an X-ray color-color diagram on time scales of hours to days. The Z sources are the most luminous;  the atoll sources cover a much wider range in luminosities \\citep[e.g. ,][and references therein]{Ford00}. For each type of source, several spectral/timing states are identified which are thought to arise from qualitatively different inner flow configurations.  In the case of atoll sources, the main three states are the extreme island state (EIS), the island state (IS) and the banana branch, the latter subdivided into lower-left banana (LLB), lower banana (LB) and upper banana (UB) states.  Each state is characterized by a unique combination of color color diagram and timing behavior. The EIS and the IS occupy the spectrally harder parts of the color color diagram (CD) corresponding to lower X-ray luminosity ($L_x$). The different patterns they show in the CD are traced out in days to weeks.  The hardest and lowest $L_x$ state is generally the \\textit{EIS}, which shows strong low-frequency flat-topped noise.  The \\textit{IS} is spectrally softer than the \\textit{EIS}.  Its power spectrum is characterized by broad features and a dominant band-limited noise (BLN) component which becomes stronger and lower in characteristic frequency as the flux decreases and the $>6$ keV spectrum gets harder.  In order of increasing $L_x$ we encounter the \\textit{LLB}, where the twin kHz QPOs are first observed, the \\textit{LB}, where dominant 10-Hz BLN occurs and finally, the \\textit{UB}, where the $<1$~Hz (power law) very low frequency noise (VLFN) dominates.  In the banana states, some of the broad features observed in the EIS and the IS become narrower (peaked) and occur at higher frequency. The twin kHz QPOs can be found in \\textit{LLB} at frequencies in excess of  1000~Hz, only one is seen in the \\textit{LB}, and no kHz QPOs are detected in the \\textit{UB} \\citep[see reviews by][for detailed descriptions of the different states]{Hasinger89,Vanderklis00,Vanderklis04,Vanderklis06}.  4U~1636--53 is an atoll source \\citep{Hasinger89} which has an orbital period of $\\sim3.8$ hours \\citep{Vanparadijs90} and a companion star with a mass of $\\sim0.4 \\ M_{\\odot}$ \\citep[assuming a NS of $\\sim1.4 \\ M_{\\odot}$, see][for a discussion]{Giles02}.  It was first observed as a strong continuous X-ray source (Norma X-1) with Copernicus \\citep{Willmore74} and Uhuru \\citep{Giacconi74}.  4U~1636--53 is an X-ray burst source \\citep{Hoffman77} which shows asymptotic burst oscillation frequencies of $\\sim581$~Hz \\citep[see e.g.][]{Zhang97,Giles02}.  This is probably the approximate spin frequency; although \\citet{Miller99} presented evidence that these oscillations might actually be the second harmonic of a neutron star spin frequency of $\\sim290$~Hz, this was not confirmed in further work by \\citet{Strohmayer01}.  \\citet{Prins97} studied the aperiodic timing behavior of 4U~1636--53 with the EXOSAT Medium Energy instrument up to frequencies of $\\sim100$~Hz both in the island and the banana state. \\citet{Wijnands97}, using observations with RXTE, discovered two simultaneous quasi-periodic oscillations (QPOs) near 900~Hz and 1176~Hz when the source was in the banana state.  The frequency difference $\\Delta\\nu$ between the two kHz QPO peaks is nearly equal to half the burst oscillation frequency, similar to what has been observed in other sources with burst oscillations or pulsation frequency $>400$~Hz.  To the extent that this implies $\\Delta\\nu \\sim \\nu_{spin}/2$, this is inconsistent with spin-orbit beat-frequency models \\citep{Wijnands03} for the kHz QPOs such as proposed by \\citet{Miller98}.  Other complications for beat frequency models include the fact that $\\Delta\\nu$ is neither constant \\citep[e.g. in Sco X-1, ][]{Vanderklis97} nor exactly equal to half the burst oscillation frequency.  Generally, $\\Delta\\nu$ decreases as the kHz QPO frequency increases, and in 4U~1636--53, observations have shown $\\Delta\\nu$ at frequencies lower as well as higher than half the burst oscillation frequency \\citep{Mendez98,Jonker02}.  \\citet{Straaten02,Straaten03} compared the timing properties of 4U~0614+09, 4U~1608--52 and 4U~1728--34 and conclude that the frequencies of the variability components in these sources follow the same pattern of correlations.  \\citet{Disalvo03}, based on five detections of kHz QPOs in 4U~1636--53 in the banana state was able to show that at least in that state the source might fit in with that same scheme of correlations. The detailed investigation of 4U~1636--53 is important because it is one of the most luminous atoll sources \\citep{Ford00} that shows the full complement of island (this paper) and banana states and that also shares other timing features with often less luminous atoll sources. For example, \\citet{Revnivtsev01} found a new class of low frequency QPOs in the mHz range which they suggested to be associated with nuclear burning in 4U~1636--53 and 4U~1608--52.  \\citet{Mendez00a} and \\citet{Mendez01} compared the relations between kHz QPOs and inferred mass accretion rate in 4U~1728--34, 4U~1608--52, Aql~X-1 and 4U~1636--53, and showed that the dependence of the frequency of one of the kHz QPOs upon X-ray intensity is complex, but similar among sources.  \\citet{Jonker00} discovered a third kHz QPO in 4U~1608--52, 4U~1728--34, and 4U~1636--53 which is likely an upper sideband to the lower kHz QPO.  Recently, \\citet{Jonker05} found in 4U~1636--53 an additional (fourth) kHz QPO, likely the corresponding lower sideband.  In this paper, we present new results for low frequency noise with characteristic frequencies $1-100$~Hz and QPOs in the range $100-1260$ Hz, for the first time including RXTE observations of the island state of this source. These results better constrain the timing behavior in the various states of 4U~1636--53.  We compare our results mainly with those of the atoll sources 4U~0614+09, 4U~1608--52 and 4U~1728--34 and find that the frequency of the hectohertz component may not be constant as previously stated, but may have a sinusoidal like modulation within its range from $\\sim100$ to $\\sim250$~Hz.  Our results also suggest that the mechanism that sets the frequency of the hHz QPOs differs from that for the other components, while the amplitude setting mechanism is common. Finally, we demonstrate that it is not possible to clearly distinguish between two harmonics of the low-frequency QPO $L_{LF}$ across different sources, as was previously thought \\citep{Straaten03}.   \\begin{figure*}[!hbtp] \\center \\resizebox{0.7\\columnwidth}{!}{\\rotatebox{0}{\\includegraphics{./f1.eps}}} \\caption{Hard color versus soft color normalized to the Crab Nebula as explained in Section~\\ref{sec:intro}. Each circle represents the average soft/hard color of one of the observations used for this paper. The filled triangles mark averages of 1 to 32 observations and are labeled with letters, in order from ($A$) the island state, Lower Left Banana ($J$) to the Lower banana ($N$). } \\label{fig:ccd} \\end{figure*}  \\begin{figure*}[!hbtp] \\center \\resizebox{0.7\\columnwidth}{!}{\\rotatebox{0}{\\includegraphics{./f2a.eps}}} \\resizebox{0.7\\columnwidth}{!}{\\rotatebox{0}{\\includegraphics{./f2b.eps}}} \\caption{Soft color vs. intensity (left) and hard color vs. intensity (right) in units of the Crab Nebula as explained in Section~\\ref{sec:intro}. Symbols as in Figure~\\ref{fig:ccd}. For clarity, the dashed line separates the observations corresponding to the IS (left), from the observations corresponding to the BS (right). } \\label{fig:cvsint} \\end{figure*} ",
        "conclusions": "\\label{sec:discussion}  In this paper, we report a detailed study of the time variability of the atoll source 4U~1636--53 using RXTE that includes, for the first time, observations of this source in the (low-luminosity) island state. We divided the data into 15 intervals, A to N, based on the position of the source in the color diagram.  Based on the fact that, (i) intervals A1, B, D and E show a narrow QPO with a characteristic frequency between $\\nu_b$ and $\\nu_h$ which was previously seen in other atoll sources when they were in their island state \\citep[e.g.][]{Straaten02,Altamirano05}; (ii) intervals A1, A2, B, D, E and F do not show either $L_{\\ell}$ or power-law VLFN at frequencies lower than $0.5$~Hz, which would be expected to be present in the banana state \\citep{Hasinger89,Vanderklis04,Vanderklis06}; (iii) the intensity of the source starts to increase from interval G (see Figure~\\ref{fig:cvsint}) and (iv) intervals A1 to F occupy the hardest loci in the color diagram (see Figures \\ref{fig:nuvsnu} and \\ref{fig:cvsint}), we conclude that intervals A1 to F show the source in the island state, representing the first RXTE observations of 4U~1636--53 in this state.  Interval G may represent the transition between the IS and the LLB since its power spectrum is very similar to that of the first five intervals, but with the difference that a weak ($\\sim1.2~\\%$ rms) VLFN is present at a frequency lower than $0.1$~Hz (see Figure~\\ref{fig:nuvsnu}).  Along the color color diagram we find all seven power spectral components that were already seen in other sources in previous works \\citep[see][for a review]{Vanderklis04}: $L_u$ is detected in all of our power spectra, $L_{\\ell}$ is unambiguously detected starting from power spectrum H ($\\nu_u \\sim 800$~Hz), $L_{hHz}$ is observed in 11 out of 15 power spectra at frequencies between $100$ and $270$~Hz, $L_h$ and $L_{LF}$ are detected mainly in the island state, $L_b$ is always observed and finally $L_{b2}$ is detected when $L_b$ becomes peaked from interval G.     Previous works have shown that the frequencies of the variability components observed in other atoll sources follow a universal scheme of correlations when plotted versus $\\nu_u$ \\citep[][and references therein]{Straaten03}.  We have found that the noise and QPO frequencies of the time variability of 4U~1636--53 follow similar correlations as well (see Section~\\ref{sec:results}) confirming the predictive value of this universal scheme.  However, we also found some differences between 4U~1636--53 and other atoll sources which we discuss below. As 4U~1636--53 is one of the most luminous atoll sources showing the full complement of island and banana states (full atoll track), the object is of interest in order to investigate the luminosity dependence of spectral/timing state behavior. This is of particular importance to the ongoing effort to understand the origin of the difference between the atoll sources and the more luminous Z sources \\citep[see e.g.][]{Homan06}. \\begin{figure}[!hbtp] \\center \\resizebox{1\\columnwidth}{!}{\\rotatebox{-90}{\\includegraphics{./f12.eps}}} \\caption{$\\nu_{b2}$ versus $\\nu_u$ for the atoll source 4U~1636--53 (this paper) and the characteristic frequencies of the low-frequency noise (LFN) versus $\\nu_u$ for the Z sources GX~17+2 \\citep{Homan02}, Cyg X-2 \\citep{Kuznetsov02}, GX~340+0 \\citep{Jonker00} and GX~5--1 \\citep{Jonker02} -- see text and \\citet{Straaten03} for details. (See also Figure~\\ref{fig:b2})} \\label{fig:b2z} \\end{figure}     \\subsection{The broad components in 4U~1636--53 and  Z-source LFN} As can be clearly seen in Figure~\\ref{fig:b2}, where we show $\\nu_{b2}$ versus $\\nu_u$, the behavior of the $L_{b2}$ component differs significantly between sources.  For 4U~1636--53 and 4U~0614+09, $\\nu_{b2}$ increases with $\\nu_u$, while this is not seen for 4U~1608--52 and 4U~1728--34. This frequency behavior is different from that observed for all other components (see Section~\\ref{sec:results}), which instead is very consistent between sources, even for the case of the hHz QPO, which has not been seen to correlate with other components \\citep[see][for a review]{Vanderklis04}. This unusual, somewhat erratic behavior of $L_{b2}$ may be related to the fact that it is usually detected as a relatively weak wing to a much stronger $L_b$, so that small deviations in the time-averaged shape of $L_b$ have a large effect on $L_{b2}$.  In order to investigate the relation of $L_{b2}$ to the well-known low frequency noise (LFN) which occurs in the same frequency range in Z sources, in Figure~\\ref{fig:b2z} we plot the results for 4U~1636--53 together with those for the LFN in the Z sources \\citep{Hasinger89} GX~17+2 \\citep{Homan02}, Cyg X-2 \\citep{Kuznetsov02}, GX~340+0 \\citep{Jonker00} and GX~5--1 \\citep{Jonker02b}.  Note that the broad-band noise in these Z sources was not fitted with a zero-centered Lorentzian but with a cutoff power law or a smooth broken power law. We used the results of the conversion from power laws to zero-centered Lorentzians done by \\citet{Straaten03}. Previous works \\citep[e.g.][]{Psaltis99,Straaten03} compared the time variability of Z sources with that of atoll sources and tried to associate variability components among these sources.  Based on frequency-frequency plots, only the kHz QPOs and the horizontal branch oscillations (HBO) found in the Z sources can be unambiguously identified with atoll source components, the latter with $L_h$. \\citet{Straaten03} suggested that the LFN might be identified with $L_{b2}$ and noted that (like in the case of $L_{b2}$) the characteristic frequency of the LFN, when plotted versus $\\nu_u$, does not follow exactly the same relations between Z sources.  By comparing the different frequency patterns in Figure~\\ref{fig:b2z}, we find that the behavior of the LFN component of GX~17+2, and that of $L_{b2}$ of 4U~1636--53 are similar, which might indicate that the physical mechanism involved is the same. Perhaps this is related to the fact that 4U~1636--53 is a relatively luminous atoll source \\citep[see][]{Ford00} while GX 17+2 may be a relatively low luminosity Z-source \\citep{Homan06}. Hence, 4U~1636--53 might be relatively close in $L_x$ to GX~17+2 and differ more in $L_x$ from the other two sources introduced above. Note that the time variability of GX~17+2 is different from that of the other Z sources plotted in Figure~\\ref{fig:b2z}.  For instance, the characteristic frequency of its LFN is rather low and it appears as a peak, it shows a flaring branch oscillation (FBO) and the harmonic of the HBO is relatively strong, whereas the other Z sources plotted show a flat LFN, no FBO and a weak harmonic to the HBO \\citep[][]{Jonker02b}.  As previously noted by \\citet{Kuulkers97}, these properties set GX~17+2 apart from the 'Cyg-like' Z sources GX~5--1, GX~340+0 and Cyg~X-2 and, associate it with the 'Sco-like' Z sources Sco~X-1 and GX~349+2, not plotted in Figure \\ref{fig:b2z} because no systematic study of the of the LFN and QPO behavior of these sources in terms of $\\nu_{max}$ is available as yet.  We further investigated the frequency similarities between 4U~1636--53 and GX~17+2 by plotting our results for the two sources. No clear component associations were found. GX~17+2 is the only Z source that had shown an anti-correlation between the frequency of one of its components (HBO) and the kHz QPOs at high $\\nu_u\\gtrsim1050$~Hz \\citep{Homan02}. A similar effect was seen in the atoll source 4U~0614+09 between $\\nu_b$ and $\\nu_u$ \\citep{Straaten02}.  As can be seen in Figure~\\ref{fig:nuvsnu}, a similar decrease of $\\nu_b$ with $\\nu_u$ at high frequency may occur in 4U~1636--53. However, the error bars on $\\nu_b$ are rather large in the relevant range, and the data are still consistent with $\\nu_b$ being constant at $\\nu_u\\gtrsim1100$~Hz, and marginally, even with a further increase in frequency.  It is interesting to note, that while 4U~0614+09 has a much lower $L_x$ than 4U~1636--53, both sources might show this same turnover in $\\nu_b$ versus $\\nu_u$.  Of course, these results need confirmation.  \\subsection{The low frequency QPO}  With respect to the low-frequency QPO $L_{LF}$, \\citet{Straaten03,Straaten05} observed that in their data there were two groups of sources, one where the $L_{LF}$ feature was visible, and a second one, were a QPO was detected which they suggested to be the sub-harmonic of $L_{LF}$ and therefore, called $L_{LF/2}$.  Following \\citet{Straaten03}, the upper continuous line in Figure~\\ref{fig:bumps} indicates a power law fitted to the $\\nu_{LF}$ versus $\\nu_h$ relation of the low luminosity bursters 1E~1724--3045 and GS~1826--24, and the BHC GX~339--4.  If we reproduce the fit where we take into account the errors in both axes, we find a best fit power-law index $\\alpha = 0.97\\pm0.01$ and $\\chi^2/dof = 80/19 \\sim 4.2$. If we fix $\\alpha=1$, the fit gives a $\\chi^2/dof=83/20 \\sim 4.1$.  According to the F-test for additional terms, there is a $<1\\sigma$ improvement of the fit when $\\alpha$ is set free, so we conclude that $\\nu_{LF}$ is consistent with being linearly related to $\\nu_{h}$. The lower dashed line is a power law with the same index $\\alpha=0.97$, but with a normalization half of that of the dashed line.  As can be seen in Figure~\\ref{fig:bumps}, in 4U~1636--53 the $L_{LF}$ component does not follow either of the two power-law fits. If we fit the points for 4U~1636--53, we find that the power-law index is $\\alpha_2 =1.40\\pm0.09$ ($\\chi^2/dof = 0.14/2$), significantly different from that of the other sources.  Given the above, it is probably incorrect to think that the difference in the $\\nu_{LF}$ vs. $\\nu_h$ relation between GX~339--4, GS~1826--24 and 1724--3045 on one hand and 4U~1608--52, Cyg X-1 and XTE J0929-314 on the other is associated with harmonic mode switching \\citep{Straaten03}. This conclusion is supported by the work of \\citet{Manu05} who also found a different correlation ($\\alpha=0.58\\pm0.06$, see also Figure~\\ref{fig:bumps}) in XTE~1807--294 over a much wider range of frequencies than we obtained for 4U~1636--53, by the high $\\chi^2/dof$ for the $\\nu_{LF}-\\nu_h$ fit on the data of the low luminosity bursters 1E~1724--3045, GS~1826--24, and the BHC GX~339--4 (see previous paragraph), by the fact that if we use the centroid frequencies instead of $\\nu_{max}$, the relations worsen \\citep[see][]{Straaten03}, and by the fact that the points for 4U~1728--38 \\citep{Straaten02} fall in between the two power laws, (solid and dashed line in Figure~\\ref{fig:bumps}).  Nevertheless it is interesting that the data for 4U~0614+09, 4U~1728--34, 4U~1636--53, 4U~1820--30, 4U~1608--52 XTE~1807--294, SAX~J1808.4--3658, XTE~J1814--338 and XTE~J0929--314 all fall on, or in between, the two previously defined power laws, i.e., do not deviate from a single relation by more than a factor of 2. We note that all the $\\nu_{LF}$ values discussed here could in principle have been affected by smearing in the averaging process discussed in Section~\\ref{sec:dataanalysis}. However, for smearing to shift a frequency-frequency point away from its proper value, large systematic differences are required between the two components in the dependence of amplitude or Q on frequency, and in the case of $L_{LF}$ and $L_{h}$ there is no evidence for this.  From Figure~\\ref{fig:bumps} it is apparent that the frequency range in which the $L_{LF}$ component has been identified is rather large (up to 2.5 decades).  Clearly, which frequency ranges $L_{LF}$ covers is not related to source spin frequency, angular momentum or luminosity of the object. The sources 4U~1608--52, 4U~1820--30, 4U~1636--53 and 4U~1728--34 all show $L_{LF}$ when they are in their island state, but with $\\nu_{LF}$ $\\lesssim2.6$~Hz for 4U~1608--52, and $\\gtrsim30$~Hz for the other three sources. The accreting millisecond pulsar XTE~J1807--294 shows $\\nu_{LF} \\gtrsim 12$~Hz while the AMP XTE~J0929--314 shows $\\nu_{LF} \\lesssim 1$~Hz, while both have very similar spin frequencies (191~Hz, \\citealt{Markwardt03C} and 185~Hz, \\citealt{Remillard02C}, respectively).  4U~1820--30 and 4U~1636--53 are at least one order of magnitude more luminous than XTE~1807--294 and SAX~J1808.4--3658 at their brightest \\citep[see][]{Ford00,Wijnands05}, but all four sources show $\\nu_{LF} > 5$~Hz.  The only systematic feature in the LF QPO frequencies is that while frequencies up to 50~Hz are seen in neutron stars, black holes have not been reported to exceed 3.2~Hz, nor did atoll sources in the EIS exceed 2.6 Hz. So, BHCs and NS in the extreme island state are similar in this respect; (this may be related to an overall similarity in power spectral shape for such sources in these states that was noted before; see, e.g., \\citealt{Psaltis99,Nowak00,Belloni02,Straaten02}).  \\subsection{The X-ray luminosity dependence of rms}  It has been suggested that an anti-correlation may exist between the average X-ray luminosity of different sources and the rms amplitude of their power spectral components \\citep[see discussion in][and references therein]{Jonker01,Straaten02,Straaten03}.  From Figure~\\ref{fig:rmsvsnu} we find differences in kHz QPO rms amplitudes of no more than a factor 2 between sources which differ in average luminosity by a factor up to 10, except for one point of 4U~0614+09 at $\\nu_u \\sim 1140$~Hz, where the rms of the upper kHz QPOs is a factor $\\sim7$ higher than that of the other atoll sources. \\citet{Mendez01}, \\citet{Jonker01} and \\citet{Straaten02} have already noted that the data are inconsistent with a model in which the absolute amplitudes of the kHz QPOs are the same among sources, and the decrease in rms with luminosity between sources is only caused by an additional source of X-rays unrelated to the kHz QPOs.   From Figure~\\ref{fig:rmsvsnu} it can also be seen that the largest rms amplitude differences are found in the hHz QPOs (excluding 4U~1608--52, which is a transient source covering a large $L_x$ range). For this component we find (1) 4U~0614+09, which has the strongest $rms_{hHz}$ ($>15$\\% when $\\nu_u\\lesssim 800$~Hz); (2) 4U~1636--53, which has the weakest $rms_{hHz}$  ($<10$\\% when $\\nu_u\\lesssim 800$~Hz) and (3) 4U~1728--34 which has $rms_{hHz}$ generally between those of (1) and (2) [between 10 and 15\\% when $\\nu_u\\lesssim 800$~Hz].  (At $\\nu_u\\gtrsim 800$~Hz, the groups can still be differentiated as the rms amplitude decreases with $\\nu_u$). From figure 1 in \\citet{Ford00}, it can be seen that 4U~0614+09 is the faintest X-ray source of our sample, while 4U~1636--53 is the brightest.  4U~1728--34 show luminosities between the first two.  This suggests an X-ray luminosity--rms anti-correlation for $L_{hHz}$ that is not as clear in the other components (see also Figure~\\ref{fig:LF}).    The fact that the rms of $L_{hHz}$ starts to decrease at the same $\\nu_u$ as that of $L_u$ and $L_b$, while $\\nu_{hHz}$ does not correlate with $\\nu_u$ as all other frequencies do, suggests that the frequency setting mechanism is different for $L_{hHz}$ compared with the other components, while the amplitude setting mechanism is common. As pointed out in Section~\\ref{sec:results}, the drop in rms in $L_u$, $L_{hHz}$ and $L_b$ starts at $\\nu_u$ between $700$ and $800$~Hz. For the case of 4U~1636--53 shown here, this corresponds to interval G. The power spectrum of this interval may represent the transition between the island and the banana state, when the geometric configuration of the system is thought to change \\citep[e.g.][]{Jonker00b,Gierlinski02}. For example, the appearance of a puffed-up disk could smear out the variability coming from the inner region where the oscillations are produced.  \\subsection{The nature of the hectohertz QPOs}   While our results indicate that the characteristic frequency of the hHz QPO may oscillate as a function of $\\nu_u$, $\\nu_{hHz}$ remains constrained to a limited range of frequencies (100--250~Hz) for 4U~1636--53 and for the other sources used in Figure~\\ref{fig:nuvsnu}. A similar result has been reported for $\\nu_{hHz}$ in several other atoll sources such as in MXB~1730--335 \\citep{Migliari05}, 4U~1820--30 \\citep{Altamirano05} and in the atoll source and millisecond accreting pulsar SAX~J1808.4--3658 \\citep{Wijnands98,Straaten05}.  Interestingly, the presence of $L_{hHz}$ has not been confirmed for Z-sources \\citep{Vanderklis06}, possibly due to the intrinsic differences between atoll and Z-sources such as  luminosity.   \\citet{Straaten02} have suggested a link between the $\\lesssim100$~Hz QPOs reported by \\citet{Nowak00} in the black holes Cyg~X-1 and GX~339--34 and $L_{hHz}$. \\citet{Straaten02} also suggested that $L_{hHz}$ could be related to the $\\sim67$~Hz QPO in the black hole GRS 1915+105 \\citep{Morgan97} and the $\\sim300$~Hz QPO in the BHC GRO J1655--40 \\citep{Remillard99b} which also have stable frequencies. \\citet{Fragile01} made a tentative identification of the $\\sim9$~Hz QPO in the BHC GRO~J1655--40 \\citep{Remillard99b} with the orbital frequency at the Bardeen-Petterson (B--P) transition radius \\citep{Bardeen75} and suggested the same identification for $L_{hHz}$ in neutron star systems. In this scenario, the orbital frequency at the radius where a warped disk is forced to the equatorial plane by the Bardeen--Petterson effect can produce a quasi-periodic signal \\citep[see][for an schematic illustration of the scenario]{Fragile01}.  Attempts have been made to theoretically estimate the B--P transition radius from accretion disks models in terms of the angular momentum and the mass of the compact object \\citep[e.g.][]{Bardeen75,Ivanov97,Hatchett81,Nelson00}. \\citet{Fragile01} propose a parameterization involving a scaling parameter A, which according to them  lies in the range $10\\lesssim A \\lesssim300$.  These authors write the B--P radius as $R_{BP}=A \\cdot a_{\\star}^{2/3} \\cdot R_{GR}$, where $a_{\\star}=Jc/GM^2$ is the dimensionless specific angular momentum (J and M are the angular momentum and the mass the compact object, respectively) and $R_{GR}$ is $GM/c^2$. The Keplerian orbital frequency associated with the B--P transition radius can be written as $\\nu_{kep,BP}= c^3 \\cdot (2 \\pi G)^{-1} \\cdot (M a_{\\star} A^{3/2})^{-1}$. If we assume that the atoll sources plotted in Figure~\\ref{fig:nuvsnu} all have masses between $1.4$ and $2M_{\\odot}$, that $0.3<a_{\\star}<0.7$ \\citep[see e.g.][ and references within]{Salgado94,Cook94} and that the central frequency of the hHz QPOs is between $\\sim100$ and $\\sim250$~Hz, we can constrain the scaling factor A for these source to be between $\\sim20$ and $\\sim84$. If A only depends on the accretion disk (i.e. does not depend on the central object), this can be used to constrain the frequency range in which we expect to observe $\\nu_{kep,BP}$ in black holes. For example, the black hole BHC GRO~J1655--40, whose mass is estimated from optical and infrared investigations as $M=6.3\\pm0.5M_{\\odot}$ \\citep{Greene01} and whose specific angular momentum $a_{\\star}$ can be estimated to be between 0.5 and 0.95 \\citep{Cui98,Fragile01}, would have $\\nu_{kep,BP}$ between $\\sim6.5$ and $\\sim127$~Hz, which would exclude the 300~Hz QPO observed in GRO~J1655--40 but would be consistent with the the 9~Hz QPO as proposed by \\citet{Fragile01}. If one assumes the 450~Hz QPO in GRO~J1655--40 is associated with orbital motion at the last stable orbit, then $a_{\\star}$ could be as low as $\\sim0.15$ \\citep{Strohmayer01a}. In this case, $\\nu_{kep,BP}$ can be as high as $\\sim425$~Hz for a black hole mass of $5.7 M_{\\odot}$."
    },
    "0806/0806.0824_arXiv.txt": {
        "abstract": "The accretion-powered (non-X-ray burst) pulsations of XTE J1814-338 are modeled to determine neutron star parameters and their uncertainties. The model is a rotating circular hot spot and includes: (1) an isotropic blackbody spectral component; (2) an anisotropic Comptonized spectral component; (3) relativistic time-delays and light-bending; and (4) the oblate shape of the star due to rotation. This model is the simplest possible model that is consistent with the data. The resulting best-fit parameters of the model favor stiff equations of state, as can be seen from the 3-$\\sigma$ allowed regions in the mass-radius diagram. We analyzed  all data combined from a 23 day period of the 2003 outburst, and separately analyzed data from 2 days of the outburst. The allowed mass-radius regions for both cases only allow equations of state (EOS) that are stiffer than EOS APR \\citep{APR}, consistent with the large mass that has been inferred for the pulsar NGC 6440B \\citep{Fre08}. The stiff EOS inferred by this analysis is not compatible with the soft EOS inferred from a similar analysis of SAX J1808. ",
        "introduction": "\\label{s:intro}  The accretion-powered millisecond-period X-ray pulsars are promising targets for constraining the neutron star equation of state (EOS) through the modeling of emission from hot spots on the pulsar's surface. The first pulsar discovered in this class, \\saxj \\citep{Wij98}, has a spectrum consistent \\citep{Gie02} with emission from a hot spot on the star's surface. Pulse shape modeling of rapidly rotating neutron stars relies on two relativistic effects: the gravitational bending of light rays reduces the modulation of the pulsed emission and depends on the mass to radius ratio $M/R$; and the Doppler boosting due to the star's rotation creates an asymmetry in the pulse shape and depends on the star's radius $R$. These features, combined with reasonable models of the emission properties at the neutron star's surface can be used to constrain the neutron star's mass and radius and hence the EOS of supra-nuclear density matter.  XTE J1814-338 (hereafter XTE~J1814) was discovered during outburst in June 2003 \\citep{MS03}, and is an accretion powered millisecond pulsar with spin frequency 314.36 Hz and orbital period of 4.3 hr \\citep{Metal03}. A detailed timing analysis for \\xte was performed by \\citet{Pap07} to obtain accurate values for orbital period, projected semi-major axis, pulse spin frequency and spin down rate. A similar analysis of the pulse arrival times was carried out by \\citet{Wat06} and \\citet{Chu07}, which both included an analysis of phase lags. Soft lags were found in the 2-10 keV energy band, similar to those for SAX J1808-3658 and consistent with  an origin in Doppler boosting of a Comptonized pulse component with a much broader emission pattern than the blackbody component.  \\citet{Str03} found the same frequency in the X-ray bursts as was found in the persistent emission, but with a lower second harmonic content. \\citet{Wat08} showed that the X-ray burst oscillations are tightly phase-locked with the non-burst pulsations. \\citet{Bhat05} modeled the oscillations during X-ray bursts with a hot spot model for a spherical star and for 2 equations of state. Using a large grid of models they found an upper limit on compactness $R_{S}/R<0.48$, with $R_{S}$, the Schwarzschild radius.  There are pulse shape models for a few other X-ray pulsars. The 1.2 s period X-ray pulsar \\herx was modeled by \\citet{Leahy04} using a model that includes accretion columns. The model for \\herx constrains the neutron star EOS to a fairly moderate stiffness \\citep{Leahy04}. \\citet{Zav98} and \\citet{Bog07} have modeled the X-ray emission from the 5.8 ms period radio pulsar PSR~J0437-4715 using a Hydrogen atmosphere model. In the case of PSR~J0437-4715, \\citet{Bog07} found that a simple isotropic blackbody model is inconsistent with the data. In their models, \\citet{Bog07} showed that the radius of PSR~J0437-4715 must be larger than 6.7 km if the mass is $1.4 M_\\odot$. Unfortunately the mass of this pulsar is not well-constrained. \\citet{Bog08} have shown that constraints on radius for a number of other ms radio pulsars are also possible.  Constraints on \\saxj (with a spin period of 2.5 ms) were made by \\citet{PG03} using data from the 1998 outburst. The modeling done by \\citet{PG03} included blackbody emission from a hot spot that is Compton scattered by electrons above the hot spot. Their model makes use of a spherical model for the star's surface and does not include the effects of relative time-delays caused by the different time of flights for photons emitted from different parts of the star's surface. More recently \\citet{CLM05} and \\citet{CMLC07} have shown that time-delays and the star's oblate shape are important factors that can affect the outcome of pulse-shape modeling for rapidly rotating pulsars such as \\saxj. The 1998 outburst data for \\saxj was revisited using a models that included time-delays and oblateness \\citep{LMC08} with the result that the EOS for \\saxj is constrained to be very soft.  In this paper we model the accretion-powered pulsations of \\xte using a hot-spot model. The hot-spot model allows for one or two circular hot spots with a two-component spectrum. The spectral model includes isotropic blackbody emission and an anisotropic Compton-scattered component described by a powerlaw. The photons are propagated to the observer using the oblate Schwarzschild approximation \\citep{MLCB07} which allows the photon initial conditions to be placed on an oblate-shaped initial surface determined by an empirical formula. The Schwarzschild metric is used to compute the photon bending angles and time delays since it has been shown \\citep{CMLC07} that the corrections induced by the Kerr black hole metric or a numerical metric for a rotating star are insignificant compared to the corrections induced by the oblate shape. In order to do the pulse-shape modeling, we construct light curves in two narrow energy bands, 2-3 keV and 7-9 keV. We first analyse a composite pulse-shape constructed from 23 days of data and then consider pulse-shapes constructed from single days of data in order to determine whether variations of the pulse-shape with time are significant.  The outline of this paper is as follows. In Section \\ref{s:method} the method used to construct the light curves and analyse them is outlined. The results of the best-fit models are presented in Section \\ref{s:results}. A discussion of the results is presented in Section \\ref{s:discuss}. ",
        "conclusions": "\\label{s:discuss}  We have analyzed the accretion-powered (non-X-ray burst) pulsations of \\xte using a hot spot model. Our modeling includes (1) an isotropic blackbody spectral component; (2) an anisotropic Comptonized component; (3) relativistic time-delays; (4) the oblate shape of the star due to rotation. The model presented in this paper is the simplest possible model that is consistent with the data. The resulting best-fit models favor stiff equations of state, as can be seen from the 3-$\\sigma$ allowed regions in Figures \\ref{fig:b1} and \\ref{fig:a}. In Figure \\ref{fig:b1} all data from a 23 day period of the 2003 outburst were included, while for Figure \\ref{fig:a} data from only 2 days were included. The allowed regions for the two data sets differ slightly, but both only allow equations of state that are stiffer than EOS APR \\citep{APR}.  It is interesting that a large mass has been inferred for the pulsar NGC 6440B \\citep{Fre08} through measurements of periastron precession. Assuming that the observed periastron precession is purely from relativistic effects, the pulsar's mass is $M = 2.74 \\pm 0.21 M_\\odot$ (1-$\\sigma$ error bars) \\citep{Fre08}. If the mass really is this high, it would be consistent with the stiff equations of state allowed by our analysis of \\xte. However, it is still possible that the large periastron precession observed for NGC 6440B could be caused by a very rapidly rotating companion \\cite{Fre08}, in which case the pulsar's mass would be smaller and compatible with more moderate equations of state. A high mass for \\saxj is also inferred by observations during its quiescent state \\citep{Hei07}. Modelling of the neutron star X7 in 47 Tuc \\citep{Hei06} allows for a high mass neutron star, although for X7 a  low mass neutron star is also allowed.  Similar hot spot models of \\saxj imply a soft equation of state and a column model for Her X-1 also implies a soft EOS (see \\cite{LMC08} and \\cite{Leahy04} for details). The best-fit pulse-shape models found for \\xte have mass and radius incompatible with the 3-$\\sigma$ allowed regions of mass and radius found for \\saxj or Her X-1. Some possible interpretations could be (1) time-variations in the pulse profile of \\saxj led to an underestimate of the star's radius; (2) the equation of state for dense matter has a two-phase nature allowing both large and small compact stars;  or (3) the simple hot-spot model doesn't describe one or more of these pulsars.   The first reason is a factor for \\saxj. It was demonstrated in \\cite{LMC08} that the analysis of the pulse profile averaged over a long observation gave significantly different results than the pulse profile obtained from a much shorter observation. Supporting this conclusion is a recent analysis by \\cite{Har08} of data from the 1998, 2002 and 2005 outbursts of \\saxj showing a great deal of variation in the pulse shape over time (see Figure 3 of \\citet{Har08}). The pulse-shape analysis by \\citet{LMC08} only made use of the data from the 1998 outburst. The 1998 data is very sinusoidal in nature and has very little harmonic content. The results of \\citet{Har08} show that the later outbursts have a stronger harmonic content. Since a larger radius star can produce a stronger harmonic content, it is possible that the addition of data from 2002 and 2005 will alter the conclusions of \\citet{LMC08} about SAX J1808. But the effect of pulse shape variability is not a factor for Her X-1 where the pulse shape has high stability.  The second reason above, i.e. a bimodal equation of state, is a possible, but speculative, solution to the greatly different allowed regions for M and R. In this scenario, there is still only one baryonic equation of state and one quark matter EOS, but above a certain critical density, $\\rho_{crit}$, the whole star makes a transition from baryonic  matter to quark matter. Then for stars with central density $\\rho_c$ having $\\rho_c<\\rho_{crit}$, the M vs. R relation follows a stiff baryonic EOS, somewhat like EOS L, whereas for $\\rho_c>\\rho_{crit}$ the star has converted to a quark star and lies on a quark matter EOS curve in the mass-radius diagram. In the case of EOS L, the $3 \\sigma$ region with $M<2.7 M_\\odot$ would require a value of $\\rho_{crit} \\sim 10^{15} g/cm^3$. Since quark matter EOS curves have a lower maximum mass than baryonic EOS curves, any baryonic star that makes the transition and has mass above the quark star maximum mass must lose mass to end up as a stable quark star. The mass loss depends on the physics of the transition process and is likely vary from star to star. In this scenario, we interpret \\xte to lie on the baryonic branch of the mass vs. radius diagram and \\saxj and Her X-1 to lie on the quark matter branch.  The third listed reason for the discrepancy, that the emission region models are too simple to represent the actual emission regions on the stars, is a definite possibility. For instance, if the emission is coming from the magnetosphere, then our models are incorrect. Alternatively, the emission may arise from surface spots, but the region's shape might be more complicated than a circle. This can only be tested by constructing more complex models and applying them to the observed pulse shapes. However with a more complex model with more parameters describing the model, better pulse shape data is required to constrain the model parameters. Future work is planned to explore more complex emission models, and to test whether these resolve the apparent discrepancy in mass and radius for different pulsars."
    },
    "0806/0806.0015_arXiv.txt": {
        "abstract": "We present the discovery of the widest known ultracool dwarf -- white dwarf binary. This binary is the first spectroscopically confirmed widely separated system from our target sample. We have used the 2MASS and SuperCOSMOS archives in the southern hemisphere, searching for very widely separated ultracool dwarf -- white dwarf dwarf binaries, and find one common proper motion system, with a separation of $3650-5250$\\,AU at an estimated distance of $41-59$\\,pc, making it the widest known system of this type.  Spectroscopy reveals 2MASS J$0030-3740$ is a DA white dwarf with $T_{\\rm eff}$=$7600\\pm100$K, log($g$)$=7.79-8.09$ and M$_{WD}=0.48-0.65$\\,M$_{\\odot}$. We spectroscopically type the ultracool dwarf companion (2MASS J$0030-3739$) as M9$\\pm$1 and estimate a mass of $0.07-0.08$\\,M$_{\\odot}$, $T_{\\rm eff}=2000-2400$\\,K and log($g$)=$5.30-5.35$, placing it near the mass limit for brown dwarfs. We estimate the age of the system to be $>$1.94\\,Gyrs (from the white dwarf cooling age and the likely length of the main sequence lifetime of the progenitor) and suggest that this system and other such wide binaries can be used as benchmark ultracool dwarfs. ",
        "introduction": "The rise of deep large area surveys such as 2MASS (the Two Micron All Sky Survey), SDSS (the Sloan Digital Sky Survey) and UKIDSS (the UK Infrared Deep Sky Survey) has increased the number of known brown dwarfs to several hundred, since the discovery of Gliese 229B \\citep{nakajima95} and Tiede 1 \\citep{rebolo95} just over a decade ago. These populations have helped shape our understanding of ultracool dwarfs (UCDs; with spectral type $\\ge$M7 e.g. \\citealt{jones01}) and revolutionized the classification system for sub-stellar objects including the creation of two new spectral types L and T. The latest M dwarfs ($\\sim$M7-9) have effective temperature ($T_{\\rm eff}$) reaching down to $\\sim$2300K. At lower $T_{\\rm eff}$ ($\\sim$2300-1400\\,K) are the L dwarfs, which have very dusty upper atmospheres and generally very red colours. T dwarfs are even cooler having $T_{\\rm eff}$ in the range $\\sim$1400-650\\,K, where the low $T_{\\rm eff}$ limit is currently defined by the recently discovered T8$+$ dwarfs, ULAS J0034-0052 \\citep{warren07} and CFBDS J005910.90-011401.3 \\citep{delorme08}. T dwarf spectra are dominated by strong water vapour and methane bands, and generally appear very blue in the near infrared (\\citealt{geballe03}; \\citealt{burgasser06}).  The physics of ultracool atmospheres is complex, and very difficult to accurately model. Atmospheric dust formation is particularly challenging for theory (\\citealt{allard01}; \\citealt{burrows06}), and there are a variety of other important issues that are not well understood, including the completeness of CH${_4}$/H${_2}$O molecular opacities, their dependence on $T_{\\rm eff}$, gravity and metallicity (e.g. \\citealt{jones05}; \\citealt{burgasser06a}; \\citealt{liu07}), as well as the possible presence of vertical mixing in such atmospheres \\citep{saumon07}. The emergent spectra from ultracool atmospheres are strongly affected by factors such as gravity and metallicity (e.g. \\citealt{knapp04}; \\citealt{burgasser06a}; \\citealt{metchev06}), and we must improve our understanding of such effects if we are to be able to spectroscopically constrain atmospheric and physical properties (such as mass, age and composition) of low-mass/sub-stellar field populations.  Discovering UCDs whose properties can be inferred indirectly (without the need for atmospheric models) is an excellent way to provide a test-bed for theory, and observationally pin down how physical properties affect spectra. We refer to such UCDs as {\\it benchmark} objects (e.g. \\citealt{pinfield06}). A population of benchmark UCDs with a broad range of atmospheric properties will be invaluable in the task of determining the full extent of spectral sensitivity to variations in UCD physical properties. However, such benchmarks are not common, and the constraints on their properties are not always particularly strong.  One variety of benchmark UCD that could yield accurate ages and surface gravities over a broad age range are in binary systems with a white dwarf (WD) companion. In particular, if the WD is relatively high mass, then the main sequence progenitor star will have a short lifetime, and the age of the binary system (and the UCD) will essentially be the same as the cooling age of the WD, which can be well determined from theory.  There have been several searches to find UCD companions to WDs. Despite this, only a small number of detached UCD--WD binaries have been identified; GD 165B(L4) \\citep{zuckerman92}, GD 1400(L6/7) (\\citealt{farihi04}; \\citealt{dobbie05}), WD\\,$0137-349$(L8) (\\citealt{maxted06}; \\citealt{burleigh06}) and PG1234+482 (L0) (\\citealt{steele07}; \\citealt{mullally07}).  The two components in GD\\,165 are separated by 120\\,AU, the separation of the components in GD\\,1400 and PG\\,$1234+482$ are currently unknown and WD\\,$0137-349$ is a close binary (semi-major axis $a=0.65$\\,R$_\\odot$). \\citet{farihi05} and \\citet{farihi06} also identified three late M companions to white dwarfs; WD$2151-015$ (M8 at 23\\,AU), WD$2351-335$ (M8 at 2054\\,AU) and WD$1241-010$ (M9 at 284\\,AU). The widest system previously known was an M8.5 dwarf in a triple system -- a wide companion to the M4/WD binary LHS 4039 and LHS 4040 \\citep{scholz04}, with a separation of 2200AU. There are several other known UCD--WD binaries, however these are cataclysmic variables (e.g. SDSS 1035; \\citealt{littlefair06}, SDSS1212; \\citealt{burleigh06b}, \\citealt{farihi08}, EF Eri; \\citealt{howell01}) but these are unlikely to provide the type of information that will be useful as benchmarks, as they have either evolved to low masses via mass transfer or their ages cannot be determined because of previous evolutionary phases (e.g. common envelope). WD0137-349 is also not likely to be useful as a benchmark due to the large uncertainty in age caused by the common envelope stage and the heating from the white dwarf.  Recent analysis from \\citet{farihi08b} shows that the fraction of brown dwarf companions (L and T) at separations within a few hundred AU of white dwarfs is $<$0.6 per cent.  However, despite this UCDs in wide binary systems are not uncommon (revealed through common proper motion) around main sequence stars at separations of $1000-5000$\\,AU (\\citealt{gizis01}; \\citealt{pinfield06}). However, when a star sheds its envelope post-main sequence, we may expect a UCD companion to migrate outwards to even wider separation (\\citealt{jeans1924}; \\citealt{burleigh02}), and UCD--WD binaries could thus have separations up to a few tens of thousands of AU. Although some of the widest binaries may be dynamically broken apart quite rapidly by gravitational interactions with neighbouring stars, some systems may survive, offering a significant repository of benchmark UCDs.  We present here the first results from our search of 2MASS and the digitized Schmidt plate archive SuperCOSMOS, for wide UCD--WD binaries in the southern sky, and present our first discovery of a wide UCD--WD binary confirmed through common proper motion and spectroscopy. In Section 2 we describe the techniques we have developed to select candidate UCD--WD binary systems using database photometry and proper motion. Section 3 describes our second epoch imaging of candidates, and followup infrared and optical spectroscopy along with proper motion analysis. We present spectroscopic analysis in Section 4. Section 5 discusses the newly discovered system, and Section 6 presents a discussion of future work. ",
        "conclusions": "\\subsection{A randomly aligned pair?}  In order to determine if the new system is a bonafide UCD-WD binary, we have statistically assessed the likelihood that two such objects could be a line-of-sight association with photometry and proper motion consistent with binarity by random chance. To do this we began with the UCD luminosity function of Cruz et al. (2007), which gives a number density of (4.9\\,$\\pm$0.6)$\\times 10^{-3}$ UCDs per $\\rm {pc^{-3}}$. We then calculated a volume associated with 1532 circular areas on the sky (one for each of our WD candidate sample), with radii of 89 arcsec (separation of the components) and a line-of-sight depth of 58\\,$\\pm$17pc (approximate distance to the new M9 UCD, using relations from \\citealt{dahn02}).  This volume equates to 58\\,$\\pm$26 $\\rm {pc^{-3}}$, giving a total expected number of 0.28 UCDs to be within 89 arcsec of one of our WD candidates.  To factor in the probability that two objects might have a common proper motion at the level as our measurements, we downloaded a magnitude-limited sample ($R$\\,$<$20) from the SuperCOSMOS Science Archive, applying the same minimum proper motion requirement that was used to create our WD candidate sample. This sample of 160 sources was centred on our WD and selected from a large circular sky area of radius 90 arcmin. We then constructed a proper motion vector-point-diagram, and counted sources that were found to be within the 2\\,$\\sigma$ uncertainty circle of the measured UCD proper motion. We found that four of the 160 sources had proper motion consistent with our UCD, suggesting a probability of 2.4\\,$\\pm$1.2 per cent that such common proper motion could occur by random chance (where we assume Poisson uncertainties associated with this and other samples considered in this discussion).  An additional requirement fulfilled by our UCD-WD system is that the colour-magnitude information must be consistent with a common distance (see Fig.~\\ref{ld_cmd}), and we found that 53 per cent ($786\\,\\pm28$) of our WD candidate sample were photometrically consistent with being at the same distance as the UCD.  Finally we consider what fraction of our WD candidates might be spurious, and thus not able to contribute to non binary line-of-sight associations where the WD has been confirmed spectroscopically. In the magnitude range $R$\\,$<14$, where MS99 is thought to be essentially complete, we find that 66 per cent of our WD candidates are included in the MS99 catalogue. This suggests that, at least for brighter magnitudes, our WD candidates are relatively free from contaminating objects, and that our selection techniques are robust. While we cannot be sure that the same low-level of contamination applies to the full magnitude range, we take the conservative approach and consider the full WD candidate sample as potentially contributing to non-binaries that appear to be UCD-WD pairs.  Taking into account all these factors, we estimate that we would expect $0.0036\\,\\pm0.0025$ randomly aligned UCD-WD pairs with $\\le$\\,89 arcsec separation and proper motion and photometry consistent with binarity at the level of our observations. The likelihood of the system being merely a line-of-sight association is thus vanishingly small, and we can assume that the UCD-WD pair is a gravitationally bound binary system.   \\subsection{Binary age}  The age of the binary system can be constrained from the white dwarf mass and cooling age. For our corrected fit, we suggest that 2MASSJ$0030-3740$ is 0.48\\,M$_\\odot$ and has a cooling age of 0.94\\,Gyr. We access the IFMR determinations of \\citet{weidemann00}, \\citet{dobbie04}, \\citet{dobbie06}, \\citet{ferrario05}, \\citet{catalan08} and \\citet{kalirai08} to estimate a likely, initial-mass constraint for the main-sequence progenitor star of $1-2$\\,M$_\\odot$.  The main sequence lifetime of a progenitor of this mass is likely $>$1Gyrs (probably several Gyrs) and thus is not useful when trying to constrain the upper age limit of the system. Thus the upper age limit of our system remains uncertain. Note that if the WD mass were higher, for example if the helium enrichment of the atmosphere is lower than typical (see Section 4.2) then the WD's log($g$) could be as high as 8.1\\,dex, with a WD mass of $\\sim$0.65\\,M$_\\odot$ and cooling age of $\\sim$1.5\\,Gyrs. This would allow progenitor mass to be constrained to a likely range of $<2.7$\\,M$_\\odot$, giving a main-sequence progenitor lifetime of $<0.83$\\,Gyrs \\citep{monteiro06}, yielding a binary age constraint of $1.5-2.3$\\,Gyrs. This possibility is instructive at least, in demonstrating the level of age constraints (with accompanying UCD constraints) that may be placed on benchmark binaries of this type. However, it is not possible to judge the helium content of the WD's atmosphere (if any), and we can only confidently place a lower limit on the age of this binary from our best fit cooling age for the WD and the likely length of the main sequence lifetime of the progenitor, which is likely equal to or larger than the white dwarf cooling age. The age of the binary is thus $>1.94$\\,Gyrs.   \\subsection{UCD properties}  We have estimated $T_{\\rm eff}$, mass and log($g$) from the Lyon group dusty models (\\citealt{chabrier001}; \\citealt{baraffe02}), using the minimum age of the system (0.94\\,Gyrs) and our estimated $M_J$ of 2MASSJ$0030-3739$. The models indicate that 2MASSJ$0030-3739$ has $T_{\\rm eff}=2000-2400$\\,K, mass=$0.07-0.08$\\,M$_{\\odot}$ and log($g$)=$5.30-5.35$, placing it close to the limit for hydrogen burning. Note that our $T_{\\rm eff}$ is consistent with the semi-empirical estimates of \\citet{golimowski04} for an M9\\,$\\pm$1 dwarf, which use well measured luminosities and a model constraint on radius (which changes by $<$10 per cent for ages of $1-5$\\,Gyrs), to determine $T_{\\rm eff}$ values spanning a wide range of spectral type. The full list of properties for the binary are listed in Table~\\ref{properties}.    \\begin{table} \\caption{Parameters of the binary and it's components.} \\centering \\begin{tabular}{|l|l|c|} \\hline Parameter &  & Value \\\\ \\hline Separation on sky & ............ & 89 arcsec\\\\ Estimated distance & ............ & $41-59$\\,pc\\\\ Estimated line-of-sight \\\\ separation & ............ & $3650-5250$\\,AU\\\\ Minimum age of system & ............& $>1.94$ \\,Gyrs\\\\ & & \\\\ \\underline{Ultracool Dwarf} &  &  \\\\ RA & ............ & 00 30 06.26\\\\ DEC & ............ & -37 39 48.2\\\\ 2MASS designation & ............ & 2MASSJ$0030-3739$\\\\ Distance & ............ & 41 -- 75pc\\\\ 2MASS $J$  & ............ & 15.2\\,$\\pm$ 0.05\\\\ 2MASS $H$  & ............ & 14.4\\,$\\pm$ 0.05\\\\ 2MASS $K_s$  & ............ & 13.8\\,$\\pm$ 0.06\\\\ DENNIS $I$  & ............ & 18.4\\,$\\pm$ 0.23\\\\ DENNIS $J$  & ............ & 15.06\\,$\\pm$ 0.14\\\\ SuperCOSMOS $I$  & ............ & $\\sim$18.3\\\\ $\\mu$ RA  & ............ & -130\\,$\\pm$30\\,mas yr$^{-1}$\\\\ $\\mu$ DEC  & ............ & -70\\,$\\pm$20\\,mas yr$^{-1}$\\\\ Spectral Type  & ............ & M9\\,$\\pm$1\\\\ Mass & ............ & $0.07-0.08$\\,M$_\\odot$\\\\ $T_{\\rm eff}$ & ............ & $2000-2400$\\,K\\\\ log($g$) & ............  & $5.30-5.35$\\,dex \\\\ & & \\\\ \\underline{White Dwarf} &  &  \\\\ RA & ............ & 00 30 11.9\\\\ DEC & ............ & -37 40 47.2\\\\ 2MASS designation & ............ & 2MASSJ$0030-3740$\\\\ Distance & ............ & $27-59$\\,pc\\\\ 2MASS $J$  & ............ & 16.1\\,$\\pm$ 0.11\\\\ 2MASS $H$  & ............ & 15.8\\,$\\pm$ 0.15\\\\ DENNIS $I$  & ............ & 16.2\\,$\\pm$ 0.07\\\\ DENNIS $J$  & ............ & 15.9\\,$\\pm$ 0.22\\\\ SuperCOSMOS $B$  & ............ & $\\sim$16.77\\\\ SuperCOSMOS $R$  & ............ & $\\sim$16.35\\\\ SuperCOSMOS $I$  & ............ & $\\sim$15.97\\\\ $\\mu$ RA  & ............ & -83\\,$\\pm$30\\,mas yr$^{-1}$\\\\ $\\mu$ DEC  & ............ & -70\\,$\\pm$12\\,mas yr$^{-1}$\\\\ Spectral Type & ............& DA\\\\ $T_{\\rm eff}$  & ............ & 7600$\\pm$175\\,K\\\\ log($g$)  & ............ & $7.79-8.09$\\,dex\\\\ Mass  & ............ & $0.48-0.65$\\,M$_{\\odot}$\\\\ WD cooling age & ............ & $0.94 - 1.5$\\,Gyrs\\\\ WD progenitor age & ........... & $>1$\\,Gyr\\\\   \\hline \\end{tabular} \\label{properties} \\end{table}"
    },
    "0806/0806.3988.txt": {
        "abstract": "We report on the discovery of three new dwarf galaxies in the Local Group. These galaxies are found in new CFHT/MegaPrime $g,i$ imaging of the south-western quadrant of M31, extending our extant survey area to include the majority of the southern hemisphere of M31's halo out to 150\\,kpc. All these galaxies have stellar populations which appear typical of dwarf spheroidal (dSph) systems. The first of these galaxies, Andromeda~XVIII, is the most distant Local Group dwarf discovered in recent years, at $\\sim 1.4$\\,Mpc from the Milky Way ($\\sim 600$\\,kpc from M31). The second galaxy, Andromeda~XIX, a satellite of M31, is the most extended dwarf galaxy known in the Local Group, with a half-light radius of $r_h \\sim 1.7$\\,kpc. This is approximately an order of magnitude larger than the typical half-light radius of many Milky Way dSphs, and reinforces the difference in scale sizes seen between the Milky Way and M31 dSphs (such that the M31 dwarfs are generally more extended than their Milky Way counterparts). The third galaxy, Andromeda~XX, is one of the faintest galaxies so far discovered in the vicinity of M31, with an absolute magnitude of order $M_V \\sim -6.3$. Andromeda~XVIII, XIX and XX highlight different aspects of, and raise important questions regarding, the formation and evolution of galaxies at the extreme faint-end of the luminosity function. These findings indicate that we have not yet sampled the full parameter space occupied by dwarf galaxies, although this is an essential pre-requisite for successfully and consistently linking these systems to the predicted cosmological dark matter sub-structure. ",
        "introduction": "Edwin Hubble first coined the term ``Local Group'' in his 1936 book ``The Realm of the Nebulae'', to describe those galaxies that were isolated in the general field but were in the vicinity of the Galaxy. In recent years, the galaxies of the Local Group have been at the focus of intense and broad-ranging research, from providing laboratories for the investigation of dark matter properties (e.g., \\citealt{gilmore2007} and references therein) to determinations of the star formation history of the Universe (e.g., \\citealt{skillman2005} and references therein). Understanding individual galaxies in the Local Group offers important contributions to galaxy structure and evolution studies; understanding the properties of the population is central to galaxy formation in a cosmological context.  Hubble originally identified nine members of the Local Group: the Galaxy and the Large and Small Magellanic Clouds; M31, M32 and NGC205; M33, NGC6822 and IC1613; along with three possible members NGC6946, IC10 and IC342. The distances of the latter three were highly uncertain due to heavy extinction; IC10 has since been confirmed as a member (\\citealt{sakai1999}) although the other two lie outside the Local Group (NGC6946; \\citealt{sharina1997}; IC342: \\citealt{krismer1995}).  The discovery of new Local Group members continued at a relatively constant rate up to the start of 2004 (e.g., \\citealt{ibata1994,whiting1997,whiting1999, armandroff1998,armandroff1999,karachentsev1999}), at which point the discovery rate has increased sharply. This has mostly been due to large area photometric CCD-based surveys of the Milky Way and M31 stellar haloes: by searching for overdensities of resolved stars in certain regions of colour-magnitude space, it is possible to identify very faint dwarf satellites which have previously eluded detection.  Around the Milky Way, this technique has so far lead to the discovery of 9 new satellites since 2005 (including possible diffuse star clusters) (\\citealt{willman2005,willman2006,belokurov2006,belokurov2007,zucker2006b,walsh2007}). All of these discoveries have been made using the Sloan Digitized Sky Survey (SDSS). In addition, two new isolated dwarf galaxies have been identified: Leo T, more than 400kpc from the Milky Way (\\citealt{irwin2007}), was discovered in the SDSS, and a revised distance estimate for the previously known UGC4879 has moved this galaxy from $> 10$\\,Mpc to being placed on the periphery of the Local Group (a scant $\\sim 1.1$\\,Mpc from the Milky Way; \\citealt{kopylov2008}).  Around M31, 9 new dwarf galaxy satellites have been discovered since 2004 (not including results presented herein). Two of these galaxies (Andromeda IX, X) were found in special SDSS scans of M31 (\\citealt{zucker2004a,zucker2007}) and one (Andromeda~XIV) was discovered serendipitously by \\cite{majewski2007} in Kitt Peak 4m imaging of fields in the south-east halo of M31. The remaining new dwarf galaxies have been discovered as part of our ongoing photometric survey of this galaxy and its environs using the INT/WFC (Andromeda~XVII, \\citealt{irwin2008}) and CFHT/MegaPrime (Andromeda~XI, XII and XIII, \\citealt{martin2006}; Andromeda XV and XVI, \\citealt{ibata2007}). Despite its name, Andromeda~XVII is only the fifteenth dwarf spheroidal satellite of M31 to be discovered; Andromeda~IV is a background galaxy (\\citealt{ferguson2000}) and Andromeda~VIII was originally identified using planetary nebulae (\\citealt{morrison2003}) which were later shown to belong to M31 and not to a separate entity (\\citealt{merrett2006}). Additionally, only thirteen of these dwarfs are actually located in the constellation of Andromeda (Andromeda~VI $\\equiv$ the Pegasus dSph; Andromeda~VII $\\equiv$ the Cassiopeia dSph).  The unique, panoramic, perspective of the resolved stellar populations of galaxies provided by Local Group members make them ideal targets for observational programs aimed at understanding the detailed structure of galaxies, their formation processes and their evolutionary pathways. Dwarf galaxies are of particular interest, given that they are thought to be the lowest mass, most dark matter dominated systems which contain baryons (e.g., \\citealt{mateo1998a}). They are therefore particularly sensitive probes of external processes, such as tides and ram pressure stripping (e.g., \\citealt{mayer2006,mcconnachie2007c,penarrubia2008b}), and internal processes such as feedback from star formation (e.g., \\citealt{dekel1986,dekel2003}). Further, their potential as probes of dark matter (e.g., \\citealt{gilmore2007,strigari2007a}) and their probable connection to cosmological sub-structures (e.g., \\citealt{moore1999,bullock2000,kravtsov2004,penarrubia2008a}) give them an importance to galaxy formation not at all in proportion to their luminosity.  Here we report on the discovery of three new dwarf galaxies in the Local Group, all of which have been found as part of our ongoing CFHT/MegaPrime photometric survey of M31. This new imaging extends our survey area from the south-eastern quadrant discussed in \\cite{ibata2007} to the west, and currently includes an additional $49$\\,sq.\\,degrees of M31's halo out to a maximum projected radius of 150\\,kpc. Section~2 summarises the observations and data-reduction procedures and Section~3 presents a preliminary analysis of the new dwarfs and quantifies their global properties. In Section~4, we discuss our results in relation to some of the key questions which have been prompted with the discoveries of so many new low luminosity galaxies in the Local Group. Section~5 summarises our results. ",
        "conclusions": "Andromeda~XVIII, XIX and XX have a range of relatively unusual properties. In particular, Andromeda~XVIII is one of the most distant Local Group galaxies discovered for several years, and is one of the most isolated systems in the Local Group. Andromeda~XIX is extremely extended, with a very large half-light radius and extremely faint central surface brightness. Andromeda~XX, on the other hand, is one of the lowest luminosity dwarf galaxies so far discovered around M31, with a magnitude of $M_V \\simeq -6.3^{+1.0}_{-0.7}$, comparable to the luminosity of Andromeda~XII ($M_V = -6.4 \\pm 1.0$; \\citealt{martin2006}). In this section, we discuss the properties of these galaxies in the larger context of the main science questions raised by the recent discoveries of so many new dwarf galaxies.  \\subsection{Completeness}  Prior to 2004, there were 15 dSph galaxies known in the Local Group (nine Milky Way satellites, six M31 satellites and two isolated systems, Cetus and Tucana). Since this time, 22 new dwarf galaxies (including possible diffuse star clusters around the Milky Way) have been discovered in the Local Group, the overwhelming majority of which are dSph satellites of the Milky Way and M31. For the Milky Way, the SDSS has been responsible for all the discoveries to date, and most of the galaxies discovered have been extremely faint; no new Milky Way satellites with $M_V \\lesssim -8$ have been found. Thus, apart from satellites hidden by the Milky Way disk, our satellite system is probably complete to this approximate magnitude limit, as originally argued by \\cite{irwin1994}.  Around M31, it is more difficult to identify extremely faint dwarf galaxies since we cannot probe as far down the stellar luminosity function. Andromeda~XII and Andromeda~XX are the two faintest M31 satellites found so far, both with $M_V \\sim -6.3$. For comparison, the faintest Milky Way satellite found to date is probably Willman I, with $M_V \\sim -2.7$ (\\citealt{willman2006,martin2008}).  Andromeda~XVIII is considerably brighter than Andromeda~XX, and has a central surface brightness similar to or brighter than Andromeda~V ($S_o = 25.6 \\pm 0.3$\\,mags\\,arcsec$^{-2}$). Andromeda~XVIII is clearly visible in the POSSII/UKSTU (Blue) survey image which we retrieved through the Digitized Sky Survey and which is reproduced in the lower panel of Figure~5. However, its identification is made more complicated by numerous nearby bright stars and nebulosity in its vicinity, which may act to explain why it was not discovered using these data. We have also confirmed that it is visible in the original POSSI (Blue) survey. Its belated discovery indicates that previous surveys for relatively {\\it bright} dwarf galaxies around M31 were incomplete and that some dwarfs were missed. Variable and unknown completeness is problematic for studies of satellite distributions and highlights the vital need for more systematic studies such as those now being conducted.  It is fortuitous that Andromeda~XVIII lies within our survey area given its considerable distance from M31. Indeed, even as current and future surveys help improve the completeness of the M31 and Milky Way satellite systems, many isolated Local Group galaxies can be expected to continue to elude detection: unlike the Milky Way satellites, they are not nearby, and unlike the M31 satellites, they are not necessarily clustered in an area amenable to systematic searches. PanStarrs $3\\pi$ will survey a large fraction of the sky a magnitude deeper than SDSS, and should discover isolated Local Group galaxies, particularly those within 500\\,kpc or so from the Milky Way. However, very faint galaxies much further away than this ($\\sim 1$\\,Mpc) may prove more difficult to spot. Exactly how many very faint dwarf galaxies are to be found at the periphery of the Local Group is likely to remain uncertain for some time yet.  \\subsection{Spatial distribution}  Several recent studies of the spatial distributions of satellites around the Milky Way and M31 (\\citealt{willman2004,kroupa2005,mcconnachie2006a,koch2006,metz2007,irwin2008}) have generally concluded that the distributions appear anisotropic: \\cite{mcconnachie2006a} highlight the fact that (at the time) 14 out of the 16 candidate satellites of M31 are probably on the near side of M31, while others (\\citealt{kroupa2005,koch2006,metz2007,irwin2008}) conclude that many of the Milky Way and M31 satellites are aligned in very flattened, disk-like, distributions (an observation originally made by \\citealt{lyndenbell1976,lyndenbell1982}).  Andromeda~XVIII, XIX and XX do not lie near any of the principle satellite planes previously proposed to exist around M31. As discussed in the previous sub-section, the census of Local Group galaxies is clearly not complete, and it is too early to draw definitive conclusions regarding the distributions of satellites. This is particularly true around M31, where relatively bright satellites are still being discovered. For the Milky Way, the SDSS covers roughly one-fifth of the Milky Way halo in the direction of the north galactic cap; depending upon how many satellites are found in future surveys at lower latitudes, the statistical significance of the proposed streams of satellites may change substantially.  In terms of spatial distributions, Andromeda~XVIII is unusual insofar as it is very distant - roughly 1.4\\,Mpc from the Milky Way, and roughly 600\\,kpc from M31. Thus it is probably not a satellite of M31, although kinematics may help reveal whether it is approaching M31 and the Local Group for the first time (like Andromeda~XII, \\citealt{chapman2007}) or if it has been thrown out from M31 following an interaction (like Andromeda~XIV, \\citealt{majewski2007,sales2007}).  \\subsection{Environment and structures}  \\subsubsection{Andromeda~XVIII, position and morphology}  Andromeda~XVIII appears to possess stellar populations typical of dSph galaxies. If it is subsequently confirmed to be gas poor, then it will be the third dSph galaxy found in isolation in the Local Group (in addition to Cetus and Tucana). The fact that isolated galaxies are preferentially more gas-rich compared to satellites (\\citealt{einasto1974}) has lead to the proposition that satellite galaxies are stripped of their gas via ram-pressure stripping and tidal harassment in the halo of the host galaxy (e.g., \\citealt{mayer2006}). However, for isolated systems such as Andromeda~XVIII, Cetus and Tucana, prolonged interactions with massive galaxies are unlikely to have occurred. Likewise, the gas-deficient satellite Andromeda~XII is not believed to have undergone any past interactions with a large galaxy since it appears to be on its first infall into the potential of M31 (\\citealt{chapman2007}). Further, the most compelling case of a dwarf galaxy thought to be undergoing ram-pressure stripping is Pegasus (DDO216; \\citealt{mcconnachie2007c}), an {\\it isolated} galaxy more than 400\\,kpc from M31. Clearly, understanding if these observations are consistent with the present models for dwarf galaxy evolution requires a more complete inventory of nearby galaxies and their properties than we currently possess.  \\subsubsection{Andromeda~XIX, tides and substructure}  The half-light radius of Andromeda~XIX is $6.2$\\,arcmins. At the distance we derive for it, this corresponds to $r_h \\simeq 1.7$\\,kpc, which is the largest value yet recorded for any dSph in the Local Group. The average half-light radius for Milky Way dSphs is an order of magnitude less, at $r_h \\sim 150$\\,pc, and none have half-light radii larger than $r_h \\simeq 550$\\,pc (with the exception of the tidally disrupting Sagittarius dSph; \\citealt{majewski2003}) . M31 dSphs, on the other hand, have typical half-light radii of $r_h \\sim 300$\\,pc, with the previous extremes being Andromeda~II, with $r_h \\simeq 1.1$\\,kpc, and Andromeda~VII, with $r_h \\simeq 750$\\,pc (\\citealt{mcconnachie2006b}).  The extremely diffuse and extended nature of Andromeda~XIX is reminiscent of the ``outer component'' of Andromeda~II, as traced by horizontal branch stars by \\cite{mcconnachie2007a}.  It is tempting to attribute the diffuse structure of Andromeda~XIX to tidal interactions. In this respect, it is relevant to note that Andromeda~XIX lies very close to the major axis substructure identified by \\cite{ibata2007}. No independent distance estimate to this substructure currently exists; \\cite{ibata2007} assumed it to be at the distance of M31 but if it is at the same distance as Andromeda~XIX then the photometric metallicity estimates of these features will be very similar. Figure~7 shows the surroundings of Andromeda~XIX as a stellar density map; the first two contour levels are $2$ and $3 - \\sigma$ above the background, and the levels then increase by $1.5\\,\\sigma$ over the previous level. As well as showing Andromeda~XIX as a prominent overdensity, there is some evidence of stellar material in its outskirts (also visible in the contours of Figure~3). Whether or not Andromeda~XIX is the source of the major axis substructure identified in \\cite{ibata2007}, or is being tidally perturbed, will require detailed kinematics in this region. We note that \\cite{penarrubia2008b} show that the effect of tides on dwarf galaxies in cosmological haloes is to decrease the central surface brightness and {\\it decrease} the half light radius of the bound component. This would argue against tidal effects explaining the structure of Andromeda~XIX.  The large scale-size of Andromeda~XIX reinforces the difference in scale-size between the Milky Way and M31 satellites first highlighted in \\cite{mcconnachie2006b}, such that the M31 dSphs are more extended than their Milky Way counterparts. \\cite{penarrubia2008a,penarrubia2008b} have investigated the cause of this disparity in an attempt to relate it to either differences in the underlying dark matter properties of the dwarfs or differences in their evolution around their hosts. They conclude that tidal effects are insufficient to explain the magnitude of the effect. However, if the different scale sizes reflect intrinsic differences between the Milky Way and M31 sub-haloes then this should reveal itself in the kinematics of the two populations (with the M31 dwarfs being dynamically hotter than their Milky Way counterparts). Whatever the cause, the comparison of Andromeda~XIX and the other M31 satellites to the Milky Way population highlights the importance of sampling dwarfs in a range of environments so as to obtain a fuller appreciation of the range of properties that these systems possess. In turn, this helps us understand the physical drivers behind the differences and similarities we observe. We note that studies of the star clusters of M31 (\\citealt{huxor2005,huxor2008}) have already extended the known parameter space for these objects, with the M31 population containing extended star clusters not found in the Milky Way population.  \\subsection{Satellites that are missing and ``the missing satellites''}  Andromeda~XX is an exceptionally faint galaxy with a very poorly populated RGB. This makes an accurate derivation of its properties particularly difficult. However, the star formation history of Andromeda~XX and the other ultra-faint satellites is particularly relevant to the ``missing satellites'' question (do all the thousands of dark matter sub-haloes predicted to exist in the haloes of galaxies like the Milky Way and M31 contain stars and, if they do, where are they?).  Until recently, only a dozen or so dwarf satellites were observed, and it was noted that the cumulative mass distribution of these satellites was dramatically different to that of predicted dark matter sub-haloes, even at relatively large masses (\\citealt{moore1999,klypin1999}). To solve this discrepancy without altering the underlying cosmology, it was suggested that either there were a large number of luminous satellites awaiting discovery or that not all sub-haloes have a luminous component.  Despite many new galaxies in the Local Group being discovered, and many more undoubtedly awaiting discovery, we consider it very unlikely that these discoveries will resolve the discrepancy between theory and observation. The original comparison between the observed and predicted satellite mass functions shows that the discrepancy sets in for dwarfs as luminous as the Small Magellanic Cloud ($M_V \\simeq -16$) and Fornax ($M_V \\simeq -13 $). Finding thousands of very faint (and presumably less massive?) satellites would not solve the disagreement at the more massive end and there is no evidence to suggest that a dozen galaxies the luminosity of Fornax have been missed (e.g., \\citealt{irwin1994}). Further, as higher resolution dark matter simulations make clear (e.g., \\citealt{diemand2007}), the sub-halo mass function appears to continue to increase at the low mass end. It seems reasonable, therefore, that at some point these haloes will not be massive enough to be able to accrete and/or retain baryons and form stars, and this implies that there is a minimum mass halo which can host a luminous component (\\citealt{kravtsov2004}).  A re-analysis of the observed dynamics of the dwarf galaxies by \\cite{penarrubia2008a,penarrubia2008b} within the $\\Lambda-$CDM \u00e7framework has shown that few-if-any of these galaxies (including recent discoveries) occupy a halo with a circular velocity less than $\\sim 10 - 20$\\,km\\,s$^{-1}$. Further, these estimates bring the cumulative distribution of luminous satellites and dark matter sub-haloes into good agreement at the high-mass end. Using a different technique, \\cite{strigari2007b} find a similar result. Given that these authors find good agreement between observations and theory down to a certain mass limit, their results support the idea of a mass threshold in dark matter haloes below which star formation becomes highly inefficient. Therefore, by continuing to identify new, ultra-faint dwarfs, we probe the astrophysics of galaxy formation at low mass limits where the sensitivity to complex feedback mechanisms - such as star formation (\\citealt{kravtsov2004}) and reionization (\\citealt{bullock2001}) - is greatest."
    },
    "0806/0806.3220_arXiv.txt": {
        "abstract": "{The Pierre Auger Collaboration has reported 27 Ultra-High Energy Cosmic Ray Events (UHECRs) with energies above 56$\\times 10^{18}$ eV (56~EeV) and well determined arrival directions as of 2007 August 31. They find that the arrival directions of these UHECRs are not isotropic, but instead appear correlated with the positions of nearby Active Galactic Nuclei (AGNs) from the catalog of V{\\'e}ron-Cetty \\& V{\\'e}ron. } {Our aim was to determine the sources of these UHECRs by comparing their arrival directions with more complete and/or comprehensive astronomical source catalogs. } {We have cross correlated the arrival directions of the UHECRs with the positions of supernovae, radio supernovae, galaxies, active galaxies, radiogalaxies, and clusters of galaxies, all at distances within $\\sim$200~Mpc. } {Four (eight) of the 27 UHECRs with energy greater than 56~EeV detected by the Pierre Auger Observatory have arrival directions within 1.5\\arcdeg\\ (3.5\\arcdeg) of the extended ($\\ge$ 180~kpc) radio structures of nearby radiogalaxies or the single nearby BL~Lac with extended radio structure. Conversely the radio structures of three (six) of all ten nearest extended radiogalaxies are within 1.5\\arcdeg\\ (3.5\\arcdeg) of a UHECR; three of the remaining four radiogalaxies are in directions with lower exposure times. This correlation between nearby extended radiogalaxies and a subset of UHECRs is significant at the 99.9\\% level. For the remaining $\\sim$20 UHECRs, an isotropic distribution cannot be ruled out at high significance. The correlation found by the Auger Collaboration between the 27 UHECRs and AGNs in the V{\\'e}ron-Cetty \\& V{\\'e}ron catalog at D$\\lesssim$71~Mpc has a much lower significance when one considers only the $\\sim$20 UHECRs not `matched' to nearby extended radiogalaxies. No clear correlation is seen between UHECRs and supernovae, supernova remnants, nearby galaxies, or nearby groups and clusters of galaxies. } {Nearby extended radiogalaxies are the most likely source of at least some UHECRs detected by the Pierre Auger Observatory. The remaining UHECRs are not inconsistent with an isotropic distribution; their correlation to nearby AGNs is much less significant than earlier estimated. This is the first direct observational proof that radio galaxies are a significant source of UHECRs. The primary difference between the UHECR detections at the Pierre Auger Observatory and previous experiments, e.g. AGASA, may thus primarily be that the Southern Hemisphere is more privileged with respect to nearby radiogalaxies with highly extended radio jets and lobes. } ",
        "introduction": "\\label{secintro}  Ultra-High Energy Cosmic Rays (UHECRs) are protons or light nuclei with energies greater than about 10$^{19}$ eV (10 EeV). When these highly energetic particles enter the earth's atmosphere they produce a shower of secondary particles and excite the Nitrogen molecules in the atmosphere; both effects can be detected from the ground. Recent advances in ground shower detectors allow a precise measurement of both the initial energy and the arrival direction of the UHECR above the earths atmosphere, e.g. AGASA \\citep{taket99} and the Pierre Auger Observatory \\citep{abret04,demet07}. While UHECRs are rare they have a unique astronomical potential: UHECRs with energies $\\gtrsim$50~EeV are not expected to be deflected significantly by Galactic or intergalactic magnetic fields. Their measured arrival directions can therefore be traced directly back to the originating source within the observatory measurement uncertainties.  Suspected sources of high energy cosmic rays include core collapse Supernovae (SN II), Supernova Remnants, Pulsars, Active Galaxies - especially radio jets and radio lobes, and gamma rays bursts (for reviews and lectures see \\citet{kat08,hil99} and references therein). In these sources, the cosmic rays are posited to undergo successive accelerations via scattering off energetic charged particles and/or in shocks. More exotic explanations include an origin in dark matter anhilation.  The main energy loss for UHECR propagating over cosmological distances is expected to be pion-production, the so called Greisen Zatsepin Kuzmin  (GZK) effect \\citep{gre66,zatkuz66}. In this process, a UHECR interacts with a CMB photon and loses an estimated $\\sim$30\\% of its energy. Both the energy loss and the mean free path between interactions depend on several details; energy loss length predictions vary between 20 and 100~Mpc (see e.g. \\citet{sta04},\\citet{kat08},PA08). If the GZK effect is at work then UHECRs with energy above $\\sim$50~EeV are expected to originate primarily in sources closer than this predicted energy loss length.  The Pierre Auger Observatory (hereafter PAO) in Argentina \\citep{abret04,demet07}, an array of 1600 Cerenkov detectors spread over 3000 km$^2$ plus six optical telescopes, is designed to measure arrival directions and energies of cosmic rays via their secondary particle showers and their atmospheric flouresence. Latest details on the instrumentation and exciting science results being produced by the Auger Observatory and Collaboration can be found at the Auger web-site http://www.auger.org. The array has been in partial operation since January 2004. The Auger collaboration has reported 81 events between 2004 January 1 and 2007 August 31 with reconstructed energies above 40~EeV and zenith angles smaller than 60\\arcdeg; of these, 27 have energies above 56~EeV \\citep[][hereafter PA07,PA08]{abret07,abret08}. The origins of the latter 27 UHECRs are the focus of this paper.  The Auger Collaboration found that the arrival directions of the 27 UHECRs with energies above 56~EeV are not isotropic at a 99\\% significance level (PA07,PA08). Instead, they found that the arrival directions are correlated with the positions of AGNs within $\\sim$71~Mpc (PA07,PA08,\\citet{molet07}). For this analysis they used all AGNs (and galaxies with \\hii\\ nuclei) in the catalog of V{\\'e}ron-Cetty \\& V{\\'e}ron, 12th Edition \\citep{verver06}. In their statistical correlation of all 81 UHECRs at energy above 40~EeV with the V{\\'e}ron-Cetty \\& V{\\'e}ron catalog of AGNs, PA07,PA08 found that the highest correlation between UHECRs and AGNs was obtained for a maximum AGN distance of 71~Mpc, an angular separation of 3.2\\arcdeg\\ between ``matched'' AGNs and UHECRs, and an energy threshold for UHECRs of 56~EeV. This maximum distance, 71~Mpc, is in line with the expectations of the GZK effect. \\citet{goret07}, in an astro-ph comment, disagreed with this finding, pointing out that weighting each AGN with its distance would predict a very different distribution of UHECRs than that found by the PAO.  Our interest in this topic was provoked when a quick examination of the 27 UHECR events showed that the direction of at least five of the 27 UHECR events were directly in the line of sight to nearby extended radiogalaxies. In this article, we attempt a more comprehensive correlation between the arrival directions of the 27 UHECRs with energies above 56~EeV and catalogs of potential sources of UHECRs, both Galactic and Extragalactic. Sec. \\ref{secdata} describes the data used in our study, Sec. \\ref{secres} describes the principle results obtained, and Sec. \\ref{secdis} contains a brief discussion and conclusions of our study. Distances to galaxies are calculated using 72 $\\H0$, except for relatively nearby galaxies for which we use distances as referenced. ",
        "conclusions": "\\label{secdis}  The arrival directions a subset of the 27 UHECRs with energies above 56~EeV detected by the PAO are statistically most closely related to nearby radiogalaxies with extended radio jets and/or lobes. It is thus likely that nearby (due to the GZK effect) radiogalaxies with extended (radio extent $\\ge$ 180~kpc), two sided jets and lobes are the source population of a subset of the UHECRs detected by the PAO. Additionally, there is weak evidence that hosting a jet with morphology closer to FR~II makes a radiogalaxy more likely to be a UHECR source. Interestingly, all these factors are consistent with the lack of UHECR detections towards M~87.  The results of the previous section were based only on UHECRs detected by the PAO. We have also tested for correlations against the UHECRs with energy above 56~EeV detected by AGASA (red circles with radius 3.5\\arcdeg\\ in Fig.\\ref{figjetned}. Among the nearby (D $\\leq$ 75~Mpc) galaxies with extended radio jets, the only clear match is to CGCG~514-050, a radiogalaxy at D=72.2~Mpc with radio extent $\\sim$250~kpc. The remaining three nearby galaxies with extended radio galaxies, NGC~315, NGC~383, and NGC~5127, are not matched to within 3.5\\arcdeg\\ of any AGASA UHECR event, though there are several UHECRs at slightly larger distances.  Additionally, One AGASA UHECR is close to the location of NGC~7626 and another is centered on the radiogalaxy 3C~120, at D=135~Mpc and total radio extent $\\sim$315~kpc.  Among the UHECRs detected by PAO, an isotropic distribution of the $\\sim$20 UHECRs which are not in the line of sight to nearby extended radiogalaxies cannot be ruled out at high significance. The significance of the correlation found by PA07 and PA08 between UHECR arrival directions and nearby AGNs from the catalog of V{\\'e}ron-Cetty \\& V{\\'e}ron is much lower when one only considers the $\\sim$ 20 UHECRs not matched to nearby extended radiogalaxies and even lower when H~II nuclei in the the AGN Catalog are deleted. The findings of PA07,PA08 were probably largely influenced by the correlation between UHECRs and nearby radiogalaxies which we have shown above.  The UHECR arrival directions are not strongly correlated with either supernovae, extragalactic radio-supernovae, or nearby groups and clusters of galaxies.  The main difference between the results of the Pierre Auger Collaboration and those from previous studies of UHECRs, e.g. AGASA, could primarily be that the southern location of the PAO is more privileged with respect to nearby extended radiogalaxies."
    },
    "0806/0806.4894_arXiv.txt": {
        "abstract": "We have conducted a systematic search for multiperiodic pulsations in the RR~Lyrae-type stars of the galactic globular cluster $\\omega$~Cen. Secondary periodicities close to the primary pulsation frequency have been detected in 17 out of 70 studied fundamental mode (RRab) pulsators and in 35 out of 81 overtone (RRc) pulsators. Because of the observed period ratios, these newly detected periodicities must correspond to nonradial modes. Their beating with the primary radial pulsation leads to a slow amplitude and phase modulation, commonly referred to as the Blazhko effect. The incidence rate of Blazhko modulation in ${\\omega}$~Cen RRab stars ($24\\pm 5\\%$) is similar to that observed in the Galactic Bulge. In the case of $\\omega$~Cen RRc stars, the incidence rate of Blazhko effect is exceptionally high ($38\\pm5\\%$), more than 3 times higher than in any other studied population.  In addition to Blazhko variables, we have also identified two RR~Lyr variables exhibiting first overtone/second overtone double-mode pulsations, and a triple-mode High Amplitude $\\delta$ Scuti variable. ",
        "introduction": "Omega Centauri ($\\omega$ Cen) is the largest globular cluster of the Galaxy. It contains about one million stars of which almost 500 are known to be variable (Kaluzny et al. 2004, Weldrake et al. 2007). Recently, precise $B$ and $V$ CCD photometry of variables in $\\omega$ Cen were obtained by Cluster AgeS Experiment (CASE) team (Kaluzny et al. 2004). Among others, they published precise light curves of 151 RR Lyr stars belonging to the cluster. These light curves contained from 594 to 761 points and were collected in period from 1999 February 6/7 to 2000 August 9/10. ",
        "conclusions": ""
    },
    "0806/0806.3199_arXiv.txt": {
        "abstract": "We show data from the Survey of Ionization in Neutral Gas Galaxies (SINGG) and Survey of Ultraviolet emission in Neutral Gas Galaxies (SUNGG) which survey the star formation properties of \\HI\\ selected galaxies as traced by \\Halpha\\ and ultraviolet emission, respectively. The correlations found demonstrate a strong relationship between the neutral ISM, young massive stars, and the evolved stellar populations. For example the correlation between $R$ band surface brightness and the \\HI\\ cycling time is tighter than the Kennicutt-Schmidt Star Formation Law. Other scaling relations from SINGG give strong direct confirmation of the downsizing scenario: low mass galaxies are more gaseous and less evolved into stars than high mass galaxies.  There are strong variations in the \\Halpha\\ to UV flux ratios within and between galaxies.  The only plausible explanations for this result are that either the escape fraction of ionizing photons or the upper end of the IMF varies with galaxy mass.  We argue for the latter interpretation, although either result has major implications for astrophysics.  A detailed dissection of the massive star content in the extended \\HI\\ disk of NGC~2915 provides a consistent picture of continuing star formation with a truncated or steep IMF, while other GALEX results indicate that star formation edges seen in \\Halpha\\ are not always apparent in the UV.  These and other recent results settle some old questions but open many new questions about star formation and its relation to the ISM. ",
        "introduction": "Strong correlations between the star formation rate (SFR) of galaxies and their \\HI\\ content have been known for some time.  For example Kennicutt (1998a) showed that globally averaged star formation intensity, \\SSFR, in galaxies correlates more strongly with the \\HI\\ than with the CO surface density.  This is puzzling since stars form out of the molecular not the neutral interstellar medium (ISM).  We are working on two surveys meant, in part, to examine the nature of the \\HI\\ - star formation connection: the Survey of Ionization in Neutral Gas Galaxies (SINGG), and the Survey of Ultraviolet emission in Neutral Gas Galaxies (SUNGG).  These image nearby galaxies selected blind to optical properties from the \\HI\\ Parkes All Sky Survey (\\HIPASS) in the light of two star formation tracers: \\Halpha\\ (SINGG) and the far and near ultraviolet (FUV and NUV) continuum (SUNGG). \\Halpha\\ emission traces the presence of ionizing O stars having masses $\\Mstar \\gapeq\\ 20\\, \\Msun$, while UV emission is sensitive to both O and B stars having masses down to $\\Mstar \\gapeq\\ 3\\, \\Msun$.  Meurer et al.\\ (2006) discuss the SINGG observations and measurements. An initial description of the SUNGG survey can be found in Wong (2007).  The full description of SUNGG is currently being written by Wong et al.\\ while the preliminary results presented here are being written-up by Meurer et al.; both should be submitted for publication by the middle of 2008.  Some of the open questions we aimed to address with these surveys include What is the best form of the Star formation Law (SFL)?  Is it constant?  Is the Initial Mass Function (IMF) universal?  What is the heating source for the \\HI\\ dominated disks that are often seen to extend well past the apparent optical extent of galaxies.  There have been many papers on the SFL.  Probably the most influential have been the papers of Kennicutt and collaborators (Kennicutt, 1989; Kennicutt 1998a, Martin \\&\\ Kennicutt 2001).  They show that the \\SSFR\\ has a power law dependence on the total ISM surface density $\\SSFR \\propto \\Sigma_g^N$ where $N \\approx 1.4$, but only where the $\\Sigma_g$ is large enough for the ISM disk to be self-gravitating.  Thus extended \\HI\\ disks are thought to result if the ISM is not dense enough to form stars.  The standard assumption in much of astronomy is that the IMF is constant, which certainly seems to hold for stars clusters (Kroupa, 2001).  By using two star formation tracers we can test this assumption and probe whether the same SFL that holds for O stars also works for B stars.  The parameters derived from the SINGG and SUNGG data discussed here are based on integrated fluxes measured from concentric elliptical apertures.  In particular, \\SSFR\\ and the $R$ band surface brightness \\SR\\ are measured within the half-light radius and corrected for inclination.  The gas cycling time \\tgas\\ is derived from the ratio of the \\HI\\ and \\Halpha\\ fluxes with a crude uniform correction for helium and molecular gas.  Star formation rates are calculated using the calibrations of Kennicutt (1998b) which adopt a Salpeter (1955) IMF over the mass range of 0.1 to 100 \\Msun.  The SINGG data are corrected for dust absorption (and \\fion{N}{II} contamination) using the relationships of Helmboldt et al.\\ (2004) and validated with FIR data as shown by Meurer et al. (2006).  Dust absorption is removed from the UV fluxes based on the FUV -- NUV colors (similar to Gil de Paz et al. 2007, for example). ",
        "conclusions": "Recent work by our teams as well as others has cleared up several open questions.  The star formation, \\HI\\ and stellar light properties of galaxies are tightly correlated as shown by global scaling relations, indicating that an improved form of the star formation law is within reach.  This must have the star formation rate dependent on the stellar as well as ISM mass density, perhaps similar to the form suggested by Dopita \\&\\ Ryder (1994).  It makes sense that the stellar mass density should contribute to regulating star formation since in most galaxies stars are the major contributor to the galactic potential of the optically bright portion of galaxies, and hence are key to setting the hydrostatic pressure of galactic disks.  It appears that extended \\HI\\ disks are not empty of stars but have sparse populations of B stars that heat the disk.  It is also becoming clear that FUV and \\Halpha\\ properties are different between galaxies and even have different distributions within galaxies.  This probably indicates that the B/O ratio varies within and between galaxies, and the most likely explanation for that is the upper end of the IMF is not universal. The only alternative is that the escape fraction of ionizing photons is much larger at the low surface brightness (low mass) end of the star forming galaxy sequence.  While this is not ruled out by observations, it is contrary to naieve expectations.  Whatever the cause of the systematic \\Fhafuv\\ variations there are major implications and many new questions to resolve.  What is \\fesc\\ in low surface brightness galaxies?  If the IMF varies, which parameters vary and what drives the variations?  What is the best way to measure the SFR of galaxies?  Other open questions relate more directly to the \\HI\\ - star formation connection.  On the most basic level what is the nature of the connection?  Is \\HI\\ a tracer for the ISM that fuels the star formation, or does it represent the byproduct of the young stellar populations photo-dissociating the molecular ISM they formed out of (e.g.\\ Tilanus \\&\\ Allen 1993)?  Finally, what form of the SFL best explains the inter-relationship between star formation, the existing stars, and the \\HI\\ content of galaxies?  \\begin{theacknowledgments}  Combined the SINGG, SUNGG and ACS Instrument Definition Teams includes 69 members, all of which contributed to making these large projects function.  Unfortunately there is insufficient space to acknowledge them fully here.  This work was supported by NASA LTSA grant NAG5-13083, NASA Galex Guest Investigator grant GALEXGI04-0105-0009, and NASA grant NAG5-7697. \\end{theacknowledgments}"
    },
    "0806/0806.3975_arXiv.txt": {
        "abstract": "We report on the results of novel global high-resolution three-dimensional simulations of disk-planet interaction which incorporate simultaneously realistic radiation physics and the self-gravity of the gas, as well as allowing the planet to move. We find that thermodynamics and radiative physics have a remarkable effect on both migration and accretion of Jupiter mass planets. In simulations with radiative transfer adopting flux-limited diffusion, inward migration can be decreased by about $30\\%$ relative to the isothermal case, while in adiabatic runs migration nearly shuts off after a few tens of orbits. Migration varies because the relative strength of the inner and outer spiral perturbations is affected by thermodynamics, thus changing the net torque acting on the planet. Mass accretion rates on the planet can be reduced by more than an order of magnitude going from isothermal to radiative transfer and adiabatic simulations. A circumplanetary disk always forms except in adiabatic runs. With radiative transfer the disk is sub-keplerian ($v_{rot}/v_{kep} \\sim 0.7$) owing to significant pressure support. We discuss the effect of circumplanetary disk structure on the drift of embedded dust grains and planetesimals and thus on the formation of the rocky satellites of giant planets. ",
        "introduction": "In the conventional picture, giant planets form via a two-stage process, known as core-accretion. In such a model, first a massive rocky core is assembled via gravitational accumulation of planetesimals in the protoplanetary disk and then the core begins to accrete the surrounding gas once it has grown massive enough \\citep{Pollack96,Ida04,Alibert05}. The growing planet excites density waves as it moves through the disk. Such waves exert a torque on the planet \\citep{Lin93} whose net effect is to extract angular momentum from its orbit  for standard protosolar nebula models  \\citep{Ward97}.   Different regimes of migration have been identified by numerous studies depending on the planet's mass and on the importance of co-rotation torques \\citep[see][for a review]{PPV}. In particular, when the planet has a mass comparable to Saturn or larger the interaction with the disk becomes markedly nonlinear % (type II migration) % and the planet carves a gap in the disk as a result of the planetary tide \\citep{Lin93,Crida07}. Then, the planet decouples dynamically from the disk % and migrates inward at a pace determined by the local viscous timescale \\citep{Lin86}.  A major assumption of almost all simulations published so far is that the disk is locally isothermal. This means that any heating is immediately radiated away, which would only be true if the disk was optically thin. The planet moves supersonically in the disk, shock-heating the gas, and gas accretion onto the planet generates compressional heating. With typical densities $10^{-11}$g/cm$^3$, the midplane of a minimum solar nebula disk is indeed optically thick ($\\tau \\approx 10$) \\citep{d'Angelo03}. Recently, Paardekooper \\& Mellema (2006, 2008, hereafter \\citet{Paardekooper06,Paardekooper08}) have used both 2D and 3D adaptive mesh refinement simulations with radiative transfer modeled via flux-limited diffusion finding that for low mass planets both inward migration and gas accretion can be strongly suppressed. Klahr \\& Kley (2006, hereafter \\citet{KK06}) studied Jupiter-sized planets with a static 3D grid code and flux-limited diffusion. They did not report significant differences on migration compared with the isothermal case but found the structural evolution of the circumplanetary gas distribution to be strongly affected by radiation physics as noticed earlier in the 2D nested-grid calculations by \\citet{d'Angelo03}.   Yet, even these recent simulations lack several important ingredients. First, they do not treat self-consistently the dynamics of the disk, planet and star;  the planet and the star cannot move, and in some cases a gap is introduced already at the beginning of the simulation (KK06). Second, except for the 2D simulations of d'Angelo et al. (2003), they adopt inviscid disks. Recently, Edgar (2007) found that giant planet migration in an isothermal viscous disk does not obey the standard type II regime. A deep gap is never produced and hence migration does not proceed on the viscous timescale. Finally, all these simulations neglect the self-gravity of the gas, which is required when simulating freely moving planets and might affect disk torques (Baruteau \\& Masset 2008). Self-gravity has been previously included only in a few isothermal calculations \\citep{Nelson03a,Nelson03b,Lufkin04}.  In this Letter we present the first high-resolution three-dimensional hydrodynamical simulations of the interaction between a massive, Jupiter-sized planet and a surrounding viscous protoplanetary disk that include simultaneously radiative transfer, shock heating, self-gravity of the gas and fully self-consistent dynamics.  We study migration, mass flow towards the planet and the circumplanetary gas distribution, exploring also the implications on the formation of satellites of giant planets. We compare the results with those obtained for locally isothermal disks as well as other simulations with simplified disk thermodynamics.  \\begin{table} \\scriptsize \\caption{Simulations parameters\\label{tab:M}} \\begin{tabular}{rrrrrr} \\tableline Run   & Number & Mean      & Partial   & Mean      & Partial   \\\\ Name  & of     & accretion & accretion & migration & migration \\\\ & orbits & rate      & rate      & rate      & rate      \\\\ \\tableline Model M1 & M$_d$= & 0.004 M$_{Sol}$ &  &   &        \\\\ \\tableline IsoT1M1       &  39.8 & 6.34 10$^{-4}$  & 7.59 10$^{-4}$  & 1.26 10$^{-3}$ & 5.02 10$^{-4}$ \\\\ Adia1M1       &  80.4 & 1.02 10$^{-6}$ & 5.12 10$^{-6}$  & 7.29 10$^{-5}$ & 3.24 10$^{-4}$ \\\\ $\\star$ IsoT200K1   & 173.4 & 5.53 10$^{-5}$  & 2.46 10$^{-4}$  & 2.87 10$^{-4}$ & 6.09 10$^{-4}$ \\\\ $\\star$ FLD200K1    &  20.6 & 2.19 10$^{-5}$  & 2.19 10$^{-5}$  & 4.22 10$^{-4}$ & 4.22 10$^{-4}$ \\\\ $\\star$ NIsoT200K1  &  82.3 & 0               & 0               & 6.07 10$^{-5}$ & 3.16 10$^{-4}$ \\\\ $\\star$ Adia200K1   &  93.2 & 0               & 0               & 3.73 10$^{-5}$ & 1.83 10$^{-4}$ \\\\ \\tableline Model M2 & M$_d$= & 0.01 M$_{Sol}$ &    &     &      \\\\ \\tableline IsoT1M2       &  38.9 & 1.63 10$^{-3}$  & 1.68 10$^{-3}$  & 1.34 10$^{-3}$ & 4.01 10$^{-4}$ \\\\ FLD1M2        &   2.9 & 9.77 10$^{-4}$  & \\nodata         & 1.15 10$^{-3}$ & \\nodata        \\\\ NIsoT1M2      &  14.3 & 3.02 10$^{-4}$  & \\nodata         & 1.22 10$^{-3}$ & \\nodata        \\\\ Adia1M2       &  17.4 & 2.29 10$^{-5}$  & \\nodata         & 9.01 10$^{-4}$ & \\nodata        \\\\ IsoT100K2     & 167.7 & 1.27 10$^{-3}$  & 1.91 10$^{-3}$  & 6.48 10$^{-4}$ & 9.60 10$^{-4}$ \\\\ FLD100K2      &  27.9 & 1.14 10$^{-3}$  & 1.14 10$^{-3}$  & 8.33 10$^{-4}$ & 8.33 10$^{-4}$ \\\\ NIsoT100K2    & 111.8 & 0               & 6.77 10$^{-7}$ & 9.61 10$^{-5}$ & 4.00 10$^{-4}$ \\\\ Adia100K2     & 133.8 & 0               & 0               & 1.43 10$^{-4}$ & 2.24 10$^{-4}$ \\\\ \\tableline \\end{tabular} \\tablenotemark{k} \\tablenotetext{k}{The Table shows the accretion and migration rates for the various runs. The legend for the run names is as follows: IsoT refers to Isothermal runs, FLD to runs with the Flux Limited Diffusion approximation, NIsoT are adiabatic runs with $\\gamma$ = 6/5 and Adia are adiabatic runs with $\\gamma$ = 7/5. 100K corresponds to 100 000 particles, 200K to 200 000 particles and 1M to 1 Million particles. The runs are normally done with a planet gravitationnal softening of R$_H$/5, where R$_H$=0.35~AU is the Hill radius for a Jovian planet, except for those marked with $\\star$ for which the softening is equal to R$_H$. Accretion rates are given in M$_{J}$ yr$^{-1}$ and migration rates in AU yr$^{-1}$. Mean rates are computed by averaging over time over the full extent of the simulation. Partial rates are computed after a number of orbits equivalent to the maximum number of orbits for the FLD runs, namely after 20.6 orbits for model M1 and 27.9 orbits for model M2. When the accretion rate is 0, it means that the mass inside the Hill radius is below the SPH mass resolution.} \\end{table}   \\begin{figure}% \\plotone{f1.eps} \\caption{Temperature map with overplotted midplane density contours of run FLD1M2 after 2.9 orbits. The shock along the spiral arms triggered by the planet is evident. The planet is seen as a hot spot in the disk.} \\label{fig:ArmFLD} \\end{figure} ",
        "conclusions": "We have shown the disk thermodynamics has an impact on all the most important aspects of disk-planet interaction for Jupiter-mass planets. The way migration and accretion are affected is qualitatively consistent with the findings by \\citet{Paardekooper06} for low-mass planets, although in that work the effect was quantitatively much stronger.  Migration is slowed down relative to isothermal runs because heating modifies the relative strength of the inner and outer spiral arm, or stifles the spiral perturbations completely as in the case of adiabatic runs. This effect partially counterbalances the faster migration found in our isothermal disks relative to standard estimates that neglect viscosity, self-gravity and the motion of the planet. If the slower migration seen in radiative transfer runs persists over $> 100$ orbits one would expect the mean migration rate to differ from the estimates of standard type II migration by no more than a factor of 2, confirming the earlier conclusions reached by d'Angelo et al. (2003) using 2D simulations. The mass transport towards the planet is hampered by the increasing pressure gradient in non-isothermal runs. This will have to be investigated further by including a proper accretion model in the simulations \\citep{d'Angelo03,KK06}.   The different structure of the circumplanetary envelope depending on disk thermodynamics that we find is qualitatively in agreement with the results of \\citet{KK06}, although our disks with flux-limited diffusion are less sub-keplerian than theirs. Satellite formation is not easy in the thick circumplanetary disk, but it may still be possible provided that the solid component is dominated by large boulders."
    },
    "0806/0806.4622_arXiv.txt": {
        "abstract": "{We continue our investigation on how the cluster environment affects the evolution of galaxies.} {By examining both galaxy structure and internal kinematics of cluster galaxies at lookback times of $\\sim5$\\,Gyr we study the nature and impact of possible interactions at the peak epoch of cluster assembly. } {Going beyond our previous measurements of two-dimensional rotation curves, we here observe the whole velocity field of the galaxies of the sample.  We achieve a complete coverage and optimal spatial sampling of galaxy sizes by placing three adjacent and parallel FORS2 MXU (Mask eXchange Unit) slits onto each object yielding simultaneously several emission and absorption lines.  We reconstruct the gas velocity field and decompose it into circular rotation and irregular motions using a harmonic decomposition method called kinemetry. To measure the irregularity in the gas kinematics, we define 3 parameters: $\\sigma_{PA}$ (the standard deviation of the kinematic position angle within a galaxy), $\\Delta \\phi$ (the average misalignment between kinematic and photometric position angles) and $k_{3,5}$ (squared sum of the higher order Fourier terms).} { We present the analysis of the velocity fields and morphology of 22 distant galaxies in the MS\\,0451.6--0305 field with 11 members at $z=0.54$ and a local sample from SINGS.  Using local, undistorted galaxies the three parameters $\\sigma_{PA}$, $\\Delta \\phi$ and $k_{3,5}$ can be used to establish the regularity of the gas velocity fields.  Among the galaxies for which we could measure these parameters, we find both field ones (4 of 8)  and  cluster members (3 of 4), which have a velocity field that we consider both irregular and asymmetric.  We show that these fractions are underestimates of the total number of objects with irregular velocity fields.  The values  of the irregularity parameters for cluster galaxies are not very different from those of the field galaxies, implying that there are isolated field galaxies that are as distorted as the cluster members.  None of the deviations in our small sample correlate with photometric/structural properties like luminosity or disk scale length in a significant way. } {We have demonstrated that our 3D-spectroscopic method successfully maps the velocity field of distant galaxies.  Together with a structural analysis the importance and efficiency of cluster specific interactions can be assessed quantitatively.} ",
        "introduction": "The motion of stars and gas clouds within a galaxy are important measurable characteristics representative of the whole system.  Since the internal kinematics are subject to the overall gravitational potential they provide a proxy for the total mass.  In addition to the baryonic mass of gas and stars, that may be inferred from photometric observations, velocities trace the dark matter distribution \\citep[e.g.][]{SR01}.  The (ir)regularity of its 3-dimensional velocity field can provide clues about possible distortions of a galaxy, such as warps.  Peculiar velocity fields may also be indicators of recent or ongoing interaction processes.  This is particularly important in the environment of galaxy clusters, where specific interactions occur in addition to merging and accretion events that are observed in the field population and that are elements of the hierarchical growth of structure in the Universe \\citep[e.g.][]{Gorko04,P04}.  Cluster specific processes are presumably rather frequent at redshifts $z\\approx0.4-1$, because the assembly of galaxy clusters is expected to peak at these epochs under the conditions of the concordance cosmology \\citep[e.g.][]{B91,KSABLPST07}.  That is also the reason why fundamental properties of galaxies at these redshifts are measured and compared to those of local objects in order to explore galaxy formation and evolution.  There are considerable advantages in using 3-dimensional spectroscopy over conventional long-slit data.  Apart from a better assessment of the interaction origin of galaxies it also improves the accuracy of the establishment of scaling relations.  For example, an important tool to measure the cosmological evolution of disk galaxies is the Tully--Fisher relation \\citep[TFR,][]{TF77} where the parameter next to the intrinsic luminosity is the maximum velocity $V_{\\rm max}$ of the flat part of the rotation curve.  Only for very regular Rotation Curves (RC's), where the turn-over and the flat part are clearly visible, $V_{\\rm max}$ can be derived with sufficient accuracy to include the galaxy under scrutiny into a TF analysis.  This is particularly important in the case of distant, small and faint galaxies \\citep[e.g.][]{ZBFJN02,BZSBF03,CBRVP05,BAM06,KWFKL07,BZ07}.  The assessment of RC quality may also be the main reason why recent studies of the cluster TFR evolution differ from each other with respect to sign and amount of offset between distant galaxies in the cluster and the field environment \\citep{ZBJHM03,BMAS05,NAMAI05,MKSP06}.  Apparently smooth RCs can nevertheless result in wrong estimates of $V_{\\rm max}$, if for example the position angle of the major axis, as measured from the photometry, is not the same as the kinematic position angle.  The same holds if the photometric center does not agree with the center measured from the kinematics.  A way to resolve the situation is to obtain 3-dimensional kinematical information.  For example, \\citet{MAPB03} use velocity field of galaxies in compact groups obtained using Fabry--P\\'{e}rot spectroscopy  to show that smooth RCs can be derived for most galaxies that were previously judged to be distorted on the basis of the limited  information given by 2D-spectra \\citep{RHF91}.  Alternatively, in some cases the rotation curve along the photometric axis may look regular even though the velocity field is distorted.  Because of this, 3-dimensional information is clearly preferable to long-slit data.  It remains a requirement that the observed velocity fields should cover the flat part of the rotation curve.  Several comprehensive studies of galaxies in the Local Universe exist that explore velocity fields using optical observations.  Fabry--P\\'erot interferometry of the H$\\alpha$ emission line is the basis of the GHASP \\citep{GMABG05}, SINGS \\citep{DCAHCBK06} and Virgo \\citep{CBCCA06} surveys, for example.  3D--spectroscopy is regularly performed with the SAURON integral field unit, delivering spectra within a limited wavelength range, with which stellar absorption lines can be investigated in addition to gaseous emission lines \\citep[e.g.][]{GFPCE06,FBBCD06,SFDBBCZEFKKMP06}.  Integral--field spectroscopy of HII regions in nearby disk galaxies was established, as another example, with DensePak on the WIYN 3.5m telescope \\citep{ABSGW06}.  While spatially resolved H\\,\\textsc{i} measurements with radio telescopes are quite common locally, \\citep[e.g.][]{BOFHS05,NHSSA07} such observations at higher redshift are just becoming feasible with new instrumentation (e.g. EVLA, APEX).  At $z\\gtrsim0.2$, studies of the global velocity field of galaxies in the optical and near-infrared regimes are also quite challenging.  While in the NIR a high spatial resolution can be achieved thanks to the combination of 3D--spectroscopy with adaptive--optics techniques \\citep{FGBVESSDL06}, one is hampered by seeing effects in the optical \\citep{KKSZ07}.  One of the more comprehensive optical studies has made use of the 15 deployable small IFUs (integral field units) of the FLAMES instrument at the VLT. That way, 35 field galaxies at $0.4<z<0.7$ were examined and analyzed using the TFR \\citep{FHPAB06}.  Eliminating the galaxies that have perturbed or complex kinematics, they found no evolution in the Tully Fisher Relation since $z=0.6$.  Other investigations are restricted to a small number of objects.  Using the GMOS integral field spectrograph and exploiting the light magnification of a foreground cluster, \\citet{SBSSK06} for example, probed the emission--line properties of six $z=1$ field galaxies in the background.  Investigating where these galaxies lie on the B and I band TFR compared to local galaxies, they reported that in the B-band the galaxies are brightened by $0.5\\pm0.3$ mag, while in the I band they are in agreement with the local relation.  To determine the nature and efficiency of interaction processes in the cluster environment we continue our study of galaxies at $z\\approx0.5$ with new spectroscopy using FORS2 at the VLT and HST/ACS imaging.  Our project involves observations of four different cluster fields (MS0451.6--0305 at $z=0.54$, MS1008.1--1224 at $z=0.301$, F1557.19TC at $z=0.510$, MS2137.3--2353 at $z=0.313$) with about 20 galaxies in each of them.  While in our previous campaigns we derived the internal kinematics along the photometric major axis of the galaxies \\citep[]{ZBJHM03,JZBHM04} this time we measure the three-dimensional velocity field of the galaxies.  We achieve this by placing three adjacent, parallel FORS2 MXU (Mask eXchange Unit) slits onto the same galaxy.  In HST Cycle\\,14 the clusters were observed using the ACS camera (PID 10635), covering the full field-of-view of the spectroscopy by a $2\\times2$ mosaic allowing a detailed morphological and structural analysis of all target galaxies.  In this paper, we present the results of our kinematic and photometric analysis of the cluster MS0451 at $z=0.54$ \\citep{D96}, classify the galaxies according to the regularity of their gas kinematics and investigate whether this is related to the environment (cluster/field) and some other properties like galaxy type, luminosity, disk scale length, etc.  In follow-up papers, similar data for the other three clusters will be investigated.  In addition, we will compare the kinematic and structural analysis with hydrodynamic N--body simulations of both isolated \\citep{KKSBZ05} and interacting galaxies \\citep{KKSBK06}.  Using these simulations we already examined systematically what is the effect of observational constraints like a low spatial sampling in the case of distant small galaxies on the measured velocity fields \\citep{KKSZ07}.  In a final paper we will combine all our results to address the question of the interaction history of cluster galaxies.  In Sect.2, we outline our unique approach of  ``matched IFU simulations'' using MXU masks with the FORS2 spectrograph at the VLT and give an account of our spectral reduction, explain how the velocity fields are constructed and describe the observed sample in the field of MS0451.  Our kinematic analysis of the gas velocity fields and how we measure the stellar kinematics are also explained in this Section.  In Sect.3, we describe how we derive structural parameters from HST images and luminosities in different filters from ground-based photometry.  In Sect.4, we explain how we quantify the deviations from regularity in the gas kinematics, present the results of that analysis, make a comparison between the field and cluster galaxies and investigate the existence of a relation between the kinematic regularity and some photometric properties.  In Sect.5, we discuss our results.  In Sect.6 we summarize the paper and draw our conclusions.  In Appendix A, we discuss the effects of the spatial resolution on our kinematical analysis.  Kinematic and photometric information on individual galaxies in the MS0451 cluster field are given in Appendix B.  Throughout this paper we use a standard cosmology with: $H_0 = 70$\\,km\\,s$^{-1}$\\,Mpc$^{-1}$, $\\Omega_{\\rm m}=0.3$, $\\Omega_{\\lambda}=0.7$ \\citep{TSBCC03}. ",
        "conclusions": "In this paper we present the kinematic and photometric analysis of 22 galaxies, 11 of which are members of MS0451 at $z=0.54$.  All cluster members are within 2 Mpc from the center and most of them are closer then 1 Mpc.  Among the cluster objects, there are two in interaction with each other and 5 with weak or no emission.  For 8 field and 6 cluster galaxies we were able to obtain the gas velocity fields.  We used kinemetry to characterise these velocity fields using Fourier coefficients.  The maximum radius for which we could measure the kinematic parameters (position angle, ellipticity, etc.) of each galaxy is given in Table \\ref{tab2}.  We use 3 parameters to quantify the regularity of the velocity fields: $k_{3,5}/k_{1}$, $\\sigma_{PA}$ and $\\Delta \\phi$ (Sect.4) and introduce a criterion for each parameter to distinguish between regular and irregular gas kinematics.  To do that, we also measured each parameter for nearby galaxies in the SINGS survey and investigated the distribution of these galaxies in the parameter space.  The regularity criteria we define are: $\\Delta \\phi \\leq 25$, $\\sigma PA \\leq 20$ and $k_{3,5} / k_{1} \\leq 0.15$.  We calculated the fraction of galaxies in our sample that have irregular gas kinematics according to each criterion (without subdividing as field galaxies and cluster members).  $27 \\pm 13\\%$ (according to $\\sigma_{PA}$ criterion), $50 \\pm 16\\%$ ($\\Delta \\phi$ criterion), $40 \\pm 15\\%$ ($k_{3,5} / k_{1}$ criterion) and $58 \\pm 14\\%$ (any criterion).  Among these objects, there are four cluster galaxies for which we were able to measure these parameters.  We find that field and cluster galaxies span similar range of parameter values (Fig.\\ref{cop1}).  In particular, we find irregular cases in both environmental classes: $29 \\pm 17\\% / 25 \\pm 22\\% $ (according to $\\sigma_{PA}$ criterion), $50 \\pm 20\\% / 50 \\pm 25\\%$ ($\\Delta \\phi$ criterion), $33 \\pm 19\\% / 50 \\pm 25\\%$ ($k_{3,5} / k_{1}$ criterion) and $50 \\pm 18\\% / 75 \\pm 22\\%$ (any criterion). However, we note that the number of galaxies we have analyzed is not statistically sufficient to give robust results.  In a future paper we will analyse a much larger sample in the same way, increasing the statistics considerably.  We investigated the spatial resolution effects on our kinematic analysis and found that the fraction of irregular velocity fields in a sample determined from low spatial resolution data is an underestimate of the real fraction.  We also found out that for some combinations of a large $k_{3,5}/k_{1}$ and sampling, the measured values of $\\Delta\\phi$ and $\\sigma_{PA}$ will be boosted, so that the different types of irregularity become degenerate in the rebinned data. Hence, identifying irregular galaxies without distinguishing between the three types of irregularity gives more reliable results in case of low spatial resolution.  6 galaxies (5 of them are cluster members) have strong absorption lines (mainly CaII H,K and Balmer lines) which enabled us to derive their stellar rotation curves. For three of them (C5, C7 and C9), stellar and gas curves show different behaviours.  These are galaxies that are classified to have irregular gas kinematics (C5 is interacting with its companion).  C7 and C9 also have relatively large photometric asymmetries.  We have performed a surface photometry analysis of each galaxy (Table \\ref{tab4}), and calculated their rest-frame absolute B band magnitude (Table \\ref{tabrun2}). We also measured some other photometric properties (e.g. asymmetry and concentration index) and investigated the existence of a correlation of these parameters with the kinematic regularity. Although we do not find any significant correlation, these should be re-investigated with a larger sample."
    },
    "0806/0806.3002_arXiv.txt": {
        "abstract": "We present new data obtained with the Submillimeter Array for a sample of fourteen nearby luminous and ultraluminous infrared galaxies. The galaxies were selected to have distances $D_L < 200$ Mpc and far-infrared luminosities $\\log L_{\\rm FIR} > 11.4$. The galaxies were observed with spatial resolutions of order 1 kpc in the CO J=3-2, CO J=2-1, $^{13}$CO J=2-1, and HCO$^+$ J=4-3 lines as well as the continuum at 880 $\\mu$m and 1.3 mm. We have combined our CO and continuum data to measure an average gas-to-dust mass ratio of $120 \\pm 28$ (rms deviation 109) in the central regions of these galaxies, very similar to the value of 150 determined for the Milky Way. This similarity is interesting given the more intense heating from the starburst and possibly accretion activity in the luminous infrared galaxies compared to the Milky Way. We find that the peak H$_2$ surface density correlates with the far-infrared luminosity, which suggests that galaxies with higher gas surface densities inside the central kiloparsec have a higher star formation rate. The lack of a significant correlation between total H$_2$ mass and far-infrared luminosity in our sample suggests that the increased star formation rate is due to the increased availability of molecular gas as fuel for star formation in the central regions. In contrast to previous analyses by other authors, we do not find a significant correlation between central gas surface density and the star formation efficiency, as trace by the ratio of far-infrared luminosity to nuclear gas mass. Our data show that it is the star formation rate, not the star formation efficiency, that increases with increasing central gas surface density in these galaxies. ",
        "introduction": "Ultra-luminous infrared galaxies (ULIRGs) contain the regions of most intense star formation in the local universe. Although their high rates of star formation and accretion appear to be triggered by the merger of two gas-rich galaxies \\citep{s88a,v02}, the detailed physical connection between galaxy mergers and star formation and, in particular, the time evolution of this process, is not well understood. Relating numerical hydrodynamical models \\citep{miho96,cox04} to observations is complicated by the difficulty in identifying the precise stage of the merger \\citep{murp01}. In addition, while high resolution imaging has found that most ULIRGs have nuclear separations from $<$0.3 kpc to 48 kpc \\citep{murp96}, other strongly interacting galaxies with these nuclear separations which are {\\it not} ULIRGs have also been found \\citep{brai04}. These observations suggest that the onset of the intense star formation and accretion which produces a ULIRG is not a simple function of the age of the merger and leaves open the question of whether all luminous infrared galaxies (LIRGs\\footnote{$L_{\\rm FIR} = 4 \\pi D_L^2 F_{\\rm FIR}$ $L_\\odot$, where $F_{\\rm FIR} = 1.26 \\times 10^{-14} (2.58 f_{60} + f_{100})$ erg cm$^{-2}$ s$^{-1}$ and $f_{60}$ and $f_{100}$ are the IRAS fluxes in Jy at 60 and 100 $\\mu$m \\citep{sanders96}.}; $11 \\le \\log (L_{\\rm FIR}/L_\\odot) < 12$) will pass through a ULIRG phase ($\\log (L_{\\rm FIR}/L_\\odot) \\ge 12$) at some point in their evolution.  Local ULIRGs are also important as the closest analogs to the high-redshift submillimeter galaxies \\citep[SMGs;][]{blai02}: both populations have high infrared luminosities, large amounts of molecular gas \\citep{fray98,fray99,neri03,grev05,t06}, and morphological evidence of recent or ongoing mergers \\citep{v02,cons03}. Since galaxy merger rates are substantially higher in the early universe \\citep{lefe00,gott01}, understanding the physical and dynamical properties of nearby ULIRGs is also important for understanding the processes in the early universe which give rise to the very luminous submillimeter galaxy population.  Because molecular gas is the fuel for current and future star formation, the physical properties and distribution of the warm, dense molecular gas are crucial for understanding the processes and timescales controlling star formation in galaxy mergers. Previous high-resolution studies of molecular gas in luminous infrared galaxies have used primarily the ground-state rotational transition of CO, which is sensitive to gas as cold as 10~K \\citep{s91,ds98,bs99}, with a few galaxies observed in the CO J=2-1 line \\citep{bs96,ds98,s99,t99}. However, since the CO J=3-2 line traces the warmer and denser gas, it is more likely to be directly associated with the starburst activity and/or fueling of the active galactic nuclei (AGN) in these galaxies. Indeed, observations of the CO J=3-2 emission in two luminous infrared galaxies,  VV 114 \\citep{iono04} and NGC 6090 \\citep{wang04}, reveal that the large-scale distribution and kinematics of the CO J=3-2 line can be significantly different from those of the CO J=1-0 line.  In this paper, we present new data obtained with the Submillimeter Array (SMA) for a sample of fourteen luminous and ultraluminous infrared galaxies in the CO J=3-2, CO J=2-1, $^{13}$CO J=2-1, and HCO$^+$ J=4-3 lines. In addition, we present new high-resolution observations of continuum emission at 880 $\\mu$m and 1.3 mm, which allow us to study the dust properties in the central kiloparsec of the galaxies. This SMA legacy survey aims to address five broad scientific questions:  \\begin{enumerate}  \\item {\\it What are the distributions, kinematics, and physical conditions of dense molecular gas in U/LIRGs?} The high resolution CO J=3-2 data cubes trace the distribution of the warm and dense gas that feeds the starburst (and any accretion) activity in these luminous galaxies. The new SMA CO J=3-2, J=2-1, and $^{13}$CO J=2-1 data can be combined with published CO J=1-0 and J=2-1 data for a detailed investigation of the physical properties of the molecular gas using Large Velocity Gradient \\citep[LVG,][]{ss74,gk74} and Monte Carlo \\citep{juve97} models as our primary diagnostic tools. The CO J=3-2 kinematics  allow us to study the detailed gas dynamics in the inner few hundred parsecs, yielding measurements of the total enclosed mass and of the local linewidth that is a parameter in models of disk turbulence. The combination of morphology and kinematics offers clues to the geometry of each merger via comparison of the separation and orientation of the galaxy nuclei with the results from numerical simulations \\citep[see, e.g.,][for an analysis of Arp 220]{m01} .  \\item {\\it What is the distribution of the dust in U/LIRGs?} The 880 $\\mu$m  continuum images trace the spatial distribution of the cold and warm (10--70 K) dust, which reflects both the local rate of star formation activity and the available mass of gas. The submillimeter dust emission is often significantly more compact than that of CO \\citep[e.g.,][]{s99,mt01,s06} % arising exclusively from deep in the gravitational potential wells of the galactic nuclei. High-resolution continuum images from the SMA can be combined with spectra from the {\\it Spitzer Space Telescope} \\citep[e.g.,][]{a07} to estimate the dust temperature via the mid-infrared to submillimeter spectral energy distribution (SED), the dust mass % \\citep[including both small and large grains, e.g.,][]{m07} and, indirectly, the gas mass based on 3D radiative transfer modeling \\citep{j03}. In addition, these spatially resolved SEDs of local U/LIRGS will improve our interpretation of the templates used for determining photometric redshifts of high-redshift submillimeter galaxies \\citep{yun02,a03,a05}.   \\item {\\it Do the gas properties change as the interaction progresses?} Our sample of fourteen U/LIRGs covers a range from mid to late merger stages and should be sufficiently large that we can establish a merger sequence through comparison of global morphologies with numerical models \\citep{miho96,cox04}. This data set allows us to determine how the distribution and kinematics of the gas change as a function of physical conditions such as density and temperature, or vice versa, and to correlate those changes with the stage of the merger. This type of detailed gas physics on small scales still poses challenges for numerical simulations. Thus, on large scales, where the dynamics of the system are well described by the models, the numerical simulations can help with the interpretation of the data, while on smaller scales, the data can drive the development of more accurate descriptions for the gas physics in the simulations.  \\item {\\it How do the properties of the dense gas in local U/LIRGs compare to those of the gas in high-redshift submillimeter galaxies?} Armed with a robust local sample of fourteen U/LIRGs, we can make a rigorous comparison of the properties of the gas with those in higher redshift galaxies \\citep{grev05,t06}. Changes in gas characteristics over the age of the universe will reveal important information about the process of star formation and ultimately the formation and evolution of galaxies at early times.  \\item {\\it  What is the origin of nuclear OH megamasers?} Bright 1667 MHz OH megamaser emission is observed in the nuclei of some luminous infrared galaxies, including five galaxies in our sample (IRAS 17208-0014, Mrk 231, Mrk 273, UGC 5101 and Arp 299). These extremely bright masers are promising tracers of dust-obscured star formation and mergers at high redshifts, and could ultimately be used to estimate the merger rate as a function of redshift \\citep{darl02}. However, in order to use OH megamasers as high-redshift probes, we need to understand whether there is a specific type or stage of merger that leads to OH maser emission. Whether maser emission occurs is likely governed by the physical, chemical, and kinematic conditions in the molecular gas in the nuclear regions of the mergers. For example, using global CO and HCN luminosities, \\citet{darling07} concludes that OH megamasers are associated with high mean molecular gas densities and high dense gas fractions. Our sample, which contains galaxies with and without megamasers, is well suited for identifying any unique nuclear conditions that produce OH megamasers in luminous infrared galaxies. \\end{enumerate}  In this paper, we describe the sample selection, observations, and reduction of the survey data (\\S~\\ref{sec-obs}). We also examine the gas to dust mass ratio in the central kiloparsec (\\S~\\ref{sec-gasdust}) and correlations between the central gas mass, gas surface density, infrared luminosity, dust temperature, and CO J=3-2/2-1 line ratio (\\S~\\ref{sec-corr}). A companion paper presents a detailed analysis of the galaxy VV 114 \\citep{p08}; a detailed analysis of NGC 6240 is given in \\citet{i07}. Future papers  will compare the results from this survey with similar observations of high-redshift submillimeter galaxies \\citep{i08}; examine the physical properties of the molecular gas by combining molecular line observations with radiative transfer models; use the molecular gas data to place constraints on the origin of the OH megamaser activity seen in some of the galaxies in our sample; compare the properties of the cold gas and dust as seen with the SMA with the properties of the warm dust derived from {\\it Spitzer} data; and compare the molecular gas and dust properties with the predictions of numerical simulations to place the galaxies into a merger sequence. ",
        "conclusions": "In this paper, we have presented new data obtained with the Submillimeter Array for a sample of fourteen luminous and ultraluminous infrared galaxies selected to have distances $D_L < 200$ Mpc and far-infrared luminosities $\\log L_{\\rm FIR} > 11.4$. We have obtained data in the CO J=3-2, CO J=2-1, $^{13}$CO J=2-1, and HCO$^+$ J=4-3 lines as well as continuum data at 880 $\\mu$m and 1.3 mm with spatial resolutions of order 1 kpc or better in all but one of the target galaxies. We present integrated intensity, velocity field, and velocity dispersion maps for the $^{12}$CO lines, integrated intensity maps for the continuum, $^{13}$CO, and HCO$^+$ lines, and peak and integrated spectra for all the detected lines.  We have compared our CO J=3-2 and 880 $\\mu$m continuum fluxes detected with the SMA with published, archival, and new data from the James Clerk Maxwell Telescope. This comparison shows that the interferometric data miss a significant fraction (typically 50\\%) of the CO J=3-2 emission for eight of the galaxies in our sample and also a significant fraction (typically 50-80\\%) of the continuum flux for nine of the galaxies. This large amount of missing continuum flux suggests that a significant fraction of the 880 $\\mu$m emission in these systems occurs on moderately large spatial scales. The good agreement between the percentage of missing flux seen in the CO J=3-2 line and the 880 $\\mu$m continuum suggests that the missing continuum flux comes from dust emission associated with molecular gas in the more extended disks of the galaxies.  We have combined our CO and continuum data to determine the gas-to-dust mass ratio in the central regions of these galaxies. We adopt the smaller value of the  CO-to-H$_2$ conversion factor from \\citet{ds98} and calculate the dust temperature by fitting a modified blackbody function as in \\citet{k01}. Because of the lower signal-to-noise ratio in the continuum data, we find that we obtain more consistent measurements of the gas-to-dust mass ratio if we use a single beam to probe the central region of each galaxy or galaxy component. We find an average gas-to-dust mass ratio of $120 \\pm 28$ (rms deviation 109), very similar to the value of 150 determined for the Milky Way. This similarity between the gas-to-dust ratio in these luminous systems and that in the Milky Way is somewhat surprising, given that the dust is subject to more intense heating from the starburst and possibly accretion activity compared to typical regions in the Milky Way.  We have searched for correlations among nine physical and observational quantities for the galaxies in our sample. We find five correlations that appear to be statistically significant as well as robust to small changes in the exact galaxy sample. The most interesting correlation is that of peak H$_2$ surface density with the far-infrared luminosity. Since the far-infrared luminosity can be used to estimate the star formation rate, these correlations suggest that galaxies with higher gas surface densities inside the central kiloparsec have a higher star formation rate. We do not see a significant correlation of total H$_2$ mass with the far-infrared luminosity, which suggests that the increase in star formation rate is due to the increased availability of molecular gas as fuel for star formation in the central regions, rather than the total amount of gas available on somewhat larger scales.  Our data do not show any evidence of a significant correlation between central gas surface density and the ratio of far-infrared luminosity to nuclear gas mass. This lack of correlation is different from the results of \\citet{s91}, who interpreted their observed correlation as indicating that higher star formation efficiencies result from higher gas surface densities. We suggest that the correlation seen by \\citet{s91} was produced by the wider distribution of spatial resolutions in their data set  and is not an intrinsic property of these very luminous galaxies. To reiterate, our new data show that it is star formation {\\it rate}, not star formation {\\it efficiency}, that increases with the central gas surface density in luminous and ultraluminous infrared galaxies  There are a number of additional papers in preparation or planning that will present more detailed analysis of various aspects of the data. We will compare the results from this survey with similar observations of high-redshift submillimeter galaxies to study the gas properties of a wide range of luminous galaxies using the CO J=3-2 line to trace the molecular gas content \\citep{i08}. A detailed analysis of the molecular gas properties of NGC 6240 has already been presented in \\citet{i07}; we will present similar detailed analyses of the gas, dust, and star formation properties individual galaxies such as VV 114 \\citep{p08} and Arp 299. We will examine the physical properties of the molecular gas for the entire sample using spatially resolved radiative transfer models, similar to what has been done for NGC6240 \\citep{i07}, as well as carry out dynamical analysis and modeling of the galaxies, both of which can give an independent estimate of the CO-to-H$_2$ conversion factor, similar to the analysis of \\citet{ds98}. This physical analysis will also allow us to place constraints on the origin of the OH megamaser activity in luminous infrared galaxies \\citep{darling07}. We will combine our high-resolution SMA data with {\\it Spitzer} data to compare the properties of the warm and cold dust \\citep[see also][]{a07,m07}. Finally, we will make detailed comparisons between the molecular gas and dust properties of these U/LIRGs and  the predictions of numerical simulations of merging galaxies \\citep[e.g.,][]{c06,c07}."
    },
    "0806/0806.2835_arXiv.txt": {
        "abstract": "We use baryon acoustic peak measurements by \\citet{eisensteinetal} and \\citet{percivaletal07a}, together with the WMAP measurement of the apparent acoustic horizon angle, and galaxy cluster gas mass fraction measurements of \\citet{allen08}, to constrain a slowly rolling scalar field dark energy model, $\\phi$CDM, in which dark energy's energy density changes in time. We also compare our $\\phi$CDM results with those derived for two more common dark energy models: the time-independent cosmological constant model, $\\Lambda$CDM, and the XCDM parametrization of dark energy's equation of state. For time-independent dark energy, the \\citet{percivaletal07a} measurements effectively constrain spatial curvature and favor a close to spatially-flat model, mostly due to the WMAP CMB prior used in the analysis. In a spatially-flat model the \\citet{percivaletal07a} data less effectively constrain time-varying dark energy. The joint baryon acoustic peak and galaxy cluster gas mass constraints on $\\phi$CDM model are consistent with but tighter than those derived from other data. A time-independent cosmological constant in a spatially-flat model provides a good fit to the joint data, while the $\\alpha$ parameter in the inverse power law potential $\\phi$CDM model is constrained to be less than about 4 at 3$\\sigma$ confidence level. ",
        "introduction": "About a decade ago type Ia supernova (SNIa) observations provided initial evidence that the cosmological expansion is accelerating \\citep{riess98, perlmutter99}. If general relativity is valid on the scales of current cosmological observations, more recent SNIa data as well as results of various other cosmological tests including large-scale structure tests and observations of the cosmic microwave background (CMB) anisotropy can be reasonably well reconciled if we assume that about two-thirds of the cosmological energy budget is in the form of dark energy.  Different theoretical models of dark energy have been proposed over the years. A few models try to do away with the need for an exotic dark energy component by modifying general relativity on large scales \\citep[see, e.g.,][]{wang07a, demianski07, tsujikawa08, capozziello08, wei08, gannouji08}. If general relativity is valid we need a substance that has negative pressure, $p<-\\rho/3$ (where $\\rho$ is the energy density), to have accelerated cosmological expansion. The simplest standard cosmological model is $\\Lambda$CDM \\citep{pee84} in which the cosmological constant $\\Lambda$ has negative pressure and powers the current accelerated expansion of the universe. Although $\\Lambda$ has a quantum field theory motivation as vacuum energy, $\\Lambda$CDM has a number of apparent problems. The most celebrated is the fact that the value of vacuum energy density calculated from field theory with a Planck scale cutoff is many orders of magnitude larger than the measured value. Because of this other models have been developed, despite the fact that the simple $\\Lambda$CDM model provides a fairly good fit to most cosmological data. In our paper we also study the slowly rolling scalar field dark energy model  \\citep[$\\phi$CDM,][]{pee88, rat88}. In the $\\phi$CDM model the small (classical) value of the current vacuum energy density is a consequence of the scalar field dynamics. The third model we consider is the XCDM parametrization. XCDM parametrizes dark energy's equation of state as $p_{\\rm x}=\\omega_{\\rm x}\\rho_{\\rm x}$, where $\\omega_{\\rm x}$ is a negative constant. This approximation is not accurate in the scalar field dominated epoch \\citep{rat91}.\\footnote{For recent reviews of dark energy see, e.g., \\citet{rat08}, \\citet{linder08}, \\citet{frieman08}, and \\citet{martin08}.}$^,$\\footnote{We assume that dark energy and dark matter only couple gravitationally. For discussion of models with other couplings see, e.g., \\citet{costa08}, \\citet{mainini07}, \\citet{brookfield08}, \\citet{he08}, and \\citet{olivares08}. For other models of dark energy see, e.g., \\citet{grande07}, \\citet{neupane08}, \\citet{mathews08}, \\citet{usmani08}, and \\citet{ichiki08}.}  In the $\\phi$CDM model one can explain the accelerated expansion of the universe by introducing a scalar field $\\phi$ minimally coupled to gravity. The action for such a term is  \\begin{equation} S_{\\phi}=\\int d^4x\\sqrt{-g}\\left[\\frac{1}{2}g^{\\mu\\nu}\\partial_{\\mu}\\phi\\partial_{\\nu}\\phi-\\frac{V(\\phi)}{G}\\right], \\end{equation}  \\noindent where $G$ is the gravitational constant. If the scalar field is close to homogeneous on cosmological scales, then to leading order it's energy density and pressure are given by  \\begin{equation} \\rho_{\\phi}=\\frac{1}{2}\\left(\\frac{d\\phi}{dt}\\right)^2+\\frac{V(\\phi)}{G}, \\end{equation}  \\begin{equation} p_\\phi=\\frac{1}{2}\\left(\\frac{d\\phi}{dt}\\right)^2-\\frac{V(\\phi)}{G}, \\end{equation}  \\noindent When the scalar field changes only slowly in time, the effective equation of state parameter $\\omega_{\\phi}=p_{\\phi}/\\rho_{\\phi}$ is negative and the scalar field acts like a time-dependent cosmological constant. To specify the $\\phi$CDM model one has to pick a specific form of potential energy density $V(\\phi)$. Neither cosmological observations nor fundamental particle physics theory can provide significant motivation for a specific form of potential energy and a lot of different cases have been studied. In our paper we work with the inverse power law potential energy density $V(\\phi)\\varpropto \\phi^{-\\alpha}$, because it has been well studied and it provides a practical way of parametrizing the slowly rolling scalar field with one nonnegative dimensionless parameter $\\alpha$. Physically, large values of $\\alpha$ correspond to rapid time evolution, while the limit of $\\alpha=0$ gives a time-independent cosmological constant. \\citet[Fig.\\ 2]{podariu00} relate this $\\phi$CDM model to the XCDM parametrization and discuss how $\\alpha$ and effective $\\omega_{\\phi}$ are related. For large values of $\\alpha$ the time-dependent equation of state parameter changes very fast and the XCDM parametrization fails to provide a good phenomenological description of the scalar field. Figure~1 shows the residuals between comoving distance calculated in $\\phi$CDM and XCDM models. Already for $\\alpha=2.0$ the predictions of XCDM differ significantly at high redshifts. If $\\alpha$ is very close to zero the scalar field equation of state changes very slowly and becomes more and more difficult to distinguish from $\\Lambda$CDM and the XCDM parametrization is reasonable. If $\\alpha$ turns out to be a very small but nonzero number, a lot of independent high precision cosmological measurements supplemented with the better understanding of underlying high energy physics will be necessary to discriminate between different dark energy models.  Assuming the cold dark matter (CDM) model of structure formation \\citep[for a discussion of apparent problems with this model see][and references therein]{pee03}, and assuming that the dark energy is a time-independent cosmological constant \\citep[see, e.g.,][]{wang07, gong07, ichikawa08, virey08}, CMB anisotropy data combined with independent dark matter density measurements \\citep[see, e.g.,][]{che03b} are consistent with negligible spatial curvature \\citep[see, e.g.,][]{pod01b, page03, spergel07, doran07}. CMB anisotropy data in combination with the low measured density of nonrelativistic matter then require the presence of dark energy and so are consistent with the SNIa results.  Many different observational tests have been used to constrain cosmological parameters. An issue of great current interest is whether dark energy is Einstein's cosmological constant or whether it evolves slowly in time and varies weakly in space. Current SNIa data are unable to resolve this \\citep[see, e.g.,][]{mignone07, wuyu08, lin08, dev08, liu08, kowalski08}, but future SNIa data will improve the constraints \\citep[see, e.g.,][]{pod01a} and, unless $\\alpha$ has a very small value, might be able to detect time variation of dark energy. Current SNIa and CMB data are consistent with the $\\Lambda$CDM model, but it is not yet possible to reject other dark energy models with high statistical confidence \\citep[see, e.g.,][]{rap05, wilson06, davis07}.  To tighten the constraints on cosmological parameters, it is important to have many independent tests of dark energy models. Comparison of constraints from different tests can help uncover unknown systematic effects, and combinations of constraints from different tests can better discriminate between models. Other observational tests under recent discussion include the angular size of radio sources and quasars as a function of redshift \\citep[see, e.g.,][]{chen03a, pod03, dal07, santos08}, strong gravitational lensing \\citep[see, e.g.,][]{lee07, oguri08, zhang07, zhu08}, weak gravitational lensing \\citep[see, e.g.,][]{takada08, fu07, dore07, lavacca08}, measurements of the Hubble parameter as a function of redshift \\citep[see, e.g.,][]{samushia06, lazkoz07b, weizhang08, szydlowski08}, large-scale structure baryon acoustic oscillation peak measurements \\citep[see, e.g.,][]{xia07, lim07, sapone07} and galaxy cluster gas mass fraction versus redshift data \\citep[see, e.g.,][]{allen04, che04, sen08}. For recent reviews of the observational constraints on dark energy see, e.g., \\citet{kur07} and \\citet{wang07b}.  Many different observational test have been used to constrain the slowly-rolling scalar field dark energy model. The constraints are getting tighter as the quality and quantity of new measurements is increasing. Constraints on the $\\alpha$ parameter from different tests are shown in Table\\ 1. In our paper we use baryon acoustic oscillations (BAO) peak measurements to constrain the $\\phi$CDM model of dark energy and compare our results to the constraints on $\\Lambda$CDM model and XCDM parametrization. Since the peak has been measured at only two redshifts, $z=0.2$ and $z=0.35$ \\citep{eisensteinetal, percivaletal07a}, BAO data alone can not tightly constrain the models. To more tightly constrain the dark energy models, we perform a joint analysis of the BAO data with new galaxy cluster gas mass fraction versus redshift data \\citep{allen08}.\\footnote{The galaxy cluster gas mass fraction test was proposed by \\citet{sasaki96} and \\citet{pen97}.} The resulting constraints are consistent with, but typically more constraining than, those derived from other data (see Table\\ 1).  In Sec.\\ 2 we briefly describe the BAO method we use. In Sec.\\ 3 we summarize the BAO and galaxy cluster gas mass fraction data and computations. We discuss our results in Sec.\\ 4. ",
        "conclusions": "The \\citet{eisensteinetal} BAO peak measurement has been used in conjunction with other data to place constraints on various cosmological models \\citep[see, e.g.,][]{alam06, nesseris07, movahed07, zhangwu07, wright07, shafieloo07}. The more recent \\citet{percivaletal07a} data has also been used for this purpose \\citep{ishida08, lazkoz07a}.  Constraints from BAO peak measurements and galaxy cluster gas mass fraction data are shown in Figs.\\ 2, 3, and 4. The solid line contours in Figs.\\ 2 and 3 show the constraints on $\\Lambda$CDM and XCDM derived from the \\citet{percivaletal07a} BAO data and are comparable to those shown with dashed lines in Fig.\\ 12 in their paper. The dashed contours in Fig.\\ 3 are comparable to those shown in Fig.\\ 11 in \\citet{eisensteinetal}. \\citet{eisensteinetal} do not show contours for $\\Lambda$CDM (see our Fig.\\ 2) and the BAO contours we show in Fig.\\ 4 have not previously been presented. Figure\\ 2 shows that the \\citet{percivaletal07a} constraints, which make use of the WMAP measurement of the apparent acoustic horizon angle, constrain the sum of parameters $\\Omega_\\Lambda$ and $\\Omega_{\\rm m}$ to be very close to one ($\\Omega_{\\rm k}=1-\\Omega_{\\rm m}-\\Omega_\\Lambda \\approxeq 0$) and favor a close to spatially flat model if dark energy is time independent. The spatial curvature is constrained so well mainly because we use the WMAP measurement of the apparent acoustic peak angle. BAO measurements by themselves can not effectively constrain dark energy parameters very well \\citep[see shaded areas in Fig.\\ 12 of][]{percivaletal07a}. In spatially-flat models BAO peak measurements put tight constraints on the $\\Omega_{\\rm m}$ parameter; they do not well constrain the ``orthogonal'' cosmological parameter $\\Omega_\\Lambda$ and in particular they allow dark energy to vary in time (see Figs.\\ 3 and 4).   The BAO constraints are significantly tighter than the Hubble parameter versus redshift data ones \\citep[see, e.g.,][]{samushia07} and the strong gravitational lensing ones \\citep[see, e.g.,][]{cha04}. They are, in general, about as constraining as the SNIa results and constrain roughly the same linear combination of cosmological parameters \\citep[see, e.g.,][]{wilson06}. In Figs.\\ 2 and\\ 3 the best fit values from \\citet{percivaletal07a} measurement and WMAP prior are more then 3$\\sigma$ away from the best fit of the cluster gas mass fraction constraints. This is most probably due to unknown systematic errors in one or both of the measurements or an effect of poor statistics and should change when more and better data are available.  The joint BAO peak and cluster gas mass fraction constraints are shown in Figs.\\ 5, 6, and 7. They are fairly restrictive and favor a spatially-flat $\\Lambda$CDM model with $\\Omega_{\\rm m} \\sim $ 0.25 and $\\alpha < 0.5$ on $1\\sigma$ confidence level. Since the predictions of $\\phi$CDM for a very small value of the $\\alpha$ parameter are very close to the predictions of the spatially-flat $\\Lambda$CDM model, current observational tests are unable to discriminate between a time-independent cosmological constant and a slowly varying scalar field with $\\alpha$ of order 1. All three models considered here give about the same $\\chi^2 \\sim 52$ for 41 degrees of freedom and there is no reason to favor one model over another based on Bayesian statistics. The constraints from the joint analysis on all three dark energy models are comparable to the constraints derived from a joint analysis \\citep{wilson06} of earlier SNIa data \\citep{riess04} and earlier cluster gas mass fraction data \\citep{allen04}. Constraints on $\\alpha$ derived from the joint analysis are stronger than the results quoted in previously published papers (see Table\\ 1) The joint BAO peak and gas mass fraction data constraints on $\\Lambda$CDM and XCDM derived here are a little weaker than those derived from BAO peak and more recent SNIa \\citep{astier06} data, see Figs.\\ 13 of \\citet{percivaletal07a}. In the joint analysis done here the uncertainties on $h$ and $\\Omega_{\\rm b}h^2$ play a less significant role than they do in the cluster gas mass fraction analysis, i.e., the contours for the two prior sets are closer to each other in Figs.\\ 5,\\ 6, and\\ 7, than in Figs.\\ 2,\\ 3, and\\ 4. The contours in Figs.\\ 5 and\\ 6 are in agreement with tighter joint results from other data sets considered by \\citet{wangetal07}.  From Fig.\\ 1 it is clear that for large values of $\\alpha$ $\\phi$CDM and XCDM models predict different cosmological evolution. For small values of $\\alpha$, however, if parameters are chosen appropriately, different dark energy models will at low redshifts predict very similar background evolution. Because of that, low redshift distance measurements have to be complemented with high redshift CMB and large scale structure measurements to discriminate between dark energy models. Better quality BAO peak data at a number of redshifts and more gas mass fraction measurements, along with tighter priors on nuisance parameters like the Hubble parameter and the density of baryonic matter, will allow for tighter constraints on dark energy parameters and could soon either detect a time dependence in dark energy or constrain it to a very small value."
    },
    "0806/0806.2232_arXiv.txt": {
        "abstract": "Only by incorporating various forms of feedback can theories of galaxy formation reproduce the present-day luminosity function of galaxies. It has also been argued that such feedback processes might explain the counter-intuitive behaviour of `downsizing' witnessed since redshifts $z\\simeq$1-2. To examine this question, observations spanning $0.4 < z < 1.4$ from the DEEP2/Palomar survey are compared with a  suite of equivalent mock observations derived from the Millennium Simulation, populated with galaxies using the {\\sc Galform} code. Although the model successfully reproduces the observed total mass function and the general trend of  `downsizing', it fails to accurately reproduce the colour distribution and type-dependent mass functions at all redshifts probed. This failure is shared by other semi-analytical models which collectively appear to  ``over-quench'' star formation in intermediate-mass systems.  These mock lightcones are also a valuable tool for investigating the reliability of the observational results in terms of cosmic variance.  Using variance estimates derived from the lightcones we confirm the significance of the decline since $z \\sim 1$ in the observed number density of massive blue galaxies which, we argue, provides the bulk of the associated growth in the red sequence.  We also assess the limitations arising from cosmic variance in terms of our ability to observe mass-dependent growth since $z \\sim 1$. ",
        "introduction": "The physical picture of how galaxies assemble has changed markedly over the past decade. A pure `hierarchical dark matter model' in which gas cooling and subsequent star formation occurs in synchronisation with the growth, via gravitational instability, of  their parent dark matter halos fails to reproduce the local luminosity function of galaxies \\cite{Benson03,Somerville99,Kauffmann99} and has been challenged by the presence of massive ($\\simeq10^{11}\\,M_{\\odot}$) galaxies at redshifts $z\\simeq$2 \\cite{Glazebrook04,Cimatti04,vanDokkum06}. As a result, a new paradigm has emerged which argues for the importance of  `feedback' processes that serve to govern the star formation rate in a galaxy. As the efficacy of these processes depends on the mass of the host galaxy, so it is possible to reconcile the predictions of the standard CDM model with the local galaxy luminosity function \\cite{Croton06,Bower06}.  Despite this progress, the physical basis of the feedback processes incorporated into the recent semi-analytic models remains largely untested. The most effective way of suppressing star formation and hence inhibiting further growth in massive galaxies is ``radio mode'' feedback \\cite{Croton06,Bower06}, where additional gas cooling in halos in which the cooling time is longer than the dynamical time is prevented by low levels of accretion onto central supermassive black holes. \\scite{Bower06} have argued that such a process can lead naturally to a characteristic mass scale associated with the transition between cooling on a hydrostatic timescale and more rapid cooling.  Although such a feedback mode can be arranged to match the break in the present day luminosity function, a key issue is whether it explains the trajectory of star formation in galaxies over the past 5-10 Gyr.  Similar progress has been made observationally in measuring the evolving stellar mass function of galaxies over 0$ < z < $1.5, where large and complete samples can be obtained \\cite{Fontana04,Drory04,Bundy06,Borch06,Pozzetti07}.  The advent of large-format near-infrared detectors used in conjunction with deep, spectroscopic and multi-wavelength surveys has characterized the evolving stellar mass function \\cite{Fontana04}, illuminated the bimodal nature of local galaxies \\cite{Kauffmann03}, demonstrated the presence of morphological evolution associated with assembly since $z \\sim 1$ \\cite{Brinchmann00} and revealed how the quenching of star formation in massive galaxies produces the downsizing signature \\cite{Bundy06}.  The time is therefore ripe for a direct confrontation between recent simulations which incorporate ``radio mode'' feedback to fit the local luminosity function and the history of mass assembly over $0<z<1.5$ from the new generation of deep surveys. By using stellar mass estimates as a bridge between theory and observations we can gain significant insight not only into the success of the physical prescriptions employed by these models \\cite{Kitzbichler06}, but also the utility of observations in answering the questions posed above.  A key issue in comparing theoretical predictions with observations is the reliability of the latter. Previous comparisons \\cite{KW07,Bower06,Cole01} have been content to compare the {\\it total} stellar mass function. However, the signature of downsizing in star formation is most readily tested by examining the mass function partitioned into star-forming and quiescent populations \\cite{Bundy06}. While their fractional contributions are relatively easily measured, the absolute numbers of these sub-populations are less certain because of the limitations of cosmic variance.  This must be properly understood in any comparison, whether it be theory versus data or one dataset against another.  As an example of the uncertainties arising from cosmic variance, we note that while it is generally accepted in the community that the number of red-sequence galaxies grows with time, confusion remains over whether this growth arises at the expense of a decline in the blue population \\cite{Bundy06}, or is dominated by ``dry'' mergers occurring within the red sequence, leaving the population of blue galaxies mostly invariant \\cite{Faber07}. In observational samples broken by redshift bin and galaxy class, the impact of cosmic variance becomes a critical issue, sufficient perhaps to explain differences in empirical interpretations. In this paper we will use numerical simulations not only to test popular models of feedback against observations, but also to evaluate rigorously the limitations in the data arising from cosmic variance.  Our work follows logically from the earlier study of \\scite{KW07}. Those authors considered the output of lightcones drawn from semi-analytical models incorporating radio mode feedback as applied to the Millenium Simulation \\cite{Springel05a} and  compared predictions with various observables over $0<z<5$. They concluded broad agreement except in the abundance of high-$z$ galaxies with large stellar masses. In this paper, we will focus on how star formation is distributed within the evolving galaxy population.  We construct an ensemble of 20 {\\em lightcones} (\\S\\ref{Lightcones}) drawn from the {\\sc Galform} semi-analytic model \\cite{Bower06} as applied to the Millennium Simulation.  These lightcones are designed to mimic the Palomar/DEEP2 survey observations presented in \\scite{Bundy06} which we argue in \\S2 below is currently the best dataset for addressing the questions of differential mass assembly and evolution of star forming and quiescent galaxies.  These lightcones include the detailed survey geometry, optical and $K$-band magnitude limits, and photometric as well as stellar mass uncertainties described in \\scite{Bundy06}. Critically, by studying the variation across our 20 realizations, we are also able to determine reliable estimates of the effect of cosmic variance. As we will show, this is key to an accurate comparison between models and data.  A plan of the paper follows. In \\S2 we justify our choice of the DEEP2/ Palomar spectroscopic stellar mass catalog \\cite{Bundy06} as the comparison dataset. In \\S3 we discuss the Millenium Simulation and the associated {\\sc Galform} galaxy formation model. We review the various feedback mechanisms and illustrate how they lead, in principle, to the concept of downsizing. We then describe how we construct suitable lightcones and implement the effect of mass errors in the context of the observational data. We compare the predicted colour distribution with the observations and discuss the uncertainties associated with various ways of selecting active and quiescent galaxies. In \\S4 we compare our mass functions with the observations of \\scite{Bundy06}. We also discuss our findings with regard to how cosmic variance may limit such conclusions, not only in the context of the current survey \\cite{Bundy06} but also in other extant and projected surveys. ",
        "conclusions": "Lightcones derived from the Millenium Simulation, designed specifically to match the geometry and selection parameters of the DEEP2/Palomar survey of Bundy et al (2006), reproduce reasonably well the evolving stellar mass function over 0.4$<z<$1.2. These mock observations are populated with galaxies using the {\\sc Galform} code, which incorporates a prescription for `radio mode'  feedback. Such feedback processes are thus an adequate explanation of the broad trend of `downsizing' seen in recent redshift surveys.  However, the {\\sc Galform} code can not satisfactorily reproduce the evolving form of the bimodal colour distribution, or the stellar mass functions of the corresponding two populations. An alternative model \\cite{DeLucia07},  proposed originally by \\scite{Croton06}, produces similar discrepancies; both models appear to over-quench star formation in intermediate mass galaxies.  This approach also allows us to address the important question of how cosmic variance may limit the validity of various conclusions drawn from observational surveys. Using variance estimates derived from the lightcones, we confirm the significance of the decline since $z\\simeq$1 in the number density of massive blue galaxies claimed by Bundy et al (2006). We argue that the transformation of these blue galaxies must, necessarily, provide the bulk of the associated growth in red sequence galaxies given the near-constant total mass function.  This discussion of cosmic variance is extended to demonstrate the limitations of other, more ambitious, surveys ongoing at the present time in terms of detecting the {\\it mass-dependent} growth of galaxies since $z\\simeq$1."
    },
    "0806/0806.0371_arXiv.txt": {
        "abstract": "The X-ray transient 080109, associated with SN 2008D, can be attributed to the shock breakout emission from a normal Type Ib/c supernova. If the observed emission is interpreted as thermal emission, the temperature and radiated energy are close to expectations, considering that scattering dominates absorption processes so that spectrum formation occurs deep within the photosphere. The X-ray emission observed at $\\sim 10$ days is attributed to inverse Compton scattering of photospheric photons with relativistic electrons produced in the interaction of the supernova with the progenitor wind. A simple model for the optical/ultraviolet emission from shock breakout is developed and applied to SN 1987A, SN 1999ex, SN 2008D, and SN 2006aj, all of which have optical emission observed at $t\\sim1$ day. The emission from the first three can plausibly be attributed to shock breakout emission. The photospheric temperature is most sensitive to the radius of the progenitor star core and the radii in these cases are in line with expectations from stellar evolution. The early optical/ultraviolet observations of SN 2006aj cannot be accommodated by a nonrelativistic shock breakout model in a straightforward way. ",
        "introduction": "The shock wave generated by the central explosion in a supernova accelerates through the outer layers of the star, giving rise to the peak photospheric temperature at the time of shock breakout. The resulting X-ray/ultraviolet flash was predicted long ago \\citep{KC78,Fal78,MM99}. GRB 060218, discovered by its X-ray radiation \\citep{Cam06}, was associated with the fairly normal Type Ic SN 2006aj \\citep{Maz06}, and its early emission has been interpreted as due to shock breakout \\citep{Cam06,Wax07}. However, mildly relativistic motion is needed and it is likely that a central engine played a role in the emission.  The X-ray transient XRT 080109 was serendipitously discovered during {\\it Swift} observations of SN 2007uy in NGC 2770 \\citep{soder08}. The supernova associated with the burst, SN 2008D, was found to be spectroscopically similar to the Type Ib SN 2005hg \\citep{modjaz08}. Early radio emission was detected, but not at the high luminosity that would suggest relativistic motion, and no gamma ray burst emission was detected by {\\it Swift} \\citep{soder08}. In view of these properties, we further examine the possibility that the early ($\\la 1$ day) emission from SN 2008D was entirely due to the supernova. This view  was advocated by \\cite{soder08}, but \\cite{xu08} and \\cite{li08} find that a central engine flow is needed. We treat the X-ray emission in \\S~2 and the optical/ultraviolet emission in \\S~3. We take the distance to NGC 2770 to be 27 Mpc. ",
        "conclusions": ""
    },
    "0806/0806.3287_arXiv.txt": {
        "abstract": "We use the radiative transfer code \\phx\\ to study the line formation of the wavelength region 5000--7000~\\ang.  This is the region where the SNe~Ia defining Si~II feature occurs. This region is important since the ratio of the two nearby silicon lines has been shown to correlate with the absolute blue magnitude. We use a grid of LTE synthetic spectral models to investigate the formation of line features in the spectra of SNe~Ia. By isolating the main contributors to the spectral formation we show that the ions that drive the spectral ratio are \\FeIII, \\FeII, \\SiII\\ and \\SII. While the first two strongly dominate the flux transfer, the latter two form in the same physical region inside of the supernova. We also show that the na\\\"ive blackbody that one would derive from a fit to the observed spectrum is far different than the true underlying continuum. ",
        "introduction": "Type Ia \\supernovae\\ have been used as ``standardizable candles'' for more than 10 years, thanks to the correlation between their light curve shape and their absolute blue magnitude at maximum light. This correlation is well matched by a spectroscopic sequence \\citep{nugseq95} defined by the ratio, \\RSi, of the depth of two absorption features usually identified as \\SiII\\  \\SiIIblue\\ and \\SiIIred\\ lines. Using the radiative transfer code \\phoenix\\ \\citep[][and references therein]{hbjcam99,bhpar298,hbapara97,phhnovetal97,phhnovfe296}, this sequence has been matched to a temperature sequence that could physically be related to the amount of \\nni\\ produced in the context of Chandrasekhar-mass white dwarf explosions.  \\citet{hathighv00} showed that \\RSi\\ and the velocity derived from the minimum of the $\\approx 6100$~\\AA\\ feature correlated poorly, suggesting that a one-parameter description of \\SNeIa\\ was insufficient.  This velocity derived from the P-Cygni feature associated with the \\SiII\\ \\SiIIred\\ line, and its time evolution break \\SNeIa\\ into subclasses depending upon more than one parameter \\citep{benetti05}.  \\citet{branchcomp206} and \\citet{hach06} also showed that different spectroscopic indicators exist which can be used to distinguish variations even within the ``normal'' \\SNeIa.  In the context of Chandrasekhar-mass explosions, these variations cannot be interpreted only in terms of differences in the amount of \\nni\\ produced.  The velocity scatter could be due to properties of the progenitor system that impact on the kinetic energy released during the explosion. For example, the ratio ($^{54}$Fe+$^{58}$Ni)/$^{56}$Ni has been proposed as a second physical parameter which could account for this dispersion \\citep{mazzpod06}.  This paper is an attempt to unravel the spectral formation in the $5000-7000$~\\AA region (henceforth the \\RSi\\ wavelength region), as a way to address the physical reasons for their correlation with luminosity, and especially the link between this correlation and the temperature sequence of \\SNeIa. A previous attempt using  the radiative transfer code \\synow\\ \\citep{fisher00}, was performed by \\citet{garn99by04} who explained the counter-intuitive correlation of \\RSi\\ with luminosity by the temperature sensitivity of a \\TiII\\ line. On the other hand synthetic spectral investigations  \\citep{stehle02bo05,branchcomp206} failed to find any important \\TiII\\ line contributions in this part of the spectrum.   We use the radiative transfer code \\phx\\ differently than it has usually been used previously. Instead of fitting a supernova by varying the time and bolometric luminosity, we use a grid of simulated models to probe the line formation process. Our input model is the parametrized deflagration model W7 \\citep{nomw7}, homologously expanded to calculate the abundances at any given time.  The knowledge of the abundance structure of the explosion model, and the self-consistent physical structure provided by \\phoenix\\ are used to describe more precisely how the spectrum is formed in \\SNeIa.  \\citet{bongard06a} introduced the line ratio \\RSiS, which is defined as the ratio of two integrals of the flux over a chosen wavelength range, where each integral is centered around the two lines used in the original definition of \\RSi\\ (see \\S\\ref{rsisdef}). It is designed to be useful for spectra which have lower signal to noise than well-observed nearby SNe~Ia.  This understanding of the formation of the spectrum allows us to explain why the line ratio \\RSiS\\ is a good probe of the temperature sequence, and reinforces the link between this sequence and the absolute blue magnitude of \\SNeIa.  These results also allow us to rule out a \\TiII\\ contribution to the correlation of \\RSi\\ with luminosity.  We emphasize that even though the synthetic spectra presented here \\emph{do not} do a good job of reproducing observed SNe~Ia spectra, we consider W7 to be the fiducial SNe~Ia model. \\citet{bbbh06} showed that reasonably good synthetic spectra can be produced with W7 near maximum light. We make differential comparisons based on that model, keeping in mind that \\citet{bbbh06} emphasized that W7 did not reproduce the wavelength region that we examine in detail here.    We will address time evolution only qualitatively. The goal of this paper is to understand the temperature sequence of SNe~Ia spectra around maximum light, thus we concentrated on day 20 after explosion, since maximum light in $B$ in SNe~Ia has been shown to occur 18--25 days after explosion. ",
        "conclusions": "We used a grid of LTE \\phoenix\\ synthetic spectra using the \\nomw\\ model, to study the spectral formation in \\SNeIa\\ around the time of maximum light. Instead of focusing on fitting synthetic to observed spectra, we probed the detailed line formation in the synthetic spectrum to try to understand which ion was responsible for which feature, while taking detailed line blending into account.  We analyzed the synthetic spectra in detail to study the line formation in the \\RSi\\ wavelength region, showing that the \\SiIIred\\ peak and the 6100~\\AA\\ trough were dominated by \\SiII\\ lines forming over a blend of  \\FeII/\\FeIII\\ lines. The intermediate 5800~\\AA\\ peak has been shown to be a complex blend of  \\SiII\\ \\SiIIred, \\SiII\\ \\SiIIblue, \\FeIII, and \\FeII\\ weak lines. The \\TiII\\ contribution to the 5500~\\AA\\ trough has been ruled out. We showed that the redder edge was dominated by the \\SiIIblue\\  \\SiII\\ line whereas the bluer one was dominated by \\SiII\\ 5455~\\AA\\ line or \\FeII\\ 5170~\\AA\\ depending on luminosity. Based on our unraveling of the line formation in this wavelength region, we are able to illustrate our motivation for the definition of \\RSiS, isolating the \\SiII\\ and \\SII\\ contributions of the \\RSi\\ region.  We described a multi-layer model where the observed pseudo-continuum is formed throughout the entire supernova and thus strong features from multiple ionization stages can occur in the observed spectrum. We stress the importance of the numerous weak iron lines and show that their blends dominate the flux transfer. As a corollary, in the multi-layer model, the inner spectrum is not close to a Planck function but contains much more structure background. We showed that the ionization stage of the iron core dominates the \\SNIa\\ colors, explaining the brighter-bluer relation at maximum light. Even though a detailed study of time evolution is beyond the scope of this paper, we showed that our results where also qualitatively consistent in this regard, our earlier spectra being systematically much bluer than our later ones."
    },
    "0806/0806.4833_arXiv.txt": {
        "abstract": "The first astrophysical results of the VLTI focal instrument AMBER have shown the importance of the differential and closure phase measures, which are supposed to be much less sensitive to atmospheric and instrumental biases than the absolute visibility. However there are artifacts limiting the accuracy of these measures which can be substantially overcome by a specific calibration technique called Beam Commutation. This paper reports the observed accuracies on AMBER/VLTI phases in different modes, discusses some of the instrumental biases and shows the accuracy gain provided by Beam Commutation on the Differential Phase as well as on the Closure Phase. ",
        "introduction": "\\label{sec:intro}  %  In Optical long baseline interferometry, the phase of the object complex visibility is lost because of the random and fast phase variations introduced by the atmosphere.% Three main techniques exist to overcome this strong limitation and recover a part of the phase information: phase referencing\\cite{1998SPIE.3350..807Q}, where an off axis source in the isoplanetic angle is used as a phase reference, closure phase\\cite{1986Natur.320..595B, 2005AJ....130..246K}, which is a combination of the phases from three telescopes and is by nature free of any disturbance introduced in any of the beams, and then differential phase\\cite{1999ASPC..194...89A, 2006MNRAS.367..825V}, which is the same as phase referencing but with the source itself in some particular wavelength as a phase reference.  The phase referencing technique is affected by atmospheric effects, such as chromatic optical path difference (OPD), which is not exactly the same for the reference and the science target. Instrumental effects also affect the phase. Phase referencing technique uses a high precision metrology and a technique of commutation between the science star and the off axis star to overcome these effects. The OPD effects completely disappear in the phase closure, which explain why this phase-measurement technique has been used up to now with great success. However, for high dynamic and high precision measurements, closure phase remains sensitive to all non-OPD instrumental effects. The differential phase is affected by the same effects as the phase referencing technique. The instrumental effects can be reduced by a careful instrumental design, making the conditions of measure of all spectral channels (for example) as identical as possible, but they are almost impossible to eliminate completely. Therefore, a commutation technique would allow one, just like for phase referencing, to overcome in great part the instrumental effects.  To calibrate the differential and closure phases of AMBER, we installed a specific tool called the Beam Commutation Device (BCD). It has been inspired by our experience in differential speckle interferometry (Cuevas and Petrov, 1994) \\cite{1994SPIE.2200..501C} and consists in commuting two of the input beams without otherwise modifying the images pupils positions and OPD values seen by the instrument. The main advantage of such a commutation is that it can be performed very rapidly, to the contrary of usual calibration techniques.  This paper presents a first set of tests of this calibration method on the AMBER instrument and compares it results with those of a standard calibration using a reference star observed at a different time than the science source. ",
        "conclusions": "We have shown that the AMBER differential and closure phases display evolving patterns much larger than expected and specified. With a standard calibration using a reference star, it seems difficult to trust phase features smaller than 0.1 radians. This is confirmed by data on other stats and in other bands (MR-H and LR-H) that we could fit in this paper. Since our observations, we have learned that the main perturbation results from Perrot Fabry effects in various glass and air blades, both within AMBER and before it (IRIS feeding dichroic). The major cause for this perturbations has been proved to be the AMBER polarizing filters which will be soon replaced and relocated to minimize this wavelength beating effects.  The specific phase calibration function that we have developed improves the situation but quite not enough. With our Beam Commutation Device we can approach the 0.01 radian limit in the best cases, but this remains much larger than the fundamental noise limits, which are at least 10 times smaller with the ATs and 100 times smaller with the UTs.  The BCD clearly improves other sources of phase errors, more likely to be found in the calibration of the detector or of the position of the beams on the detector. This gives us good hope that the combination of Perrot Fabry effect suppression and of BCD calibration will allow us to approach the 0.001 radians limit, that we already approach in LR-K differential phase.  With the current hardware and without using the BCD, we had 0.1 radian (5 degrees) accuracy on phase, which was enough for the first science application on large disks ($\\alpha$ Arae cite{}) or on binaries with almost equal components ($\\gamma^{2}$ Velorum cite{}). With the BCD and a carefully data processing, we can approach the 0.01 radian specification sufficient for most circumstellar physics (disks and wind in stars of all ages) and by combining hardware and software upgrades with the BCD calibration we still confidently hope that the milliradian limit needed for extra solar planets and stellar activity work will be achieved."
    },
    "0806/0806.2138_arXiv.txt": {
        "abstract": "We have combined data from the NOAO Deep Wide-Field Survey in Bo\\\"{o}tes and the {\\em Spitzer Space Telescope} to determine basic properties  for sixteen optically ``invisible'' MIPS 24 $\\mu$m (OIMS) and thirty-five optically ``invisible'' radio (OIRS) sources, including their spectral energy distributions (SED) and luminosities. Most OIMSs possess steep power-law SEDs over $\\lambda_{\\rm rest} = 1-10 ~\\mu m$, indicating the presence of obscured AGN in agreement with {\\em Spitzer} spectroscopy. These objects are extremely luminous at rest-frame near and mid-IR ($\\nu L_{\\nu}(5 ~\\mu m) \\approx 10^{38}-10^{39}$ W), consistent with accretion near the Eddington limit and further implying that they host buried QSOs. The majority of the IRAC detected OIRSs have flat 3.6 to 24 $\\mu$m SEDs, implying comparable emission from stellar photospheres and hot AGN illuminated dust. This may reflect relatively small amounts of dust close to the central engine or current low mass accretion rates. A small subset of OIRSs appear to be starburst dominated with photometric redshifts from 1.0 to 4.5. The OIMSs and OIRSs with significant starburst components have similar L$_{\\rm K}$ and stellar masses ($M_* \\approx 10^{11}$ M$_{\\odot}$) assuming minimal AGN contribution.  Roughly half of the OIRSs are not detected by {\\it Spitzer's} IRAC or MIPS. These are most likely $z \\ga 2$ radio galaxies.  The IRAC detected OIRSs are more likely than OIMSs to appear non point-like in the 3.6 $\\mu m$ and 4.5 $\\mu$m images, suggesting that interactions play a role in triggering their activity. The AGN powered OIMSs may represent sub-millimeter galaxies making the transition from starburst to accretion dominance in their evolution to current epoch massive ellipticals. ",
        "introduction": "Sensitive multi-band surveys in the optical and near-infrared portion of the electromagnetic spectrum have proven extremely successful in identifying large and diverse galaxy populations in the high-redshift universe. Deep images obtained through optical {\\it U$_n$GR} filters, for example, readily find $z \\ga 3$ galaxies by the presence of a strong 912 \\AA ~Lyman break (e.g., Steidel et al. 1996).  Lyman Break Galaxies (LBGs) tend to be massive galaxies with star formation rates (SFR) between $\\approx 4-25$ M$_{\\odot}$ yr$^{-1}$ and relatively low dust obscuration.  Similarly, unobscured star forming disk (S-BzK galaxies) and passively evolving spheroidal galaxies at $z \\sim 1.4 - 3$ can be discerned from the myriad background sources by their optical/near-infrared colors or other magnitude-color criteria (see, e.g., Daddi et al. 2004; Adelberger et al. 2004). The S-BzK galaxies in particular tend to possess large stellar masses ($M_* \\approx 10^{11}$ M$_{\\odot}$) with SFR $\\approx 200$ M$_{\\odot}$ yr$^{-1}$. Nevertheless, it had been long suspected that a substantial fraction of luminous $z > 1$ galaxies powered by star formation or accretion onto super-massive black holes would not be identified using these and related techniques because of heavy dust obscuration, but would turn up instead at infrared, sub-millimeter, and radio wavelengths (see, e.g., Chary \\& Elbaz 2001 and references within). Indeed, deep but small-area surveys at sub-millimeter wavelengths by SCUBA (e.g., Smail et al. 1997) revealed a population of extremely dusty galaxies at $z \\approx 1-3$, whose high luminosities are generally though to originate primarily in star formation (SFR $\\ga 1000$ M$_{\\odot}$ yr$^{-1}$), though a small fraction may be AGN dominated. Galaxies powered by starbursts or active nuclei (AGN) can also be strong emitters at radio wavelengths through either large populations of  supernova remnants or accretion disk phenomena. Even though this emission is unaffected by obscuration, roughly $10-15\\%$ of compact radio sources identified in deep radio surveys have either extremely faint optical counterparts or none at all (e.g., Richards et al. 1999; Fomalont et al. 2002), and may represent distant luminous obscured starburst or AGN powered galaxies.  Such objects make a still unknown contribution to the luminosity, chemical, and accretion history of the universe. Determining an accurate census of heavily obscured sources - their redshifts, luminosities, space densities, and dominant power source - is one of the key issues in observational astrophysics and a primary objective of the {\\it Spitzer Space Telescope} ({\\em Spitzer}; Werner et al. 2004).\\footnotemark[1] \\footnotetext[1]{The {\\em Spitzer Space Telescope} is operated by JPL, California Institute of Technology for the National Aeronautics and Space Administration. Information on {\\em Spitzer} can be found at http:$//$ssc.spitzer.caltech.edu$/$.}   A number of ambitious large area multi-wavelength surveys incorporating {\\it Spitzer} data have been conducted to date. These include the {\\it Great Observatories Origins Deep Survey} (GOODS), the {\\it Spitzer Wide-area InfraRed Extra-galactic} survey (SWIRE), and the {\\it First Look Survey} (FLS). This paper is concerned with the nature of infrared and radio selected sources with extremely faint optical counterparts originally selected from the {\\it  NOAO Deep Wide-Field Survey} (NDWFS) in Bo\\\"{o}tes (Jannuzi \\& Dey 1999)\\footnotemark[2] with the expectation that they are either heavily extincted, at high redshift, or both. \\footnotetext[2]{The NOAO Deep Wide-Field Survey is supported by the National Optical Astronomy Observatory, which is operated by AURA, Inc., under a cooperative agreement with the National Science Foundation. Information on the NDWFS can be found at http:$//$www.noao.edu$/$noao$/$noaodeep$/$.} Higdon et al. (2005; hereafter, Higdon05) identified thirty-six {\\em Optically ``Invisible'' Radio Sources} (hereafter, OIRSs) out of 377 compact or unresolved radio sources found in a VLA A-array 20 cm survey covering 0.5 deg$^{2}$ in the NDWFS Bo\\\"{o}tes field. These observations reached a  flux density limit of $\\approx 80$ $\\mu$Jy ($5 \\sigma$) at the three overlapping pointing centers. To be considered optically ``invisible'', a radio source must have no $B_W$, $R$, or $I$ counterpart within a 1.5\\arcsec~ radius. This corresponds to limiting $3 \\sigma$ magnitudes of approximately 26.9, 25.6, and 24.9 (Vega; 2\\arcsec ~aperture) respectively, with the precise value depending upon location within the optical survey. The NDWFS region was surveyed at 24 $\\mu$m and 70 $\\mu$m with the Multiband Imaging Photometer for {\\em Spitzer} (MIPS, Rieke et al. 2004). Analysis of the OIRS's {\\it q}-parameter led Higdon05 to conclude that they are a population powered by relatively unobscured radio-loud active nuclei.\\footnotemark[3] While none of the OIRSs have measured redshifts, Higdon05 argued that they are likely to be at $z > 1$ based on the faint optical limits set by the NDWFS survey. \\footnotetext[3]{{\\it q~}$ \\equiv log(F_{24 \\mu m}/F_{20 cm})$. Systems powered by star formation or radio-quiet active nuclei possess {\\it q} $= 0.5-1.1$.  Smaller, and in particular, negative values of {\\it q} indicate the presence of increasingly radio-loud active nuclei (Appleton et al. 2004).} An independently defined sample of seven optically faint ($R \\ge$ 24.5) and ten optically ``invisible'' sources from the NDWFS Bo\\\"{o}tes region, subject to the additional constraints that $F_{\\rm 24 \\mu m} >$ 0.75 mJy and $\\nu f_{\\nu}(24 \\mu m)/ \\nu f_{\\nu}(0.8 \\mu m) \\ga$ 100, was observed using {\\it Spitzer's} Infrared Spectrometer (IRS, Houck et al. 2004). These {\\em Optically ``Invisible'' MIPS Sources} (hereafter, OIMSs) are a high-z population, with $z \\sim 1.6-2.7$. Comparisons with mid-IR spectra of local starburst and AGN dominated galaxies suggested that the primary energy source in thirteen of the seventeen OIMSs is a heavily obscured active nucleus (Houck et al. 2005, hereafter Houck05). All seventeen have inferred $8 - 1000 ~\\mu m$ luminosities of $\\approx$10$^{13}$~ L$_{\\odot}$, placing them in the ``hyper''-luminous class.    It is not known how the radio and infrared selected ``invisible'' populations in Higdon05 and Houck05 are related, apart from the fact that the majority of both appear to possess powerful AGN. We do not know, for example, if the OIRS population lies at systematically higher redshifts than the OIMSs with similar L$_{\\rm IR}$, or for that matter, if they are a sub-L$^{*}$ population at $z \\sim 0.5$. Key to addressing these issues is estimating photometric redshifts for the OIRSs, which are generally too faint for {\\em Spitzer's} IRS. In this paper, we present results obtained by combining optical $B_W$, $R$, and $I$, and MIPS 24 $\\mu$m, 70 $\\mu$m,  and 160 $\\mu$m flux densities (or limits) with data from {\\em Spitzer's} Infrared Array Camera (IRAC, Fazio et al. 2004) Shallow Survey of the NDWFS Bo\\\"{o}tes field (Eisenhardt et al. 2004). The IRAC Shallow Survey covers 8.5 deg$^{2}$ in four bands centered at 3.6, 4.5, 5.8, \\& 8.0 $\\mu$m. For $z \\approx 1-3$, these bands measure emission in the rest-frame near-infrared.  Because this wavelength regime is much less sensitive to extinction relative to the optical, the IRAC data can constrain the evolved stellar content of galaxies at this epoch. Photometric redshifts are also possible by virtue of the rest-frame 1.6 $\\mu$m peak arising from the H$^{-}$ opacity minimum in the photospheres of evolved stars. It is our aim to derive basic properties for both optically ``invisible'' source populations, including rest-frame near-infrared luminosities and SEDs, and to determine, for example, if their rest-frame near-infrared emission is dominated by starlight or an AGN. Only then would it be possible to relate OIMSs and OIRSs to other high-z populations. It is also our intent to use the IRAC images to explore the near environments of OIMSs and OIRSs for additional clues to their nature and evolution.  A brief description of the extraction of IRAC flux densities for the infrared and radio selected ``invisible'' sources  is given in $\\S$2. Photometric redshift estimates for selected OIRSs are presented in $\\S$3, along with determinations of the average rest-frame SEDs,  luminosities, and near environments of the OIMS and  OIRS samples detected with IRAC. The nature of these two populations is discussed in $\\S$4. These results are summarized in $\\S5$. Throughout this paper we will refer to the OIRSs and OIMSs by number (i.e., OIRS \\#97 or OIMS \\#13), corresponding to their entries in Table  2 of Higdon05 and Table 1 of Houck05. Source coordinates can be found from those papers directly, or through their SIMBAD designations, [HHW2005]~\\# and [HSW2005]~\\# for the OIRSs and OIMSs respectively. We assume a flat $\\Lambda$CDM cosmology, with $\\Omega_{\\rm M}$=0.27, $\\Omega_{\\rm \\Lambda}$=0.73, and a Hubble constant of 71 km s$^{-1}$ Mpc$^{-1}$. ",
        "conclusions": "OIMSs can be split into two populations based on their rest-frame 1-10 $\\mu$m SEDs and IRAC colors in concordance with their IRS spectra: most (12/16) are dominated by a heavily obscured radio-quiet active nucleus, while the remainder are powered by either a starburst or a composite starburst/active nucleus. The AGN dominated {\\it obsSy1} and {\\it obsSy2} OIMSs in particular are extremely luminous in the rest-frame near and mid-IR, with $\\nu$L$_{\\nu}$(5 $\\mu$m) comparable to the most luminous local NLRGs and quasars. They can be regarded as ``buried'' QSOs and likely represent the predecesors of current epoch massive elliptical galaxies. OIMSs are also distinct from other high-z source populations routinely selected using UV/optical or optical/near-infrared criteria such as LBGs and BzKs in their dominant power source, levels of obscuration, and mid-IR luminosity. Compared with other optically faint mid-IR selected populations (e.g., Yan et al. 2007), OIMSs represent extremes in both obscuration and AGN luminosity. This follows from their respecive selection criteria, which for OIMSs favor steeper mid-IR continuua and higher levels of extinction. There appears to be significant overlap between OIMSs and SMGs, with the {\\it sb/comb} OIMSs appearing largely indistinguishable from SMGs in their optical, mid-infrared, and radio properties. Moreover, the {\\it obsSy1} and {\\it obsSy2} OIMSs may represent a brief obscured phase in the transition of a ``cold'' SMG to a QSO, and eventually to a massive current epoch elliptical galaxy.  The OIRSs do not represent a single source population. The minority (6/35) that are detected by both IRAC and MIPS at 24 $\\mu$m have SEDs, mid-IR colors, and {\\it q}-values indicative of either starburst or composite starburst/AGN powered systems. For these sources we find $z_{\\rm phot}$ in between 1.0 and 4.5, implying rest-frame K-band luminosities, maximum star formation rates, and stellar masses virtually identical to the ``cold'' (i.e., starburst dominated) SMGs in the Lockman Hole East region.   The remaining OIRSs - which comprise $83\\%$ of the parent population in Higdon05 - fall into two classes depending upon whether or not they are detected by IRAC. Those that are detected have flat mid-IR SEDs implying comparable luminosities from stellar photospheres and hot AGN illuminated dust. As a group, their average L$_{\\rm 24 \\mu m}$/L$_{\\rm 3.6 \\mu m}$ ratio is most consistent with either Low Excitation Radio Galaxies (LERGs) or objects like the ``red and dead'' radio galaxy LBDS 53W091. Those that are not detected by IRAC must be at $z \\ga 2$ if they are as massive as the hosts of local radio sources. Both of these groups are characterized by negative values of {\\it q} and thus may represent a population of {\\it relatively} unobscured radio galaxies at high redshift. Both represent populations that are highly distinct from the {\\it obsSy1} and {\\it obsSy2} OIMSs.  Differences between the optically {\\it invisible} populations detected through observations at sub-millimeter, mid-IR and radio wavelengths can be understood in terms of selection effects: (1) extreme optical/mid-IR luminosity ratios appears to ensure highly obscured AGN dominated sources ({\\it obsSy1} and {\\it obsSy2} OIMSs), (2) extreme optical/sub-millimeter luminosity ratios will select highly obscured sources primarily powered by star formation (SCUBA/MAMBO sources), and (3) compact sub-mJy radio sources lacking optical counterparts appears to preferentially choose distant BL Lac-like objects (OIRSs).  We find no evidence that OIMSs or OIRSs inhabit the cores of rich clusters. Nor do we find significant differences in local galaxy density between the two on $\\la 100$ kpc scales. However, unlike the OIMSs, a large fraction of the IRAC detected OIRSs appear to possess close and massive companions, though higher angular resolution studies will be needed to reach firm conclusions. This suggests that the luminosity of OIRSs (and conceivably OIMSs) may be triggered by tidal interactions, as appears to be the case for low-$z$ radio galaxies and quasars. The IRAC detected OIRSs may thus represent the formation of very massive galaxies at high redshift through major mergers."
    },
    "0806/0806.0417_arXiv.txt": {
        "abstract": "Asteroseismology, as a tool to use the indirect information contained in stellar oscillations to probe the stellar interiors, is an active field of research presently. Stellar age, as a fundamental property of star apart from its mass, is most difficult to estimate. In addition, the estimating of stellar age can provide the chance to study the time evolution of astronomical phenomena. In our poster, we summarize our previous work and further present a method to determine age of low-mass main-sequence star. ",
        "introduction": "Due to the frequencies of these oscillations depend on density, temperature, gas motion, and other properties of the stellar interior, it can take the window to ``see\" the interior of stars and help us to know the stellar internal structure and understand the stellar evolution. With the advance of observational technique, % several stars have been detected the solar-like oscillations. Using the latest asteroseismic data, we reconstruct the model of $\\alpha$ Cen B and 70  Ophiuchi A (Tang et al. 2008a, 2008b). In additional, \\cite{Bi} have performed preliminary seismological analysis of two MOST targets. % ",
        "conclusions": "1. The ($\\langle\\Delta\\nu\\rangle$, $r_{01}$) diagram as a new asteroseismic diagnostic tool can estimate the mass and the age of solar-like stars. The virtue is that the age of stars can be marked in the diagram, so we can obtain the mass and age directly.  2. We will discuss the effects of the assumed initial abundance of helium and the mixing length parameter on the asteroseismic diagram in future work. \\begin{figure} \\includegraphics[height=1.5in,width=2.5in,angle=0]{fig1.eps}% \\includegraphics[height=1.5in,width=2.5in,angle=0]{fig2.eps} \\caption{[a]: The ratio of small separations adjacent in $l$ vs. age for each of 129 stellar models described in \\cite{Tanga}. [b]: ($\\langle\\Delta\\nu\\rangle$, $r_{01}$) diagram for stellar models. The vertical lines are evolutionary tracks, labeled by the mass in the top, whereas the transverse lines are isopleths with constant age, labeled by the age from 0.5 Gyr to 7.0 Gyr increasing with 0.5 (unit is Gyr )}\\label{fig:contour} \\end{figure}"
    },
    "0806/0806.0621_arXiv.txt": {
        "abstract": "{} % {We analyse the properties of the early-type dwarf galaxy population ($M_V>-17$ mag) in the Hydra\\,I cluster. We investigate the galaxy luminosity function (LF), the colour--magnitude relation (CMR), and the magnitude--surface brightness relation down to $M_V\\sim-10$ mag. Another goal of this study is to find candidates for ultra-compact dwarf galaxies (UCDs) in Hydra\\,I.} {Two spectroscopic surveys performed with Magellan~I/LDSS2 at Las Campanas Observatory and VLT/VIMOS, as well as deep VLT/FORS1 images in $V$ and $I$ bands, covering the central parts of the cluster, were examined. We identify cluster members by radial velocity measurements and select other cluster galaxy candidates by their morphology and low surface brightness. The candidates' total magnitudes and central surface brightnesses were derived from the analysis of their surface brightness profiles. To determine the faint-end slope of the LF, the galaxy number counts are completeness corrected.} {We obtain radial velocities for 126 objects and identify 32 cluster members, of which 5 are previously uncatalogued dwarf galaxies. One possible UCD candidate with $M_V=-13.26$ mag is found. Our sample of $\\simeq 100$ morphologically selected dwarf galaxies with $M_V>-17$ mag defines a CMR that extends the CMR of the giant cluster galaxies to the magnitude limit of our survey ($M_V\\sim-10$ mag). It matches the relations found for the Local Group (LG) and the Fornax cluster dwarf galaxies almost perfectly. The Hydra\\,I dwarf galaxies also follow a magnitude--surface brightness relation that is very similar to that of the LG dwarf galaxies. Moreover, we observe a continuous relation for dwarf galaxies and giant early-type galaxies when plotting the central surface brightness $\\mu_0$ of a S\\'ersic model vs. the galaxy magnitude. The effective radius is found to be largely independent of the luminosity for $M_V>-18$ mag. It is consistent with a constant value of $R_{\\mathrm{e}}\\sim 0.8$ kpc. We present the photometric parameters of the galaxies as the Hydra\\,I Cluster Catalogue (HCC). By fitting a Schechter function to the luminosity distribution, we derive a very flat faint-end slope of the LF ($\\alpha = -1.13 \\pm 0.04$), whereas fitting a power law for $M_V>-14$ mag gives $\\alpha = -1.40 \\pm 0.18$.} {Our findings of a continuous CMR and $\\mu_0$ -- $M_V$ relation for dwarf and giant early-type galaxies suggests that they are the same class of objects. The similarity of those relations to other environments like the LG implies that internal processes could be more important for their global photometric properties than external influences.} ",
        "introduction": "Dwarf galaxies are the most abundant type of galaxy in the universe. They are most commonly found in galaxy clusters, such as Virgo, Coma and Fornax \\citep[e.g.][]{Sandage1984, Binggeli1985, Ferguson1988, Secker1996, Roberts2007}. Also in the Local Group (LG) a large number of dwarf galaxies have been identified \\citep[e.g.][and references therein]{Mateo1998, vandenBergh1999, vandenBergh2000, Grebel2003}.  The classification of dwarf galaxies is not standardised in the literature. They are usually distinguished from giant elliptical and spiral galaxies by their low luminosities and low surface brightnesses and, just as for giant galaxies, one refers to early-type and late-type dwarf galaxies (dwarf irregulars). In this paper we focus on \\emph{early-type} dwarf galaxies, comprising dwarf elliptical galaxies (dEs) and dwarf spheroidal galaxies (dSphs). In \\citet{Grebel2001} dEs are defined as objects with low luminosities ($M_V\\gtrsim -17$ mag) and typical central surface brightnesses of $\\mu_V \\lesssim 21$ mag arcsec$^{-2}$. Dwarf spheroidal galaxies have even lower luminosities ($M_V\\gtrsim-14$ mag) and central surface brightnesses ($\\mu_V \\gtrsim 22$ mag arcsec$^{-2}$). Unless stated otherwise, we use the term dwarf galaxy to refer to both types (dEs and dSphs).  Another probable type of early-type dwarf galaxies has been identified in the nearby galaxy clusters Fornax, Virgo and Centaurus - the so-called ultra-compact dwarf galaxies (UCDs) \\citep[e.g.][]{Hilker1999b, Drinkwater2000, Hasegan2005, Jones2006, Mieske2007b}. UCDs are of intermediate nature between dwarf elliptical galaxies and globular clusters in terms of their morphology, being characterised by sizes of $10<r_{\\mathrm{eff}}<100$ pc and luminosities between $-13.5<M_V<-11$ mag. The luminosities of the brightest UCDs in Virgo and Fornax are comparable to those of nuclei of dwarf ellipticals and late-type spirals. In order to understand the importance of dwarf galaxies for the evolution of larger galaxies and galaxy clusters, it is essential to investigate their properties in different environments.  Current hierarchical cold dark matter (CDM) models of galaxy formation claim that dwarf-sized galactic fragments were very abundant in the early universe. They represent the building-blocks of larger galaxies. The models predict that Milky Way-sized dark matter haloes as well as cluster-sized haloes contain a large number of dark matter subhaloes \\citep{Moore1999}. If every small DM subhalo contained luminous matter, hundreds of dwarf galaxies are expected to be observed in the LG, but already there the number of observed satellite galaxies is too low by about one order of magnitude. This discrepancy is known as the \\emph{missing satellites problem} \\citep{Klypin1999} and its origin is still a matter of debate. Either the predictions of the hierarchical models are not reliable, or there is a large number of luminous satellites being not yet discovered, or a major fraction of low-mass DM haloes has not formed any stars. In fact, the analysis of SDSS data recently led to the discovery of several extremely faint ($M_V\\gtrsim -8$ mag) LG dwarf galaxies \\citep[e.g.][]{Willman2005a, Willman2005b, Zucker2006a, Zucker2006b, Belokurov2006, Irwin2007, Simon2007}, proving that such hardly detectable objects do exist. But they are still not abundant enough to explain the expected DM sub-structure.  \\subsection{The galaxy luminosity function} A well known way to quantify the discrepancy between the number of observed and the number of predicted dwarf galaxies is the determination of the galaxy luminosity function (LF). An analytic function for the LF was proposed by \\citet{Schechter1976} as \\[ \\phi(L)dL =\\phi^* \\left(\\frac{L}{L^*}\\right)^{\\alpha} e^{-(L/L^*)} d\\left(\\frac{L}{L^*}\\right).\\] The logarithmic faint-end slope $\\alpha$ can be contrasted with the predicted slope of $\\simeq-2$ for the mass spectrum of cosmological dark-matter haloes \\citep[e.g.][]{Press1974, Moore1999, Jenkins2001}. Inconsistent with the predictions, rather flat slopes of $-1.0 \\lesssim \\alpha \\lesssim -1.5$ have been derived from studies of various environments ranging from the LG to the Coma cluster \\citep[e.g.][]{vandenBergh1992, Trentham2002, CZ2003, Hilker2003, Mobasher2003, Trentham2005, Chiboucas2006, Mieske2007, Rines2007, Penny2008}. Other authors, however, reported on steeper slopes \\citep[e.g.][and references therein]{Sabatini2003, Milne2007}. The LF of the Hydra\\,I cluster has previously been determined by \\citet{Yamanoi2007} who find $\\alpha \\sim -1.6$ in the magnitude range $-20<M_{B,R}<-10$. This is a steeper slope than reported by \\citet{Yagi2002} who give $\\alpha=-1.31$ for the faint end slope of the composite LF of 10 nearby clusters (including Hydra\\,I) at $-23<M_R<-16$. \\citet{CZ2003} find $\\alpha=-1.21$ for the composite LFs of six clusters (also including Hydra\\,I) at $-22<M_R<-14$, derived from deep spectroscopic samples.    \\subsection{Photometric scaling relations for early-type galaxies} Correlations among global parameters of early-type galaxies can provide insight into the physical processes that have impact on the formation mechanisms and the evolution of these galaxies. For example, luminosity, colour, surface brightness, central velocity dispersion and the $\\mathrm{Mg}_2$ absorption line index are related to each other \\citep[e.g.][]{Faber1976, Kormendy1977, Kormendy1985, Djorgovski1987, Ferguson1988, Bender1992, Bernardi2003b, Bernardi2003c, Bernardi2003d, Chang2006}. In particular, the colour--magnitude relation and the magnitude--surface brightness relation connect the physical properties of the underlying stellar population and the structural properties with the galaxy masses. Investigating those scaling relations in multiple environments sets constraints for galaxy formation and evolutionary models of early-type galaxies.  \\subsubsection{Colour--magnitude relation} \\label{sec:cmrintro} A tight colour--magnitude relation (CMR) for early-type cluster galaxies has been known for a long time \\citep{Visvanathan1977}. It is most commonly explained by an increase of the mean stellar metallicity with increasing galaxy mass, leading to redder colours of the more luminous galaxies \\citep[e.g.][]{Kodama1997, Ferreras1999}. With a large sample from the Sloan Digital Sky Survey, \\citet{Gallazzi2006} confirmed this perception. The metallicity of a galaxy strongly depends on the fraction of gas that has been turned into stars. Due to their deeper potential well, massive galaxies are able to retain their interstellar gas and stellar ejecta longer and more effectively than low-mass galaxies since the escape velocity of stellar yields depends on the galaxy mass. Hence, subsequent generations of stars will be formed out of already enriched gas. Because giant elliptical galaxies with their high star formation rate (SFR) consume their gas very fast, the mean stellar metallicity can reach high values in a short time. That accounts for the reddest colours of the most luminous galaxies.  \\citet{Koeppen2007} demonstrated that a variable integrated galactic initial mass function (IGIMF) that depends on the SFR can also explain the observed mass--metallicity relation. For a low SFR the IGIMF differs from the standard IMF in the sense that fewer high mass stars are expected to form. It has been found that galaxies with a low current SFR contain star clusters with a lower maximum mass \\citep{Weidner2004}, and less massive star clusters less likely contain very massive stars \\citep{Kroupa2003}. Since less massive galaxies are expected to have lower star formation rates, the effective upper mass limit for stars in such galaxies is lower. This reduces the number of type II supernovae as the main source of $\\alpha$-elements. Hence, the dependence of the IGIMF on the SFR implies a dependence of the metal abundance on the mass of the galaxy.  As an alternative to the aforementioned scenarios, \\citet{Worthey1994} and \\citet{Poggianti2001} suggested that a change in the mean age of a stellar population could at least in part explain the CMR. A stellar population will gradually redden as stars with increasing age evolve off the main sequence towards the red giant branch. In this picture, redder colours of more massive galaxies imply systematically older ages. Observational evidence for this was given by \\citet{Rakos2004}, who reported on younger ages of dwarf elliptical galaxies in the Coma and Fornax clusters. Hence, an increase of age at fixed metallicity has the same effect on galaxy colours as an increase of metallicity at fixed age. This ambiguous interpretation of the CMR is based on the well-known \\emph{age-metallicity degeneracy} of integrated optical colours.  \\subsubsection{Magnitude--surface brightness relation} It has been known for decades that elliptical galaxies follow a relation between effective radius and absolute magnitude \\citep{Kormendy1977}. Dwarf elliptical galaxies, on the other hand, follow a $r_{\\mathrm{eff}}$-$M_V$ relation unequal in slope to that of brighter giant ellipticals \\citep{Bender1992}. It was believed that a dichotomy between dwarf elliptical and giant elliptical galaxies also exists in the magnitude--surface brightness plane (based on exponential and de Vaucouleurs fits). According to this, dwarf galaxies follow a relation in the sense that brighter dEs have higher surface brightnesses, whereas the surface brightness of giant elliptical galaxies decreases with increasing luminosity \\citep[e.g.][]{Kormendy1985, Ferguson1988, Bender1992}.  \\citet{Graham2003}, however, showed that there does not exist a dichotomy when relating the magnitude $M$, the profile shape index $n$ and the central surface brightness $\\mu_0$ of \\citet{Sersic1968} profile fits with each other (except maybe for the most luminous giant ellipticals that deviate from a smooth relation due to their low surface brightness cores). In this study, the interpretation of different galaxy formation mechanisms between giant and dwarf galaxies is not supported.  \\citet{Phillipps1988} and \\citet{Irwin1990} argued against the existence of a magnitude--surface brightness relation for dEs. They claimed that the relation in eye-selected samples of dwarf galaxies is merely a result of selection effects, but \\citet{Karick2003} and \\citet{Mieske2007} observed such a relation for Fornax dEs for which cluster membership was assured either by an un-biased radial velocity survey (Karick) or by SBF measurements (Mieske).   \\subsection{The Hydra\\,I galaxy cluster (Abell 1060)} \\begin{figure} \\resizebox{\\hsize}{!}{\\includegraphics{hydracluster9_final_small.eps}} \\caption{$45'\\times 45'$ ($540 \\times 540$ kpc at the cluster distance) image of the Hydra\\,I cluster centred on NGC 3311, extracted from the Digital Sky Survey. The small squares are the fields observed with LDSS2. The large square marks the VIMOS-pointing. The circle indicates the cluster core radius $r_{\\mathrm{c}}=170\\ h^{-1}$ kpc \\citep{Girardi1995}, adopting $h=0.75$.} \\label{fig:hydracluster} \\end{figure}   Hydra\\,I is a relatively poor cluster [richness class $R=1$ \\citep{Abell1958}, BM classification III \\citep{Bautz1970}]. A pair of bright galaxies is located near the cluster centre. NGC 3309 is a regular giant elliptical galaxy (E3). The brighter and larger galaxy NGC 3311 possesses an extended cD halo and an extraordinarily rich globular cluster system \\citep{vandenBergh1977, McLaughlin1995, Mieske2005, Wehner2008}.  The cluster is the prototype of an evolved and dynamically relaxed cluster, being dominated by early-type galaxies and having a regular core shape. From X-ray measurements, \\citet{Tamura2000} derived  an isothermal distribution of the intracluster medium within $\\sim 160\\ h^{-1}\\ \\mathrm{kpc}$. They give $2.1 \\times 10^{14}\\ h^{-1}\\ \\mathrm{M_{\\sun}}$ as the cluster virial mass. From optical studies, \\citet{Girardi1995} found a core radius of $r_{\\mathrm{c}} = 170\\ h^{-1}\\ \\mathrm{kpc}$. Applying the virial theorem to the member galaxies, \\citet{Girardi1998} calculated a virial mass of $1.9 \\times 10^{14}\\ h^{-1}\\ \\mathrm{M_{\\sun}}$.  \\citet{Mieske2005} estimated the distance to Hydra\\,I from $I$-band surface brightness fluctuations (SBF). They found the distance to be $41.2 \\pm 1.4\\ \\mathrm{Mpc}$ (distance modulus $(m-M)=33.07 \\pm 0.07\\ \\mathrm{mag}$). This is at the low end of distance estimates by other authors whose mean is $\\sim 15$\\% higher \\citep[see discussion in][]{Mieske2005}. From a deep spectroscopic sample of cluster galaxies, extending to $M_R=-14\\ \\mathrm{mag}$, \\citet{CZ2003}, hereafter CZ03, derived $\\overline{cz}=3683\\pm46\\ \\mathrm{km\\ s^{-1}}$ as the mean cluster redshift with a velocity dispersion of $\\sigma = 724\\pm 31\\ \\mathrm{km\\ s^{-1}}$. This corresponds to a distance of $51.2 \\pm 5.7\\ \\mathrm{Mpc}$, assuming $H_0=72\\pm 8\\ \\mathrm{km\\ s^{-1}\\ Mpc^{-1}}$ \\citep{Freedman2001}. Already earlier, it was discovered that the cluster is clearly isolated in redshift space, having no foreground galaxies and no background galaxies up to $cz\\sim 8000\\ \\mathrm{km\\ s^{-1}}$ \\citep{Richter1982, Richter1987}. This implies huge empty regions of space of about 50 Mpc path length in front and behind the cluster.  The Hydra\\,I cluster is close enough for current 8-m class telescopes like the VLT to resolve faint dwarf galaxies under good seeing conditions. Only a few images of about $7\\times7$ arcmin (the typical field-of-view size for most CCD cameras) are needed to observe the cluster centre as well as areas out to about one core radius (see Fig. \\ref{fig:hydracluster}).  In this paper we present a study of the early-type dwarf galaxy population in the Hydra\\,I cluster. The study is based on two spectroscopic surveys and a photometric survey. One spectroscopic survey was performed with LDSS2 at Magellan~I and aimed at searching for normal cluster dwarf galaxies. The other spectroscopic survey -- performed with VLT/VIMOS -- focused on the search for UCDs. The photometric survey, based on deep VLT/FORS1 images, was used to identify other cluster dE/dSph candidates. The paper is organised as follows. In Sect. \\ref{sec:spectroscopy} we describe the observations, the candidate selection and the data reduction for the two spectroscopic surveys. In Sect. \\ref{sec:photometry} the photometric analysis of the candidate cluster dEs/dSphs, as selected from the VLT/FORS1 images, is addressed. Our results are presented in Sect. \\ref{sec:results}. We summarise and discuss our findings in Sect. \\ref{sec:discussion}. ",
        "conclusions": "\\label{sec:discussion} In this paper we presented a spectroscopic and photometric study of the early-type dwarf galaxy population in the Hydra\\,I cluster. Two spectroscopic surveys were analysed, one executed with Magellan~I/LDSS2 at Las Campanas Observatory, the other with VLT/VIMOS at ESO/Paranal. The imaging data were obtained with VLT/FORS1. Seven fields were observed in Johnson $V$ and $I$ filters with a seeing between $0.5''$ and $0.7''$.  \\subsection{Massive UCDs in Hydra\\,I?} From the two spectroscopic data sets we identify three compact cluster members being candidates for bright GCs and/or UCDs. Our rather incomplete VIMOS data set reveals the existence of one UCD candidate with $M_V=-13.26$ mag, corresponding to a mass of $\\sim 10^8\\ \\mathrm{M_{\\sun}}$. However, due to its comparatively low S/N spectrum its cluster membership has to be reconfirmed by follow-up observations.  This \\emph{possible} detection is intriguing, since such sources are extremely rare in other environments. The complete spectroscopic surveys for Fornax and Virgo have revealed only one single such bright UCD in either cluster \\citep{Jones2006}, with all other UCDs being at least one magnitude fainter. Does the possible detection of this bright UCD point to a large number of massive UCDs in Hydra\\,I? The cluster centre is indeed dominated by the cD galaxy NGC 3311 with its very pronounced diffuse light component, which implies that the cluster environment in Hydra\\,I could be more important in shaping the central region than in Fornax or Virgo. A large number of massive UCDs would therefore strengthen the idea of environment dependent UCD formation processes, such as the dynamical transformation of dwarf galaxies \\citep{Bekki2003, Goerdt2008}. Further insight into the nature of the possible UCD candidate -- for example its mass-to-light ratio and size -- will require space based imaging and high resolution spectroscopy. With only one UCD candidate, it is not possible to make a meaningful statement about the abundance and properties of UCDs in the Hydra\\,I cluster. For this, extended spectroscopic surveys are required.    \\subsection{Global photometric properties of dwarf galaxies} By radial velocity measurements we confirm the cluster membership of 29 galaxies, of which 5 are previously uncatalogued dwarf galaxies. The confirmed cluster galaxies define a tight colour--magnitude relation (Fig. \\ref{fig:ldssprops}b). From the visual inspection of the images, we identify 82 additional cluster dE/dSph candidates. Our sample of $\\simeq 100$ early-type cluster dwarf galaxies fainter than $M_V=-17$ mag follows a colour--magnitude relation being the extension of the CMR of the brighter cluster galaxies (see Fig. \\ref{fig:cmdhydra}). This is consistent with earlier studies where early-type dwarf galaxies are reported to follow the CMR of giant early-type galaxies \\citep[e.g.][]{Secker1997, Conselice2003, Hilker2003, LopezCruz2004, Adami2006}.  Our sample of early-type dwarf galaxies defines a magnitude--surface brightness relation in the sense that the central surface brightness (as estimated from an exponential law) increases with luminosity. Local Group dEs and dSphs follow reasonably well the same relation. By fitting S\\'{e}rsic models to the galaxies surface brightness profiles, we find a common relation for dwarf galaxies and E/S0 galaxies in a $\\mu_0$ -- $M_V$ diagram, consistent with the results of \\citet{Graham2003}, \\citet{Ferrarese2006} and \\citet{Cote2006, Cote2007, Cote2008}. The slope of the relation is found to be equal to one. Moreover, in a $\\mu_{\\mathrm{e}}$ -- $M_V$ diagram, galaxies fainter than $M_V = -18$ mag define to a good approximation a relation of constant effective radius $R_{\\mathrm{e}}\\sim 0.8$ kpc (see Fig. \\ref{fig:sersic}).  Our results, namely the \\emph{common} relations for dwarf galaxies and giant elliptical galaxies in a colour -- magnitude and a  $\\mu_0$ -- $M_V$ diagram over a wide range of magnitudes, suggest that dEs are not a separate class of objects but rather the low mass counterparts of massive early-type galaxies. The almost constant effective radius of galaxies with magnitudes $-18<M_V<-10$ mag might indicate the same.  We find a similarity of the Hydra\\,I CMR to both the Fornax and the LG relation. Since the environments in Hydra\\,I and Fornax are different from the Local Group in terms of mean density and strength of the gravitational potential \\citep{Girardi1993, Girardi1998, Karachentsev2006}, this could imply that internal evolution is more important for the global photometric properties of dwarf galaxies than external influences due to different environments. The almost identical magnitude--surface brightness relations of the Hydra\\,I and the LG dwarfs may support this hypothesis. \\citet{SmithCastelli2008} argue in the same manner based on the results of their photometric study of the galaxy population in the Antlia cluster. Further deep photometric studies in various clusters of different mass and dynamical state have to be performed in order to investigate the influence of different environments on the photometric and structural properties of dwarf galaxies.    \\subsection{The Hydra\\,I LF} We derive a very flat luminosity function ($\\alpha \\sim -1.1$, as derived from a Schechter fit) for the Hydra\\,I cluster. Recently, \\citet{Yamanoi2007} found $\\alpha \\sim -1.6$ for the cluster LF. Our observed fields overlap with their central observed region and we reach a similar limiting magnitude. The seeing of our images is slightly better. However, \\citet{Yamanoi2007} do not fit Schechter functions to their data but use a power-law model to determine the parameter $\\alpha$ at the faint end of the LF. In an analogous manner we derive a slope of $\\alpha \\sim -1.4$ from our data, being more consistent with the results of Yamanoi et al. Qualitatively, we find the same behaviour for the Fornax LF \\citep{Hilker2003, Mieske2007}. When fitting a power-law to the faint end, we derive $\\alpha \\sim -1.3$ for $-14.0<M_V<-9.8$ mag, as opposed to $\\alpha \\sim -1.1$ when fitting a Schechter function. These findings can be interpreted in the way that the description by a power-law model seems to result in steeper slopes than the description by Schechter functions. The slope of the power-law model might be considered as an upper limit in this context.  We point out that the differences in $\\alpha$ could also arise from the different methods used to construct the LF. The dwarf galaxy candidates in \\citet{Yamanoi2007} were selected only by setting lower limits in FWHM to the SExtractor detections. Contaminating background galaxies were statistically subtracted. No curve-of-growth analysis or surface brightness profile fitting was performed to cross-check the SExtractor results. By means of the surface brightness profile and the colour we could exclude a number of objects from our sample (see Sect. \\ref{sec:hydracmd} and Fig. \\ref{fig:cmdhydra}). Our sample should therefore not be contaminated by many background galaxies that only resemble cluster dwarf galaxies.  Being consistent with our results, \\citet{Trentham2002} measured the LF in different local environments (including the Virgo cluster) and found shallow logarithmic slopes of $\\alpha \\sim -1.2$. They applied selection criteria similar to ours (based on surface brightness and morphology). \\citet{Trentham2002} stated that it is unlikely that they have missed a large number of dwarfs, since they were sensitive to very low surface brightness galaxies and the seeing was good enough to distinguish high surface brightness dwarfs from background galaxies. The same is true for our data. We could potentially have missed very compact cluster members that resemble objects like M32, but these objects are rare \\citep{Drinkwater2000b, Mieske2005b, Chilingarian2007, Chilingarian2008} and do not significantly contribute to the LF.  Clearly, for the determination of the faint-end of the LF it is crucial which method is used to identify contaminating background galaxies. Different methods in constructing the LF can lead to different results. One example for this is the Fornax cluster. Applying morphological selection criteria, \\citet{Hilker2003} found a flat LF, whereas \\citet{Kambas2000} reported on a very steep LF ($\\alpha \\simeq -2$) for the same cluster. However, by comparing the data sets \\citet{Hilker2003} could show that most of Kambas' dwarf galaxy candidates were non-members of the Fornax cluster.  \\vspace{0.25cm} \\noindent We conclude that at least for nearby galaxy clusters that are close enough for large telescopes like the VLT to resolve even faint dwarf spheroidal galaxies under good seeing conditions, the membership assignment by means of morphology and surface brightness seems to be an appropriate way to construct the galaxy luminosity function and constrain photometric scaling relations."
    },
    "0806/0806.1951_arXiv.txt": {
        "abstract": "We perform a detailed investigation into the disruption of central cusps via the transfer of energy from sinking massive objects. Constant density inner regions form at the radius where the enclosed mass approximately matches the mass of the infalling body. We explore parameter space using numerical simulations and give an empirical relation for the size of the resulting core within structures that have different initial cusp slopes. We find that infalling bodies always stall at the edge of these newly formed cores, experiencing no dynamical friction over many dynamical times. As applications, we consider the resulting decrease in the dark matter annihilation flux due to centrally destroyed cusps, and we present a new theory for the formation of close binary nuclei -- the `stalled binary' model. We focus on one particularly interesting binary nucleus system, the dwarf spheroidal galaxy VCC~128 which is dark matter dominated at all radii. We show that its nuclei would rapidly coalesce within a few million years if it has a central dark matter cusp slope steeper than $r^{-1}$. However, if its initial dark matter cusp is slightly shallower than a log slope of $-0.75$ at $\\sim 0.1$\\% of the virial radius, then the sinking nuclei naturally create a core equal to their observed separation and stall. This is close to the log slope measured a recent billion particle CDM halo simulation. ",
        "introduction": "\\label{sec:intro} Massive objects orbiting within a cuspy mass distribution are expected to lose momentum and sink via dynamical friction \\citep{chandrasekhar, white, hernquist, capuzzo}. While spiralling inwards the massive perturber transfers momentum to central particles/stars etc, moving them to a larger orbital radius. One effect of this process is to make an initially cuspy dark matter distribution shallower \\citep{amr, amr2, tonini, ma, jardel}.  In the prevailing $\\Lambda$CDM cosmology, dark matter haloes are cuspy, having an inner density slope $\\rho(r) \\propto r^{-\\gamma}$ with $\\gamma > 1$ beyond $\\approx$ 1\\% of the virial radius \\citep[e.g.][]{dubinski, navarro, diemand}. By contrast, observations of dwarf galaxies seem to indicate that they have a cored mass distribution \\citep[e.g.][]{deblok, sanchez, kleyna}, while controversial evidence for cored mass distributions in dwarf spiral galaxies has been debated for over a decade \\citep{moore2}.  Previously published work has demonstrated numerically \\citep{amr, amr2, romano, ma, merritt, justin, jardel} and semi-analytically \\citep{tonini} that a sinking massive compact object -- a {\\it perturber} -- will transfer energy and angular momentum to the background via dynamical friction, creating a central constant density core from an initially cuspy density distribution. Once a core has formed, dynamical friction is no longer effective \\citep{mich, justin, inoue}. Dynamical arguments show that sinking perturbers will stall at the outer edge of a core \\citep{justin,inoue, inoue2}.  Here we quantify this stalling behaviour as a function of perturber mass $M_{\\rm pert}$ and central cusp slope $\\gamma$. We consider a much larger range in $M_{\\rm pert}$ and $\\gamma$ than in previous papers \\citep{mich, justin} and find that stalling persists even at very high perturber mass, as found also recently by \\citet{gualandris}. We also investigate how the cusp is physically transformed into a core.  The core formation mechanism we study is just one of several ways in which central cores can be formed. Core formation can also proceed by three-body encounters with a supermassive black hole binary \\citep{milosavljevic}, as a result of rapid mass loss due to supernovae outflows \\citep[e.g.][]{navarro2, justin3}, or as a result of the rapid ejection of a central supermassive black hole due to anisotropic gravitational radiation recoil \\citep[e.g.][]{boylankolchin}. If these mechanisms play an important role then our derived stalling radii as a function of $M_{\\rm pert}$ and $\\gamma$ will be lower bounds.  We consider two applications of our results. The first is the effect of cusp-destruction on the expected dark matter annihilation signal from galaxies and dwarf galaxies. For a wide range of popular dark matter particle models, dark matter can self-annihilate to produce $\\gamma$-rays \\citep{gunn}. Since the signal goes as the dark matter density squared it is sensitive to the central density distribution \\citep{silk, lake}. For our second application, we present a new theory for the formation of close binary nuclei -- the `stalled binary' model. Close binary nuclei with projected separation $< 100$\\,pc have been observed on a range of scales in the Universe \\citep[e.g.][]{lauer1, lauer, lauer2, bender, houghton, mast, victor}. The standard model for these has become the \\cite{tremaine} eccentric disc model originally proposed to explain M31 \\citep{lauer2}. However, when the double nucleus of M31 was first discovered, \\cite{lauer1} speculated that the two bright nuclei really were just that -- one from M31 and the other the cannibalised centre of a smaller merged galaxy. The primary argument against this was that dynamical friction would cause the nuclei to rapidly coalesce. We show that, as binary nuclei sink via dynamical friction, they create a central constant density core. They then stall at the edge of this core experiencing no further friction over many dynamical times. We apply our new `stalled binary' model to one particularly interesting binary nucleus system -- the dwarf spheroidal galaxy VCC~128 discovered by \\citet{victor}. We show that this galaxy is dark matter dominated at all radii. As a result -- if its binary nucleus is a `stalled binary' -- VCC~128 gives us a unique opportunity to constrain the central log-slope of the dark matter density profile on very small scales.  This paper is organised as follows: In \\S2 we describe our analytical framework, which is supported using $N$-body simulations. In \\S3 we apply our findings to dark matter annihilation, to binary nuclei, and to the special case of VCC~128. Finally, in \\S4 we present our conclusions. ",
        "conclusions": "We have performed a detailed investigation into the disruption of central cusps via the transfer of energy from sinking massive objects. Constant density inner regions form at the radius where the enclosed mass approximately matches the mass of the infalling body. We explored parameter space using numerical simulations and gave an empirical relation for the size of the resulting core within structures that have different initial cusp slopes. We went on to demonstrate that infalling bodies always stall at the edge of these newly formed cores, experiencing no dynamical friction over many dynamical times.  As applications, we considered the resulting decrease in the WIMP annihilation flux due to centrally destroyed cusps; and we presented a new theory for the formation of close binary nuclei -- the `stalled binary' model. Our key results are as follows:  \\begin{enumerate} \\item Core formation due to sinking massive objects can soften a central dark matter cusp reducing the expected WIMP annihilation flux (predicted by structure formation simulations that model dark matter in the absence of baryons) by up to a third.  \\item Core formation due to sinking massive objects could help to alleviate the long-standing cusp-core problem \\citep[see eg:][]{spekkens,moore2,deblok}. From equation (\\ref{eq:hash}), a $\\sim$1\\,kpc sized core will form from perturber having $\\sim$1\\% of the mass of the host. This recovers the earlier results of \\citet{amr} and \\cite{jardel}. However, such massive infalling perturbers must later disrupt or be removed in order to be consistent with the low surface density of stars and gas observed in galaxies where the cusp-core problem is most apparent.  \\item Infalling nuclei at the centres of galaxies will evacuate a core and stall indefinitely, provided that the initial background density is not significantly steeper than $r^{-1}$. This could explain a number of binary nuclei systems in the Universe.  \\item We focused on the special binary nucleus system VCC~128 since it is dark matter dominated at all radii. Assuming that its binary nucleus can be explained by our `stalled binary' model, we found that the initial inner log density slope $\\gamma$ of the dark matter halo of VCC~128 must be $0.5 < \\gamma < 0.75$ at $\\sim 0.1\\%$ of the virial radius. For $\\gamma > 0.75$ initially, the dynamical friction sink-in time is so small in comparison to the lifetime of the galaxy that we run into a fine tuning problem. For $\\gamma < 0.5$ initially, the nuclei stall far beyond their current projected separation of 32\\,pc. For $0.5 < \\gamma < 0.75$ initially, the nuclei create a central constant density core of separation $\\sim 40$\\,pc after which they stall indefinitely. Our preferred inner slopes are consistent with those found in the recent billion particle CDM halo simulations of \\citet{ghalo}, \\citet{vialactea} and \\citet{aquarius}. \\end{enumerate}"
    },
    "0806/0806.1817_arXiv.txt": {
        "abstract": "{In this letter we present a morphological comparison between giant radio halos and radio mini-halos in galaxy clusters based on radio--X-ray luminosity, $P_{1.4}$-$L_{\\rm X}$, and radio luminosity-size, $P_{1.4}$-$R_{\\rm H}$, correlations. We report evidence that $P_{1.4}$-$L_{\\rm X}$ and $P_{1.4}$-$R_{\\rm H}$ trends may also exist for mini--halos: mini--halo clusters share the same region of giant halo clusters in the $(P_{1.4},L_X)$ plane, whereas they are clearly separated in the $(P_{1.4},R_H)$ plane. The synchrotron emissivity of mini-halos is found to be more than $50$ times larger than that of giant halos, implying a very efficient process for their origins. By assuming a scenario of sporadical turbulent particle re-acceleration for both giant and mini halos, we discuss basic physical differences between these sources. Regardless of the origin of the turbulence, a more efficient source of injection of particles, which eventually takes part in the re-acceleration process, is required in mini-halos, and this may result from the central radio galaxy or from proton-proton collisions in the dense cool core regions. ",
        "introduction": "The intra--cluster medium (ICM) consists not only of hot gas emitting in X-rays but also of non-thermal components. The major evidence for this comes from observations in the radio band where Mpc--scale diffuse synchrotron emission from the ICM is detected in a number of clusters (\\egg Feretti 2005; Ferrari et al. 2008), indicating the presence of relativistic electrons and magnetic fields.  These radio sources are generally referred to as giant radio halos when located at the cluster center and radio relics when located at the cluster periphery.  There are also some examples of diffuse radio emission on smaller scales ($\\sim 200$-$500$ kpc), referred to as mini radio halos, extending around powerful radio galaxies at the center of some cool core clusters, \\ie clusters characterized by a very peaked surface brightness profile and short central cooling time formerly known as `cooling flow' clusters (\\egg Peterson \\& Fabian 2006). Galaxy clusters hosting giant halos are found to always be characterized by a peculiar dynamical status indicative of very recent or ongoing merger events (\\eg Buote 2001; Schuecker et al 2001), whereas clusters hosting mini halos are characterized by a cool core, with or without signs of moderate dynamical activity. Several statistical studies reveal that radio halos are not common in clusters (Giovannini et al. 1999; Kempner \\& Sarazin 2001; Venturi et al. 2008; Brunetti et al. 2007; Cassano et al. 2008); instead, the statistics for mini halos is much poorer.   The main difficulty in understanding the origin of the synchrotron emitting electrons in both giant halos and mini halos is related to the fact that the diffusion length of the relativistic electrons is much shorter than the typical scale of the radio emission (\\eg Brunetti 2003). Therefore both giant halos and mini halos cannot be explained in terms of simple diffusion of the relativistic electrons from one or more cluster radio galaxies. Two main possibilities have been proposed so far to explain the origin of both giant radio halos and mini radio halos: {\\it i)} {\\it re-acceleration} models, whereby relativistic electrons injected in the ICM are re-energized {\\it in situ} by various mechanisms associated with turbulence in the ICM. Turbulence in radio halos is supposed to be generated by massive merger events (\\eg Brunetti et al. 2001; Petrosian et al. 2001). In mini radio halos, a seed large--scale turbulence frozen into the flow could be amplified by the compression of the ICM in the cool core (Gitti et al. 2002, 2004); {\\it ii)} {\\it secondary electron} models, whereby the relativistic electrons are secondary products of the hadronic interactions of cosmic rays (CR) with the ICM (\\eg Dennison 1980; Blasi \\& Colafrancesco 1999, Pfrommer \\& En\\ss lin 2004). Although the properties of giant halos and mini halos are clearly different (different size, different dynamical state of the hosting clusters), it is not clear whether they are different astrophysical phenomena or if they might share similar physics. In this letter we carry out a morphological comparison between giant radio halos and mini radio halos aimed at studying the differences between their physical properties. We also consider the case of diffuse cluster sources with intermediate properties between giant halos and mini halos. To do this we investigate the presence of scaling relations between the main properties of these sources. A $\\Lambda$CDM cosmology ($H_{o}=70\\,\\rm km\\,\\rm s^{-1}\\,\\rm Mpc^{-1}$, $\\Omega_{m}=0.3$, $\\Omega_{\\Lambda}=0.7$) is adopted. ",
        "conclusions": "\\begin{figure} \\begin{center} \\includegraphics[width=0.3\\textwidth]{E_B_new2.ps} \\caption[]{Ratio between turbulence energy densities of MHs and GHs normalized to the thermal ones as a function of $B_{MH}$. The calculations are reported for $f\\approx1$, $z=0.1$, $T_{GH}/T_{MH}\\approx 3$ and in the case of $B_{MH}\\approx 3 B_{GH}$ (solid line) and in the case $B_{MH}\\approx 6 B_{GH}$ (dashed line).} \\label{Fig.stupido} \\end{center} \\end{figure}  In this letter we have compared the observed properties of mini radio halos (MHs) and giant radio halos (GHs) in clusters of galaxies.  GHs are the most prominent evidence of non-thermal components in the ICM and several correlations between thermal and non-thermal properties have been explored for these sources, including those relating $P_{1.4}$ to $L_{\\rm X}$ and to $R_{\\rm H}$ (\\eg Cassano et al. 2006, 2007; Brunetti et al. 2007). On the other hand, an extensive investigation of the statistical properties of MHs is presently not possible since only a few clusters host well-studied MHs.  We collected a sample of MH and compared their behavior with that of GHs in the ($P_{1.4},L_{\\rm X}$) and ($P_{1.4},R_{\\rm H}$) planes. We find that $P_{1.4}-L_{\\rm X}$ and $P_{1.4}-R_{\\rm H}$ trends may also exists for MHs. While in the ($P_{1.4},L_{\\rm X}$) plane MHs and GHs share the same region, in the ($P_{1.4},R_{\\rm H}$) plane MHs do not follow the same correlation of GHs at smaller radii, but are clearly separated. We find that the typical synchrotron emissivity of MHs is at least $50$ times larger than that of GHs. This implies a very efficient mechanism at the origin of the emitting electrons in MHs. For completeness we also consider the few cases of smaller scale emission in non-CC (and without central radio galaxy) merging clusters. These sources are morphologically intermediate between GH and MH and may be GH at some early evolutionary stage.  The distribution in the ($P_{1.4},L_{\\rm X}$) plane of a small sample of CCCs with available radio observations suggests that MHs are not ubiquitous in CCCs, with upper limits for CCCs without diffuse radio emission well below the radio power of MHs in clusters with similar $L_X$. Those CCCs without MHs also appear to be more relaxed than that with MHs. All these findings, if confirmed, would point in favor of sporadic turbulent re-acceleration as the origin of the emitting particles. In addition to the possibilities already explored in the literature (Gitti et al. 2002; Mazzotta \\& Giacintucci 2007), minor mergers (see also Gitti et al. 2007a) and/or the central AGN outbursts may contribute to the injection of turbulence in the ICM of CCCs. By adopting this scenario, under the assumption that magnetosonic waves drive the particle acceleration process, we find that the larger synchrotron emissivity of MHs can be explained by assuming that the energy density of the relativistic particles that interact with turbulence is about one order of magnitude higher than in GHs, and that this does not necessarily imply a larger amount of turbulence in MHs. The extra amount of relativistic particles in these sources may be provided by the central cluster galaxy or by secondary electrons injected in the dense cool core region."
    },
    "0806/0806.0399_arXiv.txt": {
        "abstract": "We survey \\ion{N}{5} absorption in the afterglow spectra of long-duration gamma-ray bursts (GRBs) with the intent to study highly ionized gas in the galaxies hosting these events. We identify a high incidence (6/7) of spectra exhibiting \\ion{N}{5} gas with $z \\approx z_{GRB}$ and the majority show large column densities $\\N{N^{+4}} \\gtrsim 10^{14}\\cm{-2}$. With one exception, the observed line-profiles are kinematically `cold', i.e.\\ they are narrow and have small velocity offset ($\\delta v \\lesssim 20 \\mkms$) from absorption lines associated with neutral gas. In addition, the \\ion{N}{5} absorption has similar velocity as the UV-pumped fine-structure lines indicating these high ions are located within $\\approx 1$\\,kpc of the GRB afterglow. These characteristics are unlike those for \\ion{N}{5} gas detected in the halo/disk of the Milky Way or along sightlines through high $z$ damped \\lya\\ systems but resemble the narrow absorption line systems associated with quasars and some high $z$ starbursts. We demonstrate that GRB afterglows photoionize nitrogen to \\nion\\ at $r \\approx 10$\\,pc. This process can produce \\ion{N}{5} absorption with characteristics resembling the majority of our sample and and we argue it is the principal mechanism for \\nion\\ along GRB sightlines. Therefore, the observations provide a snapshot of the physical conditions at this distance. In this scenario, the observations imply the progenitor's stellar wind is confined to $r < 10$\\,pc which suggests the GRB progenitors occur within dense ($n > 10^3 \\cm{-3}$) environments, typical of molecular clouds. The observations, therefore, primarily constrain the physical conditions -- metallicity, density, velocity fields -- of the gas within the (former) molecular cloud region surrounding the GRB. ",
        "introduction": "Long-duration gamma-ray bursts (GRBs) are believed to have massive star progenitors arising in active star-forming regions of high $z$ galaxies \\citep[e.g.][]{wb06}.  Roughly half of these events have associated UV/optical afterglows and a subset of these have apparent magnitudes sufficient for high-resolution spectroscopy using 10m-class telescopes \\citep[e.g][]{fdl+05,cpb+05}.  In principle, the power-law afterglow spectrum has imprinted within it features from gas throughout the interstellar medium (ISM) of the host galaxy. This stands in contrast to studies of quasars whose integrated photon output ionizes their ISM\\footnote{The obvious exception is the gas identified as broad absorption lines in quasar spectra which is located very close to the quasar.} and surrounding gas out to many tens of kpc. % Furthermore, although quasar sightlines frequently penetrate foreground, star-forming galaxies \\citep[the so-called damped \\lya\\ systems, QSO-DLA;][]{wgp05}, these are probed according to gas cross-section and quasar sightlines should only rarely intersect the small, dense regions undergoing active star-formation \\citep{zp06}.  In these respects, GRB afterglow spectra allow one to probe a diversity of phases in the ISM of star-forming galaxies: the circumstellar material from the massive star progenitor, the \\ion{H}{2} region produced by the progenitor and neighboring OB stars, the neutral ISM of the host galaxy, and any diffuse gas within the galactic halo.  Unfortunately, even though these phases arise at distinct distances along the sightline, the observed spectrum resolves only the relative velocities of the gas.  To focus analysis on a specific phase, one is generally forced to isolate a unique ion and/or material associated with a specific velocity.   GRB afterglow spectra reveal large column densities of \\ion{H}{1} gas and metals associated with the host galaxy ISM \\citep{sff03,vel+04,jfl+06}.  The analysis of the metal-line transitions have localized this neutral gas within the ambient ISM of the host galaxy.  Specifically, the detection of fine-structure lines of Si$^+$ and Fe$^+$ ions places the gas within $\\approx 1$\\,kpc of the GRB afterglow while the detection of \\ion{Mg}{1} absorption requires the gas to lie at distances greater than $\\approx 100$\\,pc \\citep{pcb06}.  These conclusions are supported by direct distance determinations based on analysis of line-variability in fine-structure lines \\citep[$\\approx 100$\\,pc and 2\\,kpc from GRB~020813 and GRB~060418 respectively;][]{dcp+06,vls+07}.  The majority of afterglow spectra also show high-ion absorption (e.g.\\ \\ion{C}{4}) that is offset by several tens to hundred \\kms\\ from the peak optical depth of the neutral ISM \\citep{cpr+07}.  This diffuse, ionized gas is also traced by strong transitions of low-ions (e.g.\\ \\ion{Si}{2}~1526) but without coincident fine-structure absorption.  Therefore, the gas must lie at distances greater than a few kpc from the GRB.  These characteristics identify the clouds as partially ionized gas within the halo of the GRB host galaxy \\citep{pcw+08}.  While studies of the gas in the neutral ISM and galactic halo are valuable for studying the physical conditions in star-forming galaxies, these phases offer only indirect constraints on the nature of the GRB progenitor \\citep{rtb02}. Of great interest is to identify gas located within the star-forming region or even gas shed by the progenitor itself. To date, however, no study has presented compelling evidence for gas within $\\approx 100$\\,pc of the GRB: neither circumstellar material \\citep{cpr+07}, the molecular cloud that presumably beget the progenitor \\citep{tpc+07}, nor material associated with a pre-existing \\ion{H}{2} region \\citep{wph+08}.  Regarding the latter phase, most of the key diagnostics (e.g.\\ \\ion{Si}{4}, \\ion{Al}{3}, \\ion{Si}{3}, \\ion{C}{4}) are either compromised by blending with the \\lya\\ forest or can be confused with a galactic halo component.  Presently, there is only indirect evidence for significant column densities of ionized gas near GRB: a number of GRB sightlines exhibit X-ray absorption with implied metal column densities that significantly exceed the neutral ISM column densities measured from the optical spectra \\citep{gw01,whf+07}. This result\\footnote{There have also been claims of temporal variation in the X-ray absorption spectrum \\citep[e.g.][]{clr+07} which would also imply significant metals near the GRB afterglow \\citep{lp02}, but these variations are more naturally explained by the temporal evolution of the intrinsic X-ray spectrum \\citep{bk07}.} hints at a large reservoir of highly ionized gas near the GRB which has not yet been revealed by the rest-frame UV spectra acquired by ground-based facilities.  The challenge to identify and study gas close to the GRB has motivated us to survey GRB afterglow spectra for the presence of \\ion{N}{5} absorption.  Because the \\nthr\\ ion has a large ionization potential (IP=77eV), it is difficult to produce \\nion, especially using stellar radiation fields.  In the ISM of local galaxies, \\nion\\ is generally believed to trace collisionally ionized gas either in equilibrium at a high temperature ($T > 10^{5}$K) or out of equilibrium due to a post-shocked gas cooling from $T > 10^6$\\,K \\citep[e.g.][]{is04a}.  In terms of GRB studies, however, the GRB event itself and its bright afterglow emit sufficient numbers of $h\\nu \\approx 80$\\,eV photons to produce N$^{+4}$ gas near the progenitor.  Observationally, the \\nion\\ ion is notable for exhibiting an alkali doublet at $\\lambda\\lambda 1238,1242$ in the rest-frame which lies redward of the \\ion{H}{1} \\lya\\ transition.  In GRB sightlines, therefore, this transition (unlike \\ion{O}{6} and \\ion{S}{6} doublets) does not blend with \\ion{H}{1} absorption from the \\lya\\ forest, and GRB afterglow spectra that cover \\lya\\ will often provide an analysis of the \\ion{N}{5} doublet.  Indeed, previous studies have reported the detection of the \\ion{N}{5} doublet along individual sightlines \\citep{vel+04,cpb+05,twl+07}.  In this paper we perform a systematic search and analysis of \\nion\\ gas for a modest sample of $z>2$ GRBs.  We report on the incidence of its detection, its characteristic column density, and compare the line-profiles with other transitions identified along the sightline. Finally, we investigate the origin of this gas and explore constraints on the nature of the GRB progenitor environment.  \\begin{deluxetable}{lcccccc} \\tablewidth{0pc} \\tablecaption{GRB-NV SAMPLE\\label{tab:obs}} \\tabletypesize{\\footnotesize} \\tablehead{\\colhead{GRB} &\\colhead{RA} & \\colhead{DEC} & \\colhead{Instrument} & \\colhead{$R$} & \\colhead{Ref}} \\startdata \\object{021004}&00:26:54.68&+18:55:41.6&UVES&52,000&1\\\\ \\object{030323}&11:06:09.40&$-$21:46:13.2&FORS2&2,600&2\\\\ \\object{050730}&14:08:17.14&$-$03:46:17.8&MIKE&30,000&3\\\\ \\object{050820}&22:29:38.11&+19:33:37.1&HIRES&30,000&4\\\\ \\object{050922C}&19:55:54.48&$-$08:45:27.5&UVES&30,000&5\\\\ \\object{060206}&13:31:43.42&+35:03:03.6&ISIS&4,000&6\\\\ \\object{060607}&21:58:50.40&$-$22:29:46.7&UVES&43,000&7\\\\ \\enddata \\tablerefs{ 1: \\cite{fdl+05}; 2: \\cite{vel+04}; 3: \\cite{cpb+05}; 4: \\cite{pcb+07}; 5: \\cite{pwf+07}; 6: \\cite{twl+07}; 7: \\cite{gcn5237}}  \\end{deluxetable}   \\begin{deluxetable*}{lcccccccccccc} \\tablewidth{0pc} \\tablecaption{SURVEY SUMMARY\\label{tab:summ}} \\tabletypesize{\\footnotesize} \\tablehead{\\colhead{GRB} &\\colhead{$z_{GRB}$} & \\colhead{log \\nhi} & \\colhead{[M/H]$^a$} & \\colhead{$W_{1238}^b$} & \\colhead{$\\log \\N{N^{+4}}$} & \\colhead{$\\delta$(NV)$^c$} \\\\ & & ($\\cm{-2}$) &  & (\\AA) & ($\\cm{-2}$) & (\\kms) } \\startdata GRB021004&2.3291&19.00&$ 0.0$&$ 0.307 \\pm 0.008$&$ 14.64 \\pm 0.04$&$-10\\pm 10$\\\\ GRB030323&3.3720&21.90&$>-0.9$&$ 0.325 \\pm 0.043$&$> 14.35$&$ 20\\pm 30$\\\\ GRB050730&3.9686&22.15&$-2.3$&$ 0.142 \\pm 0.021$&$ 14.09 \\pm 0.08$&$  5\\pm  5$\\\\ GRB050820&2.6147&21.00&$-0.6$&$ 0.045 \\pm 0.007$&$ 13.45 \\pm 0.05$&$-90\\pm 10$\\\\ GRB050922C&2.1990&21.60&$-2.0$&$ 0.197 \\pm 0.026$&$> 14.19$&$-20\\pm  5$\\\\ GRB060206&4.0480&20.85&$-0.9$&$ 0.093 \\pm 0.010$&$ 13.73 \\pm 0.15$&$  0\\pm 15$\\\\ GRB060607&3.0748&16.80&$ 0.0$&$-0.009 \\pm 0.003$&$< 12.61$&$$\\\\ \\enddata \\tablenotetext{a}{Gas metallicity derived from low-ion absorption. See \\cite{pcd+07} and \\cite{dpc+08} for details.} \\tablenotetext{b}{Rest-frame equivalent width of the \\ion{N}{5}~1238 transition.} \\tablenotetext{c}{Estimated velocity offset between the rough centroid of the \\ion{N}{5} line-profile and the peak optical depth of the fine-structure lines.}  \\end{deluxetable*}  \\begin{figure*} \\epsscale{0.8} \\begin{center} \\includegraphics[height=6.8in,angle=90]{f1.eps} \\end{center} \\caption{Velocity profiles of the \\ion{N}{5} doublet (top two panels in each sub-figure) compared against a low-ion resonance and fine-structure transition from our sample of GRB afterglow spectra with resolution $R > 2500$. We do not show the data for GRB~060206 (also included in the analysis);  the interested reader should see \\cite{twl+07}. With the exception of GRB~060607, the data reveal a positive detection of the \\ion{N}{5} doublet. Aside from GRB~050820, the \\nion\\ line-profiles are relatively narrow and are roughly aligned with the low-ion resonance and fine-structure profiles. } \\label{fig:data} \\end{figure*} ",
        "conclusions": "We have performed a survey of \\ion{N}{5} absorption along seven GRB sightlines and reported six positive detections within 100\\kms\\ of the neutral gas associated with the host galaxy.  Aside from the GRB~050820 sightline (where the \\ion{N}{5} absorption is broad, weak and offset by $\\delta v \\approx -100\\mkms$), the \\nion\\ gas has large column density and kinematically `cold' line-profiles.  The latter characteristic refers to a low velocity dispersion and a small offset $|\\delta v| < 20 \\mkms$ from the neutral gas.  The \\ion{N}{5} profiles are also coincident in velocity with fine-structure absorption which suggests the gas is located within $\\approx 1$\\,kpc of the GRB afterglow.  We have explored several scenarios that could produce \\ion{N}{5} absorption along GRB sightlines.  Models related to the halo of the host galaxy or material shock-heated by the progenitor's stellar wind are disfavored by the observations.  In contrast, a scenario where the \\nion\\ gas is material photoionized by the GRB afterglow naturally reproduces the observations provided the gas at $r \\approx 10$\\,pc is cold ($T \\approx 10^4$\\,K), and has a modest density ($n \\approx 1 \\cm{-3}$), a non-negligible metallicity ([N/H]~$>-2$), and a similar velocity as the ISM at $r \\gtrsim 100$\\,pc.  The afterglow photoionization model places several important constraints on the progenitor and its environment.  In particular, this scenario requires that the stellar wind of the progenitor terminates at less than $r \\approx 10$\\,pc.  This suggests the GRB progenitor has a weak, main-sequence stellar wind owing to a low mass-loss rate, a low wind speed, and/or a short lifetime.  These characteristics may be a natural consequence of the progenitors that favor the GRB phenomenon, e.g.\\ higher mass, lower metallicity stars. The wind can also be confined by introducing a dense $(n \\gtrsim 10^3 \\cm{-3})$ external medium.  Perhaps GRBs are preferentially embedded within the dense regions of molecular clouds as opposed to a violent, starburst region.  On the other hand, we note that the \\ion{N}{5} absorption detected along the GRB~050820 sightline has characteristics consistent with shock-heated gas provided a shock with $v \\approx 100\\mkms$.  In this case, we may be seeing the signatures of a starburst galaxy.   Before concluding, we wish to comment on a few directions for future research.  One aspect to explore is how the soft X-ray absorption observed in the afterglow spectroscopy compares with the \\nion\\ observations.  To zeroth order, both measurements are sensitive to the column density of metals near the GRB progenitor, although likely at somewhat different radii.  A comprehensive model of these observations may constrain the density profile of gas close to the GRB.  Another implication of our research is that one predicts recombinations of the \\nion\\ gas (and other high-ions, e.g.\\ O$^{+5}$) once the afterglow fades.  For low $z$ GRB, it may be possible to observe this line emission with a sensitive ultraviolet telescope \\citep{prl00}.  Finally, one predicts that other high ionization states will be produced by the afterglow (e.g.\\ O$^{+5}$, S$^{+5}$) that could be studied in a similar fashion to constrain the relative abundances of gas near the progenitor star."
    },
    "0806/0806.2355_arXiv.txt": {
        "abstract": "{Recent analyses of several isospin effects in heavy-ion reactions have allowed us to constrain the density dependence of nuclear symmetry energy at sub-saturation densities within a narrow range. Combined with constraints on the Equation of State (EOS) of symmetric nuclear matter obtained previously from analyzing the elliptic flow in relativistic heavy-ion collisions, the EOS of neutron-rich nuclear matter is thus partially constrained. Here we report effects of the partially constrained EOS of neutron-rich nuclear matter on the mass-radius correlation, moment of inertia, elliptical deformation and gravitational radiation of (rapidly) rotating neutron stars.}  \\keyword{Symmetry Energy, Equation of State, Neutron-Rich Nuclear Matter, Neutron Stars, Gravitational Waves} \\PACS{21.65.Cd., 21.65.Ef,25.70.-z,21.30.Fe,21.10.Gv,21.60-c.}  \\begin{document} ",
        "introduction": " ",
        "conclusions": ""
    },
    "0806/0806.3090_arXiv.txt": {
        "abstract": "The ionizing ultraviolet background (UVB) during reionization can suppress the gas content of low-mass galaxies, even those capable of efficient atomic cooling (i.e. with virial temperatures $\\Tvir \\gsim 10^4$ K). This negative radiative feedback mechanism can thus reduce the star formation efficiencies of these halos, which may delay the completion of reionization. In this work, we explore the importance of UV radiative feedback on $\\Tvir \\gsim 10^4$ K halos during the middle and late stages of reionization.  We do not try to self-consistently model reionization; instead, we explore a large parameter space in an attempt to draw general, robust conclusions. We use a tiered approach. Using 1-D hydrodynamical simulations, we model the ability of gas to collapse onto halos of various masses under UVBs of various intensities.  We then generate realistic, parametrized maps of the inhomogeneous UVB, using large-scale semi-numeric simulations. By combining these results, we find that under all reasonably conservative scenarios, UV feedback on atomically-cooled halos is not strong enough to notably delay the bulk of reionization. Such a delay is only likely if ionizing efficiencies of $z\\gsim10$ sources are much higher ($\\sim$ two orders of magnitude) than $z\\sim6$ data seem to imply.  Towards the end of reionization, star formation can be quenched only in halos in a narrow mass range close to the atomic-cooling threshold. This result depends only weakly on the intensity of the UVB: quenching star formation in halos just twice as massive requires an order of magnitude increase in source ionizing efficiencies. This implies that the natural time-scale for the bulk of reionization is the growth of the global collapsed fraction contained in $\\Tvir\\gsim 10^4$ K halos. Thus the likely reionization scenario would involve a small HII filling factor ``tail'' extending to high redshifts, governed by more complicated feedback on $\\Tvir \\lsim 10^4$ K objects, followed by a period of relatively rapid evolution in the HII filling factor. Furthermore, our results underscore the importance of extended dynamical ranges when modeling reionization.  Simulations must be capable of resolving halos with mass $\\gsim 10^8 \\Msun$, even when modeling the late stages of reionization, while at the same time being large enough to capture HII regions several tens of Mpc in size. ",
        "introduction": "\\label{sec:intro}  The epoch of reionization, when light from early generations of astrophysical objects began flowing through the intergalactic medium (IGM), remains one of the most compelling frontiers of modern cosmology, offering a wealth of information about cosmological structure formation and physical processes in the early universe.  Only recently have we begun to gather clues concerning this epoch.  The presence of flux in the Lyman line absorption regions of $z\\lsim6$ quasars discovered in the Sloan Digital Sky Survey (SDSS) indicates that reionization is mostly completed by this redshift (though note that large-scale fluctuations in the ionization field can yield long ionized sightlines even before the completion of reionization, e.g. \\citealt{Lidz07}).  Spectra from some higher redshift quasars have even been interpreted as evidence of a partially neutral IGM \\citep{MH04, WL04_nf, Fan06, MH07}, though most of these interpretations are controversial \\citep{LOF06, BH07, Maselli07, BRS07, MF08damp}. The {\\it Wilkinson Microwave Anisotropy Probe} ({\\it WMAP}) measured a Thomson scattering optical depth of $\\tau_e = 0.087\\pm0.017$ \\citep{Dunkley08}.  Assuming instantaneous reionization, this value of $\\tau_e$ corresponds to a reionization redshift of $z = 11.0 \\pm 1.4$.  However, this is only an integrated measurement and can tell us very little about the evolution of reionization.  With the interpretation of even the sparse existing data being subtle and complicated, much effort has been invested in improving the modeling of reionization.  We think we can accurately model the abundances and clustering properties of dark matter halos at high redshift.  Combined with some analytic relation of halo mass $\\rightarrow$ effective ionizing photon emissivity, one can get an estimate of how the global neutral fraction evolves; this technique has been used is countless reionization studies.  Aside from the seemingly magical ``$\\rightarrow$'' in the previous sentence, one of the main additional uncertainties are the feedback processes: how do sources impact the current and future generations of sources?  Answering this question is non-trivial, especially in the early universe.  Radiative, chemical and mechanical feedback can effect the ionizing efficiencies of the first, highly biased sources.  Molecular hydrogen (H$_2$) cooling can allow very small (with virial temperatures, $\\Tvir$, of several hundred Kelvin) halos to host astrophysical sources.  At later times (though still prior to the bulk of reionization, \\citealt{HRL97}), as the H$_2$ dissociative background builds-up and the contribution of molecularly-cooled halos declines, feedback enters the regime of atomically-cooled halos, $\\Tvir \\gsim 10^4$ K.  This regime is more straightforward to model as negative radiative feedback emerges as a single dominant mechanism, especially on large scales.  In the presence of an ionizing background radiation, the IGM is photo--heated to a temperature of $\\gsim 10^4$K, raising the cosmological Jeans mass, which could suppress gas accretion onto small-mass halos (e.g. \\citealt{Efstathiou92, SGB94, TW96, HG97}). Early work on this subject (so-called ``Jeans mass filtering'') suggested that an ionizing background would completely suppress star formation in low--redshift ``dwarf galaxy'' halos with circular velocities $v_{\\rm circ} \\lsim~35~\\kmps$, and partially suppress star--formation in halos with 35 $\\kmps$ $\\lsim$ $v_{\\rm circ}$ $\\lsim$ 100 $\\kmps$ \\citep{TW96}.  Many reionization studies assume prescriptions of gas suppression based on these results (e.g. \\citealt{RS03, OM04, WC07, Iliev07, WBH08}).  However, more recent studies \\citep{KI00, Dijkstra04} find that at $z \\gsim 3$, more compact halo profiles, increased cooling efficiencies, and shorter exposure times to the ultraviolet background (UVB) could lessen the importance of negative radiative feedback.  We expand on these works by exploring a wider parameter space, placed in a broader context of an inhomogenious reionization with a patchy UVB due to both source clustering and modulation by HII regions.  \\citet{Gnedin00filter} studied the effects of such radiative feedback in cosmological simulations; however, these simulation boxes were of necessity very small and only included a single reionization model.  {\\it The purpose of this paper is to explore the importance of UV radiative feedback during the middle and late stages of reionization}. We do not attempt to self-consistently model feedback during reionization; such a thing is beyond the capacity of current simulations, especially given our poor understanding of the first generations of astrophysical sources and their environments.  Instead we statistically present the effects of an inhomogeneous UVB on the suppression of gas content in low-mass galaxies {\\it capable of atomic cooling}, i.e. $\\Tvir \\gsim 10^4$ K.  We do this using a tiered approach: using numerical simulations \\citep{TW95, Dijkstra04} to calibrate very large scale, high resolution ``semi-numerical'' simulations \\citep{MF07}.  In order to keep things general, we explore wide swaths of parameter space keeping assumptions minimal.  This paper is organized as follows. In \\S \\ref{sec:mark_sims} and \\S \\ref{sec:my_sims}, we describe our hydrodynamic collapse simulations and semi-numerical cosmological simulations, respectively.  In \\S \\ref{sec:fcoll} we present the results from our collapse simulations covering a wide range of parameter choices, while in \\S \\ref{sec:dist} we present parametrized distributions of UV fluxes from our cosmological simulations. In \\S \\ref{sec:feedback} we combine these results to quantify the importance of UV feedback during reionization.  In \\S \\ref{sec:ass} we discuss the assumptions and uncertainties in our approach. Finally, in \\S \\ref{sec:conc} we present our conclusions.  We quote all quantities in comoving units, with the exception of flux, where we denote proper units with a prefix 'p'. We adopt the background cosmological parameters ($\\Omega_\\Lambda$, $\\Omega_{\\rm M}$, $\\Omega_b$, $n$, $\\sigma_8$, $H_0$) = (0.76, 0.24, 0.0407, 0.96, 0.76, 72 km s$^{-1}$ Mpc$^{-1}$), matching the three--year results of the {\\it WMAP} satellite \\citep{Spergel07}, which in turn are consistent with the recent five-year data release \\citep{Komatsu08}. ",
        "conclusions": "\\label{sec:conc}  In this paper, we quantify the importance of UV radiative feedback during the middle and late stages of reionization.  Specifically, we concern ourselves with halos capable of atomic cooling, $\\Tvir\\gsim10^4$ K. We first run suites of spherically-symmetric halo collapse simulations using various values of $z$, $M$, and $J_{21}$. We then generate parametrized UV flux distributions at $z=$ 7, 10, and 13, using semi-numerical, large-scale simulations of halo and ionization fields.  Combining these two results, we estimate how efficient is radiative feedback at hindering the progress of reionization, during its advanced stages (when $\\Tvir \\lsim 10^4$ K halos become subdominant contributors of ionizing photons).  This tiered approach allows us to explore an extremely wide range of parameter space, which is necessary to make any robust conclusions, given our poor knowledge of the properties of high-$z$ sources.  We find that under all reasonably conservative scenarios, UV feedback on atomically-cooled halos is not strong enough to notably delay the bulk of reionization.  For fiducial choices of source ionizing efficiencies (calibrated to match $z\\sim6$ LAE and LBG luminosity functions) and turn-on redshift, fewer than $40$\\% of $\\Tvir \\gsim 10^4$ K halos are left without gas at $(z, \\avenf)$ $\\approx$ (7, 0); this number drops to $\\sim10\\%$ when the distribution is mass-weighted (as is more appropriate for estimating the global star formation rate). Suppressing more than half of such halos requires a factor of $\\sim$ 3--4 increase in fiducial ionizing efficiencies at $(z, \\avenf)$ $\\approx$ (7, 0) and over two orders of magnitude increase for the same fraction at $(z, \\avenf)$ $\\approx$ (10, 0)!  We find that feedback is very strongly dependent on halo mass.  For example, at $(z, \\avenf)$ $\\approx$ (7, 0) and $\\toteff=0.25$, the fraction of halos left without gas decreases from 1 to 0.2 as the halo mass is increased only from $10^8$ to $1.7\\times10^8$ $\\Msun$.  The fraction of affected halos only decreases as $\\avenf$ is increased.  Accurate quantitative estimates will have to wait for a break-through in our understanding of the UV emission properties of high-redshift sources, and their dependence on halo mass.  Nevertheless, the inability of photoionization feedback to delay the middle and late stages of reionization is compelling, especially given that we do not include radiative transfer effects which would only decrease the relevance of UV feedback (see \\S~\\ref{sec:ass} for a more detailed discussion).  It seems that delaying the advanced stages of reionization through UV feedback on $M\\gsim10^8\\Msun$ halos is only likely if ionizing efficiencies of $z\\gsim10$ sources are much higher ($\\sim$ two orders of magnitude) than $z\\sim6$ data seem to imply.  An evolution in $\\toteff$ could be obtained by changing the fiducial values in eq. (\\ref{eq:toteff}), possibly as a result of a top-heavy IMF or a star-burst dominated epoch.  However, an increase in the fiducial efficiency, $\\toteff$, by $\\sim$ two orders of magnitude seems unlikely.  Our results also suggest that the natural time-scale for the bulk of reionization is the growth of the collapsed fraction contained in $\\Tvir\\gsim 10^4$ K halos.  The natural timeframe for half of the universe to be ionized in such a prescription, using a toy model for the evolution of the HII filling factor \\citep{FL05, MJH06, HB06, Lidz07} is $\\Delta z_{\\rm re} \\sim$ 2 -- 3, with a small filling factor tail extending to higher redshifts. This tail is governed by more complicated feedback on $\\Tvir \\lsim 10^4$ K objects.  Such a scenario would be consistent with results implied by both the {\\it WMAP} optical depth measurements \\citep{Dunkley08, Komatsu08} and the SDSS quasars \\citep{MH04, WL04_nf, Fan06, MH07}, especially given the inhomogeneous nature of reionization and the associated difficulties in interpreting present data \\citep{LOF06, BH07, Maselli07, BRS07, MF08damp}.  This scenario is also in agreement with recent observational claims that faint galaxies at or below current detection thresholds dominate the ionizing photon budget at $z\\gsim6$ \\citep{YW04b, YW04a, Bouwens06}.  Finally, our results underscore the importance of extended dynamical ranges when modeling reionization.  Simulations must be capable of resolving halos with mass $\\gsim 10^8 \\Msun$, even when modeling the late stages of reionization, while at the same time being large enough to capture HII regions several tens of megaparsecs in size.  \\vskip+0.5in  We are grateful to Steven Furlanetto and Zolt\\'{a}n Haiman for helpful comments on the manuscript. MD is supported by Harvard University Funds."
    },
    "0806/0806.1095_arXiv.txt": {
        "abstract": "High-resolution SiO observations of the NGC 1333 IRAS 4A star-forming region showed a highly collimated outflow with a substantial deflection. The deflection was suggested to be caused by the interactions of the outflow and a dense cloud core. To investigate the deflection process of protostellar outflows, we have carried out three-dimensional hydrodynamic simulations of the collision of an outflow with a dense cloud. Assuming a power-law type density distribution of the obstructing cloud, the numerical experiments show that the deflection angle is mainly determined by the impact parameter and the density contrast between the outflow and the cloud. The deflection angle is, however, relatively insensitive to the velocity of the outflow. Using a numerical model with physical conditions that are particularly suitable for the IRAS 4A system, we produce a column-density image and a position-velocity diagram along the outflow, and they are consistent with the observations. Based on our numerical simulations, if we assume that the initial density and the velocity of the outflow are $\\sim 10 \\cm3$ and $\\sim 100 \\kms$, the densities of the dense core and ambient medium in the IRAS 4A system are most likely to be $\\sim 10^5 \\cm3$ and $\\sim 10^2 \\cm3$, respectively. We therefore demonstrate through numerical simulations that the directional variability of the IRAS 4A outflow can be explained reasonably well using the collision model. ",
        "introduction": "Radio interferometric observations in molecular lines show directional variability of jets in some star forming regions. The understanding of the variability of protostellar jets provides us with clues to a variety of physical processes involved in the interaction of jets and surrounding clouds. To explain the variability of protostellar jet, several models were proposed \\citep[]{em97,gcr99,choi2001a}. These models can be classified into two categories. The first category includes external perturbation models, such as density gradients in the environment, flow instabilities (e.g., Kelvin-Helmholtz instability), side winds, and magnetic fields. The second category includes intrinsic variability models, such as the motion of the jet source perpendicular to the jet direction and the precession of the jet. Moreover, to understand the bending of the jet, especially in the context of young stellar objects (YSOs), numerical simulations have been carried out by several authors \\citep[]{dp99,hurka99,rc95,raga2002}. Three-dimensional smoothed particle hydrodynamic simulations of the penetration of a jet into a dense stratified cloud were carried out by \\citet{dp99}, and it was found that a radius ratio and a density ratio between the jet and the cloud are important parameters in determining the outcome. \\citet{hurka99} studied non-radiative jets bent by magnetic fields with a strong gradient. They showed that a low-velocity jet can be deflected by magnetic fields with a modest strength, but a high-velocity jet needs very strong fields to be deflected. \\citet{rc95} presented analytical and two-dimensional numerical studies of the collision of an Herbig-Halo jet with a dense molecular cloud core. \\citet{raga2002} studied the interaction between a radiative jet and a dense cloud through three-dimensional hydrodynamic simulations and generated H$\\alpha$ and H$_2$ $1-0$ S(1) emission maps.  The NGC 1333 molecular cloud contains numerous YSOs and outflows and is a well-studied star formation region. A highly collimated outflow with a substantial deflection angle was observed from high resolution SiO observations of NGC 1333 IRAS 4A \\citep{choi2005a}. To investigate the overall structure of the IRAS 4A outflow system, \\citet[]{choi2001b,choi2005a,choi2005b} and \\citet{choi2006} observed the outflow in the C$^{18}$O, $^{13}$CO $J=1 \\rightarrow 0$, HCO$^+$, HCN $J=1 \\rightarrow 0$, SiO $v=0$ $J=1 \\rightarrow 0$, and H$_{2}$ $1-0$ S(1) lines. \\citet{choi2005a} suggested that the sharp bend was caused by a collision between the northeastern outflow and a dense core in the ambient molecular cloud. And then, \\citet{choi2005b} confirmed that the obstructing cloud with a dense core is located just north of the bending point in images of the HCO$^+$ and HCN $J=1 \\rightarrow 0$ lines. A sharp bending of the northeastern outflow in the NGC 1333 IRAS 4A system is unusual structure and provides a good example to study the bending mechanism of outflows in protostellar environment.  In this study we have carried out three-dimensional hydrodynamic simulations on the interactions of an outflow with a dense molecular cloud surrounded by an ambient medium, in order to understand the deflection mechanism of the outflow and to construct a detailed model of deflected outflows. We performed many simulations with different sets of physical parameters that can change the outflow direction, then picked up one model that gives a similar deflection angle with the observed one in the IRAS 4A system. Finally, we make a column-density image of shocked gas and a position-velocity diagram along the outflow. They are reasonably consistent with the observed data of the IRAS 4A northeastern outflow. In \\S 2, we describe a numerical setup for the interaction of an outflow and a cloud. Simulation results are presented in \\S 3, followed by the summary and discussion in \\S 4. ",
        "conclusions": "We have performed three-dimensional hydrodynamics simulations to study the interaction of an outflow with a dense molecular cloud. Since the most evident information from observations is the deflection angle, we focus on the deflection angle in the simulations by varying four parameters: the impact parameter, the outflow density contrast with the ambient medium ($\\eta$), the contrast between outflow and cloud density ($\\chi$), and the velocity of outflow. The main results can be summarized as follows. \\begin{itemize} \\item When the impact parameter is smaller than 0.6 times the cloud radius, the outflow bores a hole into the uniform cloud.  On the contrary, the outflow with the impact parameter as small as 0.3 is deflected by the cloud with the power-law type density distribution. \\item When the obstructing cloud has a power-law density distribution, the deflection angle of the outflow is determined by the impact parameter and the density contrast between the outflow and the cloud ($\\chi$). \\item The deflection angle is relatively insensitive to the velocity of outflow. \\item The initially light outflow ($\\eta<1$) can be more collimated than the heavy one ($\\eta>1$). \\end{itemize}  \\citet{choi2005a} suggested that the NGC 1333 IRAS 4A outflow was deflected as a result of a collision with a dense cloud core and provided several lines of evidence. They are (1) the asymmetric morphology of the bipolar outflow, (2) the good collimation and complicated kinematics of the deflected flow, (3) the low-velocity emission from the molecular gas near the bend, and (4) the enhancement of SiO emission in the deflected flow. In this study, we confirmed these evidences through numerical experiments. Especially, the deflection angle ($\\sim 30\\arcdeg$) of collimated outflow and the enhancement and complicated kinematics of shocked gas in the deflected region produced from our numerical simulations are very similar to those seen in the SiO observations of the IRAS 4A region. If we take model PC02 as the most appropriate numerical model for the IRAS 4A outflow, we can infer that this outflow was significantly deflected by the dense molecular core $\\sim 2000$ years ago and the ratio of the outflow density to the density at the cloud center is very low $\\sim 10^{-4}$. Moreover, if we assume that the initial density and the velocity of the outflow are $\\sim 10 \\cm3$ and $\\sim 100 \\kms$, the densities of the dense core and ambient medium in the IRAS 4A system are most likely to be $\\sim 10^5 \\cm3$ and $\\sim 10^2 \\cm3$, respectively.  Although our simulations do not consider the radiative cooling and chemical reactions, our results provide some insights to the interactions of an outflow with a dense cloud. To make a more sophisticated numerical model compatible with the SiO observations, we need to include chemistry, especially related to the SiO molecule. In the near future, we will execute numerical simulations which include the radiative cooling and chemical reactions related to the SiO molecule and then try to directly compare the simulated SiO emission image with the observed one."
    },
    "0806/0806.4679_arXiv.txt": {
        "abstract": "A knowledge of the recombination time on the grain surfaces has been a major obstacle in deciding the production rate of molecular hydrogen and other molecules in the interstellar medium. We present a numerical study to compute this time for molecular hydrogen for various cloud and grain parameters. We also find the time dependence, particularly when a grain is freshly injected into the system. Apart from the fact that the recombination times seem to be functions of the grain parameters such as the activation barrier energy, temperature etc, our result also shows the dependence on the number of sites in the grain $S$ and the effective accretion rate per site $a_s$ of atomic hydrogen. Simply put, the average time that a pair of atomic hydrogens will take to produce one molecular hydrogen depends on how heavily the grain is already populated by atomic and molecular hydrogens and how fast the hopping and desorption times are. We show that if we write the average recombination time as $T_r \\sim S^\\alpha/A_H$, where, $A_H$ is the hopping rate, then $\\alpha$ could be much greater than $1$ for all astrophysically relevant accretion rates. Thus the average formation rate of $H_2$ is also dependent on the grain parameters, temperature and the accretion rate. We believe that our result will affect the overall rate of the formation of complex molecules such as methanol which require successive hydrogenation on the grain surfaces in the interstellar medium. ",
        "introduction": "It has long been suggested that the dust grains play a major role in the formation of molecular hydrogen in the interstellar medium (ISM) (Gould \\& Salpeter, 1963). Considerable studies were made since then to understand the real physical processes which are taking place both theoretically (e.g., Hollenbach, Werner \\& Salpeter, 1971; Takahashi, Matsuda \\& Nagaoka, 1999; Biham et al. 2001) as well as experimentally (e.g., Pirronello et al. 1997a,b, 1999). More recently, Biham et al. (2001), and Green et al. (2001) have computed $H_2$ production rate by physisorption. It was found that a significant production is possible in cooler ($\\sim 10-25$K) clouds. Cazaux \\& Tielens (2002, 2004) use both physisorption and chemisorption, to demonstrate that $H_2$ production is possible at high temperatures ($\\sim 200-400$K) also. The goal is to study the rate at which the $H$ atoms combine together on the surface of the grains to form $H_2$ and then they are desorbed into the gas phase to react with other atoms. When compared with the average mass fractions of various molecular species obtained through gas phase reactions (see, Chakrabarti \\& Chakrabarti, 2000ab; Das et al. 2006), it was found that the observed abundances of more complex species, such as methanol, are much higher. It is possible that methanol as well as its precursors also have to be formed on grain surfaces through successive hydrogenation. Our finding for molecular hydrogen has thus important bearings on the formation of more complex molecules on grains. These molecules would, in turn, desorb into the gas phase and would be expected to produce more complex species such as amino acids in due course.  One of the most challenging problems is to determine the average rate at which the recombination of atomic hydrogen takes place on a grain surface. In theoretical investigations which are prevalent in the subject (See, Acharyya and Chakrabarti, 2005; hereafter Paper I;  Acharyya, Chakrabarti and Chakrabarti, 2005; hereafter ACC05), the diffusion rate $A_H$ (inverse of the diffusion time $T_d =1/A_H$) is divided  by $S$, the number of sites on the grain surface (e.g., Biham et al. 2001) to get the recombination rate. The argument for reducing the rate by a factor of $S$ is this: on an average, there are $S^{1/2}$ number of sites in each direction of the grain. Since the hopping is random, it would take square of this, i.e., $S$ number of hopping to reach a distance located at $S^{1/2}$ sites away, where, on an average, another $H$ is available. Thus, the effective recombination rate was chosen to be $A_H/S$. It is an empirical factor and needs more careful treatment. In our present paper, we replace $S$ by an 'unknown' quantity $S^\\prime = S^\\alpha(t)$, where $\\alpha(t)$ may be time dependent (if the grain and cloud parameters change) and it could also deviate from unity. Let $\\alpha=\\alpha_0$ when a steady state is reached. Higher the accretion rate, lesser should be the value of $\\alpha_0$ as the effective surface area $S^\\prime$ $(t \\rightarrow \\infty)$ gets smaller and smaller. The opposite is true for smaller accretion rates. In fact, in the limit, if the accretion rate is so low, that a lone $H$ sweeps around the grain several times to find another $H$, one would get $\\alpha_0>1$ since the effective site number is higher than $S$. Using our simulation, we determine how the effective site number deviates from $S$, one way or the other when the accretion rate is varied. Our result is likely to have important consequences for the formation of other hydrogenated species, such as water, methanol on grain surfaces. This will be discussed elsewhere. Some preliminary results with steady state $\\alpha_0$ have been presented in Das et al. (2005) and  Chakrabarti et al. (2006). In the current paper, we discuss the time and temperature dependence of $\\alpha (t)$ and studied the cases for more varied astrophysically important accretion rates.  In the next Section, we present the modified typical equations which govern the molecular hydrogen production rate on a grain surface. Incorporating the {\\it physical aspects} of these equations, we perform a numerical simulation to determine the numbers on the grains. In \\S 2, we present the procedure for the simulation and in \\S 3 we present the results for two types of commonly used grains, namely, olivine and amorphous carbon. We show how $\\alpha(t)$ depends on time and how it settles into a number (generally, $>1$) when steady state is reached. Finally, in \\S 4, we present our concluding remarks. ",
        "conclusions": "Fig. 1 gives the variation of $n_H$ and $n_{H2}$ residing on the grain surface since the beginning of the simulation. We chose a grain with $10^4$ sites at $8$K. For clarity, we plot average numbers in every $\\sim 5.5 \\times 10^6$s bin after the initial transient period of $\\sim 2 \\times 10^6$s is over. The simulation was carried out near about  $10^{10}$s. The effective accretion rate is assumed to be $\\phi_H=7.98 \\times 10^{-4}$s$^{-1}$. In this case, the steady state has clearly been reached by this time and we can compute $\\alpha_0$ by using Eq. 4. This is done below.  \\begin {figure} \\vskip 0.5cm \\includegraphics[width=5.1cm]{fig1.eps} \\centering \\vskip 0.5cm \\caption{: Variation of the number of H and $H_2$ on an olivine grain of $10^4$ sites kept at $8$K and exposed to an effective accretion rate per site of H $7.98 \\times  10^{-8}$ $s^{-1}$. We carried out our simulation near about $10^{10}$s. For cleanliness, after the initial transient period of about $2 \\times 10^6$s is over, we took time average at every $5.5\\times 10^6$s before plotting the numbers. } \\end {figure}  \\subsection {Olivine grains} We start with the olivine grains which are kept at 8 K. The evaporation rate of $H_2$ is $W_{H_2}= 0.834844 \\times 10^{-5}$ s$^{-1}$ which corresponds to a time scale of $\\sim 119782.8$s and similarly the evaporation rate of $H$ is $W_H=0.59 \\times 10^{-8}$s$^{-1}$ which corresponds to $\\sim 169585908$s. The hopping time for hydrogen on this grain is $3680.58$s (data taken from Katz et al. 1999).  In Fig. 2a, we present the computed $\\alpha_0$ (Eq. 4) as a function of $a_s$ -- the effective accretion rate per site. The solid, dot-dashed and the dashed curves are for $S=10^4$, $9\\times 10^4$ and for $10^6$ sites respectively. No spontaneous desorption has been included (i.e., $\\mu=1$). We extrapolated our curve to very low accretion rates which would have taken a very long computation time, just to show the trend of the result and extreme conditions. We note that $\\alpha_0$ is generally higher than unity in the region of our interest. This is because $H$s are scattered few and far between and it takes a longer time (generally more than one sweeping) for one $H$ to locate another. $\\alpha_0$ monotonically drops as the accretion rate goes up. In Fig. 2b, we show the variation of $\\beta_0$ (Eq. 6) for the same case. Here too, we see that $\\beta_0$ is very high compared to unity for low rates, but becomes $\\sim 0.5$ or lower for higher rates as expected. Note that $\\alpha_0$ and $\\beta_0$ go down with increasing site number $S$ also. Since for a smaller grain, the possibility of getting it filled at a high rate is higher, one would have expected an opposite result.  However, it is to be remembered that for a larger grain, the accretion rate itself ($\\phi_H=S a_s$) is also large. Hence the plots are to be compared carefully. For instance, the result of $a_s=0.0005\\times 10^{-8}$$s^{-1}$ for $9\\times 10^4$ sites is to be compared with that of $a_s=0.0045\\times 10^{-8}$s$^{-1}$ for $10^4$ sites in order to make a meaningful comparison. In any case, for reasonable $\\phi_H$ values with number densities up to $10^6$ cm$^{-3}$, the relevant $a_s$ would be below $10^{-6}$ where $\\alpha_0>1$ in general.  \\begin {figure} \\vskip 0.5cm \\includegraphics[width=5.1cm]{fig2a.eps} \\centering \\vskip 0.7cm \\includegraphics[width=5.1cm]{fig2b.eps} \\vskip 0.5cm \\caption{\\small{\\bf (a-b)}: Variation of (a) $\\alpha_0$ and (b) $\\beta_0$ as a function of $a_s$, the effective accretion rate per site for various olivine grains kept at $8$K. The dashed, dot-dashed and solid curves are for $S=10^6,\\ 9\\times 10^4, and \\ 10^4$ respectively. $\\alpha_0$ is clearly a function of the accretion rate. For rates relevant in molecular clouds $\\alpha_0$ and $\\beta_0$ are much larger than unity. For very high rates $\\beta_0$ comes down to $0.5$ or lower. The deviation is highlighted using dotted curves by extrapolating at very low accretion rates. } \\end{figure}  \\begin {figure} \\includegraphics[width=12.7cm]{fig3.eps} \\centering \\vskip -.9cm \\caption{\\small{\\bf (a-d)}: In (a-b) snapshots of the grain surface with $H$ (hollow squares) and $H_2$ (filled squares) at two arbitrarily chosen times (a) $ 8 \\times 10^8$s and at (b)$ 10^9$s. Here an olivine grain (at $8$K) with $900$ sites has been chosen. This is bombarded with an accretion rate per site of H $3.02 \\times 10^{-7}$ per sec. No spontaneous desorption has been assumed here. In (c-d) spontaneous desorption has been included and plotted for the same time as before. Thus, numbers of $H_2$ residing on the grain at any instant are lesser.}  \\end{figure}    \\begin{figure} \\vskip 0.5cm \\includegraphics[width=5.1cm]{fig4.eps} \\centering \\vskip 0.5cm \\caption{: Temperature dependence of $\\alpha_0$ for the olivine grains at $10$K (solid), $9$K (dot-dashed) and $8$K (dashed). The deviation is highlighted using dotted curves by extrapolating at very low accretion rates.} \\end{figure}  \\begin{figure} \\vskip 0.5cm  \\includegraphics[width=5.1cm]{fig5.eps} \\centering \\vskip 0.5cm \\caption{: A comparison between the recombination efficiency obtained from the rate equation (solid) and that obtained from our simulation (dashed). We use the accretion rate per site $1.8 \\times 10^{-9}$ per second for a grain of diameter $0.1 \\mu$m. The difference can be attributed to the temperature dependence of the $\\alpha_0$ as shown in Fig. 4 above.} \\end{figure}  \\begin{figure} \\vskip 0.5cm \\includegraphics[width=5.1cm]{fig6.eps} \\centering \\vskip 0.5cm \\caption{: A comparison of the simulation results (dark circles) with those obtained from analytical considerations (dashed curves) when suitable modification of the average recombination rate is made. An Olivine grain of $10^4$ sites at a temperature of $8$K has been chosen in this comparison. Dotted curves are drawn using analytical results for $\\alpha_0$ extrapolated to very low accretion rates.} \\end{figure}  In Fig. 3(a-b), we present snapshots of the occupancy of $H$ (hollow squares) and $H_2$ (filled squares) at two instants of time on the grain containing only $900$ sites at two arbitrarily chosen times (a) t=$8 \\times 10^8$s (b) t=$10^9$s respectively. Here an olivine grain at ($8$K) with $900$ sites has been chosen. This grain is facing an accretion rate of $F_H/S$ of $3.02 \\times 10^{-7}$s$^{-1}$ per site. No spontaneous desorption has been assumed here. This is to be compared with Fig. 3(c-d) for the same simulation and for the same time as before when the spontaneous desorption has been included. The numbers of $H_2$ are fewer since after formation of $H_2$ molecules some part of the $H_2$ are spontaneously desorbed in to the gas phase.  When temperatures of the grain is increased, all the rates go down exponentially. As a result, we expect $\\alpha_0$ to rise with temperature for a given accretion rate. In Fig. 4 we show this behaviour for $T=8, \\ 9, \\ 10$K respectively for olivine grains. We choose $S=10^4$ in this case. This behaviour affects the recombination efficiency $\\eta$ as defined by, $$ \\eta=\\frac{2R_{H_2}}{F_H} \\eqno{(7)} $$ itself. In Fig. 5, we compare the temperature dependence of $\\eta$ as obtained from the rate equations (solid) curve with that obtained from our simulations. The accretion rate per site of $1.8 \\times 10^{-9}$ sec$^{-1}$ was used (same in both the cases). We note that for  $T\\lsim7.5$K  the simulation results are higher and for $T\\gsim 7.5$K the simulation results are lower. This is because $\\alpha_0$ itself is strongly temperature dependent as shown in Fig. 4 while the rate equation uses $S=S^\\prime$  (i.e., $\\alpha_0=1$) for all temperatures.  In Fig. 6 we compare our results with the analytically obtained results from the rate equation method provided correct $S^\\prime$ was chosen (Eqs. 2ab). The simulation results are shown by the dark circles and those obtained from the analytical considerations are shown by the dashed curves. An Olivine grain of $10^4$ sites at a temperature of $8$K has been chosen in this comparison. Dotted curves are drawn using analytical results for $\\alpha_0$ extrapolated to very low accretion rates.  It is interesting to compare the results of our simulation with those obtained from the analytical considerations with and without our $\\alpha_0$ factor. In Table 1, we present this comparison. We take an Olivine grain of $10^4$ sites at $8$K and vary the accretion rates. In Column 1, we give the accretion rate per site of the grain. In Column 2 we present the coefficient $\\alpha_0$ which we derive from our simulation. In Columns 3-5, we present the number of $H$ as obtained by our simulation and the modified equation (Eq. 2a) and the standard equations (Paper I) respectively. Columns 6-8, we present similar results for $H_2$. We find that our simulation matches more accurately with the analytical results provided $S^\\prime$ is chosen as the surface area. If the standard equation is used, the deviation is very significant. Indeed, the number of $H$ on the grain could be roughly half as much when simplistic analytical model is used. What observe is that on the grains we tend to have more $H$ and less $H_2$ than what analytical work suggests.  \\begin{table} \\caption{\\label{table} Comparison of $H$ and $H_2$ abundances in various methods} \\vskip 0.5cm \\hskip -1.2cm \\begin{tabular}{|ll|lll|lll|} \\hline Accretion Rate&  $\\alpha_0$ &  & $H$ with &  &   & $H_2$ with  &  \\\\ \\cline{3-5} \\cline{6-8} per site  $A_s(S^{-1})$ & & simulation &  $\\alpha_0\\ne 1$ & $\\alpha_0=1$ & simulation & $\\alpha_0\\ne 1$& $\\alpha_0=1$\\\\ \\hline\\hline  $  6.79 \\times 10^{-7}$ &  1.04 & 403.11&  407.80&  340.57&  361.71&  383.82 & 377.48\\\\ $  2.72  \\times 10^{-7}$&  1.05&  275.78&  278.02&  219.31 & 149.96&  158.16&  156.53\\\\ $ 2.72 \\times 10^{-8}$ & 1.07 & 99.03 & 99.31 & 70.35 & 15.40 & 16.14 & 16.11\\\\ $  1.36 \\times 10^{-8}$ & 1.08 & 72.27 & 72.42&  49.80&  7.71 & 8.08 & 8.07\\\\ $  5.43 \\times 10^{-9}$& 1.09 & 47.30&  47.37 & 31.52&  3.11 & 3.23 & 3.23\\\\ $  2.72 \\times 10^{-10}$&  1.11 & 11.62&  11.62&  7.01&  $1.49 \\times 10^{-1}$ & $1.59 \\times 10^{-1}$ & $1.60 \\times 10^{-1}$\\\\ $  1.09\\times 10^{-10}$ & 1.12 & 7.46&  7.48&  4.42 & $5.81 \\times 10^{-2}$ &  $6.24 \\times 10^{-2}$ & $ 6.35\\times 10^{-2}$\\\\ $  2.72 \\times 10^{-11}$&  1.13 & 3.86 & 3.82 & 2.18 & $1.31 \\times 10^{-2}$ & $1.49 \\times 10^{-2}$ & $1.55 \\times 10^{-2}$\\\\ $^*  2.04\\times 10^{-11}$ & 1.13 & 3.36&  3.33 & 1.89 & $9.7 \\times 10^{-3}$ & $1.11 \\times 10^{-2}$ & $1.17 \\times 10^{-2}$\\\\ $^*  2.92\\times 10^{-12}$ & 1.14 & 1.32&  1.30 & 0.70 &  $1.3 \\times 10^{-3}$&  $1.5 \\times 10^{-3}$ & $1.6 \\times 10^{-3}$\\\\ $^*  4.16\\times 10^{-13}$ & 1.16& 0.52 & 0.51&  0.26 &  $2.00 \\times 10^{-4}$&  $2.00 \\times 10^{-4}$&  $2.00 \\times 10^{-4}$\\\\ $^*  5.95\\times 10^{-14}$ & 1.17 & 0.20 & 0.20&  0.10  & $2.11 \\times 10^{-5}$&$  2.81\\times 10^{-5}$& $ 3.13 \\times 10^{-5}$\\\\ \\hline\\hline \\end{tabular} \\vskip 0.3cm $*$ represents the extrapolated value. \\end{table}  \\subsection {Amorphous carbon grains}  Carbon grains produce significant $H_2$ at a higher temperature than olivine grains because all the barrier energies are higher. We plot in Fig. 7a the variation of $\\alpha_0$ with $a_s$ at temperature $14$K. The nature of variation of $\\alpha_0$ remains the same, namely, $\\alpha_0$ goes down with $a_s$. The solid, dot-dashed and the dashed curves are for  $S=10^6$, $9\\times 10^4$ and for $10^4$ sites respectively. In Fig. 7b, we show the variation of $\\beta_0$ and as expected its value can become as low as $0.5$ for very large accretion rate. These results are representative as they are strongly temperature dependent as in the case of olivine (see, Fig. 4).  \\begin{figure} \\vskip 0.5cm \\includegraphics[width=5.1cm]{fig7a.eps} \\centering \\vskip 0.7cm \\includegraphics[width=5.1cm]{fig7b.eps} \\centering \\vskip 0.5cm \\caption{\\small{\\bf(a-b)}: Variation of (a) $\\alpha_0$ and (b) $\\beta_0$ as a function of $a_s$, the effective accretion rate per site, for various amorphous carbon grains kept at $14$K. The solid, dot-dashed and dashed curves are for $S=10^4, \\ 9\\times 10^4, \\ 10^6$ respectively The exponents $\\alpha_0$ and $\\beta_0$ are strong functions of the accretion rate. For low rates $\\beta_0$ is higher than unity, while for higher rates it is close to $0.5$ or even lower. The deviation is highlighted using dotted curves by extrapolating at very low accretion rates.} \\end{figure}"
    },
    "0806/0806.4353_arXiv.txt": {
        "abstract": "With 40 or more transiting exoplanets now known, the time is ripe to seek patterns and correlations among their observed properties, which may give important insights into planet formation, structure, and evolution. This task is made difficult by the widely different methodologies that have been applied to measure their properties in individual cases. Furthermore, in many systems our knowledge of the planet properties is limited by the knowledge of the properties of the parent stars. To address these difficulties we have undertaken the first comprehensive analysis of the data for 23 transiting planets using a uniform methodology. We revisit several of the recently proposed correlations, and find new ones involving the metallicity of the parent stars. ",
        "introduction": "Most of our knowledge about the structure, atmospheric properties, and other physical characteristics of extrasolar planets has come from the study of those that transit their parent stars. The pace of discovery of transiting planets has increased rapidly over the last year or so, and at the time of this writing there are nearly 40 systems with detailed studies in the literature, along with several more that have been announced recently.  The time is ripe to seek patterns and correlations among their observed properties, which may give important insights into planet formation, structure, and evolution.  Several such relations have already been proposed. Unfortunately, our ability to gauge their reliability or to find new ones is made difficult by the widely different methodologies that have been applied by individual investigators to measure the properties of the planets and their parent stars. Furthermore, in many cases our knowledge of the planet properties is limited by the knowledge of the properties of the stars themselves, as surprising as this may seem. The latter properties are usually determined with the help of stellar evolution models, but not always have the best constraints been applied consistently. In particular, for the majority of transiting systems without a parallax determination, the weakly constrained surface gravity of the star determined spectroscopically has often been used as a proxy for luminosity. A much better constraint related to the mean stellar density is available directly from the light curves (\\cite[Sozzetti et al.\\ 2007]{Sozzetti:07}), but has generally been overlooked.  To address these difficulties we have undertaken the first comprehensive analysis of the data for 23 transiting planets using a uniform methodology (\\cite[Torres et al.\\ 2008]{Torres:08}). We describe our procedures here, along with a few highlights of our findings. ",
        "conclusions": ""
    },
    "0806/0806.1740_arXiv.txt": {
        "abstract": "Using analytic arguments and numerical simulations, we examine whether chondrule formation and the FU Orionis phenomenon can be caused by the burst-like onset of gravitational instabilities (GIs) in dead zones.  At least two scenarios for bursting dead zones can work, in principle.  If the disk is on the verge of fragmention, GI activation near $r\\sim4$ to 5 AU can produce chondrule-forming shocks, at least under extreme conditions.  Mass fluxes are also high enough during the onset of GIs to suggest that the outburst is related to an FU Orionis phenomenon.  This situation is demonstrated by numerical simulations.  In contrast, as supported by analytic arguments, if the burst takes place close to $r\\sim1$ AU, then even low pitch angle spiral waves can create chondrule-producing shocks and outbursts.  We also  study the stability of the massive disks in our simulations against fragmentation and find that although disk evolution is sensitive to changes in opacity, the disks we study do not fragment, even at high resolution and even for extreme assumptions. ",
        "introduction": "\\subsection{GIs, MRI, FU Orionis Outbursts, and Chondrules}  Gravitational instabilities can activate in a disk when the Toomre (1964) parameter $Q={c_s\\kappa}/{\\pi G\\Sigma}\\lesssim 1.7$ for a thick disk (see Durisen et al.~2007), where $c_s$ is the sound speed, $ \\Sigma$ is the surface density, and $\\kappa$ is the epicyclic frequency.   As indicated by $Q$, a disk is unstable against GIs when it is cold and/or massive.  The resulting spiral waves driven by self-gravity efficiently transfer angular momentum outward and mass inward (e.g., Lynden-Bell \\& Kalnajs 1972; Durisen et al.~1986).  Another mechanism that can efficiently transfer angular momentum outward is the magnetorotational instability (MRI; see Balbus \\& Hawley 1991; Desch 2004).  In contrast to GIs, the MRI only requires a weak magnetic field coupled to the gas.  These mechanisms, either separately or in combination, are likely to be the principal way T Tauri stars accrete gas from a disk (Hartmann et al.~2006).  In order for the MRI to occur, ionized species must be present in the gas phase.   Thermal ionization of alkalis occurs wherever $T\\gtrsim1000$ K, but depletion of ions by dust grains may move the temperature threshold closer to $T\\sim1700$ K (Desch 1999; Sano et al.~2000), where the dust sublimates completely. Elsewhere, the ionization must be driven by a nonthermal source, e.g., energetic particles (EPs).  For this discussion, EPs refers to any particles that could ionize the gas, e.g., X-rays.  Gammie (1996) proposed that disks may have active and inactive MRI layers due to attenuation of EPs by the gas.  In the inner regions of a disk  (at radii $\\sim$ few AU) where the column densities are large, MRI may only be active in a thin layer, resulting in {\\it layered accretion}.  As one moves outward, column densities drop, and the entire disk can become MRI active.  The region where the MRI is mostly absent is called the {\\it dead zone}.   EPs are attenuated by a surface density of only about 100 g cm$^{-2}$ (Stepinski 1992), and so even a minimum mass solar nebula (MMSN) will likely exhibit layered accretion (Desch 2004).  Even if mass accretion is only reduced and not altogether halted as a result of Reynolds stresses (Fleming \\& Stone 2003; Oishi et al.~2007), mass may still pile up in the dead zone.  If enough mass accumulates, then even for an otherwise low-mass disk, GIs can activate.  The FU Orionis phenomenon is characterized by a rapid (1-10s yr) increase in optical brightness of a young T Tauri object, typically by 5 magnitudes, and is driven by sudden mass accretion of the order $10^{-4} M_{\\odot}~\\rm yr^{-1}$ from the inner disk onto the star (Hartmann \\& Kenyon 1996). Because FU Ori objects appear to have decay timescales of about 100 yr, the entire mass of a MMSN ($\\sim0.01~M_{\\odot}$) can be accreted onto the star.  To date, the best explanation for the optical outburst is a thermal instability (e.g., Bell \\& Lin 1994; Kley \\& Lin 1999; see discussions in Hartmann \\& Kenyon 1996, Green et al.~2006; and Zhu et al.~2007). Armitage et al.~(2001) suggested that GIs in a bursting dead zone might be able to trigger an FU Ori outburst by rapidly increasing the accretion into the inner disk ($\\lesssim 0.1$ AU) and initiating an MRI through thermal ionization. Hartmann (2007, in private communication) and Zhu et al.~(2007) also suggest that the heating due to gravitational torques might drive an MRI for $r>0.1$ AU, which would then feed mass inside 0.1 AU until a thermal instability sets in.   The FU Ori phenomenon may be a result of a {\\it cascade} of instabilities, starting with a burst of GI activity in a dead zone, followed by accretion due to an MRI, followed finally by a thermal instability (cf Kley \\& Lin 1999).  Indeed, recent observations of FU Ori indicate that very large mass fluxes are present out to at least $r\\sim0.5$ AU (Zhu et al.~2007).  Although details are still being debated, most chondrules formed in the first 1 to 3 Myr of the Solar Nebula's evolution (Bizzarro et al.~2004; Russell et al.~2005).   Chondrule precursors were flash melted from solidus to liquidus, where high temperatures $T\\sim1700$ K were experienced by the precursors for a few minutes. The melts then cooled over hours, with the actual cooling time depending on chondrule type. Chondrule collisional histories and isotopic fractionation data, chondrule-matrix complementarity, fine-grained rim accumulation, and petrological and parent body location arguments  (Krot et al.~2005) suggest that chondrules formed in the Solar Nebula (Wood 1963) in strong, localized, repeatable heating events.  The shock wave model for chondrule formation can accommodate these observational constraints and reproduce heating and cooling rates required to form chondrule textures (Iida et al.~2001; Desch \\& Connolly 2002; Cielsa \\& Hood 2002; Miura \\& Nakamoto 2006). One plausible source of chondrule-producing shocks is a global spiral wave (Wood 1996).  Harker \\& Desch (2002) suggest that spiral waves could also explain thermal processing at distances as large as 10 AU, which may be necessary to explain  observations of comets (Wooden et al.~2005) and recent {\\it Stardust} results (e.g., McKeegan et al.~2006). It has been suggested that bursts of GIs may be able to produce the required shock strengths  (Boss \\& Durisen 2005) and provide a source of turbulence and mixing (Boss 2004b; Boley et al.~2005).  Global spiral shocks are appealing because they fit many of the constraints above.  They may be repeatable, depending on the formation mechanism for the spiral waves; they are global, but produce fairly local heating; they can form chondrules in the disk; and they can work in the inner disk as well as the outer disk.  \\subsection{Fragmentation}  Knowing under what conditions protoplanetary disks can fragment is crucial to understanding disk evolution inasmuch as a fragmented disk may produce gravitationally bound clumps.  This has become known as the {\\it disk instability}  hypothesis  for the formation of gas giant planets (Kuiper 1951; Cameron 1978; Boss 1997, 1998).   The strength of GIs is regulated by the cooling rate in disks (Tomley et al.~1991, 1994; Pickett et al.~1998, 2000, 2003), and if the cooling rate is high enough in a low-$Q$ disk, a disk can fragment (Gammie 2001). Gammie quantified that a disk will fragment when $t_{\\rm cool}\\Omega\\lesssim 3$ for a disk with a  $\\Gamma=2$, where $\\Gamma$ is the two-dimensional adiabatic index, such that  $\\int p dz = P\\sim\\Sigma^{\\Gamma}$, where $p$ is the gas pressure and $z$ is the vertical direction in the disk. Here,  $t_{\\rm cool}$ is the local cooling time and $\\Omega$ is the angular speed of the gas.  This criterion was approximately confirmed in 3D disk simulations by Rice et al.~(2003) and Mej\\'ia et al.~(2005).  Rice et al.~(2005) showed through 3D disk simulations that this fragmentation criterion depends on the 3D adiabatic index $\\Gamma_1$ and, for $\\Gamma_1=\\gamma=5/3$ or 7/5, the fragmentation limit occurs when $t_{\\rm cool}\\Omega\\lesssim6$ or 12, respectively.  These results show that a change by a factor of about 1.2 in $\\gamma$ has a factor of two effect on the critical cooling time.  In addition, these results indicate that the cooling time must be roughly equal to the dynamical time of the gas for the disk to be unstable against fragmentation when $\\gamma=5/3$.  Do such prodigious cooling rates occur in disks when realistic opacities are used with self-consistent radiation physics? This question is heavily debated in the literature (e.g., Nelson et al.~2000; Boss 2001, 2004a, 2005; Rafikov 2005, 2007; Boley et al.~2006, 2007b; Mayer et al.~2007; Stamatellos \\& Whitworth 2008).  The simulations to date use a wide variety of numerical methods, including very different approximations for radiation physics, and only two groups have published results of radiative transfer tests appropriate for disks (Boley et al.~2007b; Stamatellos \\& Whitworth 2008).  Nelson et al.~(2000) used 2D SPH simulations with radiation physics to study protoplanetary disk evolution.  Because their simulations were evolved in 2D, they assumed that the disk at any given moment was in vertical hydrostatic equilibrium.  Using a polytropic vertical density structure and Pollack et al.~(1994) opacities, they cooled each particle according to an appropriate effective temperature.  In their simulations, the cooling rates are too low for fragmentation.   In contrast, Boss (2001, 2005) employed radiative diffusion in his 3D grid-based code; fragmentation occurs in his simulated disks.  Besides the difference in dimensionality of the simulations, Boss assumed a fixed temperature structure for Rosseland mean optical depths less than 10, as measured along the radial coordinate (cf recent flux-limited simulations by Boss 2008). Boss (2002) found that the fragmentation in his disks is insensitive to the metallicity of the gas and attributed this independence to fast cooling by convection (Boss 2004a).  However, it must be noted that Nelson et al.~(2000) assumed a vertically polytropic density structure.  Because the specific entropy $s\\sim \\ln K$, where $p=K\\rho^{\\gamma}$ is the polytropic equation of state and $\\rho$ is the gas density, the Nelson et al.~approximation effectively assumes efficient convection.  Except for extremely massive and extended disks (Stamatellos et al.~2007), recent simulations with radiative physics by Cai et al.~(2006), Boley et al.~(2006, 2007b), and Stamatellos \\& Whitworth (2008) find long cooling rates and no efficient cooling by convection.  Cai et al.~also show that the strength of GIs is dependent on the metallicity.  Furthermore, Boley \\& Durisen (2006) suggest that shock bores, which can cause a rapid vertical expansion in the post-shock region of spiral shocks, could be misidentified as ``convective'' flows.  In both contrast and support of these studies, Mayer et al.~(2007) use SPH simulations with 3D flux-limited diffusion and find that fragmentation only occurs once the mean molecular weight of the gas $\\mu\\gtrsim2.7$. However, the simulations presented by Mej\\'ia (2004), Cai et al.~(2006), and Boley et al.~(2006) unintentionally were evolved with $\\mu\\approx2.7$ due to an error in the inclusion of He in the opacity tables, and their disks do not fragment.  The issue of fragmentation in radiatively cooled disks thus remains unsettled.  \\subsection{Current Study}  For this study, we adopt the hypothesis, which we refer to as the {\\it unified theory}, that bursts of GI activity in dead zones drive the FU Ori phenomenon and produce chondrule-forming shocks.  This hypothesis is an amalgamation  of ideas presented in Wood (1996, 2005), Gammie (1996), Armitage et al.~(2001), and Boss \\& Durisen (2005). In order to investigate this scenario, we designed a numerical experiment to evolve a massive, highly unstable disk with an initial radial extent between 2 and 10 AU.  Commensurately, we investigate the stability of these massive, gravitationally unstable disks against fragmentation to assess the feasibility of the disk instability hypothesis for gas giant formation. ",
        "conclusions": "In this section, we review the implications of the results of this study for disk fragmentation, the effects of opacity on disk cooling, the FU Ori phenomenon, and chondrule formation.  We remind the reader that the simulations presented here are meant to be a numerical experiment that explores the possible connection between chondrules, FU Ori outbursts, and bursts of GI-activity.  Moreover, this experiment provides a systematic study of the effects of opacity on disk cooling, which complements the Cai et al.~(2006) metallicity study and provides a test bed for disk fragmentation criteria.   \\subsection{Fragmentation}  After the onset of the burst, none of the disks fragments, and only the 512 $10^{-4}$ $\\kappa$ simulation shows dense knot formation during the peak of the burst (10-30 yr).  These knots do not break from the spiral wave even in the 1024 simulation, and so clump formation does not seem to be missed due to poor resolution. One should also keep in mind that for these simulations, the disk is first relaxed to equilibrium with standard opacity and then the opacity is suddenly dropped by a factor of 10$^{4}$ at $t=0$.  In effect, the dust settling is treated as instantaneous.  Knot formation may not occur in a disk with more realistic settling timescales.  As discussed below, this is also a caveat for the chondrule formation results.  On the other hand, it does suggest that disk fragmentation might occur inside 10 AU under even more extreme, but perhaps unphysical, conditions than we have modeled.  The stability of these disks against fragmentation is supported by Figures 13 through 15.  The cooling rates for the standard and $10^{-2}~\\kappa$ simulations are too low to cause disk fragmentation.  The $10^{-3}$ and  $10^{-4}~\\kappa$ simulations do have areas where $\\zeta = t_{\\rm cool}\\Omega/f(\\Gamma_1) \\lesssim 1$, but the disks approach stability against GIs in those regions ($Q\\gtrsim1.7$; Durisen et al.~2007).   As noted in \\S 4.2, $\\zeta \\lesssim 1$ is not a strict instability criterion, but it does serve as an estimate for disk stability against fragmentation.  For no simulation is $\\zeta$ well below unity where $Q$ is also $\\lesssim1.7$.  Lowering the opacity increases the cooling rates.  The $10^{-4}$ $\\kappa$ simulations exhibit the most rapid cooling because the midplane optical depths are near unity, which results in the most efficient radiative cooling possible. {\\it Changes in dust opacity have a profound effect on disk cooling}.  Although not modeled, if the opacity were to continue to drop such that the midplane optical depth becomes well below unity, cooling would once again become inefficient.  In addition, supercooling of the high optical depth disks (standard and $10^{-2}~\\kappa$) by convection does not occur.  Based on our results, we believe that hydraulic jumps as a result of shock bores (Boley \\& Durisen 2006) rather than convection are a better explanation for the upwellings around spiral arms reported by Boss (2004a) and Mayer et al.~(2007).   These findings are consistent with analytic arguments by Rafikov (2005, 2007), with numerical studies of disk fragmentation criteria by Gammie (2001), Johnson \\& Gammie (2003), and Rice et al.~(2005), and with global disk simulations where tests of the radiation algorithm and/or careful monitoring of radiative losses were performed (Nelson et al.~2000; Boley et al.~2006, 2007b; Stamatellos \\& Whitworth 2008).  The radiative algorithm used in CHYMERA has passed a series of radiative transfer tests that are relevant for disk studies, including permitting convection when expected.  Our conclusions regarding fragmentation are based on mulitple analyses: surface density rendering (Figures 4-7), an energy budget analysis (Figures 9 and 10), cooling temperature and brightness maps (Figures 11-12 ), and a cooling time stability analysis (Figures 13-15).   It is also important to point out that these results do not contradict Stamatellos et al.~(2007) or Krumholz et al.~(2007), who find fragmentation in massive, extended disks ($>100$ AU).  It is also pertinent to demonstrate that the CHYMERA code can detect fragmentation when cooling rates are high and $Q$ is low.  Figure 17 shows a snapshot for a simulation similar to the 512 $10^{-4}~\\kappa$ simulation, but with the divergence of the fluxes artificially increased by a factor of two.  In the normal simulation, kinks in the spiral waves form during the onset of the burst. As discussed above, the disk is very close to the fragmentation limit, but the knots do not break from the spiral wave.  Figure 15 suggests, although for a later time, that increasing the cooling rates by a factor of two would drop $\\zeta$ well below unity in low-$Q$ regions. Figure 17 confirms that the disk fragments with such enhanced cooling.   However, we remind the reader that the $10^{-4}~\\kappa$ simulation has the optimal optical depth for cooling ($\\tau\\sim1$).  We cannot imagine any physical process that could cause a factor of two enhancement in cooling.  Three clumps form, one for each spiral wave, between 4 and 5 AU.  The location of clump formation is consistent with the prediction by Durisen et al.~(2008) that a spiral wave is most susceptible to fragmentation near corotation.  One of the clumps survives for several orbits and eventually passes through the inner disk boundary.  Resolution is always a concern for simulations.  Because these results are consistent with analytic fragmentation limits and numerical fragmentation experiments, we conclude that CHYMERA can detect fragmentation at the resolutions employed for this study.  As discussed above, when the opacity is abruptly decreased by a factor of 10${^4}$, the disk does approach fragmentation-like behavior.  A similar effect is reported by Mayer et al.~(2007), who find that their disk only fragments when the mean molecular weight is suddenly switched from $\\mu=2.4$ to $\\mu=2.7$.  However, the simulation presented by Boley et al.~(2006) was accidentally run with $\\mu=2.7$, and Cai et al.~(2006, 2008) purposefully ran their simulations at the high $\\mu$ for comparison with the Boley et al.~results.  None of the studies reported disk fragmentation.  This may indicate that when a disk approaches fragmentation shortly after a sudden switch in a numerical parameter, e.g., opacity here and $\\mu$ in Mayer et al., the fragmentation may be numerically driven rather than physically, especially because sudden changes in cooling rates make disks more susceptible to fragmentation (Clarke et al.~2008). Regardless, the results do indicate that disk fragmentation by GIs may be possible under very extreme conditions. Whether such conditions are physically possible or realistic remains to be shown.  Based on our results here and in earlier papers, disk fragmentation for $r \\lesssim 10$s of AU appears to be a yet-unproven exception rather than the norm.   \\subsection{FU Ori Outbursts}  In these simulations, the strong bursts of GI activity provide high mass fluxes  ($\\dot{M}\\gtrsim10^{-4}~M_{\\odot}$ yr$^{-1}$, see Fig.~16) throughout each disk.  Even though corotation is at $r\\sim4$ AU for the major spiral arms, the 2 AU region of the disk is strongly heated.  Fluid elements approach peak temperatures of $T\\sim1000$ K in all simulations (Fig.~18). It is not difficult to speculate that if a larger extent of the disk were modeled, the temperatures due to shocks would be large enough to ionize alkalis thermally and possibly sublimate dust (1400-1700 K).  If such a condition is met, then an MRI could activate and rapidly carry mass into the innermost regions of the disk. From these simulations it appears to be plausible, but by no means proven, that a burst of GI activity as far out as 4 AU can drive mass into the inner regions of the disk and create strong temperature fluctuations, which may then be responsible for a thermal instablility.  Although we are simulating very massive disks, we note that such systems may exist during the Class I YSO phase.  Additional studies need to be conducted, preferably with a self-consistent buildup of a dead zone, to address the efficiency of this mechanism in low mass disks.  There are at least three observable signatures for this mechanism.  First, if a GI-bursting mass concentration at a few AU ultimately results in an FU Ori phenomenon, then one would expect to see an infrared precursor from the GI burst, with a rise time of approximately tens of years.  Second, one would also expect for a large abundance of molecular species, which would normally be frozen on dust grains, to be present in the gas phase during the infrared burst due to shock heating (Hartquist 2007, private communication).  Third, approximately the first ten AU of the disk should have large mass flows if the burst takes place near $r\\sim4$ to 5 AU.  We speculate based on these results that if the burst were to take place at 1 AU, then high outward mass fluxes should be observable out to a few AU.  \\subsection{Dust Processing} Each simulation shows that a large fraction of material goes through shocks with $\\mathcal{M}^2\\ge2$. Although these shocks are weak, their abundant numbers may result in the processing of dust to some degree everywhere in the 2-10 AU region. In fact, such processing may be necessary for prepping chondritic precursors for strong-shock survival  (Ciesla 2007, private communication).  Based on the arguments presented in \\S 2, the intent was to produce spirals with large pitch angles by constructing a disk biased toward a strong, sudden GI-activation near 5 AU.  Even with this bias, the pitch angles of the spirals remain small, with $i\\approx10^{\\circ}$. Why are the pitch angles so small?  According to the WKB approximation, $\\cot i = \\mid k_r r/m\\mid$ (Binney \\& Tremaine 1987), where $k_r$ is the radial wavenumber for some $m$-arm spiral.  The most unstable wavelength for axisymmetric waves (Binney \\& Tremaine 1987) is roughly $\\lambda_u\\approx 2\\pi c_s/Q\\kappa\\approx2\\pi h/Q$ for disk scale height $h$.   For a disk unstable to nonaxisymmetric modes, $\\lambda_u$ corresponds to some $m$-arm spiral (e.g., Durisen et al.~2008).  By relating $k_r=2\\pi\\beta m/\\lambda_u$, \\begin{equation} \\cot i\\approx \\beta Qr/h \\end{equation} in the linear WKB limit ($\\mid k_r r/m\\mid \\gg 1$), where $\\beta$ is a factor of order unity that relates $\\lambda_u$ to the $m$-arm spiral.  Because $\\beta Qr/h \\sim10$  in gravitationally unstable disks, the linear WKB analysis may be marginally applicable.  Equation (7) predicts that linear spiral waves in these disks should have pitch angles $i\\approx6$ to 11$^{\\circ}$ for $\\beta$ between 1 and 1/2, respectively.  It appears that this estimate for $i$ extends accurately to the nonlinear regime, a result we did not expect.  The GI spirals are efficient at heating the disk and transporting angular momentum, but {\\it not} at producing chondrules.  Only for the $10^{-4}$ $\\kappa$ simulation are multiple candidate chondrule-producing shocks detected, and there is only one chondrule-producing shock in the 1024 $10^{-3}$ $\\kappa$ simulation. These low opacity disks have relatively flocculent spiral morphologies, including kinks in spiral waves.  The $10^{-4}~\\kappa$ simulations are on the verge of fragmentation. When interpreting these results, it is important to remember that the opacities were lowered abruptly and in a quite unphysical way. In addition, these detections are based on the fractional pressure change $\\eta$.  If the temperature change is used, no chondrule-producing shocks are detected.  This discrepancy may be due to efficient radiative cooling, which may make the assumption of adiabatic shock conditions incorrect, and/or to confusion with additional waves induced by shock bores.  For the rest of this section, we assume that the pressure difference is the reliable shock identifier in order to discuss the implications of detecting chondrule-forming shocks in these disks.  A thermal and spatial history for a fluid element that experiences a possible chondrule-forming shock is shown in Figure 19.  To estimate the occurrence of a chondrule-forming shock, we employ a generous $u_1$-$\\rho$ criterion (Fig.~20).  We do not take into account the optical depth criterion of Miura \\& Nakamoto (2006) on grounds that the large scale over which these shocks take place can allow for chondrules to equilibrate with their surroundings (Cuzzi \\& Alexander 2006). However, one should be aware that the optical depth criterion of Miura \\& Nakamoto will likely exclude all chondrule-producing shocks inside 4 AU in the $u_1$-$\\rho$ plane for these simulations.  All candidate chondrule-forming shocks occur between $r\\sim3$ and 5 AU and at altitudes that are roughly less than a third of the gas scale height.  Because a large degree of settling is assumed, these shocks are located in regions that may be consistent with the dusty environments in which chondrules formed (Wood 1963; Krot et al.~2005).  We estimate that  $\\sim 1~M_{\\oplus}$ of dust is processed through chondrule-forming shocks.  Because these shocks are limited to the onset of the GI burst, a few$\\times10~M_{\\oplus}$ of dust would be processed in the $10^{-4}$ $\\kappa$ disks if they went through about ten outbursts.  If more outbursts take place than those that lead to an FU Orionis event, then more chondritic material could be produced.  In order to produce these shocks, the disk was pushed toward fragmentation by suddenly dropping the opacity by a factor of $10^{4}$.  We conclude that for bursts of GIs near 4 or 5 AU to produce chondrules, the disk must be close to fragmentation."
    },
    "0806/0806.1430_arXiv.txt": {
        "abstract": "{} {An important question about Be stars is whether Be stars are born as Be stars or whether they become Be stars during their evolution. It is necessary to observe young clusters to answer this question.} {To this end, observations of stars in NGC\\,6611  and the star-formation region of Eagle Nebula have been carried out with the ESO-WFI in slitless spectroscopic mode and at the VLT-GIRAFFE (R$\\simeq$ 6400--17000). The targets for the GIRAFFE observations were pre-selected from the literature and our catalogue of emission-line stars based on the WFI study. GIRAFFE observations allowed us to study accurately the population of the early-type stars with and without emission lines. For this study, we determined the fundamental parameters of OBA stars thanks to the GIRFIT code. We also studied the status of  the objects (main sequence or pre-main sequence stars) by using IR data, membership probabilities, and location in HR diagrams.} {The nature of the  early-type stars with emission-line stars in NGC\\,6611 and its surrounding environment is derived. The slitless observations with the WFI clearly indicate a small number of emission-line stars in M16. We observed with GIRAFFE 101 OBA stars, among them 9 are emission-line stars with circumstellar emission in H$\\alpha$.  We found that: W080 could be a new He-strong star, like W601. W301 is a possible classical Be star, W503 is a mass-transfer eclipsing binary with an accretion disk, and the other ones are possible Herbig Ae/Be stars. We also found that the rotational velocities of main sequence B stars are 18\\% lower than those of pre-main sequence B stars, in good agreement with theory about the evolution of rotational velocities. Combining adaptive optics, IR data, spectroscopy, and radial velocity indications, we found that 27\\% of the B-type stars are binaries.  We also redetermined the age of NGC\\,6611 found equal to 1.2--1.8 Myears in good agreement with the most recent determinations.} {} ",
        "introduction": "The origin of the Be phenomenon, i.e. periods of spectral emission due to the presence of a circumstellar envelope around Be stars, is still debated. Rapid rotation seems to be a major key in triggering this phenomenon. To understand the Be phenomenon, it is important to know at which phase of the stellar evolution on the main sequence (MS) it appears.  According to a statistical study of Be stars in clusters, \\citet{fabregat2000} concluded that it may occur in the second half of the MS phase. Taking into account effects due to fast rotation, \\citet{zorec2005} estimated that the appearance of the Be phenomenon among field early-type stars is probably mass-dependent and that it may appear at any time during the MS phase. \\citet{marta2006b,marta2007a} showed that the appearance of Be stars is mass-, and metallicity-dependent. In the  field of Milky Way (MW), they found that less massive Be stars only appear during the second part of the MS, while massive Be stars appear mainly during the first part of the MS, and the intermediate-mass Be stars appear during the whole MS.  To confirm this, it is necessary to observe young or very young open clusters with emission-line stars (ELS).   \\addtocounter{figure}{+1} \\begin{figure*}[t] \\begin{tabular}{cc} \\centering \\includegraphics[angle=0, width=8cm]{wfi_W483.ps} &  \\includegraphics[angle=0, width=8cm]{wfi_w371absbeau.ps} \\\\ \\includegraphics[angle=0, width=8cm]{newHaW483.ps} &  \\includegraphics[angle=0, width=8cm]{w371Ha.ps} \\\\ \\end{tabular} \\caption{Left: H$\\alpha$ spectra of the true ELS star W483. Top-left: spectrum, scaled to the mean value of its continuum, obtained with the WFI in slitless spectroscopic mode. Bottom-left: spectrum,  normalized to its continuum, obtained with the VLT-GIRAFFE. In this last spectrum, the circumstellar and nebular H$\\alpha$ emissions are visible. Right: H$\\alpha$ spectra of the star W371. Top-right: spectrum, scaled to the mean value of its continuum, obtained with the WFI in slitless spectroscopic mode. Bottom-right: spectrum with nebular emission line, normalized to its continuum, obtained with the VLT-GIRAFFE. Only nebular emission is present.} \\label{specWFI1} \\end{figure*}             \\object{NGC\\,6611} in \\object{M16} is a very young cluster previously known to contain a large number of ELS \\citep{hillenbrand1993,dewinter1997}, which have been included in the reference database SIMBAD and WEBDA. However, the investigation of the ELS character was made from low and moderate resolution spectra, which did not allow to distinguish intrinsic stellar emission from nebular emission as noted by \\citet{hillenbrand1993}. Using deep objective prism spectroscopy, which is not sensitive to nebular lines, \\citet{herbig2001} only identified a small number of ELS at the opposite of the other studies. This was recently confirmed by \\citet{evans2005} who observed the more massive population of NGC 6611 with high-resolution spectroscopy (FLAMES and FEROS at ESO). With the ESO-WFI in slitless spectroscopic mode and the VLT-GIRAFFE, we performed observations from late O to early A type stars in NGC 6611 and surrounding fields, which are thought to be still in formation stages \\citep[see e.g.][]{indeb07}, in order to analyse the B star population with and without emission lines.  In the present paper we report on the detection of new ELS in the NGC 6611 region, we determined the fundamental parameters and studied: (i) the evolution of rotational velocities between pre-main sequence phase and main sequence, (ii) the age-distributions of objects in M16, (iii) the evolutionary status of each B-type star, as well as the nature % of the emission/absorption stars: pre-main sequence stars or HAeBe (ELS), or main-sequence stars or classical Be stars (ELS). ",
        "conclusions": "Thanks to our observations with 2 different instrumentations, the ESO-WFI in slitless spectroscopic mode which is not sensitive to the ambient nebula and the VLT-GIRAFFE fibre multiobjects high-resolution spectrograph, we were able to find a small number of true circumstellar ELS (in H$\\alpha$) among the brightest population of the very young cluster NGC 6611 and the star formation region of Eagle nebula. With spectra obtained at the VLT, we were able to study accurately their nature: Herbig Ae/Be or classical Be star. We also conducted the same study for the other non-ELS. Finally, only 11 true ELS with circumstellar or wind emission were found. The other previous potential Be stars from the literature are actually stars with a strong nebular emission pollution in H$\\alpha$.  We determined the fundamental parameters for 85 stars and gave general information for several others. Among our sample of B-type stars, we found 27\\% of them as binaries. Concerning rotational velocities, we found that the B-type MS stars rotate 18\\% slower than B-type PMS objects, in good agreement with published theoretical models at the ZAMS. This value could be used to constrain the models currently developed for the stellar evolution with rotation from the younger (PMS phase) to the older ages (G. Meynet, private communication). With IR data, we found that the low-mass stars are mainly PMS stars without circumstellar emission.  We redetermined the age of NGC\\,6611, found equal to 1.2--1.8 Myears, in good agreement with recent estimates. With clues from spectroscopy, IR, HR ages, membership probabilities, RV, and evolutionary status, we found that: there is a MS population and a PMS population in NGC 6611 itself but also in the surrounding ambient star-formation region of the Eagle Nebula. Among the true circumstellar H$\\alpha$ ELS, we found that: WFI017, W245, W494, W235, W483, W500 are Herbig Ae/Be stars; W301 is a classical Be star, W503 a binary with an accretion disk, and W080 is a possible He-strong magnetic star like W601. This study confirms that the appearance of Be stars is mass- and age-dependent."
    },
    "0806/0806.2575_arXiv.txt": {
        "abstract": "Gaia will provide parallaxes and proper motions with accuracy ranging from 10 to 1000 microarcsecond on up to one billion stars. Most of these will be disk stars: for an unreddened K giant at 6 kpc, it will measure the distance accurate to 2\\% and the transverse velocity to an accuracy of about 1 km/s.  Gaia will observe tracers of Galactic structure, kinematics, star formation and chemical evolution across the whole HR diagram, including Cepheids, RR Lyrae, white dwarfs, F dwarfs and HB stars. Onboard low resolution spectrophotometry will permit -- in addition to an effective temperature estimate -- dwarf/giant discrimination, metallicity measurement and extinction determination. For the first time, then, Gaia will provide us with a three-dimensional spatial/properties map and at least a two-dimensional velocity map of these tracers. (3D velocities will be obtained for the brighter sources from the onboard RV spectrograph). This will be a goldmine of information from which to learn about the origin and evolution of the Galactic disk.  I briefly review the Gaia mission, and then show how the expected astrometric accuracies translate into distance and velocity accuracies and statistics.  I then briefly examine the impact Gaia should have on a few scientific areas relevant to the Galactic disk, specifically disk structure and formation, the age--metallicity--velocity relation, the mass--luminosity relation, stellar clusters and spiral structure. Concerning spiral arms, I note how a better determination of their locations and pattern speed from their OB star population, plus a better reconstruction of the Sun's orbit over the past billion years (from integration through the Gaia-measured gravitational potential) will allow us to assess the possible role of spiral arm crossings in ice ages and mass extinctions on the Earth. ",
        "introduction": "\\vspace*{2ex}  Gaia is a high-accuracy astrometric satellite to be launched by ESA at the end of 2011.  By measuring the positions of stars tens of times over a five year baseline, it can derive the mean position, the parallax and two-dimensional proper motion of each star. These are five components of the six dimensional phase space, the sixth component -- the radial velocity -- being provided for the brighter stars by an onboard spectrograph.  Gaia delivers absolute parallaxes tied to an inertial (quasar--based) reference frame. Gaia, which is currently under construction, represents a major step beyond the enormously successful Hipparcos mission, and is currently the only large-scale astrometric mission beyond the planning stage.  Gaia will perform a survey of the entire sky complete to magnitude G=20 (V=20--22). This covers some 10$^9$ stars, a million quasars and a few million galaxies.  The sky coverage is complete to this magnitude, bar about 1\\% of the sky area where Gaia is confusion limited.  Gaia will achieve an astrometric accuracy of 12--25\\,\\uas\\ at G=15 (providing a distance accuracy of 1--2\\% at 1 kpc) and 100--300\\,\\uas\\ at G=20.  These numbers are also the approximate parallax accuracy in \\uas\\ and the proper motion accuracy in \\uas/year.  The accuracy range reflects a colour dependency: better accuracy is achieved for redder sources (as more photons are collected). Astrometry and photometry are done in a broad (``white light'') band (G). Gaia will also measure radial velocities to a precision of 1--15\\,km/s for stars with V=17 via R=11\\,500 resolution spectroscopy around the CaII triplet (the ``Radial Velocity Spectrograph'', RVS).  To characterize all sources (which are detected in real time), each is observed via low dispersion prism spectrophotometry over 330--1000\\,nm with a dispersion between 3 and 30 nm/pixel (the ``BP/RP'' instrument). From this we will estimate the ``usual'' astrophysical parameters, \\teff, \\logg\\ and \\feh, but also the line-of-sight extinction to stars individually, plus perhaps also [$\\alpha$/Fe] in some cases.  Spectra from the radial velocity spectrograph will help the parameter determination for some stars brighter than about $V=14$, and will also allow the detection of emission lines and abundance anomalies.  Gaia has a nominal mission duration of five years (plus a possible one year extension), and 2--3 years following the end of operations are planned to complete the data processing. (The astrometry is self-calibrating, so the data must be reduced globally/simultaneously to get the final solutions and best astrometric accuracy.) The final catalogue will be available in about 2020, proceeded by earlier data releases. For more information on the satellite, science and data processing see {\\tt http://www.rssd.esa.int/Gaia} and the proceedings volume by \\cite[Turon et al.\\ (2005)]{turon05} (also available from the website). ",
        "conclusions": ""
    },
    "0806/0806.2275.txt": {
        "abstract": "{}{}{}{}{} % 5 {} token are mandatory  \\abstract % context heading (optional) % {} leave it empty if necessary {} % aims heading (mandatory) {We have performed a set of high- and low-spectral resolution phase-resolved X-ray observations of the magnetic B star \\bcep, for which theoretical models predict the presence of a confined wind emitting X-rays from stationary shocks. \\refcom{Some of the} models predict, given the peculiar geometry of \\bcep, strong rotational modulation of the X-ray emission, \\refcom{while other models predict a much lower amplitude modulation at 90 deg phase shift from the modulation predicted from the first group of models}. Our observations were designed to provide a stringent test of such models.} % methods heading (mandatory) {We obtained four observations spaced in rotational phase with \\xmm\\ (using both the EPIC cameras and the RGS spectrograph) and with \\emph{Chandra} (using the LETG spectrograph). A detailed analysis of the data was performed to derive both photometric and spectral parameters from the EPIC data, searching for rotational modulation, and to derive the location of the X-ray plasma from the line ratios in the He-like triplets of N, O and Ne from the RGS data. The LETG data were used to constrain the presence of bulk motions in the plasma. } % results heading (mandatory) {The strong rotational modulation predicted by the \\refcom{early, static} magnetically confined wind model for the X-ray emission is not observed in \\bcep. The small modulation present goes in the opposite direction, pointing to the absence of any optically thick disk of neutral material, \\refcom{and showing a modulation consistent with the later, dynamic models of magnetically confined wind models in B stars}. The lack of observed bulk motion points to the plasma being confined by a magnetic field, but the low plasma temperature and lack of any flaring show that the plasma is not heated by magnetic reconnection. Therefore, the observations point to X-ray emission from shocks in a magnetically confined wind, with no evidence of an optically thick, dense disk at the magnetic equator.} % conclusions heading (optional), leave it empty if necessary {} ",
        "introduction": "Early-type stars were established as strong soft X-ray sources already during the very first survey of stellar X-ray emission conducted with the \\emph{Einstein} observatory (\\citealp{vcf+81}). O and B stars lack the external convection layer which is an essential component of the dynamo in late-type dwarfs, and thus are expected to lack the highly structured magnetic fields which confine and heat solar-type coronae. The observed X-ray emission is thermal in nature, and in general of lower temperature than observed in active late-type stars. The mechanism initially proposed to explain the observed X-ray emission was self shocking in the radiatively driven fast, strong stellar winds which are a characteristic of massive stars (\\citealp{lw80}). This scenario makes a well-defined prediction, i.e.\\ that emission lines should be broad and blue-shifted. The advent high spectral resolution made possible by the launch of \\xmm\\ and \\emph{Chandra} has shown that some massive stars do indeed show the expected signature of X-ray emission from the wind (\\refcom{e.g.\\ $\\zeta$~Pup,} \\citealp{cmw+2001}). Other stars however show narrow spectral lines \\refcom{(e.g.\\ $\\theta^1$ Ori C, \\citealp{goc+2005}).} \\refcom{Narrow lines imply a low-velocity plasma which has been interpreted as a confined plasma, likely requiring magnetic fields (e.g.\\ \\citealp{sch+2003}, \\citealp{sth+2006} and reference therein). Alternative scenarios not requiring magnetic confinement for the X-ray emission from early-type stars showing narrow lines have also been put forward (\\citealp{clg+2006}; \\citealp{lpk+2006}). A recent survey of early-type high-resolution spectra has been performed by \\cite{wc2007}.}  Recently, dipolar fields have been detected in a number of massive stars, in some cases with the magnetic axis at a significant angle from the rotation axis. Scenarios implying magnetic confinement of the wind from a dipolar magnetic field have been developed, and (as discussed below) they succeed in explaining several of the observed characteristics in massive stars with a measured magnetic fields and X-ray emission. For example, \\cite{bm97a} developed a magnetically confined wind shock model (MCWS) to explain the characteristics of the X-ray emission from Ap and Bp stars, and \\citet{bm97b} applied the MCWS model to explain some of the salient characteristics of the X-ray emission of   $\\theta^1$ Ori C, one of the only two O-type stars with a detected magnetic field. $\\theta^1$ Ori C shows strong modulation of the X-ray emission with the rotational period, with a nearly sinusoidal light curve and a $\\simeq 50\\%$ peak-to-peak modulation amplitude (\\citealp{gcs+97}; \\citealp{sfm+2005}); such strong modulation agrees very well with the predictions of the MCWS scenario. Interestingly, the other O star on which a magnetic field has recently been detected, HD 191612 (\\citealp{dhb+2006}) has an unusually slow rotation rate, and it also has a field with a strong dipolar component.  As discussed in Sect.~\\ref{sec:bcep}, $\\beta$\\,Cep has a significant dipolar magnetic field \\refcom{($\\simeq 360$ G, \\citealp{hjd+2000})}, with the magnetic axis at almost 90 deg from the rotational axis, and it therefore is a particularly interesting star for the study of the magnetically confined wind model. Its X-ray emission had already been detected by the \\emph{Einstein} observatory, and \\cite{dwb+2001} (hereafter D01) devised a detailed model for it, based on the MCWS scenario \\refcom{of \\citet{bm97b}} and the later ROSAT observations. \\refcom{They make} clear and verifiable predictions on the temporal variability of the X-ray emission and on its spectrum, as well as on the spatial location of the X-ray plasma.  \\refcom{Later work based on dynamical modeling of the wind-magnetic field interaction however showed that the thick disk predicted by \\citet{bm97b} would be unlikely to form around a star like \\bcep\\ (\\citealp{goc+2005}; \\citealp{uo2002}; \\citealp{to2005}). }  We have carried out a campaign of X-ray observations using both \\emph{Chandra} and \\xmm, specifically designed to provide a stringent test of MCWS-based model of the X-ray emission of \\bcep. The observations were designed to determine whether the X-ray emission presents the variability with rotational phase predicted by the D01 model, and (by using triplet ratios and Doppler shifts) at which distance from the star's photosphere the bulk of the X-ray emission is concentrated. Our interest in the X-ray emission of \\bcep\\ was initially also driven by its being a Be star, although, as discussed in Sect.~\\ref{sec:bcep}, recent observations have suggested that the Be phenomenology is due to the secondary (presumably much less X-ray active) companion in the \\bcep\\ system, rather than to the magnetic, X-ray active primary.  The present paper is structured as follows: after the Introduction, the characteristics of our target star, \\bcep, are discussed in Sect.~\\ref{sec:bcep}. \\xmm\\ and \\emph{Chandra} observations are discussed in Sect.~\\ref{sec:obs} with their analysis, and the relative results are presented in Sect.~\\ref{sec:res}. Finally, our conclusions are presented in Sect.~\\ref{sec:concl}. ",
        "conclusions": "\\label{sec:concl}  \\subsection{X-ray modulation and location of the plasma}  \\bcep\\ has a close B6-8 companion which could be the component on which the H$\\alpha$ emission is located (\\citealp{sho+2006}). The companion is very unlikely to be contributing to X-ray emission in the \\bcep\\ system: in the regime in which X-ray emission is originating in shocks in the wind (i.e. for OB stars) the emission increases with stellar mass and spectral type; therefore the B1 primary is very likely dominating the X-ray emission from the system and thus all our conclusions are unaffected by the presence of the late B-type companion and by the fact that the \\bcep\\ primary may no longer be considered a Be star.  Our X-ray campaign on \\bcep\\ shows limited evidence of modulation of its X-ray emission level, both short-term (i.e.\\ within a rotational period) and long-term (i.e.\\ from one observation to the other), with for example the ROSAT and \\xmm\\ observations showing very similar flux levels across several years. Evidence of low-level modulation of the X-ray emission with stellar rotation is present in the pn low-resolution X-ray data, while the LETG and RGS high-resolution spectra appear remarkably similar to each other. For instance, as shown in Fig.~\\ref{fig:orgs}, the \\ovii\\ triplet shows no evidence of variation from one rotational phase to the other. For the \\emph{Chandra} LETG spectra little if any variability is present in the line fluxes as well as in the centroid of the line positions.  Such high degree of constancy in the X-ray emission is in contradiction with the high ($\\simeq 50\\%$) level of rotational modulation expected in the MCWS framework \\refcom{of} D01: with the geometry of \\bcep\\ the disk is seen alternatively face-one and edge-on. If indeed X-ray emission is due to stationary shocks on either side of the disk (which is expected to have a high optical thickness at soft X-ray wavelengths), during face-on configurations about 50\\% of the flux should be absorbed.  In contrast to what was expected from the MCWS scenario \\refcom{by D01}, we find a slightly ($\\simeq 10$\\%) higher emission level for the face-on configuration with respect to the edge-on configuration. This allows us to exclude the simple scenario proposed by D01 and in particular it definitively rules out the presence of the optically thick layer in the magnetic equator. \\refcom{Such smaller modulation, with the observed phase, is along the lines predicted by the dynamical MCWS models of \\cite{uo2002} and \\cite{goc+2005}, which would indeed predict for the \\bcep\\ configuration, no thick disk and a modulation of approximately 5\\%, fully consistent with the observed values.}  The measurements of the line intensity in the He-like triplets, and in particular the \\ovii\\ one, allows us to determine the location of the bulk of the emitting plasma. The temperature of formation of the \\ovii\\ triplet peaks at 0.17 keV, not very different from the characteristic temperature of the bulk of the emitting plasma, as determined from the pn spectra. Therefore, the characteristic distance from the star's photosphere determined from the \\ovii\\ triplet should be representative of the majority of the plasma at X-ray temperatures. The detection at $\\sim$3$\\sigma$ of the $f$ line in the \\ovii\\ triplet allows us to determine a distance range from the photosphere (rather than an upper limit), estimated to be $d \\simeq 4\\,R_{\\star}$ (see sec.~\\ref{RGS data}). \\refcom{The derived location for the X-ray emitting plasma is also consistent with the location predicted by the \\cite{uo2002} and \\cite{goc+2005} models, which locate the plasma between the Alfven and Kepler radii, i.e. between $R \\simeq 5 \\,R_{\\star}$ and $R \\simeq 7 \\,R_{\\star}$.}  An additional constraint on the spatial location of the X-ray plasma in \\bcep\\ comes from the relation between the plasma pressure and the confining magnetic field, under the assumption that the X-ray emitting plasma is indeed magnetically confined, in agreement with the lack of line broadening otherwise expected for a standard wind shock model. The plasma density can be derived from the emission measure, derived either from the global fit to the pn spectrum, or taking the emission in a given line. If we take the flux in the strongest ($r$) \\ovii\\ line, we derive $E\\!M = 1.5 \\times 10^{55}$ cm$^{-3}$. Assuming the plasma to be confined in a spherical shell comprised between 4 and $6\\,R_{\\star}$ (as derived from the triplet ratios), the emitting volume is $V = 7.5\\times 10^{37}$ cm$^3$, resulting in an average plasma density $n = 4.5\\times 10^8$ cm$^{-3}$. Such low density is consistent with no density effect on the He-like triplets of the elements being considered. The magnetic field needed to confine this plasma (taking into account the average temperature $T \\simeq 0.3$ keV) is $B = \\sqrt{8\\pi 2 n k T} \\simeq 3$\\,G. Assuming the stellar magnetic field to dipolar, and taking the best estimate for the polar magnetic field of 355 G, and scaling the dipolar field as $d^3$ ($d$ being the distance from the star), the magnetic field at $5\\,R_{\\star}$ from the photosphere is $355/5^3 \\simeq 3$\\,G, similar to the pressure required to confine the emitting plasma. This indicates that the emitting plasma is weakly confined by the magnetic field, with $\\beta \\simeq 1$.  Using the location of the plasma found above ($R_{\\rm out} = 6\\,R_{\\star}$ and $R_{\\rm in} = 4\\,R_{\\star}$) and assuming a geometrically thin disk, we can estimate the variation of the X-ray flux from edge-on to face-on configuration to be of the order of 7\\%. This number is compatible within the error bars with the observed variation level. This scenario would rule out the presence of X-ray emitting plasma at high latitudes above the magnetic equator \\refcom{and is again compatible with the predictions from the more recent, dynamical MCWS models.}  \\subsection{Magnetic confinement}  Standard models of X-ray emission from massive stars, originating in shocks in an unconfined wind, predict both a general blue-shifting of the line together with its broadening. The magnitude of the expected  broadening is comparable with the wind terminal velocity, which in the case of \\bcep\\ is $\\simeq 800$ - $1500$ km/s. The blue-shift depends on the amount of absorption of the line's red wing, and is expected to be less than the broadening.  The analysis of the LETG line profiles rule out the presence of any broadening above the instrumental line profile, which has a FWHM of approximately 600 km/s, and of blue-shifts greater than 160 km/s. While they are not extremely constraining relative to the wind terminal velocity of \\bcep, these elements indicate a lack of any significant bulk motions in the X-ray emitting plasma, and point toward its being magnetically confined, therefore ruling out, for \\bcep, X-ray emission from shocks in an unconfined wind.  The relatively low temperature of the X-ray emitting plasma (with the dominant component being at $\\approx 3$ MK) together with the lack of significant short-term variability (flares) point however to a lack of magnetic heating (i.e.\\ due to magnetic reconnection): in active cool stars, the magnetic reconnection that dominates the heating of the plasma results in  much higher temperatures and stochastic variability. The high temperatures observed for $\\theta^1$ Ori C also points to the presence of magnetic heating in some massive stars.  In the case of \\bcep\\ the magnetic heating only has the apparent function of confining the plasma, which is likely heated to the observed 3 MK by shocks. Also, the bulk of the X-ray plasma appears to be confined by a relatively weak magnetic field, close to the limit where the magnetic pressure becomes too weak to confine the plasma.  \\subsection{Comparison with other stars}  The behavior of \\bcep\\ as observed in our \\emph{Chandra} and \\xmm\\ observations appears to be significantly different from the other well studied example of magnetically confined wind in a massive star, $\\theta^1$ Ori C. In that case, the X-ray emission is strongly modulated at the rotational period, by approximately 50\\%, as predicted by the MCWS model. At the same time, the observed plasma temperatures are, for $\\theta^1$ Ori C, much higher than for \\bcep, reaching up to 30 MK. Furthermore, the analysis of the triplets in $\\theta^1$ Ori C show that the emitting plasma is located much closer to the star's photosphere than in the case of \\bcep. In $\\theta^1$ Ori C the bulk of the plasma is located at $\\simeq 1.5\\,R_{\\star}$ while in \\bcep\\ the cooler plasma traced by the O triplet is located at $\\simeq 5\\,R_{\\star}$, while the hotter plasma traced by the Ne triplet is located at $\\la 2\\,R_{\\star}$ from the photosphere, pointing also to a stratification of the plasma as a function of temperature. Finally, the metal abundances determined for the $\\theta^1$ Ori C plasma are much higher than for \\bcep. In \\bcep\\ all elements appear to be depleted with the exception of Si, while many elements are enhanced in the  $\\theta^1$ Ori C spectrum, showing a different process in operation.  \\cite{goc+2005} have modeled the observed emission from $\\theta^1$ Ori C with a dynamic version of the MCWS, showing that both infall to the photosphere and outflow make it unlikely that the thick disk predicted by \\cite{bm97a} would actually form. They interpret the observed modulation (which goes in the opposite direction from the one predicted by the D01 model, as also observed by us for \\bcep) as occultation of the X-ray emitting plasma (located, in their model, close to the equator) by the stellar photosphere.  A similar situation, with the dynamics preventing the thick disk from forming, is possibly present in \\bcep, so that also in this case the X-ray emission is produced by a magnetically confined wind, with the lower observed temperatures explained by the lower wind velocity.  \\subsection{Conclusions}  The \\emph{Chandra} and \\xmm\\ observations discussed here have failed to display the signatures expected by the static MCWS model \\refcom{of D01}, in particular the strong rotational modulation of the X-ray emission due to the presence of the cool disk around the star. However the X-ray plasma appears confined, and a small amplitude modulation is visible in the \\xmm\\ EPIC data. Together with the low temperature and the lack of flaring, this can be interpreted as emission from a magnetically confined wind but without the cool, optically thick disk predicted by the D01 model. \\refcom{The observed modulation is fully compatible with the emission scenario predicted by the more recent dynamical MCWS models for a star with the characteristics of \\bcep.}"
    },
    "0806/0806.0570_arXiv.txt": {
        "abstract": "Massive early-type galaxies are observed to lie on the Mass Plane (MP), a two-dimensional manifold in the space of effective radius $\\Re$, projected mass $\\Metp$ (measured via strong gravitational lensing) and projected stellar velocity dispersion $\\sget$ within $\\Re/2$.  The MP is less `tilted' than the traditional Fundamental Plane, and the two have comparable associated scatter. This means that the dimensionless structure parameter $\\cet=2G\\Metp/(\\Re\\sgetsq)$ is a nearly universal constant in the range $\\sget=175-400$ km s$^{-1}$. This finding can be used to constrain the mass distribution and internal dynamics of early-type galaxies: in particular, we explore the dependence of $\\cet$ on light profile, dark-matter distribution, and orbital anisotropy for several families of spherical galaxy models.  We find that a relatively wide class of models has values of $\\cet$ in the observed range, because $\\cet$ is not very strongly sensitive to the mass distribution and orbital anisotropy. The degree of fine tuning required to match the small intrinsic scatter of $\\cet$ depends on the considered family of models: if the total mass distribution is isothermal ($\\propto r^{-2}$), a broad range of stellar luminosity profile and anisotropy is consistent with the observations, while Navarro, Frenk \\& White dark-matter halos require more fine tuning of the stellar mass fraction, luminosity profile and anisotropy. If future data can cover a broader range of masses, the MP could be seen to be tilted by the known non-homology of the luminosity profiles of early-type galaxies, and the value of any such tilt would provide a discriminant between models for the total mass-density profile of the galaxies. ",
        "introduction": "The origin of empirical scaling laws is a key open issue in observational cosmology. Galaxies do not come in all sizes, shapes, colours, but rather tend to to live in lower-dimensional manifolds, which represent a stringent testing ground for theories of galaxy formation and evolution.  Early-type galaxies obey a particularly tight scaling law: the so-called Fundamental Plane \\cite[FP;][]{Djo87,Dre87}.  In the space of effective radius $\\Re$, central velocity dispersion $\\sgee$ and effective surface brightness $\\Ie{\\equiv}L/(2\\pi\\Resq)$ (where $L$ is the total luminosity of the galaxy), they lie on the following relation with remarkably small scatter \\citep[$\\lsim 20 \\%$ in $\\Re$;][]{Ber03b}: \\begin{equation} \\log\\Re=a\\log\\sgee+b\\log\\Ie+\\const, \\label{eqFP} \\end{equation} where the numerical value of $a$ and $b$ depends somewhat upon the wavelength of observations and upon the sample and the fitting method \\citep{Pah98,Ber03b}. The FP is said to be `tilted', in the sense that the coefficients $a$ and $b$ differ significantly from the values $a=2$ and $b=-1$ expected {\\it for structurally and dynamically homologous systems with luminosity-independent stellar mass-to-light ratio and dark-matter distribution}. Several explanations have been proposed for the tilt, including a systematic dependence of stellar mass-to-light ratio or dark-matter content and distribution upon luminosity (and hence presumably upon mass), structural non-homology and orbital anisotropy \\citep[e.g.][]{Fab87,Ben92,RenC93,CioLR96,Bor03}.  Recently, \\cite{Bol07,Bol08b}, using a sample of strong gravitational lenses, have shown that early-type galaxies lie on a Mass Plane (MP) \\begin{equation} \\log\\Re=\\am\\log\\sget+\\bm\\log\\Sget+\\const, \\label{eqMP} \\end{equation} where $\\sget$ is the projected velocity dispersion within an aperture radius $\\Re/2$ and $\\Sget$ is the surface mass density within $\\Re/2$, with $\\am=1.82 \\pm 0.19$, $\\bm=-1.20 \\pm 0.12$ and RMS orthogonal scatter of 1.24 when normalized by the observational errors.  The fact that $(\\am,\\bm)$ are close to $(2,-1)$ and that the scatter is small can be expressed in terms of structural and dynamical homology of the lenses, by defining the dimensionless structure parameter \\begin{equation} \\cet\\equiv{2G\\Metp\\over \\Re\\sget^2}, \\label{eqcet} \\end{equation} where $\\Metp$ is the total projected mass within $\\Re/2$.  For their sample of lens early-type galaxies from the Sloan Lens ACS (SLACS) Survey \\cite{Bol08b} find on average \\begin{equation} \\langle\\log \\cet\\rangle=0.53\\pm0.057, \\label{eqobsrange} \\end{equation} which throughout the paper we will refer to as the ``observed range'' of $\\cet$. We note that the observed scatter on $\\langle\\log \\cet\\rangle$ is 0.08, but here we consider the estimated intrinsic scatter 0.057 \\citep[see][]{Bol08b}.  As discussed in several papers \\citep{Bol06,Tre06,Bol08a,Tre08} the SLACS lenses are found to be indistinguishable from control samples of Sloan Digital Sky Survey (SDSS) galaxies with the same stellar velocity dispersion and size, in terms of luminosity/surface brightness, location on the FP, and environment.  This inspires some confidence that the results found for the lens sample, including the MP, are generic properties of the overall class of early-type galaxies.  Independent of its origin and theoretical interpretation, the existence of the MP is a powerful empirical tool to estimate galaxy mass by using information on size and velocity dispersion only \\citep{Bol07}. In addition, it is clear that the very existence of the MP may be used to improve our understanding of galaxy formation and evolution.  What is the origin of such a strong correlation among measurable galaxy quantities? Or, in other words, what kinds of galaxy models can be ruled out by the existence of a tight MP? Although this question has been asked before in regards to the traditional FP, the MP provides an additional powerful tool. In fact, there are a few differences between the FP (equation~\\ref{eqFP}) and the MP (equation~\\ref{eqMP}):  \\begin{enumerate}  \\item The FP is sensitive to the galaxy {\\it stellar} mass-to-light ratio, while the MP is not.  This implies that, e.g., the role of stellar populations in establishing the tilt and scatter of the FP can be disentangled by looking at the MP.\\footnote{Strictly speaking the MP depends on the properties of the stellar populations through $\\Re$ and $\\sget$. However, $\\Re$ and $\\sget$ do not depend on the value of the stellar mass-to-light ratio, but only on its radial variation. This variation is expected to be small based on observed colour gradients and is generally neglected in dynamical studies \\citep[e.g.][]{Kro00,Cap06}. For simplicity, in this study we assume uniform stellar mass-to-light ratios within each galaxy.}  \\item The FP is traditionally based on the central velocity dispersion $\\sgee$ (measured within $\\Re/8$), while the MP has been constructed using $\\sget$, which is measured within $\\Re/2$.  This is a consequence of the fixed spatial observing aperture of the SDSS spectrograph; an MP based upon $\\sgee$ could be constructed using spatially resolved spectroscopy of the SLACS lens sample.\\footnote{In general, the larger the aperture radius $R$ used to measure the aperture velocity dispersion $\\sgasq$, the less $\\sgasq$ is sensitive to the orbital anisotropy. We recall that for any stationary, non-rotating, spherically symmetric system with constant mass-to-light ratio $\\sgasq(R)\\to\\sgvsq/3$ for $R\\to\\infty$, where $\\sgvsq$ is the virial velocity dispersion \\citep[e.g.][]{Cio94}.}  \\item The FP combines quantities evaluated on different scales ($\\Re$, $\\Re/8$), while MP combines quantities evaluated within the same radius $\\Re/2$.  Again, this is partially due to the fixed SDSS spectroscopic aperture, though the apertures of the lensing mass measurements are fixed by the cosmic configuration of the individual strong-lens systems.  \\end{enumerate}  Each of the points above can contribute to make the MP less tilted (and presumably with less scatter) than the FP. For example, given the relatively large spectroscopic aperture used to define $\\cet$, we expect it to be robust with respect to changes in the detailed properties of galaxy structure, internal dynamics, and dark-matter content. Similarly, by replacing surface brightness with surface mass density we expect that tilt and scatter due to diversity of chemical composition or star formation history be reduced in MP. Furthermore, having all but removed the effects of stellar population the MP is potentially a cleaner diagnostic than the FP of the structural and dynamical properties of early-type galaxies.  In this paper, we exploit the existence of the MP to constrain important properties of early-type galaxies, such as orbital anisotropy and dark-matter distribution. We achieve this goal by constructing observationally and cosmologically motivated families of galaxy models and finding the range of parameter spaces consistent with the observed range of $\\cet$.  For the sake of simplicity, in the present investigation we limit ourselves to spherically symmetric models. As with the FP \\citep{Fab87,Sag93,Pru94,Lan03,Ric05}, deviation from spherical symmetry is expected to increase the scatter of the MP because of projection effects. Thus, a natural follow-up of the present work would be the extension to non-spherical models.  The paper is organized as follows. In Section~\\ref{secmod} we describe the models, in Sections~\\ref{secres} and~\\ref{secser} we present our results, and in Section~\\ref{seccon} we conclude. ",
        "conclusions": "\\label{seccon}  In the range $\\sget=175-400$ km s$^{-1}$ the MP of early-type galaxies has no significant tilt and small associated scatter. This means that the dimensionless structure parameter $\\cet$ defined in equation~(\\ref{eqcet}) is a nearly universal constant. In other words, the range of values of $\\cet$ ``allowed'' by the observational data is remarkably small. Even for spherical galaxy models, $\\cet$ is expected to depend on the stellar density profile, orbital anisotropy of stars, and total (dark plus luminous) mass distribution.  In this paper we explored the constraints posed by the existence of the MP on several relevant families of galaxy models.\\footnote{ Throughout the present paper we considered the MP in the standard context of Newtonian gravity with dark matter. See \\cite{San08} for an interpretation of the MP in the context of Modified Newtonian Dynamics.}  Limiting to spherically symmetric models, we found that $\\cet$ is not very strongly dependent on galaxy structure and kinematics, so a relatively wide class of models have values of $\\cet$ within the observed range.  Therefore, strictly speaking, the massive early-type galaxies lying on the MP are not necessarily structurally and dynamical homologous.  However, not all the studied models behave in the same way when compared to the observational data.  Models in which the total density profile is a SIS are consistent with the observed range of $\\cet$ for a wide class of stellar density profiles, and only models with extremely radial or tangential anisotropies are excluded. The light-traces-mass hypothesis is not excluded by the observational constraints here considered, apart for the case of high-$m$ \\Sersic models, which cannot be reconciled with the MP within the observed scatter.  (However, LTM models are known to fail other observational constraints: see Section~\\ref{sectot}). We also considered cosmologically-motivated models with NFW dark-matter halos (with or without adiabatic compression), finding that they are consistent with the MP only for a relatively limited range of values of their parameters, so a degree of fine-tuning between light profile and anisotropy is required.  Among these NFW models, those with adiabatically contracted halos and those that are baryon dominated seem to require slightly less fine tuning than those with non-contracted halos and those that are dark-matter dominated.  This work has focused on the average value of $\\cet$ in the SLACS sample, along with its intrinsic scatter.  With the exception of Section~\\ref{secser}, we have not explored the implications of the fact that this intrinsic scatter is not correlated with either mass or with the ratio of Einstein radius to $\\Re$ \\citep{Bol08b}.  These observational results indicate a degree of structural homogeneity across a range in mass.  In future works, we will explore these mass-dependent results in the context of mass-dynamical models such as those considered here.  We also plan to refine these analyses based on the results of forthcoming velocity-dispersion measurements of higher signal-to-noise ratio and in smaller and more uniform spatial apertures.  We also explored the possibility that the observed dependence of the \\Sersic index $m$ on the galaxy luminosity could tilt the MP when a sufficiently large mass range is considered. In this respect, SIS models behave differently from all other models: a slightly tilted MP is predicted in the early type galaxies have SIS total density distribution, while a ``bent'' MP is predicted in all the other explored cases.  The effect is large enough to be measurable with sample of lenses comparable to SLACS in size and quality and extending a further decade in galaxy mass.  In conclusion, our results are consistent with the hypothesis that massive early-type galaxies have isothermal ($\\propto r^{-2}$) total mass density distribution, though alternative hypotheses cannot be excluded on the basis of the existence of the MP alone, although in some cases they require a degree of fine tuning.  In any case, the process of formation of early-type galaxies lead to systems with a combination of total mass distribution, luminosity profile, and orbital anisotropy such that they lie close to the MP. It will be interesting to quantify whether the observed fine tuning is quantitatively consistent with the range of simulated properties of early-type galaxies in the standard hierarchical model of galaxy formation."
    },
    "0806/0806.1191.txt": {
        "abstract": "In a previous paper~\\cite{Khlopov:2007ic}, we showed how the minimal walking technicolor model (WTC) can provide a composite dark matter candidate, by forming bound states between a $-2$ electrically charged techniparticle and a $^4He^{++}$. We studied the properties of these \\emph{techni-O-helium} $tOHe$ ``atoms'', which behave as warmer dark matter rather than cold. In this paper we extend our work on several different aspects. We study the possibility of a mixed scenario where both $tOHe$ and bound states between $+2$ and $-2$ electrically charged techniparticles coexist in the dark matter density. We argue that these newly proposed bound states solely made of techniparticles, although they behave as Weakly Interacting Massive Particles (WIMPs), due to their large elastic cross section with nuclei, can only account for a small percentage of the dark matter density. Therefore we conclude that within the minimal WTC, composite dark matter should be mostly composed of $tOHe$. Moreover in this paper, we put cosmological bounds in the masses of the techniparticles, if they compose the dark matter density. Finally we propose within this setup, a possible explanation of the discrepancy between the DAMA/NaI and DAMA/LIBRA findings and the negative results of CDMS and other direct dark matter searches that imply nuclear recoil measurement, which should accompany ionization.  %Minimal walking technicolor models can provide a nontrivial solution %for cosmological dark matter, if the lightest technibaryon $UU$ and technilepton $\\zeta$ are %doubly charged. Technibaryon and technilepton asymmetry generated in the early %Universe is related to baryon asymmetry and it is possible to %create excess of technileptons with charge ($-2$) larger, %than the excess of technibaryons with charge ($+2$). It provides all the %positively charged technibaryons to bind with technileptons in composite atom-like Weakly Interacting Massive Particle (WIMP) states $[(UU)\\zeta]$, %while the excessive technileptons are all captured by $^4He$, creating %\\emph{techni-O-helium} $tOHe$ ``atoms'', as soon as $^4He$ is %formed in Big Bang Nucleosynthesis (BBN). In a wide range of techniparticle masses below %few TeV this solution corresponds to %a fixed negative value of $L/B$ ratio, but at larger masses this ratio grows exponentially %by absolute value, leading %to strong lepton asymmetry in the period of BBN. The model predicts nonzero weak charge %for $[(UU)\\zeta]$ atoms and severe Cryogenic Dark Matter Search (CDMS) constraints admit no more than few percents for the contribution of %this WIMP component into total dark matter density, in which nuclear interacting techni-O-helium dominates. %Nuclear interactions with matter slow down cosmic %techni-O-helium in Earth %below the threshold of underground dark matter detectors, thus %escaping their direct detection. However, annual modulation of inelastic electromagnetic processes in the bodies %of detectors is possible, what can be interesting for interpretation of results of DAMA/NaI and DAMA/LIBRA experiments.  {Walking technicolor theories} ",
        "introduction": "Walking technicolor theories (WTC) have regained a lot of interest recently. This is because they can naturally break the electroweak symmetry without violating experimental constraints by the electroweak precision measurements. Several technicolor theories that have techniquarks transforming under higher representations of the gauge group, require a small number of colors and flavors in order to become quasi-conformal~\\cite{Sannino:2004qp,Hong:2004td,Dietrich:2005jn,Dietrich:2006cm}. Because of this property, the Higgs particle can be composed of two techniquarks and be able to couple even to the heavier Standard Model particles like the top quark. On the other hand, the fact that these theories become conformal only for a small number of colors and flavors, differentiates them from the old baroque technicolor models that are excluded by the electroweak precision measurements. In addition, the possibility of unification of the couplings makes the walking technicolor theories legitimate candidates for the Large Hadronic Collider (LHC)~\\cite{Gudnason:2006mk}.  Among the walking technicolor theories, special interest has been drawn to the minimal case. This particular model contains only two techniquarks that transform under the adjoint representation of the $SU(2)$ technicolor group and a new lepton family in order to cancel the Witten global anomaly. This minimal model has been investigated thoroughly in~\\cite{Gudnason:2006ug,Foadi:2007ue}. A holographic description of the theory was presented in~\\cite{Dietrich:2008ni}, where several predictions regarding the mass spectrum were made. Lattice methods have also been used recently for the study of gauge theories with fermions that transform under higher dimensional representations~\\cite{Catterall:2007yx,DelDebbio:2008wb,DelDebbio:2008zf}. Although simple in nature, this minimal walking technicolor model can provide several possibilities for dark matter. In particular, the theory can admit as dark matter particles technibaryons~\\cite{Gudnason:2006yj}, bound states between a neutral techniquark and a technigluon~\\cite{Kouvaris:2007iq}, heavy leptons of the fourth lepton family~\\cite{Kainulainen:2006wq}, or bound states between a $-2$ electrically charged techniparticle and a $He^{++}$~\\cite{Khlopov:2007ic}. In the latter case, WTC offers a new exciting realization of a composite dark matter scenario, which was earlier considered in different aspects in the model of teraparticles \\cite{Glashow:2005jy,Fargion:2005xz}, in the AC model \\cite{Fargion:2005ep,Khlopov:2006uv}, based on the approach of an almost commutative geometry \\cite{Connes:1994yd,Stephan:2005uj}, and in the model of 4th generation \\cite{Khlopov:2005ew,Belotsky:2006fd,Belotsky:2006pp,4Q}, assuming existence of stable heavy $U$ quark \\cite{Belotsky:2005ui}.  In all these recent models (see review in \\cite{Khlopov:2006dk,Khlopov:2007zza,Khlopov:2008rp,Khlopov:2008rq}), the predicted stable charged particles form neutral atom-like states, composing the dark matter of the modern Universe and escaping experimental discovery. It offers new solutions for the physical nature of the cosmological dark matter. The main problem for these solutions is to suppress the abundance of positively charged species bound with ordinary electrons, which behave as anomalous isotopes of hydrogen or helium. This problem remains unresolved, if the model predicts stable particles with charge $-1$, as it is the case for tera-electrons \\cite{Glashow:2005jy,Fargion:2005xz}.  The possibility of stable doubly charged particles $A^{--}$ and $C^{++}$, revealed in the AC model, offered a candidate for dark matter in the form of elusive (AC)-atoms. In the charge symmetric case, when primordial concentrations of $A^{--}$ and $C^{++}$ are equal, their binding in the expanding Universe is not complete due to freezing out and a significant fraction of free   relic $C^{++}$, which is not bound in (AC)-atoms, is left in the Universe and represents a potential danger of anomalous helium overproduction. The suppression of this fraction in terrestrial matter involves a new long range interaction between A and C, making them to recombine in (AC)-atoms inside dense matter bodies~\\cite{Fargion:2005ep,Khlopov:2006uv}.  In the asymmetric case, corresponding to excess of $-2$ charge species, as it was assumed for $(\\bar U \\bar U \\bar U)$ in the model of stable $U$-quark of a 4th generation, their positively charged partners annihilate effectively in the early Universe. The dark matter is in the form of nuclear interacting O-helium - atom-like bound states of $-2$ charged particles and primordial helium, formed as soon as $He$ is produced in Big Bang Nucleosynthesis (BBN). Such an asymmetric case was realized in \\cite{Khlopov:2007ic} in the framework of WTC, where it was possible to find a relationship between the excess of negatively charged anti-techni-baryons and/or technileptons and the baryon asymmetry of the Universe.  The minimal walking technicolor model we use is the same as in our previous paper~\\cite{Khlopov:2007ic} (and references therein). It contains two techniquarks that transform under the adjoint representation of an $SU(2)$ gauge group, i.e. up $U$ and down $D$, with electric charges $1$ and $0$ respectively. There is also a new fourth family of leptons $\\nu'$ and $\\zeta$ with charges $-1$ and $-2$ respectively. This hypercharge assignment is not unique, however it is consistent, since it makes the theory gauge anomaly free. It was already noticed in \\cite{Khlopov:2007ic} that since two types of stable doubly charged particles (technibaryon $(UU)^{++}$ and technilepton $\\zeta^{--}$) can exist, the excess of positively charged $(UU)^{++}$ together with the excess of negatively charged $\\zeta^{--}$ is also possible, giving rise to atom-like $[(UU)\\zeta]$ WIMP species.  Here we analyze the ability of WTC to provide this WIMP solution for composite dark matter. It is evident that the predicted abundance and cosmological role of $[(UU)\\zeta]$ are determined by the relation between the excess of its constituents $(UU)^{++}$ and $\\zeta^{--}$. Their excess can be different, although the case where the excess of $(UU)^{++}$ is larger, leads to the unresolved problem of anomalous helium overproduction.  One can find a similar problem in the case where the excess of $(UU)^{++}$ is equal to the excess of $\\zeta^{--}$. In full analogy with the cosmology of the AC model \\cite{Fargion:2005ep,Khlopov:2006uv}, most of $(UU)^{++}$ and $\\zeta^{--}$ are bound in this case in $[(UU)\\zeta]$ ``atoms\", but the remaining fraction of unbound $(UU)^{++}$ is still up to ten orders of magnitude larger than the experimental upper limits on anomalous helium in terrestrial matter \\cite{exp3}. Since the minimal WTC can not offer new long range interactions between $(UU)$ and $\\zeta$, ordinary atoms of anomalous helium $[(UU)ee]$ and nuclear interacting techni-O-helium $[He^{++}\\zeta^{--}] $, having different mobilities in matter, inevitably fractionate. It prevents their recombination in $[(UU)\\zeta]$, which might reduce the concentration of anomalous helium in terrestrial matter below experimental upper limits.  Therefore to solve the problem of anomalous helium in the framework of minimal WTC, we are left with the only option to have the excess of negatively charged $\\zeta^{--}$ larger than the excess of $(UU)^{++}$. This provides complete binding of $(UU)^{++}$ in $[(UU)^{++}\\zeta^{--}]$, while the residual excessive $\\zeta^{--}$ bind with helium in techni-O-helium. This solution can be effective even if the excess of $\\zeta^{--}$ exceeds the excess of $(UU)^{++}$ by relative amount of $\\sim 10^{-8}$. Therefore it seems that the WIMPs $[(UU)\\zeta]$ can be the dominant dark matter component, making the nuclear interacting techni-O-helium dynamically negligible, as it was the case for the AC model \\cite{Fargion:2005ep,Khlopov:2006uv}.  However, we'll show here that unlike the neutral $(AC)$ atoms, having zero electroweak charge, the weak charge of $[(UU)\\zeta]$ is non-zero and its interaction with nuclei, mediated by ordinary $Z$-boson, should lead to an observable effect in the CDMS experiment  \\cite{Akerib:2005kh,Ahmed:2008eu}, unless the contribution of $[(UU)\\zeta]$ to the total dark matter density is restricted to be a few percent.  An interesting feature of the considered scenario is that in a wide interval of masses of $(UU)$ and $\\zeta$, the generation of excess corresponding to the saturation of the observed dark matter by techniparticles, predicts a fixed negative value for the ratio of lepton number $L$ over the baryon number $B$. This ratio is constant for masses below few TeV and then rapidly grows by absolute value for larger masses and exceeds $10^{8}$, when they approach 10 TeV. A large negative value of $L/B$ corresponds to strong lepton asymmetry and to the excess of antineutrino in the period of BBN, which leads to a corresponding growth of primordial $He$ abundance. This argument provides an upper limit on masses of techniparticles.  The paper is organized as follows. After a brief description of the general chronological framework for the considered techniparticle Universe (Section \\ref{Chronology}), we study the relation between baryon asymmetry and techniparticle excess, fixing the value of $L/B$ ratio (Section \\ref{Excess}). %and discuss the mechanisms of composite dark matter formation (Section \\ref{Formation}) In Section \\ref{Detection}, we deduce an upper limit on possible contributions of $[(UU)\\zeta]$ WIMPs in the total dark matter density, which follow from the most recent severe constraints of the CDMS experiment \\cite{Ahmed:2008eu}. We also speculate on the possibility to explain the positive results of DAMA/NaI (see for review \\cite{Bernabei:2003za}) and DAMA/Libra \\cite{Bernabei:2008yi} experiments by ionization effects of inelastic processes, induced by techni-O-helium in the matter. We consider the main results of the present work in Section \\ref{Discussion}. ",
        "conclusions": "Discussion} In this paper we explored the cosmological implications of a walking technicolor model with doubly charged technibaryons $UU^{++}$ and technileptons $\\zeta^{--}$. We studied a possibility for a WIMP-like composite dark matter in the form of heavy ``atoms\" $[UU^{++}\\zeta^{--}]$. To avoid overproduction of anomalous isotopes (related to $UU^{++}$, which are not bound in these atoms), the excess of $-2$ charged technileptons $\\zeta^{--}$ should be larger than the excess of $UU^{++}$ generated in the Universe. The residual doubly charged $\\zeta^{--}$ bind with $^4He$ in the techni-O-helium neutral states.  In all the previous realizations of composite dark matter scenarios, this excess was put by hand to saturate the observed dark matter density. In our paradigm, the abundance of techibaryons and technileptons is connected naturally to the baryon relic density. Moreover, in a rather wide window of techniparticle masses below few TeV, a robust prediction follows for the ratio $L/B$ of lepton and baryon asymmetries. At further increase of techniparticle mass, this ratio grows rapidly. It provides an upper limit on the mass of techniparticles from the condition that large negative value of $L/B$ does not lead to overproduction of primordial $^4He$ in BBN.  Since techni-O-helium binds some fraction of $^4He$, an interesting possibility appears that is at large values of $L/B$, the excessive $^4He$ is hidden in the techni-O-helium. However, due to the non-zero weak isospin charge of $[UU^{++}\\zeta^{--}]$, the presence of this dark matter component should lead to observable effect in underground dark matter detectors. The CDMS constraints reduce the allowed fraction of this component to a few per cent, making techni-O-helium the dominant form of composite dark matter in the considered scenario. On that reason, a possibility to hide the excessive $^4He$ in the techni-O-helium is elusive. On the contrary, even having taken into account possible systematic errors in the determination of primordial helium, to provide its abundance within the observed limits, one should constraint the amount of helium bound with $\\zeta^{--}$. Since this amount is determined by the techni-O-helium number density, the condition that techni-O-helium saturates the observed dark matter density leads to a lower limit for the mass of $\\zeta^{--}$.  We come to the conclusion that in the minimal WTC model, contrary to the case of the AC-model, WIMP-like component of composite atom-like dark matter should be sparse, so that the formation of large scale structure should follow a warmer than cold dark matter scenario of the techni-O-helium Universe considered earlier.  In addition to the detailed description of a warmer than cold dark matter model, another challenging problem that is left for future work is the nuclear transformations catalyzed by techni-O-helium. The question about their consistency with observations remains open since special nuclear physics analysis is needed to reveal what are the actual techni-O-helium effects in BBN and in terrestrial matter.  The latter effects inside the body of underground dark matter detectors can experience annual modulation and lead to ionization events with a few keV energy release. It can make techni-O-helium (as well as any other form of O-helium) an interesting candidate, which might explain the difference between the positive result of DAMA/NaI (DAMA/Libra) and negative results of other experiments on direct dark matter search.  The destruction of techni-O-helium by cosmic rays in the Galaxy releases free charged technileptons, which can be accelerated and contribute to the flux of cosmic rays. In this context, the search for techniparticles at accelerators and in cosmic rays acquires the meaning of a crucial test for the existence of the basic components of the composite dark matter. At accelerators, techniparticles would look like stable doubly charged heavy leptons, while in cosmic rays, they represent a heavy $-2$ charge component with anomalously low ratio of electric charge to mass. If it has the same energy spectrum as ordinary cosmic rays, it can be observed in the PAMELA experiment.  To conclude, the minimal walking technicolor cosmology can give a robust cosmological scenario of composite dark matter, giving rise to a set of exciting observable effects."
    },
    "0806/0806.1754_arXiv.txt": {
        "abstract": "We present the results of a search for all embedded protostars with internal luminosities $\\le$ 1.0 \\lsun\\ in the full sample of nearby, low-mass star-forming regions surveyed by the \\emph{Spitzer Space Telescope} Legacy Project ``From Molecular Cores to Planet Forming Disks'' (c2d).  The internal luminosity of a source, \\lint, is the luminosity of the central source and excludes luminosity arising from external heating.  On average, the \\emph{Spitzer} c2d data are sensitive to embedded protostars with \\lint\\ $\\geq 4 \\times 10^{-3}$ $(d/140 \\, \\rm{pc})^2$ \\lsun, a factor of 25 better than the sensitivity of the \\emph{Infrared Astronomical Satellite (IRAS)} to such objects.  We present a set of selection criteria used to identify candidates from the \\emph{Spitzer} data and examine complementary data to decide whether each candidate is truly an embedded protostar.  We find a tight correlation between the 70 \\um\\ flux and internal luminosity of a protostar, an empirical result based on both observations and detailed two-dimensional radiative transfer models of protostars.  We identify 50 embedded protostars with \\lint\\ $\\le$ 1.0 \\lsun; 15 have \\lint\\ $\\le$ 0.1 \\lsun.  The intrinsic distribution of source luminosities increases to lower luminosities.  While we find sources down to the above sensitivity limit, indicating that the distribution may extend to luminosities lower than probed by these observations, we are able to rule out a continued rise in the distribution below \\lint\\ $= 0.1$ \\lsun.  Between $75-85$\\% of cores classified as starless prior to being observed by \\emph{Spitzer} remain starless to our luminosity sensitivity; the remaining $15-25$\\% harbor low-luminosity, embedded protostars.  We compile complete Spectral Energy Distributions for all 50 objects and calculate standard evolutionary signatures (\\lbol, \\tbol, and \\lbolsmm), and argue that these objects are inconsistent with the simplest picture of star formation wherein mass accretes from the core onto the protostar at a constant rate. ",
        "introduction": "Recently, the \\emph{Spitzer Space Telescope} Legacy Project ``From Molecular Cores to Planet Forming Disks'' (c2d; Evans et al. 2003) completed an extensive $3.6-160$ \\um\\ imaging survey of nearby, low-mass star-forming regions.  One of the results to come out of this survey is the discovery of very low luminosity objects (VeLLOs; Young et al. 2004).  If the internal luminosity of a source, \\lint, is the total luminosity of the central protostar and circumstellar disk (if present), a VeLLO is defined to be an object embedded within a dense core with $\\lint \\leq 0.1$ \\lsun\\ (Di Francesco et al. 2007).  The bolometric luminosity of an embedded protostar, an observable quantity that can be calculated by integrating over the full Spectral Energy Distribution (SED), is composed of both internal and external luminosity ($\\lbol\\ = \\lint\\ + \\lext$).  The external luminosity is usually that arising from heating of the circumstellar envelope by the Interstellar Radiation Field (ISRF), and will add, on average, a few tenths of a solar luminosity to \\lbol\\ (e.g., Evans et al. 2001).  Thus, the distinction between \\lbol\\ and \\lint\\ is most relevant for embedded protostars with $\\lint \\la 1.0$ \\lsun, where the external luminosity can be a significant fraction of the observed \\lbol.  For VeLLOs, the external luminosity can dominate the observed \\lbol.  Radiative transfer modeling of the SEDs of embedded protostars, including both the emission from the envelope at submillimeter and millimeter wavelengths and the emission from the central source itself at infrared wavelengths, is required to decouple internal and external luminosities (e.g., Shirley et al. 2002; Young et al. 2004; Dunham et al. 2006).  Several VeLLOs have been discovered in cores that were previously classified as starless prior to being observed by \\emph{Spitzer}.  In fact, the very first starless core observed by c2d, L1014, was found to harbor a VeLLO with \\lint\\ $\\sim$ 0.09 \\lsun\\ (Young et al. 2004).  This discovery reinforces that the known sample of embedded protostars is not complete.  This sample has been assembled primarily by two methods:  (1) searching for \\emph{IRAS} sources that are associated with dense cores and have colors consistent with those expected for embedded protostars (e.g., Myers et al. 1987), and (2) identifying molecular outflows and radio point sources associated with dense cores indicating the presence of protostars too deeply embedded to detect with \\emph{IRAS} (e.g., \\andre\\ et al. 1993).  Myers et al. (1987) found that the \\emph{IRAS} data could detect protostars with \\lint\\ $\\ga 0.1$ $(d/140 \\, \\rm{pc})^2$ \\lsun, where $d$ is the distance to the protostar, although this does not include the younger, more deeply embedded protostars that were only identified on a case-by-case basis by the second method.  The regions surveyed by c2d with \\emph{Spitzer} are located at distances ranging from $125-500$ pc.  Even in the closest of these regions VeLLOs are likely to fall below the \\emph{IRAS} sensitivity limit.  In the more distant regions, no protostars with \\lint\\ $\\la$ 1 \\lsun\\ would be detected.  While some of these protostars might have been identified on a case-by-case basis as described above, the full sample of embedded protostars with \\lint\\ $\\le$ 1 \\lsun\\ is clearly incomplete.  Constructing a complete sample of embedded protostars with \\lint\\ $\\le$ 1 \\lsun\\ is important for studies of low-mass star formation.  Despite substantial progress in recent decades, the details of the physical processes regulating mass accretion from the envelope to the protostar remain poorly understood (see McKee \\& Ostriker [2007] for a recent review).  Several authors have attempted to constrain evolutionary models of the formation of low-mass stars by determining the observational signatures of these models and comparing them to the properties of known protostars (e.g., Myers et al. 1998; Young \\& Evans 2005).  A result common to all such studies is a substantial population of protostars with luminosities below model predictions.  An idea proposed to explain this discrepancy is that the mass accretion is episodic in nature and the protostars with the lowest luminosities are those observed in quiescent accretion states (e.g., Kenyon \\& Hartmann 1995; Young \\& Evans 2005; Enoch 2007; Enoch et al. 2008a, in preparation).  Theoretical studies have provided several mechanisms by which such a process may occur, such as material piling up in a circumstellar disk until gravitational instabilities drive angular momentum outward and mass inward in short-lived bursts (Vorobyov \\& Basu 2005, 2006).  Alternatively, quasi-periodic magnetically driven outflows in the envelope can cause mass accretion onto the protostar to occur in magnetically controlled bursts (Tassis \\& Mouschovias 2005).  Indeed, the evidence for non-steady mass accretion in young protostellar systems still in the embedded phase is steadily growing (e.g., Hartmann \\& Kenyon 1985; Dunham et al. 2006; Acosta-Pulido et al. 2007; K\\'{o}sp\\'{a}l et al. 2007; Fedele et al. 2007).  However, as the sample of embedded, low luminosity protostars is incomplete, the true nature of the discrepancy between evolutionary models and observations of protostars is unknown.  Future work devoted to assessing the validity of various models by comparing their predictions to the properties of known protostars depends on the existence of a sample that is as complete and unbiased as possible.  The VeLLOs are a particularly interesting subset of embedded, low-luminosity protostars; in essence, they are an extreme case of the problem discussed above.  To date, only three VeLLOs have been studied in detail:  L1014-IRS (\\lint\\ $\\sim$ 0.09 \\lsun; Young et al. 2004), L1521F-IRS (\\lint\\ $\\sim$ 0.06 \\lsun; Bourke et al. 2006), and IRAM 04191-IRS (\\lint\\ $\\sim$ 0.08 \\lsun; Dunham et al. 2006).  Despite the fact that all three have similar internal luminosities, they differ greatly in envelope and outflow properties, as discussed by Bourke et al. (2006).  IRAM 04191 drives a large, bipolar molecular outflow, features bright molecular line and dust continuum emission, and shows evidence for infall, depletion, and deuteration (\\andre\\ et al. 1999; Belloche et al. 2002).  L1521F is also bright in molecular line and dust continuum emission and also shows evidence for infall, depletion, and deuteration (Crapsi et al. 2004), but the envelope is not as centrally condensed as IRAM 04191 and the presence of an outflow is uncertain (Crapsi et al. 2004; Bourke et al. 2006).  L1014 does not show significant evidence for infall, depletion, or deuteration (Crapsi et al. 2005a), but it does drive a compact, weak molecular outflow detected only in interferometer observations (Bourke et al. 2005; Crapsi et al. 2005a).  A systematic search for all VeLLOs in the regions surveyed by c2d will allow their properties to be examined in detail both on a case-by-case basis and as a class of objects.  Identifying the complete sample of VeLLOs will also allow us to determine how many cores classified as starless prior to being observed by \\emph{Spitzer} truly are starless, a question with important implications for estimates of the lifetime of starless cores (e.g., Kirk et al. 2005).  In this paper, we present the results of a search for all embedded protostars with \\lint\\ $\\leq$ 1.0 \\lsun\\ in the full c2d imaging dataset.  Depending on the distance to each individual source, some will already have been detected by \\emph{IRAS}, while others will be new sources.  We consider this work to be complementary to several related studies:  A search by Kirk et al. (2007) for embedded protostars in 22 cores classified as starless prior to being observed by \\emph{Spitzer}; a search by J\\o rgensen et al. (2007; 2008, in preparation) for all embedded protostars in Perseus and Ophiuchus, regardless of luminosity, conducted by combining \\emph{Spitzer} and SCUBA 850 \\um\\ dust continuum emission data; and a search by Enoch (2007) and Enoch et al. (2008a, in preparation) for all embedded objects in the Perseus, Ophiuchus, and Serpens molecular clouds, regardless of luminosity, conducted by combining \\emph{Spitzer} and Bolocam 1.1 mm dust continuum emission data.  The key difference between the work presented here and the searches for embedded protostars listed above is that we do not start by identifying dense cores from their millimeter dust continuum emission and then look for associated \\emph{Spitzer} sources embedded within them.  Instead, we develop a set of criteria to identify candidate embedded, low-luminosity protostars based on the $3.6-70$ \\um\\ \\emph{Spitzer} data.  This way, we are able to identify all of the candidates in the full c2d dataset, regardless of the availability and quality of millimeter wavelength observations for each region.  Only after we identify all candidates based on \\emph{Spitzer} data alone do we turn to other observations to distinguish the objects of interest from various contaminants masquerading in our sample.  Our method identifies candidates for further examination once large-scale surveys of nearby star-forming regions are completed with SCUBA-2 (Ward-Thompson et al. 2007) and Herschel, and the method can easily be extended to search for embedded, low-luminosity protostars in the additional nearby, low-mass star-forming regions being surveyed by the \\emph{Spitzer} Gould Belt Legacy Project (L. Allen et al. 2008, in preparation).  The organization of this paper is as follows:  In \\S \\ref{observations}, we provide a brief description of the c2d observations and data reduction, emphasizing those aspects relevant to this work.  The criteria for identifying candidate embedded, low-luminosity protostars from the $3.6-70$ \\um\\ \\emph{Spitzer} data are discussed in \\S \\ref{id}, along with the possibilities for estimating source internal luminosities directly from observable quantities.  A general proof-of-concept demonstrating the validity of these criteria is given in \\S \\ref{proofofconcept}.  In \\S \\ref{confirmation}, we discuss the contamination expected in the list of candidates, both from background extra-galactic sources and from more evolved Young Stellar Objects (YSOs) no longer embedded within their dense cores.  We discuss the necessary requirements to prove that a candidate is truly an embedded protostar in \\S \\ref{prove}, we apply these requirements to our candidate list in \\S \\ref{groups}, and we discuss the difficulties in including regions lacking good quality 70 \\um\\ data in \\S \\ref{need70}.  We discuss several general results of this work in \\S \\ref{discussion}.  Finally, we present our conclusions in \\S \\ref{conclusions}. ",
        "conclusions": "We have conducted a search for all embedded protostars with \\lint\\ $\\le$ 1.0 \\lsun\\ in the c2d dataset of nearby, low-mass star-forming regions.  We identify 218 candidates from the \\emph{Spitzer} data alone; examining all available complementary data for each candidate results in a sample of 50 objects that show at least some evidence that they are indeed embedded within dense cores.  A summary of our major results is as follows:  \\begin{itemize} \\item On average, the \\emph{Spitzer} c2d data are sensitive to embedded protostars with \\lint\\ $\\geq 4 \\times 10^{-3}$ $(d/140 \\, \\rm{pc})^2$ \\lsun, a factor of 25 better than the sensitivity of the \\emph{Infrared Astronomical Satellite (IRAS)} to such objects. \\item The 70 \\um\\ flux and internal luminosity of a protostar are tightly correlated.  As the former is a directly observable quantity but the latter is not, this correlation gives a powerful method for estimating protostellar internal luminosities when detailed radiative transfer models for each source are lacking. \\item Of the 50 objects in our sample, 15 (30\\%) have \\lint\\ $\\leq$ 0.1 \\lsun\\ and are thus classified as VeLLOs.  The distribution of source luminosities is not uniform and instead increases with decreasing luminosity.  Accounting for incompleteness arising from non-uniform distances to the observed regions, we find sources down to the above sensitivity limit, indicating that the intrinsic luminosity distribution may extend to lower luminosities than probed by these observations.  Despite this, we are able to rule out a continued rise in the distribution below \\lint\\ $= 0.1$ \\lsun. \\item Between $75-85$\\% of cores classified as starless prior to being observed by \\emph{Spitzer} remain starless down to the above luminosity sensitivity; the remaining $15-25$\\% harbor low-luminosity, embedded protostars.  This is in general agreement with Kirk et al. (2007), who examined archival \\emph{Spitzer} data of 22 starless cores and found only one to be harboring a low-luminosity protostar.  However, with our larger sample size, we are able to better constrain the fraction of cores previously classified as starless that in fact harbor such objects.  We confirm that recent estimates of starless core lifetimes (e.g., Kirk et al. 2005; Enoch et al. 2008b) do not feature large errors introduced by previously undetected, low-luminosity protostars. \\item The observed luminosity distribution for embedded objects with \\lint\\ $\\le$ 1.0 \\lsun\\ is inconsistent with the simplest picture of star formation wherein mass accretes from the core onto the protostar at a constant rate.  Combining this result with other studies that find clear indications of episodic outflow activity strongly suggests that protostellar mass accretion is episodic in nature. \\end{itemize}  We have outlined several avenues of future work that must be pursued now that relatively complete and unbiased samples of embedded, low-mass protostars are being compiled.  Only with such future studies can we begin to build a coherent picture of low-mass star formation consistent with the growing observational database provided by systematic, large-scale surveys of low-mass star forming regions."
    },
    "0806/0806.3798_arXiv.txt": {
        "abstract": "Thanks to the availability of high-resolution high-sensitivity telescopes such as the Very Large Array, the \\emph{Hubble Space Telescope}, and the \\emph{Chandra X-ray Observatory}, there is now a wealth of observational data on relativistic jets from active galactic nuclei (AGN) as well as galactic sources such as Black-Hole X-ray Binaries.  Since the jet speeds cannot be constrained well from observations, but are \\Rev{generally} believed to be relativistic, physical quantities inferred from observables are commonly expressed in terms of the unknown beaming parameters: the bulk Lorentz factor and the line-of-sight angle, usually in their combination as relativistic Doppler factor.  This paper aims to resolve the discrepancies existing in the literature about such ``de-beaming'' of derived quantities, in particular regarding the minimum-energy magnetic field estimate.  The discrepancies arise because the distinction is not normally made between the case of a fixed source observed with different beaming parameters and the case where the source projection on the sky is held fixed. The former is usually considered, but it is the latter that corresponds to interpreting actual jet observations.  Furthermore, attention is drawn to the fact that apparent superluminal motion has a spatial corollary, here called ``retardation magnification'', which implies that most parts of a relativistic jet that are actually present in the observer's frame (a ``world map'' in relativity terminology) are in fact hidden on the observer's image (the ``world picture'' \\Rev{in general, or ``supersnapshot'' in the special case of astronomy}). ",
        "introduction": "\\label{s:intro}  For over 50 years from the appearance of the seminal \\emph{Zur Elektrodynamik bewegter K\\\"orper} \\citep{Einstein05}, only the Lorentz transformations of the 4-coordinates of ``events'' were considered in the literature, but not how relativistically moving bodies would \\emph{appear} when looked at or photographed.  This was first done independently and (relatively) simultaneously by \\citet{Pen59} and \\citet{Terrell59}; the former showed that the projected outline of a relativistically moving sphere is always a circle, while the latter provided a more extensive discussion of the appearance of moving bodies and pointed out the key features of observing relativistically moving bodies: they appear both \\emph{rotated} and \\emph{scaled} (more details will be given below).  The motivation for writing the present paper is work on interpreting observations of relativistic jets \\citep[e.g.,]{JHMM06}, where the need arises to infer physical properties of the jet fluid in its own rest frame from observations, subject to corrections due to relativistic beaming, whose magnitude is, however, not known from observations.  A particular quantity of interest is the rest-frame minimum-energy magnetic-field estimate for a synchrotron source \\citep{Bur59}, and there are different opinions in the literature about how the true rest-frame minimum-energy field scales with the relativistic Doppler factor compared to that inferred assuming a non-relativistic source \\citetext{compare eqn.~A3 of \\citealt{SSO03} to eqn.~A7 of \\citealt{HK02}}.  Some of the argument revolves around whether the observed morphological features of jets are ``blobs'' or ``jets'', and their (apparently) different beaming properties.  Given these differences of opinion on how to ``de-beam'' properly, it is perhaps surprising that NASA's Astrophysics Data System lists only three papers on interpreting jet observations as citing \\citet{Terrell59}, and his results do not seem to be part of the common knowledge of jet researchers.  One of the citing papers is \\citet{LB85}, who consider the implications of relativistic beaming on the difference between observed and intrinsic source counts and give detailed formulae for relating observed and jet-frame fluxes and emissivities.  Some of these formulae had already been presented in the seminal paper by \\citet{BK79}.  It appears that the difficulties in interpreting jet observations arise because the problem under consideration is ill-posed. As will be argued in detail below, what matters for interpreting jet observations subject to unknown beaming paramters is that we have observed the 2-dimensional projection of a source's appearance onto the plane of the sky and try to infer the source's rest-frame properties from this projection.  Confusion arises because most formulae in the literature consider what happens to the observed quantities when a \\emph{fixed source} moves with different Lorentz factors and at different line-of-sight angles to the observer, while in observations, it is the \\emph{projection} of the source which is held constant. \\Rev{Furthermore, the effects of light-travel time delays along the line of sight are typically only mentioned explicitly in work comparing jet simulations to observations, e.g., in \\citet{AMGea03} and \\citet{SwiftHughes08}, but not in the observational literature.} This and the preference for adopting the fixed-source view may be related to the fact that Lorentz transformations are usually covered in great detail in a typical course on special relativity, but \\citet{Pen59} and \\citet{Terrell59} are hardly mentioned in relativity textbooks.  This leads to the present paper with the following outline: the remainder of the introduction summarizes the results by \\citet{Pen59} and \\citet{Terrell59} and sets out some basic definitions and terminology.  The appearance of relativistic objects, and in particular of astrophysical jets, is discussed in \\S\\ref{s:appearance}, both from a theoretical point of view and using a simple ray-tracer. Ready-to-use formulae for relating jet-frame quantities to observables are given in \\S\\ref{s:beaming}, including the minimum-energy field.  The discussion and summary are given in \\S\\ref{s:disc}, while Appendix~\\ref{s:illustr.lor} describes some \\emph{Gedankenexperimente} on non-conventional world-map measurements that lead to length expansion and time acceleration.  \\subsection{World Pictures and Supersnapshots} \\label{s:intro.worldsnap}  As first noted by \\citet{Terrell59}, there is a fundamental difference in relativity between the \\emph{locations} (4-coordinates) of events as judged by observers that are local to the events and equipped with sets of clocks that are synchronized in their rest frame, and the \\emph{appearance} of relativistically moving bodies as judged by distant observers by means of photons that are received simultaneously; a little earlier, \\citet{Pen59} had considered the special case of the observed outline of a relativistically moving sphere.  The set of event locations is a \\emph{world map}, while the picture that is taken of the events is a \\emph{world picture}.  In the special case of photons arriving at right angles to the detector taking the world picture, it is called a \\emph{supersnapshot} \\citep{Rindler77}.  Astronomical observations clearly fall under the definition of a supersnapshot.  The appearance of relativistically moving objects in a supersnapshot is governed by two aspects of photon paths in special relativity \\citep{Terrell59,Rindler77,LB85}: \\begin{enumerate} \\item \\Rev{Two photons traveling abreast with a separation $\\Delta s$ in one frame (i.e., photons traveling ``alongside each other'' with $\\Delta s$ measured perpendicular to their direction of motion) do so in \\emph{all} frames. This is the case because $|\\Delta s|^2$ is invariant under Lorentz transformations and $\\Delta s$ is a space-like interval.} \\item If a photon is traveling at an angle $\\theta\\pr$ to the direction of motion of some frame that is moving with speed $\\beta c$ and Lorentz factor $\\Gamma = (1-\\beta^2)^{1/2}$ with respect to an observer, the angle between the direction of motion and the photon direction in that frame is related to the angle $\\theta$ between the direction of motion and the photon direction in the observer's frame by \\begin{equation} \\mu\\pr = \\frac{\\mu - \\beta}{1-\\beta\\mu}, \\label{eq:costrans} \\end{equation} where $\\mu = \\cos\\theta$ etc., or, equivalently, \\begin{equation} \\sin \\theta\\pr = \\D \\sin\\theta, \\label{eq:sintrans} \\end{equation} where \\D\\ is the relativistic Doppler factor \\begin{equation} \\D = \\left[\\Gamma (1-\\beta \\mu) \\right]^{-1}. \\label{eq:Dopplerdef} \\end{equation} \\end{enumerate} The latter phenomenon is the well-known angle aberration; the former is perhaps less well-known, but essential for the analysis of images of relativistically moving objects, and implies that the supersnapshot is a \\emph{scaled} version of the rest-frame image.  Taken together, they yield Terrell's result that the appearance of such an object in a supersnapshot is simply the object's appearance as seen from the aberrated angle $\\theta\\pr$ in its rest frame, with its apparent size along the direction of motion scaled by the Doppler factor $\\D$.  As a consequence of eq.~(\\ref{eq:costrans}), even approaching objects appear to be seen ``from behind'' unless $\\mu < \\beta$, i.e., $\\D > \\Gamma$; in the limiting case $\\mu=\\beta \\Leftrightarrow \\D=\\Gamma \\Leftrightarrow \\sin\\theta = 1/\\Gamma$, a relativistic object is seen exactly side-on in its rest frame and with exactly its rest-frame length as its ``projected'' length.     \\subsection{Terminology: ``Blobs'' \\emph{versus} ``jets'' \\emph{versus} ``shocks''  \\Rev{-- at rest in different frames}} \\label{s:intro.terminology}  It is useful clearly to set out the terminology for the remainder of the paper, because the brightness pattern observed in astrophysical jets can be at rest in frames that are different from both the observer frame, and the fluid rest frame, as discussed in detail by \\citet{LB85}.  \\Rev{Their discussion and notation is adopted here. It distinguishes between ``blob'',  ``jet'' and ``shock'' features, which are defined by being at rest in one of three frames relevant to the problem.  Thus, it is useful to give the definitions of the relevant frames together with those of the morphological terms:} \\begin{enumerate} \\item \\Rev{The ``observer frame'' is that in which the astronomer is at rest.  Once appropriate cosmological corrections are applied, the observer frame is conceptually identical to the frame in which the jet source and its host are at rest.}  \\Rev{A ``jet'' feature is then a \\emph{resolved} brightness pattern whose outline is at rest in the observer frame. Observer-frame quantities have no primes, e.g. $j$ for volume emissivity.} \\item \\Rev{The ``fluid frame'' is the rest frame of the emitting fluid, which is taken to be moving through the observer frame at relativistic speed.  The term ``rest frame'' is used interchangeably with ``fluid frame''.}\\footnote{The emitting fluid is not necessarily identical with that carrying the bulk of the jet's kinetic energy, nor are those two fluids necessarily moving at the same speed \\protect\\citep{HK07}. However, this distinction does not affect the relation between observables and physical quantities in the rest frame of the emitting fluid, which is the subject of this paper. Nevertheless, it needs to be kept in mind when interpreting fluid-frame quantities.}  \\Rev{A ``blob'' or ``plasmoid'' is a brightness pattern whose outline is at rest in the fluid frame.  Fluid-frame quantities will be designated by double primes, e.g. $j\\dpr$.} \\item A ``pattern'' or ``shock'' feature is a brightness pattern whose outline is at rest neither in the fluid nor in the observer frame, \\Rev{e.g., a shock traveling through the jet fluid.  It defines the third frame, the frame in which this pattern is at rest.}  Pattern-frame quantities have single primes, e.g. $j\\pr$.  The emissivity of the fluid traveling through such a ``shock'' transforms according to the fluid's Doppler factor $\\D\\dpr$, while its projected appearance and morphology are governed by the shock's Doppler factor $\\D\\pr$. \\end{enumerate} These are fairly intuitive definitions. \\Rev{Nevertheless, the difference between the ``blob'' and ``jet'' formulae in eqn.~C7 of \\citet{BBR84}} is just one of \\emph{choice of integration boundaries}, and in particular whether the integration boundaries are held fixed in the observer frame when $\\D$ is changed (jet case) or are allowed to vary according to the different projected morphology of a ``blob'' under changes of $\\D$. Thus, it is possible to apply a ``jet'' formula to a small segment of a blob as long as the integration boundaries are held fixed in the observer frame.  In \\S\\S\\ref{s:beaming.blobs.restframe} and \\ref{s:beaming.jets.restframe} below, I will present detailed formulae for converting observed to fluid-frame properties in each case, with expressions for the minimum-energy field in \\S\\ref{s:disc.Bmin.obsfixed}.  \\subsection{Basic definitions and beaming formulae} \\label{s:intro.beamformulae}  This section summarizes the basic definitions of surface brightness/intensity, flux density and luminosity of astronomical sources, as well as the beaming properties of blobs, jets, and shocks. I will give explicit formulae for observed surface brightness and total flux in terms of source parameters for simple geometries, as well as ray-tracing images showing the appearance of such sources in supersnapshots.  For the computation of surface brightness and flux, I use the notation and formulae as given by \\citet{BK79} and \\citet{LB85}, assuming an optically thin, isotropic emission with a power-law emissivity $j_{\\nu} \\propto \\nu^{\\alpha}$ that is constant within the emitting region.  All observables will be expressed in terms of the emissivity $j\\dpr$ in the fluid rest frame and the source size in the pattern frame $\\Sigma\\pr$, which is identical to the fluid and observer frame for a ``blob'' and ``jet'', respectively.  \\Rev{Cosmological transformations, however, are not always given explicitly in order to simplify the notation; they can be re-incorporated in the usual way by inserting appropriate powers of $(1+z)$ for cosmological redshifts, and using the appropriate cosmological distance measures.}  The surface brightness or intensity $I$, flux density $S_\\nu$ and luminosity $L$ of a source are given by \\begin{eqnarray} I_\\nu &=& \\int_0^{s} j_\\nu \\de x,\\\\ S_\\nu &=& \\int_A I_\\nu \\de A\\\\ &=& \\dL^{-2} \\int_V j_\\nu \\de V,\\\\ L_\\nu &=& 4 \\pi \\dL^2 S_\\nu \\nonumber \\\\ &=& 4\\pi \\int_V j_\\nu \\de V, \\label{eq:Ldef} \\end{eqnarray} where \\dL\\ is the luminosity distance to the source, which has specific emissivity $j$, volume V and projected surface area A.  The transformation properties of these quantities then follow from the relativistic invariance of $I_\\nu/\\nu^3$ and the volume transformation \\citep[taken from Appendix C of][]{BBR84}: \\begin{eqnarray} \\nu &=& \\D\\dpr\\; \\nu\\dpr \\label{eq:nutrans} \\\\ \\de\\Omega &=& \\D\\dpr^{-2}\\;\\de\\Omega\\dpr\\\\ I_\\nu(\\nu) &=& \\D\\dpr^3\\;I\\dpr_{\\nu\\dpr}(\\nu\\dpr)\\\\ j_\\nu(\\nu) &=& \\D\\dpr^2\\;j\\dpr_{\\nu\\dpr}(\\nu\\dpr) \\label{eq:jtrans} \\end{eqnarray} Assuming optically thin emission makes the discussion appropriate for arcsecond-scale jets, where sources are not compact enough for self-absorption to become important.  The difficulties of interpreting observations of optically thick sources, such as compact cores and milli-arcsecond scale jets, have been highlighted by \\citet{BK79} and \\citet{LB85}. \\Rev{The essential point here is that the appearance of optically thick sources varies as function of viewing direction, and the relativistic angle aberration implies that the beamed appearance is governed by this intrinsic viewing angle dependence in addition to the flux and surface brigthness beaming.} The volume transformation deserves separate consideration.  \\subsection{Volume transformation of relativistic objects in astronomical images} \\label{s:intro.voltrans}  The fact that the outline of the different kinds of brightness pattern is at rest in different frames has led some authors to write down different volume transformation formulae for astronomical observations of ``jets'' and ``blobs'' \\citetext{see \\citealp{SMMea97}, Appendix~A, and \\citealp{SSO03}, Appendix~A, e.g.}.  However, what matters for the volume transformation of a feature identified in an astronomical image or radio map is only that the image is a supersnapshot.  What matters for the interpretation of the supersnapshot is the volume of fluid whose photons arrive simultaneously on the \\emph{supersnapshot}, not the volume of fluid that is located within the jet volume in the \\emph{world map}.  Hence, the correct volume transformation for any fluid volume $V\\dpr$ observed by means of a supersnapshot is \\begin{equation} V = \\D\\dpr V\\dpr, \\label{eq:Vtrans} \\end{equation} where $\\D\\dpr$ is the Doppler factor of the fluid in the observer frame.  If the decisive criterion was not the fact that astronomical observations are supersnapshots, one could argue with equal justification that the correct volume transformation formula for the fluid in a jet section is $V\\dpr = V/\\Gamma$ because the jet volume is at rest in the observer frame and hence appears contracted in the rest frame of the fluid, or alternatively that the correct transformation is $V\\dpr = V \\times \\Gamma$ because the fluid is moving through the observer frame, and therefore \\emph{it} is contracted.  Both can of course be correct, depending on whether one is \\Rev{judging the jet volume with the help of events that are simultaneous in the fluid or the observer frame}. However, a supersnapshot corresponds to neither world-map case --- the supersnapshot criterion is photons \\emph{arriving} simultaneously at the observer, which nearly always does not correspond to photons \\emph{being emitted} simultaneously in any frame.  That eq.~(\\ref{eq:Vtrans}) is correct for both the ``jet'' and ``blob'' case can be seen also by considering a section of a ``jet'' as a collection of infinitesimal blobs that are each at rest in the fluid frame. Alternatively, a ``jet'' can be considered as a section of a ``blob'' that is moving through a transparent gap in obscuring material that is at rest in the observer frame --- if 90\\% of a blob's volume is covered in the observer frame, the rest-frame volume of the visible part is 10\\% of the blobs's total observer-frame volume, and hence must also be 10\\% of the blob's rest-frame volume.  \\Rev{As an alternative derivation of eq.~(\\ref{eq:Vtrans}), consider that the observer-frame volume of a ``jet'' or ``blob'' (or an infinitesimal element of it) is given by \\begin{displaymath} V = s \\times l \\times h, \\end{displaymath} where $s$ is its extent transverse to the line of sight in the plane of its motion, $h$ is the extent perpendicular to both the line of sight and the direction of motion, and $l$ is along the line of sight.  The individual factors of $V$ transform into the fluid rest frame as follows.}  \\Rev{First, since $h$ is perpendicular to the direction of motion, it is not affected by relativity in any way, and $h\\dpr = h$. Next, recall from \\S\\ref{s:intro.worldsnap} above that the transverse separation $\\Delta s$ of two photon paths, i.e., light rays, is Lorentz-invariant.  The transverse extent $s$ is defined by two such parallel light rays and therefore it is also Lorentz-invariant, hence $s\\dpr = s$.  Finally, to determine the transformation properties of $l$, consider the following argument. The optical depth $\\tau$ along $l$ has to be Lorentz-invariant since it encodes the fraction $e^{-\\tau}$ of photons that are absorbed by the jet material, which is independent of the motion of any observer \\citep[p.\\,147]{RL79}.  By definition, the optical depth is \\begin{displaymath} \\tau = l \\, \\kappa_\\nu, \\end{displaymath} where $\\kappa_\\nu$ the absorption coefficient of the material. The Lorentz invariance of $\\tau$ therefore implies that $l$ transforms inversely to $\\kappa_\\nu$. From the Lorentz invariance of $\\nu \\kappa_\\nu$ \\citep[again see][]{RL79}, it follows that $l$ transforms as $\\nu$, i.e., $l = \\D\\dpr l\\dpr$. Hence $ V = s \\, l h = s \\, \\D\\dpr l\\dpr \\, h = \\D\\dpr \\, s\\dpr \\, l\\dpr \\, h\\dpr = \\D\\dpr \\, V\\dpr$, again yielding eq.~(\\ref{eq:Vtrans}).}  Thus, the relation between rest-frame and observer-frame volume for supersnapshots is always given by eq.~(\\ref{eq:Vtrans}), no matter whether we are considering a ``jet'', ``blob'' or even ``shock'' feature.  As noted at the end of the preceding section, the well-known apparent difference between the beaming formulae for a blob and a jet \\citep[$\\D^{2-\\alpha}$ versus $\\D^3$, such as in App.\\ C7 of][]{BBR84} is in fact just a difference of \\emph{integrands}; since the \\emph{integration boundaries} differ depending on whether an object is considered as blob or jet, the final answer is independent of the assumed geometry. In other words, \\textbf{jets and blobs have the same beaming properties if identical source volumes are considered}. The equivalence of jet and blob formulae will be shown explicitly in \\S\\ref{s:beaming.jets.restframe} and \\ref{s:disc.Bmin.obsfixed} below.  While eq.~(\\ref{eq:Vtrans}) appears straightforward to interpret, the supersnapshot is merely a projection of the observer-frame volume onto the plane of the sky, so that \\textbf{the observer-frame volume $V$ is not a direct observable} (see Fig.~\\ref{f:ray_sph} below). Therefore, the volume formula can only be used for interpreting astronomical images if an assumption is made about the geometry of the source. However, its use in determining observables from known rest-frame quantities is straightforward.  \\begin{figure*} \\includegraphics[width=168mm]{f1.eps} \\caption{\\Rev{Illustration of retardation magnification and hiding, showing a sequence of events in which a relativistically moving ``blob'' is emitted by some source (e.g., an accretion disk around a black hole), and the picture recorded by a distant observer at the corresponding time.  The upper frames give the \\emph{world map} in the $(x,y)$ plane with the true locations of all events; for an infinite speed of light, the world map corresponds to the ``top view'' of the events as seen by an observer at 90\\degr\\ to the blob's direction of motion. The lower frames give the \\emph{supersnapshot}, the image projected onto the $(x,z)$ plane as recorded by a distant observer looking along the $+y$ axis by means of simultaneously arriving photons, i.e., photons that are crossing the dashed ``screen'' line simultaneously.  Panel \\textbf{(0)} shows the setup, with the black dot marking the location of the source (``black hole'') ejecting the relativistic blob. \\textbf{(1)} The front end ``A'' of the blob is ejected. \\textbf{(2)} The rear end ``B'' of the blob is ejected, and at the same location as ``A'' was in frame (1).  ``A'' itself has travelled some distance from the black hole. The curved line illustrates the current location of the wavefront by which the observer will later imply that ``A'' has been ejected.  \\textbf{(3)} The wavefront from the ejection of the front end ``A'' reaches the ``screen'' location and appears on the observer's picture.  The second wavefront carrying the information about the ejection of ``B'' is lagging behind. \\textbf{(4)} The light from the ejection of the rear end ``B'' reaches the screen location.  At the same time, the front end ``A'' crosses the screen location.  Therefore ``B'' and ``A'' appear at the shown locations on the supersnapshot. The separation B--A on the supersnapshot is greater than it is in the world map, and the observer records a magnified image of the blob.  If any further material is ejected after ``B'' (and hence occupied the region indicated by the dotted line), it will not yet be visible to the observer.  Hence, the apparent magnification of the blob's extent implies that any further ejections will be unobservable until the light emitted by them has had time to reach the observer, thus (temporarily) being hidden from view.}} \\label{f:magnif} \\end{figure*} ",
        "conclusions": "\\label{s:disc}  In \\S\\ref{s:appearance}, I have illustrated the difference between the \\emph{world map} of a relativistic jet, what is actually there, and the \\emph{world picture}, or its special case, the \\emph{supersnapshot}, that corresponds to what is observable by distant astronomers.  For the quantitative interpretation of world pictures in the presence of unknown beaming parameters (Lorentz factor $\\Gamma$ and angle $\\theta$ between the fluid motion and the line of sight), what matters is that the \\emph{projected appearance} of the jet is kept fixed and not the \\emph{intrinsic volume}.  This gives rise to de-beaming formulae that are slightly different from those in the existing literature.  \\subsection{Implications for interpretation of flux and surface brightness of jet features}  The most important conclusion for the quantitative analysis of jet observations is that the scaling relations relating rest-frame quantities (volume emissivity, intrinsic source size) to observables (projected source size, surface brightness, total flux) \\emph{cannot} be stated as function of the line-of-sight angle and Lorentz factor in a general way, but depend on the details of the source geometry.  It is possible to write down explicit scaling relations for certain simple geometries such as spheres and elongated, rotationally symmetric blobs of constant cross-section. For other shapes, such as ellipsoidal blobs or blobs with non-symmetric cross-sections, the projected appearance is affected by edge effects, and additionally by observational effects such as the contrast between the faintest parts of the source and the sky background, as well as the available signal-to-noise level.  Edge effects are properly taken into account by the ray-tracing in \\S\\ref{s:illust.rays}, and such ray-tracing modeling of observables is probably the most accurate route to interpreting observations of relativistic objects. Indeed, it is part of the prediction of observables from jet simulations such as those by \\citet{AMGea03}, e.g. \\Rev{The work of \\citet{SwiftHughes08}, which explicitly considers the relation between the jet appearance in a supersnasphot and the underlying physical quantities, taking into account the retardation along the line of sight.}  \\subsection{Implications for interpretation of morphological information in jet images}  Most radio, optical and X-ray maps of relativistic jets \\citep[a list of radio jets is given by][]{LiuZhang02}\\footnote{See \\url{http://home.fnal.gov/~jester/optjets/} and \\url{http://hea-www.harvard.edu/XJET/} for lists of optical and X-ray jets.} show a series of well-separated, distinct features usually referred to as ``knots'', with diffuse emission linking them. Given that relativistic beaming favours the detection of objects with jets at small angles to the line of sight, and the superluminal motions detected in the cores of many such sources, it is plausible that the jet material itself is still relativistic even at large separations (and indeed, this is required in models accounting for the X-ray emission from powerful radio jets as beamed inverse-Compton scattering of cosmic microwave background photons; see \\citealp{Tav00,Cel00} for the original development of the idea, as well as the recent review by \\citet{HK06}).  However, what is not clear is whether the knots themselves are stationary shock features, or themselves moving relativistically.  Referring to Figs.~\\ref{f:ray_sph} and \\ref{f:ray_slab}, the prevalence of well-separated knots in jet images seems to suggest that the knots are moving at least with mildly relativistic Lorentz factors -- otherwise, there should be \\Rev{\\emph{some} jets observed at small angles (favoured by Doppler boosting)} where different knots overlap along the line of sight, washing out any individual morphological features.  \\Rev{In this case, the knots are subject to retardation magnification and hiding (illustrated in Figs.~\\ref{f:magnif} and \\ref{f:angles}), and we are not seeing \\emph{all} of the jet features which are actually present between core and hot spot, but just a small fraction (whose magnitude is given by Fig.~\\ref{f:visfrac}).  An observation of apparent superluminal motion of individual knots would be a direct confirmation that they are moving relativistically, as in the case of parsec-scale knots in VLBI observations.  The kiloparsec-scale knots are resolved out in VLBI observations, so that very long-term monitoring programmes at sub-arcsecond spatial resolution are required to make a potential superluminal motion observable.}  If jet knots are indeed moving relativistically, the retardation magnification and hiding need to be taken into account when interpreting morphological observables such as the ratio between knot separation and jet width, which is important for addressing the question of the origin of the jets' morphological features, e.g., whether they arise from instabilities \\citep{Hardee03} or as manifestation of a stable magnetohydrodynamical configuration \\citep{KC85}.  \\subsection{Conclusion}  It becomes clear once more that relativistic effects are counter to our non-relativistic intuition, and that familiarity with \\emph{world maps}, Lorentz transformations and the resulting phenomena of length contraction and time dilation is not sufficient for interpreting \\emph{world pictures}. When considering the beaming properties of quantities expressed in terms of integrals, such as a surface brightness, flux or the minimum-energy magnetic field estimate, one needs to consider the transformation properties both of the integrand and the integration volume.  Apparent differences between the beaming properties of ``blobs'' and ``jets'' disappear when the same source volume is considered.  Finally, the de-beaming of astronomical observations needs to be done not for a fixed source, but for the fixed \\emph{projection} of the source.  Doing so resolves some conflicts (again only apparent ones) between different de-beaming formulae in the literature.  Given that astronomy provides only world pictures, the concepts first laid out by \\citet{Pen59} and \\citet{Terrell59} deserve more attention in the interpretation of jet observations."
    },
    "0806/0806.3751_arXiv.txt": {
        "abstract": "{}% {We determine the components of the $\\Lambda$-effect tensor that quantifies the contributions to the turbulent momentum transport even for uniform rotation.} {Three-dimensional numerical simulations are used to study turbulent transport in triply periodic cubes under the influence of rotation and anisotropic forcing. Comparison is made with analytical results obtained via the so-called minimal tau-approximation.}% {In the case where the turbulence intensity in the vertical direction dominates, the vertical stress is always negative. This situation is expected to occur in stellar convection zones. The horizontal component of the stress is weaker and exhibits a maximum at latitude $30\\degr$ --- regardless of how rapid the rotation is. The minimal tau-approximation captures many of the qualitative features of the numerical results, provided the relaxation time tau is close to the turnover time, i.e.\\ the Strouhal number is of order unity. }{} ",
        "introduction": "Differential rotation plays a crucial role in dynamo processes that sustain large-scale magnetic activity in stars like the Sun (e.g.\\ Moffatt \\cite{Moffatt1978}; Krause \\& R\\\"adler \\cite{KrauRaed1980}). The internal rotation of the Sun is familiar from helioseismology (e.g.\\ Thompson et al.\\ \\cite{Thompsonea2003}), but the processes sustaining the observed rotation profile are not understood well. The angular momentum balance of a rotating star is determined by the conservation equation% \\EQ% \\frac{\\pd}{\\pd t} (\\rho s^2 \\mean\\Omega) + \\bm\\nabla \\cdot (\\rho s^2 \\mean\\Omega\\; \\meanv{U} + \\rho s \\mean{u_\\phi \\bm{u}}) = 0\\;,% \\EE% where $\\meanv{U}$ is the meridional flow, $s$ the cylindrical radius, $\\rho$ the density (neglecting however its fluctuations), $\\mean\\Omega=\\mean{U}_\\phi/s$ the local angular velocity, and $\\mean{u_\\phi \\bm{u}}$ the zonal component of the Reynolds stress. Overbars denote averages over the azimuthal direction.  The meridional flow can also be directly affected by the Reynolds stresses (e.g.\\ R\\\"udiger \\cite{Ruediger1989}), but it is more strongly determined by the baroclinic term that arises if the isocontours of density and pressure do not coincide. Such a configuration can appear because of latitude-dependent turbulent heat fluxes that arise naturally in rotating convection (e.g.\\ R\\\"udiger \\cite{Ruediger1982}; Pulkkinen et al.\\ \\cite{Pulkkinenea1993}; K\\\"apyl\\\"a et al.\\ \\cite{Kaepylaeea2004}; R\\\"udiger et al.\\ \\cite{Ruedigerea2005a}) or from a subadiabatic tachocline (Rempel \\cite {Rempel2005}) which is likely to occur in the Sun (Rempel \\cite{Rempel2004}; K\\\"apyl\\\"a et al.\\ \\cite{Kaepylaeea2006c}). The flows due to thermodynamic effects are likely to be needed to avoid the Taylor--Proudman balance in the Sun (e.g.\\ Durney \\cite{Durney1989}; Brandenburg et al.\\ \\cite{Brandenburgea1992}; Kitchatinov \\& R\\\"udiger \\cite{KitRued1995}; Rempel \\cite{Rempel2005}). The overall importance of the meridional flow in the angular momentum balance of the Sun is, however, still unclear since no definite observational information about it is available below roughly 20\\,Mm depth (e.g.\\ Zhao \\& Kosovichev \\cite{ZhaoKoso2004}).  Although not much more is known about the Reynolds stresses from observations, already this limited knowledge can be used to gain insight into the theory of turbulent transport. Solar surface observations indicate that there is an equatorward flux of angular momentum, as described by the Reynolds stress component $Q_{\\theta \\phi} = \\overline{u_\\theta u_\\phi}$, of several $10^3$\\,m$^2$\\,s$^{-2}$ in the latitude range where sunspots are observable (e.g.\\ Ward \\cite{Ward1965}; Pulkkinen \\& Tuominen \\cite{PulkTuo1998}; Stix \\cite{Stix2002}). In mean-field theory the simplest approximation that can be made concerning the Reynolds stresses is to assume them proportional to the gradient of mean velocity (the Boussinesq ansatz): \\begin{equation} Q_{ij} \\equiv \\overline{u_i u_j} = - \\mathcal{N}_{ijkl} \\overline{U}_{k,l}\\;. \\end{equation} In the Sun this ansatz turns out to be insufficient because the observed $Q_{\\theta \\phi}$ and solar surface differential rotation profile indicate that the turbulent viscosity is negative. Thus, in analogy to mean-field dynamo theory, additional contributions to the Reynolds stress were conjectured to appear (e.g.\\ Wasiuty$\\acute{\\rm n}$ski \\cite{Wasiutynski1946}; Krause \\& R\\\"udiger \\cite{KrauseRued1974}), leading to the present description \\begin{equation} Q_{ij} =  \\Lambda_{ijk} \\mean\\Omega_k - \\mathcal{N}_{ijkl} \\overline{U}_{k,l}\\;, \\end{equation} where $\\Lambda_{ijk}$ is a third-rank tensor describing the $\\Lambda$-effect that contributes to the Reynolds stress even in the case of rigid rotation. These terms are often referred to as ``non-diffusive'' contributions to the Reynolds stress. The zonal components of the stress can be written in the form (e.g.\\ Stix \\cite{Stix2002})% \\EQA% Q_{\\theta \\phi} &=& \\Lambda_{\\rm H} \\cos \\theta\\, \\mean\\Omega - \\nut \\sin \\theta \\frac{\\pd \\mean\\Omega}{\\pd \\theta}\\;,\\\\% Q_{r \\phi} &=& \\Lambda_{\\rm V} \\sin \\theta\\, \\mean\\Omega - \\nut (1 - \\epsilon)  r \\sin \\theta \\frac{\\pd \\mean\\Omega}{\\pd r}\\;,% \\EEA% where $\\Lambda_{\\rm H}$ and $\\Lambda_{\\rm V}$ describe the non-diffusive transport and where $\\nut$ is the turbulent viscosity. The factor $1-\\epsilon$ in the latter equation indicates that the turbulent viscosity can be anisotropic. Furthermore, the coefficients $\\Lambda_{\\rm H}$, $\\Lambda_{\\rm V}$, and $\\nut$ can vary as functions of the spatial coordinates.  Much effort has been devoted to determining Reynolds stresses from convection simulations (Hathaway \\& Somerville \\cite{HathaSomer1983}; Pulkkinen et al.\\ \\cite{Pulkkinenea1993}; Rieutord et al.\\ \\cite{Rieutordea1994}; Brummell et al.\\ \\cite{Brummellea1998}; Chan \\cite{Chan2001}; K\\\"apyl\\\"a et al.\\ \\cite{Kaepylaeea2004}; R\\\"udiger et al.\\ \\cite{Ruedigerea2005b}; Hupfer et al.\\ \\cite{Hupferea2005}, \\cite{Hupferea2006}; Giesecke \\cite{Giesecke2007}). These studies have confirmed the existence of the $\\Lambda$-effect and revealed some surprising results that are at odds with theoretical considerations (e.g.\\ Kitchatinov \\& R\\\"udiger \\cite{KitRued1993}, \\cite{KitRued2005}) derived under the second-order correlation approximation (SOCA). The discrepancies are most prominent in the rapid rotation regime, $\\Cost \\approx 10$, where \\EQ% \\Cost = 2\\, \\Omega_0 \\tauto\\;, \\label{equ:Corioliscom}% \\EE% is the Coriolis (or the inverse Rossby) number. Here, $\\Omega_0$ is the angular momentum-averaged rotation rate and $\\tauto$ the convective turnover time. In the solar convection zone, $\\Cost$ varies between 10$^{-3}$ near the surface to ten or more in the deep layers. Convection simulations in the latter regime show that the horizontal angular momentum flux is directed toward the equator, corresponding to positive $Q_{\\theta \\phi}$ in the northern hemisphere, and that it peaks very sharply near the equator (Chan \\cite{Chan2001}; K\\\"apyl\\\"a et al.\\ \\cite{Kaepylaeea2004}; Hupfer et al.\\ \\cite{Hupferea2005}). On the other hand, the vertical stress $Q_{r \\phi}$ can be directed outward (K\\\"apyl\\\"a et al.\\ \\cite{Kaepylaeea2004}), contradicting the theory for vertically dominated turbulence (e.g.\\ Biermann \\cite{Biermann1951}; R\\\"udiger \\cite{Ruediger1980}, \\cite{Ruediger1989}). So far, these results remain without proper explanation.  Often the Reynolds stress realized in the simulation is taken to solely represent the $\\Lambda$-effect. This approach seems like a reasonable starting point but in an inhomogeneous system large scale mean flows are generated when rotation becomes important. These flows affect the Reynolds stresses via the turbulent viscosity. Furthermore, in the presence of stratification, heat fluxes can also significantly affect the stresses (Kleeorin \\& Rogachevskii \\cite{KleeRoga2006}). In the present study we simplify the situation as much as possible in order to disentangle the effect of the turbulent velocity field from other effects. Thus we neglect stratification and heat fluxes by adopting a periodic isothermal setup. Turbulence is driven by external forcing, which provides clear scale separation between the turbulent eddies and the system size, which is typically not achieved in convection simulations. Further insight is sought from comparison of simple analytical closure models with numerical data.  Preliminary results on the $\\Lambda$-effect are presented in K\\\"apyl\\\"a \\& Brandenburg (\\cite{KaBr2007}). In the present paper, numerical datasets covering a significantly larger part of the parameter space are analyzed, and a more thorough study of the validity and results of the minimal tau-approximation are presented.  The remainder of the paper is organized as follows. Section~\\ref{sec:model} summarizes the model and the methods of the study, and in Sects.~\\ref{sec:results} and \\ref{sec:conclusions} the results and the conclusions are given. \\vfill ",
        "conclusions": "\\label{sec:conclusions} Turbulent momentum fluxes, which are described by the Reynolds stresses, were determined from numerical simulations of homogeneous rotating anisotropic turbulence. Since no large-scale shear is present, the generated Reynolds stresses correspond to contributions that are already present for uniform rotation. The resulting term is known as the $\\Lambda$-effect (Krause \\& R\\\"udiger \\cite{KrauseRued1974}). The component responsible for the horizontal transport, $\\qxy$, is positive and peaks around latitude $30\\degr$ regardless of the Coriolis number. The vertical component is predominantly negative and it always peaks at the equator.  Although the numerical results for the $\\Lambda$-effect broadly agree with analytical SOCA calculations (Kitchatinov \\& R\\\"udiger \\cite{KitRued1993}, \\cite{KitRued2005}), the MTA-model seems to reproduce certain features of the numerical results somewhat more closely. The present numerical results do not show the enigmatic results, such as the extreme latitude distribution of $\\qxy$ or a positive $\\qyz$ for rapid rotation, which have been reported from convection simulations (e.g.\\ Chan \\cite{Chan2001}; K\\\"apyl\\\"a et al.\\ \\cite{Kaepylaeea2004}). The difference lies most likely in our neglecting stratification and heat fluxes. The exact manner in which they affect the Reynolds stresses is not within the scope of the present paper, but should be investigated more closely in the future.  By applying the minimal tau approximation closure relation to the Reynolds stress equation, qualitatively similar results are obtained, but the magnitude of the stresses is in general too large. The vertical flux in the MTA-model, however, has a maximum at mid-latitudes for intermediate and rapid rotation. Adding an empirical rotational isotropization term (motivated in Sect.~\\ref{sec:diaresults}) also brings the magnitude in line with the 3D simulations. Although adding this term with this particular form has no rigorous theoretical basis, we can see that phenomenological effects of isotropization of turbulence due to rotation are indeed seen in the simulations and that the term is thus justified.  Another drawback of the MTA-model is that the diagonal components of the Reynolds tensor are rather badly reproduced since the nonlinear effects of rotation manifest in the numerical simulations are not explicitly taken into account. The empirically added rotational isotropization term augments the magnitudes, but not the latitude distribution. Furthermore, no direct evidence of the validity of the MTA-assumption $\\qij = -T_{ij}/\\tau$ was found in the numerical simulations. Contrasting the behavior of the diagonal components to the fairly good correspondence between the numerical simulations and the MTA-model for the off-diagonal components leads us to conclude that, where the behavior of the diagonal components is dominated by the inadequately modeled nonlinear effects, the off-diagonals are fairly well presented by the linear terms.  A Strouhal number of order unity in the MTA-model gives best fits to the numerical results. Fitting the numerical results to expressions derived under the MTA, similar values of $\\St$ are found for slow rotation. For Coriolis numbers of order unity or larger, however, the Strouhal number obtained in this manner decreases rapidly. In the passive scalar case, the situation is somewhat more complex, although a similar decreasing trend of the Strouhal number is recovered for rapid rotation, see Fig.~\\ref{fig:pstrouhal_pscalar}. These results are in accordance with earlier results from convection simulations (K\\\"apyl\\\"a et al.\\ \\cite{Kaepylaeea2005}, \\cite{Kaepylaeea2006a}) using Reynolds stresses or correlation analysis of the velocity field.  A related aspect in turbulent transport that requires closer study is the turbulent viscosity (see preliminary results in K\\\"apyl\\\"a \\& Brandenburg \\cite{KaBr2007}) and the possibility of a $\\Lambda$-effect due to the anisotropy induced by a large-scale shear flow (Leprovost \\& Kim \\cite{LeproKim2007}). These matters will be considered in more detail in a future publication."
    },
    "0806/0806.3084_arXiv.txt": {
        "abstract": "{ Molecular hydrogen is the main constituent of circumstellar disks and could be an important tracer for the evolution and structure of such disks. So far, H$_2$ has only been detected in a few disks and only through spectroscopic observations, resulting in a limited knowledge of the spatial distribution of the H$_2$ emitting gas. } {We report the detection of quiescent H$_2$ emission in a spatially resolved ring-like structure within 100 AU of T Tau N. We present evidence to show that the emission most likely arises from shocks in the atmosphere of a nearly face-on disk around T Tau N.} {Using high spatial resolution 3D spectroscopic K-band data, we trace the spatial distribution of several H$_2$ NIR rovibrational lines in the vicinity of T Tau N. We examine the structure of the circumstellar material around the star through SED modeling. Then, we use models of shocks and UV+X-ray irradiation to reproduce the H$_2$ line flux and line ratios in order to test how the H$_2$ is excited.} {We detect weak H$_2$ emission from the v=1-0 S(0), S(1), Q(1) lines and the v=2-1 S(1) line in a ring-like structure around T Tau N between 0\\farcs1 ($\\sim15$AU) and 0\\farcs7 ($\\sim$100AU) from the star. The v=1-0 S(0) and v=2-1 S(1) lines are detected only in the outer parts of the ring structure. Closer to the star, the strong continuum limits our sensitivity to these lines. The total flux of the v=1-0 S(1) line is $1.8\\times10^{-14}$ergs s$^{-1}$cm$^{-2}$, similar to previous measurements of H$_2$ in circumstellar disks. The velocity of the H$_2$ emitting gas around T Tau N is consistent with the rest velocity of the star, and the H$_2$ does not seem to be part of a collimated outflow. Both shocks impinging on the surface of a disk and irradiation of a disk by UV-photons and X-rays from the central star are plausible candidates for the H$_2$ excitation mechanism. However, irradiation should not create a large degree of excitation at radii larger than 20 AU. Most likely the H$_2$ emission arises in the atmosphere of a flared disk with radius 85-100 AU and mass 0.005-0.5M$_{\\odot}$, where the gas is excited by shocks created when a wide-angle wind impinges on the disk. The H$_2$ emission could also originate from shock excitation in the cavity walls of an envelope, but this requires an unusually high velocity of the wide-angle wind from T Tau N. } {} ",
        "introduction": "The study of circumstellar disks around young stars is essential to understanding their evolution from gaseous disks to planetary systems. In this paper, we examine the spatial distribution of molecular hydrogen, the main constituent of disks around young stars. Disks have been observed in a wide range of wavelengths ranging from optical to millimeter, although only a few studies have concentrated on the H$_2$ component. Many investigations have focused on the broad band spectral energy distribution, which reflects the disk geometry and the structure of the dust content. Disks have also been observed more directly via optically thick dust lanes blocking the scattered light from young stars, and as near-infrared images of the scattered light of the disk itself \\citep[e.g.][]{mccabe2002,weinberger2002}. Molecular line emission from species such as CO or HCO$^+$ is also used as a tracer for disks. The use of such tracers is, however, subject to some uncertainty. Heavy element molecules may freeze out on dust grains, which likely settle to the midplane of the disk and/or get bound in larger rocks or planetesimals. Thus, molecules such as CO can become undetectable even if the disk still exists.  \\subsection{Molecular Hydrogen in Disks} Examining the H$_2$ content in disks has many advantages. Hydrogen and helium are the last parts of the gas to be bound up when planets form, and will therefore remain in the disk after CO and dust have become undetectable. Observations of molecular hydrogen directly trace the gas mass of the disk without making assumptions about the dust-to-gas or CO-to-H$_2$ ratios. Furthermore, molecular hydrogen will remain in the surface layers of the disk when the dust settles to the midplane and is more directly accessible to incoming light than the dust and heavier elements. As a result, H$_2$ may prove to be a better tracer for exploring the evolution and structure of circumstellar disks, since it may be observable for a longer period of time.    Direct observations of H$_2$ in disks have been undertaken by several groups. We focus here on the observations of the IR rovibrational lines, although some studies have concentrated on pure rotational lines in the MIR \\citep[e.g.][]{lahuis2007} as well as fluorescent H$_2$ in the UV \\citep[e.g.][]{walter2003,herczeg2006}. Emission from the H$_2$ v=1-0 S(1) line at 2.1218$\\mu$m has been detected in the disks of several T Tauri stars, classical as well as weak-line \\citep{bary2003, bary2008,itoh2003,weintraub2005,ramsay2007,carmona2007}. These detections are made through longslit spectroscopic observations, and they do not reveal much about the spatial distribution of the molecular hydrogen, beyond indicating that the emitting gas is located within 50 AU of the central star. \\cite{chen1998} presented images of H$_2$ v=1-0 S(1) emission from photoevaporating disks in Orion and showed that the emission arises on the disk surface. In this case, the disks were externally irradiated, and the H$_2$ emission was found from a region $\\sim 200$AU in size.  In this paper, we present spatially resolved images of H$_2$ emission from a ring around T Tau N obtained with the integral field spectrograph SINFONI on the ESO-VLT. The presence of H$_2$ emission in the T Tau system has been known for decades, but this is the first time that the weak emission within 100 AU of T Tau N has been resolved and analyzed.  \\subsection{T Tau} T Tau is a triple star system with an age of $\\sim 1$Myr \\citep{white2001}. The binary component T Tau S, consisting of T Tau Sa and T Tau Sb (separation $\\sim$0\\farcs1), is currently $\\sim$0\\farcs7 south of T Tau N. All three stars are actively accreting and believed to host disks \\citep{duchene2005}. T Tau S shows heavy extinction (A$_V=15$), which is attributed to a circumbinary structure \\citep{duchene2005}. Another possibility is that T Tau S is obscured by the disk around T Tau N \\citep{hogerheijde1997,beck2001}.  T Tau N is a $\\sim$2M$_{\\odot}$ star \\citep{white2001} and is believed to have a disk that is seen nearly face-on \\citep{akeson1998}. Based on photometric periodicity and assumed stellar radius, \\cite{herbst1997} derive an inclination of 19$^{\\circ}$. \\cite{stapelfeldt1998} suggest an outflow and disk with the axis at position angle 300$^{\\circ}$ and with inclination of $\\sim 45 ^{\\circ}$ in order to explain the morphology of scattered optical light. \\cite{akeson2002} find the inclination to be $20-40^{\\circ}$ from SED fitting.  This paper is organized as follows. In Sect.~\\ref{obs}, we describe the observations and data reduction. In Sect.~\\ref{results}, we present the spatial distribution of molecular hydrogen around T Tau N and the velocity distribution of the gas. Section~\\ref{geometry} discusses the geometry of the star-disk-envelope system and Sect.~\\ref{excitation} examines the H$_2$ excitation mechanism. In Sect.~\\ref{ttaus}, we consider the possible implications for T Tau S, and finally, we draw conclusions in Sect.~\\ref{conclusion}. ",
        "conclusions": "\\label{conclusion}  We detect emission from the H$_2$ v=1-0 S(1) rovibrational line at 2.12$\\mu$m in a ring-like structure very close to T Tau N. We find that the weak H$_2$ emission is most likely linked to a nearly face-on flared disk. Another possible solution is that the H$_2$ emission originates from shocks impacting on the lower density walls of an envelope cavity. This scenario, however, requires that the high velocity jet in T Tau N be less collimated than in other T Tauri stars. The radius of the disk is $\\sim 85-100$AU, based on SED modeling and the extent of the H$_2$ emission. The velocity in the vicinity of T Tau N is consistent with the rest velocity of the star to within the errors. Both shocks associated with a wide-angle wind impinging on the disk and UV + X-ray irradiation from the central star onto the disk are plausible excitation mechanisms which can reproduce the H$_2$ flux. Both these mechanisms require a substantial disk around T Tau N. However, models and observations indicate that irradiation from the central star cannot excite H$_2$ at radii much larger than 20 AU. Thus, the most likely excitation mechanism of H$_2$ is that of a wide-angle wind impinging on a flared disk. A PDR created by irradiation may exist within $\\sim$ 6AU from T Tau N."
    },
    "0806/0806.0491_arXiv.txt": {
        "abstract": "{The recent downward revision of the solar photospheric abundances now leads to severe inconsistencies between the theoretical predictions for the internal structure of the Sun and the results of helioseismology. There have been claims that the solar neon abundance may be underestimated and that an increase in this poorly-known quantity could alleviate (or even completely solve) this problem. Early-type stars in the solar neighbourhood are well-suited to testing this hypothesis because they are the only stellar objects whose absolute neon abundance can be derived from the direct analysis of photospheric lines. Here we present a fully homogeneous NLTE abundance study of the optical \\ion{Ne}{i} and \\ion{Ne}{ii} lines in a sample of 18 nearby, early B-type stars, which suggests $\\log \\epsilon$(Ne)=7.97$\\pm$0.07 dex (on the scale in which $\\log \\epsilon$[H]=12) for the present-day neon abundance of the local interstellar medium (ISM). Chemical evolution models of the Galaxy only predict a very small enrichment of the nearby interstellar gas in neon over the past 4.6 Gyr, implying that our estimate should be representative of the Sun at birth. Although higher by about 35\\% than the new recommended solar abundance, such a value appears insufficient by itself to restore the past agreement between the solar models and the helioseismological constraints.} ",
        "introduction": "\\label{sect_intro} State-of-the-art spectral analyses of solar photospheric lines using time-dependent, 3-D hydrodynamical models (Asplund \\etal \\cite{asplund}, and references therein; hereafter AGS05) have recently led to a reduction of the commonly accepted abundances of the dominant metals in the Sun (Grevesse \\& Sauval \\cite{grevesse_sauval}; hereafter GS98). This, in turn, greatly affects the input physics of the standard solar models (e.g. radiative opacities) and considerably worsens the agreement between the theoretical predictions and the results of helioseismic inversions. In particular, the sound speed and density profiles in the solar interior are no longer well reproduced, while the convective zone is predicted to be too shallow and with a helium abundance that is too low (see, e.g. Basu \\& Antia \\cite{basu_antia} for a comprehensive review and an account of the various solutions proposed to solve this problem).  Neon is one of the most important contributors to the opacity at the base of the convective zone, after oxygen and iron. Contrary to these latter two elements whose abundance can be estimated from the analysis of photospheric lines (or is even accurately known from meteoritic data in the case of Fe), the Ne abundance is not well constrained. As noble gases are not retained in CI chondrite meteorites and \\ion{Ne}{i} lines are completely lacking in the solar spectrum because of their high excitation energies, one has to rely instead on indirect estimates based on observations of coronal lines or high-energy particles, which are by themselves prone to large uncertainties (see below). The Ne abundance of the Sun is usually based on measurements of the [Ne/O] abundance ratio in the solar upper atmosphere and has been scaled down (by 0.17 dex) to account for the decrease in the solar oxygen content (an additional correction amounting to --0.07 dex arises from the adoption of a different neon-to-oxygen abundance ratio, as determined from energetic particles; Reames \\cite{reames}). This leads to a value lowered from 8.08 (GS98) to 7.84 dex (AGS05). An upward revision of this uncertain quantity has therefore been invoked as a possible way to compensate for the decrease in opacity brought about by the lower abundances of the other chemical elements. Standard, full solar models constructed with different chemical mixtures suggest that an increase of the Ne abundance by 0.4--0.5 dex, along with a possible adjustment of the other metal abundances within their uncertainties, is required (Bahcall \\etal \\cite{bahcall}). Models with values outside this range do not simultaneously reproduce the helioseismic constraints that are the He abundance/depth of the convective zone and the sound-speed/density profiles in the interior (see also Delahaye \\& Pinsonneault \\cite{delahaye_pinsonneault}). It was also shown that an increase of this magnitude provides a better match to the properties of the solar core, as probed by low-degree {\\it p}-modes (Basu \\etal \\cite{basu}; Zaatri \\etal \\cite{zaatri}). Observations of a large sample of active stars with the {\\em Chandra} X-ray observatory indeed seemed at first sight to support an upward revision of the newly adopted solar Ne abundance at the required levels (Drake \\& Testa \\cite{drake_testa}), but much lower values were subsequently inferred for solar-like stars (e.g. Liefke \\& Schmitt \\cite{liefke_schmitt}) or quiescent solar regions (e.g. Young \\cite{young}), thus leaving this question still open.  The abundance patterns observed in stellar coronae appear at odds with the solar mixture, in particular, with neon being strongly enhanced in active stars by some still poorly-understood mechanisms (e.g. Drake \\etal \\cite{drake01}). An attractive alternative to constrain the neon content of the Sun is, however, offered by the direct analysis of \\ion{Ne}{i} and \\ion{Ne}{ii} photospheric lines in nearby OB stars. Although neon has traditionally been largely neglected in past abundance studies of massive stars, this 'solar model crisis' has indeed renewed interest in determining the abundance of this element in nearby objects. Since the publication of AGS05 results, however, only a single study addressing this issue appeared in the literature (Cunha \\etal \\cite{cunha}). A relatively high mean Ne abundance was found in a sample of 11 B-type dwarfs in the Orion association (0.27 dex above AGS05 value), but this still falls short (by a factor $\\sim$1.5) of completely solving the controversy discussed above. To shed more light on this issue, here we present a fully homogeneous NLTE abundance analysis of a sample of 18 early B-type stars in the solar neighbourhood. The results presented in this paper supersede previous preliminary reports (Morel \\& Butler \\cite{morel_butler07}, \\cite{morel_butler08}), as significant improvements in the model atom have been made during this interval. ",
        "conclusions": "A mean, absolute neon abundance, $\\log \\epsilon$(Ne)=7.97$\\pm$0.07 dex, has been inferred from our combined NLTE abundance analysis of the photospheric \\ion{Ne}{i} and \\ion{Ne}{ii} lines in a sample of 18 nearby, early B-type stars. This indicates a value for the Sun $\\sim$35\\%  higher than the new recommended solar abundance (7.84$\\pm$0.06 dex; AGS05). In contrast, an increase of the Ne abundance by a factor $\\sim$3 in the solar interior is needed to restore the agreement between the solar models and the helioseismological data. Our results therefore clearly suggest that solving this problem by simply adjusting the solar Ne abundance would probably require an increase of this quantity well beyond the range of plausible values (note that an enhancement of the other metal abundances to within their uncertainties may also be necessary; Bahcall \\etal 2005). We point out that this is a robust conclusion attained regardless of the ion or $T_{\\rm eff}$ scale chosen (Fig.\\ref{fig_teff}).  Neon is produced during carbon burning in the final stages of the evolution of massive stars and one may naturally expect the Ne abundances of young, B-type stars to be higher than the solar value because of chemical enrichment over the past 4.6 Gyr. However, the predicted enhancements in the solar neighbourhood are likely to be very small according to Galactic chemical evolutionary models ($\\sim$0.04 dex; Chiappini \\etal \\cite{chiappini}). Our mean Ne abundance should therefore be very close to the value prevailing in the protosolar nebula. Unexpectedly, however, metal abundances derived for nearby B stars are often found to be slightly (but significantly) below the most recent estimates for the Sun or the meteoritic values (e.g. Morel \\cite{morel08}, and references therein). This discrepancy may be related to missing physics or unaccounted systematic errors in the B star analyses, and may question the assumption that the neon abundance derived from hot stars is directly transposable to the Sun. The fact that our mean Ne abundance is indistinguishable from the values recently determined for the ionized gas in the Orion nebula (8.05$\\pm$0.07 dex; Esteban \\etal \\cite{esteban}), within supergranules (7.89$\\pm$0.13 dex; Young \\cite{young}) or from {\\it in situ} observations of the solar wind (7.96$\\pm$0.13 dex; Bochsler \\cite{bochsler}) nevertheless suggests that the abundances we derived are representative of the solar value (see Fig.\\ref{fig_ne_literature})."
    },
    "0806/0806.4617_arXiv.txt": {
        "abstract": "We present hydrodynamical models for the Cassiopeia A (Cas A) supernova remnant and its observed jet / counter-jet system. We include the evolution of the progenitor's circumstellar medium, which is shaped by a slow red supergiant wind that is followed by a fast Wolf-Rayet (WR) wind.  The main parameters of the simulations are the duration of the WR phase and the jet energy. We find that the jet is destroyed if the WR phase is sufficiently long and a massive circumstellar shell has formed.  We therefore conclude that the WR phase must have been short (a few thousand yr), if present at all. Since the actual jet length of Cas A is not known we derive a lower limit for the jet energy, which is $\\sim 10^{48}$~erg. We discuss the implications for the progenitor of Cas A and the nature of its explosion. ",
        "introduction": "Over the last decade evidence has emerged that suggests that at least some core-collapse supernovae are intrinsically non-spherically symmetric explosions. The evidence is strongest for supernovae of stars that have lost most of their outer (hydrogen-rich) envelopes, i.e., the type Ib/c supernovae \\citep{2001Wangetal}. For those explosions, the inner layers are exposed early on, and asymmetries in the core more easily survive the interactions with the outer layers. This implies that departures from spherical symmetry originate from deep inside the explosion.  A better understanding of the explosion geometries is needed to provide further insights into what powers core-collapse  supernovae. In the canonical explosion model the explosion is driven by deposition of neutrino energy into the region just outside the proto-neutron star. However, up to now, computer simulations of this core collapse do not self-consistently predict supernova explosions  \\citep{2007Jankaetal}. In those simulations the role of magnetic fields and stellar rotation is usually neglected. According to \\citet{1999Khokhlovetal} and \\citet{2000Wheeleretal}, magnetic fields and rotation may play a crucial role in the explosion mechanism, and may lead to bipolar explosions. Additionally, other explosion mechanisms that are based on acoustic and hydrodynamic instabilities can result in asymmetric, albeit not necessarily bipolar explosions \\citep[e.g.][]{2007BlondinShaw,2007Burrowsetal}. However, if one considers the most energetic supernova explosions; those associated with the long duration gamma-ray bursts (LGRBs), it is very likely that these explosions are truly bipolar. The associated supernovae are of type Ic, see \\citet{2006DellaValle} for a review.  The engines that drive the explosions associated with LGRBs may or may not be related to those of ``normal'' core-collapse supernovae. In the collapsar model, LGRBs are the result of black hole formation \\citep{1999MacFadyenWoosley} and thus have a distinctly different engine from that producing normal supernovae. Alternative models that consider LGRBs to be powered by highly magnetic, rapidly rotating, neutron stars \\citep[e.g.][]{2004Thompsonetal} or by trans-relativistic blast waves in supernovae \\citep{2001Tanetal}, allow for a continuum of bipolarity and explosion energies. In those cases, the amount of rotation and magnetic field strength, and the line of sight, determine whether we observe a ``normal'' supernova or one associated with a LGRB. This study aims to shed some light on the intermediate case of a supernova that shows distinct bipolarity, but is energetically in the range of regular supernovae and does not have relativistic ejecta. We try to provide some insight in the requirements on the energy in the asymmetric part of the supernova, i.e. it's ``jets'', and on the type of progenitor that was responsible for the circumstellar medium (CSM) at time of explosion.  \\begin{figure}[!htbp] \\centering \\plotone{f1.eps}  \\caption[ ] {Three color image showing the location of the jets \\citep[red, see][]{2004Vink,2004Hwangetal} with respect to the bright X-ray shell of ejecta (green, X-ray Si~\\textsc{XIII} emission) and radio synchrotron emission (VLA archival data). The jet image is obtained by taking the ratio of Si~{\\sc XIII} over Mg~{\\sc XI} X-ray line emission. (Public domain image based on the 1~Ms Chandra observation of  Cas A \\citep{2004Hwangetal} [http://www.astro.uu.nl/$\\sim$vinkj/casa\\_jet\\_si\\_radio.jpg]). \\label{fig:casa_SiMg}} \\end{figure}  Two likely examples of bipolar supernovae are known in the local neighborhood: SN1987A \\citep{2002Wangetal} and the supernova remnant (SNR) Cassiopeia A, the subject of this paper. The bipolarity of Cas A has only recently been established from optical \\citep{2001Fesen,2006Fesenetal}, X-ray \\citep{2004Vink, 2004Hwangetal,2006Lamingetal} and infrared \\citep{2004Hinesetal} observations. These observations show that, apart from the long known ``jet'' region in the northeast, a somewhat less prominent protrusion is located in the southwest (Fig.~\\ref{fig:casa_SiMg}). In Figure 1 we show in red the image of the jet as shown in \\citet{2004Hwangetal}. In order to show the jet in the context of the overall emission it is combined with images in silicon (green) and radio (blue). For details and a discussion on the jet and its abundances we refer to \\citet{2004Vink}, \\citet{2004Hwangetal}, and \\citet{2006Fesenetal}.  The jets extend out to a radius of at least 3.8~pc, for a distance of 3.4~kpc \\citep{1995Reedetal}. The reason to believe that these jets are the results of a bipolar explosion, rather than being caused by a bipolar structure in the CSM \\citep{1996Blondinetal},  is the distinct elemental abundance patterns in both jet regions, with the jet material coming from deeper layers inside the star. The X-ray and optical data indicate that the jet material is rich in oxygen burning products (Si, S, Ar, Ca), while it lacks carbon- and neon-burning products (O, Ne, Mg). This is the reason why the jet, of which the emissivity is weak compared to the rest of Cas A, stands out by taking the ratio of the Si/Mg line emission.  Interestingly, the supernova that caused Cas A seems to have some shared characteristics with the supernovae associated with LGRBs and X-ray flashes. Cas A's progenitor probably had lost most of its hydrogen envelope, given the lack of hydrogen rich, optically identified, ejecta. Although we are not claiming that Cas A was an LGRB or even an X-ray flash, the two non-relativistic jets suggest that it may be related and there may be a continuum of bipolarity in the explosion of supernovae, thus providing a possible link between LGRBs and normal supernovae. The  total mass of shocked ejecta is 2-4~M$_\\odot$  \\citep{1996Vinketal}, and the explosion energy is about a factor of two more than the canonical explosion energy of $10^{51}$~erg  \\citep{2003HwangLaming}. The total oxygen ejecta mass of 1-2~M$_\\odot$ suggests a main sequence mass of $18-22$~M$_\\odot$ \\citep{2004Vink}. These properties are reminiscent of the parameters derived for SN2006aj, the supernova associated with the X-ray flash XRF 060218  \\citep{2006Mazzalietal}, and similar to SN2003jd, the one suggested to relate to a LGRB\\citep{2007Valentietal}.  The large amount of swept up mass in Cas A and the dynamic properties of the blast wave suggest that the blast wave is currently moving through the high-density red supergiant (RSG) wind \\citep{2003ChevalierOishi, 2004Vink}. However, the lack of H-rich ejecta suggests that Cas A exploded as a Wolf-Rayet (WR) star. Moreover, the presence of slow moving N-rich knots has been explained as originating from the hydrodynamical instabilities between the fast WR wind and the dense, slow moving, RSG wind \\citep{1996GGSetal_II}.  In this paper we present hydrodynamical simulations of the jets in the context of the progenitor's mass loss history, which we take to be a RSG phase, possibly followed by a WR phase. There are two main reasons for pursuing this problem. First of all, the energetics of the jets can be better estimated using a realistic mass loss history in the hydro-simulations. Secondly, the survival of the jets depends strongly on the mass-loss history of the progenitor. Therefore, the jets in Cas A can be used as a diagnostic on both the properties of the bipolar explosion, and on the progenitor-shaped circumstellar medium (CSM) at the time of explosion. ",
        "conclusions": "We have simulated the evolution of axisymmetric ejecta, such as may result from a bipolar supernova explosion, in the context of Cas A and the presence of a jet / counter-jet in this SNR.  For the initial conditions we used a realistic progenitor evolution for a $\\sim 20$~M$_\\odot$ star, consisting of a RSG wind, followed by a WR wind. We find that the presence of a WR shell limits the survival of the jets. The survival depends critically on the energy available in the jet region and on the mass contained in the WR shell. The latter is determined by the duration of the WR phase and the properties of the progenitor winds. For the parameters chosen for the Cas A progenitor, we find that if the WR phase is longer than $2000-5000$~yr, the jets do not protrude through the shell, in which case the situation does not correspond to the presence of jets in Cas A. Therefore, either the progenitor went through a very short WR phase, or it did not have one and exploded as a RSG. In general however, this means that, also if a SNR appears symmetric, the explosion may still have been accompanied by jets.  In order to match the length of the observed jets of Cas A, an energy of at least $2.0 \\times 10^{48}$~erg per jet is required. In case a WR shell is formed, higher energies are needed. However, for a WR duration in excess of maximally 5000~yr, the properties of the remnant and the jet do not match the observations. The upper limit we find corresponds to the findings of van Veelen \\& Langer (2008, in preparation), where they find that the properties of the forward and reverse shock do not agree with observations for a WR phase that lasts more than 5000~years.   The question now arises if a scenario involving a very short WR phase is realistic. If the progenitor was not in a binary, the chance of having a very short WR phase is very low. The main reason for invoking a WR phase at all is that it explains the clumpiness in the remnant, the lack of hydrogen in the ejecta, the presence of metal clumps far out in the ejecta, and the high N/H ratio in the CSM. A single star at the end of the RSG phase with a clumpy mass loss history \\citep{2003ChevalierOishi} may also be able to partly explain the above, but requires equally coincidental circumstances. \\citet{2007LamingHwang} favor a short WR phase because, based on the temperature and ionization age of the X-ray emitting gas, they find that the ejecta expanded in a bubble of around $\\sim 0.2$~pc., in our case corresponding to a WR phase of about 1000 yrs. They however do not take into account the presence of a shell around the bubble.  As an alternative to the single star model, a model with a binary companion has been proposed \\citep{2006Youngetal}. A common envelope (CE) phase in a close binary solves a number of problems: It explains the low ejecta mass in conjunction with a MS mass of $\\sim 20$~M$_\\odot$, and provides a natural explanation for a very short  WR phase of the primary star (Podsiadlowski, private communication). However, the details of CE evolution are not well understood and no companion star has currently been found.  The simulated jets resemble the observed jet of Cas A in opening angle for a variety of parameters. This reinforces the idea that the explosion itself was intrinsically bipolar. It remains an interesting question what mechanism is responsible for such an asymmetric explosion, and whether it is related to other bipolar explosion phenomena such as X-ray flashes or LGRBs. Although our case does not resemble the relativistic scenarios as are invoked in models for gamma ray bursts \\citep{2005Piran}, it does not seem unreasonable that {once again rotation is involved in creating the asymmetry in a low-energy explosion like this}. Unless the rotation was created during the explosion \\citep{2007BlondinMezzacappa}, rotation may have also left an imprint on the CSM and a combination of asymmetric CSM and an asymmetric explosion could have been responsible for the Cas A morphology. The observed point source near the center of the remnant is likely to be a neutron star. However, the absence of a bright pulsar wind nebula suggests that the present rotation period of the neutron star is relatively low \\citep[$> 160-330$~ms][]{1988SewardWang,2007Vink}. This seems at odds with a rapid rotation of the stellar core as a mechanism to create a bipolar explosion. However, this discrepancy may be solved if the point source is a magnetar that has considerably slowed down.  In conclusion, we would like to emphasize that the presence of jets in Cas A, together with considerable knowledge about nucleosynthesis yields, explosion energy, and compact object, makes this SNR a unique object to investigate the mechanism behind bipolar explosions, and the type of progenitors that produce bipolar explosions."
    },
    "0806/0806.1518_arXiv.txt": {
        "abstract": "This template, along with associated style files, can be used to approximate typeset \\textit{Optics Letters} (OL) pages for purposes of length check. With a few command changes, the two-column version can be disassembled into a single-column double-spaced version suitable for production and submission to OSA. Examples are given of how to account for some of the factors that affect the accuracy of the length estimate: figures, tables, equations, and author affiliations.  \\textbf{Authors should note the new affiliation style (shown above) as well as the change to bracketed reference callouts.} ",
        "introduction": " ",
        "conclusions": ""
    },
    "0806/0806.3347_arXiv.txt": {
        "abstract": "{}{Observations of Kepler's supernova remnant (G4.5+6.8) with the H.E.S.S. telescope array in 2004 and 2005 with a total live time of 13 h are presented.}{Stereoscopic imaging of Cherenkov radiation from extensive air showers is used to reconstruct the energy and direction of the incident gamma rays.} {No evidence for a very high energy (VHE: $>$100 GeV) gamma-ray signal from the direction of the remnant is found. An upper limit (99\\% confidence level) on the energy flux in the range $230 \\, \\mbox{GeV} - 12.8 \\, \\mbox{TeV}$ of $8.6 \\times 10^{-13} \\, \\mbox{erg} \\, \\mbox{cm}^{-2} \\, \\mbox{s}^{-1}$ is obtained.} {In the context of an existing theoretical model for the remnant, the lack of a detectable gamma-ray flux implies a distance of at least $6.4 \\, \\mbox{kpc}$. A corresponding upper limit for the density of the ambient matter of $0.7 \\, \\mbox{cm}^{-3}$ is derived. With this distance limit, and assuming a spectral index $\\Gamma = 2$, the total energy in accelerated protons is limited to $E_{p} < 8.6 \\times 10^{49} \\, \\mbox{erg}$. In the synchrotron/inverse Compton framework, extrapolating the power law measured by RXTE between $10$ and $20 \\, \\mbox{keV}$ down in energy, the predicted gamma-ray flux from inverse Compton scattering is below the measured upper limit for magnetic field values greater than $52 \\, \\mu \\mbox{G}$.}  \\offprints{dominik.hauser@mpi-hd.mpg.de} ",
        "introduction": "\\label{intro} It is widely believed that the bulk of the Galactic cosmic rays (CR) with energies up to at least several $100 \\, \\mbox{TeV}$ originates from supernova explosions (see for example \\cite{1994A&A...287..959D}). This implies copious amounts of very high energy (VHE: $>$100 GeV) nuclei and electrons in the shells of supernova remnants (SNRs).  These particles can produce VHE gamma rays in interactions of nucleonic cosmic rays with ambient matter, via inverse Compton (IC) scattering of VHE electrons off ambient photons, as well as from electron Bremsstrahlung on ambient matter. Therefore SNRs are promising targets for observations of VHE gamma rays. \\par In October 1604 several astronomers, among them Johannes Kepler, observed a ``new star'' which today is believed to have been a bright supernova (SN) at the Galactic coordinates $l = 4.5 ^{\\circ}$ and $b = 6.8 ^{\\circ}$.  The remnant of this supernova has since been a target of observations covering the entire electromagnetic spectrum.  In the radio regime, \\cite{1988ApJ...330..254D} determined a mean angular size of $\\sim 200''$ and a mean expansion law $R \\propto t^{0.50}$, where $R$ is the radius and $t$ is the time. However, the expansion parameter $x=\\dot{R} t/R$ varies considerably around the SNR shell, $0.35 < x < 0.65$, possibly indicating spatial inhomogenities in the circumstellar gas density. In a very recent paper by \\cite{2008arXiv0803.4011V} these properties, and the general asymmetry of the remnant, have been basically confirmed through X-ray measurements. They also allowed the analysis of a high-velocity synchrotron filament in the eastern part of the remnant with $x=0.7$. \\par In addition, the distance $d$ to the SNR is still under debate.  \\cite{1999AJ....118..926R} report on an HI absorption feature in VLA data and use the Galactic rotation model of \\cite{1989ApJ...342..272F} to calculate a lower limit $d > (4.8 \\pm 1.4) \\, \\mbox{kpc}$.  They also give an upper limit on the distance due to the lack of absorption by an HI cloud at $6.4 \\, \\mbox{kpc}$. The authors remark that these values involve uncertainties because of the proximity of Kepler's SNR to the Galactic center. In contrast, \\cite{2005AdSpR..35.1027S} and subsequently \\cite{2007ApJ...662..998B} have given a lower source distance of $d = 3.9 (+1.9 -0.9)$~kpc, from an absolute shock velocity $\\sim 1660 \\pm 120 \\, \\mbox{km} \\, \\mbox{s}^{-1}$ derived from the H$\\alpha$ emission line width of a Balmer-dominated filament that is located in the northwestern region. The line broadening, taken as an indication of the downstream thermal gas temperature, was used to determine the shock velocity. We shall return to this question in the discussion section. \\par Finally, the type of the supernova is not undisputed. From the reconstructed light curve \\cite{1943ApJ....97..119B} claimed that it was a type Ia SN, but \\cite{1985AJ.....90.2303D} argued that the light curve is also consistent with a type II-L.  \\cite{1989ApJ...347..925S} and \\cite{1999PASJ...51..239K} observed a relative overabundance of heavy elements that agrees with type Ia nucleosynthesis models, while \\cite{1994A&A...287..206D} saw more evidence that Kepler's SNR is the remnant of a core-collapse SN.  Its position, $500 - 750 \\, \\mbox{pc}$ above the Galactic plane, is more consistent with a type Ia than a type II SN, as a SN of the latter type is expected to be confined to the region of high gas density found in the plane. However, in the case of a core-collapse event this might be explained through the model of a runaway star, as proposed by \\cite{1987ApJ...319..885B}.  More recently, theoretical modeling of the detailed thermal line spectra obtained with \\textit{XMM} (\\cite{2004A&A...414..545C}) led \\cite{2005ApJ...624..198B} to the conclusion that the X-ray spectrum is best fit by a type Ia SN, a view also expressed by \\cite{2005ASPC..342..416B}. Most recently \\cite{2007ApJ...668L.135R} reported on deep \\textit{Chandra} observations and argued from the high abundance of iron and the very low abundance of oxygen that the progenitor of Kepler's SNR has been a type Ia SN. Therefore it appears that the observational evidence is finally converging on a type Ia event. \\par In this paper observations of Kepler's SNR with the H.E.S.S. telescope array are described.  An upper limit on the integrated energy flux above $230 \\, \\mbox{GeV}$ is derived. Combining this H.E.S.S. result with the theoretical predictions of \\cite{2006A&A...452..217} suggests a lower limit on the distance, close to the upper limit given by \\cite{1999AJ....118..926R}, if Kepler's SN is a priori assumed to be of type Ia. ",
        "conclusions": "Observations of Kepler's SNR with H.E.S.S. result in an upper limit for the flux of VHE gamma rays from the SNR. In the context of an existing theoretical model (BKV) for the remnant, and assuming an ejected mass of $1.4 \\, M_{\\odot}$ and an explosion energy of $10 ^{51} \\, \\mbox{erg}$ in agreement with type Ia SN explosion models, the lack of a detectable gamma ray flux implies a distance of at least $6.4 \\, \\mbox{kpc}$, which is the same as the upper limit derived by \\cite{1999AJ....118..926R} from radio observations. Given that the gamma-ray flux effectively scales with $E_{\\mbox{\\tiny SN}} ^{2}$, a significantly higher explosion energy is excluded; a theoretically acceptable lower explosion energy of $0.8 \\times 10 ^{51} \\, \\mbox{erg}$ would lower the distance limit to $6 \\, \\mbox{kpc}$. \\par Assuming a purely hadronic scenario, a standard type Ia SN explosion, and using $6.4 \\, \\mbox{kpc}$ as a lower limit for the distance, the H.E.S.S. upper limit implies that the total energy in accelerated protons is less than $8.6 \\times 10 ^{49} \\, \\mbox{erg}$. \\par In a synchrotron/IC scenario no strong constraints on the magnetic field can be obtained."
    },
    "0806/0806.1342_arXiv.txt": {
        "abstract": "We present results of new photoionization calculations for investigating gaseous regions that represent potentially expected stages of nuclear processing in the Crab Nebula supernova remnant.  In addition to gas resulting from CNO-processing and oxygen-burning, as previously reported, a large component of the nebula appears to be carbon-rich. These results suggest that the precursor star had an initial mass $\\gtrsim$~9.5~M$_{\\odot}$. ",
        "introduction": "The Crab Nebula (M1~=~NGC1952) is generally recognized as the remnant of the core-collapse supernova SN1054. Estimates of the precursor star's initial mass have been in the range $8-12$~M$_{\\odot}$ (see Davidson \\& Fesen 1985 and references therein).  The visible remnant contains at least $1-2$~M$_{\\odot}$ of He-rich line-emitting gas (MacAlpine \\& Uomoto 1991)\\footnote{Fesen et al. 1997 used global line photometry to suggest a higher nebular mass, approaching that once postulated by MacAlpine et al. 1989.  However, as discussed in MacAlpine \\& Uomoto 1991, large spatial variations in the He/H ratio cause spuriously large mass estimates based on global line photometry.  More accurate emitting gas estimates should take into account spatially resolved line photometry in order to allow for the effects of helium abundance variations on N(S)/N(H), which is used along with $[$S~II$]$ line intensities to estimate the amount of neutral gas present.}, the neutron star probably has about 1.4~M$_{\\odot}$, and the outer layers of the precursor may have left in a pre-supernova wind (Nomoto et al. 1982) or a shockwave (e.g., Chevalier 1977; Davidson \\& Fesen 1985).  Knowledge of the chemical composition of the remnant is vitally important for a better understanding of the precursor star's initial mass and associated details of the core-collapse event.  According to Nomoto et al. (1982), an $8-9.5$~M$_{\\odot}$ star would likely have undergone an O-Ne-Mg core collapse following electron capture, whereas a star more massive than 9.5~M$_{\\odot}$ would have developed an Fe-rich core.  A distinguishing factor between these two possible scenarios would be the nebular carbon abundance.  As pointed out by Nomoto (1985), for a precursor with initial mass less than about 9.5~M$_{\\odot}$, high N and low C mass fractions (compared with solar) would exist from CNO-cycle processing in the He-rich gas. On the other hand, for precursor mass $>$~9.5~M$_{\\odot}$, helium-burning and nitrogen-processing via $^{14}$N($\\alpha$,$\\gamma$)$^{18}$F($\\beta$$^{-}$$\\nu$)$^{18}$O($\\alpha$,$\\gamma$)$^{22}$Ne ought to have produced high C and low N.  In addition, models for a roughly 10~M$_{\\odot}$ precursor (Woosley et al. 1980; Woosley \\& Weaver 1986) have suggested the existence of off-center neon and oxygen flashes just prior to the core collapse, whereas such events may not be expected in lower mass stars (Hillebrandt et al. 1984; Nomoto 1984).  The necessary chemical composition information should be obtainable.  Because of its young age and location roughly 180~pc from the Galactic plane, the Crab Nebula is relatively uncontaminated by interstellar material. In addition, its line-emitting gas is ionized and heated primarily by locally-generated synchrotron radiation (rather than shock heating), so the ionization, thermal, and chemical characteristics of the gas can be accurately analyzed using numerical photoionization codes.  Although there has been a long history of observational and theoretical investigations aimed at understanding abundances in the Crab Nebula, to some extent the results have been inconclusive or contradictory. As summarized by Henry (1986), Henry and MacAlpine (1982) and Pequignot and Dennefeld (1983) compared measured line-intensity data available at the time with photoionization calculations. In both studies, the deduced C and N mass fractions were roughly solar or below, but there were problems with interpretation of near-infrared $[$C~I$]$ emission lines (see below). Davidson et al. (1982) and Blair et al. (1992) also obtained satellite observations of the ultraviolet C~IV and $[$C~III$]$ lines for certain locations, and they concluded no definitive evidence for a carbon excess, although the latter study left open the possibility.  More recently, using extensive long-slit spectroscopy covering much of the nebula, MacAlpine et al. (1996) deduced high nitrogen mass fractions in some filamants and also high abundances for products of oxygen-burning, such as sulfur and argon, at other locations.  All of the above results, taken together, have generally been interpreted as suggestive of a precursor in the $8-9.5$~M$_{\\odot}$ range, although high sulfur and argon abundances would appear to argue for a more massive star.  In this paper, we report on the results of new photoionization computations (summarized by MacAlpine et al. 2007b), which were developed using new line-intensity measurements which extend to 1~\\micron\\ (MacAlpine et al. 2007a, hereinafter Paper~1). These calculations suggest that strong $[$N~II$]$~$\\lambda$$\\lambda$6548,6583 emission lines measured at many locations by MacAlpine et al. (1996, 2007a) do {\\it not} necessarily imply high nitrogen abundances, while strong measured $[$C~I$]$~$\\lambda$9850 emission {\\it can} be interpreted as indicative of high carbon mass fractions.  In addition, high mass fractions for S and Ar in some locations are confirmed.  These results argue {\\it consistently} for a precursor star initial mass $\\gtrsim$~9.5~M$_{\\odot}$. ",
        "conclusions": ""
    },
    "0806/0806.1497_arXiv.txt": {
        "abstract": "I present details of the variations of several hundred red giant stars on time scales of a few hours to a few days from {\\em Hubble Space Telescope (HST)} observations of a low-extinction galactic bulge sample from an intensive seven day campaign.  Variations in the red giants are shown to be a strong function of position within the color-magnitude diagram (CMD) in accord with general expectations from theory. Amplitudes are greater for stars with larger radii, whether this results from higher luminosity at the same effective temperature or lower temperature at a fixed apparent magnitude.  Likewise, characteristic time scales for the variations increase to the upper right in a CMD as does the ratio of amplitudes measured at 606 nm compared to 814 nm. Characteristic variation time scales are well matched by low-order radial pulsation modes. The effective sample discussed here extends from about two magnitudes above the bulge turnoff at which red giant radii are $\\sim$7 $R/R_{\\odot}$ at 5,000 $^\\circ$K with typical amplitudes of $\\sim$0.5 mmag to $\\sim$40 $R/R_{\\odot}$ at 4,000 $^\\circ$K with amplitudes of $\\sim$3.5 mmag. Variability characteristics are quite similar at any given position in the CMD, and at levels in the CMD where oscillations are easily detected nearly all red giants show such. If these variations represent oscillations with sufficient lifetimes to derive accurate mode frequencies more extensive observations, e.g.\\ as should soon be provided by the {\\em Kepler Mission}, would provide a rich asteroseismic return. ",
        "introduction": "Red giants are attractive targets for the study of stellar oscillations. From the observational perspective, the amplitude of stochastically-driven p-mode oscillations is expected to rise with increasing $L/M$ \\citep{kje95}. Assuming that coherent oscillations exist with sufficiently long lifetimes to allow studies analogous to those from solar oscillations, asteroseismology of red giants promises probes of great interest for testing the theory of these highly evolved stars with complicated interior structures. For the red giant phase, where evolutionary tracks from different progenitor masses converge in a CMD, asteroseismology can, in principle, provide valuable constraints on the intrinsic parameters of individual stars \\citep{ste08,kal08}.  That red giants show characteristic variations is well known. Less clear (and the subject of increasing study) is what the physical nature of these variations are.  A central question for variations in any group of stars is whether these follow from normal modes of oscillation either driven by nonlinear feedbacks such as in RR Lyrae stars, or from stochastic driving as in the Sun.  The common alternative to oscillations would be variations arising from giant cell convection as suggested by \\citet{sch75}, see also \\citet{dzi01} and \\citet{lud06}, in which case their utility as probes of interior structure likely do not exist as with true oscillations. A secondary question is whether the variations, if oscillations, have lifetimes long enough to allow the accurate determination of frequencies required to support asteroseismic interpretations. This paper, while not resolving these questions, will more firmly establish the overall characteristics of red giant variations by utilizing a serendipitous data set that provides detections on several hundred stars within a narrow confine of the CMD.  After a decade in which the promise was recognized \\citep{bro94}, but results did not materialize despite concerted efforts, asteroseismology has provided solid and exciting results in recent years manifested by precise measurements of solar-like oscillations in a number of stars. See \\citet{bed07} for a recent general review and references. For red giants in particular a number of radial velocity studies requiring dedicated use of at least moderately large telescopes and excellent spectrographs have provided evidence for oscillations in red giants. \\citet{hat94a,hat94b} found 14--54 m s$^{-1}$ variations in $\\alpha$~Boo, following up on the \\citet{smi87} detection,  and $\\beta$ Oph respectively using observations over typically eight consecutive nights on the 2.1~m McDonald telescope. These velocity amplitudes would correspond to photometric variations of 0.5--2.0 mmag using the \\citet{kje95} scaling relation. With observations over 12 nights and precisions better than those successful for the K1~III star $\\alpha$ Boo, \\citet{hor96} failed to detect evidence of oscillations on any of four G8~III to K2~III stars studied. More recently \\citet{fra02} detected solar-like oscillations in the G7~III star $\\xi$ Hya peaking at about 2~m s$^{-1}$ from use of the CORALIE spectrograph and 1.2-m Swiss telescope (ESO). \\citet{der06} report $\\sim$3~m s$^{-1}$ oscillations in the G9.5~III star $\\epsilon$~Oph. The power spectrum for the $\\xi$ Hya radial velocities has maximum peaks corresponding to periods near 0.15 days, and a generally complex structure with evidence for even mode spacing expected of solar-like, coherent oscillations.  The predicted mode lifetimes from \\citet{hou02} are 17 days for this star, while \\citet{ste06} find a value of about two days from analysis of the extensive \\citet{fra02} observations. The latter value, if true and common, could well limit the overall utility of asteroseismology of red giant stars.  Photometric observations have been quite effective for elucidating variations in extreme red giants to supergiants, e.g.\\ \\citet{kis06} find ubiquitous variations in supergiants from multi-decade visual observations compiled by the AAVSO that are suggestive of stochastically driven oscillations.  The OGLE-II data base with extensive data over three years also provides \\citep{kis03} clear detection of multiple variability sequences, presumably related to different radial orders of oscillation for thousands of stars below the tip of the red giant branch from the LMC. For these upper red giant branch variables amplitudes are commonly 1--4\\% with periods of 15--20 days with excitation mechanisms related to either Mira-like pulsators or stochastically-driven pulsations from convection. Farther down the RGB \\citet{edm96} claimed that red giants often showed variations on time scales of $\\sim$2--4 days with amplitudes of 0.5--1.5\\% for K~giants in 47~Tuc based on a 40-hour U-band sequence obtained with {\\em HST}. \\citet{jor97} used extensive Str\\\"{o}mgren photometry of unusually high photometric stability to detect variations in K~giants down to $\\sim$0.5\\%, and, significantly, to establish that a clear level of variability onset exists as a function of stellar parameters above which all red giants are variable. \\citet{hen00} reported on extensive automatic telescope observations of 187 red giants with sensitivity to about 0.2\\% confirming the \\citet{jor97} results and further establishing that most red giants earlier than G2 and later than K2 are variables.  At the 2~mmag precision level available \\citet{hen00} found that most red giants over G3--K1 were not variable. The variables to the red were interpreted as radial pulsations while those in the blue set had characteristic time scales too long for radial pulsations and non-radial g-modes were proposed.  Extensive ground-based campaigns attempting to detect oscillations in red giants over luminosity, radius ranges covered by the current data have largely fallen short of success. \\citet{ste07} observed several giants in M67 and reported indications of excess power broadly consistent with expectations. \\citet{fra07} observed many giants in the globular cluster M4 and reported no consistent evidence of oscillations.  Given the extreme range of stellar properties pertaining to red giants, it is not surprising that multiple physical mechanisms may be responsible for subsets of these. With deep surface convection zones red giants are expected to show stochastically excited solar-like pulsations \\citep{chr83,hou99,sam07}. \\citet{xio07} argue that the hotter type~G giants fall within the classical Cepheid--$\\delta$~Scuti instability strip with $\\kappa$ mechanism driving, and that a second instability strip for the cooler K and M~giants at log($T_\\mathrm{eff}$) $<$ 3.7 is driven by coupling between convection and oscillations. The \\citet{xio07} result is broadly consistent with the \\citet{hen00} observational results in finding that an intermediate temperature domain is pulsationally stable.  The promise of space-based photometry has recently been demonstrated by \\citet{ste08} with the detection of power spectrum excess in 11 red giants observed with the {\\em WIRE} satellite.  The {\\em WIRE} observations of 15--61 days reach noise levels in amplitude spectra to a best of 7~ppm with a median of just over 20~ppm. Even more impressive are results from the {\\em MOST} satellite on $\\epsilon$~Oph for which a 28~day observation set for this red giant with $L/L_{\\odot} \\sim60$, $T_\\mathrm{eff}\\sim4900$ provides clear and unambiguous detection of multiple, equally spaced modes from which strong asteroseismic constraints are quoted by \\citet{bar07} and \\citet{kal08}. Arguments for relatively long mode lifetimes of 10--20 days are encouraging.  The data available for the present analyses were obtained with {\\em HST} ACS as a nearly continuous time series of alternating F606W and F814W exposures for seven days in February 2004. These data were collected for the purpose of extrasolar planet detection for Hot Jupiters orbiting upper main sequence stars in the bulge \\citep{sah06}. The results here will be of particular interest for extending detections to smaller amplitudes than previously possible from ground-based observations, and for a much larger sample than possible with previous space-based results.  Section 2 will provide an introduction to the properties of these {\\em HST} data and discuss the techniques used for time series extractions on saturated stars. The general noise properties of stars on the RGB and main-sequence will be compared at comparable brightness levels and used to show in \\S3 that variations in the red giants are ubiquitous at some magnitude levels. The distribution of oscillation characteristics over the CMD will be developed in \\S4. Section~5 will provide mappings of RGB position on the CMD to physical stellar parameters of mass, radius, luminosity and temperature and comparison of theoretically expected oscillation properties with those observed. Section~5 will also quantify how the variations detected along the galactic bulge RGB scale with $L/M$, compare with theory, and use this to predict expected variations that will be detectable by the {\\em Kepler Mission}. ",
        "conclusions": "Observations over a seven day span with the ACS imager on {\\em HST} have been used to demonstrate that most red giants with luminosities between 30 and 350~$L_{\\odot}$  show evidence of excess noise in power spectra in comparison to main sequence control stars at the same brightness. Amplitudes and frequencies depend upon position in the CMD, and thus differences of stellar parameters of $L$, $M$ and $T_\\mathrm{eff}$ in rough agreement with theoretical expectations.  The best match of frequency for observed power excess in the red giants is consistent with interpretation as low-order radial modes. Since the frequency resolution of these data are not sufficient to resolve expected oscillation characteristics over much of the sampled domain, asteroseismic applications are not provided.  The upcoming {\\em Kepler Mission} will provide improvements of more than an order of magnitude in both photometric precisions and frequency resolution allowing definitive understanding of the nature of variations in red giants all along the branch."
    },
    "0806/0806.3989_arXiv.txt": {
        "abstract": "\\vskip 3pt \\noindent  We study whether spin-independent scattering of weakly-interacting massive particles (WIMPs) with nuclei can account for the annual modulation signal reported by DAMA.  We consider both elastic and inelastic scattering processes.  We find that there is a region of WIMP parameter space which can simultaneously accommodate DAMA and the null results of CDMS, CRESST, and XENON.  This region corresponds to an ordinary, elastically-scattering WIMP with a standard Maxwell-Boltzmann distribution, a mass $3 \\mbox{ GeV} \\lesssim m_{DM} \\lesssim 8 \\mbox{ GeV}$, and a spin-independent cross section with nucleons $3 \\times 10^{-41} \\mbox{ cm}^2 \\lesssim  \\sigma_p^{SI} \\lesssim 5 \\times 10^{-39} \\mbox{ cm}^2$.  This new region of parameter space depends crucially on the recently discovered effect of channeling on the energy threshold for WIMP detection in the DAMA experiment; without the inclusion of this effect, the DAMA allowed region is essentially closed by null experiments.  Such low-mass WIMPs arise in many theories of Beyond the Standard Model physics, from minimal extensions of the MSSM to solutions of the baryon-dark matter coincidence problem. We find that inelastic scattering channels do not open up a significant parameter region consistent with all experimental results.  Future experiments with low energy thresholds for detecting nuclear recoils, such as CDMSII-Si and those utilizing ultra-low energy germanium detectors, will be able to probe the DAMA region of parameter space. ",
        "introduction": "\\label{sec:intro}  Recently, the DAMA collaboration has provided further evidence for the observation of an annual modulation in the rates of nuclear recoil in their experiment \\cite{Bernabei:2008yi}.  Such a signal arises naturally from postulating Weakly Interacting Massive Particles (WIMPs) in the galactic halo that scatter from target nuclei in detectors.  The annual modulation of the interaction rate comes from the variation in the relative velocity of the earth with respect to the galactic dark matter halo as the earth orbits the sun.  This changes the flux of dark matter particles and the size of their interaction cross-sections, with expected extrema occurring at June 2 and December 2.  The DAMA experiment observes a maximum at low nuclear recoil energies on May 24, plus or minus 8 days, and they have accumulated enough data to put the significance of the observed modulation at approximately $8\\sigma$.   Both the phase and amplitude of the signal are highly suggestive of WIMP interactions.  The collaboration has not been able to identify other systematic effects capable of producing this signal, and have claimed that the annual modulation is a discovery of dark matter.  This claim has been controversial, partly because a number of other experiments seem to be in direct contradiction.  In particular, the original DAMA allowed region with WIMP mass $30 \\mbox{ GeV} \\lesssim m_{DM} \\lesssim 200 \\mbox{ GeV}$ and dark matter-nucleon interaction cross-section $\\sigma_p \\simeq 10^{-41}-10^{-42} \\mbox{ cm}^2$ had been quite conclusively ruled out by the CDMS~\\cite{CDMSI,CDMSII,CDMSIISi} and XENON~\\cite{XENON} experiments for the case of an elastically scattering WIMP.  Methods of reconciling the DAMA signal with the results of other experiments have been proposed in the past.  Inelastic scattering processes, $\\chi_1 N \\to \\chi_2 N$, where $\\chi_1$ is the dark matter particle and $\\chi_2$ is another new state with mass splitting $\\delta$ between the states, have been proposed in the context of supersymmetric models~\\cite{Hall:1997ah,Smith:2001hy,TuckerSmith:2004jv}.  Inelastic scattering of MeV dark matter particles to lighter states was investigated as a possible solution in Ref.~\\cite{Bernabei:2008mv}.  Mirror states from a hidden-sector copy of the Standard Model have been proposed as a candidate consistent with all experimental constraints~\\cite{Foot:2005ic,Foot:2008nw}, as have various models with heavy composite states~\\cite{Khlopov:2008ki}.  A model-independent study of spin-independent elastic scattering noted that scattering from the sodium component of the NaI DAMA scintillators allowed a small window of dark matter masses in the $5-9\\,{\\rm GeV}$ to be consistent with current experimental constraints~\\cite{Gelmini,Gondolo:2005hh}.  This study noted that since sodium nuclei are lighter than germanium nuclei, the threshold for scattering off sodium could be lower than that for germanium for light dark matter states.  It was also shown previously that spin-dependent scattering may open up additional parameter space consistent with DAMA and other experiments~\\cite{kami}.  This model-independent study of elastic scattering in Ref.~\\cite{Gondolo:2005hh} is no longer applicable, as several new results have appeared recently.  New experimental constraints and improved understanding of scattering processes in the DAMA apparatus have drastically altered both the excluded parameter space and the physics underlying the DAMA modulation signal.   We perform a model-independent study of both elastic and inelastic scattering mechanisms accounting for all recent experimental measurements.  We find that completely ordinary, spin-independent elastically scattering WIMPs with masses in the range $3-8\\,{\\rm GeV}$ and scattering cross sections in the range $3\\times 10^{-41}\\,{\\rm cm^2}$ to $5\\times 10^{-39}\\,{\\rm cm^2}$ are consistent with all experimental constraints.  No additional dark matter stream is needed; a simple Maxwell-Boltzmann distribution allows this parameter space.  Inelastic scattering no longer opens up a significant region of additional allowed parameter space.  We summarize here the important features and conclusions of our analysis.  \\begin{itemize}  \\item We include the effect of channeling in the NaI crystal scintillators of DAMA, an effect recently noted in Ref.~\\cite{Drobyshevski:2007zj} and studied by the DAMA collaboration~\\cite{Bernabei:2007hw}.  Channeling occurs in crystalline detectors where the only signal measured is the light output, and when recoiling nuclei interact only electromagnetically with the detector material because of either their direction of motion or incident energy.  The effect of channeling is to remove the quenching factor usually required to convert between nuclear recoil energy and electron-equivalent energy, and in the context of dark matter searches it effectively lowers the energy threshold for detection of nuclear recoils of DAMA below that of CDMS and XENON.  This effect is crucial in reconciling elastically scattered WIMPs consistent with all experimental constraints.  In particular, the lower threshold of DAMA means that it can detect lighter dark matter particles than the higher threshold experiments like CDMS and XENON.  This effect opens the region of light WIMP parameter space for DAMA.  We note that the presence of channeling in the energy regime studied by DAMA has not been conclusively established, although it has been observed in NaI crystals at higher energies~\\cite{channel}.  \\item We include constraints from CDMS-SUF, CDMS-II, CRESST-I scattering from sapphire targets~\\cite{CRESST}, % and XENON.  Inclusion of experimental results from multiple target nuclei is necessary to correctly elucidate the allowed parameter region.   In the elastic scattering case we also include the recent results from the CoGeNT collaboration~\\cite{cogent}.  \\item We study spin-independent elastic scattering, and inelastic scattering with either positive or negative mass splitting $\\delta$ between the incident and scattered dark matter particle.  Inelastic scattering of either sign opens up only a very small window of parameter space; roughly, inelastic scattering to heavy states is ruled out by XENON and the germanium data from CDMS-II, while scattering to lighter states is ruled out by CRESST and the silicon data from CDMS.  The preferred parameter space is for light mass dark matter with elastic scattering.  Future results from ultra-low noise germanium detectors~\\cite{Barbeau:2007qi}~\\footnote{We thank J. Collar for correspondence regarding the ability of these experiments to probe this region.}, and the lower threshold silicon data from CDMS will be vital in exploring this region.  The low threshold germanium experiment TEXONO~\\cite{TEXONO} may also be able to probe this region, though its sensitivity has recently been called into question~\\cite{Avignone:2008xc} (inclusion of current TEXONO constraints does not change our results).  \\end{itemize}  There are many possible models which could give rise to such a comparatively light WIMP.  In extensions of the Minimal Supersymmetric Model, for example, GeV mass WIMPs with the right relic abundance arise~\\cite{Gunion,Barger}.  It has been shown that hidden sectors in the context of supersymmetric models give rise naturally to WIMPs with GeV or even lighter masses, as observed in Refs.\\cite{Hooper,Feng}.  It was shown explicitly that the models of this sort can account for the DAMA signal~\\cite{Feng:2008dz}.  Supersymmetric models with non-unified gaugino masses at the grand unified scale give rise to light neutralino dark matter candidates~\\cite{Bottino:2002ry}; the importance of the channeling effect for these models was noted in Ref.~\\cite{Bottino}, as was the effect of having the light sodium component of the DAMA target.  Lastly, solutions to the baryon-dark matter asymmetry problem also predict a WIMP with a mass in the range, $m_{dm} \\approx \\Omega_{dm}/\\Omega_b m_p \\approx 5 m_p$~\\cite{Kaplan,Farrar,Kitano,fw}.  It is clear that the light dark matter paradigm suggested by the direct detection experiments raises numerous theoretical questions and has phenomenological impact on a broad array of experiments.  We leave the potential implications of light WIMPs to future work, and focus here on clarifying the experimental situation.  The outline of the paper is as follows.  In Section~\\ref{sec:form} we review the formalism of direct detection of dark matter.  We review the characteristics of the relevant experiments, discuss the physics and implications of the channeling effect in DAMA, and discuss our analysis method.  We then apply these techniques to derive the allowed parameter space consistent with all experimental measurements for both elastically and inelastically scattered WIMPs in Section~\\ref{numresults}.  Finally, we conclude and discuss future directions. ",
        "conclusions": "We have studied the consistency of the dark matter interpretation of the annual modulation signal observed by DAMA with the results of the null experiments CDMS, CRESST, and XENON.  Recent work has shown the presence of a channeling effect in the crystal scintillators utilized by DAMA which drastically changes the interpretation of the experimental results.  The presence of the channeling effect opens a window in dark matter parameter space between 3 and 8 GeV where the DAMA signal is consistent with all of the null experiments.  This consistency requires no exotic dark matter physics--a vanilla, elastically scattering dark matter candidate interacting through spin-independent channels is sufficient to explain both the signal and the null results from the other experiments.  We have also examined whether possible inelastic processes can accommodate all experimental results.  Inelastic scattering of dark matter particles to heavier final states renders dark matter masses up to approximately 13 GeV consistent with all measurements.  However, the largest range of permissible dark matter masses occurs for elastic scattering candidates, indicating that inelastic processes do not open up significant regions of parameter space.  Future measurements from ultra-low energy germanium detectors and silicon results from CDMS are needed to explore this light-mass window.  The light dark matter window suggested by the DAMA results motivates many new directions for model-building and phenomenology with low-mass WIMP candidates.  Although more model-dependent, the implications of light dark matter  for indirect detection and collider experiments should be explored.  It would be interesting to also consider whether spin-dependent scattering allows a larger range of dark matter masses to be consistent with DAMA and the null experiments."
    },
    "0806/0806.4641_arXiv.txt": {
        "abstract": "We carry out high-resolution FUSE spectroscopy of the nuclear region of NGC 1068. The first set of spectra was obtained with a 30\\arcsec\\ square aperture that collects all emission from the narrow-line region. The data reveal a strong broad \\ion{O}{6} component of FWHM $\\sim 3500$ $\\rm km s^{-1}$ and two narrow \\ion{O}{6} $\\lambda\\lambda 1031/1037$ components of $\\sim 350\\ \\rm km s^{-1}$. The \\ion{C}{3} $\\lambda 977$ and \\ion{N}{3} $\\lambda 991$ emission lines in this spectrum can be fitted with a narrow component of FWHM $\\sim 1000$ $\\rm km s^{-1}$ and a broad one of $\\sim 2500$ $\\rm km s^{-1}$. Another set of seven spatially resolved spectra were made using a long slit of $1\\farcs 25 \\times 20 \\arcsec$, at steps of $\\sim 1$\\arcsec\\ along the axis of the emission-line cone. We find that (1) Major emission lines in the FUSE wavelength range consist of a broad and a narrow component; (2) There is a gradient in the velocity field for the narrow \\ion{O}{6} component of $\\sim 200$ $\\rm km s^{-1}$ from $\\sim 2 \\arcsec$ southwest of the nucleus to $\\sim 4 \\arcsec$ northeast. A similar pattern is also observed with the broad \\ovi\\ component, with a gradient of $\\sim 3000\\ \\rm km s^{-1}$. These are consistent with the HST/STIS findings and suggest a biconical structure in which the velocity field is mainly radial outflow; (3) A major portion of the \\ion{C}{3} and \\ion{N}{3} line flux is produced in the compact core. They are therefore not effective temperature diagnostics for the conical region; and (4) The best-fitted UV continuum suggests virtually no reddening, and the \\ion{He}{2} 1085/1640 ratio suggests a consistently low extinction factor across the cone. At $\\sim 2 \\arcsec$ northeast of the nucleus there is a region characterized by (a) a strong Ly$\\alpha$ flux, but normal \\ion{C}{4} flux; (b) a broad \\ion{O}{6} line; and (c) a significantly enhanced \\ion{C}{3} flux. ",
        "introduction": "NGC 1068 is a prototypical Seyfert 2 galaxy. Because of its proximity ($z=0.0038$) and brightness, it has been studied in nearly every possible detail. The polarimetric observation by \\citet{antonucci}, which reveals a Seyfert 1 spectrum in scattered light, suggests that the nucleus and its associated broad-line region (BLR) are obscured. This finding provides strong evidence for the unified theory in which viewing angles account for the differences between various active galactic nuclei \\citep{antonucci2}  The nuclear region of NGC 1068 harbors a variety of astrophysical phenomena. At the very center of the nucleus there is a bright compact ($<$0\\farcs 3) region commonly referred as ``the hot spot''. Within a few arcseconds from the nucleus, there are several bright and compact clouds that coincide with knots in the radio jets \\citep{wilson,evans}. The narrow-line region (NLR) is conical in shape toward the northeast (Fig. 1), along a position angle of $\\sim 200 \\deg$ and with an opening angle of $\\sim 40 \\deg$. Beyond a 6\\arcsec\\ radius the surface brightness drops dramatically, and emission is dominated by two ring-like filaments at $\\sim10$\\arcsec\\ and 15\\arcsec\\ from the nucleus.  High-spatial-resolution spectroscopy of NGC 1068 has been carried out in the optical \\citep{caganoff,unger,inglis,emsellem,gmos} as well as in the UV \\citep{caganoff,krm2,krm,cecil,groves}. From approximately 2\\arcsec\\ southwest of the nucleus to 4\\arcsec\\ northeast, emission lines exhibit multiple components \\citep{cecil90,krm0}: (1) major emission lines consist of narrow and broad lines; (2) broad lines are approximately 2500-4000 \\kms\\ wide, which may be linked to those that are found in polarized light and believed to be reflected light from the inner BLR, and (3) narrow lines consist of a pair of red and blue components. The [\\oiii] and [\\nii] line profiles suggest that the separation of these two components varies across the conical NLR. In addition to an overall biconical ionization configuration, there are compact knots whose optical spectra resemble kinematically the associated absorption line systems in quasars \\citep{cecil,ckg,krm4}. These line-emitting knots have blueshifted radial velocities up to 3000 \\kms\\ relative to the galaxy's systemic velocity, contributing mostly to the emission-line flux but not the continuum. Between $\\sim$2\\farcs 5 and 4\\farcs 5 northeast from the nucleus, UV line emission is redshifted relative to the systemic value, a pattern that is interpreted as the expansion of the plasma in the radio lobe \\citep{axon}.  Several important emission lines in the far-UV (FUV) region between 912 and 1150~\\AA\\ are observable only with specially crafted UV instruments. During the Astro-1 mission the Hopkins Ultraviolet Telescope (HUT) observed NGC 1068 with 18\\arcsec\\ and 30\\arcsec\\ apertures. The most striking features in the wavelengths below 1150 \\AA\\ are the strong \\ciii\\ \\lm 977 and \\niii\\ \\lm 991 lines. The line intensity ratios of \\ciii\\ I(\\lm 1909)/I(\\lm 977) and \\niii\\ I(\\lm 1750)/I(\\lm 991) are temperature sensitive, and the derived temperature is $> 25$ 000 K \\citep{kriss}, higher than the values expected for a region producing \\ion{C}{3} and \\ion{N}{3} emission by photoionization, The line ratios in \\ngc\\ are similar to those of the Cygnus Loop supernova remnant \\citep{blair}, suggesting a significant contribution from shock-heating mechanisms. Astro-2 observations of NGC 1068 were obtained with a 12\\arcsec\\ aperture at three different positions. The results \\citep{grimes} suggest that the emission lines observed with HUT likely arise in the inner nuclear region imaged with HST. However, the poor angular resolution of the HUT instrument ($> 10\\arcsec$) does not allow a study of the NLR in terms of spatial details.  The line-emitting mechanisms in the NLR, {\\em i.e.} photoionization from the nucleus or shocks produced by jets, have long been under debate. \\citet{dopita} and \\citet{bicknell} proposed that emission in the NLR may be entirely caused by shocks. Velocity splitting over 1000 \\kms, reported by \\citet{axon}, in the vicinity of some of the bright emission-line knots provides evidence that fast shocks exist in the NLR of NGC~1068. However, Seyfert galaxies host powerful nuclear sources of ionizing radiation, and the situation can be more complex with shocks or nuclear photoionization prevailing in different environments \\citep{allen,morganti}. \\citet{morse} suggested that photoionizing shocks are important when a radio jet interacts with the interstellar medium, but not in the objects where sharp, straight-edged ionization cones are observed. \\citet{ferguson} argued that strong \\ciii\\ \\lm 977 and \\niii\\ \\lm 991 emission may arise from fluorescence in photoionized gas if turbulent velocities exceed $\\sim$1000 \\kms.  However, \\citet{grimes} found that such a velocity would lead to extreme physical conditions that are inconsistent with the Astro-2 data. More recent HST data \\citep{krm,cecil} found that the emission-line ratios  are consistent with photoionization instead of shock heating mechanism.  We have carried out observations with the {\\em Far Ultraviolet Spectroscopic Explorer} (FUSE) to study the spatial distribution of FUV emission lines. In this paper we present the results of both large-aperture and spatially resolved spectroscopic observations with FUSE. For the first time, we are able to study the position dependence of several important diagnostic lines in the FUV band. ",
        "conclusions": "Many previous studies have gradually unraveled the kinematical complexities of gas in the nuclear region of NGC 1068. The X-ray data \\citep{xmm} suggest that emission lines are formed mainly in a photoionied plasma of a temperature around a few eV. The spatial resolution afforded by HST leads to a picture of decelerating jet \\citep{das}: a biconical outflow from the nuclear region sweeps up denser, ambient clouds in the interstellar medium (ISM) of NGC 1068. Other possibilities include overlapping, discrete ejection that gradually dissipates \\citep{axon,capetti}. In addition, a high-velocity radio jet impinges on some of the clouds. Some gas expands perpendicularly to the axis of the jet, and the expanding radio lobe at the end of the jet also pushes on the ambient ISM. Near the nucleus, kinematic components span several thousand kilometers per second in velocity, and the continuum hot spot visible in HST images reflects a polarized view of the broad lines in the active nucleus.  Our spatially resolved FUSE observations, while not at the resolution of HST, add information from major emission lines shortward of the HST bandpass at high spectral resolution. We plot the fluxes of major UV emission lines from our observations at seven slit positions in Fig. 8, and list the fitted line properties in Table 4. In this section, we discuss the four emission lines in the FUSE spectral range, along with a comparison to emission lines observed in the STIS spectra.  \\subsection{\\ovi\\ emission}  The most prominent feature in the FUSE spectra is the \\ovi\\ \\lm\\lm 1031,1037 emission line. In the data taken with LWRS, the two narrow \\ovi\\ components are of FWHM $\\sim 350$ \\kms\\ and a separation of $\\sim 200$ \\kms. In the seven spectra taken with a narrow slit, we resolve this blend into one narrow component with FWHM of $\\sim 350$ \\kms\\ and one broad component. In Fig. 9 we plot the \\ovi\\ profiles at the seven different slit positions. The high spectral resolution of FUSE data enables us to compare with the results of optical Fabry-Perot spectroscopy \\citep{cecil90}, which reveal a narrow core of $\\sim 300$ \\kms. According to the optical data, nearly 75\\% of the [\\nii] \\lm 6583 flux is from components of $\\sim 1500$ \\kms\\ wide. Line widths at such a scale are consistent with that derived from HST UV spectroscopy. The narrow-line flux in the FUSE spectra is highly concentrated (60\\%) at the compact core, suggesting that the ``true'' NLR probably remains unresolved, at a sub-arcsecond scale. While the narrow O VI emission line is not the dominant component, its distribution is different from its broad counterpart, while only 40\\% of the broad \\ovi\\ line flux is from slit position C.  It is surprising that the \\ovi\\ emission is dominated by a component that is broader than those seen in polarized light: in the LWRS spectrum, more than 3/4 of the \\ovi\\ flux is from a component with FWHM $\\sim 3500$ \\kms\\ that is blue-shifted relative to the narrow component by $\\sim 500$ \\kms. This broad component is present in all seven FUSE slit spectra with considerable strength. It may arise from the reflected emission from the hidden BLR, and/or may be the result of a significant velocity dispersion in the NLR. The HST FOC data \\citep{axon} reveal that emission lines near the hot knots 2\\arcsec\\ northeast (FUSE slit positions D and E) are split into two velocity systems separated by $\\sim 1500$ \\kms. The STIS spectra discussed by \\citet{groves} show [\\oiii] emission knots spanning such a broad velocity range in the immediate vicinity of the nucleus, but not at distances of several arcseconds. Since the FUSE slit collects emission from a block of regions spanning several arcseconds perpendicular to the conical axis, the total line emission from these regions may be blended into one broad component.  The ratios of \\civ/\\ovi\\ may provide insight into the physical conditions of the line-emitting regions. The value is higher for the narrow components than their broad counterparts, implying a range of the ionization parameter $U \\sim 0.05-1.0$ in a typical photionization calculation. High values of $U>1$ are consistent with models that assume the same origin for the associated absorbers and BLR \\citep{n7469}, suggesting that the clouds that produce the broad emission components may be of the same origin as associated absorbers.  \\subsection{Velocity field}  The narrow \\ovi\\ line exhibits a systematic velocity shift from position A to G by approximately 220 \\kms. This gradient in the spatially resolved spectra explains why there are two narrow-line components in the integrated flux from the large-aperture FUSE spectrum, where it is unresolved. \\citet{krm0} reported a similar velocity pattern in their HST spectra. \\citet{das} successfully modeled this as a biconical outflow in which radial velocity changes as a function of the distance to the central nucleus: the emission line knots show evidence for radial acceleration to a projected distance of 2\\arcsec\\ to the northeast direction, followed by deceleration up to 4\\arcsec. The \\ovi\\ line widths also increase at $\\pm 2$\\arcsec\\ from the nucleus, probably implying a larger dispersion in these regions.  The broad component of \\ovi\\ exhibits a qualitatively similar kinematic pattern, but at a larger amplitude: its centroid shifts by $\\sim 1500$ \\kms\\ (Fig. 10) across the same spatial region. Fig. 11 shows similar trends for the narrow and broad components of \\civ\\ in the HST spectra which have not been explicitly noted in previous studies: the narrow component follows the kinematic pattern of the optical lines modeled by \\citet{das}; the broad component of \\civ\\ shows behavior similar to that of \\ovi\\ in the FUSE spectra. Prior observations that noted this blue-shifted broad component invariably attributed it to the reflection of the BLR. The blue shift and width (at the position of the hot spot, and in integrated light) are comparable to the broad polarized H$\\beta$ line observed by \\citet{antonucci}.  In a scattered BLR picture, a blue shift is caused by the outflowing wind from the torus along our line of sight, and a line width is due to the intrinsic broad line width convolved with the thermal width of the hot reflecting wind. As a broad component is present in the FUSE spectra at all seven positions, it is possible to assume that this is reflected light from the hidden BLR. However, a large covering factor is needed to explain the observed fluxes. Assuming a covering factor of 0.1, the intrinsic flux of the broad \\ovi\\ emission in NGC 1068 would exceed that in NGC 4151. A more reasonable explanation for the observed broad line widths is the large velocity dispersion between bright knots. At approximately 2\\arcsec\\ from the core (position A and E), the FWHMs of broad components are the broadest at 3200 \\kms. These maxima coincide with the widest spliting of velocity in bright knots \\cite{krm0}. With an intrinsic dispersion of $\\sim$800 \\kms\\ and a separation in velocity of $\\sim$2500 \\kms\\ between bright knots (from STIS results), data collected by the long FUSE slit would exhibit a broad component of $\\sim$3200 \\kms, which is what observed at positions A and E. Extended regions of hot, photoionized gas are seen in X-ray images of NGC 1068 \\citep{young,xmm} that could be visible manifestations of this hot outflow. It is natural to assume that these high-velocity clouds may be related to the associated absorbers \\citep[][and references therein]{ckg} in AGN, which are mostly blueshifted. As with the lower-velocity, lower ionization emission-line gas, the acceleration of this high ionization gas eventually is brought to a halt by an unknown deceleration mechanism, which might plausibly be interaction with the ambient ISM of NGC 1068. The evidence for deceleration at arcsecond scale may suggest that acceleration of the outflow materials may take place at sub-arcsecond scales.  \\subsection{\\ciii\\ \\lm 977}  The flux ratio of \\ciii\\ \\lm 1909 to 977 is extremely sensitive to temperature, and the value measured in N1068 suggests a high temperature that is consistent with shock-heating \\citep{kriss}. The FUSE data taken with a large aperture (Table 3) reveal that this emission line consists of a narrow and a broad component. The broad component is weak and hence cannot be well separated in the data segments taken with a narrow FUSE slit. The STIS data reveal that \\ciii\\ \\lm 1909 emission can be fitted with a narrow and a broad component of FWHM $\\sim 900$ and 3500 \\kms, respectively. In principal, the ratio of \\ciii\\ \\lm 1909 to 977 should be calculated only between the narrow components.  As shown in Fig. 8, the flux of \\ciii\\ \\lm 977 is highly concentrated in position C and D. High temperatures implied by this line emission may therefore be associated with the compact core. In a large portion of the ionization cone, the line ratios (\\lm 1909/\\lm 977) are considerably higher, suggesting a lower temperature. However, the line ratios are considerably lower than that derived from the HUT data, where broad \\ciii\\ \\lm 1909 component was included in the calculation.  \\subsection{\\niii\\ \\lm 991}  The line ratio of \\niii\\ \\lm 1750 to \\lm 991 is also temperature dependent, and its value has also been used to derive a high temperature in the ionization cone. The FUSE data also reveal a pair of components in the data taken with a large aperture. Unlike \\ciii\\ \\lm 977, the broad component of \\niii\\ \\lm 991 , as shown in Table 3, is stronger than its narrow counterpart. The FUSE data taken with a narrow slit shows that nearly a half of this broad component is in position C. In other positions, the \\niii\\ \\lm 991 emission can only be fitted with one component of FWHM $\\sim 1000$ \\kms. In combination  with the trend for \\ciii\\ \\lm 977, we conclude that a bulk of flux in these temperature-sensitive emission lines is from the compact core of NGC 1068, and their intensity is not directly tied to the physical conditions in the ionization cone.  \\subsection{\\he\\ 1085}  The distribution of \\he\\ \\lm 1085 is different from other emission lines in the FUSE spectra: as shown in Fig. 8, the \\he\\ flux varies smoothly across the NLR region, like that of \\lya, \\civ\\ and other lines in the STIS spectra. Since \\he\\ emission is believed to be produced mainly by recombination and is insensitive to the gas temperature, this pattern of variation may simply reflect the distribution of NLR gas across the ionization cone. The ratio of \\he\\ emission \\lm 1640 to \\lm 1085 is believed to be a reddening indicator. As shown in Fig. 12, the ratio is nearly constant ($\\sim 5.5$) between slit positions B and F. This is consistent with the HUT result of $5.8 \\pm 1.6$, suggesting an insignificant level of extinction ($E_{B-V} \\le 0.05$). The low values at positions A and G may suggest that the data at these positions are not reliable.  The \\he\\ I(1640)/I(1085) ratio is slightly lower than that expected from recombination, and this may suggest that the \\he\\ \\lm 1085 may be slightly contaminated by other weak UV emission lines. \\ion{N}{2} \\lm 1085 is a likely candidate, and the trend of decreasing \\he\\ (1640)/I(1085) ratio with distance from the nucleus supports this as there is a noticeable decrease in ionization state at larger distances from the nucleus \\citep{axon,krm2,krm3}. As with \\ovi\\ and \\civ, there appears to be a broad component that is considerably blueshifted. However, the relevant wavelength range is in a gap between the LiF channels, and only SiC2B data are available. Because of the low S/N ratio, the reality of such a broad component is still questionable, as no such counterpart is found in the corresponding HST spectra of \\he\\ \\lm 1640."
    },
    "0806/0806.3899_arXiv.txt": {
        "abstract": "{% I have compiled observations of \\otres line and 5 GHz radio emission for a large sample of GPS, CSS and FR sources. Several properties were studied and compared. The most relevant findings are that the FWHM and the luminosity of the \\otres line are correlated with the size of the radio source. I present the data and discuss the correlations, with special focus on jet-host interaction, triggering and enhancing of \\otres emission.} ",
        "introduction": "Current models for the evolution of powerful radio galaxies suggest that these sources propagate from the $\\sim 10$ pc to Mpc scales at roughly constant velocity through an ambient medium which declines in density as $\\rho(R) \\propto R^{-2}$ while the sources decline in radio luminosity as $L_{rad} \\propto R^{-0.5}$ (e.g., O'Dea 2002, and references therein). In this scenario, GPS would evolve into CSS and these into large -supergalactic sized- sources\\footnote{Recent work added the High Frequency Peakers to the sequence, as possible progenitors of GPS (e.g., Orienti 2007 et al. and references therein).}. Such a scenario is consistent with the observed number densities of powerful radio sources as a function of linear size (e.g. O'Dea \\& Baum 1997, Fanti et al. 2001). However, to match observations, the radio jets of the young sources must slow as they cross the host ISM and dim faster than predicted. The most likely mechanism to produce these effects is interaction of the radio source with the host environment.  Interaction was found (e.g., Labiano et al. 2008;  Holt et al. 2006; Axon et al. 2000). However, the studies focused in small samples or even just a few sources. Until now, there was no study of interaction of a large representative sample of GPS and CSS sources. Furthermore, the sparseness of samples studied did not allow a general study of the consequences of interaction on the models or the host.  In order to study the interaction between the gas clouds and the radio source I collected a sample of almost one hundred sources, including GPS and CSS galaxies and quasars, as well as FR sources. I study the properties of the \\otres line and 5 GHz radio emission looking for traces of interaction and the mechanisms responsible for it. ",
        "conclusions": "\\label{sec:results}  The following properties of the \\otres line and 5 GHz radio emission were compared: FWHM, luminosity, asymmetry, kurtosis (\\otres line), power, size and turnover frequency (radio emission) in the sample.  The \\otres FWHM shows no correlation with radio power, turnover frequency and \\otres luminosity. Therefore, the shock velocity (closely related to the FWHM, see e.g., Bicknell et al. 1997) is independent of the strength radio source and does not affect the luminosity of the ionized gas or the radio spectral properties of the source. However, the data suggests a possible correlation between \\otres FWHM and radio source size (Figure \\ref{FWHMLS}) suggesting a possible deceleration of the jet as it crosses the host:   \\begin{center} log FWHM $\\simeq  2.89(\\pm0.04) - 0.08(\\pm0.05) \\times$ log LS \\end{center}  The correlation is not very strong. However, the deceleration of the jet is required to predict observations (e.g., O'Dea 1998) and the latest models are predicting decelerations in the jet (see e.g., Kawakatu et al. 2009, in this volume). Unfortunately, there is not much data available of \\otres FWHM.  \\begin{figure} \\centering \\includegraphics[width=\\columnwidth]{FWHMLS2.eps} \\caption{Plot of the \\otres FWHM of the sources in Gelderman \\& Whittle 1994 (diamonds), de Vries et al. 2000 (triangles) our Kitt Peak (circles),  STIS (shaded diamonds) observations, and a linear fit to the data. The ionization and kinematics of \\object{3C~67}, \\object{3C~277.1} and \\object{3C~303.1} are studied in Labiano et al. (2005) and O'Dea et al. (2002).  \\label{FWHMLS}} \\end{figure}   \\begin{figure} \\centering \\includegraphics[width=\\columnwidth]{RPLS.eps} \\caption{Radio power at 5GHz versus linear size for GPS, CSS and large radio sources. \\label{RPLS}} \\end{figure}  As expected (e.g., O'Dea \\& Baum 1997, the sample also shows no correlation between radio power and radio size (Figure \\ref{RPLS}), showing that GPS, CSS and FR2 sources are equally powerful (log Power$_{\\mathrm{5 GHz}}\\sim10^{26-27}$) while FR1 sources tend to be fainter.  It is also clear that quasars tend to be brighter in \\otres radio than galaxies, consistent with Unification scenarios: some of the \\otres may be hidden by the torus in radio galaxies ()e.g. Hes et al. 1993). However, the difference in \\otres emission could be due to selection effects: quasars are usually found at higher redshifts. The sample could be missing fainter quasars with luminosities similar to radio galaxies. %  \\begin{figure} \\centering \\includegraphics[width=\\columnwidth]{O3RP.eps} \\caption{Plot of the \\otres luminosity versus radio power at 5GHz. Data for the CSS and GPS sources from Gelderman \\& Whittle (1994); O'Dea (1998), de Vries et al. (2000), the 2-Jy sample and our Kitt Peak observations. Data for the large radio sources from the 2-Jy sample. The IDs of the sources have been updated with Zirbel \\& Baum (1995). I use \"C/G\" to name those sources with no clear ID as CSS or GPS. \\label{O3RP}} \\end{figure}  The sample shows the known relation (e.g., Rawlings \\& Saunders 1991;  Baum \\& Heckman 1989) between \\otres luminosity and radio power (Figures \\ref{O3RP} and \\ref{O3RPzoom}) for powerful (log Power$_{\\mathrm{5 GHz}}> 10^{25}$) radio sources. This relation is usually explained as the AGN powering both the ionized gas and radio emission. For our sample, the different correlations are:  \\begin{flushleft} GPS: \\end{flushleft} \\begin{center}  log L$_{[\\ion{O}{iii}]}$ = 32($\\pm$4) + 0.4($\\pm$0.1) $\\times$ log P$_{5GHz}$ \\end{center} CSS: \\begin{center}  log L$_{[\\ion{O}{iii}]}$ = 22($\\pm$4) + 0.8($\\pm$0.2) $\\times$ log P$_{5GHz}$\\end{center} GPS + CSS: \\begin{center} log L$_{[\\ion{O}{iii}]}$ = 25($\\pm$3) + 0.7($\\pm$0.1) $\\times$ log P$_{5GHz}$ \\end{center} GPS + CSS + large: \\begin{center} log L$_{[\\ion{O}{iii}]}$ = 12($\\pm$2) + 1.13($\\pm$0.07) $\\times$ log P$_{5GHz}$ \\end{center}   It should be noted that the correlation is not present in fainter radio sources (SDSS, Best 2008, in this volume), suggesting that a different mechanism may be producing the radio and \\otres emission in these sources.  There is no evident relation between kurtosis or asymmetry of the \\otres line with other properties of the line. Most of the sources have kurtosis lower than 1, suggesting the presence of broad wings in the gas. Whittle (1985) also finds the same for his sample of Seyfert galaxies. Most sources shows asymmetry values close to 0 in their \\otres profile, suggesting the cocoon widens or narrows in a symmetric way  or there are no major variations between both sides. However, the ground spectra may lack of enough resolution to separate more complex structures that could change the profile of the line.   O'Dea (1998) discovered that in the Gelderman \\& Whittle (1994) sample, GPS galaxies tend to have lower \\otres luminosity than CSS galaxies. However, it was not clear if the trend would be followed by a larger sample with different selection criteria, what was the behaviour for quasars and large FR sources, or the possible consequences for general radio source evolution\\footnote{A direct consequence is that it is more difficult to find optical counterparts of GPS than CSS.}. To asses these issues, I tested the trend with the current sample, which also includes GPS and CSS quasars and supergalactic-sized sources. I find that the GPS and CSS sources (galaxies and quasars) show a strong correlation between \\otres line luminosity and size of the radio source\\footnote{Note that the trend is also suggested by Figure \\ref{O3RPzoom} and the different correlations between radio power and \\otres luminosity for GPS and CSS sources.}:   %   \\begin{flushleft} GPS + CSS: \\end{flushleft} \\begin{center} log L$_{[\\ion{O}{iii}]}$ = 42.43($\\pm$0.09) + 0.46($\\pm$0.09) $\\times$ log LS$_{5GHz}$ \\end{center}  \\begin{flushleft} GPS$^*$ + CSS:  \\end{flushleft} \\begin{center} log L$_{[\\ion{O}{iii}]}$ = 42.44($\\pm$0.08) + 0.4($\\pm$0.1) $\\times$ log LS$_{5GHz}$ \\end{center}  GPS$^*$ means the complete sample of GPS except the smallest source.   In principle, this correlation could be due to the AGN enhancing both the radio and emission gas luminosities by photoionization. However, for the same radio power, small sources (GPS) are systematically fainter in \\otres than larger (CSS) sources.  I propose a scenario where the expansion of the radio source through the host ISM is triggering and/or enhancing the \\otres line emission through direct interaction. Some contribution from the AGN light must be present, however, AGN light alone, would not produce a correlation with size. Furthermore, the fact that the correlation disappears for supergalactic-sized sources supports this model: once the radio lobes leave the host, the \\otres luminosity drops (Figure\\ref{O3LS}). %  This scenario is also supported with previous observations finding evidence of strong interaction between the jet and surrounding ISM, as well as proof of shock-ionized \\otres (e.g., Labiano et al. 2005; O'Dea et al. 2003). Jet-ISM interaction is also found through  \\ion{H}{i} studies (e.g., Labiano et al. 2006; Holt et al. 2006, and references therein) and predicted by jet expansion models (e.g., Jeyakumar et al. 2005; Saxton et al. 2005)   Another interesting enhancing mechanism to consider is that the jet could enhance \\otres, at least partly, through indirect mechanisms such as jet-induced star formation (Labiano et al. 2008, in prep.). However, new data are needed to study the possible contribution of recently formed stars. The hosts of GPS and CSS sources are usually elliptical galaxies so it is unlikely that the average/normal stellar population of the host has a strong to the \\otres emission. These ``normal'' stars would however not create a correlation with radio jet size (and jet-induced stars would).  The correlation divided by source type is:\\\\  GPS: \\begin{center} log L$_{[\\ion{O}{iii}]}$ = 42.8($\\pm$0.2) + 0.8($\\pm$0.2) $\\times$ log LS$_{5GHz}$ \\end{center}  GPS$^*$: \\begin{center} log L$_{[\\ion{O}{iii}]}$ = 42.5($\\pm$0.4) + 0.5($\\pm$0.4) $\\times$ log LS$_{5GHz}$ \\end{center}  CSS: \\begin{center} log L$_{[\\ion{O}{iii}]}$ = 42.6($\\pm$0.2) + 0.2($\\pm$0.3) $\\times$ log LS$_{5GHz}$\\end{center}  The correlation tends to disappear when the sample is divided in different types of sources. This could be due to low statistics or to the smaller range of sizes covered by each type. GPS could show lower \\otres due to high obscuration. However, it is more likely that young compact sources are too small to strongly affect their environment (this effect has also been observed in star formation histories of GPS hosts, Labiano et al. 2008). The UV luminosities of GPS seem to be as high as the CSS luminosities. Furthermore, the UV luminosity of GPS sources could be correlated with their radio power (Labiano et al. 2008. These two effects suggest that obscuration is not too strong in GPS sources or, at least, similar to CSS.   \\begin{figure} \\centering \\includegraphics[width=\\columnwidth]{O3RPzoom.eps} \\caption{As Figure \\ref{O3RP}, showing only the GPS and CSS sources. \\label{O3RPzoom}} \\end{figure} \\begin{figure} \\centering \\includegraphics[width=\\columnwidth]{O3LS.eps} \\caption{Plot of the [\\ion{O}{iii}] luminosity versus linear size of the radio source, showing the correlation for GPS and CSS sources. Data from the same references as figure \\ref{O3RP}. \\label{O3LS}} \\end{figure}  Concerning the overall scenario of radio source evolution, where GPS and CSS evolve into the large, supergalactic sized, sources, the visual inspection of Figure \\ref{O3LS} suggests that CSS would evolve into FR2. Some authors also found a possible decreasing trend linking FR2 to FR1 (Best 2008, private communication) with increasing size. However, our sample needs more supergalactic-sized sources to address evolution beyond $\\sim15-20$ kpc. Extensive discussions on the FR2 - FR1 connection can be found on the literature: see e.g., M{\\\"u}ller et al. 2004; Best et al. 2005; Wold et al. 2007 and references therein, or classical papers such as Baum et al. (1995) and Zirbel \\& Baum (1995)."
    },
    "0806/0806.3582_arXiv.txt": {
        "abstract": "We present ground-based optical and \\textit{Spitzer Space Telescope} infrared imaging observations of the ecliptic (Jupiter-family) comet 21P/Giacobini-Zinner, the parent body of the Draconid meteor stream, during its 2005 apparition. Onset of nucleus activity occurred at a pre-perihelion heliocentric distance, $r_{h} \\simeq 3.80$~AU, while post-perihelion 21P was dusty (peak $Af\\rho = 131$~cm$^{-1}$) and active out to heliocentric distances $\\gtsimeq 3.3$~AU following a logarithmic slope with $r_{h}$ of $-2.04$. Coma colors, $V - R = 0.524 \\pm 0.003, R - I = 0.487 \\pm 0.004$ are redder than solar, yet comparable to colors derived for other Jupiter-family comets. A nucleus radius of $1.82 \\pm 0.05$~km is derived from photometry at quiescence. \\spitzer{} images post-perihelion exhibit an extensive coma with a prominent dust tail, where excess emission (over the dust continuum) in the 4.5~\\micron \\ IRAC image arises from volatile gaseous CO and/or CO$_{2}$. No dust trail was detected ($3\\sigma$ surface brightness upper-limit of 0.3 MJy~sr$^{-1}$~pixel$^{-1}$) along the projected velocity vector of comet 21P in the MIPS 24~\\micron{} image suggesting that the number density of trail particles is $\\ltsimeq 7 \\times 10^{-11}$~m$^{-3}$. The bolometric albedo of 21P derived from the contemporaneous optical and \\spitzer{} observations is $A(\\theta=22\\degr)=0.11$, slightly lower than values derived for other comets at the same phase angle. ",
        "introduction": "}  Comet nuclei formed beyond the protoplanetary disk frost line \\citep[heliocentric distances, $r_{h}$ \\gtsimeq 5~AU;][]{lunine04}, among the giant planets and were scattered into the Kuiper Belt and beyond into the Oort Cloud (OC). Since their formation, the interiors and surfaces of most comets have remained at temperatures below 140 K while in ``cold storage'' in the Kuiper Belt or the OC \\citep{meech04}. Moreover, most nucleus surfaces have remained below 400 K even during perihelion passage. At such low temperatures, dust mineralogy remains stable and each comet nucleus retains a record of the minerals, ices, and volatiles extant in the comet agglomeration zones in the early solar system. Furthermore, comet nuclei may retain their primordial compositional inhomogeneities so that different topographic regions would have different compositions \\citep[e.g.,][]{dellor07}, leading to variations in volatile production rates as a function of nucleus rotation. Nucleus heterogeneities are apparent in the fly-by imagery of ecliptic comet 9P/Tempel 1. Regions of distinct topography \\citep{belton07}, and spatially distinct sites of water and CO$_{2}$ release \\citep{feaga07, ahearn05} are evident, as well as heterogeneities in surface and subsurface composition \\citep{harker07, kado07}.  There are two general dynamical families of comets, classified by derived orbital elements (i.e., Tisserand parameter, $T_{J}$) based on current observations.  Nearly-isotropic comets (NICs; $T_{J} < 2$) have orbits that are approximately uniformly distributed in inclination, and are derived from the OC. Ecliptic comets (EC) have orbits that are confined to inclinations near the ecliptic plane, $2 \\ltsimeq T_{J} \\ltsimeq 3$, and originate in the Kuiper Belt. Resulting from frequent perihelion passage over the 4.0 -- 4.5 Gyr period since their formation, EC comets have become noticeably less active than OC comets, characterized by lower gas and dust production rates. Multi-epoch spectral energy distributions (SEDs) of ECs, from which comae dust properties can be constrained, are needed to assess the possible interrelationships between their reduced activity levels and dust properties.  The study of the physical properties of cometary nuclei and comae both are equally important to our understanding of the outer solar system environment during the era of icy planetesimal formation, and complement efforts to discern conditions extant in early protoplanetary disks during the epoch of planetesimal formation. We present new optical and contemporaneous \\spitzer \\ observations of comet 21P/Giacobini-Zinner obtained during its 2005 apparition obtained as part of a larger survey of both ECs and NICs \\cite[e.g.,][]{kelley06}. Ground-based optical observations enabled us to obtain precise optical photometry to asses variations in dust productivity with heliocentric distance and to study the the comet's near-nucleus structures, including its jets and coma. \\spitzer \\ images at mid-infrared wavelengths enable the study of the spatial distributions of volatiles and dust in the coma, as well as facilitating investigation of comet trail properties. We describe our observations in \\S\\ref{obs}, and discuss analysis of our optical and infrared (IR) observations in \\S\\ref{disc}, while our conclusions are summarized in \\S\\ref{concl}. ",
        "conclusions": "}  We have presented new optical and \\spitzer{} infrared observations of comet 21P/Giacobini-Zinner obtained during its 2005 apparition. Analysis of optical imagery indicates that 21P was dusty (peak $Af\\rho = 131$~cm$^{-1}$) and active out to heliocentric distances $\\gtsimeq 3.3$~AU following a logarithmic slope with $r_{h}$ of $-2.04$. Onset of nucleus activity occurred at a pre-perihelion distance $r_{h} \\simeq 3.80$~AU ($-375$~days pre-perihelion), similar in behavior to that observed in the 1991 apparition. The derived average coma colors, $V - R = 0.524 \\pm 0.003, R - I = 0.487 \\pm 0.004$ are slightly redder than solar, comparable to colors derived for other Jupiter-family comets. Pre-perihelion observations during quiescence yields a nucleus radius of $1.82 \\pm 0.05$~km.  \\spitzer{} IRAC images obtained at $r_{h} = 2.4$~AU, post-perihelion exhibit an extensive coma with a prominent dust tail, where excess emission (over the dust continuum) in the 4.5~\\micron \\ image at cometocentric distances of $\\sim 10^{4}$~km likely arises from CO$_{2}$, although a distributed source of CO cannot be discounted. The upper limits to the production rates are $Q_{CO_{2}}\\leq (5.13\\pm0.75)\\times10^{25}$ molecules~s$^{-1}$ and $Q_{CO}\\leq(6.01\\pm0.89)\\times10^{26}$ molecules~s$^{-1}$. The surface brightness of the gas emission is observed to peak along the sun angle, while the dust tail peaks near the anti-sunward angle. A search for dust trail emission along the projected velocity vector of comet 21P using our MIPS 24~\\micron{} image ($r_{h} = 2.4$~AU), yielded no trail ($3\\sigma$ surface brightness upper-limit of 0.3 MJy~sr$^{-1}$~pixel$^{-1}$), suggesting that the number density of trail particles (typical particle size $\\sim 1$~mm) is $\\ltsimeq 7 \\times 10^{-11}$~m$^{-3}$. The bolometric albedo of 21P derived from the contemporaneous optical and \\spitzer{} observations is $A(\\theta=22\\degr)=0.11$, slightly lower than values derived for other comets at the same phase angle."
    },
    "0806/0806.3061_arXiv.txt": {
        "abstract": "%  We have obtained 8.4 GHz VLBA observations of a 31-GHz complete sample of $\\sim100$ sources between 10 and 100 mJy.  The main goals of these observations are: to determine the angular size, radio spectra and identification for a weak sample of high frequency sources; to find the fraction of sources which have sufficiently compact emission for use as calibrators for VLBI observations; and for design considerations of the proposed DSN Array.  We find that a large fraction of observed sources have VLBI detections.  A majority of these sources have most of their emission in a compact $<1$ mas radio core, with remaining sources having steep radio spectra.  The source list was provided from GBT observations to remove discrete sources in the CBI fields. ",
        "introduction": "Carrying out a VLBI survey of a complete and unbiased sample of weak radio sources at high frequencies provide fundamental astronomical information on the statistical and morphological properties for this class of astrophysical objects.  Previous surveys of the nature and structure of weak radio sources have been carried out at relatively lower frequencies, often at 1.4 GHz \\citep[e.g.][]{garrett2005}.  At the mJy level, the proportion of AGN's are decreasing and the population begins to be dominated by galaxies that have significant star forming regions.  These are typically less than 3'', with about 30\\% showing milliarcsecond emission \\citep{muxlow2005}.  However, the angular characteristics of sources above 8 GHz are not well-known at the mJy level.  We plan to determine the percentage of compact milliarcsecond emission, its orientation and accurate core position for better optical identification.  We also plan to study spectral index correlation versus galaxy type and compare our results with similar studies carried out for brighter sources and similar surveys at lower frequencies. ",
        "conclusions": "We carried out a small 8 GHz VLBA survey of a complete sample of radio sources down to a flux density of 10 mJy.  Sources were identified from a sample of NVSS sources with 31 GHz GBT detections.  We detected $\\sim$ 50\\% of the observed sources with VLBI components.  As expected sources with flatter spectra tend to exhibit mas emission.  In addition, we note that the compactness factor increases for sources with flatter spectra.  Our preliminary results also indicate a relationship between optical identification and source compactness."
    },
    "0806/0806.3127_arXiv.txt": {
        "abstract": "We use a combination of new AAOmega multi-object wide-field spectroscopic observations and literature data to define 111 spectroscopically confirmed members of the massive NGC\\,5044 group with $M_B\\leq-13.5$\\,mag, providing a three-fold increase in group members over previous analyses of this group.  We find the group to have a dynamical mass of 9.2$\\times 10^{14}$\\,M$_{\\odot}$, placing it on the border between rich groups and poor clusters.  However, comparison to the $L_X$--$\\sigma_{\\nu}$ and $L_X$--Mass relations shows it more closely follows cluster scaling relations.  Using a combination of crossing time, X-ray contours and line-of-sight velocity profile we are able to preclude growth of the NGC\\,5044 group via major sub-group mergers within the last $\\sim$1\\,Gyr.  While the majority of dynamical indicators for the group suggest it is virialized, we find evidence for a small, dynamically distinct sub-group at 1.4\\,Mpc from the group centre, suggesting that the NGC\\,5044 group is the dominant structure in its local environment, and is currently accreting smaller groups. ",
        "introduction": "\\label{intro}  It is well known that our own Galaxy is not isolated, residing with M31 and a host of smaller galaxies in the Local Group.  Large area surveys have confirmed that the Local Group is not exceptional, but rather representative of an extremely common environment (e.g.  Geller \\& Huchra 1983; Eke \\etal 2004; Weinmann \\etal 2006).  However, despite the apparent commonality of the group environment, the astrophysical processes acting within groups and their effects on observed galaxy populations are relatively poorly understood.  This is due to both the low galaxy numbers and galaxy surface densities typically associated with groups that make them observationally expensive.  Consequently, it is observations of galaxy clusters that have largely driven our current understanding of environmental effects on galaxy populations, namely that galaxies in clusters are redder (e.g.  Butcher \\& Oemler 1984) and morphologically earlier types (e.g. Dressler 1980) relative to galaxies in the field.  The expected correlation of star formation rate with density, based on the morphology-density and morphology-radius relations, has led observers to measure star formation rates as a function of environment in large samples such as the 2dFGRS (e.g. Lewis \\etal 2002) and SDSS (e.g. G\\'omez \\etal 2003).  These authors have found that star formation rates decrease as one moves from the field to cluster environment.  Perhaps more interestingly, however, is that the observed star formation truncation is associated with local density rather than physical proximity to a cluster, and therefore the physical mechanisms responsible are acting primarily in lower density, group-like environments.  This has been confirmed by studies attempting to link field and cluster galaxies, which find observed cluster populations cannot be generated directly from the accretion of field galaxies.  Galaxy groups, therefore, {\\it must} be playing a significant role as an intermediary stage in galaxy evolution (e.g. Kodama \\& Smail 2001; Fujita 2003).  After initial violent relaxation where galaxies are dominated by a collective potential (Lynden-Bell 1967), the dynamics of group and cluster galaxies are predominantly affected on more local scales. Of the several possible mechanisms driving galaxy evolution, those dependent on the size of the potential well (e.g. harassment; Moore 1996) or density of the intra-cluster medium (e.g. ram-pressure stripping; Gunn \\& Gott 1972) will be most effective in clusters.  Conversely, the low-velocity group environment favours mergers as the preferred method of relaxation due to the growing efficiency of dynamical friction at low relative velocities.  It has been shown that mergers can lead to both morphology and luminosity segregation (Fusco-Femiano \\& Menci 1998; Yepes, Dom\\'inguez-Tenreiro \\& Del Pozo-Sanz 1991), and observations suggest that galaxies are segregated in both groups (Mahdavi \\etal 1999; Girardi \\etal 2003) and clusters (Adami, Biviano \\& Mazure 1998; Biviano \\etal 2002; Lares, Lambas \\& S\\'anchez 2004).  Since mergers will also affect the relative number of galaxies at a given luminosity, one might expect the prevalence of mergers in groups to influence the shape of the group luminosity function.  This is, in fact, a possible explanation of the commonly observed  ``dip'', indicative of intermediate mass galaxies merging to form more luminous ones (e.g. Trentham \\& Tully 2002; Miles \\etal 2004).  It is also worth noting that recent observations of ram-pressure stripping in groups (e.g. Bureau \\& Carignan 2002; Kantharia \\etal 2005; Rasmussen, Ponman \\& Mulchaey 2006) support the results of numerical simulations which suggest that group-level ram-pressure stripping could also play some role in group galaxy evolution (Hester 2006).  Typical studies of poor groups may classify 10-20 galaxies as group members, the majority of which have had their membership assigned using photometric or morphological criteria.  More robust studies have included only galaxies for which recession velocity measurements are available (e.g. Zabludoff \\& Mulchaey 1998; Carlberg \\etal 2001; Brough \\etal 2006a), however this often leads to the need to stack groups in order to measure their properties due to the low numbers of redshifts typically available. The problem with stacking groups, however, is that while it is then possible to constrain the generalised global properties of groups (e.g. mass distribution, velocity dispersion profile etc.) the individual properties of any single group are washed out.  Here we aim to address this issue by establishing the properties of the NGC\\,5044 group and its galaxies independently.  By studying a single rich group we are able to examine evidence for dynamical segregation, substructure and peculiarities in its dynamical properties that are washed out when stacking multiple groups.  In this paper, we present new deep spectroscopic data that allow us to spectroscopically confirm $\\sim40$ new group members.  With the addition of these velocities, we create a new list of 111 confirmed group members that then allows a comprehensive analysis of the dynamical attributes of the NGC\\,5044 group and its constituent galaxies out to nearly two virial radii.  In future work we will use these data to examine the stellar populations of the group galaxies in relation to their position in the group, HI gas properties, star formation rates and dynamical properties.  This paper is organised as follows: in $\\S$\\ref{n5044} and $\\S$\\ref{data} we describe some general properties of the NGC\\,5044 group and the data set we have assembled.  $\\S$\\ref{fof} describes the method we have used to select group members from our list of potential candidates.  The global group properties are addressed in $\\S$\\ref{global}, and the properties related to individual galaxies are summarised in $\\S$\\ref{galaxy}.  Throughout this paper we assume $H_0$ = 70\\kms\\,Mpc$^{-1}$ where applicable and recession velocities are quoted in terms of c$z$.  We adopt the distance modulus of Tonry \\etal (2001), $(m-M)_0=32.31$ (28.99\\,Mpc), measured using surface brightness fluctuations with corrections applied to adjust for the improved Cepheid distance measurements of Jensen \\etal (2003).  Magnitudes have been corrected for galactic extinction using the dust maps of Schlegel, Finkbeiner \\& Davis (1998).  Following convention we use $R$ to denote 2D, projected radii and $r$ to indicate 3D, deprojected radii. ",
        "conclusions": "\\label{discussion}  We have used multi-object spectroscopy to obtain recession velocity measurements for galaxies in the NGC\\,5044 group.  Combining these new observations with available data from literature we are able to define the NGC\\,5044 group as containing 111 members with $M_B\\leq-13.5$\\,mags, nearly a three-fold increase over previous numbers of confirmed group members.  An analysis of common dynamical indicators such as crossing time, line-of-sight velocity distribution, position of the X-ray peak relative to the centroid suggest that the NGC\\,5044 group is relaxed and virialized despite the observed 150\\kms peculiar velocity of the brightest group galaxy NGC\\,5044.  This conclusion of virialization is consistent with XMM X-ray contours for the group, which are very regular and undisturbed.  We note, however, that none of our tests for virialization are sensitive to effects along the line-of-sight, and so the true dynamical state of the group remains somewhat unknown.  Taking the above indicators at face value, however, the group's virialization suggests that the NGC\\,5044 group has experienced no {\\it major} sub-group mergers in several crossing times ($\\sim$1\\,Gyr).  While the dynamical indicators discussed above will give hints as to the timescale of major merger activity, it is likely that low-mass sub-group mergers will not significantly disrupt the virialization of the system as determined via these indicators.  By computing the Dressler-Shectman $\\Delta$ statistic, we have been able to visually and statistically search for space-velocity substructure in the group.  In doing so, we find evidence for a low mass substructure $\\sim$1.4\\,Mpc from NGC\\,5044 group centre.  Two body interactions are expected to lead to dynamical and luminosity segregation in mature groups such as NGC\\,5044 (e.g.  Fusco-Femiano \\& Menci 1998; Lares \\etal 2004).  In examining the NGC\\,5044 group's galaxy population however, we find that galaxies are primarily segregated with respect to their morphologies, and there is no {\\it strong} evidence for segregation in either dynamics or luminosity.  While luminosity segregation is an expected outcome of mergers in groups (e.g. Fusco-Femiano \\& Menci 1998; Lares \\etal 2004), Ludlow \\etal (2007) have used N-body simulations to show that it is possible to eject sub-halos from a group, via multi-body interactions, resulting in significant numbers of associated sub-halos\\footnote{Ludlow \\etal (2007) define associated halos as those that have, at some point, passed within the virial radius of the central halo} residing in the outskirts of the group (as far as 5 times the virial radius).  In relaxed groups such as NGC\\,5044, galaxies have had significant time for two-body interactions to take place, and so the ejection of galaxies from the group centre is more likely to wash out segregation trends.  In addition, this complicates the common interpretation of group-centric radius as a tracer of accretion history as it is no longer clear that galaxies in the outskirts are actually the most recent galaxies to have been accreted into the group potential.  Early-type galaxies in the NGC\\,5044 group are well described by a linear $B-K$ colour-magnitude relation, and are consistent with previous interpretations for the slope in the CMR due to decreasing metallicity or age at lower galaxy masses.  Correlations of $B-K$ colour with local density seem to further suggest that the majority of these bluer, fainter galaxies are residing in the outskirts of the group, where densities are less than $2-3$\\,galaxies Mpc$^{-2}$.  This is consistent with findings for groups and cluster outskirts in the 2dFGRS and SDSS surveys (e.g. Lewis \\etal 2002; Gomez \\etal 2003), however our low sample size prohibits a more quantitative analysis of this effect.  Future work will include a consideration of stellar populations in the galaxies of the NGC\\,5044 group, measured using a combination of our AAOmega and 6dFGS spectra.  Combining the dynamical and kinematic data of the confirmed group members presented here with age and metallicity measurements will allow us to approach the analysis of this group and its galaxy population in significantly more detail than previously attainable.  In particular, the expectation of pre-processing in groups implies a necessary chemodynamical distinction between galaxies associated with the group and field galaxies entering the group environment for the first time."
    },
    "0806/0806.2910_arXiv.txt": {
        "abstract": "We present our discovery of a narrow-line Baldwin effect, an anti-correlation between the equivalent width (EW) of a line and the flux of the associated continuum, in 5-20$\\mu$m mid-infrared lines from a sample of 68 Active Galactic Nuclei (AGN), located at z$<$0.5, observed with the Infrared Spectrograph on the {\\it Spitzer Space Telescope}. Our analysis reveals a clear anti-correlation between the EW of the [SIV] 10.51$\\mu$m, [NeII] 12.81$\\mu$m, and [NeIII] 15.56$\\mu$m lines and their mid-IR continuum luminosities, while the Baldwin effect for [NeV] 14.32$\\mu$m is not as obvious. We suggest that this anti-correlation is driven by the central AGN, not circumnuclear star formation in the host galaxy and present a new method of analyzing this effect in mid-IR lines. We also find that the slope of the narrow-line Baldwin effect in the mid-infrared does not appear to steepen with increasing ionization potential. Examining the dependence of the EW to the Eddington Ratio ($L/L_{Edd}$) we find no strong relationship for mid-IR lines. Our study indicates that the narrow-line mid-infrared Baldwin Effect is quite different from the broad-line optical/UV Baldwin effect and it is possible that the two effects are unrelated. The discovered anti-correlations open new possibilities in the understanding the physics of the ionizing region and the continuum reprocessing by dust. ",
        "introduction": "The Baldwin Effect, first discovered by Baldwin (1977), reports the decrease of the equivalent width (EW) of the broad CIV1549$\\AA$ line with increasing UV luminosity in active galactic nuclei.  The relationship was initially established hoping quasars could be used as potential standard candles in observational cosmology. Extended examination of the relationship over the past decades for both quasars and Seyferts demonstrated that the relationship is not caused by selection effects, but its cosmological use is limited due to large scatter (see review of Osmer \\& Shields 1999 and Kinney et al. 1990; Wilkes et al. 1999; Green, Forster \\& Kuraszkiewicz 2001; Croom et al. 2002; Dietrich et al. 2002; Kuraszkiewicz et al. 2002; Shang et al. 2003).  Significant correlations also exist between the continuum emission and the EW in other UV and optical emission lines including Ly$\\alpha$, H$\\beta$, CIV, CIII, Ly$\\beta$, OIV, OI, CII, AlIII, CIII, MgII, and SiIV+OIV. It was also found that the slope of these relationships appear to increase with increasing ionization potential. In addition, an X-ray Baldwin effect has also been reported in Fe K$\\alpha$ (Iwasawa \\& Taniguchi 1993; Nandra et al. 1997; Page et al. 2004).  The physical origin of this effect is still not clear. A plausible explanation is the softening of the ionizing continuum shape with increasing {\\it L}, which would lead to weaker emission lines compared to the local continuum (Baskin \\& Laor 2004). However, this has been challenged by Wilkes et al. (1999), who found no correlation between any of the UV and optical lines with the X-ray luminosity or X-ray slope. They suggest a model in which limb darkening and the projected surface area of an optically thick, geometrically thin disk combine to cause the Baldwin effect.  Some have also argued that the Eddington ratio $L/L_{Edd}$, a tracer of AGN accretion, may drive the Baldwin effect (Boroson \\& Green 1992). The Baldwin effect may then just be a secondary correlation induced by the tendency of more luminous AGN to have a higher $L/L_{Edd}$ (Baskin \\& Laor 2004; Shang et al. 2003). Others have argued that the fundamental driver is the mass of the supermassive black hole $M_{BH}$ instead of $L$ or $L/L_{Edd}$ (Warner et al. 2003, 2008).  One further plausible explanation is that metallicity of the gas in the AGN affects the EW of the lines in a way that would generate a Baldwin Effect (Deitrich et al 2002).  Most of the discussion of the Baldwin Effect has focused on broad-lines, but a few papers have also noticed a narrow-line Baldwin effect (Green et al. 2001, Croom et al. 2002, Boroson and Green 1992). Since the narrow-line region in quasars may extend to kpc scales, the physics related to the narrow-line Baldwin effect may be different from those driving the broad-line Baldwin effect (Osmer and Shields 1999), and may simply be due to the covering factor of the narrow-line region (Page et al. 2004). This is manifest in a ``disappearing NLR'' model, where the NLR size is related to the AGN luminosity and highly luminous AGN would have weak or even non-existent NLR (Croom et al. 2002).  Complications arise when moving into the mid-infrared. Several authors have noted that in many cases, the IR spectra of AGN do not reflect their optical or UV classifications as reprocessing of the ionizing radiation by the intervening dust as well as circumnuclear star formation activity affects the mid-IR spectral features (Lutz et al. 1998, Laurent et al. 2000, Armus et al. 2007, and Spoon et al. 2007). As a result, quantifying accurately the AGN contribution to the infrared or bolometric luminosity of dust enshrouded galaxies is still largely unanswered problem (see Charmandaris 2008 for a review). Ascertaining the extent of AGN domination in the mid-IR presents unique challenges to determining the extent of a Baldwin Effect. We embarked on the detailed study of the effect using Spitzer data (Keremedjiev \\& Hao 2006) while an analysis using ground based observations with the VLT/VISIR and comparing to x-ray luminosity has been presented by H\\\"{o}nig et al. (2008).  In this paper, we report on our discovery of a narrow-line Baldwin effect in the mid-IR based on Spitzer observations of a large sample which consists of 68 optically classified AGN. Our observations and data analysis are presented in Section~2, our results and detected correlations are shown in Section~3, while the implications of those are discussed in Section~4. ",
        "conclusions": "In the original Baldwin effect, the EW of the lines was anti-correlated with the strength of the UV continuum emission, and the latter was used as tracer of the AGN intensity. This technique is not as direct in the mid-IR because the local continuum may not reflect AGN intensity and could introduce some scatter in the correlations we present in Fig. 1. It is well known that in the infrared, the spectral signatures of an AGN can be severely blended by emission originating from circumnuclear starbursts (see Laurent et al. 2000, Armus et al. 2007 and references therein). This is only partially due to limited spatial resolution of infrared telescopes and focal plane arrays.  More importantly it is the intervening dust which fully reprocesses the intrinsic radiation from the sources - both massive stars and/or an accretion disk - and re-emits it in the infrared. This may lead to a difference between the optical and mid-IR classification of a source.  To address this issue a number of diagnostic methods have been developed of the years in order to obtain a robust AGN/starburst classification of mid-IR spectra (see Lutz et al. 1998, Laurent et al. 2000, Armus et al. 2007, Spoon et al. 2007, Charmandaris 2008, Nardini et al. 2008).  It is important for our analysis to explore whether the observed anti-correlation we find is indeed related to the strength of the AGN. Ideally one could address this by examining the correlation between the line EW and the strength of the X-ray emission (see Hoenig et al. 2008). However, these data are not available for our sample so instead, one can examine how strongly the AGN contributes to the mid-IR continuum emission near the lines versus how much is contaminated by circumnuclear star formation.  It is widely believed that [NeV] 14.32 $\\mu$m, due to its high ionization potential, originates in the NLR (see Gorjian et al. 2007 and references there in). So one would have to examine whether most of the emission from [NeII], [NeIII], and [SIV] in AGN also come from the same region. Gorjian et al (2007) found that both the [NeV] 14.32 $\\mu$m and [NeIII] 15.6$\\mu$m are strongly correlated and deduced that the two must be produced in the same region. Testing the correlation between [NeV] and [NeIII] in our sample confirms this results at $>5\\sigma$ confidence level. We also examined whether [SIV] and [NeII] correlate with [NeV] and found that as with [NeIII], both possessed $>5\\sigma$ relationship with very little scatter. Altogether, this suggests that the four lines investigated in this paper are likely to arise from the same region.  To estimate the contribution of massive star formation in the circumnuclear regions of the AGN sampled by the IRS slits, we used two diagnostics which have been proposed by Genzel et al. (1998) and have also been applied by Armus et al. (2007) in a study of local ultralumninous infrared galaxies. The first diagnostic was to use the [NeV]/[NeII] ratio. This ratio is useful because it is very difficult for stellar sources to produce a significant number of photons capable of ionizing [NeV]. As a result a high [NeV]/[NeII] line ratio indicates the presence an AGN which is dominant in the mid-IR (Lutz et al. 1998).  A second diagnostic is to use the EW of the 6.2$\\mu m$ emission feature. When this feature is strong, it indicates the presence of Polycyclic Aromatic Hysdrocarbons (PAHs) which are mainly produced in photo-dissociation regions (PDRs) when they are excited by the adjacent star forming regions.  Therefore, in Figure 3 we plot the EW$(6.2\\mu m)$ versus [NeV]/[NeII] similar to Figure 5 in Armus et al. (2007). From Figure 3 we see that there are eight galaxies (NGC 7469, IC 5135, NGC 5135, Mrk1066, NGC 1275, NGC 2110, NGC 2273, and NGC 7213) that have both a low [NeV]/[NeII] value and strong PAH emission. As mentioned before, we flagged them as starburst galaxies and did not use them in our correlation analysis.  \\begin{figure} % \\begin{center} \\resizebox{\\hsize}{!}{\\includegraphics[angle=90]{f3.eps}}  \\caption{The EW of the 6.2$\\mu m$ PAH feature versus [NeV]/[NeII]. The red diamond indicates 3C273, a well known quasar, where an AGN dominates its mid-IR spectrum (Hao et al. 2005). The blue open triangles denote objects with both high PAH emission and a low neon ratio. These are the strongest candidate starburst galaxies and were flagged as such in our sample. \\label{fig3}}  \\end{center} \\end{figure}  Even after removing the most flagrant starburst contaminants, it is evident that Fig. 1 still possesses a scatter. This is largely due to the fact that star formation activity is not a binary phenomenon where it is either overwhelmingly dominant or completely dormant. The Armus et al. (2007) and Genzel et al. (1998) figures show that the diagnostics used are continuous between the two extremes and demonstrate that while there are cases where star formation or the AGN activity dominate, in many galaxies the signature of both in the mid-IR is of similar strength. Therefore, even though we removed the galaxies where we were certain starburst activity was dominating, the remaining classified as AGN still possess some scatter in our anti-correlations. This is very likely one of the reasons why our [NeV] result is not strong.  It is evident from the figures that the two diagnostics used are continuous between the two extremes and demonstrate that while there are cases where star formation or the AGN activity dominate, in many galaxies the signature of both in the mid-IR is of similar strength. Therefore, even though we removed the galaxies where we were certain starburst activity was dominating, the remaining classified as AGN still possess some scatter in our anti-correlations. This is very likely one of the reasons why our [NeV] result is not strong.  To reduce the possible contribution of star formation to the observed correlations further yet, we removed the local continuum from the analysis altogether. Laurent et al. (2000) and Nardini et al. (2008) point out that the 5.5$\\mu m$ continuum is dominated by emission from dust located near the AGN torus which heated to near sublimation temperatures, with minimal contribution from star formation activity or stellar photospheric emission. Since the flux of a mid-IR line is independent of local continuum levels, the ratio of the line flux to the 5.5$\\mu m$ continuum versus 5.5$\\mu m$ luminosity should mitigate continuum contamination altogether.  The result of this experiment is clear. The overall anti-correlations are better than those of Fig. 1.  All four lines display $>3\\sigma$ significance and [NeII] and [NeIII] actually have $>4\\sigma$ significance. The correlation values are -0.58 for [SIV], -0.65 for [NeII], -0.57 for [NeV], and -0.62 for [NeIII], with null hypothesis values of $1.03\\times10^{-4}$, $1.29\\times10^{-6}$, $6.41\\times10^{-5}$, and $3.16\\times10^{-6}$ respectively. Using this method, the slope values are $-0.32\\pm0.06$ for [SIV], $-0.36\\pm0.05$ for [NeII], $-0.31\\pm0.06$ for [NeV], and $-0.37\\pm0.06$ for [NeIII]. As expected the starburst galaxies show some of the strongest [NeII] emission in our sample and are obvious outliers. The figure further suggests that the anti-correlation is present in all four lines and is driven by the central AGN. The relationship probably only weakened by star formation in the host galaxy. These results are more meaningful than the ones from Fig. 1 principally because of the inclusion of the [NeV] line which, as mentioned before, we expect most to exhibit a Baldwin effect due to its close relation to the central AGN. Therefore, its inclusion in this analysis points to the AGN as the central driver of this effect. The analysis used in Fig. 4 is shown to be a robust way to measure the decrease in line strength with increasing AGN power and may be a better diagnostic tool than the traditional analysis used in Fig. 1 for mid-IR lines.  \\begin{figure*} % \\begin{center} \\resizebox{\\hsize}{!}{\\includegraphics[angle=90]{f4.eps}} \\caption{Plot of line flux divided by 5.5$\\mu m$ flux versus 5.5$\\mu m$ Luminosity for [SIV], [NeII], [NeV], and [NeIII]. The green open circles denote Seyfert 2 galaxies while the red filled circles denote Seyfert 1 galaxies. Starburst galaxies are plotted as open blue triangles. Overplotted are lines representing the best least-squares fits. \\label{fig4}} \\end{center} \\end{figure*}  Altogether, our data support the possibility that the driver of the narrow-line Baldwin effect is a dynamic covering factor of the narrow line region. If this factor changes as a function of AGN luminosity where more luminous AGN tend to drive the NLR outward, then it would explain the decrease in relative lines strength with increasing AGN power seen in Fig. 1 and 4."
    },
    "0806/0806.0581_arXiv.txt": {
        "abstract": "{The analysis of unresolved stellar populations demands evolutionary synthesis models with realistic physical ingredients and extended wavelength coverage.} {To obtain a quantitative description of the first CO bandhead at 2.3~$\\mu$m, to allow stellar population models to provide improved predictions in this wavelength range.} {We have observed a new stellar library with a better coverage of the stellar atmospheric parameter space than preceding works. We have performed a detailed analysis of the robustness of previous CO index definitions with spectral resolution, wavelength calibration, signal-to-noise ratio, and flux calibration.} {We define a new line-strength index for the first CO bandhead at 2.3~$\\mu$m, D$_{\\rm CO}$, better suited for stellar population studies than previous index definitions. We compute empirical fitting functions for the CO feature as a function of the stellar parameters (T$_{\\rm eff}$, $\\log g$ and [Fe/H]), showing a detailed quantitative metallicity dependence. } {} ",
        "introduction": "One of the most important challenges in modern astrophysics is the proper understanding of the stellar content of unresolved systems, such as extragalactic globular clusters and galaxies in different environments. Since the pioneering work of \\citet{1961MNRAS.122...27C} and \\citet{Tinsley72,Tinsley78,Tinsley80}, this has been accomplished through the comparison of the photometric and spectroscopic data with so-called evolutionary stellar population synthesis models, which make use of theoretical isochrones and libraries of spectral energy distributions (SEDs), either theoretical, empirical or mixed \\citep[for more recent models see e.g.][]{2003MNRAS.340.1317V,2003MNRAS.344.1000B,Maraston05}. The most powerful approach to achieve this goal is to compare a number of observed line-strengths indices with their model predictions, providing in this way constraints to the relevant physical properties of the systems, namely age, metallicity, initial mass function (IMF), and the relative abundance of different chemical species. Since, obviously, the reliability of model predictions improves as more realistic physical ingredients are included, an important effort has been devoted to improve the quality of the SED libraries.  Theoretical libraries usually exhibit systematic discrepancies among themselves and when compared with observational data \\citep[e.g.,][]{1997A&AS..125..229L,1998A&AS..130...65L}. Although the alternative empirical libraries constitute a coarse grained, and usually incomplete (especially for non solar metallicities and non solar abundance ratios) sampling of the space of stellar atmospheric parameters, the use of empirical fitting functions \\citep[e.g.,][]{1993ApJS...86..153G,1999A&AS..139...29G,1994ApJS...94..687W,2002MNRAS.329..863C} can help to reduce these effects \\citep[e.g.,][]{1994ApJS...95..107W,2003MNRAS.340.1317V}.  Up to date, most of the observational effort has been focused to obtain complete libraries in the optical range. However, a full understanding of the physical properties of integrated stellar systems cannot be achieved ignoring other spectral windows. In this sense, the CO features in the K band have been used by many researchers to investigate the stellar content of galaxies, including ellipticals \\citep{1975ApJ...195L..15F,1978ApJ...220...75F,1980ApJ...240..785F,1996MNRAS.280..895M,2000MNRAS.316..507M,1999MNRAS.306..199J,2001MNRAS.326..745M,Fornax_red,Davidge08}, spirals \\citep{1999A&A...350..791J,Bendo04}, compact galaxies \\citep{2007AJ....133..576D,2008ApJ...677..276M}, starbursts and active galactic nuclei \\citep{1994ApJ...421..101D,1994ApJ...428..609R,Shier96,Puxley,Goldader97,Vanzi97,1997ApJ...482L.149M,2000ApJ...545..190I,1999AJ....117..111H,Riffel07}, among others. These strong absorptions are the bandheads formed in the first overtone ($\\Delta \\nu = +2$) bands of CO \\citep{KH86}. Despite the common use of these spectral features for stellar population studies, a proper characterization of the CO bands with stellar atmospheric parameters is still lacking. For that reason, we present in this work an improved study of the infrared region around 2.3~$\\mu$m, where the first bandhead of the strong CO absorptions appear. In particular, we have observed a new library of stars which clearly surpasses preceding works (see \\S~\\ref{subsec-previous-work}) in the coverage of the stellar atmospheric parameters. After a thorough analysis of previous index definitions that have been used to measure the first CO bandhead, we present a new index, D$_{\\rm CO}$, which is well suited for stellar population studies. This new index depends very little on spectral resolution (or velocity dispersion), less sensitive to uncertainties in radial velocities, and can be measured with poorer S/N ratios.  In Section~\\ref{sec-stellar-library} we present the new stellar library, highlighting the improvements over previous libraries, the sample selection as well as an overview of the observations and the data reduction. A detailed discussion of the D$_{\\rm CO}$ index definition is given in Section~\\ref{sec-index-definition}. This section also includes a comparative study of the robustness of the new index to relevant effects. The measurements of the D$_{\\rm CO}$ index for the stellar library, and their associated error estimates appear in Section~\\ref{sec-CO-measurements}. Section~\\ref{sec-atmospheric-parameters} describes the stellar atmospheric parameters used to compute the fitting functions, which are derived in Section~\\ref{sec-fitting-functions}. Finally, Appendix~A includes the tables with all the D$_{\\rm CO}$ measurements for all the stars used for the fitting functions, as well as their stellar atmospheric parameters. ",
        "conclusions": "\\label{sec-summary}  The aim of this work was to obtain an accurate empirical calibration of the behaviour of the CO feature at 2.3~$\\mu$m for individual stars, with the purpose of making it possible to obtain reliable predictions for the CO strength for stellar populations in unresolved systems with a wide range of ages and metallicities. The main results of this work can be summarized as follows:  \\begin{enumerate}  \\item We present a new stellar library in the spectral region around the first CO bandhead at 2.3~$\\mu$m. It consists of 220 stars with stellar atmospheric parameters in the range \\mbox{$2485\\leq T_\\mathrm{eff} \\leq 13404$~K}, \\mbox{$-0.34\\leq \\log g \\leq 5.30$~dex}, \\mbox{$-2.63 \\leq {\\rm [Fe/H]} \\leq 0.98$~dex}.  \\item We define a new line-strength index for the first CO bandhead at 2.3~$\\mu$m, D$_{\\rm CO}$, less sensitive to spectral resolution, wavelength calibration, signal-to-noise ratio and flux calibration than previous definitions.  \\item We compute empirical fitting functions for the D$_{\\rm CO}$ to parametrize the behaviour of the CO feature as a function of the stellar atmospheric parameters. In this work we quantify, for the first time, the metallicity dependence.  \\end{enumerate}  We expect that the work presented in this paper will help researchers to start exploiting in depth the so far poorly-explored and poorly-understood near-IR spectral region centered at 2.3~$\\mu$m, since, as we have shown, the strong CO bandhead can be employed to extract useful physical information of composite stellar populations."
    },
    "0806/0806.0853_arXiv.txt": {
        "abstract": "We study gravitational instabilities in disks, with special attention to the most massive clumps that form because they are expected to be the progenitors of globular-type clusters.  The maximum unstable mass is set by rotation and depends only on the surface density and orbital frequency of the disk. We propose that the formation of massive clusters is related to this largest scale in galaxies not stabilized by rotation. Using data from the literature, we predict that globular-like clusters can form in nuclear starburst disks and protogalactic disks but not in typical spiral galaxies, in agreement with observations. ",
        "introduction": "The study of instabilities in disks has a long history, following the seminal work of Toomre (1964), with a considerable literature on many aspects of it.  However, relatively little attention has been given to the most massive agglomerations that can form by the fragmentation of galactic gas disks.  These most massive agglomerations are of interest because they may be the precursors of the most massive star clusters known, the globular clusters.  Globular clusters were until recently viewed as exclusively old objects, and cluster formation models were therefore based on ideas about early stages of galaxy formation.  This view changed with the realization that elliptical galaxies often contain two populations of globular clusters that appear to have different origins, suggested to be a `primordial' population and a `merger' population (Ashman \\& Zepf 1992).  The existence of a merger population was confirmed by the Hubble Space Telescope with the discovery of massive young clusters in merging galaxies (Ashman \\& Zepf 1998; Schweizer 1998).  A modern theory of globular cluster formation should therefore address why globular cluster formation is common in some environments such as merging and high-redshift galaxies but not in others such as the present Milky Way.  Another question to be addressed is why some regions of galaxies are more favorable for the formation of massive clusters than others.  For example, the inner part of the Milky Way galaxy contains young clusters with masses up to several times $10^4\\,\\msun$, including some quite close to the central black hole (Krabbe et al. 1995; Figer 2008), that are an order of magnitude more massive than the typical open clusters found elsewhere in the galaxy.  The Giant Molecular Clouds (GMCs) in which open clusters currently form in the outer Milky Way have masses of order $10^6\\,\\msun$, but they produce clusters with masses of only several times $10^3\\,\\msun$.  The formation of globular clusters with masses of $10^5 - 10^6\\,\\msun$ therefore requires either much more massive super-GMCs (Harris \\& Pudritz 1994) or a much higher star formation efficiency.  In reality, a combination of these effects is probably involved.  Several scenarios have been proposed for the formation of massive super-GMCs, involving for example a hot primordial plasma (Fall \\& Rees 1985), collisions between normal GMCs (Harris \\& Pudritz 1994), or confinement by high pressures in mergers (Ashman \\& Zepf 2001).  Relatively little attention has been given to the possible role of disks in the formation of globular clusters, with the exception of Larson (1988, 1996), but this possibility now seems worth further investigation given the evidence for massive cluster formation in rotationally supported disks in merging galaxies (Ashman \\& Zepf 1998).  In this $Letter$ we study the gravitational instability of galactic gas disks, with special attention to the most massive clumps that can form in them.  We also consider which galaxies or environments will produce the largest unstable clumps and are therefore most favorable for massive cluster formation.  We start by reviewing the theory of instabilities in disks in \\S 2, and continue with its application to  the formation of globular-type clusters in \\S 3.  We discuss the most promising environments for globular-cluster formation in \\S 3.1, and in \\S 4 we summarize the main implications of our work. ",
        "conclusions": "In this paper we have studied gravitational instabilities in disks, with special attention to the most massive clumps that form because they are expected to be the progenitors of globular-type clusters.  The maximum unstable mass is set by rotation and depends only on the surface density and orbital frequency of the disk, unlike other mass scales such as the Jeans mass that depend on the complex gas physics.  The maximum unstable mass is therefore a well-defined quantity even if the interstellar medium is not well described by a simple equation of state.  The maximum clump mass can be expressed in terms of the total gas mass and the gas fraction in a galaxy, and this formulation makes it clear that environments with a high gas fraction are the most promising places to form massive clusters.  Using data from the literature, we predict that massive globular-like clusters can form in nuclear starburst disks and protogalactic disks but not in typical spiral galaxies, in agreement with observations.  The scenario proposed here relates massive clusters to the largest scale in galaxies not stabilized by rotation, which is the only scale intermediate between stars and galaxies that has a clear physical basis.  There is no well-defined `Jeans mass' on these intermediate scales, and the next smaller scale that has any clear physical basis is the thermal Jeans scale in molecular clouds, which is related to the masses of individual stars.  The next larger physical scale is that of the galaxy itself, so we predict three physically well-motivated scales corresponding to stars, massive stellar  clusters, and galaxies.  Another application of studies of the stability of nuclear gas disks, not discussed here, is to their likely role in feeding supermassive central black holes in galaxies (Escala 2007).  This requires the outward transfer of angular momentum, and if gravity is the most important force involved, the same mass concentrations that form massive clusters may also play an important role in the outward transfer of angular momentum and the growth of a central black hole.  We plan to investigate this problem further in ongoing work."
    },
    "0806/0806.1644_arXiv.txt": {
        "abstract": "We describe a subgrid model for including galaxies into hydrodynamical cosmological simulations of galaxy cluster evolution. Each galaxy construct-- or {\\em galcon}-- is modeled as a physically extended object within which star formation, galactic winds, and ram pressure stripping of gas are modeled analytically. Galcons are initialized at high redshift ($z \\sim 3$) after galaxy dark matter halos have formed but before the cluster has virialized. Each galcon moves self-consistently within the evolving cluster potential and injects mass, metals, and energy into intracluster (IC) gas through a well-resolved spherical interface layer. We have implemented galcons into the {\\em Enzo} adaptive mesh refinement code and carried out a simulation of cluster formation in a $\\Lambda$CDM universe. With our approach, we are able to economically follow the impact of a large number of galaxies on IC gas. We compare the results of the galcon simulation with a second, more standard simulation where star formation and feedback are treated using a popular heuristic prescription. One advantage of the galcon approach is explicit control over the star formation history of cluster galaxies. Using a galactic SFR derived from the cosmic star formation density, we find the galcon simulation produces a lower stellar fraction, a larger gas core radius, a more isothermal temperature profile, and a flatter metallicity gradient than the standard simulation, in better agreement with observations. ",
        "introduction": "\\label{sec:intro}  Current hydrodynamic cosmological simulations of galaxy clusters show an appreciable level of inconsistency with results from high-precision optical and X-ray observations. Discrepancies are particularly apparent in intracluster (IC) gas properties (e.g., Tornatore et al. 2003, Kay et al. 2007, Tornatore et al. 2007; for a recent review, see Borgani et al. 2008) - spatial distributions of density, temperature, and metallicity, but also in the stellar component for which simulations usually over-predict the stellar mass fraction while underpredicting the total number of galaxies (Nagamine et al. 2004 and references therein.). Physical processes, such as galactic winds, ram-pressure stripping, mergers of subclusters, energetic particle heating, and gravitational drag, affect the dynamical and thermal state of IC gas. Several attempts to implement some of these phenomena have been made (e.g., Kapferer et al. 2006, Domainko et al. 2006, Bruggen \\& Ruszkowski 2005, Sijacki \\& Springel 2006, Cora 2006, Kapferer et al. 2007), with some success in reconstructing IC gas properties. However, different combinations of these processes and their specific implementation in numerical codes generally result in quite different gas properties.  Because modeling star formation (SF) self-consistently requires prohibitively high level of spatial resolution, most current simulations use a SF prescription that follows the formation of collisionless star `particles' which feedback mass and energy to IC gas (Cen \\& Ostriker 1992; Nagai \\& Kravtsov 2005). This approach overestimates the SF rate (SFR) at low z (Nagamine et al. 2004), which leads to a higher than expected star to gas mass ratio. In addition, feedback from the star particles is unresolved, leading to unrealistically low levels of gas (including metals) and energy transfer from galaxies into IC gas, and consequently insufficient suppression of cooling and gas overdensity in cluster cores. This unsatisfactory state motivates our attempt to develop a new method that partly overcomes current numerical limitations.  In this {\\em Letter} we briefly describe a new approach (Section 2) for including galaxies in hydrodynamical cosmological simulations of cluster evolution which provides improved control over the relevant physical processes. Galaxies which are otherwise under-resolved (or absent!) are replaced with a physically-extended galaxy subgrid model which we refer to as a {\\em galcon} within which SF, galactic winds, and ram pressure stripping of gas are modeled analytically. Galcons are initialized at high redshift after galaxy dark matter (DM) halos have formed but before the cluster has virialized. Mass, metals, and energy are injected from galcons into IC gas. In Section 3 we compare the results of our galcon simulation with a standard simulation using a popular star formation and feedback recipe, and summarize our main conclusions in Section 4. ",
        "conclusions": "The combinination of our galaxy constructs and new semi-analytic modeling of the relevant physical processes yields a powerful tool that is capable of reproducing the basic properties of clusters. Our new approach successfully describes SF and the basic properties of IC gas, including its metallicity and energy feedback. The ever improving observational data motivate further development of the code and inclusion of additional physical processes previously unaccounted for, such as AGN feedback. We plan to improve the description of galactic mergers, and intend to implement an improved algorithm for replacing galactic halos with new galcons as the cluster evolves, instead of performing this replacement only at an initial redshift as has been done in the simulations reported here. Ongoing work on this project will hopefully lead to a much better understanding of the intrinsic properties of both DM and baryons in clusters."
    },
    "0806/0806.3021.txt": {
        "abstract": "Here we describe the first results of the ALHAMBRA survey which provides {\\it cosmic tomography} of the evolution of the contents of the Universe over most of Cosmic history. Our novel approach employs 20 contiguous, equal-width, medium-band filters covering from 3500 \\AA\\ to 9700 \\AA, plus the standard $JHK_s$ near-infrared bands, to observe a total area of 4 square degrees on the sky. The optical photometric system has been designed to maximize the number of objects with accurate classification by Spectral Energy Distribution type and redshift, and to be sensitive to relatively faint emission features in the spectrum. The observations are being carried out with the Calar Alto 3.5m telescope using the wide field cameras in the optical, LAICA, and in the NIR, Omega-2000. The first data confirm that we are reaching the expected magnitude limits (for a total of 100 ksec integration time per pointing) of AB $\\leq$ 25 mag (for an unresolved object, S/N = 5) in the optical filters from the blue to 8300~\\AA, and from AB = 24.7 to 23.4 for the redder ones. The limit in the NIR, for a total of 15 ks exposure time per pointing, is (in the Vega system) K$_s$ $\\approx$ 20 mag, H $\\approx$ 21 mag, J$\\approx$ 22 mag. Some preliminary results are presented here to illustrate the capabilities of the ongoing survey. We expect to obtain accurate redshift values, $\\Delta z/(1+z) \\leq 0.03$ for about 5 $\\times 10^5$ galaxies with I$\\leq 25$ (60\\% completeness level), and $z_{med}$ = 0.74. This accuracy, together with the homogeneity of the selection function, will allow for the study of the redshift evolution of the large scale structure, the galaxy population and its evolution with redshift, the identification of clusters of galaxies, and many other studies, without the need for any further follow-up. It will also provide targets for detailed studies with 10m-class telescopes. Given its area, spectral coverage and its depth, apart from those main goals, the ALHAMBRA-Survey will also produce valuable data for galactic studies. ",
        "introduction": "%  Only over the last few years it has become possible for Observational Cosmology to gather enough data on the distant universe to feed our comprehension of the evolution of the different objects that populate it. It has become almost commonplace to study protogalaxies at redshifts $z > 5$, and to observe particular objects at redshifts as high as $z \\approx 6.5$ (Becker \\etal 2001, Kashikawa \\etal 2006, Kawai \\etal 2006) or even $z \\approx 7.5$ (Bradley \\etal 2008). At the same time, samples of objects have been collected through different techniques at smaller distances (and shorter evolutionary times) from us, and the different properties of objects in separate redshift ranges have been measured and compared. However, it remains true that to this day, no homogeneous sample of objects has been collected covering a significant range of the age of the universe, even if some remarkable efforts have been devoted to the production of wide-field, shallow surveys, that cover the low-redshift end (like 2dFGRS, Colless \\etal 2001, SDSS, York \\etal 2000, VVDS, Le F\\`evre \\etal 2005 or DEEP-2, Davis \\etal 2003) while other groups have directed their efforts towards the most distant end, through very deep, small-area surveys like the HST Deep Fields or other legacy programs (Ferguson \\etal 2000, Beckwith \\etal 2006).  The Cosmological Principle implies the existence of maximally symmetric subspaces and the existence of a one-to-one relation between redshift and time. The corresponding evolutionary nature of the depicted Universe is a model-independent prediction, prior to any consideration about the value of the cosmological parameters. A direct way to tackle many of the problems posed by modern cosmology is hence to materialize a {\\sl foliation of the space-time}, producing narrow slices in the $z$-direction whereas the spatial sections are large enough to be cosmologically representative, obtaining as output a kind of {\\em Cosmic Tomography}.  From the observational point of view, to trace {\\em Cosmic Evolution}, which is a central topic in Cosmology, the genuine evolutionary effects have to be disentangled from both the physical variance at a given redshift and the details of the metric as measured in --- or, depending on the point of view, induced by --- the cosmological model. In other words, to approach the question of evolution meaningfully it is necessary to sample large volumes even at low redshift, to capture not only representative average properties but also their variance. This will necessarily imply a survey featuring a combination of wide area and depth, and a continuous spectral coverage to avoid complex selection functions that depend on the redshift and on the nature of the objects under analysis. Moreover, the quest for the necessary precision implies high enough spectral resolution and photometric accuracy.  Up to now, the largest surveys ensuring complete spectral coverage for large samples have been photometric, and done with broad-band filters. The resulting redshift precision obtained with these techniques ($\\sim 0.03$ in $\\Delta z/(1+z)$, at best, see Cucciati \\etal 2006 and Ilbert \\etal 2006) and in Spectral Energy Distribution (SED) determination are correspondingly rough. Moreover, large area photometric surveys like SDSS are necessarily shallow, whereas deeper surveys have sampled the distant and/or faint Universe in rather small areas. At the other extreme in spectral resolution, spectroscopic surveys can neither go as deep as the photometric ones nor cover large enough areas. Moreover, they are defined in order to observe a restricted spectral region, producing a selection effect that is a function of the object type and redshift that can be very intricate due to the selection effects inherent to spectroscopy (Fern\\'andez-Soto \\etal 2001).  For those scientific purposes where the detailed properties of individual objects are not the goal, the aim from an observational stand is therefore that of finding the optimal filter combination to produce the most homogeneous, deepest, and most accurate possible photometric survey. Such a survey would produce precise enough values for the redshift and SED for large numbers of objects. We present here the \\textbf{A}dvanced \\textbf{L}arge \\textbf{H}omogeneous \\textbf{A}rea \\textbf{M}edium \\textbf{B}and \\textbf{R}edshift \\textbf{A}stronomical, \\textbf{ALHAMBRA}-Survey, that intends to produce such an optimum survey for the study of cosmic evolution. It has been designed to achieve (with the facilities at hand) the best compromise between large area and depth, good spectral resolution and coverage, in order to produce an optimum output in terms of redshift and SED accuracy. The ALHAMBRA-Survey is a deep photometric survey using 20 contiguous, equal-width, medium-band optical filters from 3500 \\AA\\ to 9700 \\AA, plus the three standard broad band ($JHK_s$) NIR filters. The total area surveyed by ALHAMBRA will be 4 square degrees, being therefore placed halfway in between traditional imaging and spectroscopic surveys.  By design, the ALHAMBRA-Survey will provide precise ($\\Delta z < 0.03(1+z)$) photometric redshifts and SED classification for several hundred thousand galaxies and AGNs, allowing for different kinds of analysis regarding populations and structures, and their evolution in time.  The details of the project, including simulations and expected results, and all the related aspects are described in the ALHAMBRA-Book that can be found at \\url{http://www.iaa.es/alhambra}. Thanks to the unbiased nature of this survey (i.e. neither designed to detect a given class of objects nor to be precise only in some fixed spectral window), important problems other than cosmic evolution can be addressed. These include the study of stellar populations in the galactic halo, the search for peculiar stellar objects, ranging from very cold stars to blue stragglers, and the possible detection of debris from galactic satellites in the Milky Way halo. Moreover, the large surveyed volume and the ability to finely discriminate between different spectral energy distributions will permit the serendipitous detection of objects that could be classified as {\\em exotic} or {\\em rare}. This broad category includes very high redshift galaxies ($\\approx$ 2500 objects at $z > 5$, with $\\Delta z < 0.01$, expected from scaled HDF observations) and QSOs.  The observations are being carried out with the 3.5m telescope of the Centro Astron\\'omico Hispano-Alem\\'an, CAHA, Calar Alto Observatory (Almer\\'{\\i}a, Spain) and the wide-field imagers in the optical (LAICA) and in the NIR (Omega-2000). The collected data render possible the study of many different astronomical problems in a self-contained way and will provide with very interesting targets for individual studies with large size telescopes.  A separate article (Ben\\'{\\i}tez \\etal 2008) deals with the selection of the optical filters and the optimization of their characteristics to maximize the spectral information, while in this work we present the main characteristics that specifically define the ALHAMBRA-Survey, including some preliminary results from the data we have already accumulated.  This paper is organized as follows: in Section 2 we present the project implementation and its present status, and in Section 3 the first, preliminary results. We compare the ALHAMBRA survey with other surveys in Section 4, whereas our conclusions are presented in Section 5. % %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% % ",
        "conclusions": ""
    },
    "0806/0806.0018_arXiv.txt": {
        "abstract": "We measure the number of companions per galaxy ($N_c$) as a function of $r$-band absolute magnitude for both the Sloan Digital Sky Survey and the \\citet{croton06} semi-analytic catalog applied to the Millennium Run simulation. For close pairs with projected separations of 5-20 $h^{-1}$ kpc, velocity differences less than 500 km s$^{-1}$, and luminosity ratios between 1:2 and 2:1, we find good agreement between the observations and simulations, with $N_c$ consistently close to 0.02 over the range $-22 < M_r < -18$.  For larger pair separations, $N_c(M_r)$ instead becomes increasingly steep towards the faint end, implying that luminosity-dependent clustering plays an important role on small scales. Using the simulations to assess and correct for projection effects, we infer that the {\\em real-space} $N_c(M_r)$ for close pairs peaks at about $M^*$, and declines by at least a factor of two as $M_r$ becomes fainter. Conversely, by measuring the number density of close companions, we estimate that at least $90\\%$ of all major mergers occur between galaxies which are fainter than $L^*$.  Finally, measurements of the luminosity density of close companions indicate that $L^*$ galaxies likely dominate in terms of the overall importance of major mergers in the evolution of galaxy populations at low redshift. ",
        "introduction": "Galaxy mergers can produce dramatic changes in the morphological, nuclear and star forming properties of galaxies over relatively short timespans. As a result, mergers have been invoked to explain a number of aspects of galaxy evolution. In recent years, large redshift surveys have paved the way for systematic observational studies of candidate mergers, while theoretical modelling of structure formation has produced key advances in our understanding of the role of merging in a cosmological context. In general, these efforts have focussed on two key aspects of merging: (1) the effects of merging on the constituent galaxies, and (2) the frequency with which merging occurs, as described by the merger rate and related quantities.  An increasingly popular method of identifying candidate merging systems is the use of close galaxy pairs, which are the precursors to mergers. With careful choices of close pair criteria, and correction for projection effects (ie., contamination by non-merging pairs), the frequency of close pairs should correlate with the merger rate. Recent studies using cosmological simulations support this idea, demonstrating that most close pairs merge on relatively short timescales \\citep{kitzbichler08}.  In addition, the properties of paired galaxies can provide insight into the nature of merging galaxies both before and during the encounter. Galaxies in close pairs have higher asymmetries than galaxies in wider separation pairs or the field \\citep{hernandez05,patton05,depropris07}, confirming that interactions and mergers are prevalent in these systems. Star formation is enhanced in close pairs at low redshift \\citep{2dfpairs,2dfpairsb,nikolic04,patton05,alonso06,geller06,barton07,owers07,smith07,woods07,ellison08,li08}, implying that star formation has been triggered by mergers or interactions.  Differences in metallicities between paired and field galaxies are consistent with a scenario in which interactions funnel gas to the central regions of galaxies involved in these close encounters \\citep{kewley06,ellison08}. Most ultraluminous infrared galaxies (ULIRGs) originate from major mergers of gas rich galaxies \\citep{dasyra06a,dasyra06b}, while approximately half of the luminous infrared galaxies (LIRGs) at low redshift appear to be undergoing interactions or mergers \\citep{wang06}.  Using close galaxy pairs and/or galaxy asymmetries as indicators of imminent or recent mergers, the merger rate and its evolution has now been measured using a number of large redshift surveys \\citep{carlberg00,lefevre00,cnoc2mr,conselice03,bundy04,lin04,bell06a,kartaltepe07,kampczyk07,conselice08,hsieh08,lin08,lotz08,rawat08,ryan08}. Evolution estimates range from roughly $(1+z)^{0.5}$ to $(1+z)^3$, implying widely differing scenarios at $z \\sim 1$ and above.  At least some of these discrepancies result from the use of different pair criteria. For example, simulations and semi-analytical models of galaxy formation indicate that the merger rate and its evolution depends on factors such as galaxy mass, pair mass ratio and environment \\citep{khochfar01,berrier06,maller06,cox08,guo08,kitzbichler08}.  In order to better understand the role of merging, and to reconcile merger rate measurements from disparate samples, we need to assess how the frequency and nature of merging depends on factors such as environment, mass ratio (ie., major versus minor mergers), properties of the progenitor galaxies (e.g., dry mergers versus gas-rich mergers), and the mass (or luminosity) of the merging galaxies or merger remnants (e.g., formation of massive galaxies versus $L^*$ galaxies). Significant observational progress has been made in all of these areas in recent years. Galaxy groups appear to be an ideal environment for mergers \\citep{goto05,brough06,miles06,robotham06,zandivarez06,coziol07,nolan07}, though the infall regions of clusters \\citep{vandokkum99,tran05,moss06} and the low density field \\citep{barton07} are important too. Induced star formation appears to be strongest in major mergers, or in the lower luminosity (or mass) members of minor mergers \\citep{woods06,woods07,ellison08}. Dry mergers have been invoked to explain the assembly of massive elliptical galaxies since $z \\sim 1$ \\citep{vandokkum05,bell06b,naab06}, though this process likely cannot explain all recently formed early type galaxies \\citep{cox06,brown07,bundy07,scarlata07}.  In this paper, we investigate the luminosity dependence of the merger rate at low redshift, using close galaxy pairs in the Sloan Digital Sky Survey \\citep{york00}, and in the \\citet{croton06} semi-analytic galaxy catalogs derived from the Millennium Run simulation \\citep{springel05}. Both of these samples are large enough that, in addition to being able to measure close pair statistics as a function of absolute magnitude, we also have the luxury of being very selective in how we choose our close pairs.  In particular, we require all of our pairs to have spectroscopic redshifts for both members, projected separations less than 20 $h^{-1}$ kpc, and relative velocities less than 500 km s$^{-1}$.  In addition, we require our companion sample to be volume limited for all luminosity ratios between 1:2 and 2:1, thereby providing a cleaner match to the major merger candidates we seek to identify. Moreover, our measurements are carried out in the $r$-band, yielding absolute magnitudes that are a better proxy for stellar mass than those at shorter wavelengths.  This is beneficial for close pair studies, since merger-induced star formation is likely to affect the luminosities of galaxies in pairs more than normal (isolated) galaxies.  We begin by describing the creation of our SDSS spectroscopic and photometric samples in \\S~\\ref{data}. Section \\ref{ncsdss} outlines the calculation of the number of close companions per galaxy ($N_c$) for SDSS, including corrections for spectroscopic incompleteness. We make a direct comparison with the Millennium Run simulation in \\S~\\ref{ncmill}, and derive real space pair statistics for both SDSS and Millennium in \\S~\\ref{realspace}.  With additional assumptions, we then relate these close pair statistics to the merger rate in \\S~\\ref{mrate}. We summarize our conclusions  in \\S~\\ref{conclusions}. Throughout this study, we adopt cosmological parameters of $\\Omega_m = 0.3$, $\\Omega_\\Lambda = 0.7$, and $H_0 = 100~h$ km s$^{-1}$ Mpc$^{-1}$.  For brevity, we express all absolute magnitudes as $M_r$ instead of $M_r - 5 \\log (h)$. ",
        "conclusions": "We have measured the number of close companions per galaxy ($N_c$) as a function of absolute magnitude for both the SDSS and the Millennium simulation.  For SDSS, we construct samples of host and companion galaxies, and correct for spectroscopic incompleteness.  For Millennium, we create a suite of mock redshift catalogs, averaging over different views of the \\citet{croton06} cube. Using close pair criteria designed to identify imminent major mergers ($5 < r_p < 20~h^{-1}$ kpc, $\\Delta v < 500$ km s$^{-1}$, and $|\\Delta M_r| < 0.753$), we find general agreement between the observations and simulations. In redshift space, $N_c \\sim 0.02$ over the range $-22 < M_r < -18$. The flatness of this distribution indicates that $N_c(M_r)$ does not simply trace the number density of galaxies; instead, small scale luminosity-dependent clustering appears to counteract this effect.  Using three dimensional positions and velocities available from the Millennium simulations, we measure the degree to which the detected galaxy pairs are contaminated by projection effects, and find that the contamination is a strong function of absolute magnitude, rising from $45\\%$ at $M_r \\lesssim M^*$ to $\\sim 83\\%$ at $M_r = -18$. We remove this contamination from both Millennium and SDSS pair statistics, yielding $N_c(M_r)$ measurements in real space.  These measurements indicate that galaxies with $M_r \\lesssim M^*$ are the most likely to be undergoing major mergers at low redshift. However, by computing the number density of close companions ($n_c(M_r)$) in real space, we conclude that at least 90$\\%$ of all major mergers occur between galaxies which are {\\it fainter} than $M^*$. With additional assumptions, we also estimate the galaxy and volume merger rates, which trace the real-space $N_c(M_r)$ and $n_c(M_r)$ respectively.  We estimate that at least 8$\\%$ of $L^*$ galaxies are likely to have undergone a major merger since $z=1$, while this remnant fraction appears to be $\\sim 4$ times smaller for 0.1 $L^*$ galaxies.  Finally, our measurements of the luminosity density of close companions indicate a clear peak at $M^*$, implying that $L^*$ galaxies dominate in terms of the overall amount of stellar mass involved in major mergers at low redshift.  Together, these results indicate that the low redshift merger rate depends strongly on luminosity (and presumably mass). This has a number of important implications. For example, one would not expect the merger rates of massive galaxies (e.g., Masjedi et al. (2006)\\nocite{masjedi06}) to agree with those of $L^*$ galaxies (e.g., De Propris et al. 2007\\nocite{depropris07}). Also, the increase in projection effects for fainter galaxies indicates that samples of luminous galaxy pairs are more likely to provide bona fide merger candidates than samples of lower luminosity pairs. This is relevant if one wishes to assess the impact of merging on the constituent galaxies (e.g., enhanced star formation or disturbed morphologies). We also conclude that at low redshift, recent merging history is likely to be most important for galaxies which are relatively luminous. Finally, given the clear peak in $l_c(M_r)$, it appears that galaxies which are much more or much less luminous than $L^*$ are unlikely to play an important role in the overall evolution of galaxies via major mergers."
    },
    "0806/0806.2471_arXiv.txt": {
        "abstract": "We empirically determine effective temperatures and bolometric luminosities for a large sample of nearby M dwarfs, for which high accuracy optical and infrared photometry is available. We introduce a new technique which exploits the flux ratio in different bands as a proxy of both effective temperature and metallicity. Our temperature scale for late type dwarfs extends well below $3000\\,\\rm{K}$ (almost to the brown dwarf limit) and is supported by interferometric angular diameter measurements above $3000\\,\\rm{K}$. Our metallicities are in excellent agreement (usually within 0.2~dex) with recent determinations via independent techniques. A subsample of cool M dwarfs with metallicity estimates based on hotter \\emph{Hipparcos} common proper--motion companions indicates our metallicities are also reliable below $3000\\,\\rm{K}$, a temperature range unexplored until now. The high quality of our data allow us to identify a striking feature in the bolometric luminosity versus temperature plane, around the transition from K to M dwarfs. We have compared our sample of stars with theoretical models and conclude that this transition is due to an increase in the radii of the M dwarfs, a feature which is not reproduced by theoretical models. ",
        "introduction": "Low--mass dwarfs are the dominant stellar component of the Galaxy and have been employed in a variety of Galactic studies: tracing Galactic disk kinematics (e.g. Hawley, Gizis \\& Reid 1996; Gizis et al. 2002; Bochanski et al. 2005, 2007), studying the stellar age--velocity relations (West et al. 2006), investigating Galactic structure (e.g. Reid et al. 1997, Kerber et al. 2001, Pirzkal et al. 2005), and the Galaxy's mass and luminosity (e.g. Hawkins \\& Bessell 1988; Kirkpatrick et al. 1994; Gould, Bahcall \\& Flynn 1997, 1998; Zheng et al. 2001, 2004). An increasing number of M dwarfs are now known to host exoplanets (e.g. Delfosse et al. 1998; Butler et al. 2004; Bonfils et al. 2007, Udry et al. 2007). The determination of accurate fundamental parameters for M dwarfs has therefore relevant implications for both stellar and Galactic astronomy.  Observationally, the spectra of these stars are marked by the presence of strong molecular absorption features, in either optical (e.g. TiO and VO) and infrared regions (e.g. $\\textrm{H}_2\\textrm{O}$ and CO). Molecular lines blend with all other lines and create a pseudo--continuum, rendering all spectral analysis difficult (e.g.\\ Gustafsson 1989). However, recent advances in model atmospheres of low--mass dwarf stars (Hauschildt et al. 2003, Brott \\& Hauschildt 2005), have boosted the number of studies deriving accurate metallicities for M dwarfs (Woolf \\& Wallerstein 2005, 2006; Bean et al. 2006a, 2006b). Modeling the internal structure, atmospheric properties and magnetic activity of M dwarfs (e.g. Burrows et al. 1993, Allard et al. 1997, Baraffe et al. 1998, Hauschildt, Allard \\& Baron 1999, Allred et al. 2006, Reiners \\& Basri 2007) also presents ongoing theoretical challenges. For a small number of nearby M dwarfs, interferometry is currently providing direct angular diameter measurements (Lane, Boden \\& Kulkarni 2001; S\\'egransan et al. 2003; Berger et al. 2006) which confirm a large discrepancy between the predicted and observed radii, as has been noted in eclipsing binaries with M--type components (see Ribas 2006 for a review).  In this paper, we empirically determine the effective temperatures and the bolometric luminosities for more than 340 M dwarfs. This work is an extension of our previous study on G and K dwarfs, to which we applied the InfraRed Flux Method (IRFM, Casagrande, Portinari \\& Flynn 2006).  The effective temperature and the bolometric luminosity scales we derive are accurate at a level of a few percent and are supported by interferometric angular diameters above $\\sim 3000\\,\\rm{K}$.  We find in this study that below about $4000\\,\\rm{K}$ the monochromatic to bolometric flux ratio in different bands is a proxy of both effective temperature and metallicity. By exploiting this feature, we are able to derive not only $\\Teff$, but also the metallicities of our M dwarfs, which are found to be in very good agreement (usually within $0.2-0.3$~dex) with those inferred using the Bonfils et al. (2005) calibration or directly measured from Woolf \\& Wallerstein (2005, 2006). The technique we propose also looks promising for stars much cooler than those explored in the aforementioned studies.  We find considerable structure in the temperature--luminosity plane, especially around the transition between K and M dwarfs. Our study circumstantially confirms previous works which indicate that the radii of M dwarfs are larger by about 15\\% than model predictions. We also find strong evidence that this discrepancy, clearly observed in M dwarfs in eclipsing binary systems, is also present in nearby disk M dwarfs.  The paper is organized as follows. In Section \\ref{sample} we describe our sample of M dwarfs and in Section \\ref{phoenix} we compare it with the Phoenix model atmosphere in the two--colour plane. We then review the IRFM and present our new technique for estimating effective temperatures and metallicities below $4000\\,\\rm{K}$ in Section \\ref{mfm}. Our proposed metallicity, bolometric luminosity and effective temperature scales along with the comparison with other empirical determinations are discussed in Section \\ref{abs_cor}, \\ref{luca} and \\ref{tesca}, respectively. In Section \\ref{jump} we analyze the stars with good \\emph{Hipparcos} parallaxes in the HR diagram. We find a strong feature which marks the transition from K to M dwarfs and which is due to an increase in the observed stellar radii. We briefly discuss possible reasons, including the effect of magnetic fields and molecular opacity. We finally conclude in Section \\ref{conclu}. ",
        "conclusions": "We have determined the temperature scale of M dwarfs, using stars with very accurate multi--band photometry from optical to near-IR and the MOITE, a new method which exploits the flux ratio in different bands as a sensitive indicator of both effective temperatures and metallicities. Our proposed temperature scale extends down to $\\Teff \\sim 2100-2200\\,\\rm{K}$ i.e. to the L dwarf limit (e.g.  Leggett et al.  2002) and above $\\sim 3000\\,\\rm{K}$ is supported from interferometric angular diameters.  Our metallicities, which are ultimately calibrated on Bonfils et al.'s (2005) metallicity scale, are also found to be in very good agreement with the latest measurements from Woolf \\& Wallerstein (2005, 2006) and Bean et al. (2006a, 2006b), even if significant differences in the various effective temperature scales still exist. Cool M dwarfs with metallicities based on (hotter) \\emph{Hipparcos} common proper--motion companions also suggest our metallicities are reliable even below $3000\\,\\rm{K}$, although further data are needed. Accurate multi--band photometry for the coolest \\emph{Hipparcos} common proper--motion pairs would permit one to firmly extend the MOITE to the bottom of the main sequence, thus opening this elusive area also to galactic chemical evolution investigations. Exoplanets are found around M dwarfs, and a uniform metallicity scale for their host stars will also be very useful.  The high quality of our data allows us to identify a striking feature which marks the transition from K to M dwarfs, which appears to be due to an increase in the radii of the early M dwarfs relative to late K dwarfs. We have compared our sample of stars with theoretical isochrones for low mass stars and find that such a feature is not predicted by the models, substantially confirming the disagreement already noticed in the case of eclipsing binaries. Possible explanations including the effect of magnetic fields and molecular opacity have been discussed.  This work also highlight the potentiality of high accuracy multi--band photometry in determining fundamental stellar parameters and identifying fine details in the HR diagram.  In particular, the MOITE will hugely benefit from the existing infrared (2MASS, DENIS) and forthcoming optical surveys like SkyMapper (Keller et al. 2007), Pan-Starrs (Kaiser et al. 2002) and LSST (Claver et al. 2004) which will provide accurate and homogeneous multi--colour and multi--epoch photometry for a large number of stars."
    },
    "0806/0806.0183_arXiv.txt": {
        "abstract": "% The Hobby-Eberly Telescope Dark Energy Experiment (HETDEX) will outfit the 10 m HET with a new wide field and an array of 150 integral-field spectrographs to survey a 420 deg$^2$ area in the north Galactic cap. Each fiber-coupled unit spectrograph will cover 350-550 nm, simultaneously. This instrument, called VIRUS, will produce $\\sim$34,000 spectra per exposure, and will open up the emission-line universe to large surveys for the first time. The survey will detect 0.8 million Lyman-alpha emitting (LAE) galaxies with 1.9$<$$z$$<$3.5 and more than a million [OII] emitting galaxies with $z$$<$0.5.  The 3-D map of LAE galaxies in 9 cubic Gpc volume will be used to measure the expansion history at this early epoch using baryonic acoustic oscillations and the shape of the power spectrum. The aim of HETDEX is to provide a direct detection of dark energy at z$\\sim$3. The measurement will constrain the evolution of dark energy and will also provide 0.1\\%-level accuracy on the curvature of the Universe, ten times better than current.  The prototype of the VIRUS unit spectrograph (VIRUS-P) is a powerful instrument in its own right.  Used on the McDonald 2.7~m, it covers the largest area of any integral field spectrograph, and reaches wavelengths down to 340 nm. VIRUS-P is being used for a pilot survey to better measure the properties of LAE galaxies in support of HETDEX.  We report initial results from this survey. ",
        "introduction": "Progress in understanding the physical nature of dark energy will require precision measurements of the expansion history of the Universe over the redshift range 0$<$$z$$<$4. In order  to make progress towards this goal, very significant surveys involving new facilities are required. Two approaches are being pursued: to refine the accuracy of the measurement at low redshift, including constraints on possible evolution of the dark energy equation of state, or to constrain dark energy evolution, directly, through observations at high redshift.  HETDEX \\Citep{h2004b} has the goal of providing percent-level constraints on the expansion history of the universe (the Hubble parameter $H(z)$ and angular diameter distance $D_A(z)$) over redshifts $z$=1.9 to 3.5. HETDEX will use a combination of baryonic acoustic oscillations \\Citep[e.g.][]{se2007,ksg2007}) and power spectrum shape information to provide at least a 3-$\\sigma$ direct detection of dark energy over these redshifts, even in the event that dark energy is a cosmological constant. To achieve this, HETDEX needs an accuracy of 0.9\\% on $H(z)$ at $z$=2.8. This level of accuracy requires a volume of 9 Gpc$^3$ with a density of tracers $\\sim$$10^{-4}$ objects per Mpc$^3$, which can be achieved by surveying 420 deg$^2$ over 1.9$<$$z$$<$3.5 with 0.8 million Lyman-$\\alpha$ emitting (LAE) galaxies. LAEs have high number density and are easily detected with integral field spectroscopy (as shown in Sec. 3). In addition to a direct detection of dark energy, a 0.9\\% measurement of $D_A$ at $z$=$3$ from HETDEX will determine curvature to 0.1--0.2\\% \\Citep{k2006} a factor of ten better than currently known. Nearly all of the other dark energy missions require some knowledge of curvature to disentangle the dark energy contribution. This is because at low redshifts, dark energy and curvature provide the same expansion signature. ",
        "conclusions": ""
    },
    "0806/0806.0710_arXiv.txt": {
        "abstract": "We report a large-scale coronal wave (so-called ``EIT wave\") observed with high cadence by EUVI onboard STEREO in association with the GOES B9.5 flare and double CME event on 19 May 2007. The EUVI instruments provide us with the unprecedented opportunity to study the {\\it dynamics} of flare/CME associated coronal waves. % The coronal wave under study reveals deceleration, indicative of a freely propagating MHD wave. Complementary analysis of the associated flare and erupting filament/CME hint at wave initiation by the CME expanding flanks, which drive the wave only over a limited distance. The associated flare is very weak and occurs too late to account for the wave initiation. ",
        "introduction": "Large-scale large-amplitude waves and shocks in the solar corona occur in association with flares and coronal mass ejections (CMEs). The existence of flare-related global disturbances has been first inferred from Moreton waves \\citep{moreton60}, which appear as arc-like fronts in chromospheric H$\\alpha$ filtergrams, moving away from the ignition site with typical velocities of 500--1000~km/s. It was soon recognized that Moreton waves could not be propagating in the chromosphere, where no wave mode has such high velocity (e.g.\\ sound speed and Alfv\\'en speed are only of the order of tens of km/s). The first interpretation was by \\cite{uchida68} that Moreton waves are the surface-track of a coronal fast-mode magnetohydrodynamic (MHD) wave front.  The Extreme-ultraviolet (EUV) Imaging Telescope \\citep[EIT;][]{dela95} onboard Solar and Heliospheric Observatory (SOHO) for the first time directly imaged propagating global disturbances in the corona, and these so-called ``EIT waves'' were assumed to be the coronal counterparts of the Moreton waves \\citep{thompson98,thompson99}. Thereafter, coronal waves were found to be a quite frequent phenomenon, and it became an intense matter of debate whether EIT waves: \\begin{itemize} \\item[a)] are really the coronal counterparts of Moreton waves \\citep[e.g.][]{thompson00,klassen00,warmuth01,warmuth04a,eto02,khan02,narukage02,vrsnak02,gilbert04,veronig06}; \\item[b)] are caused by the flare explosive energy release or by the erupting CME \\citep[e.g.][]{warmuth01,warmuth04b,biesecker02,hudson03,zhukov04,cliver05,vrsnak06}; \\item[b)] are waves at all (and, if yes, which type of waves; cf.\\ \\citeauthor{wills07} \\citeyear{wills07}) or rather propagating disturbances related to magnetic field line opening and restructuring associated with the CME lift-off \\cite[e.g.][]{delannee99,wills99,delannee00,wang00,wu01,warmuth01,chen02,vrsnak02,ballai05,attrill07}. \\end{itemize} In addition, there might be different types of EIT waves, further complicating this debate. For detailed discussions we refer to the recent reviews by \\cite{chen05,vrsnak05,mann07,warmuth07}.  One important limitation of coronal wave studies so far is the low cadence of the EIT instrument (12--15~min), which makes it impossible to study wave kinematics beyond a rough velocity estimate. Observations of large-scale waves in TRACE EUV images are rare due to its limited field of view, but we note that one such event was observed with high cadence and studied in detail in \\cite{wills99}. The Extreme Ultraviolet Imagers (EUVI) on the recent Solar Terrestrial Relations Observatory (STEREO) spacecraft regularly perform EUV full-disk imaging with a cadence as good as 2.5~min. In this letter, we study for the first time the dynamical evolution of a globally propagating ``EIT\" wave in high-cadence EUVI images. ",
        "conclusions": "In Figure~\\ref{euvi_summary}, we show a summary plot comprising: a) the distance-time diagram of the coronal wave observed by EUVI/STEREO-A, b) the back-extrapolated (quadratic fit) distance-time diagram of CME1 observed with LASCO/SOHO, c) the distance-time diagram of CME2 observed with COR1/STEREO-A, d) the flare hard X-ray flux recorded by RHESSI, and e) the flare soft X-ray flux recorded by GOES. From the quadratic fit to the EUVI wave kinematics, we estimate the wave's launch time to $\\approx$12:45~UT. The real launch may happen somewhat later, since this method assumes a point-like origin of the wave. The flare 12--25~keV hard X-ray flux starts rising at 12:50~UT with the first and highest peak at 12:51:30~UT. At this time,  we already observe the first EUVI wave front. Such timing argues against a flare-origin of the wave, since the wave needs time to build up a large amplitude or shock to be observable. On the other hand, timing and direction of the erupting filaments indicate that the wave was closely associated with the fast CME1, since filament1 disappeared from the H$\\alpha$ filter at 12:46~UT, whereas filament2 remained visible until 12:55~UT.  However, the kinematics of the coronal wave is quite different from the kinematics of the CME's leading edge  (see Fig.~\\ref{euvi_summary}): the wave is slower than both CMEs and significantly decelerates, which is a typical characteristics of a large-amplitude MHD simple wave \\citep{mann95,vrsnak00}: such a freely propagating perturbation is powered only temporarily by a source region expansion, which could be due to the flare-related pressure pulse, due to small scale flare ejecta or due to the CME expanding flanks (which propagate laterally only over a limited distance). Since perturbation elements with larger amplitude travel faster than those with smaller amplitude (nonlinearity), the perturbation profile steepens until finally a discontinuity is formed. As a consequence of energy conservation, the amplitude of the perturbation decreases with distance, first due to spherical expansion ($\\sim$$R^{-2}$), and second, because the crest of the shock travels faster than its trail, causing broading of the perturbation profile  \\citep{landau87}. Consequently, the wave decays into an ordinary, i.e.\\ small amplitude wave propagating with the characteristic speed of the medium. For waves propagating perpendicular to the magnetic field, this is the magnetosonic speed $v_{\\rm ms} = (v_A^2 + c_s^2)^{1/2}$ with $v_A$ the Alfv\\'en velocity and $c_s$ the sound speed. The ``final\" velocity reached by the EUVI wave under study lies in the range $200\\pm 50$~km s$^{-1}$ (see Fig.~\\ref{euvi_kinematics}), which is a reasonable value of $v_{\\rm ms}$ in the quiet solar corona \\citep[e.g.,][]{mann99}, though we note that there is an ongoing discussion on this subject \\cite[e.g.][]{wang00,wu01,chen02,wills07}.  The observed EUVI wave deceleration together with the closely related timing of the wave and the erupting filament1/CME1 (in contrast to the flare peak which occurs too late) as well as the wave front shape which is roughly concentric with filament1, hint at an initiation of the wave by the CME expanding flanks. In such a scenario, the wave is only driven over a limited distance and then decays into an ordinary MHD wave."
    },
    "0806/0806.4844_arXiv.txt": {
        "abstract": "{ This methodological paper is part of a short series dedicated to the long-standing astronomical problem of de-projecting the bi-dimensional, apparent morphology of a three-dimensional distribution  of gas.} {We focus on the quantification and spatial recovery of turbulent motions in planetary nebulae (and other classes of expanding nebulae) by means of long-slit echellograms over a wide spectral range.} { We introduce some basic theoretical notions, discuss the observational methodology, and develop an accurate procedure disentangling all broadening components (instrumental resolution, thermal motions, turbulence, gradient of the expansion velocity, and fine structure of hydrogen-like ions) of the velocity profile in all spatial positions of each spectral image. This allows us to extract random, non-thermal motions at unprecedented accuracy, and to map them in 1-, 2- and 3-dimensions.} { We discuss general and specific applications of the method.  We  present the solution to practical problems in the multi-dimensional turbulence-analysis of a testing-planetary nebula (NGC 7009), using the three-step procedure (spatio-kinematics, tomography, and 3-D rendering) developed at the Astronomical Observatory of Padua (Italy). In addition, we introduce an observational paradigm valid for all spectroscopic parameters in all classes of expanding nebulae.} {Unsteady, chaotic motions at a local scale constitute a fundamental (although elusive) kinematical parameter of each planetary nebula, providing deep insights on its different shaping agents and mechanisms, and on their mutual interaction. The detailed study of turbulence, its stratification within a target and (possible) systematic variation among different sub-classes of planetary nebulae deserve long-slit, multi-position angle, wide-spectral range echellograms containing emissions at low-, medium-, and high-ionization,  to be analyzed pixel-to-pixel with a straightforward and versatile methodology, extracting all the physical information (flux, kinematics, electron temperature and density, ionic and chemical abundances, etc.) stored in each  frame at best.} ",
        "introduction": "The presence of circumstellar and interstellar gas in expansion is the typical signature of instability phases in stellar evolution, characterized by a large, prolonged  mass-loss rate (planetary and  symbiotic star nebulae, shells around Wolf-Rayet and luminous blue variable [LBV] stars), or even  explosive events (nova and supernova remnants). The mass-loss of evolved stars occupies a strategic ground between stellar and interstellar physics by raising fundamental astrophysical problems (e. g., origin and structure of winds, formation and evolution of dust, synthesis of complex molecules), by playing a decisive role in the final stages of stellar evolution, and because it is crucial for the galactic enrichment in light and heavy elements.  In particular, the late evolution of low-to-intermediate mass stars (1.0 M$_{\\odot}$$\\le$M$_*$$\\le$8.0 M$_{\\odot}$) is marked by the planetary nebula (PN) metamorphosis: an asymptotic giant branch (AGB) star gently pushes out surface layers in a florilegium of forms, crosses the HR diagram, and reaches the white dwarf regime.  \\footnote{The foregoing sentences are taken from \\cite{tur02} (2002), and \\cite{sab06} (2006), of whom the present paper is the ideal sequel.}  According to current hydrodynamical and radiation-hydrodynamical simulations (\\cite{mel94} 1994; \\cite{mar01} 2001; \\cite{vil02} 2002; \\cite{per04a} 2004a, 2004b; \\cite{scho05} 2005), the PN evolution  is driven by its central star through (a) the AGB mass-loss history; and (b)  the variation of ionizing flux and wind power in the post-AGB phase (without excluding the possible contribution of other factors, like magnetic fields and binarity of the central star; \\cite{sok98} 1998; \\cite{gar99} 1999; \\cite{mas99} 1999; \\cite{fra99} 1999).  At present, it is still unclear the relative importance of different shaping agents, their mutual influence, and how they affect the (global and small-scale) structure and kinematics of the ejected gas. Leaving simplistic ``modern views'' of PN evolution out of consideration, we remark that updated 1-D radiation-hydrodynamical evolutionary models (\\cite{per04a} 2004b; \\cite{scho05} 2005) still partly fail  in reproducing the positive expansion velocity gradient of the main shell (the Wilson's law; \\cite{wil50} 1950;  \\cite{wed68} 1968), and the gas distribution observed in a representative sample of  carefully de-projected ``true'' PNe (\\cite{sab06} 2006 and references therein). Such model-PN vs. true-PN discrepancies suggest that the former tend to systematically overestimate the relative importance of wind interaction (as supported by a recent revision downward of mass-loss rates in PN central stars, due to clumping of the fast stellar wind; \\cite{kud06} 2006; \\cite{pri07} 2007).  An even deeper uncertainty surrounds the turbulence of expanding gas, a fundamental parameter measuring local, random, unsteady motions due to non-linear energy transfer. On the one hand, 2-D hydrodynamical simulations of wind interaction (\\cite{fem94} 1994; \\cite{vis94} 1994; \\cite{mel97} 1997; \\cite{dwa98} 1998; \\cite{gar99} 1999, 2006) predict stratified turbulent motions peaking within shocked regions  and/or ionization fronts characterized by the growth of Kelvin-Helmholtz, Vishniac  or Rayleigh-Taylor instabilities. On the other hand, practical determinations based on high-resolution spectra (\\cite{gue98} 1998; \\cite{ack02} 2002; \\cite{gaz03} 2003, 2006; \\cite{gz07} 2003, 2007; and \\cite{med06} 2006) provide a unique, average turbulence-value for each nebula, ranging from 0 km s$^{-1}$ to 20 km s$^{-1}$.  We aim at quantifying and mapping turbulent motions in PNe by means of long-slit, wide spectral range, very-high spectral and spatial resolution echellograms at several position angles (PA), reduced and analyzed according to spatio-kinematical, tomographic and 3-D recovery methods developed at the Astronomical Observatory of Padua (\\cite{sab00} 2000, 2006, and references therein).  This introductory, methodological  paper (a) dissects all  broadening agents of a velocity profile; (b) describes an original procedure mapping turbulent motions in 1-, 2-, and 3-dimensions; (c) illustrates the solution to practical problems of turbulence-determination in a testing-PN; and (d) introduces an observational paradigm valid for all types of expanding nebulae.  In a forthcoming paper, we (i) will present multi-dimensional turbulence maps for a representative sample of PNe in different evolutionary phases; and (ii) will compare these observational results with expectations from hydrodynamical simulations and theoretical models. ",
        "conclusions": "So far, the quantification of stratified turbulent motions in the ionized gas of PNe - as predicted by hydrodynamical simulations and theoretical models - escaped a direct verification because of its strict connection with the general, long-standing astronomical problem of de-projecting the bi-dimensional, apparent morphology of a three-dimensional mass of gas. The solution is now approached by spatio-kinematical, tomographic, and 3-D rendering analyses applied to long-slit, multi-PA, high-resolution  echellograms over a wide spectral range.  We develop an accurate procedure disentangling all broadening components (instrumental resolution, thermal motions, turbulence, expansion gradient, and fine structure of hydrogen-like ions) of the velocity profile. This overcomes  main  weaknesses affecting  previous  determinations, and allows us to quantify and map in 1-, 2-, and 3-dimensions the gas-turbulence in PNe.  The multi-dimensional turbulence-analysis - soon applied to a representative sample of targets in different evolutionary phases, and combined with expectations from updated radiation-hydrodynamical simulations -  will provide new, deep insights on different shaping agents of PNe (ionization, wind interaction, magnetic fields, binarity of the central star, etc.), and on their mutual interaction.  Last, we emphasize all advantages - and encourage the adoption - of long-slit  echellograms covering the entire instrumental spectral range, since the usual insertion of an interference filter isolating a single echelle order proves superfluous and strongly limitative in most cases... it is like driving a Ferrari with the hand-brake applied."
    },
    "0806/0806.1781.txt": {
        "abstract": " ",
        "introduction": "Why was there a Big Bang? Why is the universe not featureless and barren? Why are there fluctuations in the cosmic microwave background? In stark contrast to the convincingly answered ``what\" questions of cosmology (e.g. What is the age of the universe?, What is the geometry of the universe?), these ``why\" questions may instead evoke a sense of disillusionment. Is it possible that cosmology's ``triumphs\"---its answers to the ``what\" questions---are frustratingly inadequate, or worse, incomplete?  However, what if the ``why\" questions provide tantalizing hints of the ultimate origins of the universe? Then instead of crisis, we encounter an amazing opportunity---one that might provide answers to the most enigmatic question of all: How did the universe begin?  Inflation \\cite{guth81} is a daring paradigm with the promise to solve many of these mysteries. It has entered its third decade of successfully confronting observational evidence and emerged as cosmology's theoretical touchstone. Despite its many successes, inflation remains unproven. While skeptics must resort to increasingly finely tuned attacks \\cite{khoury01,magueijo03}, inflation's proponents can only cite circumstantial evidence in its favor \\cite{turner02}. However, a conclusive detection of a primordial gravitational wave background (GWB) from inflation would be ``the smoking gun\" \\cite{kamkos1998}. No other known cosmological mechanism mimics the GWB's imprint on the cosmic microwave background (CMB).  New technological innovations poise cosmology at the threshold of an exhilarating era---one in which future CMB data will winnow down the seemingly boundless ``zoo\" of cosmological models and test the hypothesis that an inflationary expansion of the universe took place in its first moments.  Inflation's unique imprint on CMB polarization has generated considerable attention from US science policy advisors \\cite{nrc2003,nasmckee,nasurry,hepap}, who have all enthusiastically recommended measuring CMB polarization. The reason for this excitement is clear: inflation explains a host of critical cosmological observations, and CMB polarization is the most promising, and perhaps only, way to glimpse the GWB.  This chapter describes how the Cosmic Gravitational Wave Background induces a specific type of CMB polarization and describes the first experiment dedicated to testing this most-promising signature of inflation. This experiment, the Background Imaging of Cosmic Extragalactic Polarization (BICEP) project, has recently embarked on its third observing season. We show preliminary data from the BICEP's first season obtained with a novel polarization modulation mechanism called the ``Faraday Rotation Modulator\". Our discussion ends with a description of exciting new technology with the potential to probe inflation down to the ultimate cosmological limit. ",
        "conclusions": ""
    },
    "0806/0806.4769_arXiv.txt": {
        "abstract": "Previous analyses of magnetospheric accretion and outflow in classical T Tauri stars (CTTSs), within the context of both the X-wind model and other theoretical scenarios, have assumed a dipolar geometry for the stellar magnetic field if it were not perturbed by the presence of an accreting, electrically conducting disk.  However, CTTS surveys reveal that accretion hot spots cover a small fraction of the stellar surface, and that the net field polarization on the stellar surface is small.  Both facts imply that the magnetic field generated by the star has a complex non-dipolar structure.  To address this discrepancy between theory and observations, we re-examine X-wind theory without the dipole constraint. Using simple physical arguments based on the concept of trapped flux, we show that a dipole configuration is in fact not essential.  Independent of the precise geometry of the stellar magnetosphere, the requirement for a certain level of trapped flux predicts a definite relationship among various CTTS observables.   Moreover, superposition of multipole stellar fields naturally yield small observed hot-spot covering fractions and small net surface polarizations.  The generalized X-wind picture remains viable under these conditions, with the outflow from a small annulus near the inner disk edge little affected by the modified geometry, but with inflow highly dependent on the details of how the emergent stellar flux is linked and trapped by the inner disk regions.  Our model is consistent with data, including recent spectropolarimetric measurements of the hot spot sizes and field strengths in V2129 Oph and BP Tau. ",
        "introduction": "Various theoretical models have been proposed for the physical mechanisms driving the accretion and outflow processes in classical T Tauri stars (CTTSs), with the most popular being perhaps Blandford and Payne's (1982) pioneering study of magnetocentrifugally driven winds from Keplerian disks. However, the X-wind model has gained credence in recent years for a variety of theoretical and observational reasons (see also the discussion in \\S 5).  It is instructive at the outset to summarize and compare the different CTTS models, to review the rationale for X-wind theory and the evidence in its favor, and to motivate our subsequent generalization of this picture (see also Shu et al. 2000).  The first accretion models of CTTSs proposed that they were surrounded by Keplerian disks, which extend all the way to the stellar surface, with accretion occuring through an equatorial boundary layer (e.g., Bertout 1987; Kenyon \\& Hartmann 1987; Bertout et al. 1988).   Motivated by the observational finding that the driver behind some well-known bipolar outflow sources are neutral winds containing H I and CO (Lizano et al. 1988), a combination not seen in the interstellar medium but present in the photospheres of cool stars, Shu et al. (1988) suggested that the driving outflows are caused by boundary-layer disk-accretion onto a strongly magnetized young stellar object (YSO).  The accretion provides a rationale for why a protostar might rotate near break-up, and Hartmann \\& MacGregor (1982) had already shown that magnetized stars rotating near break-up could shed matter and angular momentum extremely efficiently through their equatorial zones.  However, Shu et al. (1988) obtained preliminary indications that streamline collimation in this kind of model was absent, or extremely slow, and therefore they speculated stellar jets were associated with ``ordinary'' stellar winds confined to flow along the rotation axis by the more powerful ``extraordinary'' centrifugally-driven outflow that they later called an X-wind.  When a disk abuts a fully convective star, Ekman pumping (internal circulation caused by slight pressure differences in a field of differential rotation) to the equatorial boundary layer adjoining the star and the disk eventually causes the entire star to spin near break-up if there are no countervailing spin-down torques (Galli 1990; see Fig. 10 of Shu et al. 1993).  CTTSs are fully convective, yet they usually rotate at velocities one order of magnitude below break-up (e.g., Vogel \\& Kuhi 1981; Hartmann \\& Stauffer 1989; Bouvier et al. 1993, 2007).  Shu et al. (1988) therefore speculated that CTTSs are young stars in which the spin-down by ordinary stellar winds overcame the spin-up toques of the viscous boundary-layer, and that CTTSs would not have strongly collimated outflows.  The latter expectation turned out to be false (see, e.g., Edwards et al. 1993).  Nevertheless, Matt \\& Pudritz (2005, 2008a, b) have recently resurrected the idea that magnetized stellar winds might provide sufficient torque to explain the slow rotation of T Tauri stars and, perhaps, be responsible for a part of the observed optical jets from YSOs.  The latter proposal using coronal gas to launch the stellar winds has, however, foundered on the constraints provided by X-ray observations of T Tauri stars (Bisnovatyi-Kogan \\& Lamzin 1977, DeCampli 1981, Matt \\& Pudritz 2008c), and the finding that CTTS stellar winds appear to be an order of magnitude cooler (Johns-Krull \\& Herczeg 2007) than suggested recently (Dupree et al. 2005). The issues, therefore, became {\\it (1)} how to accrete high angular momentum disk material onto a CTTS while keeping the latter rotating slowly, {\\it (2)} how simultaneously to generate winds efficiently from the star+disk system, and {\\it (3)} how to make the outflowing wind appear as a jet despite the slow streamline collimation?  K\\\"{o}nigl (1991) suggested a solution for the accretion part of this problem by adopting the theory developed by Ghosh \\& Lamb (1978; 1979a, b).  A strong stellar magnetosphere is assumed to truncate the disk some distance from the stellar surface, with accreting material flowing onto the star not through a boundary layer but via magnetic field lines threading the disk.  The angular momentum of the star is then regulated by the interaction between the stellar magnetosphere and the disk.  Field lines originating in the star and with disk footpoints within the Keplerian corotation radius ($\\rx$) are dragged forward by the disk gas, relative to the star, and thus tend to spin the star up (at the expense of the disk gas).  Stellar field lines threading the slower rotating disk beyond $\\rx$, on the other hand, tend to spin the star down.  The slow rotation of CTTS then arise if the spin-down caused by the outer disk outweighs the spin-up by the inner parts. Large turbulent diffusivities must prevail to allow magnetic fields to slip through the gas in such a picture.  In a series of papers with a precedent in the suggestions of Arons (1986) and Camenzind (1990), Shu and collaborators (Shu et al. 1994a, b; Najita \\& Shu 1994; Ostriker \\& Shu 1995, hereafter OS95; Shu et al. 1995) embraced the idea that strong YSO magnetospheres might truncate the disk before it abuts the stellar surface, but they pointed out a number of difficulties with the specific proposal of K\\\"onigl (1991).  The resolution of these difficulties turned out to provide solutions for each of the issues {\\it (1)} to {\\it (3)} listed above.  First, the picture painted by Ghosh \\& Lamb and K\\\"onigl changes considerably if turbulent resistivities are not large, but small, i.e., if field diffusion is competitive with advective flow not on dynamical time scales but secular ones.  In such cases, the long-term processes of angular-monetum transport outwards and mass-transport inwards dominate over notions of ``ram-pressure balance''.  Second, stellar fields strong enough to truncate an electrically conducting disk are automatically also strong enough to drive a magnetocentrifugal outflow along the outermost flux tubes of the stellar field, which are opened into an X-wind.  The magnetic torques in the wind cause the outflowing matter to gain angular momentum at the expense of the material still connected to it by field lines threading through the disk.  The back reactions to the X-wind and funnel-flow give a pinch of the exterior field lines inward and the interior field lines outward towards a common mid-point $\\rx$, {\\it with a net trapped flux}, that gives the X-wind model its name. This mid-point is both in Keplerian rotation and in corotation with the star, i.e., $\\Omega(R_X) = (GM_*/R_X^3)^{1/2} = \\Omega_*$, a condition that came to be called {\\it disk locking}. Third, isodensity contours become cylindrically stratified very quickly after the gas accelerates from the X-region (Najita \\& Shu 1994; Shu et al. 1995), yielding the optical illusion in the emission of forbidden-line and radio emission that X-winds achieve jet-like collimation close to the base of the flow (Shang et al. 2002, 2004).  In actual practice, streamline collimation is logarithmically slow, which has observable consequences for position-velocity diagrams obtained from long-slit spectrograms (e.g., Pyo et al. 2006).  Many numerical simulations have also been devoted to the X-wind/funnel flow problem, the most successful being that of Romanova et al. (2008).  Indeed, the progress made by the simulations is most concisely revealed by examining why they succeeded in obtaining the simultaneous existence of X-winds and funnel flows over extended durations when others failed.  As stated in the previous paragraph and detailed in \\S 2, the crucial concept in X-wind/funnel flow theory is that of the {\\it trapped flux} created by the two-sided pinching of field lines toward $\\rx$.  In the presence of non-zero resistivity $\\eta$, field diffusion will occur out of the X-region.  This diffusion must be offset by fluid advection from both the exterior and the interior of $\\rx$.  In the exterior, disk inflow is induced both from back-reaction from the X-wind (if present) and by viscous inflow from the disk proper, with the level of kinematic viscosity $\\nu$ dictating the disk accretion-rate $\\dot M_D$. If the viscosity $\\nu$ is only comparable to the resistivity $\\eta$, the disk inflow is too weak relative to diffusive penetration to produce a good inward pinch, a failure not conducive to the generation of X-winds.  To produce sufficiently outward-bending field lines from viscous accretion, it is necessary to have $\\eta$ smaller than $\\nu$ (by a factor of roughly the disk-aspect ratio $z_0/\\varpi \\ll 1$; see Lubow, Papalozou, \\& Pringle 1994 and Shu et al. 2007).  For the specific problem of an accretion disk interacting with a stellar magnetosphere, the simulations of Romanova et al. (2008) demonstrate explicitly that the condition $\\eta \\ll \\nu$ is indeed the crucial ingredient to achieving an X-type magnetic configuration and thereby generating an X-wind.  In the interior, a funnel flow involving closed field lines links the star to the disk.  If the stellar rotation rate $\\Omega_*$ is not chosen to be the same as the angular speed $\\sim \\Omega_X$ of the parts of the disk to which the filed lines are rooted, rapid transients are induced.  Some early simulations, which started with slowly rotating stellar magnetospheres linked to rapidily rotating disks in the initial state, managed to obtain temporary X-type magnetic configurations via such transients (Hayashi et al. 1996, Miller \\& Stone 1997, K\\\"uker et al. 2003).  The sudden removal of angular momentum from the disk by magnetic torques creates a dynamic onrush of material toward the star, generating a temporary magnetic pinch and outflow.  In these studies, finite resistivity is sometimes included in the code, but this resistivity plays no important role because the magnetic diffusion is slow compared to the fast inward flow caused by the large disequilibrium of the initial state.  Goodson et al. (1997, 1999), Goodson \\& Winglee (1999), and Romanova et al. (2002, 2003, 2004a, 2005, 2008; hereafter R02, R03, R04, R05, R08) pioneered the incorporation of more realistically rotating stellar magnetospheres, with the later papers of the latter group and Goodson et al. (1997, 1999) using controlled levels of resistivity and viscosity. In particular, R02 and R03 performed axisymmetric and non-axisymmetric calculations that start with an unperturbed, aligned and tilted, stellar dipole field threading a circumstellar disk that begins to accrete slowly via a postulated disk viscosity $\\nu$.  The only resistivity in the problem is numerical and gives an effective value for  $\\eta$ comparable to $\\nu$.  Although no X-winds arose as a consequence, these authors did show that in steady-state the radius $R_{\\rm co}$ where the magnetosphere corotates with the disk is close to the disk truncation radius (denoted by them as the stellar magnetopause $R_{\\rm m}$), with the best runs having $R_{\\rm co}/R_{\\rm m}$ = 1.2--1.3 (see also Long et al. 2005).  Using the same basic configuration, but rotating the star more quickly so that corotation is reached interior to the truncation radius -- the so-called ``propeller regime'' (Illarionov \\& Sunyaev 1975) -- R05 found strong disk outflows to be possible if $\\nu$ is several times larger than $\\eta$ (see also Ustyugova et al. 2006; hereafter U06).  The simulations of R04 discovered ``magnetic towers'' in the stellar corona near the rotation axis as first proposed by Draine (1983) and Lynden-Bell (1996).  With realistic levels of coronal density, such towers do not carry much matter and cannot explain YSO jets (Long et al. 2005; see also Figs. 4 \\& 5 of Allen et al. 2003b for the appearance of the \"magnetic tower\" phenomenon in the context of the collapse of a rotating, magnetized, molecular-cloud core).  Of specific interest to our present paper, the lowering of the coronal densities assumed by R04 by two orders of magnitude in the simulations of Long et al. (2005) made much easier the opening of stellar field lines, and led the way to the driving of powerful YSO jets in an X-type magnetic configuration by R08 when $\\eta \\ll \\nu$ and $R_{\\rm co} \\approx R_{\\rm m} \\approx R_X$.  The main remaining differences between ideal X-wind theory and numerical simulations concerns whether outflows can be steady.  Figure 1 shows that field lines dead to inflow or outflow (black) separate the funnel flow (red) from the X-wind (blue).  Some of the dead-zone field lines are closed and link the disk to the star; some are open and ``joined at infinity'' line by line to the magnetic field contained in the X-wind.  The reversal of field direction across the separatrix between the open dead-zone fields and the open X-wind fields of ideal MHD would be unstable to reconnection events in the presence of finite resistivity.  These reconnection events, which are likely to be episodic (Aly \\& Kuijpers 1990), would create a fluctuating X-wind (Shu et al. 1997; Romonova et al. 1998; Uzdensky, K\\\"{o}nigl, \\& Litwin 2002), and may underlie the outbursting behavior found in the simulations of Goodson et al. (1997, 1999), R04, R05, U06, and R08.  The steady assumption of ideal X-wind theory is then made for analytic simplicity, and can at best represent only the time-average of outflows that are time-variable and quasi-periodic in reality. ",
        "conclusions": "Since its inception more than a decade ago, generalized X-wind theory has made many successful predictions and survived several  key confrontations with observations (e.g., Johns \\& Basri 1995; JG02, Shang et al. 2004;  McKeegan 2006; Pyo et al. 2006; Zolensky et al. 2006; Carr 2008; Edwards 2008). Claims of large rotation rates seen in YSO jets have been used to argue against X-winds (e.g., Bacciotti et al. 2002, Coffey et al. 2004).  These claims turn out to have difficulties, e.g., indications in some cases that the inferred rotation in the jet is counter to that of the disk, or the absence of any obvious rotation in jets best oriented (i.e., in the plane of the sky) to show such rotation (e.g., Cabrit et al. 2006, Pety et al. 2006, Coffey et al. 2007, Lee et al. 2006, 2007).  Present upper limits on the rotation in the latter cases show that if any disk wind is present in the observations, they must have small launch radii that make them look very similar to X-winds (see Cai et al. 2008 for a fuller discussion).  In a similar manner, the finding that the magnetic fields on the surfaces of CTTSs are locally strong as seen in hot spots, but globally weak when averaged for the photospheric polarization signal (Valenti \\& Johns-Krull 2004), turns out not to be an argument against disk-locking via funnel flows (Matt \\& Pudritz 2005), but evidence for the phenomenon of flux trapping predicted by X-wind theory (OS95) and quantified in this paper (see also JG02).  In hindsight, we should perhaps not be surprised that fully convective CTTSs have surface distributions of magnetic field that are non-dipolar.  Such configurations complicate the theoretical modeling of funnel flows, but they break the degeneracy in the very different explanations offered by Ghosh \\& Lamb (1978, 1979a, b; see also K\\\"onigl 1991) and by Cameron \\& Campbell (1993), Shu et al. (1994a), OS95, and this paper for the standoff distance $R_X$ of the stellar magnetopause where the disk is truncated at its inner edge and to whose Keplerian rotation the star is locked in steady state.  In particular, this paper shows that the observed field strengths in funnel-flow hot spots and their fractional covering fractions $F_h$ have a rational explanation in the context of generalized X-wind theory, while they are inexplicable in other semi-analytic formulations of the steady-state problem. The strong concentrations of magnetic flux available to realistic configurations of surface fields on young stars may mean that it is harder to crush their magnetopsheres than has been estimated by naive dipole estimates.  While the results from numerical simulations (e.g., Goodson et al. 1999, K\\\"uker et al. 2003, Long et al. 2005) contain many of the elements seen in the semi-analytic theory, complete agreement has not yet been achieved, perhaps because the numerical simulations invariably contain, so far, too much resistive diffusion of the field, relative to angular momentum transport by turbulent viscosity, to give a good semblance of the phenomenon of trapped flux (for progress in this regard, see Romonova et al. 2008, p. 281).  Nevertheless, the trends are promising, and at some stage, the awesome computing power of modern machines will be able to reliably extend the solutions to the non-axisymmetric and time-dependent regimes occupied by actual systems that are inaccessible to semi-analytic techniques."
    },
    "0806/0806.4243_arXiv.txt": {
        "abstract": "{In this review, I will discuss how to characterize synchrotron X-ray variability of TeV blazars by using the observed/simulated light curves. Apparently, temporal studies provide independent and complementary information to the spectral studies, but surprisingly little attention has been paid especially for the blazar study. Only exception is a classical argument for presence of ``time lag'', which may (or may not) reflect the diffrence of synchrotron cooling timescale. Also very recently, it was suggested that the X-ray variability of TeV blazars indicates a strong red-noise, compared to a fractal, flickering-noise of Seyfert galaxies. Various temporal techniques are proposed in literature, e.g., the power spectrum density (PSD), the structure function (SF),  and the discrete correlation function (DCF) and other analysis tools, but special care must be taken if the data are not well sampled and observation is relatively short compared to a characteristic timescale of the system. Also, the situation is being more complicated for low-Earth orbit satellites, e.g., $ASCA$, $RXTE$ and $BeppoSAX$, since the light curve inevitably contains ``periodic gap'' due to the Earth occultation (every $\\simeq$ 6ksec). I will present detailed approaches to see how the \"gap\" and the \"finite length\" of the data affects the results of temporal analysis, and to what extent we can believe in our results.  Finally, I will briefly comment on the high-sensitivity X-ray observations with $MAXI$, that may shed new light on the forthcoming $GLAST$ era.}  \\FullConference{Workshop on Blazar Variability across the Electromagnetic Spectrum\\\\ April 22-25 2008\\\\ Palaiseau, France}  \\begin{document} ",
        "introduction": "Blazars are commonly variable from radio to $\\gamma$-rays. The variability timescale is shortened and the radiation is strongly enhanced by relativistic beaming. For extragalactic TeV sources, the X-ray/TeV $\\gamma$-ray bands correspond to the highest energy ends of the synchrotron/inverse-Compton emission, which are produced by electrons accelerated up to the maximum energy (e.g., Inoue \\& Takahara 1996; Krik, Rieger \\& Mastichiadis 1998). At the highest energy ends, variability is expected to be most pronounced, and in fact, such large flux variations are observed, on a timescale of hours to days (e.g., Kataoka et al. 2001; Tanihata et al. 2001) or even shorter (minutes scale; Aharonian et al. 2007; Albert et al. 2007). Thus the X-ray/TeV variability can be the most direct way to probe the dynamics operating in jet plasma, in particular compact regions of shock acceleration which are presumably close to the central engine.  `Snapshot' multiwavelength spectra principally provide us with clues on the emission mechanisms and physical parameters inside relativistic jets. On the other hand, detailed studies of time variability not only lead to complementary information for the objectives above, but should also offer us a more direct window on the physical processes operating in the jet as well as on the dynamics the jet itself. However, short time-coverage and under-sampling have prevented detailed temporal studies of blazars. Only a few such studies have been made in the past for blazars, e.g., evaluation of the energy dependent ``time-lags'' based on the synchrotron cooling picture. For example, by using $ASCA$ data,  Takahashi et al. (1996) argued the soft X-ray ($<$ 1 keV) variation of Mrk~421, observed to lag behind that of the hard X-rays ($\\ge$ 2 keV) by $\\sim$ 4 ks, that may well be ascribed to the energy dependence of the synchrotron cooling timescale. More recently, Kataoka et al. (2000) interpreted an observed soft-lag and spectral evolution of PKS~2155-304 by a newly developed time-dependent synchrotron self-Compton (SSC) model.  The above $paradigm$ of ``soft-lag'' was concerned, however, by several aspects. First, intensive X-ray monitoring of blazars has revealed not only soft lags but in some cases hard lags (Takahashi et al. 2000) which may be a manifestation of another process, e.g., energy dependent acceleration. Very recently, signature of hard lag was clearly observed in 1ES 1218+304, but this is so far an only example of manifestation of possible acceleration timescale in any TeV blazars (Sato et al. 2008; also in this volume). Second, as Edelson et al. (2001) voiced concerns, there was a question about the reliability of lags that are smaller than the orbital periods ($\\sim$ 6 ks) of low Earth orbit satellites. This was refuted by Tanihata et al. (2001) and Zhang et al. (2004) who showed that, although periodic gaps introduce larger uncertainties than evenly sampled data, lags on hour-scale cannot be the result of periodic gaps. A time resolved cross correlation analysis of uninterrupted Mrk~421 data obtained by $XMM$-$Newton$ revealed lags of both signs, changing on timescales of up to a few 10$^3$ s (Brinkmann et al. 2005). Hence the situation is very complex and still under debate.  Variability studies covering larger dynamic range and broader span of timescales have become common for Seyfert galaxies and Galactic black-holes (Edelson \\& Nandra\\ 1999; Markowitz  et al.\\ 2004; McHardy et al. 2005; 2008 in this volume). From power spectrum density (PSD) analyses, it is well known that rapid fluctuations with frequency dependences $P(f)$ $\\propto$ $f^{ -1 \\sim -2}$, are characteristic of time variability in accreting black hole systems (e.g., Hayashida et al. 1998). Although their physical origin is still under debate, some tentative scenarios have been suggested to account for these generic, fractal features (e.g., Kawaguchi et al.\\ 2000). Similar studies have also been proposed for blazars, but still underway. It has been suggested that X-ray variability of TeV blazars indicates a strong red-noise ($P(f)$ $\\propto$ $f^ {-2 \\sim -3}$) behavior, compared to a fractal, flickering-noise of Seyfert galaxies (Kataoka et al. 2001).  These temporal studies are obviously important, however, special care must be taken if the data are not well sampled and relatively short compared to the variability timescale of the system. The prime motivation of this talk is to delineate the characteristic X-ray variability of TeV blazars, using a simple Monte Carlro simulation to evaluate the possible effects caused by observing time windows. Fortunately, we have now the $GLAST$ mission successfully launched in June 2008, as well as various excellent missions/telescopes available through radio to TeV energy bands. Moreover, future X-ray missions including Monitor of All-sky X-ray Image ($MAXI$) is ready for launch early next year. A great advantage of $GLAST$ and $MAXI$ is to provide very uniform exposure all over the sky, that may shed new light on the temporal studies of blazars especially on longer timescale from a month to years. ",
        "conclusions": ""
    },
    "0806/0806.0076_arXiv.txt": {
        "abstract": "\\n In a previous work [E.M. Prodanov, R.I. Ivanov, and V.G. Gueorguiev, {\\it Reissner--Nordstr\\\"om Expansion}, Astroparticle Physics 27 (150--154) 2007], we proposed a classical model for the expansion of the Universe during the radiation-dominated epoch based on the gravitational repulsion of the Reissner--Nordstr\\\"om geometry --- naked singularity description of particles that \"grow\" with the drop of the temperature. In this work we model the Universe during the Reissner--Nordstr\\\"om expansion as a van der Waals gas and determine the equation of state. ",
        "introduction": "\\n In 1971, Hawking suggested \\cite{hawk} that there may be a very large number of gravitationally collapsed charged objects of very low masses, formed as a result of fluctuations in the early Universe. A mass of $10^{14}$ kg of these objects could be accumulated at the centre of a star like the Sun. Hawking treats these objects classically and his arguments for doing so are as follows \\cite{hawk}: gravitational collapse is a {\\it classical} process and microscopic black holes can form when their Schwarzschild radius is greater than the Planck length $(Gh/c^3)^{-1/2} \\sim 10^{-35}$ m (at Planck lengths quantum gravitational effects do not permit purely classical treatment). This allows the existence of collapsed objects of masses from $10^{-8}$ kg and above and charges up to $\\pm 30$ electron units \\cite{hawk}. Additionally, a sufficient concentration of electromagnetic radiation causes a gravitational collapse --- even though the Schwarzschild radius of the formed black hole is smaller than the photon's Compton wavelength which is infinite.  Therefore, when elementary particles collapse to form a black hole, it is not the {\\it rest} Compton wavelength $hc/mc^2$ that is to be considered --- one should instead consider the {\\it modified} Compton wavelength $hc/E$, where $E \\sim kT >\\!\\!> mc^2$ is the typical energy of an ultra-relativistic particle that went to form the black hole \\cite{hawk}. Microscopic black holes with Schwarzschild radius greater than the modified Compton wavelength $hc/E$,  can form classically and independently on competing quantum processes. \\\\ Hawking suggests that these charged collapsed objects may have velocities in the range 50 -- 10 $\\!\\!$000 km/s and would behave in many respects like ordinary atomic nuclei \\cite{hawk}. When these objects travel through matter, they induce ionization and excitation and would produce bubble chamber tracks similar to those of atomic nuclei with the same charge. The charged collapsed objects survive annihilation and, at low velocities (less than few thousand km/s), they may form electronic or protonic atoms  \\cite{hawk}: the positively charged collapsed objects would capture electrons and thus mimic super-heavy isotopes of known chemical elements, while negatively charged collapsed objects would capture protons and disguise themselves as the missing zeroth entry in the Mendeleev table. \\\\ Such ultra-heavy charged massive particles (CHAMPS) were also studied by de Rujula, Glashow and Sarid \\cite{glashow} and considered as dark matter candidates. \\\\ Dark Electric Matter Objects (DAEMONS) of masses just above $10^{-8}$ kg and charges of around $\\pm 10$ electron units have been studied in the Ioffe Institute and positive results in their detection have been reported \\cite{drob} --- observations of scintillations in ZnS(Ag) which are excited by electrons and nucleons ejected as the relic elementary Planckian daemon captures a nucleus of Zn (or S). \\\\ The DAMA (DArk MAtter) collaboration also report positive results \\cite{dama} in the detection of such particles using 100 kg of highly radiopure NaI(Tl) detector. \\\\ Such heavy charged particles can serve as driving force for the expansion of the Universe during the radiation-dominated epoch in a classical particle-scale model, which we recently proposed \\cite{pig}. Along with this type of  particles, within our model, magnetic monopoles can also play the same role for the expansion of the Universe: it has been suggested \\cite{preskill} that ultra-heavy magnetic monopoles were created so copiously in the early Universe that they outweighed everything else in the Universe by a factor of $10^{12}$. \\\\ Our {\\it particle-scale} model gives the expected prediction for the behaviour of the scale factor of the radiation-dominated expanding Universe, $a(\\tau) \\sim \\sqrt{\\tau} \\,, $ and can be considered as a complement to the {\\it large-scale} Friedmann--Lema\\^itre--Robertson--Walker (FLRW) model (see, for example, \\cite{rw, mtw}) which describes the Universe as isotropic and homogeneous, with very smoothly distributed energy-momentum sources modeled as a perfect fluid, applicable on scales much larger than galactic ones. \\\\ This recently proposed \\cite{pig} classical mechanism for the cosmic expansion models the Universe as a two-component gas. One of the fractions is that of ultra-relativistic \"normal\" particles of typical mass $m$ and charge $q$ with equation of state of an ideal quantum gas of massless particles. The other component is \"unusual\" --- these are the particles of ultra-high masses $M$ (of around $10^{-8}$ kg and above) and charges $Q$ (of around $\\pm 10$ electron charges and above) --- exactly as those described earlier. \\\\ For an elementary particle such as the electron, the charge-to-mass ratio is $q/m \\sim 10^{21}$ (in geometrized units $c = 1 = G$), while for the \"unusual\" particles, $M \\, \\lower1pt\\hbox{\\tiny $\\stackrel{<}{\\sim}$}\\, Q$. In view of this, the general-relativistic treatment of elementary particles or charged collapsed objects of very low masses also necessitates consideration from Reissner--Nordstr\\\"om (or Kerr--Newman) viewpoint --- for as long as their charge-to-mass ratio remains above unity. We also treat the \"unusual\" particles classically (in line with Hawking's arguments outlined earlier). That is, the \"unusual\" particles are modelled as Reissner--Nordstr\\\"om naked singularities and the expansion mechanism is based on their gravito-electric repulsion. Instead of the Schwarzschild radius, the {\\it characteristic length} that is to be considered now and compared to the {\\it modified} Compton length \\cite{hawk}, will be the radius of the van der Waals-like impenetrable sphere that surrounds a naked singularity (see \\cite{cohen} for a very thorough analysis of the radial motion of test particles in a Reissner--Nordstr\\\"om field). As shown in \\cite{pig}, for temperatures below $10^{31}\\! $ K, the radius of the impenetrable sphere of an \"unusual\" particle of mass $10^{-8}$ kg and charge $\\pm 10$ electron units is greater than the {\\it modified} Compton wavelength of the \"unusual\" particle itself. \\\\ Naked singularities have been subject of significant scrutiny for decades. In the 1950s, the Reissner--Weyl repulsive solution served as an effective model for the electron. Very recently, a general-relativistic model for the classical electron --- a point charge with finite electromagnetic self-energy, described as Reissner--Nordstr\\\"om (spin 0) or Kerr--Newman (spin 1/2) solution of the Einstein--Maxwell equations, --- has been studied by Blinder \\cite{blinder}. Naked singularities are disliked --- hence the Cosmic Censorship Conjecture \\cite{penrose} --- but not ruled out --- there is no mathematical proof whatsoever of the Cosmic Censorship. At least one naked singularity is agreed to have existed --- the Big Bang --- the Universe itself. Of particular importance in the study of naked singularities are the work of Choptuik \\cite{choptuik}, where numerical analysis of Einstein--Klein--Gordon solutions shows the circumstances under which naked singularities are produced, and the work of Christodoulo \\cite{chr} who proved that there exist choices of asymptotically flat initial data which evolve to solutions with a naked singularity. The possibility of observing naked singularities at the LHC has been studied in \\cite{casadio} --- for example, a proton-proton collision could result in a naked singularity and a set of particles with vanishing total charge or with one net positive charge --- an event probably undistinguishable from ordinary particle production. In a cosmological setting, naked singularities have been well studied and classified --- see, for example, \\cite{ellis}. ",
        "conclusions": ""
    },
    "0806/0806.2629_arXiv.txt": {
        "abstract": "{Abundance variations in moderately metal-rich globular clusters can give clues about the formation  and chemical enrichment of globular clusters.} {CN, CH, Na, Mg and Al indices in spectra of 89 stars of the template metal-rich globular cluster M71 are measured and implications on internal mixing are discussed.} {Stars from the turn-off up to the  Red Giant Branch (0.87 $<$ log g $<$ 4.65) observed with the GMOS multi-object spectrograph at the Gemini-North telescope are analyzed. Radial velocities, colours, effective temperatures, gravities and spectral indices are determined for the sample.} {Previous findings related to the CN bimodality and CN-CH anticorrelation in stars of M71 are confirmed. We also find a CN-Na correlation, and Al-Na, as well as an Mg$_2$-Al anticorrelation.} {A combination of convective mixing and a primordial pollution by AGB or massive stars in the early stages of globular cluster formation is required to explain the observations.} ",
        "introduction": "Globular Clusters (GCs)  provide important information on the early chemical and dynamical evolution of the Milky Way.  Star-to-star abundance variations of light elements - Li, C, N, O, Na, Mg, and Al - are extensively reported in the literature (Gratton et al. 2004 and references therein). Li variations among turnoff (TO) stars and a Li-Na anticorrelation have been reported; among giant stars,  C and N abundances, as well as the pairs O:Na and Al:Mg are also found to be anticorrelated. Such anomalies have not been obtained for heavier elements.   These abundance variations have been reported in the literature over the  past two decades, but their origin is still widely debated.  Some of the abundance variations seen in globular-cluster giants can be explained by evolutionary mixing with migration of processed material through the CNO cycle to the surface of giant stars (Iben 1964; Charbonnel 1994), whereas a primordial-enrichment scenario, which requires  early contamination of intracluster material, is claimed by some authors (e.g. Smith 1987; Kraft 1994).  With a moderately high metallicity ($\\langle$[Fe/H]$\\rangle$=$-$0.73, Harris 1996), and  an old age of 10-12 Gyr (Grundahl et al. 2002; Meissner \\& Weiss 2006), M71 is often regarded as a prototype of northern metal-rich globular clusters and considered as a suitable globular cluster to study abundance variations. It is located near the Galactic plane (b = $-$4.6$^{o}$) and has a reddening of E(B-V) = 0.27 $\\pm$ 0.05 and an apparent distance modulus of (m-M)$_{\\rm V}$ = 13.60 $\\pm$ 0.10 (Geffert \\& Maintz 2000). Dinescu et al. (1999) have obtained space velocities for a set of Galactic globular clusters. For M71 they determined velocity components (U, V, W) = (77 $\\pm$ 14, $-$58 $\\pm$ 10, $-$2 $\\pm$ 14) km/s and a low eccentricity orbit,  which characterizes M71, kinematically, as a thick-disk cluster.  Chemical abundances were discussed in several previous studies of this cluster, including DDO photometry of giant stars (Hesser et al. 1977; Briley et al. 2001) and low- and high-resolution analysis of stars in different evolutionary stages from the Main-Sequence (MS) TO to the Red Giant Branch (RGB) tip (Cohen 1980; Smith \\& Norris 1982; Leep et al. 1987; Smith \\& Penny 1989; Penny et al.  1992; Sneden et al. 1994; Cohen 1999;  Ramirez \\& Cohen 2002; Lee et al. 2004; Lee 2005;  Boesgaard et al. 2005; Yong et al. 2006).  Many of these studies show a CN bimodality, with CN-CH anticorrelation, Na-O anticorrelation and  variations in Al abundance.  In this paper we present the main results  of an analysis of high S/N, medium resolution, Gemini/GMOS spectra of a large number of M~71 stars, from the main-sequence turnoff to the tip of the giant branch.  Our goals are twofold: to improve the statistics on abundance variations in M71 stars; and to study the behaviour of 14 spectral indices for the sample stars in order to shed light on the main astrophysical processes leading to the observed star-to-star abundance variations.   We estimate the C and N abundances of one CN-strong and one CN-weak giant from spectrum synthesis, based on state-of-the-art model photospheres and an up-to-date line list.  A comparison of our results with those based on recent high-resolution abundance analysis is also presented.   This paper is structured as follows. The observations and data reductions are described in Sect. 2. The analysis  of radial velocities, photometry, atmospheric parameters and line indices is presented in Sect. 3. In Sect. 4 the results are shown, and in Sec. 5 they are discussed and  contrasted with previous studies. Our conclusions are summarized in Sect. 6. ",
        "conclusions": "In M71 previous studies such that of Smith \\& Norris (1982) that first showed a bimodal distribution of CN-strong and CN-weak stars and an anticorrelation between CN and CH. This was followed by Smith \\& Penny (1989), Penny et al. (1992), Briley et al. (2001), Cohen (1999), Ramirez et al. 2001, Ram\\'{\\i}rez \\& Cohen (2002), Boesgaard et al. (2005), and  Lee (2005).  We measured CN, CH, Ca4277, iron and magnesium indicators, H$_{\\beta}$, NaD and Al3953 spectral indice, from low-resolution spectra of 89 stars of the metal-rich globular cluster M71, observed with the Gemini Multi-Object Spectrograph (GMOS) at the Gemini-North telescope. CN and CH strengths were obtained for 89 stars, among which 33 giants. As expected from evolutionary mixing theories and additional extra-mixing (Iben 1964; Charbonnel 1994; Denissenkov \\& VandenBerg 2003), we find a CN-CH anticorrelation. We find CN-strong and CN-weak stars, with around 30\\% of CN-strong ones, similar to other clusters such as NGC 6752 with about 50\\%.  We confirm a CN-bimodality besides the CN-CH anticorrelation, a CN-Na correlation, and  Al-Na and Mg$_2$-Al anticorrelation. The interpretation of CN bimodality is instead better understood in terms of primordial variations, and possible scenarios include an early enrichment by winds from intermediate mass stars in the AGB phase, and captured by low-mass stars, early in the cluster's life (Gratton et al. 2004 and references therein; Bekki et al. 2007) or early pollution by fast rotating massive stars (Decressin et al. 2007).  CN-strong and CN-weak bimodality is only seen in relatively metal-rich globular clusters, but not in all of them. Such behaviour is well studied in particular in 47 Tucanae, NGC 6752 and M4. Such abundance variations in metal-rich globular clusters is undoubtedly one of the most intrincate challenges for the current theory of stellar evolution, and further observations with larger samples would be interesting."
    },
    "0806/0806.1888_arXiv.txt": {
        "abstract": "The first optical spectrum of an isolated polycyclic aromatic hydrocarbon large enough to survive the photophysical conditions of the interstellar medium is reported. Vibronic bands of the first electronic transition of the all benzenoid polycyclic aromatic hydrocarbon hexa-\\textit{peri}-hexabenzocoronene were observed in the $4080-4530\\,${\\AA} range by resonant 2-color 2-photon ionization spectroscopy. The strongest feature at 4264\\,\\AA\\ is estimated to have an oscillator strength of $f=1.4\\times10^{-3}$, placing an upper limit on the interstellar abundance of this polycyclic aromatic hydrocarbon at $4\\times10^{12}$\\,cm$^{-2}$, accounting for a maximum of $\\sim0.02$\\% of interstellar carbon. This study opens up the possibility to rigorously test neutral polycyclic aromatic hydrocarbons as carriers of the diffuse interstellar bands in the near future. ",
        "introduction": "Organic material such as polycyclic aromatic hydrocarbons (PAHs) are held responsible for infrared emission features in carbon rich objects (AIBs) and has been suggested to account for as much as 20\\% of interstellar carbon \\citep{Leger1984,Allamandola85,Snow1995}. In addition, PAHs are considered leading candidates as carriers of the diffuse interstellar bands \\citep{Leger1985,Crawford85,Salama1999}, a series of diffuse absorption features superimposed on the extinction curve of the interstellar medium (ISM). The identities of the carriers of these bands remain the longest unsolved problem of laboratory astrophysics \\citep{Herbig1995,Sarre2006}. A scenario linking the AIBs to the DIBs is that amorphous aromatic material formed in carbon rich stellar outflows is further processed into PAHs and eventually the DIB carriers by the interstellar photophysical environment \\citep{Goto2003, Sloan2007, pino08}.  Despite the extensive literature reporting on PAHs in the ISM, not a single PAH species has ever been observed astronomically. This may be partly due to a paucity of laboratory data with which to compare unidentified features in astronomical spectra, or to guide astronomical searches. Obtaining optical spectra of isolated molecules at low temperature (as in diffuse clouds), \\emph{in vacuo}, remains a challenge of modern laboratory astrophysics \\citep{Sharp2005}.  Modelling of generic PAHs in the diffuse interstellar medium aims to guide the laboratory search for the carriers of the DIBs. \\citet{Leger1985} found that species must contain no fewer than fifteen atoms if they are to avoid photo-thermolysis and \\citet{Lepage2003} concluded that PAHs containing more than about 30 carbon atoms would be largely stable. Furthermore, it is considered that a large proportion of interstellar PAH material would exist in its ionized form(s) \\citep{Crawford85}. As such, the challenge to laboratory-based scientists is to place large PAHs into the vacuum, cool them down, perhaps ionize them and obtain their spectra. Prior to the present work, the largest neutral PAH studied spectroscopically in such a way was ovalene (C$_{32}$H$_{14}$), the size of which is borderline for surviving the ISM intact \\citep{Amirav1980}.   An unguided laboratory search is daunting, however, as for systems containing $4-10$ fused aromatic rings there are over 20000 possible PAH structures \\citep{Dias2004}. In order to guide the laboratory search for DIB carriers, some selection mechanism is required, based on spectral intensity, or structure. Recently we suggested that all-benzenoid PAHs (ABPAHs) could represent such a selection mechanism \\citep{Troy2006}. These ABPAHs are formed from fused benzene rings which may be drawn as being separated by single bonds. They are characterized by higher ionization potentials and more energetic electronic transitions than their non all-benzenoid isomers. Of the more than 20000 possible structures containing $4-10$ fused rings, only 17 are all-benzenoid in character. The all-benzenoid structural motif thus provides a selection mechanism for the laboratory and perhaps the ISM. Since the ABPAHs are comparatively easy to prepare in the laboratory, it may also be that they form preferentially in the ISM. We obtained the gas phase spectrum of the smallest member of this family: triphenylene (C$_{18}$H$_{12}$)\\citep{Kokkin2007}. However, its small size means that even its lowest energy transitions are in the ultraviolet, far from the shortest wavelength DIBs. The ABPAHs present three optical band systems named the $\\alpha$, $p$ and $\\beta$-bands in increasing intensity and energy. We found that for the strong $\\beta$-band system to fall in the DIB region, it must contain no fewer than 84 carbon atoms \\citep{Troy2006}. However, Hexa-\\emph{peri}-hexabenzocoronene (HBC) is known to exhibit the weaker $\\alpha$-bands in the region of the DIBs. HBC , pictured in Fig. \\ref{fig}, possesses 42 carbon atoms and is thus large enough to survive the interstellar radiation field, as modelled by \\citet{Lepage2003}. Since its lowest energy transitions occur in the visible region, near the strongest DIB (4429\\,\\AA\\ in air) it was suggested as a carrier of this interstellar band by \\citet{Hendel1986} who obtained spectra in a 1,2,4-trichlorobenzene solution. Furthermore, its red phosphorescence suggested that it could also be a carrier of the \\emph{Red Rectangle} Bands, the unidentified emission features superimposed on the extended red emission of the protoplanetary nebula surrounding HD 44179 \\citep{Cohen1975,Schmidt1980}. \\citet{Hendel1986} also measured the mass spectrum of HBC and found that while HBC$^{2+}$ and HBC$^{3+}$ were abundantly produced, there was very little H-loss, suggesting it to be a very tough molecule of the sort required to survive harsh interstellar radiation. Despite these suggestions, a gas phase spectrum of HBC has never been reported. In this letter we present the gas phase spectrum of jet-cooled, neutral hexa-\\emph{peri}-hexabenzocoronene (HBC, C$_{42}$H$_{18}$). From these measurements an upper limit on the column density of HBC in the diffuse interstellar medium (DISM) is obtained. ",
        "conclusions": "We have succeeded in obtaining the excitation spectrum of hexa-\\emph{peri}-hexabenzocoronene (C$_{42}$H$_{18}$) as a jet-cooled, isolated molecule in the vacuum. Its strongest transition in the region studied was found to lie at 4261.1\\,\\AA, which does not correspond to any known diffuse interstellar band. An estimate of the oscillator strength of this band was made at $f=1.4\\times10^{-3}$, which afforded an upper limit on the interstellar abundance of HBC of $4\\times10^{12}$\\,cm$^{-2}$, or 0.02\\% of interstellar carbon. However, the same technique can be used to obtain the spectra of still larger and less symmetric systems which may exhibit much stronger transitions. The collection and analysis of the spectra of these all benzenoid PAH species builds up the database of gas phase spectra required to constrain models of PAH abundances in the ISM in concert with detailed observation. In particular, an exploration of larger PAHs with strong ($f\\sim1$) transitions will provide a rigorous test of neutral PAHs as carriers of the DIBs, thus opening the door to solving the longest standing mystery in astronomical spectrocopy, or putting to bed a long-favored class of candidate.  There remains an impetus to collect spectra of cold PAH cations. The present work establishes the ability to place PAHs of 42 carbons in the vacuum by laser desorption and supersonic expansion. The next step is to produce the cold cations by threshold ionization and obtain their spectra by multiphoton dissociation \\citep{Pino2007}."
    },
    "0806/0806.4463_arXiv.txt": {
        "abstract": "We have developed a new method to improve the transit detection of Earth-sized planets in front of solar-like stars by fitting stellar microvariability by means of a spot model. A large Monte Carlo numerical experiment has been designed to test the performance of our approach in comparison with other variability filters and fitting techniques for stars of different magnitudes and planets of different radius and orbital period, as observed by the space missions CoRoT and Kepler. Here we report on the results of this experiment. ",
        "introduction": "\\noindent We present a comparison among the performance of three methods applied to filter stellar variability for the detection of Earth-like planetary transits in the light curves of solar-like stars. This requires two steps: first, the filtering of stellar variability to remove the effects of photospheric cool spots and bright faculae, whose visibility is modulated by stellar rotation; secondly, the search for transits in the filtered light curves by means of suitable detection algorithms.  We recently proposed a filtering method based on a model of the flux variations of the Sun as a star, the so-called 3-spot model (\\cite[Lanza et al. 2003]{Lanzaetal2003}). Its performance was compared with that of another method, the 200-harmonic fitting, by \\cite[Bonomo \\& Lanza (2008)]{BonomoLanza2008}. They showed that the 3-spot model has a better performance than the latter when the standard deviation of the noise is at least 2-4 times larger than the central depth of the transit. On the other hand, the 200-harmonic fitting is better when the standard deviation of the noise is comparable to the transit depth.  Here we extend the comparison to the iterative non-linear filter by \\cite[Aigrain \\& Irwin (2004)]{AigrainIrwin2004}. A comparison among different variability filters is important since only the coupling of the best filtering method with the best planetary transit detection algorithm allows us to maximize transit detection efficiency. This is especially relevant when we want to detect small, terrestrial planets, which is a challenge to CoRoT and Kepler missions.  \\eject ",
        "conclusions": "\\noindent For $\\sigma= 100$ ppm, the filtering methods that achieve the best performance are the INL filter and the 200-harmonic fitting with detections up to 98 percent for $R=1.25$ R$_{\\oplus}$ and $P=10$ days. In most of the cases they give comparable results, although in some instances the INL filter has a slightly better performance, owing to the Gibbs phenomenon (\\cite[Morse \\& Feshbach 1954]{MorseFeshbach54}) affecting the 200-harmonic fitting.  When $\\sigma \\geq 200$ ppm, the method with the best performance is the INL filter when we use an appropriate window of 2 days for the median boxcar filter. It shows a performance comparable with that of the 3-spot model in most of the cases, even better in some instances (see Fig.~\\ref{fig1}). On the other hand, the 200-harmonic fitting has the worst performance because of the Gibbs phenomenon (see Fig.~\\ref{fig1} and ~\\ref{fig2}).   \\begin{figure}[!b] \\vspace*{-1.0 cm} \\begin{center} \\includegraphics[width=6cm,angle=90]{Bonomo_f2.ps} \\vspace*{-0.8 cm} \\caption{The shape of a transit of an Earth-like planet as it appears in the ideal case (solid black line), in the residuals obtained with the 3-spot model (dashed blue line) and in those of the 200-harmonic fitting (dotted red line). Note the reduction of the transit depth and the overshooting at the edges of the transit dip due to the Gibbs phenomenon in the case of the 200-harmonic fitting.} \\label{fig2} \\end{center} \\end{figure}    The performance of the INL filter depends critically on the adopted extension of the filter window. An optimal window of 2 days has been chosen for our analysis. Shorter windows negatively affect the transit detection since they give rise to a reduction of the transit depth in the filtered light curve (see Fig.~\\ref{fig3}), in which case the 3-spot model and the 200-harmonic fitting would prove to be the best methods for the cases with $\\sigma \\geq 200$ ppm and $\\sigma= 100$ ppm, respectively.    \\begin{figure}[!t] \\begin{center} \\includegraphics[width=8cm]{Bonomo_f3.ps} \\vspace*{-0.4 cm} \\caption{\\emph{Upper panel}: one of the light curves with transits simulated for the First CoRoT Blind test (\\cite[Moutou et al. 2005]{Moutouetal2005}; ID=460). \\emph{Middle panel}: the light curve filtered by means of the INL filter with a window of 2 days. \\emph{Bottom panel}: the filtered light curve with a 0.75 day window. Note the disappearance of the transits when the window extension is reduced.} \\label{fig3} \\end{center} \\end{figure}    The optimal width of the median filter window depends on the magnetic activity level of the star and its rotation period. Specifically, the higher the activity level and the shorter the rotation period, the shorter the optimal window, because the time scales of the flux variations decrease with increasing activity. In other words, in the case of highly active stars the window extension has to be shortened with respect to the solar case, otherwise some oscillations or transit-like features will appear in the residuals owing to a bad filtering of the variability. To fix automatically the window extension, we propose a method similar to that of \\cite[Regulo et al. (2007)]{Reguloetal2007}, computing the power spectrum of the time series and choosing an extension corresponding to the frequency where the power spectral density goes below a fixed threshold, usually set at $10^{-6}$ of the maximum power level.  We conclude that the INL filter, when applied with a suitable choice of its window, has a better performance than more complicated and computationally intensive methods of fitting solar-like variability, like the 200-harmonic fitting or the 3-spot model."
    },
    "0806/0806.3075_arXiv.txt": {
        "abstract": "We revisit the proximity effect produced by QSOs at redshifts $2.1-3.3$ applying the FLO approach \\citep{saitta} to a sample of $\\sim 6300$ \\Lya lines fitted in 21 high resolution, high signal-to-noise spectra. This new technique allows to recover the hydrogen density field from the \\HI\\ column densities of the lines in the \\Lya forest, on the basis of simple assumptions on the physical state of the gas. To minimize the systematic uncertainties that could affect the density recovering in the QSO vicinity, we carefully determined the redshifts of the QSOs in our sample and modelled in detail their spectra to compute the corresponding ionising fluxes. The mean density field obtained from the observed spectra shows a significant over-density in the region within 4 proper Mpc from the QSO position, confirming that QSOs are hosted in high density peaks. The absolute value of $\\rho/\\langle\\rho\\rangle$ for the peak is uncertain by a factor of $\\sim 3$, depending on the assumed QSO spectral slope and the minimum \\HI\\ column density detectable in the spectra. We do not confirm the presence of a significant over-density extending to separations of $\\sim 15$ proper Mpc from the QSO, claimed in previous works at redshifts $\\langle z\\rangle \\simeq 2.5$ and 3.8. Our best guess for the UV background ionisation rate based on the IGM mean density recovered by {\\small FLO} is $\\Gamma_{\\rm UVB}\\simeq 10^{-12}$ s$^{-1}$. However, values of $\\Gamma_{\\rm UVB} \\simeq 3\\times 10^{-12}$ s$^{-1}$ could be viable if an inverted temperature-density relation with index $\\alpha\\simeq-0.5$ is adopted. ",
        "introduction": "The ultraviolet radiation emitted by quasars (QSOs) is considered the dominant source of  ionisation of the intergalactic medium (IGM) at redshifts 2-4. Most of the absorption lines seen blue-ward of the \\Lya emission in QSO spectra (the so-called \\Lya forest) are ascribed to fluctuations in the low to intermediate density IGM \\citep[see][for a recent review]{meiksin}. As a consequence, \\Lya lines can be used as probes of the properties and redshift evolution of the UV  ionising background.  Observations show that the number density of \\Lya lines increases with redshift, but within single QSO spectra the number density of \\Lya lines decreases as the redshift approaches the QSO emission redshift. This effect was first noticed by \\citet{carswell82} and confirmed by later studies \\citep{murdoch,tytler}.  \\citet{bajtlik} called this deficiency of \\Lya absorptions near the background QSO ``proximity effect'' and attributed it to the increased ionisation of the \\Lya clouds near the QSO due to its ionising flux.  They used the proximity effect in 19 low resolution QSO spectra to estimate the intensity of the ultraviolet background radiation (UVB) at the Lyman limit\\footnote{The Lyman limit corresponds to the hydrogen ionisation energy $E_{\\rm ion}=13.6\\eV$ or $\\lambda_{\\rm LL}=912 $ \\AA.}, $J_{\\rm LL}$, for which they found the value $\\log J_{\\rm LL} \\simeq -21.0 \\pm 0.5$ ergs cm$^{-2}$ sec$^{-1}$ Hz$^{-1}$ sr$^{-1}$ over the redshift range $1.7<z<3.8$.  Several other authors carried out this analysis on other data sets \\citep{lueal,kulkarni,bechtold,williger,cristiani95,giallo96}. \\citet{lueal} found the same result of \\citet{bajtlik} in the same redshift interval but using 38 QSO spectra. \\citet{bechtold}, with 34 low resolution QSO spectra covering the range  $1.6<z<4.1$, found $\\log J_{\\rm LL}\\simeq -20.5$. This value, 3 times larger than that of \\citet{bajtlik} and \\citet{lueal}, was not confirmed by \\citet{giallo96}. They obtained $\\log J_{\\rm LL}\\simeq -21.3^{+0.08}_{-0.09}$, using 10 higher resolution ($R \\sim 25000$) QSO spectra in the same redshift interval as \\citet{bechtold}. More recently, \\citet{scott00} considered a sample of 74 intermediate resolution QSO spectra in the redshift interval $1.7 < z < 4.1$, from which they obtained a value $\\log J_{\\rm LL} \\simeq -21.1^{+0.15}_{-0.27}$. These authors paid a particular attention to the correct estimate of the systemic redshift of the QSOs. Indeed, if a QSO redshift lower than the true one is used in the analysis, the ionising effect of the QSO on the \\Lya clouds, and thus the derived value of $J_{\\rm LL}$, are over-estimated.  In the standard analysis of the proximity effect it is assumed that the matter distribution is not altered by the presence of the QSO. The only difference between the gas close to and far away from the QSO is the increased  photoionisation rate due to the QSO emission. A consequence of this hypothesis is that there should be a correlation between the strength of the proximity effect and the luminosity of the QSO. However, observational results are not conclusive on this subject (e.g. \\citealt{lueal,bechtold,srianand}; but see also \\citealt{liske01}). It is in fact likely that QSOs occupy over-dense regions. Hierarchical models of structure formation predict that super-massive black holes, that are thought to power QSOs, are in massive halos \\citep{granato,fontanot06,daangela} which are strongly biased to high-density regions.  The main aim of this work is to investigate the density distribution of matter close to QSOs from the proximity effect. To this purpose, we applied the {\\small FLO} (From Lines to Over-densities) technique, developed in \\citet[][Paper I]{saitta}, to a sample of 21 high resolution, high signal-to-noise ratio QSO spectra. {\\small FLO} converts the list of \\HI\\ column densities of \\Lya lines in a QSO spectrum into the underlying \\nh hydrogen density field. This method significantly reduces the drawbacks of the line fitting approach, in particular, the dependence on the fitting tool, the subjectivity of the result, and the strong dependence of statistics on the number of weak lines, which is in general poorly known. We could also constrain the value of the UVB ionisation rate, $\\Gamma_{\\rm UBV}$\\footnote{$\\Gamma_{\\rm UVB} \\equiv 4\\pi \\, \\int_{\\nu_{\\rm LL}}^{\\infty} \\sigma_{\\nu} (J_{\\nu}/h\\nu) d\\nu$ where, $\\sigma_{\\nu}$ is the photo-ionisation cross section, $\\nu_{\\rm LL}$ is the frequency at the Lyman limit and $J_{\\nu} = J_{\\rm LL} (\\nu/\\nu_{\\rm LL})^{-\\gamma}$.}, by matching the recovered density field in the \\Lya forest region with the mean cosmic density.   Similar analyses were performed in previous works using a different approach, based on the determination of the cumulative probability distribution function of pixel optical depth. \\citet{rollinde} studied the density structure around QSOs using 12 high resolution-spectra that belong also to our sample. These authors marginally detected the presence of an over-density at separations $3\\la r \\la 15$ proper Mpc, assuming an hydrogen ionisation rate $\\Gamma_{\\rm UVB} = 10^{-12}$ s$^{-1}$ (corresponding to $\\log J_{\\rm LL} \\simeq -21.4$). \\citet{guimaraes}, using the same technique, investigated the distribution of matter density close to 45 high-redshift ($z_{\\rm em} \\sim 3.8$) QSOs observed at medium spectral resolution. Their study reveals gaseous over-densities on scales as large as $\\sim 15$ Mpc, with higher over-densities for brighter QSOs.  The paper is organised as follows: in Section~\\ref{sectFLO}, we introduce the {\\small FLO} technique and describe an improvement to the method. Section~\\ref{sectOBS} reports the details of the observed data sample, the QSO emission redshift and luminosity determination, and the fitting analysis. In Section\\ref{sect_PROX}, the proximity effect of QSOs on the surrounding gaseous medium is pointed out and the over-density close to QSOs is reconstructed with {\\small FLO}. The characteristics of the cosmological simulations used to obtain the mock spectra and the comparison with observations are described in Section~6. Our conclusions and the prospect for the future are listed in Section~7.  Throughout this paper, we adopt a $\\Lambda$CDM cosmology with the following values for the cosmological parameters at $z=0$: $\\Omega_{\\rm m}=0.26,\\ \\Omega_{\\Lambda}=0.74,\\ \\Omega_{\\rm b}=0.0463$, $n_s=0.95$, $\\sigma_8 = 0.85$ and $H_0 = 72$ km s$^{-1}$ Mpc$^{-1}$. These parameters are consistent with the best fits values obtained from the latest results on the cosmic microwave background \\citep{WMAP} and from other flux statistics of the \\Lya forest \\citep[e.g.][]{viel06}. ",
        "conclusions": "\\label{sectCONC}   In this work, we used a sample of $\\sim 6300$ \\Lya\\ absorption lines obtained from high resolution, high signal-to-noise ratio spectra of 21 QSOs in the redshift range $2.1 \\la z_{\\rm em} \\la 3.3$ and  the {\\small FLO} algorithm for the reconstruction of baryon density fields (described in details in Paper I) to investigate the baryon distribution as a function of the distance from the QSO.  The effect of the QSO radiation (supposed to be isotropic) dominates over the UV background (UVB) in a sphere with radius varying from $\\sim 3$ to 15 Mpc for our sample, with an average of $\\sim 8$ Mpc. The increased ionisation flux modifies the local line number density causing the so-called {\\sl proximity effect}, that was used in the past to estimate the intensity of the UVB at the Lyman limit. However, the local line number density is determined not only by the ratio between the UV flux from the QSO and the UVB but also by a variation of the density field due to the fact that QSOs likely lie on density peaks.  The recovery is extremely sensitive to the emission redshift and to the UV flux of the QSOs. We spent a significant effort to obtain, both from the literature and from the spectra at our disposal, the best estimates for the systemic redshifts of the QSOs in our sample in order to reduce the uncertainties to less than $100-200\\ \\kms$. For the 16 QSOs for which the redshift was determined with $H_{\\beta}$, [\\OIII] (5 QSOs) or with the \\OI, \\MgII\\ UV emission lines, the average velocity difference with respect to the determinations present in the literature \\citep[e.g.][ based mainly on the high-ionisation UV lines and \\Lya]{kim04} is $\\sim -500\\ \\kms$. For the QSO UV flux, we adopted three different models: two with a power law with a fixed slope and one for which each QSO has its own slope obtained from the comparison with a spectral library. The resulting ionisation rates for the three models are within a factor of two.  We improved the performances of {\\small FLO} with respect of paper I by adopting a smoothing scale $\\sim 30$ percent smaller than the Jeans length, following the prescription by \\citet{gnedin_hui}. This contrivance allows us to solve the problem of over-estimation of the over-densities we had in paper I and to correctly recover the mean density in the IGM.  We applied the {\\small FLO} algorithm to each QSO spectrum and recovered the sample-averaged $\\delta$ field as a function of the distance from the emitting sources, in bins of 4 proper Mpc. The obtained results are described in the following sections.  \\subsection{QSO over-density from the proximity effect}  \\begin{description} \\item[i)] If only the UVB ionisation rate is adopted in eq.\\ref{tra} of {\\small FLO}, neglecting the QSO radiation, a decrease of the mean density field significant at $\\sim 2\\,\\sigma$ is observed within 8 proper Mpc of the emitting QSO, confirming the presence of a proximity effect.  \\item[ii)] When also the QSO ionisation rate is taken into account an over-density significant at more than $3\\,\\sigma$ is recovered at proper separations of less than 4 Mpc (first bin).  \\item[iii)] The absolute $\\delta$-value of the QSO over-density is uncertain by an overall factor of $\\sim 3$. A factor of $\\sim 2$ variation is due to the lower and upper-limit $\\delta$-value derived for the empty bins during the field reconstructions: corresponding to no absorber or to one absorber with the minimum \\HI\\ column density detectable in our sample, respectively. This uncertainty can be improved by using more spectra with a larger signal-to-noise ratio. Another factor of $\\sim 1.1-1.5$ is due to the different assumptions for the QSO continuum in the region blue-ward of the \\Lya\\ emission.  \\item[iv)] The uncertainty on the systemic redshift of the QSOs used to determine the density distribution in their neighbourhood is $\\la 200\\ \\kms$ corresponding to separations $\\la 800$ proper kpc. As a consequence, it should have negligible effects on the $\\delta$-value of the peak.  \\item[v)] We compared the resulting density field with the average density field in the proximity of massive haloes ($M>2\\times 10^{12}$ M$_{\\odot}\\ h^{-1}$) in a cosmological hydro-simulations. The two distributions are in reasonably good agreement.  \\item[vi)] There is a significant discrepancy between our results and the previous determinations of the matter distribution around QSOs, obtained with the optical depth statistics \\citep[ODS, ][]{rollinde,guimaraes} at redshifts $\\langle z_{\\rm em}\\rangle \\simeq 2.5$ and 3.8. The ODS method recovers gaseous over-densities extending to scales as large as $\\sim 15$ Mpc while our overdensity is limited to a region closer than 4 proper Mpc from the QSO. We would need to increase our sample, in particular with QSOs without associated systems in order to study if the brightest objects resides in more extended peaks. \\end{description}   \\subsection{Constraints on the IGM physical parameters}   \\begin{description} \\item[i)] In the hypothesis of a temperature of the gas at the mean density $T_0=1.8\\times 10^4$ K and an index of the temperature-density relation for the IGM $\\alpha=0.6$, an UVB ionisation rate of $\\Gamma_{\\rm UVB} \\simeq 10^{-12}$  s$^{-1}$ gives the correct IGM mean density. On the other hand, $\\Gamma_{\\rm UVB} \\simeq 2\\times 10^{-12}$ and $3\\times 10^{-12}$ s$^{-1}$ are excluded at more than 2 and $3\\,\\sigma$, respectively, because the recovered {\\small FLO} IGM density field overestimates the mean density.  \\item[ii)] Values of $\\Gamma_{\\rm UVB} > 10^{-12}$ s$^{-1}$ can be reconciled with the correct IGM mean density if different combination of $T_0$ and $\\alpha$ are adopted.  In particular, an inverted temperature-density relation with $\\alpha\\simeq -0.5$ used in the {\\small FLO} algorithm gives the correct IGM mean density for $\\Gamma_{\\rm UVB} \\simeq 3\\times 10^{-12}$ s$^{-1}$ and $T_0\\simeq 2.3\\times10^4$ K. Such large values of the ionisation rate could arise as a consequence of the \\HeII\\ reionisation at $z\\sim3$ and of UVB fluctations \\citep[e.g.][]{meiksin_white}. The performances of {\\small FLO} with different set of parameters than the reference one have however to be tested with numerical simulation. We defer the details of this analysis to a future paper.  \\end{description}  \\subsection{Limiting factors and future developments}  \\begin{description} \\item[i)] Most of the high redshift QSOs have redshifts determined from UV emission lines which are known to be systematically shifted with respect to systemic redshifts. This is particularly critical for proximity effect studies both along and transverse to the line of sight. An improvement in this sense is expected from the new intermediate-resolution, UV to near-IR spectrograph X-Shooter at the VLT \\citep{vernet}, that will be operative from the first trimester of 2009. \\item[ii)] The temperature of the IGM gas at the mean density is highly uncertain. Its best determinations date back to 2000, the evidence of a jump in its value at $z\\sim 3$ was marginal and should be verified with the present larger samples of high-resolution, high-signal-to-noise QSO spectra. Furthermore, those temperature estimates are used in many cosmological hydro-simulations to re-normalise the ionising background intensity. \\item[iii)] The intensity and nature of the UV background is another unsolved riddle which should require a new observationally-based determination (with the caveat in i) since the present simulations are only now starting to have the resolution and the physics (e.g. the radiative transfer) needed to derive it self-consistently. \\end{description}"
    },
    "0806/0806.3732_arXiv.txt": {
        "abstract": "We  have calculated a series of models of outflows from the obscuring torus in active galactic nuclei (AGN). Our modeling assumes that the inner face of a rotationally supported torus is illuminated and heated by the intense X-rays from the inner accretion disk and black hole. As a result of such heating a strong biconical outflow is observed in our simulations. We calculate 3-dimensional hydrodynamical models, assuming axial symmetry, and including the effects of X-ray heating, ionization, and radiation pressure.  We discuss the behavior of a large family of these models, their velocity fields, mass fluxes and temperature,  as functions of the torus properties and X-ray flux.  Synthetic warm absorber spectra are calculated, assuming pure absorption, for sample models at various inclination angles and observing times. We show that these models have mass fluxes and flow speeds which are comparable to those which have been inferred from observations of Seyfert 1 warm absorbers, and that they can produce rich absorption line spectra. ",
        "introduction": "One of the insights provided by observations of Seyfert galaxies and some quasars is the prevalence in their X-ray spectra of spectral lines and bound-free continua from ions of intermediate-Z elements . Early observations of Seyfert 1 galaxies using proportional counters and solid  state detectors revealed spectra with strong absorption features in the 0.1-10 keV range \\citep{Halpern84}. These features were attributed mostly to the edges of hydrogen and helium - like oxygen. The term \"warm absorber\" was proposed owing to the fact that the observed X-ray absorbing gas has an electron temperature lower than it would be if a similar level of ionization were produced by collisional ionization. However, more detailed spectroscopic studies were hampered by the limited X-ray resolution of the ASCA and ROSAT satellites. The grating spectrographs on the X-ray telescopes {\\it Chandra} and {\\it XMM-Newton} provide unprecedented spectral resolution up to $\\sim10\\, {\\rm keV}$. These show that X-ray spectra obtained from $\\sim$ half of low-red-shift active galactic nuclei (AGN) contain many lines from ions of Fe, Si, S, O, Mg, and Ne, and that these are generally broadened and blueshifted by 100-500 km/s  \\citep{Kaspi02, Steenbrugge05}. The presence of X-ray absorbing gas has been confirmed in the majority of AGNs which are bright enough to allow detections \\citep{Reynolds97,McKernan07}. There is also a partial correspondence between UV and X-ray absorbers \\citep{Crenshaw99}.  X-ray observations of warm absorbers are consistent with the Seyfert 1/Seyfert 2 dichotomy. For example, the properties of the X-ray emission in the Seyfert 2 galaxy NGC 1068 corresponds to the scattered emission expected from warm absorbers in Seyfert 1 galaxies \\citep{Kinkhabwala02}.  Constraints on the position and dynamics of the X-ray absorbing gas can be deduced from the observed widths and virial arguments, and also from the variability studies of these spectra \\citep{Behar03, Netzer03}. These show an absence of correlated response of the warm absorber gas to rapid changes ($\\sim$ days) of the continuum.   This implies that the ionization time scale in the warm absorber gas is long ($\\gtrsim$ months). Combined together, the line blueshifts, widths, and time variability analysis favors an origin of the warm absorber gas at $R \\gtrsim 1\\, {\\rm pc}$ away from the BH. This estimate coincides with the likely location of absorbing matter responsible for obscuration in Seyfert 2 galaxies \\citep{KrolikBegelman88}. The existence of an outflow from the torus has been suggested by \\cite{KrolikBegelman86, KrolikBegelman88}, and as the source of warm absorber flows by \\cite{KrolikKriss1,KrolikKriss2}.  It is believed that this matter is in the form of  a molecular torus which is responsible for obscuring the broad line region in Seyfert 2 galaxies, and which is thought to exist  in most low and intermediate luminosity AGN \\citep{AntonucciMiller86}. A growing body of direct observational evidences advocates for the existence of the obscuring torus. Mid-infrared  high spatial resolution studies of the nucleus of NGC 1068 using the Very Large Telescope Interferometer have resolved a dusty structure which is 2.1 pc thick and 3.4 pc in diameter ~\\citep{JaffeNATUR}. Observations support a multi-temperature model: the temperature of the warm component was established to be 300 K and  inside of it a second,  compact and hot ($>$800K) component has been found. Further studies of NGC 1068 systematically reduced estimates of the temperatures of different components \\citep{PonceletPerrin06}. Observations of the Circinus galaxy, which is among the closest prototype Seyfert 2 galaxies, also revealed a dense and warm $T\\gtrsim 300$ K component at about $0.2\\, {\\rm pc}$ from BH and cooler $T<300$ K component at $1\\,{\\rm pc}$ \\citep{Tristam07}. If the hotter component is located closer to the X-ray source, it may be attributed to the inner part of the torus, heated by the radiation of the compact nucleus. Although the evidence is strongest for nearby active galaxies, there is also a strong motivation to think that within the same obscuring torus paradigm exist those quasars whose central regions are heavily obscured by gas and dust (Type II quasars). Evidence for this comes from spectro-polarimetric observations by \\cite{Zakamska06}.  This paper is part of a series whose main goal is to test the hypothesis that the torus is the origin for the warm absorber flow. Preliminary results of this work have been reported in ~\\cite{Dora08} (Paper 1), in which we presented the results from a sample model and showed that the adopted model is promising in explaining the warm absorber phenomenon. In this paper we provide  more details of our methods, and display results of models which span the space of input parameters.  We present and discuss the hydrodynamic quantities which characterize our models:  mass fluxes, velocity fields, and temperature structure.  We also show sample X-ray spectra, which we will discuss extensively in a later paper of this series.  Our approach can be described as having three basic parts: i) setting up initial conditions, which requires defining an initial torus configuration and making assumptions about the external source of radiation; ii) implementation of the wind driving force (local heating-cooling rates and radiation pressure force) and actual 2D hydrodynamical calculations. The latter includes the numerical solution of the time-dependent 2D (so called 2.5D) system of equations, which takes into account centrifugal forces, and radiation pressure and heating terms; iii) calculating of the X-ray line spectra using a code which adopts Sobolev radiation transfer and ionization calculations for plasma in the intense X-ray field. Each of these steps is described in what follows. ",
        "conclusions": "We have studied X-ray excited winds from the putative gas-dusty torus in AGN. We approach this problem using numerical methods combining detailed hydrodynamical modeling with calculation of the warm absorber spectra. Hydrodynamical calculations include two-dimensional, axially-symmetric rotating flow, driven primarily by X-ray heating. Compton, bremsstrahlung, and photoionization heating/cooling processes were taken into account as well as the radiation pressure force, which was calculated in the Sobolev approximation. A code combining XSTAR for photoionization calculations with the Sobolev radiation transfer has been developed for the calculation of the spectra.  We find that a rotationally supported torus heated by radiation from the inner accretion disk and  black hole can indeed be a source of the material we observe in the warm absorber flow. We find that the inner throat of the torus is not only important as a source of the gas but also because it creates a funnel for the outflowing wind. This leads generally to larger velocities within the funnel, and different velocity distribution within the warm absorber flow from those derived from models based on spherically-symmetric  winds. The wind mass-loss rate within the funnel is not very sensitive to the details of the initial torus distribution and approaches $\\sim 0.02-0.09\\, {\\dot M}_\\odot \\, {\\rm yr}^{-1} $. Strong X-rays heat the gas within the funnel, producing a fast, $\\sim 1000\\, {\\rm km}\\, s^{-1}$, ionized flow near the axis, and slower, $\\lesssim 500\\, {\\rm km}\\, s^{-1}$, flow closer to the funnel walls. This is where optical depth effects become important and a warm absorber spectrum is produced. Using methods developed in studies of X-ray binaries we were able to estimate the mass-loss rate from such funnel flow, finding it to be in a good agreement with our numerical solution.  The funnel flow is found to be promising with respect to obtaining high velocity warm absorber flows. What is beyond the scope of our models is the possibility of having multiple phases in such high velocity flow, on spatial scales smaller than our grid resolution.  Our treatment of the gas thermal properties will produce two-phase behavior at our grid resolution; we do not find this behavior, owing to the fact that the cooling timescales are generally too long.  The answer to the question of whether there can be high velocity 'bullets' or  'embedded clouds ' on length scales smaller than the resolution of the grid is related to the problem of the origin of broad and narrow  UV/optical line emitting clouds, and requires different computational methods from those employed here.  Our models which have initial Compton depths $\\tau_\\bot^{\\rm C}\\gtrsim 1$, aspect ratio $R_0/H\\sim 1$, and located at $0.5\\lesssim r \\lesssim 1.5$ pc predict warm absorber spectra, thus confirming the main conclusion made in paper 1. The existence of such spectra depends on the fact that the flow is intrinsically two-dimensional, meaning that both the dynamics of the funnel flow is different from 1D models and optical depth effects are important as they strongly depend on inclination. The latter point requires that we include the entire torus in the computational domain rather than considering it as a boundary condition. The distribution of the ionization parameter, $\\xi$ depends strongly on $\\theta$, further confining the range of angles where conditions are right for the warm absorber flow to be observed. In most of our models warm-absorber-like spectra are produced in a $10^\\circ$ range, at  $\\theta\\simeq 40 \\pm 5^\\circ$.  This range is set both by the initial aspect ratio of the torus, which we take to be $\\sim1$, and by the thickness of the X-ray heated 'skin' of the torus.  More optically thin models produce warm absorber-like spectra for $\\theta\\simeq 40 \\pm 10^\\circ$ , as they potentially provide more partially optically thin gas for evaporation.  The bulk of the gas in this scenario has a terminal velocity of the order of the escape velocity at the inner torus edge. Because of the funnel mechanism part of the gas is re-distributed to lower inclinations and acquires a higher terminal speed, $\\sim 1000 \\,\\rm km\\, s^{-1}$. In a real AGN environment such flow may contain clumps and irregularities and even dust, which are not captured in our studies because of the intrinsic limitations our methods. Accounting for the multiple phases of a gas (on a subcellular level) may reveal this in more detail and may also broaden the range of angles where the warm absorbers appear.  The part of the flow that is shielded by the optically thick part of the torus body can also flow out as part of a torus global expansion.  Thus it strongly depends on the deposition of energy directly to its interior. This problem is related to one of the infrared support of the AGN torus vertical structure against gravitational collapse \\citep{Krolik07} and also requires additional investigation.  \\hbox{} This research was supported by an appointment to the NASA Postdoctoral Program at the NASA Goddard Space Flight Center, administered by Oak Ridge Associated Universities through a contract with NASA, and by grants from the NASA Astrophysics Theory Program 05-ATP05-18. We would like to thank the referee for his/her many constructive comments, which have lead to improvement of the manuscript."
    },
    "0806/0806.3218_arXiv.txt": {
        "abstract": "{Accurate mass, radius, and abundance determinations from binaries provide important information on stellar evolution, fundamental to central fields in modern astrophysics and cosmology.} {We aim to determine absolute dimensions and abundances for the three F-type main-sequence detached eclipsing binaries \\AD, \\VZ, and \\WZ\\ and to perform a detailed comparison with results from recent stellar evo\\-lu\\-tio\\-nary models.} {$uvby$ light curves and $uvby\\beta$ standard photometry were obtained with the Str\\\"omgren Automatic Telescope at ESO, La Silla, radial velocity observations at CfA facilities, and supplementary high-resolution spectra with ESO's FEROS spectrograph. State-of-the-art methods were applied for the analyses: the EBOP and Wilson-Devinney binary models, two-dimensional cross-correlation and disentangling, and the VWA abundance analysis tool.} {Masses and radii that are precise to 0.5--0.7\\% and 0.4--0.9\\%, respectively, have been established for the components, which span the ranges of 1.1 to 1.4 $M_{\\sun}$ and 1.1 to 1.6 $R_{\\sun}$. The \\feh\\ abundances are from $-0.27$ to $+0.10$, with uncertainties between 0.07 and 0.15 dex. We find indications of a slight $\\alpha$-element overabundance of \\afe$\\sim+0.1$ for \\WZ. The secondary component of \\AD\\ and both components of \\WZ\\ appear to be slightly active. Yale-Yonsai and Victoria-Regina evolutionary models fit the components of \\AD\\ and \\VZ\\ almost equally well, assuming coeval formation, at ages of about 1.75/1.50 Gyr (\\AD) and 1.25/1.00 Gyr (\\VZ). BaSTI models, however, predict somewhat different ages for the primary and secondary components. For \\WZ, the models from all three grids are significantly hotter than observed.  A low He content, decreased envelope convection coupled with surface activity, and/or higher interstellar absorption would remove the discrepancy, but its cause has not been definitively identified. } {We have demonstrated the power of testing and comparing recent stellar evolutionary models using eclipsing binaries, provided their abundances are known. The strongest limitations and challenges are set by $T_{\\rm eff}$ and interstellar absorption determinations, and by their effects on and correlation with abundance results. } ",
        "introduction": "\\label{sec:intro}  Accurate stellar mass and radius data from eclipsing binary systems are valuable empirical test data for stellar evolution models, basically because they are free of any scale dependent calibrations. Additional tests become possible when the effective temperatures for the components can also be determined, most often from well-calibrated photometry. The binaries then also serve as primary distance indicators, e.g. for nearby stars, stellar clusters, and Local Group galaxies, and they define the empirical mass-luminosity relations over a broad mass range. However, the most stringent model tests require that spectroscopic element abundances are available as well, minimising the number of free parameters in the comparison with theory (e.g. Andersen \\cite{ja91}). Unfortunately, the number of binary components with complete and accurate data is still small.  In our efforts to provide such complete data for a larger sample of eclipsing binary systems, F-G type main-sequence stars have been given high priority for two reasons. First, their spectra offer favourable conditions for accurate radial velocity determinations and abundance analyses. Second, such stars cover the mass range within which convective cores begin to develop and affect the observational determination of stellar ages, in particular due to difficulties and shortcomings related to the physics of core overshoot. Moreover, these ages cover the range of interest in studies of the evolution of the Galactic disk (e.g. Holmberg et al. \\cite{holmberg07}).  In this paper we present new analyses of the three detached F-type double-lined eclipsing binaries \\object{\\AD}, \\object{\\VZ}, and \\object{\\WZ}. From earlier studies they are known to have masses in the range 1.15--1.41 $M_{\\sun}$ and radii corresponding to the lower half of the main-sequence band, but better accuracy and - especially - abundance data are needed for meaningful tests of state-of-the art stellar mo\\-dels. We have therefore obtained new radial velocity observations, light curves, standard photometry, and high-resolution spectra covering a wide wavelength range, enabling us to determine all relevant parameters of the systems with high precision.  As pointed out by Popper (\\cite{dmp98a}), and also seen in Fig.~\\ref{fig:debs} \\footnote{{\\scriptsize\\tt http://www.astro.keele.ac.uk/$\\sim$jkt/}}, \\AD\\ is of particular interest for having components with the greatest differences in mass and radius among well-studied binaries in its mass range. Hence, it may provide stringent tests of stellar models, but the masses obtained in two recent analyses of the system disagree by slightly more than their stated errors (Lacy \\cite{lacy97}; Popper \\cite{dmp98a}). Moreover, in both studies the radii were based on reanalyses of the relatively old $B,V$ light curves (normal points) by Zhai et al. (\\cite{zhai82}), which have an accuracy of only about 0.02 mag and show out-of-eclipse variations by about 0.075 mag.  Previous determinations of the dimensions of \\VZ\\ and \\WZ\\ were also based on old material (Popper \\cite{dmp65}) and therefore had an accuracy of only 2--4\\%; see Popper (\\cite{dmp80}). As in the case of \\AD, the components of \\VZ\\ are fairly different, whereas those of \\WZ\\ are nearly identical.  The paper is structured as follows: First, Sect.~\\ref{sec:obsmet} describes our observations and analysis techniques adopted. The spectroscopic and photometric analyses of the indi\\-vi\\-dual systems are then discussed in detail in Sect.~\\ref{sec:specphot}. The resulting new, accurate absolute dimensions and abundances are presented in Sect.~\\ref{sec:absdim} and compared with three recent grids of stellar evolutionary models in Sect.~\\ref{sec:dis}.  Throughout the paper, the component eclipsed at the deeper eclipse at phase 0.0 is referred to as the primary $(p)$ component, the other stars as the secondary $(s)$.  \\begin{figure} \\epsfxsize=90mm \\epsfbox{debs_pub.ps} \\caption[]{\\label{fig:debs} Main-sequence eclipsing binaries with both components in the 1.1--1.5 $M_{\\sun}$ range having masses and radii accurate to 2\\% or better. The nine systems shown are, in order of in\\-crea\\-sing mass of the primary component: \\object{EW\\,Ori}, \\object{UX\\,Men}, \\object{HS\\,Hya}, \\object{V505\\,Per}, \\object{IT\\,Cas}, \\object{V1143\\,Cyg}, \\object{CD\\,Tau}, \\object{DM\\,Vir}, and \\object{BW\\,Aqr}. \\WZ, \\VZ, and \\AD, are shown with star symbols. The full drawn line is the $Y^2$ 0.3 Gyr isochrone for $Z=0.0181$ (Demarque et al. \\cite{yale04}). } \\end{figure} ",
        "conclusions": "\\label{sec:sum}  We have presented precise absolute dimensions and abundances for the three F-type double-lined eclipsing binaries \\AD, \\VZ, and \\WZ. The results, which are based on state-of-the-art analyses of new photometric and spectroscopic observations, were used to test the recent $Y^2$, Victoria-Regina, and BaSTI stellar evolutionary models.  Masses and radii precise to 0.5--0.7\\% and 0.4--0.9\\% have been established for the binary components. This level can only be reached if excellent observations are available and are analysed carefully using adequate tools. Special care was taken to identify and avoid possible systematic errors. Several methods were applied in order to determine reliable and accurate (100--150~K) effective temperatures and interstellar extinctions, still a challenging task whose importance can not be overestimated.  System \\feh\\ abundances were based on numerous lines and were derived relative to the Sun. We have reached precisions between 0.07 and 0.15 dex. Even ha\\-ving many Fe~I and Fe~II lines of different strength and excitation potential available, the strong correlation between abundance and effective temperature is difficult to break. $T_{\\rm eff}$ and microturbulence uncertainties contribute significantly to the final abundance precision. Abundances for other heavy elements, based on fewer lines, were also derived; in one case (\\WZ) indications of a slight $\\alpha$-element overabundance is seen. Comparing the results for the three binaries clearly demonstrates that analyses of disentangled spectra are far superior to those of compo\\-si\\-te spectra, essentially because many more lines become available and blending becomes less of an issue.  \\AD\\ has components with the greatest difference in mass and radius among well-studied F-type binaries (1.41~$M_{\\sun}$, 1.61~$R_{\\sun}$; 1.21~$M_{\\sun}$, 1.22~$R_{\\sun}$) and an abundance of \\feh\\,$=+0.10$. Slight surface activity is present on the secondary component. The system is fitted well by $Y^2$ and Victoria-Regina evolutionary models for identical component ages (as determined from masses and radii) of 1.75 and 1.5 Gyr, respectively. The fact that the se\\-con\\-dary appears marginally cooler than the corresponding evolutionary track may be related to mild surface activity. BaSTI models predict different ages of 2.25 and 1.75 Gyr for the components, and therefore do not fit \\AD\\ as well.  \\VZ\\ has somewhat different components (1.27~$M_{\\sun}$, 1.31~$R_{\\sun}$; 1.14~$M_{\\sun}$, 1.11~$R_{\\sun}$) and an abundance of \\feh\\,$=-0.20$. No signs of surface activity are seen. $Y^2$ and Victoria-Regina evolutionary tracks fit \\VZ\\ well. Nearly identical component ages (as determined from masses and radii) of 1.25 Gyr are obtained from the $Y^2$ isochrones, whereas the Victoria-Regina calculations give 1.0 and 0.7 Gyr for the primary and secondary. An even larger difference, 1.25 versus 0.75 Gyr, is found when using the BaSTI models.  \\WZ\\ consists of two nearly identical components (1.22~$M_{\\sun}$, 1.41~$R_{\\sun}$). The \\feh\\ abundance is $-0.27$, and as mentioned above a slight $\\alpha$-element overabundance is possible. Photometry and spectroscopy suggest low-level surface activity on both components. Contrary to the case for \\AD\\ and \\VZ, neither of the three model grids are able to fit the observations for \\WZ; models are significantly hotter than observed. We tentatively conclude that this is caused either by \\WZ\\ anomalies such as a low He content or decreased envelope convection, or by underestimated interstellar reddening. Further observations will be required to resolve the issue.  We believe our study is among the most detailed carried out to date for main-sequence eclipsing binaries, and that it sets new standards for critical tests of stellar evolutionary models based on accurate binary data. With only three systems available, it is premature to draw firm conclusions and suggest specific model shortcomings and/or improvements. We find, however, better agreement with the $Y^2$ and Victoria-Regina models than with the BaSTI models. In contradiction to this, Tomasella et al. (\\cite{asiago07}) recently found that dedicated BaSTI models, including core overshoot and both helium and heavy element diffusion, represent the components of the young $\\sim1.25 M_{\\sun}$ system V505\\,Per well. Their study is based on new absolute dimensions of extraordinary high formal precision.  As shown in Fig.~\\ref{fig:debs}, several additional F-type binaries are available, and we are presently working on re-analyses for some of them, including abundance determinations. We are also conducting a larger program on new F-G type systems exhibiting various levels of surface activity (Clausen et al. \\cite{jvcetal01}), and we expect some very useful insights to come from these studies, and from a parallel program on eclipsing binaries in clusters spanning a wide range in age and metallicity (e.g. Grundahl et al. \\cite{fgj06})."
    },
    "0806/0806.4262_arXiv.txt": {
        "abstract": "{} {We present an optical investigation of the Abell~85 cluster filament ($z=0.055$) previously interpreted in X-rays as groups falling on to the main cluster. We compare the distribution of galaxies with the X-ray filament, and investigate the galaxy luminosity functions in several bands and in several regions.  We search for galaxies where star formation may have been triggered by interactions with intracluster gas or tidal pressure due to the cluster potential when entering the cluster. } {Our analysis is based on images covering the South tip of Abell 85 and its infalling filament, obtained with CFHT MegaPrime/MegaCam (1$\\times$1 deg$^2$ field) in four bands ($u^*, g', r', i'$) and ESO 2.2m WFI (38$\\times$36~arcmin$^2$ field) in a narrow band filter corresponding to the redshifted \\ha\\ line and in an $R_C$ broad band filter. The LFs are estimated by statistically subtracting a reference field.  Background contamination is minimized by cutting out galaxies redder than the observed red sequence in the $g'-i'$ versus $i'$ colour-magnitude diagram.  } { The galaxy distribution shows a significantly flattened cluster, whose principal axis is slightly offset from the X-ray filament.  The analysis of the broad band galaxy luminosity functions shows that the filament region is well populated. The filament is also independently detected as a gravitationally bound structure by the Serna \\& Gerbal (1996) hierarchical method.  101 galaxies are detected in the \\ha\\ filter, among which 23 have spectroscopic redshifts in the cluster, 2 have spectroscopic redshifts higher than the cluster and 58 have photometric redshifts that tend to indicate that they are background objects. One galaxy that is not detected in the \\ha\\ filter probably because of the filter low wavelength cut but shows \\ha\\ emission in its SDSS spectrum in the cluster redshift range has been added to our sample. The 24 galaxies with spectroscopic redshifts in the cluster are mostly concentrated in the South part of the cluster and along the filament.  } {We find a number of galaxies showing evidence for star formation in the filament, and all our results are consistent with the previous hypothesis that the X-ray filament in Abell~85 is a gravitationally bound structure made of groups falling on to the main cluster. } ",
        "introduction": "\\label{sec:intro}  Cosmological simulations of the large-scale structure of the Universe display the filamentary nature of the large-scale galaxy distribution (e.g. \\citealp{Springel+05}), and the observed large-scale galaxy distribution is consistent with this picture, even though filaments are more difficult to see in redshift space (e.g. \\citealp{Pimbblet05}). X ray observations of the nearby rich cluster \\object{Abell~85} highlight a filament of hot gas extending towards the South East to near one virial radius  \\citep{Durret+98,Durret+03,DLNF05}.   One would obviously like to know if this X-ray filament can be traced in the galaxy distribution. If so, one would expect that this filament would be a preferential route for the infall of galaxies onto the cluster, even within the virialised region of the cluster.  The influence of infall is not always well understood, except for a few clusters, such as e.g. Coma \\citep{Adami+07}. One would also like to know if the filament region follows the same morphology-density relation as seen in clusters \\citep{Dressler80,Dressler+97}, or whether the filament constitutes a special environment. Similarly, do the galaxies in the filament dipslay the same specific rates of star formation as seen in other cluster regions of the same density, or is the star formation enhanced or quenched? Indeed, star formation can be triggered when the groups of the filament enter the cluster or dense areas, due to environmental effects such as ram pressure from the intracluster gas or tidal pressure due to the cluster potential (\\citealp{Bekki99}).   The galaxy population in the X-ray filament is easily traced in maps of the projected distribution of galaxies up to a given apparent magnitude limit and with a selection in redshifts to remove obvious cluster outliers. Alternatively, galaxy luminosity functions in several wavebands are a good tool to sample the history of the faint galaxy population (e.g. \\citealp{Adami+07} and references therein) including star formation history, evolutionary processes and environmental effects. In particular, the faint-end slopes of galaxy luminosity functions (LFs) in clusters of galaxies have been observed in some cases to vary with clustercentric distance and are expected to be influenced by physical processes (mergers, tides) affecting cluster galaxies (as summarized e.g. by \\citealp{Boue+08}, hereafter B08).  The \\ha\\ line is a good indicator of star formation and has been detected in a number of galaxies in nearby clusters.  The first pioneering work on this topic was due to \\citet{MWI88}, \\citet{MW93}, and \\citet{MWP98}, who performed the first \\ha\\ surveys in a sample of clusters with an objective prism. Based on this survey, \\citet{MWP98} and \\citet{MW00} analyzed tidally induced star formation in several clusters; they found spatial variations, both within a cluster and from one cluster to another: starburst emission in spirals increases from regions of lower to higher density, and from clusters with lower to higher central galaxy space density.  \\cite{MW05} were then able to show that the frequency of emission line galaxies (ELGs) is similar for field and cluster galaxies of all types, and that for galaxies of a given morphological type the fraction of ELGs is independent of environment.  A large \\ha\\ survey was performed on several nearby clusters by \\cite{BIVG02} and \\cite{Gavazzi+02,Gavazzi+06}. They analyzed several trends with radius and found in particular that luminous galaxies show a decrease in their average \\ha\\ equivalent width in the inner $\\sim 1$ virial radius, while low-luminosity galaxies do not show this trend.  Large \\ha\\ surveys have also allowed to estimate \\ha\\ luminosity functions and star formation rates in some of these clusters (\\citealp{IglesiasParamo+02,Umeda+04}). Observations of clusters in \\ha\\ have also revealed some interesting features. For example, a few \\ha\\ tails and filaments as well as intracluster HII regions have been detected in a few clusters such as Abell~1795 (\\citealp{CSF05}), Coma (\\citealp{YKYF07}) or Abell~3627 (\\citealp{SDV07}).  A starbursting compact group was also found to be falling on to Abell~1367, where complex trails of ionized gas behind the galaxies were detected (\\citealp{Cortese+06}).  We present here a detailed optical analysis of the filament region of Abell~85.  This cluster is at a redshift of 0.055 and shows a very complex structure in X-rays, with a main cluster, a South blob and an extended filament (discovered in X-rays) at least 4~Mpc in length. Based on ROSAT PSPC and XMM-Newton data, evidence was found for several merging episodes, one of these still ongoing, as suggested by the interpretation of the X-ray filament as groups falling on to the main cluster \\citep{Durret+98,Durret+03,DLNF05}. However, the optical properties of the galaxies composing the X-ray filament have not been analyzed until now; they may give us clues on the physical properties of this filament and on the likelihood of the merging scenario described above.  We have obtained two sets of data: ESO 2.2m WFI 38$\\times$36~arcmin$^2$ images in a narrow band filter corresponding to the wavelength of H$\\alpha$ at the cluster redshift and in a broad band $R_C$ filter to subtract the continuum contribution, covering the South half of Abell~85 and its filament, and deep 1$\\times$1 deg$^2$ field images obtained at CFHT with MegaPrime/MegaCam in four bands ($u^*$$g'$$r'$$i'$) covering the South tip of Abell 85 and the infalling filament.  Both sets of data sample the filament feeding the cluster from the Southeast, and the impact region where the filament is believed to be hitting the cluster itself (this impact region is indeed hotter in X-rays). The virial radius, defined as the radius where the mean mass density is 100 times the critical density of the Universe, is $0\\fdg65 $, derived by extrapolating the radius of overdensity 500 given by \\cite{DLNF05}. Thus, our Megacam images (not centered on the cluster) cover part of Abell~85 and more distant regions, well beyond the virial radius. The Sloan Digital Sky Survey (SDSS) covers the region of Abell~85.  We have retrieved all the redshifts available in the SDSS to build a large redshift catalogue for the region of Abell~85, as well as all the galaxy spectra in the region covered by our WFI data.  The paper is organized as follows. We present our Megacam and WFI data and data reduction in Section 2. In Section 3, we describe our results on H$\\alpha$ imaging and discuss the spatial distribution and properties of H$\\alpha$ emitting galaxies, together with properties derived from the SDSS data.  In Section 4, we present our results obtained for the LF in the four broad photometric bands.  In Section 5, we discuss our results concerning the LFs in terms of large scale environmental effects on the cluster galaxy populations.  Final conclusions are drawn in Section 6.  We assume a distance of 242.2 Mpc to Abell 85 ($H_0$ = 71 km s$^{-1}$ Mpc$^{-1}$, $\\Omega _m$ = 0.27 and $\\Omega_ {\\Lambda}$ = 0.73). The distance modulus is 36.92 and the scale is 1.055 kpc arcsec$^{-1}$. We give magnitudes in the AB system. At this distance, the Megacam field of view corresponds to 3.8$\\times$3.8 Mpc$^2$, while the virial radius  is 2.5~Mpc. ",
        "conclusions": "We have analyzed an \\ha\\ image and broad band images covering the South of Abell~85 and its filament, which was previously discovered in X-rays (\\citealp{Durret+98}).  The Galaxy Luminosity Functions in the South area of the cluster Abell~85 (including the impact region, where the groups constituting the filament hit the main cluster) and in the filament show the existence of a rich population of galaxies. The overall galaxy distribution in A85 is flattened with a principal axis that is parallel to the axis separating the filament seen in X-rays with the cluster center. All our results are consistent with the previous interpretation of the filament being made of groups falling onto the main cluster."
    },
    "0806/0806.2863_arXiv.txt": {
        "abstract": "We assess the potential of nuclear starburst disks to obscure the Seyfert-like AGN that dominate the hard X-ray background at $z \\sim 1$. Over 1200 starburst disk models, based on the theory developed by Thompson \\etal, are calculated for five input parameters: the black hole mass, the radial size of the starburst disk, the dust-to-gas ratio, the efficiency of angular momentum transport in the disk, and the gas fraction at the outer disk radius. We find that a large dust-to-gas ratio, a relatively small starburst disk, a significant gas mass fraction, and efficient angular momentum transport are all important to produce a starburst disk that can potentially obscure an AGN. The typical maximum star-formation rate in the disks is $\\sim 10$~\\sfr. Assuming no mass-loss due to outflows, the starburst disks feed gas onto the black hole at rates sufficient to produce hard X-ray luminosities of $10^{43}$--$10^{44}$~erg~s$^{-1}$. The starburst disks themselves should be detectable at mid-infrared and radio wavelengths; at $z=0.8$, the predicted fluxes are $\\sim 1$~mJy at 24\\micron\\ and $\\sim 10$--$30$~$\\mu$Jy at 1.4~GHz. Thus, we predict a large fraction of radio/X-ray matches in future deep radio surveys. Unfortunately, both the 24\\micron\\ and radio fluxes are comparable to those expected from the central AGN. In contrast, the starburst disks should be easily distinguished from AGN in future 100\\micron\\ surveys by the \\textit{Herschel Space Observatory} with expected fluxes of $\\sim 5$~mJy. Any AGN-obscuring starbursts will be associated with hot dust, independent of AGN heating, resulting in observable signatures for separating galactic and nuclear star-formation. This may be an explanation for the small observed $L_{2-10\\ \\mathrm{keV}}/\\nu L_{\\nu} (6\\ \\mu \\mathrm{m})$ ratios observed from both $z \\sim 0$ and $z \\sim 1$ AGN. Finally, because of the competition between gas and star-formation, nuclear starbursts will be associated with lower-luminosity AGN. Thus, this phenomenon is a natural explanation for the observed decrease in the fraction of obscured AGN with luminosity. ",
        "introduction": "\\label{sect:intro} All Active Galactic Nuclei (AGN) are powered by accretion onto a central supermassive black hole \\citep{lb69,ss73,pri81,balbus03}. However, the observational manifestation of this accretion can vary dramatically from object to object. For example, the spectral energy distributions (SED) of rapidly accreting quasars \\citep{elv94} are very different from the Galactic Center source Sgr A* \\citep{mf01}, or the black hole at the center of M87 \\citep{ho99}. In this case, the different observational properties are a result of the two distinct mechanisms through which energy is liberated in radiatively efficient (for the quasars) and inefficient (for Sgr A* and M87) accretion flows. Another important parameter is the orientation of the black hole-accretion disk system to the line of sight. As a result, blazars, AGN that are viewed down the axis of a radio jet, show significantly different SEDs and variability properties than all other AGN \\citep{bach07}. Finally, obscuration along the line-of-sight is another key parameter in determining the observational characteristics of an AGN. The vast majority of local optically-selected AGN show evidence for being obscured \\citep[e.g.,][]{mr95,hfs97}, and these are called Type 2 AGN, while the unobscured sources are defined as Type 1 AGN. Observationally, the obscuration is revealed as significant soft X-ray absorption \\citep[e.g.,][]{toz06}, the disappearance of the optical broad permitted lines \\citep{kw74}, and a reddened continuum \\citep{wilkes05}.  Despite its near ubiquity, the origin of the obscuration around AGN is still unclear. Infrared emission from dust near the sublimation temperature \\citep[e.g.,][]{rie78,neu79,bar87,san89} and variability in the X-ray absorbing column density \\citep{ris02,ris05} both point to a location $\\sim$1~pc from the central engine, although it is difficult to rule out contributions from the host galaxy at larger scales \\citep{mr95,bwm03,rig06}. It is also possible that the origin of the obscuration may arise through a variety of processes for AGN with different luminosities or at different points in their evolution \\citep{bal06b}.  Perhaps the most significant challenge facing the pc-scale absorber (or obscuring `torus' in the AGN unification model; \\citealt{ant93}) is to account for the observed 4-to-1 Type 2/Type 1 ratio of AGN required to fit the hard X-ray background \\citep{gch07}. This fact requires the obscuring material to cover $\\sim$80\\% of the sky as observed from the central black hole, and thus have a structure so that its vertical scale height is of the same magnitude as its radius, i.e. $H/r \\sim 1$. \\citet{kro07} and \\citet{sk08} have argued that the AGN emission absorbed and re-radiated in the infrared by dust in the torus can provide enough pressure support to inflate a structure to the required scale. Similarly, \\citet{cqm07} suggested that X-ray heating of the outer accretion disk may be sufficient to result in a $H/r \\sim 1$ geometry. Another explanation, and the one considered here in more detail, is feedback from a nuclear starburst disk \\citep{fab98,wn02,tqm05}.  There are many reasons to consider starburst disks as a source of obscuration in AGN. At the most basic level, gas to fuel the black hole must be funneled toward the center of the galaxy and is therefore likely to cause star formation en route. In fact, recent adaptive optics observations by \\citet{davies07} have shown the presence of a post-starburst population at parsec scales in 9 local Seyfert galaxies, indicating that this may indeed be a common phenomenon. In addition, recent models of AGN tori have suggested that the obscuring region may be clumpy \\citep{nie02,dvb05,honig06}, a common trait of star-forming regions in starburst galaxies \\citep{fs01,skv07}. Type 2 AGN also seem to be more correlated with star-forming activity in general than the unobscured Type 1 objects \\citep{cft95,gdhl01,lwh01,khi06,lacy07}. There is also evidence for a correlation between the observed intensity of star-formation in the host galaxy and the AGN nuclear luminosity \\citep{shi07,wki08}.  A more profound aspect of star-formation as a source of AGN obscuration is that it provides a connection between the black hole-accretion disk environment and the host galaxy \\citep[e.g.,][]{kw08}. It is now accepted that the correlations between the black hole mass and bulge properties in early-type galaxies \\citep{mag98,fm00,geb00,tre02} are evidence for a significant connection between the growth of the black hole and the build up of a galactic bulge \\citep[e.g.,][]{sr98,fab99,kh00,wl03,mqt05,dsh05,fvg08}. Therefore, it is expected that as the black hole and galaxy bulge are growing, there will be significant accretion on to the black hole that is obscured by star formation in the host galaxy \\citep[e.g.,][]{hop05}. Indeed, the hard X-ray background, emitted by accreting black holes throughout the history of the Universe, has a very hard spectrum, indicative of being dominated by obscured AGN \\citep{sw89,dm04,gch07}. Interestingly, the redshift distribution of the obscured AGN that dominate the X-ray background peaks at $z \\sim 1$, very similar to the peak in the cosmic star-formation history \\citep{toz01,bar02,hopk04,bar05}. This result may imply a connection between the obscuring material and the evolution of the host galaxy \\citep{fab98,fra99,bem06}. In fact, there is tentative evidence that the fraction of obscured AGN does increase with $z$ \\citep{bem06,tu06}. Thus, if star-forming disks do provide a significant amount of AGN obscuration at $z \\sim 1$ (where the X-ray background sources are most common) then by studying the properties and evolution of the absorbing material, we can directly probe the evolution of the underlying host galaxy.  In this paper, we make use of the analytic starburst disk models developed by \\citet{tqm05} to investigate the properties of starburst disks as a source of obscuring material in the Type~2 AGN found at $z \\sim 1$. In the paper by \\citet{tqm05}, the models were able to successfully describe many of the observed proprieties of Ultra-Luminous Infrared Galaxies (ULIRGs), and it was noted that under certain conditions, the photosphere of the starburst disks may reach $H/r \\sim 1$. Here, we adapt the model to consider the much less intense star-forming disks that might be expected to be found around the Seyfert-like luminosity AGN that dominate the population at these redshifts \\citep{ueda03,bar05}. We search for the range of physical parameters these starburst disks must have in order for them to provide significant obscuration. Both radio and far-IR fluxes are predicted to determine if the disks can be (or have already been) detected by current observational surveys. The next section provides a brief review of the \\citet{tqm05} models, and then describes how the theory was altered and applied to this problem. Section~\\ref{sect:results} presents the results of our calculations, and describes both the physical and observational properties of starburst disks that may obscure an average AGN. The results are discussed in Section~\\ref{sect:discuss} and we present our conclusions in Section~\\ref{sect:concl}. When necessary, the following $\\Lambda$-dominated cosmology is assumed in this paper: $H_0=70$~km~s$^{-1}$~Mpc$^{-1}$, $\\Omega_{\\Lambda}=0.7$, and $\\Omega_{m}=0.3$ \\citep{spe03}. ",
        "conclusions": "\\label{sect:concl} In this paper we proposed that nuclear starburst disks may be an important contributor to AGN obscuration, especially at $z \\sim 1$ where there is a large population of obscured AGN. The analytic model of starburst disks developed by \\citet{tqm05} was used to explore the properties of such disks and determine if they were a viable and general method to obscure accreting black holes. We found that a range of conditions can produce potentially AGN obscuring bursts of star-formation on pc scales. These include a high dust-to-gas ratio and a relatively small size scale of the disk, both of which seem to be consistent with current constraints. The pc-scale bursts of star-formation are different from traditional star-forming regions since they are produced in a region with a temperature $\\sim 1000$~K. This region is not heated by an AGN, but is a result of the assumption of a $Q=1$ stable star-forming disk in a black hole potential. Thus, the observational signature of these nuclear star-forming regions will be additional warm/hot dust emission above what is required by AGN heating. This effect may be an explanation for the observed $L_{2-10\\ \\mathrm{keV}}/\\nu L_{\\nu} (6\\ \\mu \\mathrm{m})$ ratios seen from both local and high-redshift AGN. In addition, the competition between star-formation and gas accretion results in a natural explanation for the decrease in the fraction of obscured AGN with luminosity. As an AGN-obscuring nuclear starburst will consume gas that may have been destined for accretion onto a black hole, this model predicts that nuclear starbursts should be closely associated with lower luminosity AGN. In contrast, high luminosity AGN require a significant gas supply, so a nuclear starburst would be an hindrance to fueling the black hole.  The observational signatures of these disks may be most easily found in future deep radio and far-IR surveys. Assuming the radio-IR correlations hold for the circumnuclear starbursts, the expected 1.4~GHz fluxes from these disks are $\\sim 10$~$\\mu$Jy at the redshifts where the fraction of obscured AGN is near its peak. The 100\\micron\\ fluxes are predicted to be $\\sim 1$--$10$~mJy, which will be detectable in future \\textit{Herschel} surveys.  Nuclear starbursts may prove to be a compelling method of studying the symbiotic relationship between black holes and their host galaxies in the era of rapid evolution of both populations.  While qualitatively interesting, this model needs to be explored further to make robust quantitative predictions. A proper estimation of the vertical structure of the disk is required to estimate covering fractions. Future work should also include incorporating the feedback from the central AGN. Another important future question is to determine the nature and observable properties of the starburst remnants."
    },
    "0806/0806.0866_arXiv.txt": {
        "abstract": "We present results from a mosaic of nine {\\it Chandra} observations of M86 and the surrounding field.  We detect three main diffuse components: the Virgo ICM at $\\sim$2.4~keV, the extended halo of M86 at $\\sim$1.2~keV, and the cooler central and stripped gas of M86 at $\\sim$0.8~keV.  The most striking feature is a long tail of emission, which consists of a plume $\\sim$ 4\\arcmin north of M86 and two main extensions emanating from the plume. Based on the morphology and temperature structure of the tail, we conclude that it is formed by ram pressure stripping of M86 as it falls into the Virgo cluster and interacts with the Virgo ICM, in agreement with earlier work. The tail is 150~kpc in projection, and a simple estimate gives a lower limit on the true length of the tail of 380~kpc, making this the longest ram pressure stripped tail presently known. The total gas mass in the plume ($\\sim 7 \\times 10^8 \\, {\\rm M_\\odot}$) and tail ($\\sim 1 \\times 10^9 \\, {\\rm M_\\odot}$) is about three times that in the core of M86, which supports the scenario where most of the gas was stripped rapidly and recently. The projected position of the plume can be understood if M86 has an aspherical potential, as suggested by optical isophotes.  Ram pressure stripping from an aspherical potential can also explain the split ``double tails'' seen in M86 and in other Virgo cluster galaxies in the field. The large line-of-sight velocity of M86 (1550~\\kms with respect to M87), its position relative to the Virgo cluster, and the orientation of the tail tightly constrain its orbital parameters. The data are inconsistent with a radial orbit, and imply inner and outer turning radii of $\\ri \\approx 300$~kpc and $\\ro \\ga 8.8$~Mpc, indicating that M86 is, at best, only weakly bound to the Virgo cluster. ",
        "introduction": "\\label{sec:intro}  M86 (NGC~4406) is a bright elliptical (E3/S0) galaxy in the Virgo cluster of galaxies.  It is the dominant member of one of the larger subgroups within Virgo (Binggeli \\etal\\ 1993; B\\\"ohringer \\etal\\ 1994; Schindler \\etal\\ 1999). Its line-of-sight velocity relative to M87, the dominant member of the Virgo cluster, is -1550~km~s$^{-1}$, much higher than the average cluster velocity dispersion (Smith et al.\\ 2000).  The X-ray surface brightness distribution is unusual, with a large ``plume'' extending to the northwest from M86, which was first noticed as part of a survey of Virgo cluster galaxies undertaken with the {\\it Einstein Observatory} (Forman et al. 1979). The galaxy has an optical asymmetry that extends in a direction similar to the direction of the plume (Nulsen \\& Carter 1987; Mihos \\etal\\ 2005). Several authors have interpreted this plume as arising from ram pressure stripping due to strong interactions with the Virgo ICM (Forman \\etal\\ 1979; Fabian \\etal\\ 1980; Takeda \\etal\\ 1984; Knapp \\etal\\ 1989; Bregman \\& Roberts 1990; White \\etal\\ 1991; Rangarajan \\etal\\ 1995).  Elmegreen \\etal\\ (2000) find dust streamers in the core of M86 that connect to the nucleated dwarf galaxy VCC~882, and suggest a recent interaction between the pair that may have contributed to the asymmetry of M86's optical isophotes.  At X-ray energies, M86 has been observed by  {\\it Einstein} (Forman \\etal\\ 1979; White \\etal\\ 1991), {\\it Ginga} (Takano \\etal\\ 1989), {\\it EXOSAT} (Edge 1990), {\\it ASCA} (Matsushita \\etal\\ 1994), and {\\it ROSAT} (B\\\"{o}hringer \\etal\\ 1994; Rangarajan \\etal\\ 1995). Recently, Finoguenov \\etal\\ (2004) presented {\\it XMM-Newton} observations of M86 and concluded that its unusual morphology is due to an interaction with an X-ray filament rather than the Virgo ICM.  However, their conclusions were based on the large separation between M86 and M87 along the line of sight of $2.4 \\pm 1.4$~Mpc reported by Neilsen \\& Tsvetanov (2000), which is inconsistent with the more recent result from Mei \\etal\\ (2007) who find a separation of $0.4\\pm 0.8$~Mpc.  We report here on a mosaic of nine {\\it Chandra} observations of M86 and the surrounding field, totaling 240~ksec of exposure. The observations and data reduction techniques are described in \\S~\\ref{sec:obs}.  The X-ray image is presented in \\S~\\ref{sec:img}, and results on temperature and abundance structure from spectral analysis are given in \\S~\\ref{sec:spec}.  In \\S~\\ref{sec:discuss}, we discuss some of the more interesting features of M86, and place constraints on its orbit.  In particular, we give new results on the extent of the stripped tail, and measure the density profile of an X-ray brightness edge seen to the southeast.  Our results are summarized in \\S~\\ref{sec:summary}.  We assume a distance to the Virgo cluster of 16 Mpc throughout, consistent with the latest results from Mei et al.\\ (2007), which gives a scale of 0.08 kpc/\\arcsec\\ for $\\Omega_0 = 0.3$, $\\Omega_{\\Lambda} = 0.7$, and $H_0 = 70$~km~s$^{-1}$~Mpc$^{-1}$.  All error ranges are 90\\% confidence intervals, unless otherwise stated. ",
        "conclusions": "\\label{sec:discuss}  \\subsection{The Ram Pressure Stripped Tail of M86} \\label{sec:tail}  The most striking feature in Figure~\\ref{fig:smoimg} is the long tail of emission extending to the NW from M86.  M86 has a line-of-sight velocity of -244$\\pm 5$~km~s$^{-1}$, while M87's is 1307$\\pm 7$~km~s$^{-1}$ (Smith et al.\\ 2000). Therefore, M86 is traversing the Virgo cluster at $v_{\\rm M86} > 1550$~km~s$^{-1}$ (about Mach 2 for $kT = 3$~keV). The tail naturally forms due to ram pressure stripping of the M86 corona by the Virgo cluster ICM. Finoguenov \\etal\\ (2004) suggested that the tail formed due to interactions with a filament rather than the Virgo ICM, though their result was based on older distance measurements that placed M86 outside the virial radius of the Virgo cluster. If we assume that M86 is bound to the Virgo cluster, its large line-of-sight relative velocity allows us to constrain the separation between M86 and M87. Using the M87 mass profile detailed in \\S~\\ref{sec:orbit}, we find that a free-fall velocity of 1500~km~s$^{-1}$ corresponds to a separation from M87 of 0.5~Mpc (the projected separation is 0.35~Mpc). This is an upper-limit, since the total relative velocity must be at least as large as the line-of-sight relative velocity. This separation is consistent with recent surface brightness fluctuation distances of Virgo cluster galaxies, which give a distance between M86 and M87 of $0.4\\pm 0.8$~Mpc (Mei et al.\\ 2007).  We also can place a lower-limit on the length of the long stripped tail using the mass profile of M87 and the line-of-sight velocity.  The maximum free-fall velocity of M86 from infinity at the projected separation of 0.35~Mpc is about 1680~km~s$^{-1}$.  Using this as an upper-limit on its current 3D velocity, we find that the angle between the direction of motion of M86 and the line-of-sight is $\\theta \\la 23^{\\circ}$.  Assuming that the stripped tail is aligned with M86's current direction of motion, and given that the length of the tail in the plane of the sky is 150~kpc (0.51$^{\\circ}$), we find a lower-limit on the actual length of the tail of $L_{\\rm tail} \\ga 380~$kpc, making this the longest ram pressure stripped tail presently known.   The long stripped tail originates in the plume of diffuse emission located directly north of M86.  For a continuous stripping process, one would expect the tail to extend from M86 itself.  This separation between M86 and the plume can be explained by rapid recent stripping, in which a significant fraction of the remaining gas in M86 is rapidly stripped when the ram pressure stripping condition is met.  A detailed discussion of this process is given in \\S~\\ref{sec:displace}   As seen from Figure~\\ref{fig:regmap} and Table~\\ref{tab:spectra}, the stripped tail shows a general trend of cooler gas in and near the plume (with temperatures in the 0.8--0.85~keV range) and warmer gas at the tip (in the 0.9--1.2~keV range).  The temperature structure of the tail is consistent with a ram pressure stripping model, where, as M86 falls into the Virgo cluster, the hotter, higher entropy group gas is stripped first due to interactions with the Virgo ICM, followed by the cooler, lower entropy M86 ISM, which is removed rapidly once the stripping condition is met.   \\subsection{The Extended M86 Halo} \\label{sec:halo}  The {\\it Chandra} and {\\it ROSAT} X-ray images (Figure~\\ref{fig:rosat}) show an extended halo associated with M86, with a sharp brightness edge to the southeast in the direction of M87. Since M86 is traversing the Virgo cluster supersonically (see \\S~\\ref{sec:tail}) we expect a shock to be driven in the Virgo ICM, possibly producing a brightness edge similar to that seen in the southeast. We attempted to measure the density jump associated with the shock in the following way. We extracted the {\\it Chandra} 0.5--2.0 keV surface brightness profile of this edge in two roughly equal sectors since the morphology of this feature is irregular in the south (see Figure~\\ref{fig:rosat}), and since the gas temperature differs in these two regions (see \\S~\\ref{sec:tmap}).  The emission measure profile for the northern sector is shown in Figure~\\ref{fig:nedge}.  Distance is measured from the center of curvature of the apparent edge, which is 10~kpc (2.16\\arcmin) east of M86.  This profile was fit with a spherical gas density model consisting of two power laws.  The free parameters were the normalization, the inner ($\\alpha$) and outer ($\\beta$) slopes, the position of the density discontinuity ($r_{\\rm break}$), and the amplitude of the jump ($A$).  We assumed that the gas is isothermal with $kT = 1$~keV and that the abundance is constant at 30\\% solar, consistent with results from spectral fits (see \\S~\\ref{sec:spec}). For the best fit model (see Figure~\\ref{fig:nedge}), we find $\\alpha = 0.49^{+1.14}_{-0.71}$, $\\beta = -0.81^{+0.14}_{-0.12}$, $r_{\\rm break} = 52^{+5}_{-6}$~kpc, and $A = 1.3^{+0.3}_{-0.4}$.  Similar results are found from the southern sector, but with larger errors.  The lack of a well-defined edge is consistent with our findings for the orbit of M86 (see \\S~\\ref{sec:orbit}). In particular, for all of the likely orbits, M86 is moving to the southeast, such that our lines of sight pass through the Mach cone. To see a sharp edge, our line of sight must be tangent to the shock front.    \\subsection{Constraints on the Orbit of M86} \\label{sec:orbit}  Knowing the orbit of M86 is the key to understanding its interaction with the Virgo cluster.  The ram pressure stripped tail reveals the motion of the galaxy on the sky. The length and direction of the tail, together with the large line-of-sight speed of M86, constrain its orbit.  The large luminosity of M86 suggests that it dominates the associated in-falling subgroup (Schindler \\etal\\ 1999).  In the following, we assume M86 is bound to the Virgo cluster. For the purpose of calculation, the gravitational potential of the Virgo cluster is treated as a spherical NFW potential (Navarro, Frenk \\& White 1997), with a virial radius of 1.3 Mpc (Evrard, Metzler \\& Navarro 1996) and a concentration parameter of 4.5 (Neto et al.\\ 2007), appropriate for a $\\sim 3$ keV cluster. The gravitating mass was normalized to match the total mass within 320 kpc of $4.4\\times10^{13}\\ {\\rm M_\\odot}$ (Schindler \\etal\\ 1999; scaled to a distance of 16 Mpc).  For this potential, the escape speed from a radius of 351 kpc (the projected separation of M86 from M87) is 1677 $\\rm km\\ s^{-1}$, not much larger than M86's line-of-sight speed of 1550 $\\rm km\\ s^{-1}$ relative to M87.  The empirical mass distribution of Schindler (1999) gives the same escape speed at 351~kpc as our model potential.  We note that the gravitational potential of the Virgo cluster is not well constrained at distances from the cluster center comparable to or greater than its virial radius.  Furthermore, the truncated NFW potential is highly simplified, ignoring mass beyond the virial radius, departures from spherical symmetry, and the dynamic state of the cluster due to continuing in-fall.  Since some of the limits we derive here are sensitive to the poorly known potential at large radii, they should be treated as indicative rather than quantitative (outside the context of the model). The material in this section is supplemented by a more detailed discussion in Appendix~\\ref{ap:orbit}.   We consider radial orbits first, as suggested by previous studies (e.g., Forman \\etal\\ 1979; White \\etal\\ 1991). In order for the gas tail to point away from the cluster center, the galaxy must be inbound. It can easily be shown that the line-of sight velocity, $\\vlos$, is maximized for some radius $r$ greater than the observed separation $s$ (see Appendix~\\ref{ap:orbit}). For a marginally bound (zero energy) radial orbit in the potential described above, the maximum value of $|\\vlos|$ occurs when M86 is close to twice its projected distance from the cluster center, giving $|\\vlos| = 1205\\ \\kms$, less than the observed value of $1550\\ \\kms$. Thus, unless M86 is significantly unbound from the Virgo cluster (or the potential is incorrect), its line-of-sight speed is inconsistent with radial and nearly radial orbits.  It is convenient to specify more general orbits in terms of their inner and outer turning radii, $\\ri$ and $\\ro$, respectively. We constrain the possible orbits of M86 by placing limits on these parameters.  Coarse limits can be placed on these radii by considering energy arguments alone. Assuming that M86 is bound to the Virgo cluster, we find $\\ri \\la 489$~kpc and $\\ro \\ga 3.5$~Mpc, or $\\simeq 2.7$ times the virial radius of the Virgo cluster (see Appendix~\\ref{ap:orbit}).  These limits can be further restricted by considering the range of possible viewing directions for each point on an orbit.  There can be zero, two, or four possible viewing directions that would place M86 at the observed separation from the cluster and give it the observed line-of-sight velocity (see Appendix~\\ref{ap:orbit}). Figure~\\ref{fig:mb} shows a range of marginally bound orbits ($\\ro = \\infty$), with regions that meet these conditions marked in color.  The range of an orbit where these conditions are met shrinks as $\\ro$ decreases, \\ie, as the orbit becomes more tightly bound.  For the marginally bound orbits, the full range of $\\ri$ for which $\\vlos$ can attain its observed value is $247 < \\ri < 395$ kpc (the range represented in Figure~\\ref{fig:mb}).  For more tightly bound orbits, the acceptable range of $\\ri$ is reduced.   The lower limit on $\\ro$ is increased if we consider the extent of the ram pressure stripped tail.  The projected orbit must extend to at least the distance ($\\sim 100$~kpc) that the tail projects beyond M86 in the direction away from the cluster center. Locations on the orbits where the line-of-sight speed can attain $-1550\\ \\kms$ and these conditions also are met are shown in red in Figure~\\ref{fig:mb}.  We see that the range of possible locations for M86 on these orbits is tightly constrained. We repeated this analysis for a number of values for the outer turning radius.  The range of potential orbits and locations for M86 shrinks with decreasing $\\ro$ and no orbits were found to meet these criteria for $\\ro \\lesssim 8.2$ Mpc (6.3 virial radii).  Figure~\\ref{fig:projmb} shows projections onto the sky of the marginally bound orbits (corresponding to the midpoints of the red regions of Figure~\\ref{fig:mb}) overlaid on the {\\it Chandra} 0.5 -- 2 keV image, with M86 at the representative location for each orbit. It is evident that the possible orbits for M86 can be pruned further. For example, the orbit on the lower right in Figure~\\ref{fig:projmb} does not reach far enough north to produce the remote part of the gas tail.  At the other limit, the position angle of M87 measured from M86 is $116^\\circ$ (east from north), generally eastward of these orbits.  Overdense stripped gas tends to fall towards the cluster center, in that direction, so that some stripped gas can reasonably lie to the east of the orbit.  Stripped gas cannot lie west of the orbit.  Thus the orbits that cross the region to the east of the tail and north of M86, where there is no sign of stripped gas, are unlikely candidates for the orbit of M86.  Note that the orbits in Figure~\\ref{fig:projmb} are not simply related to their inner turning radii.  The orbit corresponding to the smallest value of the inner turning radius appears second from the right at the top of the figure.  At first, orbits for increasing values of $\\ri$ lie to the left of this, but at a value of $\\ri$ approaching (but less than) $s$, the acceptable viewing direction flips from outward to inward (\\ie, the acceptable sign of $a'$ changes, see Appendix~\\ref{ap:orbit}).  At this point, the projected position of the orbit shifts from leftmost to rightmost at the top of Figure~\\ref{fig:projmb}. Whereas the apparent location where the orbit passes through the virial radius was moving outward, it now moves inward as $\\ri$ increases, coming into the field of view for the largest values of $\\ri$ here.  This results in disjoint ranges of possible orbits for M86.  Thus the marginally bound orbits with inner turning radii close to $\\ri \\simeq 260$ kpc and those with inner turning radii in the range $330 \\la \\ri \\la 380$~kpc are best suited to model the orbit of M86.   Applying the same criteria to the more tightly bound orbits yields the collection of possible orbits shown in Figure~\\ref{fig:m86orb}.  We have not attempted to be highly selective or exhaustive, since that requires a well defined model for the formation of the gas trail.  For values of the outer turning radius in the lower end of the acceptable range, the shrinking range of possible orbits and locations excludes orbits with inner turning radii $\\ri > s$.  All of the acceptable orbits for M86 are weakly bound to the Virgo cluster, with the most tightly bound having outer turning radii at $\\simeq 8.8$ Mpc.  All of the possible locations for M86 lie only a little farther from M87 than M86 does in projection.  All are close to the plane of the sky, ranging from 167 kpc closer than M87 to 263 kpc farther than M87 from the Sun at the extremes.  M86 must also be close to the pericenter of its orbit.  The direction of motion of M86 is close to our line of sight, within $16^\\circ$ -- $23^\\circ$ of it, for all of the cases illustrated.  M86 is traversing the Virgo cluster supersonically ($\\vlos = -1550$~\\kms\\ alone is almost twice the sound speed $\\sound$, which is taken to be 850 \\kms). Therefore, we expect a shock to be driven in the Virgo ICM. For the orbits we derive, the angle of inclination is smaller than the opening angle of the Mach cone ($\\simeq 33^\\circ$) and we should not expect shock fronts to be visible on the plane of the sky (\\cf\\ Rangarajan \\etal\\ 1995). This is illustrated in Figure~\\ref{fig:mach}, which shows the marginally bound orbit with $\\ri = 376$~kpc.  Each circle is drawn centered on a point where M86 was at a time $\\delta t$ in the past, with a radius of $\\sound \\, \\delta t$.  This is done at equally spaced times to indicate the shape of the Mach cone. The absence of caustics in Figure~\\ref{fig:mach} shows that we would not see the shock front in projection.  (While this statement is accurate, except possibly in the small region where the shock front is highly supersonic, the diagram is schematic.  For example, if M86 were coming directly towards the Earth, radii of sections of the Mach cone would be $(1 - \\sound^2/v^2)^{-1/2} \\simeq 1.20$ larger than drawn.)  Our lines of sight through the compressed gas behind the shock are longest where the circles pile up to the southeast of M86, consistent with the enhancement in X-ray surface brightness seen $\\sim3'$ southeast of M86 (Figure~\\ref{fig:mach}).   \\subsection{Displacement of Stripped Gas}\\label{sec:displace}  The prominent plume of gas lying 3\\arcmin -- 4\\arcmin ($\\sim 16$ kpc) north of M86, noted previously (\\eg, Forman et al. 1979), appears to be physically separated from the dense gas remaining in the galaxy and it is centered some distance east of the orbits illustrated in Figure~\\ref{fig:m86orb}. The removal of a significant fraction of the ISM in a single blob is expected from rapid ram pressure stripping (Takeda \\etal\\ 1984). A highly simplified model for ram pressure stripping treats the plume as a single particle subject to gravity and drag due to its motion through the surrounding gas. Allowing for buoyancy, the net force on the plume due to gravity is $(\\rho - \\rhoe) V \\gvec$, where $\\rho$ is its density, $V$ is its volume, $\\rhoe$ is the density of the ambient ICM, and $\\gvec$ is the local acceleration due to gravity.  The drag force on the plume is $-\\cdrag A \\rhoe v \\vvec$, where $\\cdrag$ is the drag coefficient, $A$ is the plume cross section, $\\vvec$ is its velocity, and $v = |\\vvec|$. Thus, the net acceleration of the plume is $\\avec = [1 - \\rhoe/\\rho] \\gvec - [\\cdrag\\rhoe A / (\\rho V)] v \\vvec$.  Two of the main parameters of this model are the density contrast, $\\rhoe/\\rho$, and the factor $\\cdrag A / V$.  For the density of the ICM, $\\rhoe$, we use the beta model of Schindler \\etal\\ (1999; $\\beta = 0.47$, core radius $= 2.7'$).  On the grounds that the density of the plume is likely to be lower now (due to its stripping and ejection from the confining potential of M86), we set its density contrast at the projected radius of M86 to be 18 (\\cf\\ $\\sim 16$ from the numbers above, see \\S~\\ref{sec:dspec}) and treat the density of the plume as constant.  The other factor is determined by treating the plume as a sphere with a constant radius of 10 kpc, with $\\cdrag = 0.75$.  As above, this radius is a little larger than the observed radius.  To simulate stripping, a gas blob is placed at the location and velocity of M86 at the time it is released from the cluster virial radius.  The orbit of the blob is followed, subject to the gravitational acceleration of the cluster and the moving galaxy.  This model cannot account for the current location of the plume if the gravitational potential of M86 is spherical.  The path of the blob is determined by the competition between gravity and drag. If the drag is large, the blob is ejected from M86 early on its orbit, and slows quickly until it is falling towards the cluster center at its terminal speed, and ultimately lies farther away from M86 than observed. Its early ejection also leaves it too far back along the orbit of M86.  Reducing the relative significance of the drag causes the blob to be ejected later, bringing it closer to its observed position but not far enough from the orbit in the direction of the cluster center.  This issue can be resolved if the gravitational potential of M86 is aspherical.  Consider a small gas blob in an aspherical galaxy moving at an inclined angle through the ICM, illustrated schematically in Figure~\\ref{fig:inclined}.  While ram pressure is insufficient to eject the blob, it is driven to an equilibrium position where ram pressure is balanced by gravity.  Since the direction of the gravitational field must oppose the drag, the external flow must be perpendicular to the equipotential surface.  The equilibrium position therefore lies at a location away from the axis passing through the center of the galaxy and parallel to the flow (\\eg, near the point labeled ``equipotential'' in Figure~\\ref{fig:inclined}).  As the drag starts to overwhelm gravity and displaces the blob in the direction of the flow, the equipotentials the blob encounters tilt, so that it is subject to a component of the gravitational force directed away from the axis.  Because forces in the direction of the flow are nearly balanced, this off-axis component of gravity drives the blob farther away from the axis of the flow.  Thus, a blob stripped from such a galaxy tends to emerge away from the axis of the flow. Treating the gas as a fluid, interstellar gas that is pushed inward by ram pressure along the leading edge of the galaxy flows out preferentially along the long axis of the potential, in the direction of the weakest gravitational force.  The model for the aspherical gravitational potential of M86 used here has the form $\\Phi (\\rvec) = F (w)$, where $F(w) = \\psi \\ln [(1 + w / c) / (1 + w)] / w$ is a minor modification of the standard NFW form ($\\psi \\ln [1 / (1 + w)] / w$ in the same notation).  The extra factor of $1 + w / c$, where $c$ is the concentration parameter, makes the total mass converge, avoiding the need to truncate the mass distribution at the virial radius (which would create problems for our fourth order integrator).  The coordinate $w = \\sqrt{(x^2 + y^2)/a^2 + z^2/b^2 + 1}$, where $a$ is the NFW potential scale length and $b$ is chosen to give the desired ellipticity.  The additional 1 under the square root flattens the potential at small $r$, providing a better model for the gravitational force on an extended gas blob (of size comparable to $a$) when it is close to the center of the galaxy.  This modification also makes the integrator behave better near the center of the galaxy.  The virial radius of M86 was set to 100 kpc and its concentration parameter to 8, roughly the values expected for a massive galaxy.  The normalizing factor, $\\psi$, was expressed as $\\psi = 9 \\sigma^2$, so that $\\sigma$ is a rough measure of the line-of-sight velocity dispersion for M86.  The value used below is $\\sigma = 185\\ \\kms$.  We note that for the model what matters is the ratio of the drag force on the blob to the binding force of the M86 potential, allowing for a trade off between the choice of potential and the density ratio and size of the blob. As a result, the choice of potential is not very critical.  The two remaining parameters are the axial ratio, $a/b$, and the orientation of M86.  Consistent with its S0/E3 classification, optical images show that M86 is highly flattened on scales comparable to the projected distance of the plume (e.g. $a/b \\simeq 1.56$ for the isophote with $a\\simeq 4.5'$ in the deep image of Nulsen \\& Carter 1987).  Based on the discussion above, the potential of M86 would need to have a long axis pointing roughly northward from our line-of-sight. Therefore, we assume it is oblate, using $a/b = 2$.  The orientation of M86 is then determined by the direction of its minor axis.  From the optical images, this has a position angle of $\\simeq 35^\\circ$ on the sky.  However, its tilt with respect to the plane of the sky is unknown.  We have set the minor axis to point towards us in the north, at an angle of $45^\\circ$ from the plane of the sky.  This is roughly the orientation that maximizes the transverse displacement of the blob.  Figure~\\ref{fig:blob} shows the path of the blob on the sky for this set of model parameters for the orbit with $\\ro$ = 9.1~Mpc and $\\ri$ = 314~kpc.  In this model, the blob is currently slower than M86 by 355~\\kms\\ along our line of sight and trails it by 38~kpc.  Around the edges of the gas halo in M86, where the shear in the external flow is strong (Figure~\\ref{fig:inclined}), if the effective viscosity is high, viscous stresses pull the interstellar gas out of the galaxy. Alternatively, if the viscosity is low, shear instabilities mix the interstellar gas with the ICM, also stripping it from the galaxy (Nulsen 1982).  This stripping is aided by the low pressure due to the Bernoulli effect around the edges of the galaxy, which tends to pull gas into the path of the flow.  This process works around all edges of the inclined galaxy, as seen from the direction of the flow.  However, it is expected to be greatest at the leading and trailing edges.  The shear in the external flow is expected to be greatest near the leading edge, favoring stripping there.  The large ram pressure at the leading edge displaces the edge of the interstellar gas deeper into the potential of the galaxy, pushing it towards the trailing edge of the galaxy, where the pressure is lower. Thus, gas at the trailing edge sits higher in the gravitational potential of the galaxy, favoring its removal from the galaxy.  This may account for the apparent double streams of gas seen trailing M86 and several other galaxies in the composite image.  ``Viscous'' stripping from the main body of M86 as well as the plume can explain the broad features of the gas trail, though modeling the finer features, such as the smaller blobs of gas lying to the east of the galaxy, requires a more detailed treatment of the gas dynamics (e.g., numerical simulations) and is beyond the scope of this paper.  The proton mean free path due to Coulomb collisions is given approximately by $\\lambda \\simeq 150 (kT)^2 n_{-3}^{-1}$ pc, where $kT$ is the gas temperature in keV and the electron density is $10^{-3} n_{-3}\\rm\\ cm^{-3}$.  In terms of this, the Reynolds number is $\\reynolds \\simeq v L / (s \\lambda)$, where $v$ is the flow speed, $s$ is the sound speed, and $L$ is a relevant length scale.  Taking $v = 1550\\ \\kms$ and $L = 10$ kpc gives $\\reynolds \\simeq 20$ in the 2.4 keV ICM, but $\\reynolds \\simeq 3000$ for the gas in the 0.77~keV plume (and interstellar gas) if it is exposed directly to the external flow. These values suggest that the external flow can be largely laminar, while the flow in the cool interstellar gas is relatively turbulent.  Further complicating matters, the effective mean free path may be significantly smaller than the Coulomb mean free path (e.g., Schekochinin \\etal\\ 2007).  The Reynolds number for this flow is therefore not well defined, due to the range in gas temperature and uncertainty in the mean free path.      We have argued that the plume and long tail of M86 formed due to ram pressure stripping forces generated as M86 falls into the Virgo cluster. Several studies have found a similar interpretation for the formation of this feature (e.g., Forman et al.\\ 1979; Fabian et al.\\ 1980; Rangarajan et al. 1995, however see Bregman \\& Roberts 1990; Finoguenov et al.\\ 2004). We concentrate on these main results:  \\begin{itemize}  \\item the plume and long tail observed in the diffuse emission are created by ram pressure stripping as M86 falls into the Virgo cluster.  The tail is 150~kpc in projection (a simple estimate, which assumes free-fall velocity for M86 and an NFW potential for M87, gives a lower-limt on the true length of the tail of 380~kpc), making this the longest ram pressure stripped tail presently known.  \\item based on the X-ray spectra, we detect three distinct components associated with the M86/Virgo cluster system: the Virgo ICM, with $kT \\sim 2.4$~keV; the extended halo of M86, with $kT \\sim 1.2$~keV; and the cooler central and stripped gas of M86, with $kT \\sim 0.8$~keV.  The temperature structure of the tail is consistent with ram pressure stripping, where the higher entropy M86 halo gas is stripped first and deposited in the tip of the tail, and the lower entropy M86 ISM is stripped more recently, constituting the base of the tail and the plume.  \\item the large line-of-sight velocity of M86, and its position relative to the Virgo cluster, tightly constrain its orbit, especially if it is assumed that the gas tail traces the orbit.  In particular, the observations are inconsistent with a radial orbit.  We show that M86 is at best only marginally bound to the Virgo cluster, with an inner turning radius on the order of 300~kpc as expected from its recent in-fall.  Our best-fitting orbital model requires that M86 be close to M87, less than 167~kpc closer than or 263~kpc farther than M87 along our line of sight, which is consistent with the most recent distance estimates based on surface brightness fluctuations (Mei et al.\\ 2007) which give a line-of-sight separation of $0.4\\pm0.8$~Mpc.   \\item the prominent plume of gas lying 3\\arcmin -- 4\\arcmin north of M86 appears to have been rapidly driven from M86 by ram pressure stripping.  The projected position of the plume, which does not lie directly on our best-fit model orbit for M86, can be understood if M86 has an aspherical potential (as suggested by optical isophotes).  If M86 moves through the Virgo ICM at an inclination angle relative to the local flow, the gas at the trailing edge is more easily stripped, thereby displacing the gas from the nominal orbit of M86 itself.  This model may also explain the apparent double streams of gas seen trailing M86, as well as those in other Virgo galaxies.  \\item the apparent brightness edge to the southeast seen in {\\it ROSAT} observations is also seen in the {\\it Chandra} images.  The edge is well fit with a two power law gas density model, with an abrupt jump in density by a factor of $ 1.3^{+0.3}_{-0.4}$ at the edge (consistent with no jump).  Assuming that this brightness edge is the shock generated as M86 supersonically falls into the Virgo cluster, the lack of a well-defined density jump is consistent with what is expected from our orbital model, which suggests that the orientation of the Mach cone would make it difficult to detect the shock edge.   \\end{itemize}"
    },
    "0806/0806.0327_arXiv.txt": {
        "abstract": "{ Close-in giant extrasolar planets (``Hot Jupiters'') are believed to be strong emitters in the decametric radio range. } {We present the expected characteristics of the low-frequency magnetospheric radio emission of all currently known extrasolar planets, including the  maximum emission frequency and the expected radio flux. We also discuss the escape of exoplanetary radio emission from the vicinity of its source, which imposes additional constraints on detectability. } {We compare the different predictions obtained with all four existing analytical models for all currently known exoplanets. We also take care to use realistic values for all input parameters.} { The four different models for planetary radio emission lead to very different results. The largest fluxes are found for the \\textit{magnetic} energy model, followed by the \\textit{CME} model and the \\textit{kinetic} energy model (for which our results are found to be much less optimistic than those of previous studies). The \\textit{unipolar interaction} model does not predict any observable emission for the present exoplanet census. We also give estimates for the planetary magnetic dipole moment of all currently known extrasolar planets, which will be useful for other studies. } {Our results show that observations of exoplanetary radio emission are feasible, but that the number of promising targets is not very high. The catalog of targets will be particularly useful for current and future radio observation campaigns (e.g.~with the VLA, GMRT, UTR-2 and with LOFAR).} ",
        "introduction": "In the solar system, all strongly magnetised planets are known to be intense nonthermal radio emitters. For a certain class of extrasolar planets (the so-called Hot Jupiters), an analogous, but much more intense radio emission is expected. In the recent past, such exoplanetary radio emission has become an active field of research, with both theoretical studies and ongoing observation campaigns.  Recent theoretical studies have shown that a large variety of effects have to be considered, e.g.~kinetic, magnetic and unipolar interaction between the star (or the stellar wind) and the planet, the influence of the stellar age, the potential role of stellar CMEs, and the influence of different stellar wind models. So far, there is no single publication in which all of these aspects are put together and where the different interaction models are compared extensively. We also discuss the escape of exoplanetary radio emission from its planetary system, which depends on the local stellar wind parameters. As will be shown, this is an additional constraint for detectability, making the emission from several planets impossible to observe.  The first observation attempts go back at least to \\citet{Yantis77}. At the beginning, such observations were necessarily unguided ones, as exoplanets had not yet been discovered. Later observation campaigns concentrated on known exoplanetary systems. So far, no detection has been achieved. A list and a comparison of past observation attempts can be found elsewhere \\citep[][]{Griessmeier51PEG05}. Concerning ongoing and future observations, studies are performed or planned at the VLA \\citep{Lazio04}, GMRT \\citep{Majid05,Winterhalter06}, UTR2 \\citep{Ryabov04}, and at LOFAR \\citep{Farrell04}. To support these observations and increase their efficiency, it is important to identify the most promising targets.  The target selection for radio observations is based on theoretical estimates which aim at the prediction of the main characteristics of the exoplanetary radio emission. The two most important characteristics are the maximum frequency of the emission and the expected radio flux. The first predictive studies \\citep[e.g.][]{Zarka97,Farrell99} concentrated on only a few exoplanets. A first catalog containing estimations for radio emission of a large number of exoplanets was presented by \\citet{Lazio04}. This catalog included 118 planets (i.e.~those known as of 2003, July 1) and considered radio emission energised by the kinetic energy of the stellar wind (i.e.~the \\textit{kinetic model}, see below). Here, we present a much larger list of targets (i.e.~197 exoplanets found by radial velocity and/or transit searches as of 2007, January 13, taken from http://exoplanet.eu/), and compare the results obtained by all four currently existing interaction models, not all of which were known at the time of the previous overview study. As a byproduct of the radio flux calculation, we obtain estimates for the planetary magnetic dipole moment of all currently known extrasolar planets. These values will be useful for other studies as, e.g., star-planet interaction or atmospheric shielding.  To demonstrate which stellar and planetary parameters are required for the estimation of exoplanetary radio emission, some theoretical results are briefly reviewed (section \\ref{sec:theory}). Then, the sources for the different parameters (and their default values for the case where no measurements are available) are presented (section \\ref{sec:modelling}). In section \\ref{sec:results}, we present our estimations for exoplanetary radio emission. This section also includes estimates for planetary magnetic dipole moments. Section \\ref{sec:conclusion} closes with a few concluding remarks. ",
        "conclusions": "\\label{sec:conclusion}  Predictions concerning the radio emission from all presently known extrasolar planets were presented. The main parameters related to such an emission were analyzed, namely the planetary magnetic dipole moments, the maximum frequency of the radio emission, the radio flux densities, and the possible escape of the radiation towards a remote observer.  We compared the results obtained with various theoretical models. Our results confirm that the four different models for planetary radio emission lead to very different results. As expected, the largest fluxes are found for the \\textit{magnetic} energy model, followed by the \\textit{CME} model and the \\textit{kinetic} energy model. The results obtained by the latter model are found to be less optimistic than by previous studies. The \\textit{unipolar interaction} model does not lead to observable emission for any of the currently known planets. As it is currently not clear which of these models best describes the auroral radio emission, it is not sufficient to restrict oneself to one scaling law (e.g.~the one yielding the largest radio flux). Once exoplanetary radio emission is detected, observations will be used to constrain and improve the model.  These results will be particularly useful for the target selection of current and future radio observation campaigns (e.g.~with the VLA, GMRT, UTR-2 and with LOFAR).  We have shown that observation seem feasible, but that the number of suitable candidates is relatively low. The best candidates appear to be HD 41004 B b, Epsilon Eridani b, Tau Boo b, HD 189733 b, Gliese 876 c, HD 73256 b, and GJ 3021 b. The observation of some of these candidates is in progress."
    },
    "0806/0806.4149.txt": {
        "abstract": "We have studied the rapid X-ray time variability in 99 pointed observations with the \\textit{Rossi X-ray Timing Explorer} (RXTE)'s Proportional Counter Array of the low-mass X-ray binary 1E~1724--3045 which includes, for the first time, observations of this source in its island and banana states, confirming the atoll nature of this source. We report the discovery of kilohertz quasi-periodic oscillations (kHz QPOs). Although we have 5 detections of the lower kHz QPO and one detection of the upper kHz QPO, in none of the observations we detect both QPOs simultaneously.  By comparing the dependence of the rms amplitude with energy of kHz QPOs in different atoll sources, we conclude that this information cannot be use to unambiguously identify the kilohertz QPOs as was previously thought. We find that Terzan~2 in its different states shows timing behavior similar to that seen in other neutron-star low mass X-ray binaries (LMXBs). We studied the flux transitions observed between February 2004 and October 2005 and conclude that they are due to changes in the accretion rate. ",
        "introduction": "\\label{sec:intro}  Low-mass X-ray binaries (LMXBs) can be divided into systems containing a black hole candidate (BHC) and those containing a neutron star (NS). The accretion process onto these compact objects can be studied through the timing properties of the associated X-ray emission \\citep[see, e.g., ][for a review]{Vanderklis06}.  \\citet{Hasinger89} classified the NS LMXBs based on the correlated variations of the X-ray spectral and rapid X-ray variability properties. They distinguished two sub-types of NS LMXBs, the Z sources and the atoll sources, whose names were inspired by the shapes of the tracks that these sources trace out in an X-ray color-color diagram (CD) on time scales of hours to days. The Z sources are the most luminous, but the atoll sources are more numerous and cover a much wider range in luminosities \\citep[e.g. ][and references within]{Ford00}. For each type of source, several spectral/timing states are identified which are thought to arise from qualitatively different inner flow configurations \\citep{Vanderklis06}.  In the case of atoll sources, the three main states are the extreme island state (EIS), the island state (IS) and the banana branch, the latter subdivided into lower-left banana (LLB), lower banana (LB) and upper banana (UB) states. The EIS and the IS occupy the spectrally harder parts of the color color diagram and correspond to lower levels of X-ray luminosity ($L_x$).  The associated patterns in the CD are traced out in hours to weeks.  The hardest and lowest $L_x$ state is the EIS, which shows strong \\citep[up to 50\\% rms amplitude, see][ and references within]{Linares07} low-frequency flat-topped noise also known as band-limited noise (BLN).  The IS is spectrally softer and has higher X-ray luminosity than the EIS.  Its power spectra are characterized by broad features and a dominant BLN component which becomes weaker and generally higher in characteristic frequency as the flux increases and the $>6$ keV spectrum gets softer.  In order of increasing $L_x$ we then encounter the LLB, where twin kHz QPOs are generally first observed, the LB, where 10-Hz BLN is still dominant and finally, the UB, where the $<1$~Hz (power law) very low frequency noise (VLFN) dominates.  In the banana states, some of the broad features observed in the EIS and the IS become narrower (peaked) and occur at higher frequency. % In particular, the twin kHz QPO can be found in the LLB at frequencies higher than 700~Hz \\citep[the lower of them with frequencies generally lower than a 1000~Hz, while the upper one up to 1258~Hz, see e.g.,][]{Belloni05a,Jonker07}, only one kHz QPO can be generally found in the LB, and neither of them is detected in the UB \\citep[see reviews by][]{Vanderklis00,Vanderklis04,Vanderklis06}. %       A small number of weak NS LMXBs % which do not get brighter than a few times $10^{36}$ ergs s$^{-1}$ \\citep[usually burst sources and often referred to as 'weak' or 'faint bursters', see e.g,][ and references within]{Muno03} % resemble atoll sources in the EIS, but in the absence of state transitions this identification has been tentative \\citep[see for example][]{Barret00a,Vanderklis06}.  An important clue is provided by the correlations between the component frequencies (and strengths - see e.g. \\citealt{Straaten02,Straaten03,Altamirano08}) which helps to identify components across sources. % For example, \\citet{Straaten02,Straaten03} compared the timing properties of the atoll sources 4U~0614+09, 4U~1608--52 and 4U~1728--34 % (see also \\citealt{Altamirano08}, for similar results when the atoll source 4U~1636--53 was included in the sample) % and conclude that the frequencies of the variability components in these sources follow the same pattern of correlations when plotted versus the frequency of the upper kHz QPO ($\\nu_u$). %and not versus $\\nu_{\\ell}$, as in the case of \\citet{Psaltis99}. \\citet{Straaten03} also showed that low luminosity systems extend the frequency correlations observed for the atoll sources. This last result gave further clues in the link between the atoll and the low luminosity sources. %   \\citet{Psaltis99} found an approximate frequency correlation involving a low-frequency QPO, %($L_h$ and $L_{LF}$) the lower kHz QPO frequency and two broad noise components interpreted as low-frequency versions of these features.  This correlation spans nearly three decades in frequency, where the Z and bright atoll sources populate the $>100$\\,Hz range and black holes and weak NS systems the $<10$\\,Hz range. As already noted by \\citet{Psaltis99}, because the correlation combines features from different sources which show either peaked  or broad components with relatively little overlap, the data are suggestive but not conclusive with respect to the existence of a single correlation covering this wide frequency range \\citep{Vanderklis06}. %   The low-luminosity neutron star systems can play a crucial role in clearing up this issue. Observations of different source states in such a system could connect the $<10$ and $>100$~Hz regions mentioned above by direct observation of a transition in a single source.  In the case of the pattern of correlations reported by \\citet{Straaten03}, low luminosity NS systems extend the frequency correlations observed for ordinary atoll sources down to $\\sim100$~Hz. Unfortunately, the low luminosity NS systems are usually observed in only one state (EIS), which makes it difficult to properly link these sources to the atoll sources. However, some of these objects show rare excursions to higher luminosity levels which might correspond to other states. The occurrence of these excursions are usually unpredictable. Therefore, in practice it was not possible until now to check on the frequency behavior of the different variability components as such a source enters higher luminosity states.  1E~1724--3045 is a classic low luminosity LMXB; a persistent Low-Mass X-ray binary located in the globular cluster Terzan~2 \\citep{Grindlay80} which is a metal-rich globular cluster of the galactic bulge.  Its distance is estimated to be between 5.2 to 7.7 kpc \\citep{Ortolani97}.  These values are consistent with that derived from a type I X-ray burst that showed photospheric expansion \\citep[see][ but also see \\citealt{Kuulkers03,Galloway06}]{Grindlay80}.  The type I X-ray bursts observed from this source also indicate that the compact object is a weakly magnetized neutron star \\citep{Swank77,Grindlay80}. \\citet{Emelyanov02} have shown, using $\\sim30$ years of data from several X-ray satellites, that the luminosity of Terzan~2 increased until reaching a peak in 1997, after which it started to decrease. They suggest that the evolution of the donor star or the influence of a third star could be the cause of this behavior. \\citet{Olive98} and \\citet{Barret00} have shown that during earlier observations of Terzan~2 its X-ray variability at frequencies $\\gtrsim0.1$~Hz resembled that of black hole candidates. This state was tentatively identified as the extreme island state for atoll sources. Until now, no kilohertz quasi-periodic oscillations have been reported for this source, which was attributed to the fact that the source was always observed in a single intensity state \\citep{Barret00a}.    Monitoring observations by the All Sky Monitor aboard the Rossi X-ray Timing Explorer showed that the source was weakly variable in X-rays (less than about a factor of 3 on a few day time scale for the first 8 years of the monitoring).  However, recently \\citet{Markwardt04} reported (using PCA monitoring observations of the galactic bulge - \\citealt{Swank01}) that during February 2004, 1E~1724--3045 flared up from its relatively steady $\\sim20$ mCrab to $\\sim66$~mCrab (2--10~keV). % In this paper we report a complete study of the timing variability of the source. % % For simplicity, and since only one bright X-ray source is detected in the globular cluster (See Section~\\ref{sec:samesource}), in the rest of this report we will refer to 1E~1724--3045 as Terzan~2. ",
        "conclusions": "\\label{sec:discussion}        \\subsection{Contamination by a second source in the same field of view?}\\label{sec:samesource}   As shown in Figure~\\ref{fig:lc}, the luminosity of the source slowly decreases with time during the quiet period 51214--52945 MJD. Although the rms amplitude changes up to 30\\% (see Figure~\\ref{fig:rms1000}), the X-ray timing characteristics are very similar (see Interval A to L in Figure~\\ref{fig:pds}). During the 53000--53700 MJD period, the source shows flares which show different X-ray timing characteristics consistent with the island and banana states observed in other atoll sources \\citep[see e.g.][]{Straaten03,Straaten05,Belloni02,Altamirano05,Altamirano08}. % A possible mechanism of the observed flux variations in Terzan~2 could be the emergence of a second X-ray source in this globular cluster unresolved by the $1^{\\circ}$ (FWHM) field of view of the PCA. If two sources are observed simultaneously with RXTE, then we would expect to see power spectra which are a combination of the intrinsic time variability of both sources. %  To further investigate this, we compared the absolute rms amplitude that we observe both in the quiet and flaring states. Observation 80105-10-01-00 is the last observation performed during the quiet period from which we measured an average source countrate of $\\sim200$ counts / second. The integrated power between $7.8 \\cdot 10^{-3}$ and $1$~Hz is $(3.4\\pm2) 10^{-2}$, which corresponds to a fractional rms amplitude of $\\sim18\\pm0.6$\\%, i.e. an absolute rms amplitude of $36\\pm1$ counts/second.   Observation 80138-06-01-00 is the first observation performed during the flaring state. Its average countrate is $\\sim740$ counts / second and the absolute rms amplitude in the $0.0078-1$~Hz range is $21.4\\pm4.4$ counts/second ($2.9\\pm0.6$\\% fractional rms amplitude). Clearly, the absolute rms amplitudes are different when comparing quiet and flaring periods.  If we repeat the analysis using the second RXTE observation during the flaring state (80138-06-02-00), the discrepancy is higher. This observation has an average source countrate of $\\sim405$ counts/second and the upper limit for the absolute rms amplitude in the $7.8 10^{-3}-1$~Hz frequency range is 5 counts/second.      Given the characteristics of the power spectra, flares cannot be explained by assuming that another source has emerged, unless Terzan~2 turned off at the same time that the other X-ray source turned on, which is unlikely. Therefore, we conclude that the flux transitions are intrinsic to the only low mass X-ray binary detected in the globular cluster Terzan~2: 1E~1724--30 \\citep{Revnivtsev02}.       \\subsection{The kilohertz QPOs, different states and their transitions}\\label{sec:kHz2}  The results presented in this paper show that the low luminosity source 1E~1724--3045 in the globular cluster Terzan~2 can be identified as an atoll source. This is the first time a source previously classified as weak burst source \\citep[see e.g,][ and references within]{Belloni02,Straaten03} showed other states than those of the extreme island state, confirming previous suggestions that these sources are atoll sources. We have identified the new states as the island and banana states based on comparisons between color color diagrams of different sources and the characteristics of the power spectra. We have detected at least one of the the kHz QPOs, and as explained in Sections~\\ref{sec:flaringsegment}~and~\\ref{sec:kHz}, in 5 cases we may be detecting the lower kHz QPO (intervals M and N -- see also Table~\\ref{table:kHz}) and in one case the upper one (interval Q). % No simultaneous twin kHz QPOs were detected within any of the 14 observations that sample the flares. Future observations of flares will allow us to confirm these identifications and might allow us to detect both kHz QPOs simultaneously.  We found that the frequencies of the various components in the power spectra of Terzan~2 followed previously reported relations (Figures~\\ref{fig:nuvsnu}~and~~\\ref{fig:pbk}).  Terzan~2 is a particularly important source in the context of these frequency correlations because it is one of the few neutron star sources that has been demonstrated to show power spectral features that reach frequencies as low as $\\simeq0.1$~Hz, which is uncommon for neutron star low mass X-ray binaries, but not for black holes. %(see Figure~\\ref{fig:pbk} %- \\citealt{Psaltis99}). % Our results demonstrate that in each of the flares, Terzan~2 undergoes flux transitions that, if directly observed, would probably allow us to resolve current ambiguities in the identification of components, such as the case of $L_{\\ell ow}$ component in atoll sources. This component is interpreted by some authors as a broad lower kHz QPO at very low frequencies \\citep[see][]{Psaltis99,Nowak00,Belloni02} which becomes peaked at higher frequencies, while other authors interpret $L_{\\ell}$ and $L_{\\ell ow}$ as different components \\citep[see e.g. discussion in ][]{Straaten03}. Another example is the identification of the upper kHz QPO at low frequencies. \\citet{Straaten03} have suggested that the broad component observed at $\\simeq150$~Hz in the EIS of atoll sources becomes the peaked upper kHz QPO $L_u$. These authors based their interpretation on the frequency correlations shown in Figure~\\ref{fig:nuvsnu}. Nevertheless, as \\citet{Straaten03} argue, these identifications should be taken as tentative.  One way to confirm the link between them would be to observe the gradual transformation from one to another one.  During the time between flares, the source shows intensities similar to those measured before the quasi-periodic flares started. Unfortunately there are no observations during those intervals, but we expect that then Terzan~2 shows X-ray variability similar to that reported in intervals A--L.  If this is the case, the state transition between the extreme island state and the island state should be observable in observations at the beginning or at the end of each flare. Given the relatively gradual and predictable transitions, Terzan~2 becomes the best source known up to now to study these important transitions.       \\subsection{On the $\\sim90$ days flare recurrence }\\label{sec:lightcurve}   The quasi-periodic variations over days and months observed in some LMXBs X-ray light-curves are generally associated with the possible precession period of a tilted accretion disk or alternatively long term periodic variations in the accretion rate or periodic outbursts of X-ray transients. % Some examples are the $\\simeq35$ cycle in Her~X-1 which is thought to be caused by a varying obscuration of the neutron star by a tilted-twisted precessing accretion disc; the $\\simeq170$ days accretion cycle of the atoll source 4U~1820--30, \\citep{Priedhorsky84,Simon03}; the 122-125 day cycle in the outbursts of the recurrent transient Aql X-1 \\citep{Priedhorsky84b,Kitamoto93}. % Understanding the mechanisms that trigger the long-term variability associated with variations in the accretion rate of LMXBs can allow us to better predict, within each source, when the state transitions occur. This is useful because these transitions are usually fast and therefore difficult to observe.   The power spectra of our observations of Terzan~2 during the flaring confirm that the source undergoes EIS-IS-LLB-LB-UB state transitions, as observed in other neutron star atoll systems \\citep[and not as seen for Z-sources, see reviews by][ and references within]{Vanderklis04,Vanderklis06}. As the source increased in X-ray luminosity, we found that the components in the power spectra increased in frequency which is consistent with the interpretation that the accretion disk is moving inwards toward the compact object. Therefore, the flaring with average 90~days period is most probably an accretion cycle. % We note that the modulation of the light-curve could be related to the orbital period of the system or set by the precession of a tilted disk.  However, the mechanisms involved in those interpretations are very unlikely to affect the frequency of the kHz QPOs.   If the flares are explained as an accretion cycle, then it is puzzling why the source underwent a smooth decrease of $L_x$ for $\\simeq8$~years before it started to show the flares. Terzan~2 may not be the only source that shows this kind of behavior. For example, KS~1731--260 is a low-luminosity burster that has shown a high $L_x$ phase, during which \\citet{Revnivtsev03} reported a possible $\\simeq38$ days period, and a low $L_x$ phase, during which much stronger variability was observed (which was described as red noise). After its low $L_x$ phase, KS~1731--260 has turned into quiescence \\citep{Wijnands02a,Wijnands02b}. In Figure~\\ref{fig:ks1731} we show the bulge scan light curve of the source during the low $L_x$ phase. At MJD~$\\sim51550$ the source reached very low intensities, then flared up again for $\\gtrsim 250$~days to finally turn into quiescence. The low luminosities are confirmed by the ASM light curve (not plotted). % Recently, \\citet{Shih05} reported that the persistent atoll source 4U~1636--53 has also shown a period of high $L_x$ followed by a period of low $L_x$. During high $L_x$, no long-term periodicity was found, but a highly significant $\\simeq46$~days period was observed after its $L_x$ decline.  These  similar patterns of behavior might point towards a common mechanism, which then must be unaffected by the intrinsic differences between these sources. %  For example, while Terzan~2 remained with approximately constant luminosity in its extreme island state for $\\simeq8$ years before showing long term periodicities, 4U~1636--53 and KS~1731--260 were observed with variable luminosity and in different states, including the banana state in which the kHz QPOs were found (see e.g. \\citealt{Wijnands97b} and \\citealt{Shih05} for 4U~1636--53 and \\citealt{Wijnands97a} and \\citealt{Revnivtsev03} for KS~1731--260). While Terzan~2 reached a maximum luminosity of $L_x/L_{Edd}\\simeq0.02$ during one of the flares, % KS~1731--260 has %turned into quiescence \\citep{Wijnands02a,Wijnands02b} and 4U~1636--53 shows similar luminosities only at its lowest $L_x$ levels \\citep[while it has reached $L_x/L_{Edd} \\gtrsim0.15$ -- see][]{Altamirano08}.  Further differences may be related to whether these systems are normal or ultra-compact binaries. While 4U~1636--53 is not ultra-compact (see below), \\cite{Intzand07} has recently proposed that Terzan~2 may be classified as ultra-compact based on measurements of its persistent flux, long burst recurrence times and the hard X-ray spectra. % If the luminosity behavior of these sources is related, the differences outlined above suggest that the mechanism that triggers the modulation of the light curve at low $L_x$ may not depend on the accretion history, the luminosity of the source or even whether the system is ultra-compact or not.  The modulation period may depend on these factors. %    \\begin{figure}[!hbtp] \\center \\resizebox{1\\columnwidth}{!}{\\rotatebox{0}{\\includegraphics{./f13.eps}}} \\caption{The PCA monitoring observation lightcurve of the atoll source KS~1731--260 during part of its low $L_x$ period. Unfortunately, there are no PCA monitoring observations of the source before to MJD~51200. Clearly, the source flares up similarly to Terzan~2 before it turns into quiescence. Interestingly, the data at MJD $\\sim51550$ shows that the source had a period of very low intensity, followed by a flaring up that lasted for $\\sim250$~days before the source finally turned into quiescence.  } \\label{fig:ks1731} \\end{figure}    Unfortunately, we cannot compare the orbital periods and the companions of the three systems, as these are only known for 4U~1636--53 \\citep[$\\sim3.8$~hours and $\\simeq0.4$~M$_{\\odot}$][]{Casares06}. Nevertheless, with the present data it is already possible to exclude some mechanisms. % For example, mass transfer feedback induced by X-ray irradiation \\citep{Osaki85} is unlikely.  In this model, X-ray radiation from the compact object heats the companion star surface, causing enhancement of the mass accretion rate in a runaway instability. However, in \\Citeauthor{Osaki85}'s scenario, it is not clear how the system could remember the phase of the cycle if one of the flares is missed or if the size of the flares differs much.  Flares F4 and F5 in Terzan~2, independently of the other two sources, may already raise an objection to this model.  Although we miss part of F4 due to a gap in the data, Figure~\\ref{fig:lc} shows that F4 was quite short (less than 9 days), while F5 was the longest ($\\lesssim36$~days) and strongest flare.   \\citet{Shih05} have suggested that the atoll source 4U~1636--53 may turn into quiescence after its low $L_x$ period, as was observed for KS~1731--260. Such an observation for 4U~1636--53 as well as for Terzan~2 would give credibility to the link between these sources.  To our knowledge, there is no model which predicts such behavior.  \\subsection{Energy dependence as a tool for kHz QPO identification}  \\citet{Homan00} discovered a single 695~Hz QPO in the low mass X-ray binary EXO~0748--676 and identified this QPO as the lower kHz QPO. These authors based their identification on the fact that at that time: (i) from the 11 kilohertz QPO pairs found in atoll sources, eight had ranges of lower peak frequencies that include 695 Hz, which was the case for only three of the upper peaks and (ii) the upper peaks in atoll sources generally had Q lower than $\\sim14$, while their QPO had Q$\\gtrsim38$, value more common for lower peaks. % While from Figure~\\ref{fig:nuvsnu} it can be seen that (i) is not valid anymore, since the upper kHz QPOs have been detected down to 300--400~Hz (and possibly down to $\\lesssim100$~Hz -- see Section~\\ref{pdsqp}), % at these low frequencies $L_u$ is usually much broader than they observed, which confirms their identification (see Section~\\ref{sec:tfp} and \\citealt{Barret05a,Barret05b,Barret05c}).  \\citet{Homan00} also based their identification on the comparison of the energy dependence of the QPO with that of the two kilohertz peaks in 4U~1608--52, which have rather different energy dependencies \\citep[the power-law rms-energy relation for $L_{\\ell}$ is steeper than that for $L_u$, see ][]{Berger96,Mendez98a,Mendez01}. Similarly, \\citet{Mendez01} use the same method to strengthen the identification of the single kHz QPO observed in the atoll source Aql~X-1. To further investigate if this method could be used to identify the sharp kHz QPOs we report in Section~\\ref{sec:kHz}, in Figure~\\ref{fig:rmsvse} we compare the energy dependence of the kHz QPOs in 4U~1608--52 \\citep{Berger96,Mendez98a,Mendez01} and 4U~0614+09 \\citep{Mendez97} with that of Terzan~2.  The data for Terzan~2 seem to fall in between those for $L_{\\ell}$ and $L_u$ of 4U~1608--52 but shows a completely different behavior than the data of 4U~0614+09. % The fact that the rms amplitude of the upper kHz QPO in 4U~0614+09 and 4U~1608--52 are significantly different (by up to a factor of 3) and that the data for Terzan~2 fall in between those of $L_{\\ell}$ and $L_u$ in 4U~1608--52 show that the method does not lead to unambiguous results. Mean source luminosity, instantaneous luminosity and instantaneous QPO frequency may all affect QPO energy dependence in addition to QPO type. % %"
    },
    "0806/0806.4674_arXiv.txt": {
        "abstract": "{} {In the interstellar clouds, molecular hydrogens are formed from atomic hydrogen on grain surfaces. An atomic hydrogen hops around till it finds another one with which it combines. This necessarily implies that the average recombination time, or equivalently, the effective grain surface area depends on the  relative numbers of atomic hydrogen influx rate and the number of sites on the grain. Our aim is to discover this dependency.} {We perform a numerical simulation to study the recombination of hydrogen on grain surfaces in a variety of cloud conditions. We use a square lattice (with a periodic boundary condition) of various sizes on two types of grains, namely, amorphous carbon and olivine.} {We find that the steady state results of our simulation match very well with those obtained from a simpler analytical consideration provided the `effective' grain surface area is written as $\\sim S^{\\alpha}$, where, $S$ is the actual physical grain area and $\\alpha$ is a function of the flux of atomic hydrogen which is determined from our simulation. We carry out the simulation for various astrophysically relevant accretion rates. For high accretion rates, small grains tend to become partly saturated with $H$ and $H_2$ and the subsequent accretion will be partly inhibited. For very low accretion rates, the number of sites to be swept before a molecular hydrogen can form is too large compared to the actual number of sites on the grain, implying that $\\alpha$ is greater than unity. }{} ",
        "introduction": "It is long known that the formation of $H_2$ in the interstellar cloud takes place on the grain surfaces and are then released into the gas phase. Considerable studies were made since then to understand the real physical processes which are taking place both theoretically (e.g., Hollenbach, Werner \\& Salpeter, 1971; Takahashi, Matsuda \\& Nagaoka, 1999; Biham et al. 2001) and experimentally (e.g., Pirronello et al. 1999). A quantitative estimation based on physisorption of the formation rate of $H_2$ was carried out by Biham et al. (2001) and Green et al. (2001). It was found that a significant production is possible in cooler ($\\sim 10-25$K) clouds (Rae et al. 2003). Cazaux \\& Tielens (2002, 2004) use both physisorption and chemisorption, to demonstrate that $H_2$ production is possible at high temperatures ($\\sim 200-400$K) also. Recently, Acharyya, Chakrabarti \\& Chakrabarti (2005) and Acharyya \\& Chakrabarti (2005) computed the formation of $H_2$ from $H$ assuming various cloud conditions and using rate and master equations presented by Biham et al. (2001). If one considers only $H$ and $H_2$, then, {\\it in a steady state}, the number of hydrogen atoms ($n_H$) and molecules ($n_{H2}$) on a grain surface are obtained by equating the effective influx (or, accretion) rate $\\phi_H$ of $H$ with the rate at which they are used up to form $H_2$ and/or get desorbed from the grain surface into the gas phase, i.e., $$ \\phi_H= W_H n_H + 2 a_H n_H^2 , \\eqno{(1a)} $$ and $$ W_{H2}n_{H2}= \\mu a_H n_H^2 . \\eqno{(1b)} $$ Here, $\\phi_H$ is the effective flux (in units of number/s) of $H$ on the grain surface, $W_H$ and $W_{H2}$ (both in units of s$^{-1}$) are the desorption co-efficients of  $H$ and $H_2$ respectively, $a_H$ (in units of s$^{-1}$) is the effective rate of recombination of two $H$ atoms to form one $H_2$ molecule and $1-\\mu$ is the fraction of $H_2$ that are formed on the grain which are spontaneously released to the gas phase. Solving these simple equations, we find in the steady state, number of $H$ and $H_2$ on a grain surface to be, $$ n_H=\\frac{\\sqrt{W_H^2 + 8 a_H\\phi_H}-W_H}{4a_H}, \\eqno{(2a)} $$ $$ n_{H2}= \\frac{\\mu a_H n_H^2}{W_H} \\eqno{(2b)} $$ We ignore $H_2$ formation by tunneling since the experiments indicate hopping to be the major process (see, Pironello et al. 1999; Katz et al. 1999). In the literature, it is customary to define the effective rate of recombination to be, $$ a_H = A_H/S, \\eqno{(3a)} $$ (See, Acharyya, Chakrabarti, \\& Chakrabarti, 2005  and references therein), where $A_H$ is the hopping rate (inverse of the diffusion time $T_d =1/A_H$) on a grain surface and $S$ is the number of sites on the grain surface (Biham et al. 2001). The argument for reducing the diffusion rate by a factor of $S$ is this: on an average, there are $n=S^{1/2}$ number of sites in each direction of the grain. Since the hopping is random, it would take a square of this, namely, $n^2=S$ number of hopping to reach a distance located at $n=S^{1/2}$ sites away, where, on an average, another $H$ is available. Thus, the effective recombination rate was chosen to be $A_H/S$. However, when the flux is very high in a small grain, many sites would be full and an $H$ need not hop a distance of $S^{1/2}$ to meet another one. Similarly, when the accretion rate is too low, one $H$ may have to sweep the whole grain several times before it can meet another $H$ to recombine. Thus, it is expected that the effective rate of recombination would be a function of the number of sites as well as the flux of $H$ on the grain. If we assume that the effective number of sites could be written as $S^{\\alpha_0}$, where $\\alpha_0$ is a constant for a given flux and grain, then, in principle, we could define the effective recombination rate to be, $$ a_H=A_H/S^{\\alpha_0}, \\eqno{(3b)} $$ and we expect that higher the accretion rate, lesser would be the value of $\\alpha_0$. In the opposite limit, when the accretion rate is low, one would get $\\alpha_0>1$ since the effective site number is higher than $S$. In order to see this effect we clearly need to actually compute time dependences of $H$ and $H_2$ by numerical simulation and derive $\\alpha_0$ when a 'steady state' is reached. This we do in this paper.  In the next Section, we present the logic based on which the simulation is carried out. In \\S 3, we present our results under various grain conditions and compare our results with the expectation from the steady state (Eqs. 2ab). We show that these two results match only if $\\alpha_0$ is used in the steady state equations. Finally, in Section \\S 4 we draw our conclusions. ",
        "conclusions": ""
    },
    "0806/0806.3748_arXiv.txt": {
        "abstract": "We report on our analysis of the 1 Ms Chandra observation of the supernova remnant Cas A in order to localize,  characterize and quantify its non-thermal X-ray emission. More specifically, we investigated whether the X-ray synchrotron emission from the inside of the remnant is from the outward shock, but projected toward the inner ring, or from the inner shell. We tackle this problem by employing a Lucy-Richardson deconvolution technique and measuring spectral indices in the 4.2-6 keV band.  We show that most of the continuum emission is coming from an inner ring that coincides with the location of the reverse shock. This inner ring includes filaments, whose X-ray emission has been found to be dominated by X-ray synchrotron emission. The X-ray emission from these filaments, both at the forward shock and from the inner ring, have relatively hard spectra with spectral index $> -3.1$. The regions emitting hard X-ray continuum contribute about 54\\% of the total X-ray emission in the 4.2-6 keV. This is lower than suggested by extrapolating the hard X-ray spectrum  as measured by BeppoSAX-PDS and INTEGRAL. This can be reconciled by assuming a gradual steepening of the spectrum toward higher energies. We argue that the X-ray synchrotron emission is mainly coming from the Western part of the reverse shock. The reverse shock in the West is almost at rest in our observation frame, corresponding to a relatively high reverse shock velocity of $\\sim 6000$~\\kms~in the frame of the freely expanding ejecta. ",
        "introduction": "Supernova remnants (SNRs) are the main candidates for producing Galactic cosmic rays, with energies at least up to the so-called knee of the cosmic ray spectrum at $\\sim 3 \\times 10^{15}$~eV. The first direct evidence for this is the detection of X-ray synchrotron emission caused by $\\sim 10^{14}$ eV electrons \\citep[first established for SN1006,][]{koyama}. Since the energy of electrons suffers from radiation losses, this might indicate even higher energies for ions. Moreover, hard X-ray tails up to 80 keV have been discovered for several Galactic SNRs \\citep{Allen1999}. This has been contributed to either non-thermal bremsstrahlung \\citep{Laming2001a} or to synchrotron radiation \\citep{Allen1997}. In recent years, additional direct evidence for efficient cosmic ray acceleration has come from detection of TeV $\\gamma$-rays for several SNRs by the High Energy Gamma-Ray Astronomy (HEGRA) experiment, the High Energy Spectroscopic System \\citep[H.E.S.S, e.g.][]{Aharonian2004} and MAGIC \\citep{Albert}. The $\\gamma$-ray emission is either caused by inverse Compton scattering by the same electrons that cause X-ray synchrotron emission, or by pion production caused by collisions of accelerated ions with the background plasma.  Cassiopeia A (Cas A) is one of the supernova remnants with a hard X-ray tail \\citep{The} and has recently also been detected in $\\gamma$-rays \\citep{Aharonian2001, Albert}.  This remnant was until recently\\footnote{Recent expansion measurements of G1.9+0.3 show an age for this remnant of around 100 years \\citep{Green, Reynolds2008}.}  the youngest known supernova remnant in the Galaxy; its supernova was probably around 1671 \\citep{Thorstensen}. In 2001 Chandra detected thin, X-ray synchrotron emitting, filaments at the forward shock of the remnant \\citep[Fig. \\ref{Continuum}, see also][] {Gotthelf}. This implies the presence of electrons with energies of $\\sim 10^{13}$~eV for the magnetic fields in Cas A, estimated to be 0.1 mG to 0.6 mG \\citep{Vink2003, Berezhko}. These synchrotron rims can be understood in the context of diffusive shock acceleration and synchrotron cooling downstream of the shock.  In addition, the Chandra image shows thin filaments at the inside of the remnant. Some of these inner filaments show a featureless spectrum (Fig. \\ref{Example}). \\cite{Hughes} identified one of the inner filaments at the West side of the remnant (`region D') as being the projected forward shock, based on its featureless spectrum. \\cite{DeLaney} found that the kinematics of the inner filaments, which they interpreted as projected forward shock filaments, are different from the forward shock; they have a lower velocity.  \\begin{figure}[t] \\centering \\plotone{f1_color.eps} \\caption{Million second 4-6 keV continuum image of Cas A, obtained with Chandra. The rectangle indicates region `D' in \\cite{Hughes}. See the electronic edition of the Journal for a color version of this figure.} \\label{Continuum}% \\end{figure}  \\notetoeditor{We would like to have Fig 1. in grayscale in the printed version and in color in the online version }  \\begin{figure}[b] \\centering \\epsscale{.9} \\plotone{f2.eps} \\caption{The spectrum of Cas A as observed by Chandra. Below is the spectrum of the featureless filament (`D') described by \\cite{Hughes} extracted from the megasecond observation, above is a spectrum of the whole remnant of one single observation (ObsID 4638), multiplied by 0.1.} \\label{Example}% \\end{figure}  Diffusive shock acceleration is a process which accelerates cosmic rays at a shock front \\citep[for a review, see][]{Malkov}. This mechanism accelerates charged particles of sufficient energy, which scatter on turbulent magnetic fields/plasma waves on both sides of the shock front. Each time the shock front is crossed, the particle gains energy, due to the difference in plasma velocity between both sides of the shock front. The higher the difference in the velocities is, the more energy is gained in one iteration and the higher the magnetic field and magnetic field turbulence, the more often particles cross the shock front. Since at the location where efficient particle acceleration takes place recently accelerated electrons are present, these locations show X-ray synchrotron radiation. However, further downstream from the shock front, synchrotron losses result in lower maximum energies of the synchrotron radiation.  The synchrotron spectrum can be approximated over a large range in frequencies with a power-law in flux density: $F_\\nu \\propto \\nu^{-\\alpha}$  with an index ($\\alpha$) related to the power-law index of the electron distribution ($p$) as: $\\alpha = (p-1)/2$. In what follows, we use index $\\Gamma$, which refers to the number density index $\\Gamma = -(\\alpha + 1)$ and $n(E) \\propto E^\\Gamma$. Near the maximum electron energies, the electron spectrum has an exponential cut-off, but the resulting synchrotron spectrum cuts off less abruptly, roughly as $\\exp(-{\\sqrt{\\nu/\\nu_{max}}})$  \\citep{Zirakashvili}. In contrast, the other important continuum emission process, thermal bremsstrahlung, has an exponential cut-off ($\\propto \\exp({-h\\nu/kT})$). In Cas A, the plasma temperature ranges between 0.6 and 3.6~keV \\citep{Yang}. One should take into account that these temperatures originate from a thermal model. If a partly non-thermal spectrum is fitted with a thermal model, the fitted temperature tends to increase with respect to the real temperature of the plasma. For a thermal bremsstrahlung spectrum with a temperature of 3.5~keV, the power-law index between 4.2 and 6.0 keV is -2.8. For a synchrotron spectrum at the forward shock of Cas A, we typically measure a power-law index of -2.1. We therefore expect the bremssstrahlung continuum to be steeper than the synchrotron continuum in the 4 to 6 keV continuum band.  In this paper we investigate the shape of the continuum spectrum and its spatial distribution in order to address several questions pertaining to the shock acceleration in Cas A: What is the location of the X-ray synchrotron filaments? What fraction of the overall X-ray continuum is thermal and what fraction is non-thermal? And what are the implications for the hard X-ray emission, above 10~keV. We do this by analyzing the Chandra megasecond observation of Cas A. ",
        "conclusions": "  \\begin{itemize} \\item[$\\bullet$] The power-law index of the spectrum between 4.2-6.0 keV is an indicator for X-ray synchrotron emission: there is a correlation between filaments, dominated by continuum emission and hard spectra, \\item[$\\bullet$] hard X-ray spectra are not exclusively associated with filaments, dominated by continuum emission, suggesting that non-thermal emission comes also from other regions, \\item[$\\bullet$] the non-thermal X-ray emission is likely to be synchrotron radiation, \\item[$\\bullet$] the non-thermal accounts for about 54\\% of the overall continuum emission in the 4-6 keV band, \\item[$\\bullet$] in the Western part of Cas A, most X-ray synchrotron comes from the reverse shock, \\item[$\\bullet$] the dominance of X-ray synchrotron emission from the West is probably the result of a locally higher reverse shock velocity of $v_s \\sim 6000$~\\kms (corresponding to a lower proper motion) than in the Eastern region ($v_s \\sim1900~$\\kms). \\end{itemize}"
    },
    "0806/0806.3781_arXiv.txt": {
        "abstract": "We  present evidence  that  the optically  unidentified radio  source, FIRST J121839.7+295325,  may be strongly lensing  a background galaxy. We estimate the redshift  of the assumed gravitational arc, discovered in parallel imaging with HST, from MMT-Blue Channel spectroscopy to be $z_{\\rm  arc}\\!=\\!\\zarc$.  We  present  lens models  with an  Einstein radius    of   $R_E\\!=\\!1\\farcs3$   which    contains   a    mass   of $M_{dyn}\\!=\\!10^{12\\pm0.5}$~M$_{\\odot}$,    where    the   uncertainty reflects the range  of possible lens redshifts.  The  putative lens is not    detected   to    $J_{\\rm   lim}\\!=\\!\\jlim$~mag    and   $H_{\\rm lim}\\!=\\!\\hlim$~mag in  our MMT-SWIRC imaging.  Using  the flux limits from  WFPC2 and SWIRC,  we estimate  that the  dynamical mass-to-light ratio               of               J121839.7+295325               is $M_{dyn}/L_B\\!\\gtrsim\\!10$~M$_{\\odot}$~L$_{\\odot}^{-1}$             for $A_V\\!=\\!1$~mag,   and  this  lower   limit  could   be  as   high  as $30$~M$_{\\odot}$~L$_{\\odot}^{-1}$  for   $A_V\\!=\\!0$~mag.   Since  the radio  source is  optically unidentified  ($V_{\\rm lim}\\!=\\!25.5$~mag) and has a radio flux of $S_{1.4~{\\rm GHz}}\\!=\\!33$~mJy, it is likely a massive   early-type  galaxy   which   hosts  a   radio-loud  AGN   at $0.8\\!\\lesssim\\!z_{\\rm  radio}\\!\\lesssim\\!1.5$.  However,  the present data cannot uniquely determine  the mass-to-light ratio of the lensing galaxy, and hence the possibility that this system may be a reasonably dark lens is not ruled out. ",
        "introduction": "\\label{intro}   In the $\\Lambda$-dominated, cold dark matter ($\\Lambda$CDM) cosmology, massive  galaxies form  hierarchically  inside of  dark matter  haloes \\citep[eg.][]{wf91,kauf93}.   The mass evolution  of these  CDM haloes can   place   stringent  constraints   on   the  cosmological   models \\citep[eg.][]{gre02}, and  the most reliable measurements  of CDM halo masses  are  generally derived  from  modeling gravitationally  lensed images. Furthermore, gravitational lensing  has become a powerful tool for  studying properties  of both  the  lensing objects  and the  more distant sources.   Of particular interest are constraints  on the mass profiles  \\citep[eg.][]{koop06,more}   and  the  mass-to-light  ratios \\citep[eg.][]{kkf98,rk05,tr06} of lens galaxies, and the host galaxies of active  galactic nuclei \\citep[eg.][]{kkm01,peng06}.   While all of these studies are advancing due  to increased numbers of known lensing systems   \\citep{rat99,leh00,bol06},  they   are   still  limited   by relatively small sample sizes.   In  this work,  we report  the discovery  of an  optical arc  which is $\\sim\\!4\\farcs0$ southwest of the optically unidentified radio source, FIRST \\name.  This radio source  was discovered in the Faint Images of the Radio Sky at Twenty-centimeters  (FIRST) with the Very Large Array \\citep[VLA;][]{beck,white} and later  detected at 610~MHz \\citep{reng} and  74~MHz  \\citep{cohen}.  Its  relatively  bright  flux of  $S_{\\rm 1.4~{\\rm   GHz}}\\!=\\!33$~mJy   suggests   that   it  is   at   $z_{\\rm radio}\\!\\lesssim\\!1.5$    \\citep[eg.][]{deb01,wad01}.    Additionally, \\name\\  was  \\textit{undetected}  in  the  pure-parallel  observations (PropID:~8090,  PI: S.~Casertano)  with the  Wide Field  and Planetary Camera~2   (WFPC2)   on   the   Hubble  Space   Telescope   (HST)   to F606W$\\gtrsim\\!25.5$~mag  \\citep[$V$-band  hereafter;][]{jqr07}, which indicates    it   is    likely   at    $z_{\\rm   radio}\\!\\gtrsim\\!0.8$ \\citep[eg.][]{win85,kron85,deb02,win03}.     These   two   independent constraints yield a coarse redshift estimate of $0.8\\!\\lesssim\\!z_{\\rm radio}\\!\\lesssim\\!1.5$.   By correlating  the WFPC2  archive  with the catalog of FIRST sources,  \\citet{jqr07} note an arc-like feature with $V_{\\rm   Vega}\\!=\\!24.0\\pm0.1$~mag.    Based   on   the   astrometric uncertainties for the HST-WFPC2  and FIRST imaging, they conclude that the  likelihood that  this arc  is the  optical identification  of the radio source is ${\\cal  L}\\!\\lesssim\\!10^{-10}$, following the work of \\citet{der}.  Since  \\citet{jqr07} rule  out the possibility  that the arc is  the optical identification  of \\name, we will  investigate the hypothesis that the optical arc  is a gravitationally lensed imaged by the optically unidentified radio source.   This  work is  organized as  follows: \\S~\\ref{mmt}  describes  our MMT observations,   \\S~\\ref{model}    outlines   the   lensing   analysis, \\S~\\ref{discuss}  discusses   the  lensing  interpretation   of  these observations,  and  \\S~\\ref{con}   gives  some  closing  remarks  with thoughts  toward  possible  future  observations.   Unless  explicitly stated, all magnitudes  are in the AB system  \\citep{abmag}.  We adopt the     following     cosmological     model:     $\\Omega_0\\!=\\!0.24$, $\\Omega_{\\Lambda}\\!=\\!0.76$, and $H_0\\!=\\!100h$~km~s$^{-1}$~Mpc$^{-1}$ where $h\\!=\\!0.73$ \\citep{wmap}. ",
        "conclusions": "\\label{con} The HST-WFPC2  and MMT data suggest  that \\name\\ is  an extremely dark galaxy  worthy of follow-up  observations to  confirm or  exclude this hypothesis.  In  particular, with  deeper optical spectroscopy  of the arc, our  redshift analysis  can be verified.   Moreover, observations blueward  of the  blue channel  spectroscopy will  break  the redshift degeneracy  of $z\\!\\simeq\\!0.15$ and  $z\\!\\simeq\\!2.5$.  Owing  to the necessary  wavelength  ($\\lambda_{\\rm  obs}\\!\\lesssim\\!3800$~\\AA)  and flux limit (${\\rm  AB}\\!\\gtrsim\\!26$~mag), these observations could be conducted  with the  ultraviolet channel  of the  Wide  Field Camera~3 (WFC3)  for  HST.  Since  the  radio  source  has not  been  optically identified, its  redshift can only be very  coarsely constrained.  The majority of millijansky radio sources have early-type morphologies and colors               at              $0.8\\!\\lesssim\\!z\\!\\lesssim\\!1.5$ \\citep[eg.][]{win85,jqr07},  consequently  its Balmer/4000~\\AA\\  break will     occur      at     $7000~{\\rm     \\AA}\\!\\lesssim\\!\\lambda_{\\rm obs}\\!\\lesssim\\!1~\\mu{\\rm m}$.   Furthermore, \\name\\ is  undetected to faint   flux-levels,   and   may   require  a   rather   deep   (${\\rm AB}\\!\\gtrsim\\!25$~mag)  imaging campaign.   Therefore,  medium-band or grism  observations  in  the  infrared  mode with  WFC3  would  better constrain  the lens  redshift.  Should  such observations  support our proposed lensing scenario, this system could be among the most distant known                       gravitational                       lenses \\citep{castles}\\footnote{\\url{http://www.cfa.harvard.edu/castles/}}, a dusty ERO with $M_*\\!\\lesssim\\!10^{10}$~M$_{\\odot}$, or a genuine dark lens,  which  would  be  a  unique confirmation  of  the  $\\Lambda$CDM paradigm."
    },
    "0806/0806.3054_arXiv.txt": {
        "abstract": "Here we present a self-consistent, bimodal stationary solution for spherically symmetric flows driven by young massive stellar clusters with a central supermassive black hole (SMBH). We demonstrate that the hydrodynamic regime of the flow depends on the location of the cluster in the 3D (star cluster mechanical luminosity - BH mass - star cluster radius) parameter space. We show that a threshold mechanical luminosity ($L_{crit}$) separates clusters which evolve in the BH dominated regime frome those whose internal structure is strongly affected by the radiative cooling. In the first case (below the threshold energy) gravity of the BH separates the flow into two distinct zones: the inner accretion zone and the outer zone where the star cluster wind is formed. In the second case (above the critical luminosity), catastrophic cooling sets in inside the cluster. In this case the injected plasma becomes thermally unstable that inhibits a complete stationary solution.  We compared the calculated accretion rates and the BH luminosities with those predicted by the classic Bondi accretion theory and found that Bondi's theory is in good agreement with our results in the case of low mass clusters. However, it substantially underestimates the accretion rates and BH luminosities if the star cluster mechanical luminosity, $L_{SC} \\ge 0.1 L_{crit}$. ",
        "introduction": "Intensive studies of active galactic nuclei (AGNs) in the optical, infrared (IR) and X-ray regimes during the last decade, have  unveiled the presence of massive starbursts around the central supermassive black hole (BH) in a number of Seyfert galaxies. For example, Rodr\\'\\i{}guez Espinosa et al. (1987) found that the 25, 60 and 100~$\\mu$m fluxes of classical optically selected Seyfert galaxies are not correlated with their ultraviolet (UV) and optical continuum and that far-IR colors of the selected galaxies are indistinguishable from those  of starburst galaxies. They suggested then that high far-IR luminosities associated with many Seyfert galaxies indicate  an intrinsic link between the circumnuclear star formation and the AGN activity. Baum et al. (1993) revealed a kiloparsec-scale, diffuse radio emission in 12 of the 13 observed Seyfert galaxies. They found that the intensity of the diffuse radio emission correlates with the far-IR luminosity of the host galaxy, suggesting that this emission is generated in the galactic superwind plasma driven by the nuclear starburst along the minor axis of the galaxy, and claimed that circumnuclear starbursts and starburst driven outflows may be intrinsic to many Seyferts although their relative strengths may vary from galaxy to galaxy.  Levenson et al. (2001) presented X-ray imaging and spectroscopy of a sample of 12 Seyfert 2 galaxies and concluded that in order to fit the observed X-ray spectra it is required to combine the power-law Seyfert component with a thermal starburst emission. Jim\\'enez-Bail\\'on et al. (2005) presented XMM-Newton and Chandra observations of the Seyfert 2 galaxy NGC 1808. They found the hard X-ray emission associated with an unresolved nuclear sources whereas the soft emission  is dominated by a thermal component associated with an extended starburst.  Terlevich et al. (1990) suggested to use the stellar CaII triplet absorption feature in the IR continuum as a direct indicator on the presence of young unresolved stellar population in the nuclear regions of Seyfert galaxies. Heckman et al. (1997) and Gonz\\'alez Delgado et al. (1998) found absorption line features associated with photospheres of O and B stars and their stellar winds in the ultraviolet and optical spectra of four Seyfert 2 galaxies: Mrk 477, NGC 7130, NGC 5135 and IC 3639 and thus presented direct evidence for the existence of nuclear starbursts in these galaxies. They found that the size of the nuclear starbursts in these galaxies ranges from several tens to a few hundred parsecs. The co-existence of W-R features in the optical and Ca II triplet in the near-IR part of the spectra, implies either a continuous star formation during more than $\\sim 10$~Myr or two stellar generations with ages about  5-6~Myr and 10-20~Myr, respectively. Such starbursts are likely to drive the high-velocity outflows detected in the above and in Seyfert 2/starburst ultra luminous infrared galaxies (Gonz\\'alez Delgado et al. 1998; Rupke et al. 2005).  On the other hand, compact, bright stellar clusters or nuclear star clusters, were found in the centers of $\\sim 75$\\% of local spirals and Virgo dwarf elliptical galaxies (B\\\"oker et al. 2002; C\\^ot\\'e et al. 2006). Their radii (a few parsecs) are similar to those of globular clusters, however they are 1 - 2 orders of magnitude brighter, more massive and may have complicated star formation histories with several episodes of star formation (Walcher et al. 2006). Ferrarese et al. (2006) claimed that massive galaxies (M$_{gal} > 10^{10}$\\Msol) host supermassive BHs whereas less massive  galaxies host only nuclear clusters. However Seth et al. (2007) presented evidences on the presence of super massive black holes in $\\sim 25$\\% of galaxies which host nuclear star clusters. More than half of these galaxies ($\\sim 15$\\%) present a mixed AGN-starburst optical spectra and are classified as composite.  Shlosman \\& Begelman (1989), Goodman (2003), Collin \\& Zahn (1999), Tan \\& Blackman (2005) have shown that accretion disks are gravitationally unstable outside of $r \\sim 10^{-2}$~pc and must fragment into self-gravitating clumps that eventually form stars. It was suggested that star formation reduces the rate of accretion and thus the luminosity of the central supermassive black hole (BH) by removing mass from the accretion flows and also due to radiative heating of the accretion discs. However non of these models took into consideration the negative feedback provided by the mechanical energy of the central starburst on the accretion flow.  Thus circumnuclear star formation occurs at different space scales around the SMBH in many AGN galaxies. Here we note that the mechanical power of young nuclear starbursts might prevent through the cluster winds the accretion of interstellar matter from the bulges and disks of their host galaxies onto the central BHs. In such cases  the BHs are fed with the matter injected by numerous stellar winds and SNe explosions that result from the multiple evolving sources. This implies that nuclear starbursts must strongly affect and perhaps even control the power of the central BH. In fact, it  may be the dominant factor to be consider in order to understand the physics and  relative contributions of the BH and starburst activity to the energy budget of the composite AGN/starburst galaxies. L\\'ipari \\& Terlevich (2006) have incorporated different ingredients of this physics into their evolutionary unification scenario which seems to be able to explain many properties of AGNs and QSOs.  The classic spherically-symmetric accretion model (Bondi, 1952; Frank et al. 2002) should then be modified if one is to apply it to the case of a massive BH at the center of a  young stellar cluster. First, it should  take into consideration the energy and mass supplied by massive stars within the star cluster volume. Second, it should account for radiative losses of energy from the hot thermal plasma. Third, the models should incorporate  initial and boundary conditions as described by the star cluster wind theory (Chevalier \\& Clegg, 1985; Canto et al. 2000; Silich et al. 2004; Tenorio-Tagle et al. 2007).  Here we present a self-consistent semi-analytic theory of stationary spherically-symmetric flows driven by the thermalization of the kinetic energy supplied by massive stars inside massive stellar clusters which includes the outflow of the injected matter and its accretion onto a central massive BH. The solutions account for proper initial and boundary conditions for a variety of stellar clusters and black holes and the impact of strong radiative cooling on the dynamics of the thermalized injected matter.  The paper is organized as follows. In section 2 we formulate our model and discuss the input physics and major simplifications. In section 3 we present a set of major equations. Boundary conditions and the selection of the proper solution from a family of integral curves are discussed in section 4. In section 5 we discuss the impact that a central BH provides on the star cluster driven flows. We discuss two hydrodynamic regimes separated in the star cluster mass vs star cluster radius diagram by a threshold line: the BH dominated regime below the threshold line and the radiative cooling dominated regime above the threshold line. In section 6 we use our model to calculate the accretion rates and BH accretion luminosities and compare them to those, predicted in the classical Bondi accretion theory. Section 7 summarizes our results and gives our conclusions. ",
        "conclusions": "We have developed a self-consistent, stationary solution for spherically symmetric accretion flows which are formed inside young stellar clusters with a central supermassive BH.  We have shown that the thermalization of the kinetic energy released by massive stars inside young stellar cluster results in a bimodal solution which presents an accretion of the injected matter onto a central BH in the inner zones of the cluster, and the ejection of the deposited matter from the outer zones of the cluster in the form of the fast superwind. We suggest that superwinds prevent the accretion of the ambient interstellar gas from the bulges and disks of their host galaxies onto the central BHs and that in such cases the BHs are fed with the matter re-inserted by massive stars in the form of numerous stellar winds and SNe explosions. The accretion rate and the BH luminosity are then defined by the central starburst but not by the gravity and the interstellar gas distribution in a host galaxy.  The hydrodynamics of the accreted and ejected matter depend on the location of the stellar cluster hosting a BH in the cluster mass - black hole mass - star cluster radius parameter space. There is a surface in this 3D parameter space which separates clusters evolving in the stationary regime from those which cannot fulfill stationary conditions. If the mechanical luminosity of the cluster exceeds the critical value, a hot plasma inside the cluster is thermally unstable. The flow is highly non-stationary and presents a complicate velocity patten with an outer stagnation point defined by strong radiative cooling and the inner one whose position is defined by the mass of the BH. The hydrodynamical structure and time evolution of the resulting flow in such cases must be calculated numerically.  Clusters whose mechanical luminosity is smaller than the critical value, compose stationary accretion flow in the central zones and form stationary outflows, the star cluster winds, in the outer zones of the cluster.  We used our model to calculate the accretion rates and the accretion luminosities of BHs at the centeres of young star forming regions. The classic, Bondi's accretion theory shows a good agreement with our semi-analytic model, but only in the case of low mass clusters located well below ($L_{SC} \\le 0.1 L_{crit}$) the critical luminosity in the $L_{SC} - M_{BH} - R_{SC}$ parameter space. Thus one has to use a semi-analytic approach in order to calculate the accretion rates and BH luminosities in the case of more energetic clusters.  In the case of extended starbursts, the BH luminosities fall well below the Eddington limit. However, the accretion luminosities grow rapidly for more compact clusters and for very compact clusters can approach the Eddington luminosity, as shown in Figure 9, panel b.  The model, here developed, is required in order to advance our knowledge regarding the relative contributions of supermassive BHs and central star bursts in composite, AGN/starburst galaxies and will be used for interpretation of observational properties of such objects in a further communication."
    },
    "0806/0806.4456_arXiv.txt": {
        "abstract": " ",
        "introduction": "\\label{GCsec:Intro} The advantages of using cluster white dwarfs over field white dwarfs in studying the formation, physical properties, and evolution of these stars come from a number of properties of star clusters. {\\em i) --} White dwarfs in globular or open clusters are located at the same distance and have (in general) the same reddening. Moreover and even more importantly, when moving from hot to cool white dwarfs the colours of cluster white dwarfs are systematically bluer than those of main sequence stars (see Fig.~\\ref{GCfig:cmdGC_OCen}). This means that to properly identify cluster white dwarfs we can use the colour-magnitude diagram instead of a colour-colour plane. Therefore, the identification of cool cluster white dwarfs is not hampered by the thorny problem of colour degeneracy with main sequence stars that affects field white dwarfs (\\cite{Hanlie03}).  {\\em ii) --} For cluster white dwarfs we can trace back the evolutionary properties of the progenitors, since both the chemical composition and the typical mass at the turn-off of the cluster are well-known.  Moreover, we can in general safely assume that cluster stars originate from the same star formation event. While there is some evidence that the most massive globular clusters (e.g. $\\omega$\\,Cen, NGC\\,1851, NGC\\,2808, NGC\\,6388, NGC\\,6441) may have a more complicated star formation and chemical enrichment history, most globular clusters are essentially mono-metallic systems with a negligible spread in age. This provides the opportunity to constrain the initial-to-final mass relation of white dwarfs and to improve the knowledge of the physical mechanisms governing their final fate (\\cite{Hans04}; \\cite{Kalirai07b}).  {\\em iii) --} According to current evolutionary predictions the number of white dwarfs in a globular cluster with an age of 12\\,Gyr and a Salpeter-like initial mass function ($\\alpha=2.35$) is about a factor of 300 larger than the number of horizontal branch stars (\\cite{Broc99}). This means that the expected local density of white dwarfs in a globular cluster is several orders of magnitude larger than the local densities of the halo, thick disk, and thin disk white dwarf populations, thus allowing us to observe large homogeneous samples of white dwarfs without the need for wide-field surveys.  The main drawback for white dwarfs located in globular clusters is that they are faint objects severely affected by crowding problems. Photometric observations are well possible with the Hubble Space Telescope (HST) and profit from software packages that allow to obtain precise measurements even in crowded fields. Spectroscopy, however, poses a different problem: HST is a too small telescope and lacks the multiplexing capacities that allow efficient observations in globular clusters. The disentangling of overlapping spectra in ground based observations is a rather difficult process, as the contributions from the different stars vary strongly with wavelength. Early spectroscopic investigations (e.g. \\cite{Moehler00}; \\cite{Moehler04}) were hampered by the fact that their targets were selected from HST photometry, which had usually been observed rather close to the cluster cores. Good photometry of the outer regions of globular clusters will facilitate spectroscopic observations enormously, as the problem of background subtraction is immensely reduced if there are no bright stars close to the white dwarfs.  In a recent investigation \\cite{Davies07} called attention to the evidence that the radial distribution of young white dwarfs in NGC\\,6397 is significantly more extended than the radial distribution of both older white dwarfs and the most massive main sequence stars. To account for this peculiar trend they suggest that the white dwarfs in this cluster do not experience a quiescent birth, but they receive a natal kick. This scenario has been supported by theoretical investigations of \\cite{Heyl07a,Heyl07b,Heyl07c}. In particular he found that asymmetric winds on the asymptotic giant branch can generate a kick affecting the trajectory of the resulting white dwarfs. This mechanism would explain why young white dwarfs are less centrally concentrated than their progenitors. If this effect is indeed present in other globular clusters as well, it will be very helpful for the study of white dwarfs in globular clusters, as it would put the youngest (and therefore brightest) white dwarfs into the less crowded outer regions.  Beyond the obvious observational complications, we still lack a detailed understanding of the impact of the high density environment of globular clusters on the formation and evolution of cluster white dwarfs (\\cite{Monelli05}).  As the exploration of white dwarfs in globular clusters is an extremely active and rapidly growing field we want to state here that this review contains information from papers or preprints (accepted for publication) available by March 2008. ",
        "conclusions": ""
    },
    "0806/0806.1045_arXiv.txt": {
        "abstract": "Molecular clouds are observed to be turbulent, but the origin of this turbulence is not well understood. As a result, there are two different approaches to simulating molecular clouds, one in which the turbulence is allowed to decay after it is initialized, and one in which it is driven. We use the adaptive mesh refinement (AMR) code, Orion, to perform high-resolution simulations of molecular cloud cores and protostars in environments with both driven and decaying turbulence. We include self-gravity, use a barotropic equation of state, and represent regions exceeding the maximum grid resolution with sink particles. We analyze the properties of bound cores such as size, shape, linewidth, and rotational energy, and we find reasonable agreement with observation.  At high resolution, the different rates of core accretion in the two cases have a significant effect on protostellar system development. Clumps forming in a decaying turbulence environment produce high-multiplicity protostellar systems with Toomre-Q unstable disks that exhibit characteristics of the competitive accretion model for star formation. In contrast, cores forming in the context of continuously driven turbulence and virial equilibrium form smaller protostellar systems with fewer low-mass members. Our simulations of driven and decaying turbulence show some statistically significant differences, particularly in the production of brown dwarfs and core rotation, but the uncertainties are large enough that we are not able to conclude whether observations favor one or the other. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0806/0806.3795_arXiv.txt": {
        "abstract": "%  I present some new results related to our understanding of the masses of galaxies both in the local and high-redshift Universe. At high-redshift new Spitzer data on galaxies in the Gemini Deep Deep Survey allow us a more accurate measure of stellar mass to light ratios (using rest frame near-IR light) showing a refinement of the measurements but not great discrepancies. In the local universe a new method is explored to estimate the {\\it baryonic} mass function of galaxies including contributions from unseen HI. This points to an interesting result: that the baryonic mass function of galaxies may in fact be quite steep, of comparable slope to the mass function of dark matter haloes. ",
        "introduction": "The mass assembly history of galaxies is an important probe of models of galaxy formation and evolution. In recent years near-infrared selected surveys have allowed measurements of this to high-redshift (Fontana et al, 2003, 2006, Glazebrook et al. 2004). In particular the main aspect probed is the {\\em stellar mass}, as beyond the very nearby Universe most galaxies are only revealed by their optical starlight emission. The nature of the game is to use multi-colour data to determine the stellar populations of distant galaxies and hence their mass-to-light ratio. There is of course uncertainty in this determination and to address it this is best done at redder wavelengths ($\\lambda \\ga 0.8\\micron$) where mass-to-light ratios vary less. This is due to the dominance of the near-IR light by older long-lived stellar populations. In contrast as one goes bluer ($\\lambda < 0.4\\micron$) the variation increases dramatically. In the rest-frame UV the range is many orders of magnitudes reflecting the fact that these wavelengths really measure {\\em star formation rate} rather than stellar mass. Infrared $K$-band surveys such as 2MASS have allowed very accurate measurements of the local stellar mass function (e.g. Cole et al. 2001). ",
        "conclusions": ""
    },
    "0806/0806.3276_arXiv.txt": {
        "abstract": "We present the first results from our GALEX program designed to obtain ultraviolet (UV) spectroscopy of nearby core-collapse supernovae (SNe). Our first target, SN 2005ay in the nearby galaxy NGC 3938, is a typical member of the II-P SN subclass. Our spectra show remarkable similarity to those of the prototypical type II-P event SN 1999em, and resemble also {\\it Swift} observations of the recent type II-P event SN 2005cs. Taken together, the observations of these three events trace the UV spectral evolution of SNe II-P during the first month after explosion, as required in order to interpret optical observations of high-redshift SNe II-P, and to derive cross-filter K-corrections. While still highly preliminary, the apparent UV homogeneity of SNe II-P bodes well for the use of these events as cosmological probes at high redshift. \\\\ ",
        "introduction": "In order to interpret optical observations of high-redshift supernovae (SNe), sampling the restframe ultraviolet (UV) radiation of these events, we need to have UV observations of local SNe of all types. Studies of high-redshift SNe promise, in turn, to shed light on key open questions, such as the evolution of cosmic metallicity, star formation at high redshift, and SN ``feedback'' processes shaping galaxy formation.  Type Ia SNe, famed for their cosmological utility as precision distance estimators, are the best-studied of all SN subclasses in restframe UV (e.g., Kirshner et al. 1993; Ellis et al. 2008; Foley et al. 2007, 2008). SNe Ia are broadly thought to result from thermonuclear explosions of white dwarf stars approaching the critical Chandrasekhar mass due to accretion from (or a merger with) a binary companion, and show remarkable homogeneity in their observational properties. However, UV studies may hint at unexplained dispersion in the restframe UV band (Ellis et al. 2008).  All other types of SNe (see Filippenko 1997 for a review) probably result from the gravitational collapse of short-lived massive stars (e.g., Crockett et al. 2008; Li et al. 2007, Gal-Yam et al. 2007 and references therein). In general, core-collapse SNe are extremely heterogeneous in every observational respect, and, in particular, different types of core-collapse events have diverse UV properties (e.g., UV-bright type IIn SN 1998S, Lentz et al. 2001 vs. UV-deficient type Ic SN 1994I, Millard et al. 1999). The dispersion in UV properties among objects within specific core-collapse SN subtypes are so far unknown.  Unfortunately, UV spectroscopy of reasonable quality was obtained for only a handful of core-collapse SNe (Panagia 2003; 2007 for reviews) and some of the best-observed events (notably SN 1987A, e.g., Eastman \\& Kirshner 1989) are quite peculiar. This observational deficit introduces significant uncertainties into the interpretation of high-redshift SN observations, which are forced to rely either on little-tested models for the UV spectrum of core-collapse events, or on the use of the few observed UV spectra for the entire population, neglecting possible dispersion in spectral evolution. The sparse database of UV core-collapse SN spectroscopy continues to limit the scientific utility of large samples of core-collapse SNe at high redshifts. These include both samples currently assembled using the {\\it Hubble Space Telescope} ({\\it HST}; e.g., GOODS, Dahlen et al. 2004; Riess et al. 2007) and from the ground (e.g., Poznanski et al. 2007a), and future datasets expected to emerge from a possible JDEM mission that includes a SN component and from deep infrared observations with the {\\it James Webb Space Telescope}. Furthermore, in order to probe the physics of core-collapse SNe, broad spectral coverage, particularly of the UV range, where line blanketing by iron peak elements plays a crucial role in the formation of the spectrum (e.g., Pauldrach et al. 1996), is essential. UV spectra are most sensitive to the metal content of the SN ejecta, a key probe of the nucleosynthetic evolution.  To alleviate this problem is the main motivation for our target of opportunity GALEX program (GALEX-GI-44, cycle 1; GALEX-GI-67, cycle 2; GALEX-GI-61, cycle 3; GALEX-GI-20, cycle 4; PI Gal-Yam), designed to obtain multi-epoch UV spectroscopy of nearby core-collapse SNe. Here, we report the first results from this program. In $\\S~2$ we describe our observation, and in $\\S~3$ we present our results. Discussion and conclusions follow in $\\S~4$. UT dates are used throughout the paper. ",
        "conclusions": "We are carrying out a spectroscopic survey of nearby core-collapse SNe using GALEX grism spectroscopy in target-of-opportunity mode. About $1-2$ nearby events are observed each year, increasing our knowledge of the spectral evolution of core-collapse SNe of the various subtypes. Our collaboration also provides supporting IR and optical observations of our GALEX targets. We have presented first results from this project -- spectra of the nearby type II-P SN 2005ay. Combined with additional observations of two similar objects from the literature, we trace the UV spectral evolution of SNe II-P and find a remarkable similarity among these objects, the most common type of core-collapse SNe and the only subtype with a sample of events having UV spectroscopic measurements. Such restframe-UV homogeneity, if supported by additional objects, indicates that the use of these SNe as cosmological probes is a promising prospect."
    },
    "0806/0806.3106_arXiv.txt": {
        "abstract": "We present a 1200-$\\mu$m image of the Great Observatories Origin Deep Survey North (GOODS-N) field, obtained with the Max Planck Millimeter Bolometer array (MAMBO) on the IRAM 30-m telescope. The survey covers a contiguous area of 287\\,arcmin$^2$ to a near-uniform noise level of $\\sim$0.7\\,mJy\\,beam$^{-1}$. After Bayesian flux deboosting, a total of 30 sources are recovered ($\\ge$3.5-$\\sigma$). An optimal combination of our 1200-$\\mu$m data and an existing 850-$\\mu$m image from the Submillimetre Common-User Bolometer Array (SCUBA) yielded 33 sources ($\\ge$4-$\\sigma$).  We combine our GOODS-N sample with those obtained in the Lockman Hole and ELAIS\\,N2 fields (Scott et al.\\ 2002; Greve et al.\\ 2004) in order to explore the degree of overlap between 1200-$\\mu$m- and 850-$\\mu$m-selected galaxies (hereafter SMGs), finding no significant difference between their $\\fdr$ distributions. However, a noise-weighted stacking analysis yields a significant detection of the 1200-$\\mu$m-blank SCUBA sources, $\\fdr = 3.8\\pm 0.4$, whereas no significant 850-$\\mu$m signal is found for the 850-$\\mu$m-blank MAMBO sources ($\\fdr = 0.7\\pm 0.3$). The hypothesis that the $\\fdr$ distribution of SCUBA sources is also representative of the MAMBO population is rejected at the $\\sim$4-$\\sigma$ level, via Monte Carlo simulations. Therefore, although the populations overlap, galaxies selected at 850 and 1200\\,$\\mu$m are different, and there is compelling evidence for a significant 1200-$\\mu$m-detected population which is not recovered at 850\\,$\\mu$m. These are submm drop-outs (SDOs), with $\\fdr = 0.7-1.7$, requiring very cold dust or unusual spectral energy distributions ($T_{\\rm d}\\simeq 10$\\,{\\sc k}; $\\beta \\simeq 1$), unless SDOs reside beyond the redshift range observed for radio-identified SMGs, i.e.\\ at $z > 4$. ",
        "introduction": "A decade ago, surveys undertaken at submillimetre (submm; 850-$\\mu$m) wavelengths with SCUBA (Holland et al.\\ 1999) transformed the accepted understanding of galaxy formation and evolution by revealing an unexpected yet significant population of dusty, far-infrared-luminous starburst galaxies (Smail, Ivison \\& Blain 1997; Hughes et al.\\ 1998; Barger et al.\\ 1998).  Since then, great strides have been made in understanding the nature of submm-selected galaxies -- see Blain et al.\\ 2002 for a review -- largely facilitated by accurately pinpointing SMGs in deep, high-resolution radio maps (e.g.\\ Ivison et al.\\ 1998a, 2000, 2002; Smail et al.\\ 2000) which helps overcome the coarse resolution of SCUBA ({\\sc fwhm} $\\simeq$ 15\\,arcsec). The majority (60--80 per cent) of bright ($S_{\\rm 850\\mu m}\\gs 5$\\,mJy) SMGs have been identified in this way and one of the most important steps forward was the acquisition of about 100 optical/near-infrared spectra of such radio-identified SMGs. This placed them typically in the redshift range $z\\simeq 1-3$, with a median of 2.3 (Chapman et al.\\ 2003, 2005).  However, while submm surveys are equally sensitive to sources of a fixed luminosity in the redshift range $z\\simeq 1-8$, due to the negative $k$-correction offsetting the cosmic dimming (Blain \\& Longair 1993), the detection rate of SMGs in even the deepest radio surveys drops off rapidly beyond $z\\simeq 3.5$. Thus, the spectroscopic requirement for a robust radio detection biases the known population to $z \\ls 3.5$. This led to the suggestion that the 20--40 per cent of the bright, radio-blank SMGs (1.4\\,GHz flux density, $S_{\\rm 1.4GHz} \\ls 20$\\,$\\mu$Jy) could be the high-redshift tail of the population ($z > 4$ -- see Carilli \\& Yun 1999; Ivison et al.\\ 2002; Eales et al.\\ 2003; Aretxaga et al.\\ 2003, 2007), although using photometric redshifts Pope et al.\\ (2006) argued that $<$14 per cent of the bright SMG population is at $z > 4$.  Perhaps the strongest candidates for $z > 4$ SMGs have come from submm/mm interferometry, with the IRAM Plateau de Bure Interferometer (PdBI) and the Submillimeter Array (SMA), of a few very bright SMGs which -- by virtue of their radio, submm/mm and near-infrared properties -- were deemed to reside at very high redshift (Dannerbauer et al.\\ 2002; Wang et al.\\ 2007; Younger et al.\\ 2007).  Soon after the first observations with SCUBA, surveys were being conducted at 1200\\,$\\mu$m with MAMBO (Bertoldi et al.\\ 2000) resulting in samples of sources that were thought to be identical to SMGs, except selected at mm wavelengths. Compared to submm surveys, however, observations at mm wavelengths should be sensitive to far-IR-luminous galaxies out to higher redshifts, due to the $k$-correction, as well as to systems with cooler dust temperatures.  As a result, at $z\\gs 3$ the 850-/1200-$\\mu$m flux ratio becomes a relatively strong function of redshift for a typical starburst far-infrared/mm spectral energy distribution (SED), and may therefore (if the SED is known) be used as a crude redshift estimator or -- if the redshift is known -- as an indicator of the shape (i.e.\\ temperature and spectral index) of the far-infrared/mm SED.  This fact was exploited by Eales et al.\\ (2003; hereafter E03), who obtained 850-$\\mu$m photometric measurements of a sample of 1200-$\\mu$m-selected sources and found a significant fraction to have very low 850/1200-$\\mu$m flux ratios ($\\fdr \\ls 2$). They argued that this was either indicative of SMGs at $z \\gg 3$ or -- since a large fraction of their sample was radio-identified with radio-to-submm flux ratios consistent with $z \\le 3$ -- SMGs having fundamentally different dust emission properties than local galaxies. The existence of a significant number of these SDOs, which we define here as sources with $S_{\\rm 850\\mu m}/S_{\\rm 1200\\mu m} \\ls 2$, was questioned (albeit not ruled out) by Greve et al.\\ (2004; hereafter G04), who performed the first unbiased comparison of 850- and 1200-$\\mu$m-selected sources, using catalogues extracted from SCUBA and MAMBO maps.  The degree of overlap between source populations extracted from submm and mm surveys (and in particular the idea that mm observations are tracing a significant sub-population of more distant and/or cooler sources) remains controversial, therefore, and is further complicated by the difficulty of comparing (sub)mm surveys with different depths and noise properties. Nonetheless, settling this issue is important; regardless of the outcome, it tells us about the population of dusty galaxies at the highest redshifts, with significant implications for our understanding of galaxy formation and evolution.  In this paper we present an unbiased 1200-$\\mu$m MAMBO survey of the northern field of the Great Observatories Origins Deep Survey (GOODS-N -- Giavalisco et al.\\ 2004). As well as presenting the 1200-$\\mu$m map and source catalogue, we also perform a rigorous comparison between our data and the existing 850-$\\mu$m SCUBA map of this region (Borys et al.\\ 2003; Pope et al.\\ 2005), in order to address the issues outlined above.  In a follow-up paper, we shall exploit other high-quality data in the GOODS-N field to explore the multi-wavelength properties of 1200-$\\mu$m selected sources.  Throughout this paper we have adopted a flat cosmology with $\\Omega_{\\rm M} = 0.27$, $\\Omega_{\\Lambda}=0.73$, and $H_0=71\\,\\mbox{km}\\,\\mbox{s}^{-1}\\,\\mbox{Mpc}^{-1}$ (Spergel et al.\\ 2003). ",
        "conclusions": "\\label{section:discussion}  The spectroscopic redshift survey of bright, radio-identified SMGs by Chapman et al.\\ (2003, 2005) located sources out to $z\\sim 3.6$, with an interquartile range, $1.7\\le z \\le 2.8$. I05 argued that of the bright SMG population ($S_{\\rm 850\\mu m} \\ge 5$\\,mJy), probably no more than $\\sim$10 per cent could be at $z > 3.5$, given that $\\sim$80 per cent have radio counterparts and $\\sim$10 per cent were likely to be spurious. Pope et al.\\ (2006) reached a similar conclusion, estimating that $\\ls$14 per cent of SMGs reside at $z\\gs 4$. Modeling the radio/mm/far-infrared colours of 120 SMGs from SHADES and adopting priors for the redshift probability of their radio-undetected sources, Aretxaga et al.\\ (2007) argued that more than half of the bright SMG population lies in the range $1.6\\ls z \\ls 3.4$, with little room for a $z > 4$ population. Thus, the evidence for a {\\it significant} population of SMGs at very high redshifts appears slim, but has not been ruled out completely.  If we consider sources {\\it strictly selected at 850\\,$\\mu$m}, seven have spectroscopically confirmed redshifts in the range $3 \\ls z \\ls 3.6$ (Ledlow et al.\\ 2002; Chapman et al.\\ 2003, 2005). A strong candidate for a $z\\gg 3$ SMG was reported by Knudsen, Kneib \\& Egami (2006) who inferred a likely spectroscopic redshift of $\\simeq 4$ for SMM\\,J16359$+$66130. Apart from this source, however, only one other convincing candidate for a $z > 4$ SMG has been presented in the literature at the time of writing, namely GN\\,850.10 (Wang et al.\\ 2004; Pope et al.\\ 2005). This source was deemed to reside at $z\\sim 4-6$, based on its near- and mid-infrared colours, which were facilitated by its accurate location at 890-$\\mu$m with the SMA (Wang et al.\\ 2007), at 1.1\\,mm with the IRAM PdBI and in the radio with the VLA (Biggs \\& Ivison 2006; Dannerbauer et al.\\ 2008).  Younger et al.\\ (2008) presented SMA observations of LH\\,850.2 -- one of the brightest 850-$\\mu$m (and 1200-$\\mu$m) sources in the Lockman Hole (G04; Coppin et al.\\ 2006) -- and derived an optical/infrared photometric redshift of $z\\simeq 3.3$, at the high-end of (but within) the Chapman et al.\\ redshift distribution.  Aretxaga et al.\\ (2007) presented a handful of SMGs from SHADES with photometric redshifts $z\\gs 4$, but with very large individual uncertainties.  \\begin{figure} \\begin{center} \\includegraphics[width=1.0\\hsize,angle=0]{S850S1200_z.ps} \\caption[]{Green, blue and red filled circles represent $\\fdr$ for sources identified robustly at 850 and 1200\\,$\\mu$m by E03, S02+G04 and in this work, respectively.  The filled triangles are the seven sources from Younger et al.\\ (2007), where we have converted the measured 890- and 1100-$\\mu$m flux densities to 850- and 1200-$\\mu$m flux densities assuming an optically thin grey-body with $\\beta=1.5$. Sources which have robust radio counterparts have been circled.  Where available, we have placed sources at their spectroscopic (Chapman et al.\\ 2005) or photometric redshifts (Aretxaga et al.\\ 2007; Pope et al.\\ 2006; Wall, Pope \\& Scott 2008), otherwise the sources have been placed at fixed redshifts below $z=0.3$ for clarity. Sources with spectroscopic redshifts are shown as large symbols while sources with only a photometric redshift -- or lacking a redshift estimate entirely -- are shown as small symbols. The grey shaded region corresponds to the range of $\\fdr$ spanned by grey-bodies with $\\beta = 1.0$ and $T_{\\rm d} = 10$\\,{\\sc k} (lower limit) and $\\beta = 2.0$ and $T_{\\rm d} = 70$\\,{\\sc k} (upper limit). Large red squares and stars represent the range in $\\fdr$ derived from the (centroid and peak) stacking analysis for the 850-$\\mu$m-blank MAMBO sample and the 1200-$\\mu$m-blank SCUBA sample, respectively.  The open upward- and downward-pointing triangles represent the lower and upper limits on the flux density ratio, respectively, obtained when using deboosted peak values in the stacking analysis (see \\S~\\ref{section:stacking}).  } \\label{figure:S850S1200S14-z} \\end{center} \\end{figure}  Turning our attention to sources {\\it strictly selected at mm wavelengths}, we find a larger number of possible $z > 4$ candidates. Using IRAM PdBI to accurately locate their positions, Dannerbauer et al.\\ (2002, 2004) presented three examples of bright ($S_{\\rm 1200\\mu m} \\gs 3.5$\\,mJy) MAMBO sources which, based on their radio, (sub)mm and near-infrared properties, were argued to lie at $z\\gs 4$. The non-detection of these sources with {\\it Spitzer}/IRS spectroscopy (Valiante et al.\\ 2007) is consistent with this claim.  Younger et al.\\ (2007) presented 890-$\\mu$m interferometric SMA observations of seven very bright ($S_{\\rm 1100\\mu m} \\ge 7.6$\\,mJy) AzTEC sources selected at 1100\\,$\\mu$m and argued that the five radio-dim ($S_{\\rm 1.4GHz}\\le 41\\,\\mu$Jy) sources in their sample were likely to reside at higher redshifts than the radio-identified sources due to their systematically higher submm-to-radio flux ratios, lower IRAC 3.6--8.0-$\\mu$m flux densities and non-detections at 24\\,$\\mu$m.  Common to all of these $z > 4$ candidates, whether they are SCUBA-, MAMBO- or AzTEC-selected sources, is that they are very bright (i.e.\\ luminous) and have been detected at both submm and mm wavelengths. The trend that the high-redshift candidates tend to be very bright was first noted by Ivison et al.\\ (2002), who from deep radio imaging of a large sample of SMGs found evidence that very bright ($S_{\\rm 850\\mu m}\\gs 8$\\,mJy) SCUBA sources tend to have higher submm-to-radio flux density ratios than less luminous SMGs, indicating that they are at higher redshifts (Carilli \\& Yun 1999) and that strong luminosity evolution may have taken place (see also Pope et al.\\ 2006).  From the above summary (and following \\S~\\ref{section:stat-mambo-scuba-maps}), one wonders whether submm- and mm-selected sources are drawn from identical populations? Does the preponderance of $z > 4$ candidates amongst the mm-selected sources reflect a real difference between the two populations, or is this merely a coincidence? The multi-wavelength properties of mm-selected sources have been studied far less thoroughly than SMGs, because they were discovered more recently and are often thought to be mere mm-wave analogues to SMGs. In particular, no systematic spectroscopic surveys have been undertaken of mm-selected sources and their redshift distribution is therefore not known.  A related question is: why are all of the $z > 4$ candidates amongst the brightest (sub)mm sources known? Is this merely a selection effect, reflecting the fact that we can currently detect only the brightest sources with (sub)mm and radio interferometers, or does it reflect a genuine evolutionary effect (Wall, Pope \\& Scott 2007)?  In \\S~\\ref{section:stacking} we showed that while the 1200-$\\mu$m-blank SCUBA sources are statistically detected at 1200\\,$\\mu$m (S/N $\\sim$ 10), the 850-$\\mu$m-blank MAMBO sources are not robustly detected at 850\\,$\\mu$m (S/N $\\sim$ 2.8). In Fig.~\\ref{figure:S850S1200S14-z} we plot $\\fdr$ for these two samples, using the stacking results listed in Table~\\ref{table:stacking}. The ranges of possible values for the average $\\fdr$ ratio for the two samples only barely overlaps, even when accounting for the high and low biases introduced when using the centroid and peak values, respectively (see \\S \\ref{section:stacking}).  If we deboost the peak fluxes which go into the stack, the allowed ranges do not overlap at all. On average, therefore, the 850-$\\mu$m-blank MAMBO sample has significantly lower $\\fdr$ ratios than the 1200-$\\mu$m-blank SCUBA sample.  The strongest support for this result was given in \\S~\\ref{section:stat-mambo-scuba-maps}, based on Monte Carlo simulations of the association fractions in the two maps.  Since we are working in the low-S/N regime, the fraction of sources recovered in two different maps will never be unity, so it is crucial to simulate the experiments. Careful simulations are required because of the inhomogeneous noise in the maps. The main outcome of our Monte Carlo simulations is that while the fraction of SCUBA sources detected with MAMBO is consistent with the $\\fdr$ distribution of sources found in both maps, there are more MAMBO sources {\\it not} detected by SCUBA than would be expected. Since the simulations account for the noise properties of the two maps, we were able to estimate the fraction of MAMBO sources which are undetected by SCUBA (compared with the expectation if the populations were the same, but in the presence of the observed noise properties) and found this to be $\\sim$40 per cent.  Thus, we have presented strong evidence which suggests that -- while there is substantial overlap between the 850-$\\mu$m- and 1200-$\\mu$m- selected sources in GOODS-N -- the two populations {\\it are not} identical. A significant fraction of the 1200-$\\mu$m-selected sources are unaccounted for by the 850-$\\mu$m data. Due to the statistical nature of our analysis and the careful way in which we have taken into account the uneven depths and noise properties of the two maps of the GOODS-N region, we are confident that this is a general result that reflects the existence of two distinct, albeit overlapping, populations.  We still need to establish whether SDOs are genuinely at $z \\gs 4$ or whether they have cooler dust temperatures at redshifts typical of radio-identified SMGs. Due to the degeneracy between dust temperature and redshift (Blain 1999), and without direct spectroscopic redshifts, it is not trivial to determine which of the two explanations is appropriate.  We note from Fig.~\\ref{figure:S850S1200S14-z} that if the typical SEDs of SDOs are given by $T_{\\rm d} = 10$\\,{\\sc k} and $\\beta=1.0$, which is believed to be extreme not only compared to starbursts at the present day but also to SMGs (Kov\\'{a}cs et al.\\ 2006), then their most likely average flux ratio ($\\fdr\\sim 1.2$) would indicate an average redshift of $z\\sim 1.3$. One such source is seen in Fig.~\\ref{figure:S850S1200S14-z}, namely Lock850.27/LH\\,1200.7 (G04; I05) which has a flux density ratio of $\\sim 1$ and lies at $z=1.21$. Given the general lack of spectroscopic redshifts for SMG samples, further examples of such sources may exist, making them viable candidates for SDOs. Adopting $T_{\\rm d} = 20$\\,{\\sc k} and $\\beta = 1.5$, close to the coolest dust temperature measured for SMGs so far (Kov\\'{a}cs et al.\\ 2006), results in an average redshift of $z\\sim 4.6$. If we assume that SDOs have SED properties similar to those of radio-identified SMGs, i.e.\\ $T_{\\rm d} = 35$\\,{\\sc k} and $\\beta = 1.5$ (Kov\\'{a}cs et al.\\ 2006), we find their average redshift to be $\\sim 9$, in the regime where the $k$-correction starts to become positive (Blain \\& Longair 1993).  In the first of the three scenarios outlined above, we note that the angular sizes of cool ($T_{\\rm d} \\ls 25$\\,{\\sc k}) SMGs at $z< 2$ should scale roughly as $\\theta \\sim (S_{\\rm 850\\mu m}/\\mbox{mJy})^{1/2}$ (Kaviani et al.\\ 2003) which would imply that SDOs have typical sizes of $\\theta \\sim 1.3$\\,arcsec (for $S_{\\rm 850\\mu m}=1.7$\\,mJy -- see Table~\\ref{table:stacking}), corresponding to a physical diameter of 11\\,kpc at $z=1.5$.  Thus if SDOs reside at $z\\ls 2$ they would have larger physical sizes than those observed for SMGs ($\\sim$4\\,kpc -- Tacconi et al.\\ 2006; Biggs \\& Ivison 2008), but similar angular sizes ($\\sim$1\\,arcsec). If SDOs predominantly lie at the same redshifts as SMGs ($z\\sim 2.5$) but have cooler dust temperatures ($T_{\\rm d}\\ls 25$\\,{\\sc k}), their angular sizes would be larger by a factor of two (Kaviani et al.\\ 2003).  Thus, if SDOs are in fact very cool systems with a redshift distribution similar to SMGs (or peaking at lower redshift), we would expect them to be quite extended in high-resolution radio and mm images.  If some SDOs, on the other hand, lie in the range $z \\sim 4-10$ and have typical dust temperatures in the range $T_{\\rm d} \\sim 20-40\\,${\\sc k}, the expected angular and physical sizes are $\\theta \\sim 1.5-2.3$\\,arcsec and 6.4--18.0\\,kpc, respectively, inconsistent with the typical sizes measured for SMGs.  Thus, if the Kaviani et al.\\ (2003) results are applicable, it would suggest that {\\it SDOs are very extended systems over a broad range of plausible dust temperatures and/or redshifts.}  We note that high-resolution 1.3-mm interferometry of SDOs should be able to test this and may provide a useful way of discriminating between typical SMGs -- generally compact at mm wavelengths ($\\ls$1\\,arcsec -- Tacconi et al.\\ 2006; Younger et al.\\ 2007) and SDOs, which our analysis here suggest may be larger.  Observational evidence for such extended systems was recently uncovered by Daddi et al.\\ (2008), who demonstrated the existence of gas-rich disks at $z=1.5$ with physical sizes 2--3 times those of SMGs, as implied by CO observations as well as rest-frame UV light. Some of these $BzK$ galaxies -- so called for the way they are selected (Daddi et al.\\ 2004) -- have been detected at very faint flux levels at 1200-$\\mu$m ($\\sim$1.5\\,mJy -- Dannerbauer et al.\\ 2005) and have similar star-formation efficiencies to those of local spirals, i.e.\\ an order of magnitude smaller than SMGs. They could be candidates for the cold, extended `cirrus' dust models proposed by Efstathiou \\& Rowan-Robinson (2003). Certainly, these properties seem to fit with those of SDOs, opening up the interesting possibility of an overlap between the two populations, with SDOs making up some fraction of the near-infrared-selected population at $z\\sim 1.5-2.4$.  However, it seems clear from a growing number of in-depth studies, that some of the brightest mm- (and submm-)selected sources are almost certainly at $z > 4$ (Dannerbauer et al.\\ 2002; Wang et al.\\ 2004; Younger et al.\\ 2007). The same may be true for many of the SDOs, although we stress that one cannot rule out a scenario in which SDOs are predominantly cool, low-redshift sources.  Adopting the high-redshift scenario ($z\\simeq 5$) for SDOs, and assuming $T_{\\rm d}=20\\,${\\sc k}, $\\beta =1.5$ and $S_{\\rm 850\\mu m} = 1.7-4.0$\\,mJy (Table~\\ref{table:stacking}), we find far-infrared luminosities and dust masses of $L_{\\rm FIR} = (1-2)\\times 10^{12}\\,\\Lsolar$ and $M_{\\rm d} = (3-8)\\times 10^{9}\\,\\Msolar$, respectively. In this scenario, SDOs are luminous and massive systems, despite being relatively fainter at submm wavelengths. The presence of such systems at $z\\simeq 5$ (corresponding to $\\sim$1\\,Gyr after the Big Bang) would have implications for our understanding of galaxy formation and evolution. Due to their large masses and high redshifts, SDOs and SMGs at $z > 4$ would be potential candidates for galaxies caught in the act of collapsing early in the Universe's history (Eggen, Lynden-Bell \\& Sandage 1962).    We have undertaken the first deep and uniform 1200-$\\mu$m survey of the GOODS-N region, identifying 30 sources (after flux deboosting) in our image. By combining our 1200-$\\mu$m MAMBO map with the existing 850-$\\mu$m SCUBA data, we have extracted a robust sample of 33 sources (sub)mm sources detected at a combined S/N$\\,{\\ge}\\,$4.  The principal question we wanted to address in this paper was whether the 850- and 1200-$\\mu$m selected sources constitute identical, or merely overlapping, populations. We performed three independent statistical analyses of the 850- and 1200-$\\mu$m maps and their corresponding source catalogues.  The first -- a simple comparison of the 850-/1200-$\\mu$m flux density ratio distributions for the 850-$\\mu$m- and 1200-$\\mu$m-selected samples, suggests that a larger fraction of MAMBO sources have low values of $\\fdr$ compared to the SCUBA sources.  This method suffers, however, from small sample sizes and the effects of maps with uneven noise properties, and so we were unable to see a significant difference between the two populations.  However, using the other two methods -- both of which take into account the uneven noise properties -- we do find evidence to suggest that the source populations selected by MAMBO and SCUBA are not identical. The strongest evidence comes from our Monte Carlo analysis of the fractions of SCUBA galaxies with MAMBO associations, and vice versa, which showed that while the fraction of SCUBA sources with MAMBO associations was consistent with the map properties and the $\\fdr$ distribution of sources robustly detected at both wavelengths, the fraction of MAMBO sources with SCUBA associations was significantly lower than expected. In fact, we found that about 40 per cent of the MAMBO sources were not recovered at 850\\,$\\mu$m, thus lending strong evidence to the notion that the two populations are not identical.  We have argued that the average $\\fdr$ of SDOs is significantly lower than the bulk of the SMG population, suggesting that these sources are either very cool or lie at higher redshifts.  It would require extremely cold and unusual ($T_{\\rm d} \\simeq 10$\\,{\\sc k} and $\\beta \\simeq 1$) far-infrared/mm SEDs to explain the presence of SDOs at low redshifts ($z\\ls 2$). Nonetheless, examples of such sources do exist. We hint at a link between such cool, extended mm-selected sources and the near-infrared-selected population of $BzK$ galaxies, which are known to be gas-rich yet have star-formation rates an order of magnitude lower than those of SMGs.  If we adopt more realistic SED parameters ($T_{\\rm d} \\simeq 20$\\,{\\sc k} and $\\beta \\simeq 1.5$) for SDOs, it would imply an average redshift of $z \\gs 4$, i.e.\\ beyond the range probed by spectroscopic surveys so far. If this is the case, it is tempting to speculate that SDOs could be potential candidates for galaxies collapsing very early on in the Universe's history.   It should be stressed, however, that our analysis merely shows that SDOs are at at generally higher redshifts (or are cooler) than SMGs.  Extensive comparison with detailed models will be required in order to determine what range of redshift distributions are consistent with the available data.  This was beyond the scope of the present study.  This task will also be made much easier with additional information on these sources, particularly spectroscopic or optical/IR photometric redshifts.  However, this first requires finding an identification of these SDOs at other wavelengths.   Deep interferometric observations at (sub)mm wavelengths sensitive enough to probe the faint flux levels of SDOs, are likely to become one of the most promising ways of furthering our understanding of SDOs. Such observations would allow us to identify near-/mid-infrared and optical host galaxies with which to constrain their redshifts.  Due to the limited sensitivity of current (sub)mm interferometers, however, such studies have been limited to the brightest SMGs.  With the new generation of wide-band receivers, and with ALMA on the horizon, studies of larger samples of SDOs are becoming feasible. Furthermore, the wide, deep, multi-colour surveys planned with SCUBA-2 and {\\sl Herschel\\/} (and on a somewhat longer time-scale, with CCAT) may provide useful far-IR/submm photometric redshifts for huge samples of SDOs and SMGs (Hughes et al.\\ 2002; Aretxaga et al.\\ 2003). Armed with these redshift constraints, heterodyne receivers with extremely high bandwidth could be employed to search for redshifted CO and/or C\\,{\\sc ii} line emission. For a complete census of obscured star formation and AGN activity across cosmic time, we must determine the range of SED properties for (sub)mm-selected sources, along with their true redshift distribution. Our evidence that mm-selected galaxies may extend beyond $z=4$ suggests that this may be an efficient way to select and study massive galaxy formation at the earliest times."
    },
    "0806/0806.2336_arXiv.txt": {
        "abstract": "Trigonometric parallax astrometry and BVI photometry are presented for two late-type subdwarf candidates, LSR1425+71 (sdM8.0) and LSR1610$-$00 (sd?M6pec).  For the former we measure an absolute parallax of 13.37~$\\pm$~0.51 mas yielding M$_{\\rm V}$~=~ 15.25~$\\pm$~0.09.  The astrometry for LSR1610$-$00 shows that this object is an astrometric binary with a period of 1.66$\\,\\pm\\,$0.01 yr.  The photocentric orbit is derived from the data;  it has a moderate eccentricity (e$\\,\\approx\\,$0.44$\\,\\pm\\,$0.02) and a semi-major axis of 0.28$\\,\\pm\\,$0.01 AU based on our measured absolute parallax of 31.02$\\,\\pm\\,$0.26 mas. Our radial velocity measure of $-$108.1$\\,\\pm\\,$1.6 km s$^{-1}$ for LSR1610$-$00 at epoch 2006.179, when coupled with the observation of $-$95$\\,\\pm\\,$1 km s$^{-1}$ at epoch 2005.167 by Reiners \\&\\ Basri, indicates a systemic radial velocity of $-$101$\\,\\pm\\,$1 km s$^{-1}$ for the LSR1610$-$00AB pair.  The galactic velocity components for LSR1425+71 and LSR1610$-$00AB -- (U,$\\,$V,$\\,$W)$\\,=\\,$(84$\\,\\pm\\,6$,$\\;-$202$\\,\\pm\\,$13,$\\;$66$\\,\\pm\\,$14) km s$^{-1}$ and (U,$\\,$V,$\\,$W)$\\,=\\,$(36$\\,\\pm\\,2$,$\\;-$232$\\,\\pm\\,$2,$\\;-$61$\\,\\pm\\,$2) km s$^{-1}$, respectively.  For both stars, the velocities are characteristic of halo population kinematics.  However, modelling shows that both stars have orbits around the galaxy with high eccentricity that pass remarkably close to the galactic center. LSR1425+71 has a luminosity and colors consistent with its metal-poor subdwarf spectral classification, while LSR1610$-$00 has a luminosity and most colors indicative of being only mildly metal-poor, plus a uniquely red $B-V$ color. The companion to LSR1610$-$00 must be a low-mass, substellar brown dwarf. We speculate on the paradoxical nature of LSR1610$-$00 and possible sources of its peculiarities. ",
        "introduction": "Over the last decade, spectroscopic surveys targeting faint stars with large proper motions have identified many subdwarfs with M spectral types.  These surveys generally make use of the spectral classification scheme developed by Gizis (1997) which employs measures of the CaH and TiO bands in the 6300--7100~\\AA\\ spectral region to form the indices CaH1, CaH2, CaH3, and TiO as defined in Reid et al. (1995).  The locations in the three diagrams for the CaH indices versus the TiO index serve to separate the stars into three spectral metallicity classes designated MV for dwarfs with [M/H]$\\,\\approx\\,$0.0, sdM for the cool counterparts to the classical sdF--sdG subdwarfs with [M/H]$\\,\\approx\\,-$1.2$\\,\\pm\\,$0.3, and esdM for extreme subdwarfs with [M/H]$\\,\\approx\\,-$2.0$\\,\\pm\\,$0.5 (Gizis \\&\\ Reid 1997).  The numerical spectral subclass is then determined from calibrated relations with both the CaH2 and CaH3 indices.  At the time that this classification scheme was established, there was only one star with subclass type later than sdM4.5 known -- namely LHS377, which was assigned the type sdM7 ``by definition\" (Gizis \\&\\ Harvin 2006).  It has since become routine to consolidate the essential features of the Gizis classification scheme into a single diagram consisting of [CaH2~+~CaH3] versus TiO5 (cf., L\\'epine et al. 2004), retaining the revised divisions between the three metallicity subclasses proposed by Burgasser \\&\\ Kirkpatrick (2006).  Further revision and extension of the Gizis system have been proposed by L\\'epine et al. (2007). A fourth subclass designated ``usdM\" indicating ``ultrasubdwarfs\" was added to the three-subclass scheme to recognize the relatively few M stars which apparently have metallicities as low as or perhaps below [M/H]$\\,=\\,-$2.5.  A realignment of the subclass separators was put forward based on a new index, $\\zeta_{\\rm TiO/CaH}$, which measures the calibrated TiO5 spectral index as a function of the [CaH2~+~CaH3] index and is expressed as a ratio with the solar metallicity value.  The separators are chosen so that they run parallel to the locii in the [CaH2~+~CaH3,~TiO5] plane for the observed components of wide common proper motion binaries with different spectral subtypes.  A proposed list of subdwarf spectral classification standard stars were presented covering types K7.0 through M8.5 for each of the subclasses subdwarf, extreme subdwarf, and ultrasubdwarf.  As further suggested by Burgasser et al. (2005), one really needs to simultaneously utilize spectral classification criteria from a wider range of wavelengths and low-resolution data covering the range 0.7--2.5 $\\mu$m as a possible way to proceed in the future.  However, while the number of earlier type M subdwarfs recognized in the greater solar neighborhood has grown to several hundred due to these survey efforts, subdwarfs with types sdM6 and later remain relatively rare.  Even as recently as this past year, the number of subdwarfs with types sdM7 or later numbered only around 15 (Burgasser et al. 2007; L\\'epine et al. 2007). The discovery of one of these few -- LSR1425+7102, hereafter referred to as LSR1425+71 -- was announced several years ago by L\\'epine et al. (2003b) and assigned a classification of sdM8.0 on the Gizis (1997) system, making it the coolest sdM star identified at that time.  The similarity in overall continuum slope with LSR2000+3057 (M6.0V) over the $6000-9000\\,$\\AA\\ spectral region was noted.  Employing the Baraffe et al. (1997; hereafter BCAH97) ``NextGen'' model atmospheres for a metal-poor 0.09M$_{\\sun}$ star with [Fe/H]=$-$1.3 to estimate the absolute magnitude, they derived a ``conservative distance estimate of 65$\\,\\pm\\,$15 pc.''  This distance, when combined with their measured proper motion of 0.635 arcsec yr$^{-1}$ in a position angle of 254.7 deg and their measured radial velocity of $-$65$\\,\\pm\\,$20 km s$^{-1}$, yielded space velocity components of (U,V,W)\\footnote{ Throughout this paper, U is measured radially outward toward the galactic anticenter.} =($-$65$\\pm$22, $-$177$\\pm$35, +64$\\pm$27) km s$^{-1}$, consistent with halo membership.  Note that LSR1425+71 is also adopted as the classification standard star for spectral type sdM8.0 by L\\'epine et al. (2007).  Later in 2003, the same investigators announced the discovery of LSR1610$-$0040, hereafter referred to as LSR1610$-$00, and suggested that it might be an early-type sdL subdwarf (L\\'epine et al. 2003a; hereafter LRS03). At that time the only other sdL candidate was 2MASS$\\,0532+8246$ discovered by Burgasser et al. (2003) from spectroscopic follow-up on 2MASS photometric candidates.  L\\'epine et al. (2003a) noted that while their optical spectrum of LSR1610$-$00 (coverage$\\approx\\,6000-10,000\\,$\\AA; resolution$\\,\\approx\\,$7$\\,$\\AA) showed obvious similarities to their spectrum of LSR1425+71, LSR1610$-$00 possessed a distinctly steeper pseudocontinuum (implying a cooler temperature) and, more significantly, showed strong Rb~I~$\\,7800,7947\\,$\\AA$\\;$ lines which are more typically present in L-type dwarfs rather than M-dwarfs.  With its weak TiO bands and totally absent VO bands LSR1610$-$00 did not fit anywhere in the known sequences of dwarf M or subdwarf M stars.  Again using the BGAH97 models, they ``conservatively\" estimated the distance to LSR1610$-$00 to be 16$\\,\\pm\\,$4 pc.  Centroids of the Rb~I lines together with those from the K~I doublet (7665,~7699$\\,$\\AA) yielded a heliocentric radial velocity of $-$130$\\,\\pm\\,$20 km s$^{-1}$. Using their measured proper motion of 1.46 arcsec yr$^{-1}$ in a position angle 212.0 deg yields galactic space velocity components of (U,V,W)=($-$117$\\pm$18, $-$108$\\pm$24, $-$10$\\pm$19) km s$^{-1}$, suggesting that LSR1610$-$00 is most likely a member of the old disk population.\\footnote{Note that the U,~V,~W values presented in Table 1 of LRS03 are incorrect, even for the input values these authors employ.}  Combining new moderate resolution spectroscopy covering 0.8--4.1$\\,\\mu$m with the 0.6--1.0$\\,\\mu$m data from LRS03, Cushing \\&\\ Vacca (2006; hereafter CV06) made a comprehensive study of the peculiarities and contradictions presented in the spectrum of LSR1610$-$00.  Similar spectral data for three other sdM stars (LHS~3409, sdM4.5; LHS~1135, sdM6.5; LSR2036+5059, sdM7.5) and the M6.0V star GL406 were employed for comparison. GL406 is the well-studied high proper motion star Wolf~359 which is a primary MV spectral classification standard (Kirkpatrick et al. 1991) and a well established active flare star (CN Leo).  Based on its (U,V,W) galactic kinematic components, GL406 is formally a member of the old disk population (Leggett 1992).  However, Pavlenko et al. (2006) argue that its age is most likely younger than 0.4 Gyr -- that is, younger than the Hyades cluster (age$\\,\\sim\\,$0.6 Gyr) which is often taken as the upper age limit for the young disk population. CV06 concur with LRS03 that, in terms of spectral features, LSR1425+71 is a better match with LSR1610$-$00 than is GL406.  On the other hand, in terms of overall relative spectral energy distribution (SED) over the entire wavelength range 0.6--4.1~$\\mu$, LSR1610$-$00 and GL406 are remarkably similar, despite the fact that some features in the 7640--8300~\\AA\\ spectral region are noteably dissimilar, especially the stronger TiO bandhead and the much weaker Rb~I doublet in Gliese 406.  Given all of the contradictory evidence summarized by CV06 (see Table 3), these authors conclude that LSR1610$-$00 is most likely a mildly metal-poor, mid-M dwarf but further call the star ``schizophrenic\" to acknowledge its several spectral peculiarities. In particular we note the abnormally strong Al~I doublet at 1.313~$\\mu$m which suggests that aluminum is most likely overabundant compared with solar.  In a study concurrent with that of CV06, Reiners \\&\\ Basri (2006; hereafter RB06) obtained high-resolution (R$\\,\\approx\\,$31000) spectra in several windows of the 0.7--1.0$\\,\\mu$m region for LSR1610$-$00, GL406, and 2MASS~1439+1929 (hereafter, 2M1439+19).  The latter is a L1V star with an established trigonometrically determined luminosity (Dahn et al. 2002). Comparision of the strengths of TiO bands with heads at 7050, 8430, and 8870 \\AA\\ do not preclude the possibility that LSR1610$-$00 might be slightly metal deficient.  On the other hand, comparison of metal hydride bands (e.g., CaH at 6800~\\AA\\ and FeH at 9900~\\AA) indicate that LHS1610$-$00 can not be very much more metal-deficient than is GL406.  As pointed out earlier by both LRS03 and CV06, the primary discrepancies arise for a few atomic lines.  In particular, the Rb~I lines at 7947 and 7800 \\AA\\ agree much better with those in 2M1439+19, not only in strength but in shape.  In contrast, the Cs~I lines at 8520 and 8944 \\AA\\ are much weaker in LSR1610$-$00 than in 2M1439+19, with the redder one being almost undetectable.  Not only are the CS~I lines more similar to those in GL406, the one at 8520 \\AA\\ is significantly stronger than in GL406.  The strengths of the Ca~II triplet components at 8498 and 8542 \\AA\\ imply that LSR1610$-$00 must be at least as warm as GL406.  Comparing the spectra of LSR1610$-$00 and 2M1439+19 with PHOENIX atmospheric model predictions (cf., Allard et al. 2001), RB06 found that the observed Rb~I and Cs~I lines {\\it could} be simultaneously reproduced at a satisfactory level by a slightly metal-deficient model ([Fe/H]~$\\approx$~$-$1) with T$_{\\rm eff}$~=~2800~K, which alters the pseudo-continuum level appropriately for LSR1610$-$00.\\footnote{L\\'epine et al. (2004) announced the discovery of LSR0822+1700 as a second subdwarf M star that showed both RbI and CsI lines.  In this instance, the location in the [CaH2~+~CaH3,~TiO5] plane clearly indicates low metallicity.  LSR0822+1700 has been designated as a usdM7.5 spectral classification standard star by L\\'epine et al. (2007).  USNO initiated a parallax determination for this star in March 2008.} Based on this result RB06 suggested a spectral type of sd?M6 would be appropriate.  However, given the remaining spectal peculiarities not addressed by RB06, we prefer to adopt the spectral type designation of sd?M6pec for the present.  Finally, RB06 measured a high-precision heliocentric radial velocity for LSR1610$-$00 of $-$95$\\,\\pm\\,$1 km s$^{-1}$ by cross-correlating its spectum with that of GL406 in the spectral region around 8000$\\,$\\AA.  They discuss the discrepancy between their result and that of LRS03 ($-$130$\\,\\pm\\,$15 km s$^{-1}$) measured from much lower resolution spectroscopy and suggest that the LRS03 value might require a corrective offset of +20 km s$^{-1}$.  Calibration of the extension to the Gizis subdwarf-M spectral classification system to physical parameters such as total luminosities and effective temperatures, such as proposed by L\\'epine et al. (2007), will require reliable trigonometrically determined distances.  A representative sample of such M-type subdwarfs is currently being observed at the Flagstaff Station.  Both LSR1425+71 and LSR1610$-$00 were added to the USNO CCD trigonometric parallax program back in the spring of 2003 and reliable results are now available for these two stars.  Since LSR1425+71 and LSR1610$-$00 have been discussed and compared with one another repeatedly in the literature, we present our results for both stars together here. ",
        "conclusions": "Previously, both CV06 and RB06 have used extensive spectroscopic data to conclude that LSR1610$-$00 appears peculiar (even ``schizophrenic''), but it is most likely to be mildly metal-poor with some elements (particularly enhanced aluminum) indicating an unusual composition. The absolute magnitudes and colors found in this paper support the idea that it is near solar metallicity, although the uniquely red B$-$V color may be another indicator of unusual composition.  However, the revised distance in this paper and the space velocity and galactic orbit shown in Sec. 5 are paradoxical: the space velocity and orbit are not at all disk-like, and would normally indicate (although do not require) a lower metallicity than mildly metal-poor. The finding in this paper of a binary companion with a separation of less than 1 AU is a new surprise.  Using Occam's razor as a guiding principle, we might hope to relate the peculiar abundances to some interaction or mass transfer between the two binary components.  The enhancement of aluminum found by CV06 is probably an important clue. Halo red giants and more massive AGB stars may produce aluminum enhancements due to mixing burning products of the hydrogen shell to the surface. The aluminum is produced at the expense of oxygen (Kraft 1994).  Hence, one possible scenario that might be considered is that LSR1610$-$00B, the astrometric companion to LSR1610$-$00A, is a white dwarf.  During the RGB or AGB phase of the evolving white dwarf progenitor, an excess of aluminum was produced and mixed to the surface. With a 1.7 year period, depending on the stellar mass, the red giant might be big enough to transfer Al-enriched material to the small dwarf companion.  This would result in the enhancement observed today.  It appears, however, that this scenario has fatal flaws and must be discarded.  The fundamental difficulty is the low mass of both components required by the binary possibilities in Table 4. White dwarf masses below 0.2 M$_{\\sun}$ have been found as companions to millisecond pulsars (van Kerkwijk et al. 1996) and in the Sloan Digital Sky Survey (Eisenstein et al. 2006), although these are rare. What appears to be required to form a white dwarf of extremely low mass is that the companion has a separation such that Roche lobe overflow occurs when the primary star is trying to leave the main sequence with only a small helium core.  If this happens, the star does not become much bigger than its main sequence size before its evolution and growth of the core mass are truncated.  As such, the 1.7 year period and the separation of the two stars of 0.4-0.8 AU are too long and too large for this to have occurred.  A second problem is that, even with a maximum cooling age on the order of that of the halo ($\\sim\\,$12 Gyr), most low-mass white dwarfs will still have a luminosity such that they will be visible in the spectrophotometry of LSR1610$-$00. The most stringent constraint is the observed M$_{\\rm B}$~=~22.4 and B$-V$~=~3.3, requiring M$_{\\rm B}\\, >\\,$25 for a possible white dwarf companion. For example, white dwarf models from Bergeron et al. (1995) with hydrogen atmospheres and M = 0.15 M$_{\\sun}$ reach M$_{\\rm B}$~=~19.1 at a cooling age of 12 Gyr.  It may be possible that a low-mass white dwarf with a helium atmosphere would cool quickly enough that it could still be present and not detected. Otherwise, the companion to LSR1610$-$00A must be an unevolved, lower-mass star or substellar brown dwarf.  Therefore, the B component to LSR1610$-$00A appears not to be the source of abundance anomalies in LSR1610$-$00A.  Nevertheless, peculiar abundances do point toward the accretion of mass onto LSR1610$-$00.  Most notable is the enhancement of Al found by CV06, a result suggesting accretion from a massive AGB star that has undergone hot-bottom-burning.  Some stars in globular clusters (both giants and subgiants) have enhanced Al and Na, and depleted O, indicative of external pollution by AGB stars (e.g., Sneden et al. 2004). The mechanism for Al and Na production is believed to be proton capture by Mg and Ne nuclei at high temperatures at the base of the hydrogen-burning shell.  The most extreme abundance anomalies of this type are seen in the extreme metal-poor star HE1327-2326 (Frebel et al. 2005) where C, N, Na, Mg, and Al are seen to be strongly enhanced compared to Fe, and Ca and Ti are mildly enhanced.  We speculate that accretion of $<$0.05 M$_{\\sun}$ of material with such strong enhancements onto LSR1610$-$00, which initially might have been about 0.05 M$_{\\sun}$ and [Fe/H] $\\sim -$2, followed by complete mixing in the fully convective star, can lead to the star we see now.  The accretion enhances all elements, making LSR1610$-$00 mildly metal-poor, it enhances C and the C/O ratio, almost (but not quite) turning LSR1610$-$00 into a carbon star and leaving most oxygen tied up in CO, it enhances Ti, but leaves TiO slightly depleted, and it greatly enhances Al.  The $B$ bandpass includes several bands of AlH, as well as MgH, SiH, and CaI:  these bands are likely to have enhanced strength in LSR1610$-$00 compared to most dM6 and sdM stars because of a combination of being mildly metal-poor (strengthening the hydrides and atomic lines) and abundance anomalies. We speculate that these (not yet observed) features are the cause of the unique red B$-$V color.  The AGB star proposed to be the source of mass accretion probably would have been a more distant companion in a triple system.  It would now be a cool white dwarf, and must have separated from the present LSR1610$-$00 binary when it lost most of its mass.  (If it were still bound with LSR1610$-$00, it would now be detectable as a companion, unless it was initially massive enough to become a neutron star.)  Alternatively, LSR1610$-$00 and the AGB star originally could have been part of a globular cluster when the accretion occurred, with LSR1610$-$00 later being lost from the cluster or the cluster dissipating entirely."
    },
    "0806/0806.2046_arXiv.txt": {
        "abstract": "We show that with GLAST there will be the possibility to detect, within the UHECR skimming the Earth atmosphere, the showers generated by very high energy upward and horizontal Tau. The effective area, thanks to the large area covered by the showers at 550 Km, is  less than that of AUGER, but its efficiency is comparable because the lower detection threshold and the consequent event rate may lead to a few  EeV and-or few Glashow resonant signals within a decade. ",
        "introduction": "Ultra high energy neutrinos UHE  $ \\nu_{\\tau}$, $ \\bar\\nu_{\\tau} $ and $ \\bar{\\nu}_e $ at EeV's up to GZK energies ($\\ge 10^{19}$ eV) can hit the earth crust at the horizon leading to UHE $\\tau$ which may decay in flight at high altitude. The consequent UHE  air showers  might be observable by next generation gamma-ray space missions like GLAST.  Here we show the expected  fluence and time signature considering two different complementary signals: the upward $\\tau$ air shower (UpTau) near the vertical at PeV energies and the horizontal $\\tau$ air shower ( HorTau ) at 1.4 10$^{19}$ eV.    \\subsection{Upward $\\tau$ air shower}  Assuming a given altitude ($h_1 \\sim 575 Km$) for the circular orbit of the satellite, the distance between the detector and the edge of the earth crust can be written as \\begin{equation} {d_{1U}=(R+h_{1})\\sin(\\theta_{1U})- {\\sqrt{(R+{h_{1}})^2 \\cdot {\\sin}^{2}(\\theta_{1U})-[(R_\\oplus  + h_1)^2 - R_\\oplus^2]}}} \\end{equation} where $\\theta_{1U}\\sim 70^o$ is the angle of the shower from the horizontal. In first approximation $${d_{1U}  \\sim h_1 / \\sin \\theta_{1U}  \\sim 612 Km }$$    A pictorial view of the detection method is shown in figure~\\ref{glastscheme1}.   \\begin{figure} \\vspace{0cm} \\epsfig{file=taufig2_glast.eps,width=1\\textwidth,angle=0,clip=} \\caption{\\label{glastscheme1} \\it A schematic picture (not in scale) of the possible detection method }\\label{glastscheme1} \\end{figure}  Now we can calculate the area $ A_U $ of the corresponding  front of the  upward $\\tau$ showers that is given by :  $$ A_U = {\\pi \\over 4} \\Delta \\theta_{sh}^2 d_{1U}^2 \\simeq  90 Km^2 $$   where $\\Delta \\theta_{sh} \\sim 1^o$ is the typical opening angle for showers.   The lateral density profile is, of course, more dense near the inner part, and here we assume that $\\sim 90 \\%$ of the gammas are contained in a narrow angle of $1/4 $ of degree, leading to a reduces  area:  $$A_{Ur}=5.62  Km^2$$  This areas allow us to calculate  the secondary gamma-ray flux.  We considered     $\\nu_\\tau$'s with a  primary energy of the order of  $\\sim 4  \\cdot  10^{15}$ eV  because at greater energies they are suppressed by the earth opacity, while at lower energies the cross section and the $\\tau$ propagation length are smaller \\cite{tau}. Indeed the probability $P(\\theta,\\, E_{\\nu})$ of escaping from the earth is approximately \\begin{equation} P(\\theta_{1U},\\, E_{\\nu}) \\simeq e^{\\frac{-2R_\\oplus \\sin \\theta_{1U}}{R_{\\nu_{\\tau}}(E_{\\nu})}} (1 - e^{- \\frac{R_{\\tau}(E_{\\tau})}{R_{\\nu_{\\tau}}(E_{\\nu})}}) \\, . \\end{equation}  where $\\theta_{1U}$ is nearly the complementary angle of the direction of the upcoming $\\tau$ angle with the zenith, $R_{\\tau}$ is the interaction length of the $\\tau$ and $R_{\\nu_{\\tau}}$ is the $\\nu$ interaction length.    \\begin{figure}[htbp] \\epsfig{file=Fig014IntTau.eps ,width=1\\textwidth,angle=0,clip=} \\caption{Lepton $\\tau$ (and $\\mu$) Interaction Lengths for different matter densities: $R_{\\tau_{o}} =  c\\cdot {\\tau_{\\tau}}\\cdot {\\gamma_{\\tau} } $ is the free $\\tau$ length,$R_{\\tau_{New}}$ is the New Physics TeV Gravity interaction range at corresponding densities, $R_{\\tau_{Nucl}\\cdot{\\rho}}$ ,\\cite{tau}, see also \\cite{Dutta et al.2001}, is the combined $\\tau$ Ranges keeping care of all known interactions and lifetime and mainly the photo-nuclear interaction. There are two slightly different split curves (for each density) by two comparable approximations in the interaction laws. Note also the neutrino interaction lenghts above lines $R_{Weak{\\rho}}= L_{\\nu}$ due to the electro-weak interactions at corresponding densities (see also \\cite{Gandhi et al 1998})  \\cite{tau}.} \\label{fig4} \\end{figure}    The $\\tau$ energy is typically  $\\sim 20 \\% $ less then the $\\nu_\\tau$ one's.  The number of gamma $ N_{\\gamma s}$ with energies around 100 MeV in the showers is in first approximation $$ N_{\\gamma s} \\sim { E_\\tau  \\over E_c} \\sim  4 \\cdot 10^6   $$  for $ E_c =100MeV$ with the assumption of the energy equi-partition between $\\gamma$, electron pairs ($\\sim$ 33\\% each component), as well as taking in account a partial ($\\sim$ 33\\%) opacity of the atmosphere for the $\\tau$ shower.   The number of photons per unit reduced area at the altitude of GLAST is then :  $$\\Phi_{\\gamma r} = {N_{\\gamma s} \\over A_{Ur}} =  7.14 \\cdot 10^{-5} cm^{-2}$$ in the inner$\\Delta \\theta = {1/4} ^o$ core and $$\\Phi_\\gamma =  4.4 \\cdot 10^{-6} cm^{-2}$$ in the wider $\\Delta \\theta = 1^o$ cone shower.  The characteristic time structure of the shower is $ t_s \\sim   L_s /  c  \\ge 10^{-4} ~s   $  where $L_s $ is the shower attenuation length at altitude  $\\sim 23 Km$, where upward $\\tau$ take place. (see ref.\\cite{tau2}).  Assuming the lateral GLAST detector area, $A=  1.3 \\cdot 10^4 cm^2$, an efficiency  $\\eta = 0.5 $, the total effective area  is $ A_{eff} = A \\cdot \\eta \\cdot cos\\theta  = 2.3 \\cdot 10^3 cm^2$ the number of photons for each event is, respectively for narrow and large view angle :  $$  N_{\\gamma r} (E \\sim 100 MeV) =  \\Phi_\\gamma \\cdot A_{eff} \\sim 0.16  $$ $$  N_\\gamma (E \\sim 100 MeV)   \\sim 10^{-2}  $$  So we conclude that GLAST can measure upward $\\tau$ only in coincidence with the GRB monitor approximately one over 6 upward $\\tau$ showers. Therefore the high energy $\\gamma$ detection alone is not an effective way to discriminate upward Tau Air-Showers (UpTaus) by GLAST.       \\subsection{Horizontal $\\tau$ air shower}  We can now use the above procedure to calculate the rate of events for the Horizontal $\\tau$ air shower. The distance between the detector and the edge of the earth crust is in this case  $d_{hH}=( 2R_\\oplus h_1)^{1/2} \\cdot  (1+ {h_1 \\over 2 R_\\oplus})^{1/2} \\sim 2768$ Km for the same altitude of 575 Km and where  the angle of the shower from the horizontal is $\\theta_{hH} = \\arctan(( {2h_1 \\over R_\\oplus} )^{1/2} \\cdot  (1+ {h_1 \\over 2 R_\\oplus})^{1/2} )\\sim 23.5^o$. However the Tau decay in flight and the HorTau appearence takes place at great distance ($ \\simeq 600$ km) from the Earth and the HorTau Shower has a characteristic distance of ($\\simeq 200$ km) making the real distance from the Shower front to the satellite reduced to  $d_{hHorTau} \\sim 2500$ km.   Now we can calculate the area of the corresponding  front of the showers given by : $$ A_H = {\\pi \\over 4} \\Delta \\theta_{sh}^2 d_h^2 \\sim  1510  Km^2 $$ This area is comparable with future AUGER experiment area. In analogy to previous UpTaus scenario we also consider inner Shower cone of a nominal beam angle $1/4$ of degree obtaining a reduced area $$ A_{H r} = {\\pi \\over 4} \\Delta \\theta_{sh}^2 d_h^2 \\sim 94.4  km^2$$.   This areas allow us to calculate  the secondary gamma-ray flux.   The optimal observable primary neutrino energy is  1.1 $ 10 ^{19}$ eV because of the earth crust slant  depth  combined with the horizontal atmospheric opacity \\cite{tau}.   The number of gamma $ N_\\gamma$ with energies around 100 MeV in the showers is in first approximation $$ N_\\gamma \\sim { E_\\tau  \\over E_c} \\sim  3.3  \\cdot 10^{10} $$ with the same assumption of the energy equi-partition between $\\gamma$, electron pairs but without the  opacity of the atmosphere for the $\\tau$ shower because in this case we are at the maximum of the shower with nearly no atmospheric suppression.   The number of photons per unit area (or reduced area) at the altitude of GLAST is then :  $$\\Phi_\\gamma = {N_{\\gamma s} \\over A_{Ur}} =  2.18 \\cdot 10^{-3}  cm^{-2}$$ $$\\Phi_{\\gamma r} = {N_{\\gamma s} \\over A_{Ur}} = 3.5 \\cdot 10^{-2}  cm^{-2}$$  The characteristic time structure of the shower is $ t_s \\sim L_s /  c  \\ge 10^{-3} ~s   $  where $L_s \\simeq 200 $ km,  is the shower attenuation length at high altitude $\\sim 23 $km, where air is much diluted (see ref.\\cite{tau2}).  Assuming  as before  the lateral area of the detector $A= 1.3 \\cdot 10^4 cm^2$, an efficiency  $\\eta = 0.5 $, the total effective area  is $ A_{eff} = A \\cdot \\eta \\cdot \\cos\\theta = 0.6 \\cdot 10^4 cm^2$ the number of photons for each event is for $\\Delta \\theta_{Sh}  = 1^o, {1/4}^o$:   $$  N_\\gamma (E \\sim 100 MeV)  \\sim  13.1  $$ $$  N_{\\gamma r} (E \\sim 100 MeV)  \\sim  210  $$  \\subsection {HorTau Event rate in GLAST} The number of events may be estimated by scaling  the EUSO experiment event rate at the horizons, keeping care of the different beaming angle and of the different horizontal area and duty cycle life-time $\\eta_{EUSO} \\simeq 0.1$ respect the  GLAST one $\\eta_{GLAST} \\simeq 1$ , for a nominal three years of recording. These event rate are scaled assuming a minimal, guaranteed GZK ( Greisen, Zatsepin, Kuzmin) neutrino fluence $\\Phi_{\\nu GZK} \\simeq \\Phi_{UHECR} \\simeq  3 \\cdot 10 ^{-18} cm ^{-2} s^{-1} sr^{-1}$ produced by observed Ultra High Cosmic Rays, UHECR, during their photopion scattering on Cosmic Big Bang Radiation within the GZK cut-off volumes: $$ {N_{GLAST}} = \\frac{A_{GLAST}}{A_{EUSO}}\\frac{1}{360^o} \\cdot \\frac{1}{\\eta_{EUSO}} {N_{EUSO}}\\frac{1}{2} \\sim 0.398 {N_{EUSO} \\sim {15}\\leftrightarrow{30}}  $$  The consequent reduced area (narrower beamed) event number is:  $$ {N_{GLAST r}} = \\frac{A_{GLAST}}{A_{EUSO}}\\frac{1}{5760^o} \\cdot \\frac{1}{\\eta_{EUSO}} {N_{EUSO}}\\frac{1}{2} \\sim 0.0248 {N_{EUSO} \\sim 1\\leftrightarrow 2} $$  \\subsection{HorTaus versus Other High Altitude Showers} Among these Upward-Horizontal Showers by $\\tau$ we must consider the competitive signals of more common and known UHECR showers at horizons: Horizontal High Altitude Shower Hias (\\cite{tau2} ) are observed by satellites above the horizons ($\\theta \\geq {0.8}^o$) and they behave as a  background signal respect to HorTau below the Horizons ($\\theta \\leq {0.05}^o$). Indeed their event number in three years  (at the same GZK energies $10^{19}$ eV, and flux $ \\Phi_{UHECRs}$ as  in previous section :$ \\Phi_{UHECRs} \\simeq 3 \\cdot 10^{-18}$ eV) is $$ {N_{GLAST}} \\sim {247}  $$ The consequence of this expected signal above the horizons is the necessary presence of a background Ultra High Cosmic Rays at a rate comparable to present AGASA and HIRES records. The very natural advantage is the general calibration of this UHECR physics on ground with this high quota Showering in Space. The drawback is the need of a clear angle discriminator between HorTaus and Hias. Because at the distances we are dealing the split angle is nearly one degree we may expect that a dozen or more gamma events will be enough to estimate the arrival direction within a needed accuracy (a few tenth of degree). In summary the Glast thresholds are described in the included figure below. \\begin{figure}[htbp] \\epsfig{file=GLASTHortasCR-All-26Giugno.eps,width=1\\textwidth,angle=0,clip=} \\caption{GLAST thresholds for Horizontal Tau air-shower shower, HORTAUs (or Earth Skimming Showers) over all other $\\gamma$, $\\nu$ and Cosmic Rays (C.R.) fluence and bounds. The fluence threshold for Glast has been estimated for a three year experiments lifetime. Competitive experiment are also shown as well as the Z-Shower expected spectra in two different light neutrino mass values ($m_{\\nu} = 0.04, 0.4$ eV).  \\cite{tau}, \\cite{tau2},\\cite{Kalashev:2002kx}.} \\end{figure}     \\subsection{GLAST}  The Gamma-ray Large Area Space Telescope (GLAST)\\cite{glast},  has been selected by NASA  as a  mission involving an international collaboration of  particle physics and astrophysics communities from the United States, Italy, Japan, France and Germany  for a launch in the first half of 2006. The main scientific objects are the study of all gamma ray sources such as blazars, gamma-ray bursts,  supernova remnants, pulsars, diffuse radiation, and unidentified high-energy sources. Many years of refinement has led to the configuration of the apparatus shown (see figure~\\ref{glastscheme2}), where one can see the  4x4 array of identical towers each formed by: $\\bullet $   Si-strip Tracker Detectors and converters arranged in 18 XY tracking planes for the measurement of the photon direction. $\\bullet $ Segmented array of CsI(Tl) crystals for the measurement the photon energy. $\\bullet $ Segmented Anticoincidence  Detector (ACD). The main characteristics are an  energy range between    20 MeV and 300 GeV, a field of view of $\\sim$   3  sr,  an energy resolution    of $\\sim$ 5\\% at 1 GeV, a point source sensitivity of  2x10$^{-9}$ (ph~cm$^{-2}$~s$^{-1}$)     at  0.1 GeV, an event deadtime    of 20 $\\mu s$ and a peak effective area  of  10000 cm$^2$, for a required power    of   600 W and a payload weight of  3000 Kg.  The list of the people and  the Institution involved in the collaboration together with the on-line status of the project is available at {\\sl http://www-glast.stanford.edu}.  The important number for our estimate is the lateral area of the tracker for each of the four sides that is  $A= 60 cm \\cdot 170 cm = 1.02 \\cdot 10^4 cm^2$.  The projected total area is $4 \\cdot  A* \\cos(\\theta_{1U}) = 1.4 \\cdot 10^4 $ where $\\theta_{1U} =70^o$ is the angle between the arrival $\\tau$ shower and the horizon constrained by the geometry of the servicing modules that do not allowed to see upward showers (see figure~\\ref{glastscheme2})   \\begin{figure} \\vspace{0cm} \\epsfig{file=taufig1_glast.eps,width=1\\textwidth,angle=0,clip=} \\caption{\\label{glastscheme2} \\it Scheme of the lateral view of  GLAST with the arrival directions of horizontal and upward $\\nu\\tau$ shower .} \\end{figure} ",
        "conclusions": "The gamma-ray space experiment  GLAST is just in orbit. Its clear detection of Cosmic rays secondaries, mostly single gamma and electron pairs  as well as muons must take place at a high rate (thousands of events a year). Most muons pairs will hit the detector at 400 GeV energies. More rare  bundle of X-$\\gamma$ and $0.4 $ TeV $\\mu$ as well as UHE (tens GeV) neutrons (with and without gamma-X traces) might  be also observable soon. PeVs-EeVs cosmic rays air-showering at the terrestrial atmosphere edge must occur at daily-weekly rate in GLAST. The first neutron-gamma-electrons and or muon-gamma-electrons at  associated bundles must flash soon opening a new road to UHECR astrophysics. Moreover with a high angular resolution (below $0.5^o$)it might be even possible in a future to reveal first EeV persistent gamma source as well as rarest PeVs-EeVs upgoing tau. This signals are to be distinguished from background noises whose single event or whose rare pair structure in different from rarest (tens) bundles of X-gamma-muons or X-Gamma neutron burst at 0.1 millisecond time structure. This upgoing airshowers will be the most exciting   signal of the long waited UHE Neutrino Astronomy. Similar results , but an order of magnitude below, maybe applied to AGILE detector."
    },
    "0806/0806.4040_arXiv.txt": {
        "abstract": "A recent measurement of $^4$He photodisintegration reactions, $^4$He($\\gamma$,$p$)$^3$H and $^4$He($\\gamma$,$n$)$^3$He with laser-Compton photons shows smaller cross sections than those estimated by other previous experiments at $E_\\gamma \\lesssim 30$~MeV. We study big-bang nucleosynthesis with the radiative particle decay using the new photodisintegration cross sections of $^4$He as well as previous data.  The sensitivity of the yields of all light elements D, T, $^3$He, $^4$He, $^6$Li, $^7$Li and $^7$Be to the cross sections is investigated.  The change of the cross sections has an influence on the non-thermal yields of D, $^3$He and $^4$He.  On the other hand, the non-thermal $^6$Li production is not sensitive to the change of the cross sections at this low energy, since the non-thermal secondary synthesis of $^6$Li needs energetic photons of $E_\\gamma \\gtrsim 50$~MeV. The non-thermal nucleosynthesis triggered by the radiative particle decay is one of candidates of the production mechanism of $^6$Li observed in metal-poor halo stars (MPHSs). In the parameter region of the radiative particle lifetime and the emitted photon energy which satisfies the $^6$Li production above the abundance level observed in MPHSs, the change of the photodisintegration cross sections at $E_\\gamma \\lesssim 30$~MeV as measured in the recent experiment leads to $\\sim 10$~\\% reduction of resulting $^3$He abundance, whereas the $^6$Li abundance does not change for this change of the cross sections of $^4$He($\\gamma$,$p$)$^3$H and $^4$He($\\gamma$,$n$)$^3$He.  The $^6$Li abundance, however, could show a sizable change and therefore the future precise measurement of the cross sections at high energy $E_\\gamma \\gtrsim$ 50~MeV is highly required. ",
        "introduction": "In standard cosmology, the universe is thought to have experienced big-bang nucleosynthesis (BBN) at a very early stage.  The light nuclides D, T+$^3$He, $^4$He and $^7$Li+$^7$Be are produced in the standard BBN (SBBN) at observable levels, while this model does not make appreciable quantities of $^6$Li.  The Wilkinson Microwave Anisotropy Probe (WMAP) satellite has measured the temperature fluctuations of the cosmic microwave background (CMB) radiation, and parameters characterizing the standard big bang cosmology have been deduced~\\cite{Spergel:2003cb,Spergel:2006hy} from these data. For the baryon-to-photon ratio $\\eta_{\\rm CMB}$ deduced from fits to the CMB, the BBN model predicts abundances of the light elements except for $^6$Li  and $^7$Li which are more-or-less consistent with those inferred from astronomical observations.  Spectroscopic lithium abundances have been detected in the atmospheres of metal poor stars.  Nearly constant abundances of $^6$Li and $^7$Li in metal-poor Population II (Pop~II) stars have been inferred. There is about a factor of three under-abundance of $^7$Li in metal-poor halo stars (MPHSs) with respect to the SBBN prediction when using the baryon-to-photon ratio $\\eta_{\\rm CMB}$.  This is called the $^7$Li problem~\\cite{Ryan:1999vr,Melendez:2004ni,Asplund:2005yt}.  In addition, spectroscopic measurements obtained with high resolution indicate that MPHSs have a very large abundance of $^6$Li, i.e.~at a level of about 3 orders of magnitude larger than the SBBN prediction of the $^6$Li abundance, which is called the $^6$Li problem~\\cite{Asplund:2005yt,ino05}.  Cayrel et al.~\\cite{Cayrel:2007te} studied line asymmetries to be generated by convective Doppler shifts in stellar atmospheres, and found that the convective asymmetry might mimic the presence of $^6$Li and an error of $^6$Li/$^7$Li amounts to a few percent that is roughly comparable to the values estimated from MPHSs.  The $^6$Li problem, therefore, may not exist in fact, since the convective asymmetry could give a possible solution to the $^6$Li problem within the framework of SBBN.  The possibility of $^6$Li production in non-standard BBN  triggered by the decay of unstable relic neutral massive particles $X$ has been studied~\\cite{Jedamzik:2006xz, Jedamzik:1999di,Kawasaki:2000qr,Cyburt:2002uv,Kusakabe:2006hc,Jedamzik:2004er,Jedamzik:2004ip, Kawasaki:2004qu,Cumberbatch:2007me}. Several critical constraints on the properties of $X$ particles were derived from the studies of radiative decay~\\cite{Jedamzik:1999di, Kawasaki:2000qr,Cyburt:2002uv,Kusakabe:2006hc}, hadronic decay or annihilation~\\cite{Jedamzik:2004er, Jedamzik:2004ip,Kawasaki:2004qu,Cumberbatch:2007me} of $X$ particles along with the BBN constraints on the light elements. These particle decay induces electromagnetic and/or hadronic showers triggering the destruction of preexisting nuclei and the production of different nuclear species.  A recent detailed study~\\cite{Kusakabe:2006hc} of the radiative decay and its influence on the $^6$Li production has found a parameter region of lifetime $\\tau_X\\sim 10^8-10^{12}$~s and abundance parameter $\\zeta_X\\sim 10^{-13}-10^{-12}$~GeV where the non-thermal nucleosynthesis of $^6$Li can explain the observed abundance level in MPHSs. This parameter region satisfies the two observational constraints on the CMB energy spectrum and the primordial light element abundances. Three important characteristics were found for the interesting parameter region. First, $^3$He and $t$ are the seeds for $^6$Li in the processes $^4$He($^3$He,$p$)$^6$Li and $^4$He($t$,$n$)$^6$Li. Second, the excess of $^6$Li abundance is therefore regulated by the amounts of $^3$He and $t$ which are produced by the non-thermal photodisintegration of $^4$He, i.e. $^4$He($\\gamma$,$p$)$^3$H and $^4$He($\\gamma$,$n$)$^3$He. Hence, the radiative decay model which results in $^6$Li-production above the MPHS abundance level is also reflected by an enhancement of the $^3$He abundance with respect to the SBBN value. Third, the radiative decay does not resolve the $^7$Li problem~\\cite{footnote}.  It is therefore concluded that other mechanisms such as the stellar depletion of the lithium isotopes in the atmosphere of MPHSs~\\cite{Richard:2004pj,Lambert:2004kn} or new burst of late-time BBN on the exotic $X$-bound nuclei in the case of negatively-charged leptonic particles $X^-$~\\cite{Pospelov:2006sc,Hamaguchi:2007mp,Cyburt:2006uv,Kohri:2006cn,Bird:2007ge,Kusakabe:2007fu,Kusakabe:2007fv,Kawasaki:2007xb,Kawasaki:2008qe,Jittoh:2007fr,Jittoh:2008eq,Jedamzik:2007cp,Jedamzik:2007qk,Pospelov:2007js,Pospelov:2008ta} must operate to lower the $^7$Li abundance.  A recent measurement of $^4$He photodisintegration reactions, $^4$He($\\gamma$,$p$)$^3$H and $^4$He($\\gamma$,$n$)$^3$He with laser-Compton photons~\\cite{Shima:2005ix} shows much smaller cross sections than those estimated from the other previous experiments~\\cite{Nakayama2007} and those summarized in Ref.~\\cite{Cyburt:2002uv} at the photon energies 20~MeV $\\lesssim E_{\\gamma} \\lesssim$ 30~MeV. If these non-thermal photon energies dominate the destruction of $^4$He, the production of $^3$He and $t$ and also the subsequent production of $^6$Li via $^4$He($^3$He,$p$)$^6$Li and $^4$He($t$,$n$)$^6$Li, this would change the parameter region of $\\tau_X$ and $\\zeta_X$ of massive relic particles $X$ so that the resultant non-thermal nucleosynthesis of $^6$Li can explain the abundance level observed in MPHSs. The first purpose of this article is to study the sensitivity of non-thermal BBN of all light elements D, T, $^3$He, $^4$He, $^6$Li, $^7$Li and $^7$Be to the photodisintegration cross sections of $^4$He. The second purpose is to infer the uncertainties of the two parameters $\\tau_X$ and $\\zeta_X$ of massive relic particles $X$ which would arise from the uncertainties of the measured reaction cross sections.  In Sec.~\\ref{sec2} we present a result of a new measurement of $^4$He photodisintegration cross sections.  In Sec.~\\ref{sec3} we briefly explain the model of non-thermal nucleosynthesis and the calculated result of the effect of the considered change of the photodisintegration cross sections.  In Sec.~\\ref{sec4} we summarize our conclusion and offer an outlook for measurements of $^4$He photodisintegration. ",
        "conclusions": "\\label{sec4}  A recent measurement of $^4$He photodisintegration reactions, $^4$He($\\gamma$,$p$)$^3$H and $^4$He($\\gamma$,$n$)$^3$He with laser-Compton photons shows lower cross sections at low energies than those estimated by other previous experiments. We studied the sensitivity of non-thermal BBN of all light elements D, T, $^3$He, $^4$He, $^6$Li, $^7$Li and $^7$Be to the photodisintegration cross section of $^4$He.  The change of cross sections of $^4$He photodisintegration has an influence on the non-thermal yields of light elements, D, $^3$He and $^4$He, which are related to the photodisintegration cross sections at low energy ($\\sim 30$~MeV). The upper limit of allowed regions of $X$-abundance parameter $\\zeta_X$ for these light nuclei shifts upward by $\\sim 300- 30$~\\% for $\\tau_X=10^6 -10^{10}$~s for this change of the cross sections. This arises from the upshift of $^3$He abundance contour (See Fig.~\\ref{fig3}). On the other hand, the non-thermal $^6$Li production is not very sensitive to the change of cross sections at low energy, since the non-thermal secondary synthesis of $^6$Li needs energetic photons of $E_\\gamma \\gtrsim 50$~MeV.  The non-thermal nucleosynthesis triggered by the radiative particle decay is one of candidates of the production mechanism of $^6$Li observed in MPHSs. In the interesting parameter region of $10^8$~s$\\lesssim\\tau_X \\lesssim 10^{12}$~s and $5\\times 10^{-14}$~GeV $\\lesssim \\zeta_X \\lesssim 5\\times 10^{-13}$~GeV which satisfies the $^6$Li production above the abundance level observed in MPHSs, the lowering of the photodisintegration cross sections at low energy $E_\\gamma \\lesssim 30$~MeV as measured in the recent experiment using laser-Compton photons leads to $\\sim10$~\\% reduction of resulting $^3$He abundance, whereas the $^6$Li abundance does not change for the change of the cross sections of $^4$He($\\gamma$,$p$)$^3$H and $^4$He($\\gamma$,$n$)$^3$He.  Let us briefly discuss other impacts of such a precise cross section measurement. Clarifying the effects of photodisintegrations of $^4$He will affect more strongly the $\\nu$-process in core-collapse supernova (SN) explosions through the neutrino-nucleus interactions specifically of $\\nu$+$^4$He. The weak transition rates for $^4$He($\\nu,\\nu$'), $^4$He($\\nu_e$,$e^-$), and $^4$He($\\bar{\\nu}_e$,$e^+$) are determined similarly to the giant electric dipole resonance observed in the photodisintegrations with the help of theoretical calculation~\\cite{Suzuki:2006qd}. In fact, several experiments of measuring the $^4$He photodisintegration cross sections~\\cite{Shima:2005ix} were carried out for this purpose. The precise knowledge of the $^4$He($\\nu$,$\\nu$'$p$), ($\\nu$,$\\nu$'$n$), ($\\nu_e$,$e^- p$), and ($\\bar{\\nu}_e$,$e^+ n$) cross sections is required to determine the unknown parameters for neutrino oscillations through the MSW effect on the $^7$Li and $^{11}$B production triggered by the $\\nu$+$^4$He reactions~\\cite{Yoshida:2006qz,Yoshida:2006sk}. The energy range $E_\\nu = 10 - 25$ MeV is very important for the $\\nu$-process nucleosynthesis in SNe. The mean neutrino energy of SN neutrinos is presumed to be about 10 - 25 MeV in numerical simulations of the neutrino transfer in core-collapse SNe, and the threshold energies for all neutrino-induced spallation reactions of $^4$He are $\\sim$20 MeV. Therefore, the difference between the newly measured~\\cite{Shima:2005ix} and previous $^4$He photodisintegration cross sections at 20~MeV $\\lesssim E_{\\gamma} \\lesssim$ 30~MeV could be critical. As a result, the absolute yields of $^7$Li and $^{11}$B produced in the $\\nu$-process in core-collapse SNe would be different from one another, depending on the assumed $\\nu$-process reaction rates as demonstrated theoretically~\\cite{Suzuki:2006qd,yoshida08} although the ratio of $^7$Li/$^{11}$B does not change largely.  Another recent focus of photodisintegration of $^4$He is on the mechanism of the core-collapse SNe. Most SN simulations still do not succeed in the SN explosion in spite of detailed numerical studies of the neutrino transfer calculations inside the core. Haxton~\\cite{Haxton:1988in} proposed that the neutrino-induced excitations of $^4$He and heavier nuclei could deposit extra-energy to the ejected materials and revive the shock wave, which motivated a recent theoretical study on the role of $^4$He spallation reactions in the core-collapse SNe~\\cite{Ohnishi:2005cv}. His theoretical suggestion also motivated recent experimental studies of photodisintegrations of $^4$He~\\cite{Shima:2005ix} in order to estimate the neutrino-induced reaction cross sections for $^4$He($\\nu,\\nu$'), $^4$He($\\nu_e$,$e^-$), and $^4$He($\\bar{\\nu}_e$,$e^+$).  As such, it is important and even critical to study the $^4$He($\\gamma$,$p$) and $^4$He($\\gamma$,$n$) reactions precisely for the discussions of the problem of SN-neutrino oscillation and SN explosion as well as the cosmological discussion concerning the BBN with a radiative decay of long-lived relic particles."
    },
    "0806/0806.1386_arXiv.txt": {
        "abstract": "Massive star clusters observed in galaxy mergers are often suggested to be progenitors of globular clusters. To study this hypothesis, we performed the highest resolution simulation of a gas-rich galaxy merger so far. The formation of massive star clusters of $10^5$ to $10^7$~M$_{\\sun}$, triggered by the galaxy interaction, is directly resolved in this model. We show that these clusters are tightly bound structures with little net rotation, due to evolve into compact long-lived stellar systems. Massive clusters formed in galaxy mergers are thus robust candidates for progenitors of long-lived globular clusters. The simulated cluster mass spectrum is consistent with theory and observations. Tidal dwarf galaxies of $10^{8-9}$~M$_{\\sun}$ can form at the same time, and appear to be part of a different class of objects, being more extended and rotating. ",
        "introduction": "Globular clusters (GCs) are an important fossil record of the evolution of physical conditions in galaxies \\citep{bekki08}, because the formation of massive clusters is triggered by shocks and high pressures in the interstellar medium \\citep{vdb79,AZ01}. Fundamental properties along the Hubble sequence are a higher frequency of GCs around elliptical than disk galaxies \\citep{harris91} and a bimodal population in particular around early-type galaxies, with a population of low-metallicity GCs and a population of younger, higher-metallicity ones \\citep{AZ92}.    Since the works of \\citet{schweizer} and \\citet{whitmore}, there has been increasing evidence that young massive star clusters (YMC) form in interacting and merging galaxies, and these are often proposed to be GC progenitors. Theoretically, massive clusters form in mergers because of shocks and high turbulent pressure in interacting galaxies, which favors the formation of tightly bound clusters rather than unbound associations \\citep{elmegreen-efremov1997}. If GCs can form this way, given that most elliptical galaxies formed by galaxy mergers (Naab \\& Burkert 2003, Bournaud et al. 2005), this mechanism could account for the youngest and most metallic GCs that are more frequent around early-type galaxies \\citep{AZ01}. These GCs would come in addition to those formed early in the Universe as a result of thermal instabilities in proto-galaxies \\citep{fall-rees85}, strong shocks at the epoch of the Reionization \\citep{cen01}, or stripping of nucleated dwarf galaxies \\citep{freeman90,gao}. Nevertheless, that an important population of GCs formed in galaxy mergers is still challenged by some observations \\citep{spitler}. Whether or not YMCs formed during mergers contribute to the present-day GC populations is still an open question also because the long-term evolution of such YMCs is uncertain \\citep[see][]{degrijs07}. A requirement for YMCs to become long-lived GCs is that they should be gravitationally bound, which is difficult to assess observationally, as the velocity dispersion of a cluster does not directly trace its mass because of mass segregation \\citep{fleck06}. The survival issue is further complicated by the the mass-loss of clusters if their IMF is too shallow, and by the tidal field of the parent galaxy which may progressively disrupt these clusters \\citep{miocchi}.  Numerical simulations are a powerful tool to study galaxy mergers, but resolving the formation of star clusters in self-consistent models requires a huge dynamical range, because structures smaller than 100~pc should be resolved in a simulated volume larger than 100~kpc. Models by \\citet{bekki02} and \\citet{KG} have shown that the pressure and density required to form massive bound star clusters could be reached in galaxy mergers. They could however not identify individually forming GCs. \\citet{limaclow} studied individual GC formation in an indirect manner, using a model where absorbing sink particles are assumed to form above a chosen density threshold. This allows the mass and spatial distribution of the putative GCs to be studied. However, the absorbing nature of sink particles in such a model dooms any stellar association to be endlessly bound, so that whether or not long-lived bound objects can actually form in mergers is not probed by such a model.  In this Letter, we present a simulation of a wet galaxy merger, which to our knowledge is the highest resolution model of this kind so far. We directly resolve the formation of dense structures resembling Super Star Clusters (SSCs) with typical masses of $10^{5-7}$~M$_{\\sun}$. Comparing their gravitational and kinetic energy, we argue that they are tightly bound and likely progenitors of long-lived GCs. % ",
        "conclusions": "Using a high-resolution simulation of a galaxy merger, we resolve the formation of structures down to masses of $10^5$~M${\\sun}$. This enables us to directly reproduce the formation of numerous Super Star Clusters, as observed in merging galaxies. We find that these are dense, tightly bound structures that will evolve into compact stellar systems and are likely progenitors of numerous globular clusters (GC). Tidal dwarf galaxies of a few $10^8$~M$_{\\sun}$ are also formed in our model; they appear to be a different type of objects, less concentrated and resolved as rotating disks. Our results suggest that both TDGs and SSCs can form at the same time in galaxy mergers. The mass function of GC progenitors in this model has a power-law shape with a slope between -2 and -3, likely truncated around $10^7$~M$_{\\sun}$. The formation of globular cluster is much more efficient in merging galaxies than in isolated disk models. The total mass of the GC progenitors in our merger model amounts to 4\\% of the available gas mass, which is 0.7\\% of the total baryonic mass. The efficiency could be even higher with the large gas fractions observed in high-redshift disks \\citep{daddi}. GC progenitors form in the tails and caustics with high turbulent speed and pressure around the central merger remnant. The material in such regions has been expulsed from the spiral disks where it has been previously enriched; it is known to have relatively high metallicities \\citep{weilbacher03}, so the GCs formed this way should also be metal-rich. The excess of GCs around elliptical and lenticular galaxies compared to spirals could then be explained by this mechanism, given that mergers can form early-type galaxies and GCs at the same time."
    },
    "0806/0806.3383_arXiv.txt": {
        "abstract": "{ We present optical polarization maps of a sample of four interacting pairs at different {\\it phases} of encounter, from nearly unperturbed galaxies to on-going mergers. \\\\ Only the pair RR~24 shows a linear polarization pattern  which extends in both galaxies for several kiloparsecs. The more perturbed member, RR~24b, is lineraly polarized up to the level of $\\approx$3\\%. No polarization is measured in the strongly perturbed late-type pair members of RR~23 and RR~99. Also, in the central part of the double nuclei shell galaxy ESO 2400100 there is no significant polarization.\\\\ We use the ionized gas velocity field of RR~24 to interpret its linear polarization structure. In RR~24a the quite regular gas kinematics reflect the unperturbed spiral-like polarization structure. In RR~24b a strong velocity gradient in ionized gas could be associated with the polarization structure. We suggest that the  large-scale magnetic field of the RR~24 pair members still  plays a role in shaping the polarization pattern. } ",
        "introduction": "During a galaxy-galaxy interaction, tidal forces strongly deform the gravitational potential affecting both the stellar and  gas distribution and dynamics, which in turn are connected with the magnetic field structure of the galaxy \\citep[see e.g.][and references therein]{Vallee97,Moss00, Widrow2002}. Recent studies, in the radio domain, suggest a possible connection between the magnetic field structure and gas flows in spiral galaxies \\citep[see e.g.][]{Beck99,Chyzy04}. These  authors suggest that the field is mostly frozen into the gas and follows its motion during galaxy encounters; contrary to the expectation of a field generated by a galactic dynamo.  More recent studies have focused their attention on interacting  and merging spirals \\citep[see e.g.][]{Soida06,Wezgowiec07,Vollmer07}. Interacting spiral galaxies indeed show both large departures from a symmetric spiral shape and are expected to host gas flows. Perturbed galaxies, then,  constitute a good laboratory for the study of the interrelations between peculiar gas flows and magnetic field structure, but no clear results emerge from the literature, mainly due to the lack of high resolution kinematical information with which to compare the magnetic field structure \\citep[see e.g.][]{Soida01}. \\citet{Chyzy04} found that the magnetic field in NGC~4038/4039, the Antennae, is very different from what is observed in normal spirals. In particular, they notice that various regions in this merging galaxy reveal different physical  conditions and evolutionary stages. In particular, processes related to star formation tangle the field lines, so that little polarization is observed in star-forming regions. The magnetic field association with cold, warm and hot gas, depends on the particular place and dominance of various physical processes.  This leads to another open question, i.e. which gas phase has the strongest influence on the magnetic field evolution.  Observations then, suggest that large-scale galactic magnetic fields evolve  during violent interaction episodes, which could lead to accretion/merging phenomena.  This study presents optical polarization maps of a sample of interacting galaxies. It is well known that interstellar dust grains get aligned with a magnetic field which induces a polarization of the light passing through the dust cloud, via a dichroic extinction. The study of the polarization has been used to gain information about the magnetic field in galaxies \\citep[see e.g.][and reference therein]{Widrow2002,Lazarian07}. In particular, at optical and infrared wavelengths the analysis of polarization has contributed to the mapping of large--scale magnetic field structure in spiral galaxies \\citep[see e.g.][]{Scarrott96,Jones97,Jones00,Alton00}. Optical polarization maps reveal magnetic field structures, spatially coherent on scale-lengths of many kpc, both in {\\it normal} and {\\it active} galaxies although the presence of scattering phenomena makes the interpretation of the observations difficult \\citep{Widrow2002} in particular in dusty environments, where a significant amount of scattering can occur.   \\begin{table*} \\scriptsize{ \\caption{\\bf Overview of the sample} \\label{table1} \\begin{tabular}{lllllllc} \\hline\\hline &  RR~23a    &  RR~23b      & RR~24a       &  RR~24b      &   RR~99a     &  RR 99b    & E2400100 \\\\ &E1510361    & E1510360     &E2440120      &E2440121       &  E5520490    & E5520500    &            \\\\ &            & NGC 454      &               &                 &   NGC 1738  & NGC 1739   &     a/b nuclei       \\\\ \\hline Morphol. Type                  &     Irr    &   Pec  & Sb[2] & S0-a[2] &  SB(s)bc pec: & SB(s)bc pec:& SAB0:pec  \\\\ Hel. Sys. Vel.   [km~s$^{-1}$] & 3626$\\pm$2  &   3645 & 6852$\\pm$24 [3]& 6719$\\pm$55[3]& 3978$\\pm$30 & 3982$\\pm$32   &3159$\\pm$20/3348$\\pm$14  \\\\ &                         &                      & & & && \\\\ {\\bf Apparent magnitude}     &                         &                     &  & & &&\\\\ {\\bf and colours}:              &                         &                     &  & & &&\\\\ B$_T$                            &  13.11$\\pm$0.09 &   13.12$\\pm$0.21  &15.70$\\pm$0.09 &14.43$\\pm$0.09&  13.70$\\pm$0.09  & 14.24   &12.60$\\pm$0.30$$ \\\\ $\\langle$(B-R)$_T \\rangle$       &    1.02  &                            & 1.86 & 1.37   &   1.05  & 1.02      &1.39 \\\\ (J-H)$_{2MASS}$                    &             &                             &         & 0.708  &  0.574   & 0.551   &0.721\\\\ (H-K)$_{2MASS}$                   &            &                              &         &0.388   &    0.329 & 0.343   &0.195 \\\\ &            &                              &  & & & & \\\\ {\\bf Galaxy structure}:                  &                             &             &  & & & &\\\\ Effective Surf. Bright. $\\mu_e$(B) &    22.21$\\pm$0.31   &   22.33$\\pm$0.34  &             & 22.63$\\pm$1.46 &  21.38$\\pm$0.34  & 21.44$\\pm$0.35     & 21.84$\\pm$0.33\\\\ Average  P.A.                                &  57.3                      & 80.9                     &   131.9 & 31.6  &   35.6  & 104.4 &   132.5  \\\\ Average $\\epsilon$                         &  0.26                      & 0.37                      &    0.49& 31.6 &  0.44    & 0.51  &   0.46 \\\\ &                          &                      &    & & &  &    \\\\ {\\bf Kinematical parameters}        &                         &                       &    & && &\\\\ Vel.disp. $\\sigma_0$ [km~s$^{-1}$]  stars   &         &    &  71$\\pm$31& 80$\\pm$66 &        &   &263$\\pm$19/192$\\pm$13 \\\\ Gas max. rotation V$_{max}$  [km~s$^{-1}$]      &                          &     &240 [3]& &205$\\pm$9   &     &\\\\ &                          &                    &  & &    4          &        & 189\\\\ \\hline {\\bf  Pair parameters}   &  RR~23    &      & RR~24      &     & RR~99    &   & E2400100 \\\\ \\hline Arp-Madore ident.   &     AM 0112-554       &  & AM0115-444&                      &    & &      \\\\ Projected separation [\\arcmin]   &  0.45      &               & 0.3           &    &  0.53   &         & 0.08\\\\ $\\Delta$ V  [km~s$^{-1}$]        & 19        &               & 133          &     &  4     &        & 189  \\\\ \\hline \\end{tabular}  \\medskip  \\noindent{{\\bf Note}-- The photometric and kinematical data are obtained from  {\\tt NED} while structural data are from the {\\tt HYPERLEDA} compilation. Detailed studies of RR~24 and ESO~2400100 have been performed by \\cite{Rampa05} and \\cite{L98b}. In particular, ESO 2400100,  is considered a single early-type galaxy showing a shell system but it is, rather, composed of two distinct components, indicated in the table with $a$ and $b$ \\citep{L98b,Rampazzo03}.}} \\end{table*}  With this pilot study we aim to investigate the ability of  optical polarization maps to trace the evolution of the large--scale magnetic field structure during different {\\it phases} of the interaction phenomenon. Furthermore, polarization maps may unveil important effects induced during galaxy encounters, such as the displacements of dust from its more usual stellar-disk enviroment \\cite[see e.g.][]{Alton00}. We choose to investigate pairs of galaxies in very low density environments. \\citet[][]{Wezgowiec07} have shown that the cluster environment may affect the large-scale magnetic field, not only by galaxy interactions, but also by the interaction between the Inter Galactic Medium and the Inter Stellar Medium,  e.g. due to ram-pressure.  The paper is organized as follows. Section~2  presents the sample with the aim of characterizing the interaction episode according to the  morphological and kinematical  effects induced on each member galaxy. Section~3 details the observations and the reduction techniques. Results are organized as individual notes on paired galaxies in Section~4. In Section~5, we discuss the polarization maps in the context of the kinematics of the galaxies by comparing the polarimetric maps with  the 2D gas  velocity fields. We investigate whether the polarization vectors follow gas flows or maintain/develop their own structures independently. Conclusions and future prospects are sketched in Section~6. ",
        "conclusions": "We present a linear polarization study of a pilot sample of interacting galaxies, consisting of three pairs containing late-type members and a double nucleus shell galaxy. The latter should represent the fossil evidence of an accretion event, i.e. the final phase of a bound galaxy encounter.  We find the following: \\begin{itemize}  \\item{We did not detect a significant polarization in the very distorted late-type members, RR~23a, RR~99a and RR~99b. It is unclear if the lack of a polarization pattern is connected to the high degree of galaxy disruption.}  \\item{The two early-type galaxies RR~23b and ESO~2400100  do not show any linear polarization. The lack of polarization may be due to the low gas and dust content. Early-type galaxies may lack gas and Inter Stellar Medium turbulence, hence cannot drive a dynamo and may not host large-scale magnetic fields.}  \\item{Only the pair RR~24 shows a linear polarization pattern  which extends in both galaxies for several kiloparsecs. We use the 2D velocity field of ionized gas  in RR~24 members to interpret the linear polarization structure. RR~24a has quite regular gas kinematics. The gas velocity gradient in the North and South-West direction of the RR~24b nucleus matches the ``S-shape'' polarization pattern. We suggest that the  large-scale magnetic field of the RR~24 pair still  plays a role in shaping such a polarization patterns. We suggest that late-type galaxy magnetic fields may be distorted by interaction. The observed pattern may depend on the interaction stage.}  \\end{itemize}  The case of RR~24 suggests a  coherence among the gas kinematics, the structure of the polarization and possibly of the large-scale magnetic field. We believe that this pilot study underlines the need for a combined study of gas velocity field and polarization structure to infer  the evolution of the large-scale magnetic field stucture.  We are performing a campaign of  polarization measures of nearby galaxy pairs with the 1.82m Asiago Telescope using the AFOSC camera equipped with a Double Wallaston Prism \\citep{Desidera2002}. We intend also to obtain the 2D galaxy velocity fields using the GH$\\alpha$FaS Fabry-Perot \\citep{Hernandez08} at the 4.2m William Herschel Telescope at IAC in Canary Islands, Spain."
    },
    "0806/0806.4795_arXiv.txt": {
        "abstract": "We have extended the search for luminous ($M_{b_J} \\le -10.5$) compact stellar systems (CSSs) in the Virgo and Fornax galaxy clusters by targeting with the recently commissioned AAOmega spectrograph three cluster environments -- the cluster cores around M87 and NGC 1399, intracluster space, and a major galaxy merger site (NGC 1316). We have significantly increased the number of redshift-confirmed CSSs in the Virgo cluster core and located three Virgo intracluster globular clusters (IGCs) at a large distance from M87 (154-173 arcmin or $\\sim\\!750$-850$\\; \\mbox{kpc}$) -- the first isolated IGCs to be redshift-confirmed. We estimate luminous CSS populations in each cluster environment, and compare their kinematic and photometric properties. We find that (1) the estimated luminous CSS population in the Virgo cluster core is half of that in Fornax, possibly reflecting the more relaxed dynamical status of the latter; (2) in both clusters the luminous CSS velocity dispersions are less than those of the cD galaxy GC system or cluster dE galaxies, suggesting luminous CSSs have less energetic orbits; (3) Fornax has a sub-population of cluster core luminous CSSs that are redder and presumably more metal-rich than those found in Virgo; (4) no luminous CSSs were found in a 10-20$\\; \\mbox{ arcmin}$ (60-130$\\; \\mbox{kpc}$) radial arc east of the 3 Gyr old NGC 1316 galaxy merger remnant or in the adjacent intracluster region, implying that any luminous CSSs created in the galaxy merger have not been widely dispersed. ",
        "introduction": "High-resolution simulations of structure formation in the accepted $\\Lambda$CDM Universe\\footnote{$H_0 = 73 \\; \\mbox{km} \\, \\mbox{s}^{-1}$, $\\Omega_M = 0.27$, $\\Omega_\\Lambda = 0.73$} predict the present-day survival, around isolated large galaxies such as the Milky Way or M31, of many more satellite dark matter concentrations than are observed through their radiant baryons \\citep*[see discussion by][]{Cote..2002}. Two mechanisms may explain this apparent `missing satellite' problem.  Firstly that the epoch of reionization \\citep[see discussion by][]{Cote..2002} after the first stars formed at $z=12\\pm2$ \\citep{Moore..2006} halted the collapse of baryonic matter (hydrogen gas) into the smallest dark matter gravitational potential wells -- evidence supporting this hypothesis includes detection in the Local Group of dark matter dominated, very low surface brightness star clusters or dwarf galaxies \\citep{Belokurov..2007}.  Secondly that in the dense environments of galaxy clusters small dark matter haloes, although surviving the gradual hierarchical galaxy merger process, have been disrupted through tidal interaction with the cluster core. The visible remnants of this destruction are field stars and compact stellar systems (CSSs), which comprise the globular clusters (GCs) and the more recently discovered massive but sub-galactic luminous CSSs. Simulations by \\citet*{Henriques..2007} show that galaxy cluster field stars (the detectable intracluster light) from totally and partially disrupted dwarf galaxies can explain the observed shallow slope in the faint end of the cluster galaxy luminosity function. Tidal threshing simulations (\\citealt*{Bekki..2001}; \\citealt{Bekki..2003I}) predict the formation of luminous ($M_{b_J} \\le -10.5$) CSSs through the partial disruption of nucleated dwarf galaxies, by stripping away the stellar envelope to leave the naked nucleus of an ultra-compact dwarf galaxy (UCD). This process may also account for the surviving GCs, so comparing the properties of luminous CSSs with their fainter GC counterparts may reveal whether they form a single population -- for example, in the Fornax Cluster core region the luminosity distribution of the most luminous CSSs seems to extend smoothly into the bright tail of the GC population (\\citealt*{Mieske..2004I}; \\citealt{Drinkwater..2004}).  While both these mechanisms may have played a role, the great distances to nearby galaxy clusters limit observable evidence to remnant intracluster light and CSSs. Since disruption efficiency depends on environmental factors such as density and cluster mass, we investigate the spatial distribution, kinematics and photometric properties of an increased sample of Virgo and Fornax redshift-confirmed luminous CSSs in three cluster environments.  \\begin{itemize} \\item \\textbf{Cluster Core Environment.}  In the nearby Universe CSSs are numerous in the dense environments surrounding giant elliptical galaxies at the cores of galaxy clusters. Disruption through tidal interaction with the deep gravitational potential well of these giant galaxies is greatly enhanced in the cluster core, so we expect to observe a greater density of remnant CSSs and intracluster light \\citep{Bekki..2004}. The Virgo and Fornax clusters at distances less than $20 \\; \\mbox{Mpc}$ contain the central giant elliptical galaxies M87 and NGC 1399 that have been extensively imaged for point source CSS candidates. The rich Virgo Cluster containing several thousand galaxies is embedded in an extensive supercluster, whereas the Fornax Cluster containing several hundred galaxies is smaller and isolated. Fornax has approximately two times the central galaxy density but approximately half the velocity dispersion of Virgo -- this implies that Fornax is more dynamically relaxed than Virgo and we expect a more evolved CSS population. Observations in Virgo and Fornax provide an opportunity to compare CSS populations in two differing cluster core regions, since both clusters are at similar distances and close enough to explore the faint end of their CSS distributions.  The cluster core environment is dominated by the extensive GC system of the central elliptical galaxy. These GCs are thought to have multiple origins -- most were formed during the first mergers of galaxies that formed the cD galaxy, and the remainder were stripped from galaxies accreted by the central giant elliptical galaxy as it evolved \\citep[see review of GC formation mechanisms by][]{Ashman..1998}. GCs observed now are the survivors from an even larger population exposed to tidal destruction as the cluster core evolved. The complex and crowded environment in the cluster core blurs the forensic evidence of luminous CSS origins -- their distribution, kinematics and chemistry are consistent with both the tidal threshing and the bright GC theories.  The radial extent of cluster core environments can be approximated through imaging surveys of the fall in number density of point sources, predominantly GCs of the central giant elliptical galaxy, to a `background' level at which the number density profile flattens out. The transition to a background number density is difficult to precisely locate due to the relatively smooth decline in point source density and lack of supporting point source redshift measurements. Several photometric studies trace the overdensity of point sources across the core regions of Virgo and Fornax \\citep[e.g.][]{Jordan..2002, Hasegan..2005, Bassino..2006}. For the spatially isolated Fornax Cluster, \\citet{Bassino..2006} used this method to define a radial limit of $45\\pm5 \\, \\mbox{arcmin}$ for the core region centred on  NGC 1399. The Virgo Cluster is part of a more extended and complex dynamical structure, making it more difficult to define a core region boundary -- however, several previous wide-field photometric and dynamical studies of M87's GC system \\citep[e.g.][]{Cote..2001, Tamura..2006} show that it extends to at least the $60 \\; \\mbox{arcmin}$ radius of our AAOmega field.  \\item \\textbf{Intracluster Environment.} The intracluster environment beyond the tidal influence of the dominant cluster galaxies is relatively sparsely populated, mainly by dwarf elliptical (dE) galaxies. \\citet{West..1995} proposed that galaxy clusters have a population of intracluster globular clusters (IGCs) not gravitationally bound to the central giant elliptical or other cluster galaxies. Subsequent photometry in the central region of Fornax has suggested an excess of GC candidates near dwarf galaxies \\citep{Bassino..2003} and between NGC 1399 and the nearby galaxy NGC 1387 \\citep{Bassino..2006}, while in wide-field Virgo imaging by \\citet{Tamura..2006} GC candidates were found at some distance from major galaxies. More recently \\citet{Williams..2007a} detected four faint potential IGCs (not redshift confirmed) in HST images of the outer part of Virgo's cluster core region $\\sim\\!40 \\; \\mbox{arcmin}$ from M87. None of the existing CSS redshift surveys \\citep[e.g.][]{Drinkwater..2000a, Mieske..2004I, Bergond..2007} target IGCs far outside the cluster core environment -- our AAOmega observations include both Virgo and Fornax intracluster fields starting at least $\\sim350 \\, \\mbox{kpc}$ from the central giant elliptical galaxies.  \\item \\textbf{Galaxy Merger Environment.} The galaxy merger environment differs from the cluster core environment, being marked by the recent merging of massive gas-rich galaxies rather than ongoing accretion of many relatively small galaxies. When large galaxies interact or merge they potentially disrupt satellite objects, and gas-rich mergers such as the well-known Antennae galaxy pair undergo intense starburst activity along shock fronts. It has been predicted \\citep{Kroupa..1998, Fellhauer..2002, Fellhauer..2005} that luminous CSSs can evolve from stellar superclusters created from the merging of young massive clusters (YMCs) formed in such galaxy mergers.  NGC 1316, the radio source Fornax A in the Fornax Cluster, is a giant elliptical galaxy located over 1 Mpc or $3.7^\\circ$ south-west of the cluster core \\citep[see photometric study by][]{Schweizer..1980}. It shows evidence of a 3 Gyr old major merger event which resulted in the formation of  massive star clusters \\citep{Goudfrooij..2001a, Goudfrooij..2001b}. Based on the number density profile of its GC population, NGC 1316's radial extent is estimated to be approximately $7 \\; \\mbox{arcmin}$ \\citep[see figure 14 in][]{Goudfrooij..2001b}. In order to include more spatially dispersed merger products, we define the radial extent of the NGC 1316 galaxy merger environment to be $20 \\; \\mbox{arcmin}$ ($\\simeq\\!30 \\; \\mbox{kpc}$) encompassing the faint outer traces of its stellar envelope. With AAOmega we have undertaken a redshift survey of point sources within this radius to locate additional luminous CSSs produced by the galaxy merger.  The NW border of our M87 AAOmega field encompasses the strongly distorted galaxy NGC 4438 which may be undergoing a merger or high-speed interaction with the nearby galaxy NGC 4435 \\citep[see for example][]{Panuzzo..2007}. We did not specifically target these galaxies as a merger site, but we note in passing that our AAOmega survey results show no overdensity of confirmed foreground or cluster point sources that might be associated with them. \\end{itemize}  In Section 2 we describe Virgo and Fornax observations during 2006 with the AAOmega multi-fibre spectrograph at the 3.9-m Anglo-Australian Telescope (AAT). The Virgo observations were completed in 2006, between March 28 and April 1. The scientific objectives were to (1) compare the distributions of bright and faint CSSs, in order to test the hypothesis that bright CSSs are simply the bright tail of the GC luminosity distribution \\citep[e.g.][]{Mieske..2004I}; and (2) test the hypothesis that bright CSSs are the remnant nuclei of tidally stripped dE,N galaxies, and are consequently more widely distributed than classical GCs -- tidal threshing simulations by \\citep{Bekki..2003I} predict a cut-off in the UCD distribution at $2.5^\\circ$ from M87. Our Fornax observations during 2006, over 5 nights between December 12 and 16, were designed to (1) investigate the UCD--IGC interface by measuring the radial distribution and kinematics of CSSs around NGC 1399; and (2) locate CSSs in the intracluster and galaxy merger environments near NGC 1316.  In Section 3 we add our results to previous published data in order to investigate the spatial distribution, recession velocity dispersion and photometric properties of Virgo and Fornax CSSs in the defined cluster environments. Our findings are summarised in Section 4. Table \\ref{table:parameters} lists the parameters for the Virgo and Fornax galaxy clusters that we assume throughout this paper.  \\begin{table*} \\caption{Assumed Cluster Parameters} \\label{table:parameters} \\centering{ \\small{ \\begin{tabular*}{1.00\\textwidth} {@{\\extracolsep{\\fill}}lcclcl} \\hline \\hline Cluster/Galaxy & Distance & \\multicolumn{2}{c}{Distance Modulus} & \\multicolumn{2}{c}{Recession Velocity} \\\\ & (Mpc) & (mag) & & ($\\mbox{km} \\, \\mbox{s}^{-1}$) & \\\\[3pt] \\hline \\hline Virgo (M87) & 16.7 & 30.97 & \\citet{Mei..2007} & 1307$\\pm$7 & \\citet{Smith..2000}\\\\ Fornax (NGC 1399) & 18.3 & 31.35 & \\citet{Richtler..2000} & 1425$\\pm$4 & \\citet{Graham..1998}\\\\ NGC 1316 & 22.8 & 31.79 & \\citet{Richtler..2000} & 1760$\\pm$10 & \\citet{Longhetti..1998}\\\\[3pt] \\hline \\end{tabular*} }} \\end{table*} ",
        "conclusions": "\\subsection{Cumulative CSS Datasets} Our aim is to compare the spatial distribution, kinematics and colours of spectroscopically confirmed Virgo and Fornax CSSs in different cluster environments to the predictions of CSS formation theories. To provide the largest possible spectroscopically confirmed dataset for this comparison, we have combined bright CSSs detected in our AAOmega observations with previously catalogued GCs, dEs and bright CSSs from the sources listed below. We will use the combined M87 and NGC 1399 datasets in subsection \\ref{clustercore}, and the NGC 1316 dataset in subsection \\ref{galaxymerger}.  \\begin{itemize} \\item \\textbf{M87:} (a) Using 2dF \\citet{Jones..2006} located 9 UCDs in the core region of the Virgo Cluster through a survey of colour-selected point sources ($b_J<21.5$). More precise recession velocities, together with other properties, were subsequently published for six of these UCDs by \\citet{Evstigneeva..2007}. (b) The key source of published redshift data for the M87 GC system is the survey by \\citet{Hanes..2001} confirming 286 GCs with recession velocities between 500 and 2500 $\\mbox{km} \\, \\mbox{s}^{-1}$. In addition, \\citet{Schuberth..2006} has published redshift data for 174 GCs located between 0.9 and $15.5 \\; \\mbox{arcmin}$ from M87. \\item \\textbf{NGC 1399:} (a) 60 UCDs were located with 2dF in the all-object Fornax Cluster Spectroscopic Survey \\citep{Drinkwater..2000b, Phillipps..2001} to $b_J<19.7$, and in a follow-up survey of fainter ($b_J<21.5$) blue point sources in the core region \\citep{Drinkwater..2005, Gregg..2007}. More precise recession velocities, were subsequently published for four of these UCDs using the UVES-VLT spectrograph \\citep{Hilker..2007}. (b) The dataset comprising 468 redshifts published by \\citet{Dirsch..2004} is our key source for the inner (2 to 9$\\; \\mbox{arcmin}$) GC system of NGC 1399. Approximately 160 additional redshifts of the outer GC system were described by \\citet{Schuberth..2004} but the dataset and positions have not yet been published. \\citet{Bergond..2007} has published redshifts for 149 GCs, including 61 that are potentially IGCs. However, by using photometric data \\citet{Schuberth..2008} contends that all but one of these IGCs are either bound to a local host galaxy or part of the extended GC system of NGC 1399. \\item\\textbf{NGC 1316:} \\citet{Goudfrooij..2001a} has published a catalog of 24 CSSs, including 10 with $M_V<-10.8$, found by multi-slit spectroscopy within an $8 \\times 8 \\, \\mbox{arcmin}^{-2}$ field centred on NGC 1316. \\end{itemize}  In the following analysis we avoid sampling bias when comparing between Virgo and Fornax or extrapolating results to the overall CSS population, by extracting subsets common to the different sampled parameter spaces (colours, magnitudes and areal coverage), as described below.  \\subsection{The Cluster Core Environment} \\label{clustercore} In this subsection we analyse the central CSS populations and in particular compare the Fornax and Virgo populations.  \\subsubsection{Bright CSS Populations in Virgo and Fornax} To better understand their origins, we firstly identify luminous CSSs that are gravitationally bound to non-dwarf satellite galaxies of M87 or NGC 1399 (see Fig.~\\ref{fig:aao_1010}). In the cluster core regions dominated by giant elliptical galaxies, CSSs are considered to be bound objects if they are within the tidal radius of a non-dwarf galaxy and have a relative recession velocity less than the escape velocity \\citep[see][for a description of this method]{Firth..2007}. Since only 4 out of 105 luminous CSSs in the NGC 1399 cluster core region (and none in the M87 region) are bound to satellite galaxies, our results are not significantly affected by their inclusion.  \\begin{figure*} \\centering \\includegraphics[width=8.5cm]{M87_1010.eps} \\includegraphics[width=8.5cm]{NGC1399_1010.eps} \\caption{Distribution of bound CSSs (triangles) and unbound CSSs (crosses) within the Virgo and Fornax cluster core environments. Tidal radii (circles) of non-dwarf galaxies are computed with respect to M87 and NGC 1399 respectively. Identifying numbers for the galaxies with larger tidal radii are from the Virgo and Fornax Cluster Catalogues. In intracluster environments, tidal radii are meaningless as an indicator of bound CSSs.} \\label{fig:aao_1010} \\end{figure*}  Table \\ref{table:csscore} combines previous 2dF surveys \\citep{Drinkwater..2000a, Gregg..2007, Jones..2006} with our AAOmega surveys, and compares in absolute magnitude bins the number of luminous CSSs located in the core regions surrounding M87 and NGC 1399, separating those CSS that are brighter ($M_{b_J} \\le -10.5$) than the typical magnitude range of GCs. After eliminating target duplications, we calculate completeness with respect to the APM catalogue of point sources within the respective cluster core fields and colour-magnitude limits ($15.0 < b_J < 21.8$, $b_J - R \\le 1.6$). The Palomar survey plates used for the Virgo APM catalogue have a shallower faint limit than the UKST survey plates used for the Fornax APM catalogue, so APM point sources in our M87 catalogue tail off for $M_{b_J}>-10.5$, but we restrict our CSS population estimates to the luminous CSSs. Both cluster core fields in Table \\ref{table:csscore} have the same spatial extent ($\\sim\\!630 \\; \\mbox{kpc}$) -- the $1.8^\\circ$-diameter field for the more distant NGC 1399 is spatially equivalent to the $2^\\circ$-diameter M87 field.  Within the magnitude range $M_{b_J}<-10.5$ we estimate, by adjusting for completeness and using Poisson statistics for uncertainty estimation, that M87 has $25\\pm8$ luminous CSSs, or half the number ($47\\pm12$) estimated for NGC 1399. This finding contrasts sharply with the relative size of the innermost GC populations, usually expressed as the luminosity-scaled specific frequency $S_N$\\footnote{$S_N = N_{GC} 10^{0.4(M_V+15)}$, where $N_{GC}$ is the estimated number of GCs and $M_V$ is the galaxy absolute magnitude.} \\citep[see review by][]{Elmegreen..1999}. Based on observations within $7 \\; \\mbox{arcmin}$ of M87 and NGC 1399, \\citet{Forte..2002} calculated $S_N$ of $6\\pm1$ and $3.7\\pm0.8$, galaxy luminosities of $V=8.54\\pm0.01$ and $V=9.02\\pm0.06$, and total GC populations of $4700\\pm400$ and $2300\\pm300$ respectively. If luminous CSSs are simply an extension of the central giant elliptical galaxy GC population, we would expect the Virgo cluster core to have a luminous CSS population $\\sim4$ times greater than we have estimated from our redshift surveys. Even though the $S_N$ estimates are not redshift confirmed we consider they are likely to be accurate within a factor of $\\sim\\!2$, so our contrary estimates for the luminous CSS populations may be due either to a real difference with the GC populations or to uncertainties in population estimates caused by lack of observing completeness in Virgo (see Table \\ref{table:csscore}) -- further observations would clarify this.  \\begin{table*} \\caption{Cluster Core Luminous CSS Results} \\label{table:csscore} \\centering \\scriptsize{ \\begin{tabular*}{1.00\\textwidth} {@{\\extracolsep{\\fill}}ccccccc} \\hline \\hline $M_{b_J}$ Range & APM & \\multicolumn{2}{c}{Redshifts} & Completeness & \\multicolumn{2}{c}{CSSs Found}\\\\ \\cline{3-4} \\cline{6-7} (mag) & Targets & 2dF$^a$ & AAOmega$^b$ & (per cent) & 2dF & AAOmega\\\\[3pt] \\hline \\hline\\\\ \\multicolumn{6}{l}{M87 -- $2^\\circ$ Field}\\\\ -14.0 to -14.5 & 344 & 66  &    & 19.2 & 0 & -- \\\\ -13.5 to -14.0 & 373 & 73  & 21 & 25.2 & 0 & 0 \\\\ -13.0 to -13.5 & 399 & 132 & 8  & 35.1 & 1 & 0 \\\\ -12.5 to -13.0 & 503 & 238 & 5  & 48.3 & 0 & 0 \\\\ -12.0 to -12.5 & 537 & 253 & 3  & 47.7 & 0 & 0 \\\\ -11.5 to -12.0 & 505 & 301 & 6  & 60.8 & 1 & 0 \\\\ -11.0 to -11.5 & 430 & 240 & 15 & 59.3 & 5 & 2 \\\\ -10.5 to -11.0 & 345 & 229 & 31 & 75.4 & 3 & 1 \\\\ \\\\ -10.0 to -10.5 & 18  &     & 54 & --   & -- & 5 \\\\ -9.5  to -10.0 & 0   &     & 30 & --   & -- & 8 \\\\ \\\\ \\multicolumn{6}{l}{NGC 1399 -- $1.8^\\circ$ Field}\\\\ -15.0 to -15.5 & 183  & 82  & 34  & 63.4  & 0  & 0 \\\\% AAO 34 - 2dF repeats 3 (see NGC1399_2dFCatalog_vcat4runz) -14.5 to -15.0 & 225  & 214 & 22  & 100.0 & 0  & 0 \\\\% AAO 22 - 2dF repeats 8 -14.0 to -14.5 & 281  & 255 & 14  & 95.7  & 0  & 0 \\\\% AAO 14 - 2dF repeats 8 -13.5 to -14.0 & 328  & 319 & 17  & 100.0 & 0  & 0 \\\\% AAO 17 - 2dF repeats 6 -13.0 to -13.5 & 347  & 341 & 5   & 99.7  & 0  & 0 \\\\% AAO 5 - 2dF repeats 5 -12.5 to -13.0 & 466  & 458 & 7   & 99.8  & 0  & 0 \\\\ -12.0 to -12.5 & 588  & 575 & 12  & 99.8  & 0  & 0 \\\\ -11.5 to -12.0 & 758  & 673 & 16  & 90.9  & 1  & 0 \\\\ -11.0 to -11.5 & 1040 & 705 & 23  & 70.0  & 4  & 0 \\\\ -10.5 to -11.0 & 1474 & 642 & 33  & 45.8  & 17 & 0 \\\\ \\\\ -10.0 to -10.5 & 1819 & 376 & 80  & 25.1  & 15 & 5 \\\\ -9.5  to -10.0 & 2506 & 71  & 33  & 4.2   & 2  & 1 \\\\ \\hline \\end{tabular*} } \\begin{list}{\\hspace{0.5 cm}}{} \\item a. Redshifts from 2dF surveys \\citep{Drinkwater..2000a, Gregg..2007}.\\\\ \\item b. Some AAOmega redshifts duplicate 2dF redshifts, but this has been corrected in our completeness percentages.\\\\ \\item c. We separate the absolute magnitude range $M_{b_J} \\le -10.5$ covering the luminous CSS and for which we have a reasonable level of redshift completeness in both galaxy clusters.\\\\ \\end{list} \\end{table*}   \\subsubsection{Recession Velocity Dispersion} Recession velocity dispersion is a useful method to compare CSS sub-populations that may have differing kinematical histories -- for example, how do the velocities of luminous CSSs (at least two or three magnitudes brighter than the peak of the GC luminosity function) compare with their fainter counterparts? Table \\ref{table:velocities} and Fig.~\\ref{fig:combinedcz_1030} compare the heliocentric recession velocity distributions of GCs and luminous CSSs in the Fornax and Virgo cluster core regions.  Firstly, the larger velocity dispersions of both GCs and luminous CSSs in Virgo shows that both sub-populations are dynamically more energetic in Virgo than in Fornax, as confirmed by the statistical test results shown in Table \\ref{table:velocities}. These results are expected since M87 is more massive than NGC 1399, and embedded in a more complex cluster undergoing relaxation. Secondly, in each cluster the velocity dispersion of luminous CSSs is less than that of the central galaxy's GC system; however, this difference is more significant in Fornax ($<0.01$ $F$-test probability that this is due to sampling variability) than in Virgo (0.49 $F$-test probability that this is due to sampling variability).  These results provide clear evidence in Fornax that luminous CSSs form a dynamically distinct population to the central galaxy GCs. A possible explanation is that in Fornax the majority of luminous CSSs formed from tidally stripped dwarf galaxies in preferentially less energetic orbits \\citep[as discussed by][]{Bekki..2001, Bekki..2003I}.  \\begin{table*} \\caption{Cluster Core Recession Velocity Data} \\label{table:velocities} \\centering \\scriptsize{ \\begin{tabular*}{1.00\\textwidth} {@{\\extracolsep{\\fill}}lcccccccccc} \\hline \\hline \\vspace{2pt} & \\multicolumn{6}{c}{Mean Velocity} & & \\multicolumn{3}{c}{Velocity Dispersion}\\\\ \\cline{2-7} \\cline{9-11} & \\multicolumn{2}{c}{GCs} & & \\multicolumn{2}{c}{Luminous CSSs} & $t$-test$^c$ & & GCs & Luminous CSSs & $F$-test$^d$\\\\ & (N) & ($\\mbox{km} \\, \\mbox{s}^{-1}$) & & (N) & ($\\mbox{km} \\, \\mbox{s}^{-1}$) & Probability &  & ($\\mbox{km} \\, \\mbox{s}^{-1}$) & ($\\mbox{km} \\, \\mbox{s}^{-1}$) & Probability \\\\[3pt] \\hline \\hline\\\\ Fornax (NGC 1399) & 758 & $1443\\pm37^a$ & & 109 & $1475\\pm35$ & 0.19 & & $313\\pm40^a$ & $228\\pm34$ & $<0.01$ \\\\ \\\\ Virgo (M87) & 455 & $1179\\pm74^b$ & & 26 & $1238\\pm38$ & 0.41 & & $387\\pm39^b$ & $344\\pm15$ & 0.49 \\\\ \\\\ Comparison Tests & & $<0.01$ $^c$ & & & $<0.01$ $^c$ &  &  & $<0.01$ $^d$ & $<0.01$ $^d$ &  \\\\[3pt] \\hline \\end{tabular*} } \\begin{flushleft} \\item a. GC data extracted from \\citet*{Mieske..2002}; \\citet{Dirsch..2004}; \\citet{Mieske..2004I, Bergond..2007}.\\\\ \\item b. GC data extracted from \\citet{Hanes..2001, Schuberth..2006}.\\\\ \\item c. The $t$-test probability tests the hypothesis that two sample means are drawn from the same population.\\\\ \\item d. The $F$-test probability tests the hypothesis that two sample variances are drawn from the same population.\\\\ \\end{flushleft} \\end{table*}  \\begin{figure*} \\centering \\includegraphics[width=15cm]{combinedcz_1030.eps} \\caption{CSS recession velocity distribution as a function of radial distance from the cD galaxy. GCs (points) are compared with bright/spatially dispersed CSSs (filled circles). The horizontal dashed lines mark cD galaxy recession velocities. Vertical lines show $\\pm1\\sigma$ velocity dispersions of GCs (dotted line) and luminous CSSs (dashed line). \\textsc{Left:} The Virgo Cluster core (M87) luminous CSS velocity dispersion is $344\\pm15 \\; \\mbox{km} \\, \\mbox{sec}^{-1}$. \\textsc{Right:} The Fornax Cluster core (NGC 1399) luminous CSS velocity dispersion is $228\\pm34 \\; \\mbox{km} \\, \\mbox{sec}^{-1}$.} \\label{fig:combinedcz_1030} \\end{figure*}  \\subsubsection{CSS Photometric Properties} To compare CSS colours and magnitudes in the core regions of Virgo and Fornax, we need to convert our results to the same extinction-corrected (de-reddened) photometry. We therefore obtained $gri$ de-reddened Virgo photometry from SDSS \\citep{Adelman..2006}\\footnote{The final SDSS $ugriz$ photometry obtained with the 2.5-m survey telescope at Apache Point Observatory should not be confused with the $u^\\prime g^\\prime r^\\prime i^\\prime z^\\prime$ system \\citep{Fukugita..1996} which refers to filter/grating-detector combinations mounted on other telescopes (such as the SDSS $20^{\\prime\\prime}$ photometric monitoring telescope).} and cross-matched our Fornax CSS results to de-reddened CTIO photometry \\citep{Karick..2007} covering most of the NGC 1399 one square degree field, having checked that these two photometry sources are well calibrated. We were unable to obtain CTIO photometry for 10 Fornax CSSs and have excluded them from the following analysis.  As expected, the extinction corrections \\citep*[based on][]{Schlegel..1998} for Fornax were smaller than those for Virgo, due to their very different galactic latitudes. Fig.~\\ref{fig:colmag_2} compares extinction-corrected colours and absolute magnitudes of luminous CSSs in the Virgo and Fornax cluster core environments, and shows target selection limits for the various 2dF and AAOmega observations. The luminosity distributions of Virgo and Fornax luminous CSSs in the cluster core region appear to scale with cD galaxy mass (M87 has approximately four times the mass of NGC 1399) -- for example 10 Virgo CSSs are brighter than $M_{g,0}=-11.4$ compared with only 6 Fornax CSSs. To properly compare the luminous CSS distributions in Fornax and Virgo, we define a colour-magnitude range (gray box in Fig.~\\ref{fig:colmag_2}) common to the various observation sets; on the colour index axis we use the AAOmega Virgo target selection limits, and on the magnitude axis we use the \\citet{Gregg..2007} faint limit and a bright limit of $\\mbox{M}_{g,0}=-12.4$ to exclude the unusually luminous CSSs. Our bright limit specifically excludes Fornax UCD3 \\citep{Drinkwater..2000a} and Virgo UCD7 \\citep{Jones..2006} which are atypically bright for CSSs -- the luminosity of UCD3 appears to be increased by a background spiral galaxy \\citep{Evstigneeva..2007}, while UCD7 is bright enough to be a dwarf galaxy.  \\begin{figure} \\centering \\includegraphics[width=8.5cm]{colmag_2.eps} \\caption{Colour-magnitude plot of bright CSSs in the cluster core regions of Virgo (unfilled triangles) and Fornax (solid triangles) from wide-area 2dF and AAOmega observations. Typical photometry error bars are shown at upper-left for Virgo and upper-right for Fornax. The labelled target selection limits are: in Fornax, \\citet[F1:][]{Drinkwater..2000a} and our AAOmega NGC 1399 survey with limits the same as \\citet[F2:][]{Gregg..2007}; in Virgo, \\citet[V1:][]{Jones..2006}. The dashed line shows the extinction-corrected approximate limits of our AAOmega Virgo survey (based on $0.5<(g-r)<1.0$ and $g \\le 21.5$). The grey box shows the parameter space common to the 2dF/AAOmega Fornax and Virgo observations. An apparent gap ($-11.2<M_{g,0}<-10.6$) in the luminosity distribution of M87 CSSs is due to low completeness of our AAOmega observations in this magnitude range.} \\label{fig:colmag_2} \\end{figure}  Colour is a useful measure to compare luminous CSS with GCs in order to investigate how they are related. Studies of extragalactic GC systems \\citep[see review by][]{Ashman..1998} show that, although the colour of individual GCs is not a strong function of their luminosity, `broad-band colours reflect the age and metallicity of stellar populations' and the average colours of GC systems trend redder (more metal-rich) with increasing host galaxy luminosity. Our results for M87 and NGC 1399 luminous CSS seem contrary to this trend -- in contrast to the Virgo luminous CSS, Fornax luminous CSS extend to redder colours. The mean $g-r$ colour indices of redshift-confirmed Fornax and Virgo luminous CSSs, within the parameter space common to the 2dF/AAOmega Fornax and Virgo observations (grey box in Fig.~\\ref{fig:colmag_2}), are $0.66\\pm0.02$ and $0.63\\pm0.02$ respectively (with standard deviations of $0.09\\pm0.01$ and $0.06\\pm0.01$). Although the means appear similar, the Fornax luminous CSS $g-r$ population has a significant red tail not found in the Virgo population -- assuming Gaussian colour distributions, the $t$-test and $F$-test probabilities are only 0.09 and 0.06 respectively that the colour distributions of the luminous CSS populations are indistinguishable; and the non-parametric K-S test probability of 0.10 also confirms this result.  The average colour of individual GCs has been observed to become bluer with increasing radial distance from their host galaxy \\citep[see review by][]{Ashman..1998}, although this colour gradient is not observed at larger galactocentric radii. In Fig.~\\ref{fig:virgofornaxds9} we compare the wide-area spatial and radial distributions of red ($g-r \\ge 0.7$ or $g_0-r_0 \\ge 0.68$) and blue Virgo and Fornax luminous CSS sub-populations. By eye it appears that there is a concentration of blue (presumably metal-poor) luminous CSSs near the central massive galaxy in each cluster. However, while the Virgo red sub-population is too small to give reliable statistics, a 2D K-S test confirms (probability statistic 0.1) that the red and blue luminous CSS sub-populations around NGC 1399 have similar overall spatial distributions.  \\begin{figure*} \\centering \\includegraphics[width=8.7cm]{M87_highgr.eps} \\includegraphics[width=8.7cm]{NGC1399_highgr.eps} \\caption{Comparison of spatial and radial distributions of Virgo (left: $10^\\prime \\simeq 50 \\; \\mbox{kpc}$) and Fornax (right: $10^\\prime \\simeq 60 \\; \\mbox{kpc}$) cluster core luminous CSS with respect to their $g-r$ colours. Each DSS blue plate image is 60 arcmin across; North is up and east to the left. The red CSS sub-populations ($g-r \\ge 0.7$) are shown as boxes and the blue sub-populations as circles.} \\label{fig:virgofornaxds9} \\end{figure*}  Fig.~\\ref{fig:colmag_6} compares the SDSS extinction-corrected colour indices of CSSs in the Fornax and Virgo cluster core environments with those of dEs and GCs. To aid our analysis, we separate the brighter and fainter luminous CSSs at an arbitrary deredenned magnitude of $M_{g,0} = -11$. We also indicate with circles those CSSs outside the magnitude limits (grey box) shown in Fig.~\\ref{fig:colmag_2}, and exclude them from our analysis. Fornax luminous CSSs, both brighter and fainter, closely follow the point source (mainly stellar) locus. Virgo luminous CSSs (not circled) are distributed on two best fitting lines -- the brighter CSS's (solid line: $[g-r] = (0.76\\pm0.33)[r-i] + (0.40\\pm0.10)$) more closely follow the distribution of Virgo dEs, whereas the fainter CSSs (dashed line: $[g-r] = (0.23\\pm0.19)[r-i] + (0.56\\pm0.06)$) appear more widely dispersed like the M87 GCs. We could interpret the differing distributions of brighter and fainter sub-populations of luminous CSSs in the Virgo Cluster to support the hypothesis that the brightest CSSs come from tidally-stripped dE,N galaxies while the fainter CSSs are the luminous tail of the cluster core GC distribution -- however the dashed line for faint CSS and the GC distribution are influenced strongly by the uncertainties in photometry (see typical error bars on the plot) and small number statistics (evidenced by the uncertainties in best fitting line slopes quoted above).  \\begin{figure*} \\centering \\includegraphics[width=15cm]{colmag_6.eps} \\caption{Colour-colour diagram of cluster core CSSs in Fornax (filled symbols) and Virgo (open symbols) from SDSS dereddened photometry. The brighter CSSs ($-12.4<M_{g,0}<-11$) are shown as large squares and fainter CSSs ($-11<M_{g,0}<-10$) as smaller triangles. CSSs brighter or fainter than the common Fornax and Virgo observation colour-magnitude limits discussed in the text (see also Fig.~\\ref{fig:Virgogri}) are encircled, and are not taken into account in the analysis. For Virgo we show the best fitting lines (continuous and dashed) for the bright and faint CSSs within these common colour-magnitude limits. The distribution of Virgo CSSs are compared with redshift confirmed Virgo cluster core dEs \\citep[grey crosses:][]{Binggeli..1985} and M87 GCs \\citep[small black squares:][]{Hanes..2001}. Typical 1-$\\sigma$ error bars are shown (upper left and right) for the faint and bright Virgo luminous CSS.} \\label{fig:colmag_6} \\end{figure*}  \\subsection{The Intracluster Environment} IGCs are of interest for several different reasons. By obtaining redshift confirmation of a sample of IGCs, their overall number can be estimated -- if redshift-confirmed IGCs prove to be sufficiently numerous, we can treat them as test particles to map the cluster potential (i.e. the dark matter distribution). Other properties of IGCs can also be compared with intracluster dE galaxies and objects in the cluster core environment to assess whether IGCs are a distinct sub-population with independent origins and kinematics. For example, at large distances from the cluster giant elliptical galaxies, IGC recession velocities should follow those of the overall cluster if they are `freely floating in the cluster potential' as suggested by \\citet{West..1995}.  Several CSSs have been spectroscopically confirmed by previous researchers in the Fornax cluster environment closely surrounding the cluster core region. For example, \\citet{Bergond..2007} surveyed a $500 \\times 150 \\; \\mbox{kpc}$ area centred on NGC 1399 and found 61 potential IGCs, all except one subsequently shown \\citep{Schuberth..2008} to be gravitationally associated with NGC 1399 or its surrounding galaxies. Our AAOmega observations in Virgo and Fornax intracluster fields are the first attempt to survey wide-area ($2^\\circ$-diameter) fields for candidate IGCs at least $\\sim\\!300 \\; \\mbox{kpc}$ from the cluster centres.  We found three IGCs in the Virgo intracluster field (none of them adjacent to a major galaxy -- see Fig.~\\ref{fig:aao_1010a}), but we located no IGCs in the Fornax intracluster field. The Virgo IGCs are sufficiently distant from any major cluster galaxies to be classified as gravitationally unbound. These are the first examples of redshift-confirmed IGCs that can be compared with the population predicted by \\citet{West..1995}. Two of the Virgo IGCs have recession velocities of $\\sim\\!1150 \\; \\mbox{km} \\, \\mbox{s}^{-1}$ (similar to that of GCs in the cluster core), and one has an unusually high velocity $\\sim\\!1980 \\; \\mbox{km} \\, \\mbox{s}^{-1}$. Their $g-r$, $r-i$ colours are in the range of CSSs in the Virgo cluster core region (see Fig.~\\ref{fig:Virgogri}), although this may be due to our colour selection limits. The IGCs appear as isolated point sources in SDSS images -- as noted in Table \\ref{table:aaovirgolist} the brightest one shows hints of resolution on the best SDSS image, with the red, extended source just to the north being very likely a background galaxy.  \\begin{figure} \\centering \\includegraphics[width=8.5cm]{12h30m_IGCs.eps} \\caption{Virgo IGCs and intracluster galaxies. Virgo IGCs (small crosses) are overlaid on a DSS image (POSS-II F) centred at 12:31:30, +09:44:00. North is up and east is to the left. We show potentially associated surrounding galaxies (all are dwarfs if cluster members) from the Virgo Cluster Catalogue \\citep[labelled with VCC numbers:][]{Binggeli..1985} that are either confirmed cluster members (VCC 1444) or have unknown recession velocities. The IGC at lower-right is $3 \\; \\mbox{arcmin}$ ($\\sim\\!15 \\; \\mbox{kpc}$) from the nearest dwarf galaxy (VCC 1207, a possible member).} \\label{fig:aao_1010a} \\end{figure}  Table \\ref{table:cssintracluster} summarises the results of our AAOmega point source survey of the intracluster environment. We have estimated an upper limit for the population of bright IGCs in the Virgo and Fornax intracluster environments by using the hypergeometric distribution\\footnote{The hypergeometric distribution is applicable to a binomial sampling process without replacement. We used the online calculator available at http://www.quantitativeskills.com/sisa/distributions/hypergeo.htm.} with 95 per cent confidence limit. In Fornax we found no IGCs (upper limit 2 IGCs) in observations that were 81 per cent complete to a faint limit of $M_{b_J}<-11.5$, while in Virgo we found 3 IGCs (see Fig.~\\ref{fig:aaovirgofornax_3}) to a 54 per cent complete faint limit of $M_{b_J}<-10.2$ suggesting a population upper limit of 9 IGCs.   \\begin{table} \\caption{Intracluster/Galaxy Merger CSS Results} \\label{table:cssintracluster} \\centering \\scriptsize{ \\begin{tabular}{lcccc} \\hline \\hline $M_{b_J}$ Range & Available & Redshifts & Completeness & CSSs\\\\ (mag) & Targets & AAOmega & (per cent) & Located \\\\[3pt] \\hline \\hline\\\\ \\multicolumn{5}{l}{Virgo Intracluster Field}\\\\ $<-14.0$ & 4 & 1 & 25.0 & - \\\\ -13.5 to -14.0 & 97& 43 & 44.3& - \\\\ -13.0 to -13.5 & 138 & 55 & 39.9 & - \\\\ -12.5 to -13.0 & 146 & 34 & 23.3 & - \\\\ -12.0 to -12.5 & 186 & 72 & 38.7 & - \\\\ -11.5 to -12.0 & 162 & 123 & 75.9 & - \\\\ -11.0 to -11.5 & 185 & 144 & 77.8 & 1 \\\\ -10.5 to -11.0 & 277 & 106 & 38.3 & - \\\\ -10.0 to -10.5 & 298 & 138 & 46.3 & 2 \\\\ -9.5 to -10.0 & 454 & 76 & 16.7 & - \\\\ -9.0 to -9.5 & 420 & 10 & 2.0 & - \\\\ \\\\ \\multicolumn{5}{l}{Fornax Intracluster Field}\\\\ $<-14.0$ & 521 & 444 & 85.2 & - \\\\ -13.0 to -13.5 & 199 & 148 & 74.4 & - \\\\ -12.5 to -13.0 & 247 & 213 & 86.2 & - \\\\ -12.0 to -12.5 & 255 & 197 & 77.3 & - \\\\ -11.5 to -12.0 & 304 & 228 & 75.0 & - \\\\ -11.0 to -11.5 & 350 & 74 & 21.1 & - \\\\ \\\\ \\multicolumn{5}{l}{NGC 1316 Galaxy Merger Field}\\\\ $<-14.0$ & 39 & 31 & 79.5 & - \\\\ -13.0 to -13.5 & 13 & 10 & 76.9 & - \\\\ -12.5 to -13.0 & 14 & 12 & 85.7 & - \\\\ -12.0 to -12.5 & 21 & 9 & 42.9 & - \\\\ -11.5 to -12.0 & 26 & 3 & 11.5 & - \\\\ \\\\ \\hline \\end{tabular} } \\end{table}  Our survey of intracluster fields in Fornax and Virgo was limited to the relatively luminous CSSs located at a considerable distance from the cluster centres. We have no theoretical predictions with which to compare these results -- the predictions by \\citet{West..1995} relate to the environment of the brightest cluster galaxy. Our low IGC count may be due to the areal limit of our survey, since outside the cluster core the distribution of IGCs may be non-uniform, mirroring the cluster galaxy distribution. Our findings suggest that populations of IGCs may exist in both clusters, but that we need to probe a wider area of the intracluster regions to a deeper faint limit in order to further investigate their distribution and origins.  \\subsection{The Galaxy Merger Environment} \\label{galaxymerger} Active gas-rich galaxy mergers are observed to contain close groupings of newly-condensed star clusters \\citep[such as those observed in the Antennae galaxy pair][]{Kroupa..1998}, and numerical simulations \\citep[e.g.][]{Fellhauer..2002} confirm that the largest of these should survive disruptive tidal forces and merge on relatively short time-scales ($\\sim\\!100 \\; \\mbox{Myr}$) to form stellar superclusters. Superclusters with mass of order $10^8 \\; \\mbox{M}_{\\odot}$ may then evolve into present-day bright CSSs, or potentially into spheroidal dwarf galaxies with low dark matter content. This mechanism for luminous CSS formation is theoretically supported by simulations \\citep{Fellhauer..2002, Fellhauer..2005}, whereas their formation as ancient but unusually massive GCs is not well understood. On the other hand, it is unclear whether stellar superclusters are ejected with enough efficiency from their formation sites to account for the CSS populations we find in cluster core and intracluster environments.  The NGC 1316 galaxy merger environment borders our Fornax intracluster field (see Fig.~\\ref{fig:aaovirgofornax_11}). \\citet{Goudfrooij..2001a, Goudfrooij..2001b} found evidence of both old and intermediate age (high hydrogen content) GCs within an $8^\\prime\\times8^\\prime$ region surrounding this 3-Gyr-old merger remnant galaxy, including four exceptionally bright CSSs ($M_V<-12$). Although we radially restrict our definition of the NGC 1316 galaxy merger environment to $<\\!20 \\; \\mbox{arcmin}$ ($130 \\; \\mbox{kpc}$, encompassing the faintest traces of the galaxy's stellar envelope), our AAOmega intracluster field extends east of NGC 1316 almost $2^\\circ$ ($800 \\; \\mbox{kpc}$). Above the relatively bright limit ($M_{b_J}<-11.7$) of our survey, we found no evidence of dispersed intermediate-age stellar superclusters emanating from the NGC 1316 galaxy merger. While observing conditions restricted our detection limit to the most luminous CSSs, we would have expected to find any ejected stellar superclusters since their initial luminosity ($M_B\\simeq-15$) should still be above our detection limit after $3 \\; \\mbox{Gyr}$ \\citep{Fellhauer..2002}. We conclude that there is no evidence such objects have been dispersed outside the innermost region of the NGC 1316 galaxy merger or into intracluster space.  The galaxy NGC 4438, which has been strongly distorted by interaction with NGC 4435, is marked in Fig.~\\ref{fig:aaovirgofornax_11} at the edge of our M87 field. In our survey we found no overdensity of confirmed foreground or cluster objects that might be associated with these galaxies."
    },
    "0806/0806.4937_arXiv.txt": {
        "abstract": "{} {51 Oph is one of the few young Be stars displaying a strong CO overtone emission at $2.3$ microns in addition to the near infrared excess commonly observed in this type of stars. In this paper we first aim to locate the CO bandheads emitting region. Then, we compare its position with respect to the region emitting the near infrared continuum.} {We have observed 51 Oph  with AMBER in low spectral resolution (R=35), and in medium spectral resolution (R=1500) centered on the CO bandheads. } {The medium resolution AMBER observations clearly resolve the CO bandheads. Both the CO bandheads and continuum emissions are spatially resolved by the interferometer. Using simple analytical ring models to interpret the measured visibilities, we find that  the CO bandheads emission region is compact, located at $0.15_{-0.04}^{0.07}$AU from the star, and that the adjacent continuum is coming from a region further away ($0.25_{-0.03}^{0.06}$AU).  These results confirm the commonly invoked scenario in which the CO bandheads originate in a dust free hot gaseous disk. Furthermore, the continuum emitting region is closer to the star than the dust sublimation radius (by at least a factor two) and we suggest that hot gas inside the dust sublimation radius significantly contributes to the observed 2 $\\mu$m continuum emission.} {} ",
        "introduction": "The environment of young stars is formed by a dusty and gaseous circumstellar disk in which the first stages of planet formation are thought to take place. Hence, building a comprehensive picture of the physical mechanisms governing this formation requires the understanding of the distribution and the evolution of both the dust and gas species. The content and dynamics of the gaseous component, which dominates the total mass of the disk, control the accretion-ejection processes and will define the final architecture of forming planetary systems. On the other hand, the dust, responsible for the continuum infrared excess, plays an essential role in determining the structure of the protoplanetary disk itself, since its opacity dominates in the continuum. Dust grains are also the first building blocks that will grow into kilometric planetesimals and then into eventual planets. \\\\ To study the planet formation taking place in the very inner region of the disk, one needs to combine milli-arcsecond high {\\it angular} resolution techniques with high {\\it spectral} resolution in order to {\\it directly} probe the physical processes at stake both in the continuum and in the emission lines. In this respect, near infrared spectro-interferometry that associates long baseline near infrared interferometry with spectroscopy appears perfectly suited. This technique indeed offers a unique way to spatially resolve both the continuum infrared excess and emission line components, that is to locate the regions of emission independently of any prior assumption regarding the physical mechanisms at play. In that sense, infrared spectro-interferometry has already provided important results about the origin of the hydrogen line $\\mathrm{Br}\\gamma$. Different authors have thus located the $\\mathrm{Br}\\gamma$ emission region, showing that it can arise either from the accreting columns of gas falling onto the star \\citep{eisner_1} or from outflowing winds in which ionized matter is ejected from the star \\citep{malbet_1, tatulli_2, kraus_1}. \\\\ In this paper we are interested in exploring lines which are more direct probes of gas in the disk itself. We present the first interferometric observations of the $2.3\\mu\\mathrm{m}$ CO overtone emission in the (B9) Be star 51 Oph. 51 Oph is one of the very few young stars where this emission is strong enough to be observed with infrared spectro-interferometry \\citep{thi_1,berthoud_1}. This star exhibits also a strong $\\mathrm{Br}\\gamma$ emission, which, assuming a magnetospheric origin \\citep{muzerolle_1}, indicates an accretion rate of $\\dot{M}=1-2.10^{-7}M_{\\odot}/yr$ \\citep{garcialopez_1, brittain_1}. The dust component  is rather compact (\\citet{lagage_1, liu_1}) and shows a strong silicate feature in emission \\citep{bouwman_1} but no PAH \\citep{meeus_1}, consistent with a flat and geometrically thin dusty disk \\citep{malfait_1}. All the previous analysis are pointing towards a disk seen nearly edge-on. ",
        "conclusions": "\\subsection{Sizes of emitting regions vs. modelling} The uniform brightness and the narrow ring are the two extreme cases of a unique model of uniform distribution in which the width of the emitting region is varying. We emphasize that this geometrical modelling defines the metric that allows us to unequivocally link the measured visibilities to the spatial brightness distribution of these regions, through the Zernicke-Van Cittert theorem \\citep{goodman_1}. It does not introduce assumptions regarding the mechanisms at the origin of these emissions. We present in Table \\ref{tab_res} the results for the two scenarios. All the intermediate models would also work, with intermediate values for the sizes. Note that the uniform brightness returns the largest size because emission is added at small spatial scales down to the stellar radius. The uniform brightness  has no proper physical meaning since neither the dust nor the CO gas can survive so close to the star. However, it represents the {\\it characteristic size} of the emission region. As this quantity is commonly used in infrared interferometry, it is given here for this purpose. Conversely the narrow ring model is sensitive to the location where most of the emission is arising from, usually the {\\it inner rim} of this emission. In the following, we will analyze our results in the light of the narrow ring only. We emphasize that the size of the CO emitting region is {\\it smaller than the one of the continuum}, independently of the choice of our model. \\subsection{Inclination and position angle of the system} We find a position angle of $126_{-5}^{+15}$ degrees and an inclination of $i=82_{-15}^{+8}$ degrees. This inclination is in agreement with most of the previous studies. \\citet{berthoud_1} have, however, proposed an alternative scenario in which the object would have an inclination lower than $36^{\\circ}$. Our measurements are positively ruling out this latter option. \\subsection{Locating the CO bandheads emission region} For the first time, we are spatially resolving the CO bandheads emission region in 51 Oph. We find that most of the CO bandheads are emitted at a radius of $R_l =0.15_{-0.04}^{0.07}$AU from the central star. This is fully supporting the scenario in which  the CO bandheads emission is arising from the inner region of a hot gaseous disk in Keplerian rotation, between  $0.15$AU and $0.35-0.53$AU, as derived from spectroscopic modelling of the CO bandheads \\citep{thi_1, berthoud_1}, and where the temperature is hot enough ($\\ge 2000$K) to excite the first overtone bands at 2.3$\\mu$m. More precisely, these authors, though they are fitting the CO bandheads with a broad ring, are pointing out that most of the emission comes from a region close to the inner radius of their gaseous disk model at $\\sim$0.15AU. This is in excellent agreement with the value we derive in the present analysis. We note that one single ring enables us to reproduce the visibility in the two bandheads, indicating that both the 2-0 and 3-1 bands come from the same physical region. We also note that the CO emission region is inside the near infrared continuum one, the latter being located at a radius of $R_c = 0.24_{-0.03}^{+0.06}$AU. This result corroborates that the CO is emitted from a dust free region, regardless of the true origin of the continuum (see Sect. \\ref{subsec_cont}). At such a distance from the star, the survival of the CO gas from photo-dissociation must however be considered. \\citet{thi_1} suggested that two mechanisms can circumvent the photo-dissociation process: the self-shielding of the CO molecules that occurs if the column density is high enough ($N(CO)>10^{15}\\mathrm{cm}^{-2}$, \\citet{dishoeck_1}), as well as the C + OH  $\\rightarrow$ CO + H chemical reaction, which insures a continuous replenishment of this molecule. \\subsection{The origin of the continuum near-infrared excess} \\label{subsec_cont} \\begin{figure}[!t] \\centering \\includegraphics[width=0.4\\textwidth]{sedfit.pdf} \\caption{Spectral energy distribution of 51 Oph taken from literature \\citep{waters_1} (crosses) and superimposed best fit model (solid line) using MCFOST dusty disk model. The main parameters that are fitted are the inner radius $R_{dust}=1.2$AU, the outer radius $R_{out}=400$AU scale height at 1AU $H_0 = 0.04$, the flaring exponent $\\beta=1.1$, the grain size distribution (see text), the surface density exponent $p=-1.25$ and the total mass of dust $M_{dust} = 10^{-8}M_{\\odot}$. We refer to \\citet{pinte_1} for a thorough description of the modelling.}\\label{fig_results_sed_1}% \\end{figure} \\begin{figure}[!t] \\centering \\includegraphics[width=0.45\\textwidth]{sedvis_2.pdf} \\caption{MR-K continuum visibilities (top panel) and LR-K visibilities as a function of the baseline and  and overplotted MCFOST model arising from SED fitting (lines). The color code is the same than in fig. \\ref{fig_results_MRK_alone} and \\ref{fig_results_LRK_alone}. We can see that this model, with inner radius equal to the dust sublimation radius, is incompatible with the interferometric data, with a resulting $\\chi^2$ of $\\sim$10.}\\label{fig_results_sed_2}% \\end{figure} The continuum excess of 51 Oph is weak and, according to the shape of its SED \\citep{waters_1},  appears to come from a small and tenuous disk of dust which is optically thin at most wavelengths, as already inferred by \\citet{malfait_1}. We performed SED fitting with the MCFOST dusty disk code of \\citet{pinte_1}, and found the same results with a low amount of dust present in the disk ($M_{dust} \\sim 10^{-8}M_{\\odot}$). The result of our modelling is shown in Fig. \\ref{fig_results_sed_1}. This model however requires a dust inner edge located at the dust sublimation radius, which in the case of 51 Oph would be at $R_{dust} \\sim 1.2$AU for a typical grain size distribution of spherical particles\\footnote{${\\rm d}n(a) \\propto a^{-3.7} ~{\\rm d}a$, ranging from $a_{\\rm min} = 0.03~\\mu$m to $a_{\\rm max} = 1~$mm, with optical constants from \\citet{mathis_1}} and a dust sublimation temperature of $T_{dust} \\sim 1500K$. Obviously this model, with such a large inner rim, is not compatible with our measured visibilities as shown by Fig. \\ref{fig_results_sed_2}, and for which we calculate a $\\chi^2$ of $\\sim$10.  Studying the extreme case where only big grains are present would shrink the dust sublimation radius down to $R_{dust} \\sim 0.56$AU, still at least two times further out than the distance suggested by our data. Interestingly, a similar problem was evidenced by \\citet{leinert_1} on the same star, the 0.5AU mid-infrared size of 51 Oph derived from their MIDI measurements being too small to be compatible with the shape of the SED. \\\\ Several scenarios may however circumvent this apparent disagreement of the location of the dust evaporation radius with respect to our derived inner edge position. First, the gas inside the dust sublimation may  help the dust to survive closer in by absorbing a fraction of the stellar radiation available to heat the dust.  According to \\citet{muzerolle_1}, the gas accreting onto the star becomes optically thick to stellar radiation for accretion rates higher than $\\dot{M} \\sim 10^{-7}M_{\\odot}/yr$. The accretion rate of $1-2.10^{-7}M_{\\odot}/yr$ derived  from $\\mathrm{Br}\\gamma$ luminosity would actually place the gas in 51 Oph at the frontier between the optically thin and optically thick regimes. \\\\ Another possibility comes from the rapidly rotating nature of 51 Oph. Given its measured rotational velocity $v\\sin{i}=270$km/s \\citep{dunkin_1}, and knowing that this star is seen almost edge-on, 51 Oph would rotate at $90\\%$ of its critical velocity given by $v_{crit}=\\sqrt{2GM_{\\ast}/(3R_{\\ast})} \\simeq 300$km/s. Therefore, 51 Oph is likely elongated with a drop of its gravity ($g_{eff}$) from pole to equator, and subsequently of its effective temperature \\citep{zeipel_1}. The star being cooler at the equator than at the poles, the dust distributed in the equatorial plane will be heated less efficiently and the evaporation radius moved closer to the star. From the gravity darkening law $T_{eff}^4 \\propto g_{eff}$ (solid body approximation), the effective temperature of 51 Oph at the equator can be estimated to be of the order of $75\\%$ of that of the pole, hence lowering the sublimation radius by a factor of 2.\\\\ Finally, one last explanation would be that the infrared excess is not entirely originating from dust thermal emission. The gas inside the dust sublimation radius could substantially contributes to the near infrared energy balance. \\citet{muzerolle_1} have shown that inner optically thick gaseous zones are indeed expected to emit a large continuum excess  in the near infrared through free-free emission. Interestingly, this effect has been recently shown to take place in Herbig Ae/Be stars \\citep{eisner_2, isella_2, tannirkulam_1}. \\subsection{51 Oph: a classical Be star?} 51 Oph appears to be a peculiar source in an unusual transitional state. In the frame of Herbig Ae/Be stars, its SED presents a near infrared excess which is too small to account for a classical puffed-up inner rim.  In the frame of $\\beta$-Pic like stars, 51 Oph is also lacking the far infrared-excess bump associated with the presence of an outer dusty disk, the inner disk being emptied  by a potential forming planet \\citep{malfait_1}. Instead these authors suggest that 51 Oph is undergoing an alternative evolution scheme, without forming planets. Furthermore, the presence of strong CO overtone emission bandheads also makes 51 Oph quite a puzzling case. Such emission requires large column densities of warm  gas in order to produce detectable emission. Such large column densities are rare except in sources with the largest accretion rates \\citep{najita_1}. Interestingly, the CO overtone has also been detected in the B9 star HD58647 \\citep{berthoud_2}, for which the SED profile is similar to that of 51 Oph \\citep{malfait_1}. In HD58647 however, there is less CO emission but its infrared excess is stronger. \\citet{berthoud_2} thus concluded that both stars are most likely classical Be stars surrounded by massive gaseous disk, though seen at different evolutionary stages. This scenario is in agreement with their high rotational velocities ($v\\sin{i}=270$km/s and $118$km/s for 51 Oph and HD58647 respectively) and gives credit to the hypothesis that most of the near infrared continuum  emission -- if not all -- is arising from the circumstellar gas. HD58647 has been observed with the Keck interferometer \\citep{monnier_1} and these authors have derived a rather compact size of $0.4$AU for its near infrared continuum emission, slightly larger as that of 51 Oph for an equivalent luminosity ($L=260L_{\\odot}$ for 51 Oph vs. $L=250L_{\\odot}$ for HD58647). This is again below -- or at the very lower limit of -- the dust sublimation radius, pointing towards the same gaseous origin for the continuum than that of 51 Oph. Hence, the very inner environment of this type of stars seems to follow an intriguing evolution scheme where the dust  progressively dissipates leaving behind a massive gas-rich, strongly accreting disk. As the hot dust is vanishing, the relative contribution of the gas to the continuum infrared excess increases, its region of emission moving closer to the star. How the circumstellar dust disappears remains rather unclear. Repeated studies combining spectroscopic detection of emission bandheads (CO, hydrogen) and interferometric measurements in a large sample of stars would certainly help improve our understanding on how these disks are evolving and dissipating."
    },
    "0806/0806.1515_arXiv.txt": {
        "abstract": "}[2]{{\\footnotesize\\begin{center}ABSTRACT\\end{center} \\vspace{1mm}\\par#1\\par \\noindent {~}{\\it #2}}}  \\newcommand{\\TabCap}[2]{\\begin{center}\\parbox[t]{#1}{\\begin{center} \\small {\\spaceskip 2pt plus 1pt minus 1pt T a b l e} \\refstepcounter{table}\\thetable \\\\[2mm] \\footnotesize #2 \\end{center}}\\end{center}}  \\newcommand{\\TableSep}[2]{\\begin{table}[p]\\vspace{#1} \\TabCap{#2}\\end{table}}  \\newcommand{\\FigCap}[1]{\\footnotesize\\par\\noindent Fig.\\  % \\refstepcounter{figure}\\thefigure. #1\\par}  \\newcommand{\\TableFont}{\\footnotesize} \\newcommand{\\TableFontIt}{\\ttit} \\newcommand{\\SetTableFont}[1]{\\renewcommand{\\TableFont}{#1}}  \\newcommand{\\MakeTable}[4]{\\begin{table}[htb]\\TabCap{#2}{#3} \\begin{center} \\TableFont \\begin{tabular}{#1} #4 \\end{tabular}\\end{center}\\end{table}}  \\newcommand{\\MakeTableSep}[4]{\\begin{table}[p]\\TabCap{#2}{#3} \\begin{center} \\TableFont \\begin{tabular}{#1} #4 \\end{tabular}\\end{center}\\end{table}}  \\newenvironment{references}% { \\footnotesize \\frenchspacing \\renewcommand{\\thesection}{} \\renewcommand{\\in}{{\\rm in }} \\renewcommand{\\AA}{Astron.\\ Astrophys.} \\newcommand{\\AAS}{Astron.~Astrophys.~Suppl.~Ser.} \\newcommand{\\ApJ}{Astrophys.\\ J.} \\newcommand{\\ApJS}{Astrophys.\\ J.~Suppl.~Ser.} \\newcommand{\\ApJL}{Astrophys.\\ J.~Letters} \\newcommand{\\AJ}{Astron.\\ J.} \\newcommand{\\IBVS}{IBVS} \\newcommand{\\PASP}{P.A.S.P.} \\newcommand{\\Acta}{Acta Astron.} \\newcommand{\\MNRAS}{MNRAS} \\renewcommand{\\and}{{\\rm and }} { We have surveyed a $6\\zdot\\arcm5 \\times 6\\zdot\\arcm5$ field centered on the globular cluster M56 (NGC 6779) in search for variable stars. We have detected seven variables, among which two objects are new identifications. One of the new variables is an RR Lyrae star, the third such star in M56. Comparison of the new observations and old photometric data for an RV Tauri variable V6 indicates a likely period change in the star. Its slow and negative rate of $-0.005\\pm0.003$~d/yr would disagree with post-AGB evolution, however this could be a result of blue-loop evolution and/or random fluctuations of the period. } {Hertzsprung-Russell (HR) and C-M diagrams -- Stars: variables : BL Her, RV Tau, RR Lyr -- open clusters and associations: individual: M56 (NGC 6779)} ",
        "introduction": "CURious Variables Experiment (CURVE) is a long-term project focused on observations of open clusters, globular clusters and cataclysmic variable stars in the northern hemisphere (Olech \\etal 2003, Olech \\etal 2007, Rutkowski \\etal 2007). In stellar clusters we principally search for variable objects. However, our data also allows us to estimate basic parameters of observed clusters, such as distances and ages (Pietrukowicz \\etal 2006).  The globular cluster M56 (NGC 6779) is located in a rather dense galactic field at $(l,b)=(62\\zdot\\arcd66,+8\\zdot\\arcd34)$. The most recent deep $BVRI$ photometry of the cluster was obtained by Hatzidimitriou \\etal (2004) using 1.3-m telescope at Skinakas Observatory, in Crete. They estimate the distance modulus and the reddening for M56 of $(m-M)_V=15.62\\pm0.26$ and $E(B-V)=0.32\\pm0.02$, respectively. The authors also demonstrate that M56 is one of the most metal-poor ([Fe/H]$_{CG}=-2.00\\pm0.21$ on the scale proposed by Carretta and Gratton 1997) and one of the oldest globular clusters in the Galactic halo (13 Gyrs, using the age-index calibration of Salaris and Weiss 2002).  Despite the very early discovery of the first variable star in the globular cluster M56 (the object classified now as V3, Davis 1917) and excellent position for northern hemisphere observers of the cluster, identification of its variables has proceeded very slowly. Clement \\etal (2001) lists only 12 variable stars in M56, but five of them are very likely field objects. In this contribution we present results of the search for variable stars in M56 based on new data and with the use of image subtraction method, which works much better in crowded fields than classical photometry. ",
        "conclusions": "We have presented the results of a search for variable stars in the globular cluster M56. Besides five already known variables we have identified two new objects: V13 and V14. The object V13 has a period of 38.96 days and probably is a pulsating AGB star. The variable V14 is very likely an RR Lyrae star which belongs to the cluster, the third such object in M56, after V4 and V12. The number of RR Lyrae stars in this metal-poor globular cluster seems to be very small, but there are known clusters with similar characteristics. For example, in M30 of metallicity [Fe/H]$_{CG}=-2.17\\pm0.08$ (Carretta 2003) there have been found only 5 such variables (Clement \\etal 2001, Pietrukowicz and Kaluzny 2004); in NGC 6397 of metallicity [Fe/H]$_{CG}=-2.03\\pm0.05$ (Gratton \\etal 2003) there is no known RR Lyrae star at all (Kaluzny \\etal 2006).  For previously known variables in M56 we have confirmed positive period changes in BL Her variable V1 and semi-regular nature of V3 and V5. For variable V6, the RV Tau star, we have found, for the first time, very likely period change in this star. The negative period change rate of $-0.005\\pm0.003$ d/yr seems to be in contradiction to the evolutionary status of RV Tau stars as post-AGB objects, but not with blue-loop evolution. Numerous studies of period changes in RV Tau (\\eg Percy et al. 1997, Percy and Coffey 2005) also show that $O-C$ diagrams are dominated by random cycle-to-cycle period fluctuations of typically 0.005 to 0.02 of a period. The fluctuations may mask real evolutionary period changes. Moreover, the interpretation of the diagrams depends on the specific interval involved.  The results presented here have improved our knowledge on variable stars in the globular cluster M56, but future searches will require a bigger telescope (a 1-m or 2-m class telescope) at an observatory located in a place with better seeing conditions.  \\Acknow{ The authors would like to thank Dr. W. Pych for providing some useful software which was used in the analysis. PP and AO acknowledge support from the Domestic Grant for Young Scientists of the Foundation for Polish Science and Polish MNiI grant N203 301 335, respectively. Telescope operation was supported by the BW grant to Warsaw University Observatory. }"
    },
    "0806/0806.0725_arXiv.txt": {
        "abstract": "{ We infer from different seismic observations the energy supplied per unit of time by turbulent convection to the acoustic modes of {\\acenA} (HD 128620), a star which is similar but not identical to the Sun. The inferred rates of energy supplied to the modes ({\\it i.e.} mode excitation rates) are found to be significantly larger than in the Sun. They are compared with those computed with an excitation model that includes two sources of driving, the Reynolds stress contribution and the advection of entropy fluctuations. The model also  uses a  closure model, the Closure Model with Plumes (CMP hereafter), that takes the asymmetry between the up- and down-flows ({\\it i.e.} the granules and plumes, respectively) into account. Different prescriptions for the eddy-time correlation function are also confronted to observational data. Calculations based on a Gaussian eddy-time correlation underestimate excitation rates compared with the values derived from observations for {\\acenA}. On the other hand, calculations based on a Lorentzian eddy-time correlation  lie within the observational error bars. This confirms results obtained in the solar case. With respect to the helioseismic data, those obtained for {\\acenA} constitute  an additional support for our model of excitation. We show that mode masses must be computed taking  turbulent pressure into account. Finally, we emphasize the need for more accurate seismic measurements in order to discriminate, in the case of  {\\acenA}, between the CMP closure model and the quasi-Normal Approximation as well as to confirm or not the need to include the excitation by the entropy fluctuations. } ",
        "introduction": "{\\acenA} is the most promising star after the Sun for constraining the modelling of $p$-mode excitation by turbulent convection. Indeed, due to its proximity and its binarity, the fundamental parameters of {\\acenA} (effective temperature, luminosity, metallicity, gravity, radius) are quite well known. For this reason this star and its companion ($\\alpha$ Cen B) have been extensively  studied (see for instance the most recent modelling by \\citet{Miglio05} and the references therein). As pointed out recently by \\citet{Samadi07b}, the detection of $p$-modes and the measurement of their amplitudes as well as their mode line-widths ({\\it i.e.} lifetime) from {\\acenA} enable to derive the rates at which energy is supplied to the acoustic modes for this star. These observational constraints   can then  be used to check models of $p$-mode excitation by turbulent convection.  Such comparisons have been first undertaken in the case of the Sun by different authors \\citep[see  the recent review by][]{Houdek06}. They enable to test different models of stochastic excitation of acoustic modes as well as different models of turbulent convection \\citep[see eg.][]{Samadi05b}. Among those theoretical prescriptions, we consider that of \\cite{Samadi00I} with the improvements proposed by \\cite{Samadi02II} and \\cite{Kevin06b}. It was shown by \\cite{Samadi02II} that the way the eddy time \\emph{correlation}  is modelled plays an important role on the efficiency of excitation. Indeed, calculations of the mode excitation rates, $\\mathcal P$, that use a Lorentzian eddy-time correlation function better reproduce helioseismic data than those using a Gaussian one. In addition, \\cite{Kevin06b}, in the case of the Sun, showed that excitation rates computed using an adapted closure model that takes the presence of plumes into account reproduce much better the solar observations than the calculations based on the classical Quasi-Normal Approximation \\citep{Million41}.    An alternative theoretical model of the excitation of acoustic modes by turbulent convection  proposed by \\citet{Chaplin05} differs from that by \\citet{Samadi00I} in several ways: it does not take the driving by the advection of the entropy fluctuations by the velocity field into account. They only use the classical Quasi-Normal Approximation. More importantly, these authors claim that a Gaussian eddy-time correlation function reproduces better than a Lorentzian one the frequency dependence of mode excitation rates inferred from helioseismic data. However, they are led to introduce in their model a factor by which they multiply their formulation in order to reproduce the maximum of the  solar mode excitation rates.   {\\acenA}  provides a second opportunity to test various assumptions in the modelling of the p-mode excitation: the  amplitudes of the  acoustic modes detected in {\\acenA} were derived by \\cite{Butler04} using spectrometric data. From those data,  an estimate of the averaged mode line-widths has been first proposed by \\cite{Bedding04} and more recently updated in \\cite{Kjeldsen05}. Using a different method and data from the WIRE satellite, \\cite{Fletcher06} proposed a new estimate of the averaged mode line-widths that differ significantly from the one derived by \\cite{Kjeldsen05}. Indeed, the two data sets place the mode life time between 2.2 days \\citep{Kjeldsen05} and 3.9 days \\citep{Fletcher06}. For comparison, the averaged mode life time derived for the Sun by \\cite{Bedding04} in a similar way as for {\\acenA} by \\cite{Kjeldsen05} is about two days.  \\citet{Samadi07b} have inferred from those sets of seismic constraints  the p-mode excitation rates $\\mathcal P$. They have found that they are significantly  larger than those associated with the solar p-modes. Furthermore, $\\mathcal P$ peaks  in the frequency domain $~\\sim $ 2.2  --  2.6~mHz  while it peaks at the frequency $\\nu_{\\rm max} \\sim $~3.8 mHz in the case of the Sun.   Although the spectroscopic characteristics ($T_{\\rm eff}=5810$~K, $\\log~g =$~4.305) of {\\acenA}  are close to that of the Sun ($T_{\\rm eff}=5780$~K, $\\log~g =$~4.438), the seismic signatures are quite different. Consequently, finding agreement between predicted and observed excitation rates would be a non-trivial result, providing additional support for the theory.   A preliminary comparison with theoretical calculations obtained in the manner of \\citet{Kevin06a} was carried out  by \\citet{Samadi07b}. Discrepancies between the excitation rates inferred from the observations and the theoretical calculations were found. We update here this study by proceeding in a similar way as \\citet{Rosenthal99}. Indeed, these authors have built a solar 1D model where the surface layers are taken directly from a  fully compressible 3D hydrodynamical numerical  model. We will refer here to such a 1D model as a ``patched'' model. \\citet{Rosenthal99} have obtained a much better agreement between observed and theoretical eigenfrequencies of the Sun computed for such a ``patched'' 1D model than those obtained for a ``standard'' 1D model based on the standard mixing-length theory with no turbulent pressure included. Following \\citet{Rosenthal99}, we build here such a ``patched'' model to derive  adiabatic mode radial eigen-displacements ($\\xi_r$) and mode inertia ($I$). We use them to compute the mode excitation rates, which we compare with excitation rates computed using  $\\xi_r$ and $I$ obtained with a ``standard'' 1D model.   The paper is organized as follows: in Sect.~\\ref{modelling} we describe our procedure to compute the  mode excitation  rates  for the specific case of {\\acenA}. We then describe in Sect.~\\ref{constraints} the way the mode excitation rates are inferred from seismic observations of {\\acenA}. In Sect.~\\ref{comparison}, we compare  theoretical calculations of $\\mathcal P$  with those  inferred from the seismic data obtained for {\\acenA}. We compare and explain in Sect~\\ref{differences} the differences between {\\acenA} and the Sun. Finally, Sect.~\\ref{discussion} and Sect.~\\ref{conclusion} are devoted to the discussion and conclusions, respectively. ",
        "conclusions": "  \\subsection{Differences with the Sun}  Although  {\\acenA} has an effective temperature very close to that of the Sun, we find here that  the p-mode excitation rates $\\mathcal P$ inferred from the seismic constraints obtained for {\\acenA} are  about two times larger than in the Sun. These differences are  attributed to the fact that the eddies in {\\acenA} have a larger characteristic size ($\\Lambda$) than in the Sun. This is related to the fact that {\\acenA} has a smaller surface gravity.   Furthermore, the p-mode excitation rates for {\\acenA} are maximum at lower frequencies than in the Sun. This  behaviour is related to the fact that the eddies have a longer turn-over time as a result of a larger $\\Lambda$.  The seismic characteristics of the p-modes detected in {\\acenA} significantly differ from that of the Sun. They can therefore provide additional constraints on the model of stochastic excitation.      \\subsection{Inferred versus modelled  excitation rates}  Our modelling gives rise to excitation rates within the error bars associated with the observational constraints. We stress that this modelling was undertaken for {\\acenA} {\\it independently} from the solar case, {\\it i.e.} without  using a formulation fitted on the helioseismic data  as it is the case, for instance, in the case of the Sun in \\citet{Chaplin05} or in the case of {\\acenA} in \\citet{Houdek02b}. The seismic constraints from {\\acenA} then provide a clear validation of the basic underlying physical assumptions included in the theoretical model of stochastic excitation, at least for stars not too different from the Sun.   \\subsection{Constraints on the description of turbulence: eddy-time correlation}  We find that our theoretical estimations of $\\mathcal P$,  which assume a Lorentzian eddy-time correlation function ($\\chi_k$) and the Closure Model with Plumes (CMP) proposed by \\citet{Kevin06a}, lie  in the observed domain. On the other hand, when a Gaussian function  is chosen for $\\chi_k$, $\\mathcal P$ is significantly underestimated. The comparison with the seismic data for {\\acenA} confirms the results  for  the solar case obtained by  \\citet{Samadi02II} that  $\\chi_k$ significantly departs  from a Gaussian. As in  \\citet{Samadi02II}, we attribute the departure of $\\chi_k$ from a Gaussian to diving plumes ({\\it i.e.} down-flows), which are more turbulent than granules ({\\it i.e.} the up-flows). This result confirms  that a Lorentzian function is a more adequate description for the eddy-time correlation  than a Gaussian.   \\subsection{Constraints on the modelling of turbulent convection in the equilibrium stellar model}  Calculations involving eigenfunctions computed on the basis of a global 1D model that includes a realistic description of the outer layers of the star (taken from 3D simulations) reproduce  much better (see Fig.~\\ref{figA}) the seismic data than calculations that use eigenfunctions computed with a  standard stellar model built with  the mixing-length theory (MLT) and ignoring turbulent pressure. This is because a model that includes turbulent pressure results in larger mode masses ${\\cal M}$ than a model which ignores turbulent pressure. This can be understood as follows. Within the super-adiabatic region, a model that includes turbulent pressure provides an additional support against gravity and hence has a lower gas pressure and density (see Fig.~\\ref{figC}) than a model that does not include turbulent pressure. As a consequence,  mode inertia  (and hence mode masses) are then larger in a model that includes turbulent pressure.  These conclusion are similar to that obtained in the Sun. Indeed, the mode masses considered by \\citet{Kevin06b} in the case of the Sun were obtained with a 1D model computed using \\citet{Gough77}'s  non-local mixing-length formulation of convection. The model thus includes turbulent pressure. We do not observe significant differences between excitation rates obtained  with this  non-local model and those obtained with  a ``patched'' solar computed as described here in the case of {\\acenA}. On the other hand, excitation rates computed with mode masses obtained with a ``standard''  solar model (that is with no turbulent pressure included) or with a model in which turbulent pressure is included according to the mixing-length theory under-estimate significantly the helioseismic constraints.   These results tell that one must compute mode masses from 1D models that include turbulent pressure using a 3D hydrodynamical model or using a non-local description of convection.  \\subsection{Need for improved data sets}  As shown by \\citet{Samadi02II} in the case of the Sun, contribution of the entropy fluctuations to the excitation cannot be neglected. Furthermore, recently, \\citet{Kevin06b} have shown that theoretical calculations based on the CMP result in a better agreement with the helioseismic constraints than those based on the Quasi-Normal Approximation (QNA).  However, in the case of {\\acenA}, differences between theoretical calculations that use the CMP  and those based on the QNA (see Fig.~\\ref{figB}) as well as differences between calculations including  driving by  entropy fluctuations and those that do not include it (not shown), are of the same order as the observational uncertainties associated with the two data sets. The present seismic constraints  therefore are  unable to discriminate between these assumptions.  This emphasizes the need for more accurate seismic data for {\\acenA}."
    },
    "0806/0806.0380_arXiv.txt": {
        "abstract": "{Most upcoming CMB polarization experiments will use direct imaging to search for the primordial gravitational waves through the B-modes. Bolometric interferometry is an appealing alternative to direct imaging that combines the advantages of interferometry in terms of systematic effects handling and those of bolometric detectors in terms of sensitivity.} {We calculate the signal from a bolometric interferometer in order to investigate its sensitivity  to the Stokes parameters paying particular attention to the choice of the phase-shifting scheme applied to the input channels in order to modulate the signal.} {The signal is expressed as a linear combination of the Stokes parameter visibilities whose coefficients are functions of the phase-shifts.} {We show that the signal to noise ratio on the reconstructed visibilities can be maximized provided the fact that the phase-shifting scheme is chosen in a particular way called \"coherent summation of equivalent baselines\". As a result, a bolometric interferometer is competitive with an imager having the same number of horns, but only if the coherent summation of equivalent baselines is performed. We confirm our calculations using a Monte-Carlo simulation. We also discuss the impact of the uncertainties on the relative calibration between bolometers and propose a way to avoid this systematic effect.} {} ",
        "introduction": " ",
        "conclusions": ""
    },
    "0806/0806.2666_arXiv.txt": {
        "abstract": "We report results from a systematic study of X-ray emission from black hole transients in quiescence. In this state mass accretion is thought to follow the geometry of an outer optically thick, geometrically thin disc and an inner optically thin, geometrically thick radiatively inefficient accretion flow (RIAF). The inner flow is likely also coupled to the jets near the black hole that are often seen in such systems. The goal of the study is to see whether the X-ray emission in the quiescent state is mainly powered by the accretion flow or the jets. Using data from deep {\\it XMM-Newton} observations of selected black hole transients, we have found that the quiescent X-ray spectra are, to a high precision, of power-law shape in the cases of GRO~J1655$-$40 and V404~Cyg. Such spectra deviate significantly from the expected X-ray spectrum of the RIAF at very low accretion rates. On the other hand, they can naturally be explained by emission from the jets, if the emitting electrons follow a power-law spectral distribution (as is often assumed). The situation remains ambiguous in the case of XTE~J1550$-$564, due to the relatively poorer quality of the data. We discuss the implication of the results. ",
        "introduction": "The majority of X-ray binaries that are known to contain a stellar-mass black hole are transient X-ray sources. They spend most of their time in the quiescent state, in which the mass accretion rate is thought to be extremely low. Occasionally, they undergo an outburst during which they may become the brightest X-ray sources in the sky. The exact mechanism that triggers such an outburst is not entirely understood but is thought to be related to a sudden surge in the accretion rate that is caused by a thermal instability in the accretion disc (see reviews, e.g., by King 1995 and Lasota 2001). During an outburst, the X-ray properties of a black hole transient is often described empirically in terms of spectral states (e.g., McClintock \\& Remillard 2006; Xue, Wu, \\& Cui 2008).  It is proposed that the spectral states may correspond to different configurations of the underlying accretion process at different mass accretion rates (Narayan 1996; Narayan, McClintock \\& Yi 1996; Esin, McClintock, \\& Narayan 1997). Specifically, when the accretion rate is high in the high-soft state, the accretion flow is thought to follow the geometry of the Shakura \\& Sunyaev (1973) disk (SSD), which is geometrically thin and optically thick. This can naturally explain the observed blackbody-like X-ray spectrum that is characteristic of the high-soft state. As the accretion rate decreases, the source evolves towards the low-hard state. In the process, a phase transition is thought to occur in the inner portion of the disc, in which the accreted matter is heated to nearly local virial temperatures and may also form an outflow wind (see Narayan 2005 for a comprehensive review of the models, their evolution, and their applications to black hole candidates and active galactic nuclei). The accretion flows in this region is a geometrically thick but optically thin configuration, and is radiatively inefficient. The accretion power is mostly advected into the black hole or carried away by the outflow (Blandford \\& Begelman 1999). Such a radiatively inefficient accretion flow (RIAF) is capable of producing hard X-rays by up-scattering ambient soft photons. This can naturally explain the increasing dominance of the power-law component of the X-ray spectrum, as the source approaches the low-hard state. The RIAF model predicts that the trend continues towards even lower accretion rates, as more of the accretion flow becomes advection dominated, and that the X-rays in the quiescent state originate entirely from the RIAF (Narayan, McClintock, \\& Yi 1996).  A more recent development is the realization of the potentially critical role of jets, which seem to be ubiquitous in black hole transients (see review by Fender 2006 and references therein). Yuan, Cui \\& Narayan (2005) demonstrated that it would be nearly impossible for the RIAF model, which is quite successful in explaining the X-ray emission from black hole transients in the low-hard state, to also account for the observed emission at longer wavelengths (radio and IR in particular). In order to describe the broadband spectral energy distribution (SED) in the low-hard state, they showed that contributions from both accretion flow and jets would be needed, with the former mainly responsible for emission at UV/X-ray wavelengths, the latter for emission at radio/IR wavelengths, and both for emission in between (cf. Malzac, Merloni, \\& Fabian 2004, who argued that the optical emission might be dominated by the jets), when the accretion rates are relatively high. Extrapolating the result of Yuan et al. (2005) to lower accretion rates, Yuan \\& Cui (2005) predicted that the X-ray emission from the jets would eventually exceed that from the hot flow (because the former is proportional to the accretion rate $\\dot{m}$, normalized to the Eddington rate, and the latter roughly to $\\dot{m}^2$). In the quiescent state, the X-ray emission should, therefore, be mainly powered by the jets, at variance with the prediction of the RIAF model. Falcke, Kording \\& Markoff (2004) also postulated that the X-ray emission of the quiescent state would be dominated by the jets. But, in contrast to Yuan \\& Cui (2005), they argued that this would be the case even for the low-hard state.  Is it the accretion flow or jets that power the X-ray emission from black hole transients in the quiescent state? Yuan \\& Cui (2005) proposed two observational tests to answer the question. If the quiescent X-rays are powered by the jets, they predicted: (1) the radio/X-ray correlation would steepen when the X-ray flux drops below a characteristic value and (2) the X-ray spectrum would be of power-law shape. It has been claimed that the observation of A0620-00 is at odds with the first prediction (Gallo et al. 2006). However, the conclusion hinged critically on a radio/X-ray correlation that had been thought to hold for all black hole candidates. The universality of the radio/X-ray correlation has since been brought into question (Xue \\& Cui 2007). The second prediction is a viable test because the X-ray spectrum of an RIAF would deviate strongly from power-law shape at sufficiently low accretion rates, when the density of the flows becomes so low that Comptonization is dominated by single scattering (Narayan, McClintock, \\& Yi 1996; Quataert \\& Narayan 1999; McClintock et al. 2003; Yuan, Cui, \\& Narayan 2005).  In this work, we present results from a systematic study of black hole transients in quiescence. A number of such sources had been observed and detected earlier with {\\it Chandra} and {\\it XMM-Newton} (Kong et al. 2002; Hameury et al. 2003) but none of the X-ray spectra obtained are of sufficiently high quality that would allow us to distinguish jet-based and accretion-based models. To improve the situation, we carried out deep observations of selected sources with {\\it XMM-Newton}. The results reported here are based on data from these as well as an archival {\\it XMM-Newton} observation. ",
        "conclusions": "Fender et al. (2003) argued, on the basis of the ``universal radio/X-ray correlation'', that the energetics of the quiescent state ought to be dominated by the jets, in the sense that the kinetic power of the jet is much greater than the X-ray luminosity of the accretion flows. In their jet-dominated state, however, the X-ray luminosity of the jet is not necessarily also greater than that of the accretion flows, because the radiative power of the jet is only of the order of 1\\% of the kinetic power (see Yuan \\& Cui 2005 for a more detailed discussion). Here, we have shown that the quiescent X-rays from transient black hole candidates are likely to originate from the jets, as opposed to the accretion flows.  In summary, we have, for the first time, found direct evidence that the quiescent state may be fundamentally different from the low-hard state, as far as the source of X-ray emission is concerned. Contrary to the view that the former may be a simple extension of the latter towards lower accretion rates, our results suggest that the X-ray emission from transient black holes is dominated by contribution from the jets (or other sources of non-thermal electrons) in the quiescent state, while in the low-hard state it is likely dominated by contribution from the RIAF (e.g., Esin, McClintock, \\& Narayan 1997; Yuan et al. 2007).  Finally, we would like to emphasize that the prediction of Yuan \\& Cui (2005) is insensitive to the mass of black holes and might thus also hold for active galactic nuclei (AGN). Wu, Yuan \\& Cao (2007) has recently modeled a sample composed of eight FR I galaxies and found that their X-ray spectra should be dominated by jets, rather than by RIAFs, if their luminosities are below $\\sim 10^{-6}L_{\\rm Edd}$, and vice versa, as predicted by Yuan \\& Cui (2005). Wrobel, Terashima, \\& Ho (2008) observed two low-luminosity AGN at 8.5 GHz and found that the observed radio luminosity is within a factor of 3 of the value that is predicted from the observed X-ray luminosity and the radio--X-ray--mass relation derived by Yuan \\& Cui (2005)."
    },
    "0806/0806.0624_arXiv.txt": {
        "abstract": "We consider a potentially new class of gravitational wave sources consisting of a white dwarf coalescing into a massive black hole in the mass range $\\sim 10^4-10^5\\,\\msun$. These sources are of particular interest because the gravitational wave signal produced during the inspiral phase can be detected by the {\\it Laser Interferometer Space Antenna} ({\\it LISA}) and is promptly followed, in an extended portion of the black hole and white dwarf mass parameter space, by an electro-magnetic signal generated by the tidal disruption of the star, detectable with X-ray, optical and UV telescopes. This class of sources could therefore yield a considerable number of scientific payoffs, that include precise cosmography at low redshift, demographics of black holes in the mass range $\\sim 10^4 - 10^5\\Ms$, insights into dynamical interactions and populations of white dwarfs in the cores of dwarf galaxies, as well as a new probe into the structure and equation of state of white dwarfs. By modelling the gravitational and electromagnetic radiation produced by these events, we find them detectable in both observational windows at a distance $\\approx 200$ Mpc, and possibly beyond for selected regions of the parameter space. We also estimate the detection rate for a number of model assumptions about black hole and white dwarf mass functions and dynamical interactions:  the rate is (not surprisingly) highly uncertain, ranging from $\\sim 0.01\\,\\mathrm{yr}^{-1}$ to $\\sim 100\\,\\mathrm{yr}^{-1}$. This is due to the current limited theoretical understanding and minimal observational constraints for these objects and processes. However, capture rate scaling arguments favor the high end of the above range, making likely the detection of several events during the {\\it LISA} lifetime. ",
        "introduction": "The simultaneous detection of sources in both the electro-magnetic band - which provides a measurement of the source redshift, $z$ -- and the gravitational wave (GW) window -- which yields a direct determination of the luminosity distance $D_\\mathrm{L}$ to the source -- could revolutionize cosmography by determining the distance scale of the Universe in a precise, calibration-free way. This was pointed out initially by Schutz (1986) in the context of ground-based observations of GWs from coalescing compact binaries with the network of ground-based laser interferometers now in operation (Whitcomb 2008). The observational capability of space-based instruments such as the {\\it Laser Interferometer Space Antenna} ({\\it LISA}; Bender et al. 1998), which could observe many sources at high signal-to-noise ratio (SNR) and large redshift, has attracted much attention recently. Several scenarios have been considered, primarily related to the identification of the host galaxy or galaxy cluster of massive black hole (MBH) binary systems detected in GWs (Cutler 1998; Hughes 2002; Menou 2003; Vecchio 2004; Holz \\& Hughes 2005; Lang \\& Hughes 2006, 2008; Kocsis et al. 2006, 2007a,b; Arun et al. 2007; Cornish \\& Porter 2008; Trias \\& Sintes 2008), and the possible electro-magnetic signatures produced by the pre-glow/afterglow of the MBH mergers (Milosavljevi{\\'c} \\& Phinney 2005; Dotti et al. 2006). The main obstacles to such groundbreaking observations are either the possible paucity of sources likely to produce significant gravitational and electro-magnetic radiation detectable to cosmological distances and/or the rather poor angular resolution of GW instruments (e.g. Cutler 1998; Hughes 2002; Vecchio 2004; Lang \\& Hughes 2006; Arun et al 2007; Cornish \\& Porter 2008; Trias \\& Sintes 2008), which could inhibit the electro-magnetic identification of the host.  In this paper, we discuss a new class of GW sources that have received little attention so far (Menou, Haiman \\& Kocsis 2008): the inspiral of a white dwarf (WD) around a MBH in the mass range $\\sim 10^4 - 10^5\\Ms$ followed by the tidal disruption of the star before it plunges into the MBH. As we will show, these sources may be observable at low redshift (a few hundreds Mpc) with {\\it LISA} \\emph{and} their electro-magnetic emission may be detectable with X-ray observatories and optical ground based telescopes. From the GW point of view, a MBH-WD binary is a different flavour of the so-called Extreme-Mass Ratio Inspirals or EMRIs. Traditionally, the fiducial EMRI is taken to be a stellar-mass $\\sim 10\\,\\msun$ black hole orbiting a $10^6\\,\\msun$ MBH (Barack \\& Cutler 2004). The key difference between a \"traditional EMRI\" and a MBH-WD system considered in this paper is that for a range of MBH and WD masses, the inspiral does not proceed all the way until the compact object falls into the MBH horizon, but it terminates with the tidal disruption of the WD producing an electro-magnetic signature (for a MBH mass $\\simgt 3\\times 10^5\\,\\msun$, a WD survives throughout the whole inspiral and the system behaves just like a traditional EMRI). The observation of both gravitational and electro-magnetic signals from the same source provides a direct and calibration-free measurement of the $D_L(z)$ relationship and opens new avenues for cosmography and, more directly (due to the low redshift of most of the expected sources) a completely independent determination of the Hubble parameter $H_0$ that does not depend on any distance calibration. This class of sources can also provide new insights into a number of unanswered questions in relativistic astrophysics: (i) the demographics of MBHs in the mass range $\\sim 10^4 - 10^5\\Ms$ -- only a handful of MBH candidates with masses below $10^6\\msun$ (Greene \\& Ho 2004; Barth, Greene \\& Ho 2005) is known to date, their mass estimate is rather uncertain, since it is based on the emission-line spectra of the active nuclei, and none of them have masses below $10^5\\msun$ -- (ii) the populations of WDs and the dynamical processes that take place in the cores of dwarf galaxies, that are unknown and unconstrained by observations, and (iii) the structure and equation of state of WDs -- the exact point at which tidal disruption occurs indeed depends on the WD equation of state (e.g. Magorrian \\& Tremaine 1999), and the electro-magnetic signature carries information about the WD composition.  The paper is organized as follows: in Section~\\ref{s:GW} we identify the mass range of WDs and central MBHs that lead to the tidal disruption of the star before it plunges onto the black hole and we determine the volume of the Universe that {\\it LISA} will be able to survey; in Section~\\ref{s:rate} we derive the GW detection rate of these sources and discuss its uncertainties; in Section~\\ref{e:EM} we model the electro-magnetic counterpart to MBH-WD EMRIs; finally, in Section~\\ref{e:summary} we summarize the main results and our conclusions. ",
        "conclusions": "\\label{e:summary}  In this paper we have considered a class of {\\it LISA} sources that have not been explored so far, consisting of a WD inspiralling onto a MBH in the mass range $\\sim 10^4-10^5\\,\\msun$. These sources are of particular interest because the gravitational wave signal produced during the inspiral can be detected with {\\em LISA} and serves as a precursor for an electro-magnetic flash -- likely detectable in X-ray, and at optical and ultra-violet wavelengths -- generated during the tidal disruption of the star and the subsequent distinctive accretion episode, which takes place for a considerable region of the mass parameter space.  Observations of the same source in the gravitational and electromagnetic band enable an independent and calibration-free determination of the Hubble parameter at low redshift $z\\simlt 0.1$ (through the direct measurement of the redshift {\\em and} luminosity distance of the same source), studies of the faint end of the MBH mass function in (dwarf) galaxies, measurements of the mass distribution function of extra-galactic WDs and in particular of those in galaxy cores, and provide a probe of the structure and equation of state of WDs, and of dynamical processes in the cores of galaxies.  We have determined the range of WD and MBH masses that lead to a tidal disruption before the final plunge, hence producing an electro-magnetic counterpart to the GW emission.  In order to do so we have adopted a polytropic approximation of the WD mass-radius relation. In the future it is important to explore more sophisticated models to establish the dependence of the results presented here on this assumption. We have modelled the expected GW signal using \"analytical kludge\" waveforms {\\em a-la} Barack \\& Cutler and have explored the dependence of the SNR on selected values of the eccentricity and the MBH spin. We found that the maximum distance at which {\\it LISA} can observe these events depends both on the MBH spin and the orbital eccentricity, with circular prograde orbits around a highly spinning MBH producing signals observable to larger distances. A typical 0.5$\\msun$ WD orbiting a $5\\times10^4\\msun$ Schwarzschild MBH would be observable up to $200$ Mpc, assuming a detection threshold equivalent to an optimal SNR of 30 in the two combined noise orthogonal unequal-arm Michelson observables for 5 years of integration.  Based on current (rather uncertain) estimates of the MBH population and of the rates at which they capture WDs we have computed the LISA detection rate and discussed its uncertainties. Circular binaries yield a higher detection rate (by up to a factor of $\\sim 3$) with respect to eccentric ones, due to their higher SNR. A substantial population of highly spinning MBH would enhance the detection rate, since the last stable circular orbit of a WD inspiralling on a prograde orbit would be much smaller. More massive (i.e. observable to further distances) MBHs would enhance the detection rate. The rates that we obtain span almost four orders of magnitude, from $0.01$ to $100$ events per year and are listed in table \\ref{tab:1}. Scaling arguments for the stellar density in galactic centers suggest high WD disruption rates, up to 10$^{-6}$ yr$^{-1}$, in dwarf nuclei favoring the upper end of the detection rate range quoted above. Assuming a minimum lifetime of 5 yrs, {\\it LISA} will likely see several of these events.  We have also modeled the electro-magnetic signature associated to the WD disruption. Assuming disruption outside the last stable orbit, part of the WD material will settle in an accretion disk powering a luminous X-ray source that will be detectable up to $\\sim 200$ Mpc. A small fraction of the debris will become unbound forming a thin arc or annulus that expands at nearly the escape velocity.  Photoionization of this annulus by X-ray and UV photons produced by the accretion process will result in several different (depending on the WD composition) emission lines, with relative intensities changing on a timescale of a week as the debris expands. Such lines will be observable well beyond 100 Mpc with future 30m-class telescopes and provide a unique signature of this kind of events. If disruptions during the final plunge also result in observable electro-magnetic flares, the detection rate would increase by a factor of $\\sim 3$ (see table \\ref{tab:1}) considering Schwarzschild MBHs. However whether such a disruption would produce an observable signal is an open question that we have not tried to address here and deserves further investigation.  Based on the above results it is conceivable that several MBH-WD binary systems will be observed both in gravitational wave and electro-magnetic band at a distance up to a few hundreds of Mpc; however no events as well as hundreds of events are consistent with our current understanding of the key physical processes.  The need to observe MBH-WD binaries in the electro-magnetic and gravitational window to maximise the science return raises the issue of {\\em prompt alerts} generated by {\\it LISA} to electro-magnetic observatories. So far alerts have been discussed only in the context of observations of massive-black hole binary systems ({\\em e.g.} Kocsis et al 2006, 2007a,b; Lang \\& Hughes 2008), but the scenario that we have discussed in this paper calls for a systematic study regarding extreme-mass ratio-inspirals in the relevant mass range. Unfortunately, the complexity of searches for EMRIs and the still limited understanding and maturity of end-to-end algorithms (e.g. Babak et al. 2008; Gair et al. 2008a,b; Cornish 2008) prevent at present a realistic study of such an important problem."
    },
    "0806/0806.0554_arXiv.txt": {
        "abstract": "Because of the corotation, the polarization angle (PA) curve of a pulsar lags the intensity profile by $4r/\\rlc$ rad in pulse phase. I present a simple and short derivation of this delay-radius relation to show that it is not caused by the aberration (understood as the normal beaming effect) but purely by contribution of corotation to the electron acceleration in the observer's frame. Available altitude-dependent formulae for the PA curve are expressed through observables and emission altitude to make them immediately ready to use in radio data modelling. The analytical approximations for the altitude-dependent PA curve are compared with exact numerical results to show how they perform at large emission altitudes. I also discuss several possible explanations for the opposite-than-normal shift of PA curve, exhibited by the pedestal emission of B1929$+$10 and B0950$+$08. ",
        "introduction": "In the simplest model of pulsar polarization (Komesaroff 1970; Radhakrishnan \\& Cooke 1969, hereafter RC69) the position angle of polarization does not depend on the radial distance\\footnote{The quantity $r$ represents the distance measured from the \\emph{center} of the neutron star.} of radio emission $r$. The polarization angle (PA) becomes dependent on $r$ when dynamic effects of pulsar's rotation are taken into account (dragging of electrons by the corotating magnetic field). Blaskiewicz et al.~(1991) (hereafter BCW) showed that if the emission originates from a fixed radial distance $r$, the shape of the PA swing is (approximately) preserved (ie.~it is the same as in the case of negligible $r$), but the entire swing is shifted towards later phases by $\\dphbcw\\approx4r/\\rlc$ radians with respect to the center of the profile (where $\\rlc$ is the light cylinder radius). Following BCW, I will refer to this formula with the term `delay-radius' relation. Hibschmann \\& Arons (2001) (hereafer HA) have shown that the PA curve also undergoes vertical shifts, ie.~in the PA values. Both these results appear to have interesting observational consequences (eg.~Ramachandran \\& Kramer 2003; von Hoensbroech \\& Xilouris 1997).  In the case of phase-dependent emission altitude different parts of the PA curve undergo different shifts and the PA curve assumes a distorted shape. This effect regularly happens to be employed to model observed distortions of PA curves and to derive magnetospheric emission altitudes (eg.~Krishnamohan \\& Downs 1983, hereafter KD83; Xu, Qiao \\& Han 1997; Gil \\& Krawczyk 1997; Mitra \\& Seiradakis 2004). A tool that is needed for this is an analytical formula for the PA that explicitly depends on the radial distance of the emission region. BCW and HA provide various forms of this equation. However, their formulae are not in a ready-to-use form: they are expressed through the emission time instead of the pulse longitude $\\phobs$ (hereafter called pulse \\emph{phase}\\footnote{Throughout this paper the phase is assumed to be measured in radians whenever dimensionless terms are added to it.}). It is the need for this last step of the BCW's analysis that actually sparked writing of this paper.  A strict and formal description of the PA subject is given in the superb work of Hibschman \\& Arons (2001; see their appendices) and it will not be repeated in this paper. My intention here is to provide a simple reference for those who want to use the altitude-dependent PA curves in their data modelling. The aim is to clarify some obscure aspects of the subject by trivialising the formalism and to provide practical PA equations in their final form. Accordingly, Sect.~\\ref{simpderiv} presents a very simple and short derivation of the delay-radius relation to clearly expose its origin. In Sect.~\\ref{pacurv} I introduce the fiducial phase, describe the magnitudes and directions of various relativistic shifts with respect to it, and I write down the equations for the altitude-dependent PA curve in a form that is ready for immediate use in data modelling. In Sect.~\\ref{limits} I compare the analytical approximations to exact numerical results obtained for various emission altitudes to show the validity range of the BCW theory. In Sects.~\\ref{accuracy} and \\ref{noncurv} I discuss possible explanations for the opposite-than-expected shifts of PA curve, as exemplified by the pedestal radio emission components of B1929$+$10 and B0950$+$08. These and other interpretations of the anti-BCW shifts are summarized in  Sect.~\\ref{antibcw}. ",
        "conclusions": "\\label{abornot}  It is shown in Sect.~\\ref{simpderiv} that the magnitude of the shift given by the delay-radius relation has nothing to do with the aberration (understood as the normal beaming effect). On the other hand, Hibschman \\& Arons mention a phase shift of $3r/\\rlc$ (their eq.~4) and claim that it is just the aberration (``simple beaming\") that is responsible for the remaining $r/\\rlc$ part of the total delay. Do we have a contradiction here?  Not really, though the wording of HA can easily be misunderstood. Both in the BCW and HA, the delay-radius relation is derived as a difference of two moments of emission: 1) the moment when the radiation bearing the steepest-gradient PA is directed to the observer (this happens at $\\Omega t_e=3r/\\rlc$) and 2) an earlier moment when the radiation in the middle of the profile (emitted in the $\\om$-plane in the CF) becomes directed toward the observer. HA explicitly notice that the latter does \\emph{not} take place at the moment $t_e=0$, when the dipole axis is in the $\\vec \\Omega$-observer plane: instead, because of the aberration it happens at $\\Omega t_e=-r/\\rlc$ (see Fig.~\\ref{emigeom} and eq.~\\ref{phem} with $\\phem=0$). The key point that is not mentioned in their description, however, is that the moment $\\Omega t_e=3r/\\rlc$ at which the steepest gradient radiation is directed toward the observer has already been also advanced in time by the same aberration angle of $r/\\rlc$. Thus, \\emph{the aberration, understood as the normal beaming effect, works identically both at the center of the main pulse, as well as on its trailing side and contributes practically nothing to the shift of the PA curve with respect to the intensity profile.}  Another issue related to aberration is mostly nomenclatural. In this paper I describe the origin of the delay-radius relation as the `straightening of electron trajectories' when they are transformed from CF to IOF. Or, one can say the effect is due to the transformation of electron acceleration from the non-inertial CF to IOF. In the wording of HA the phase shift caused by these effects is described as `aberrational' too (see eg.~their Appendix F and G). An argument that can (possibly) justify this is eq.~\\mref{vel} (their eq.~F1). The addition of velocities (aberration) given in this equation determines the electron velocity in IOF, which is next differentiated to obtain IOF acceleration. Of course, what really matters is not the fact that we add the corotation velocity, but the fact that the added velocity is time-dependent (or that CF is non-inertial). Since the phenomenon of aberration by itself refers only to the velocity of a reference-frame, and not to its acceleration, the term `aberration' misses the essence of the effect, which is the \\emph{time-dependent nature} of the corotation velocity in \\mref{vel}.  A word of comment on papers that have assumed that the PA curve is shifted \\emph{forward} in phase, just as the intensity profile does (eg.~KD83; Xu et al.~1997; note that the work of KD83 was published in the pre-BCW era). If one corrects the analysis of KD83 for the direction of the PA shift, the altitude order of their four emission regions (see.~their fig.~16) should be turned upside-down, ie.~the region 4 should be at the top, region 3 below 4, etc.  Finally a note related to the high-energy emission. Electrons with Lorentz factors $\\gamma \\simeq 10^6$ -- $10^7$ that move along the dipole axis trajectory with $\\rhomu = \\rlc/2$ and $P \\sim 0.1$ s emit curvature radiation that extends up to $0.1$ -- $100$ MeV, respectively. This energy is too low for pair production (see fig.~7 in Dyks \\& Rudak 2002). However, this radiation may contribute to the observed X-rays and gamma-ray pulse profiles and spectra. In the numerical models of the high-energy emission from pulsars it is therefore necessary to calculate the curvature radiation using the radius of curvature  of electron trajectory in the inertial observer frame."
    },
    "0806/0806.2883_arXiv.txt": {
        "abstract": "Stars of late-M and L spectral types, collectively known as Ultracool Dwarfs (UCDs), may be excellent targets for searches for extrasolar planets. Owing to their small radii, the signal from an Earth-size planet transiting a UCD is, in principle, readily detectable. We present results from a study designed to evaluate the feasibility of using precise near infrared (NIR) photometry to detect terrestrial extrasolar planets orbiting UCDs. We used the Peters Automated InfRared Imaging TELescope (PAIRITEL) to observe a sample of 13 UCDs over a period of 10 months. We consider several important systematic effects in NIR differential photometry and develop techniques for generating photometry with a precision of 0.01 mag and long-term stability. We simulate the planet detection efficiency of an extended campaign to monitor a large sample of UCDs with PAIRITEL. We find that both a targeted campaign with a single telescope lasting several years and a campaign making use of a network of telescopes distributed in longitude could provide significant sensitivity to terrestrial planets orbiting UCDs, potentially in the habitable zone. ",
        "introduction": "Theories of planet formation make different predictions about the frequency, mass, and separations of planetary companions to stars of different masses.  The existence of small ($M\\ll M_{\\rm{J}}$) companions to small stars at separations $<10$ AU, and a relative paucity of large ($M\\approx M_{\\rm{J}}$) companions, is a general prediction of the core accretion model of planet formation \\citep{Adams2005}. \\citet{Ida2005} found that planet formation via core accretion can be efficient around low-mass stars and that it is even possible for small planets to open a gap in the protoplanetary disk and move inwards through Type II migration. Recently, \\citet{Payne2007} studied the specific case of planet formation around Ultracool Dwarfs (UCDs\\footnote{Here we follow the convention of \\citet{Bailer-Jones2002} and define the term Ultracool Dwarf (UCD) to include objects of spectral type later than about M7, including L and T dwarfs. Such objects may be bonafide brown dwarfs or very small hydrogen-burning stars.}) and found that the formation of terrestrial planets up to 5 M$_{\\earth}$ should be possible, though migration of these planets inward may be unlikely. The formation of large companions can also be explained within the disk instability paradigm. \\citet{Boss2006} recently suggested that searching for massive companions to UCDs is an important way to differentiate between these two competing planet formation scenarios. There are initial indications from radial-velocity searches and microlensing surveys that early-M dwarfs have a relative paucity of Jupiter-mass companions compared to Sun-like stars \\citep{Butler2004, Gaudi2002}. Recently, there has been evidence that lower mass companions to M dwarf stars may be a common occurrence. The microlensing detections of super-Earth and Neptune mass planets orbiting M dwarfs by \\citet{Beaulieu2006} and \\citet{Gould2006b} indicates that these types of planets may be relatively common.  There are more than 2000 UCDs known today. These objects, many discovered with the 2-Micron All Sky Survey (2MASS, \\citealt{Cruz2003}), are a well-studied class spanning the M, L, and T spectral types. Little is known about the formation of planets around UCDs. Disks around young UCDs, which would be  a starting point for all planet formation mechanisms, appear to be quite common. \\citet{Luhman2005}  found that around 50\\% of the UCDs in two star-forming regions show evidence for infrared excess attributed to the presence of a disk. This disk fraction is consistent with the disk fraction of larger stars also observed in those star forming regions. The authors argue that this is strong evidence for a common formation mechanism of stars and UCDs and that the raw materials for planet formation are available near UCDs as often as near early M-type stars. Observations of UCD disks by \\citet{Apai2005} show that dust settling and grain growth are taking place in these disks, indicating that an important first step in planet formation is occurring in disks surrounding UCDs.  Current planet searches pay relatively little attention to stars of spectral type later than about M4. There are seven known M dwarf planetary systems (GJ 876: \\citealt{Delfosse1998,Rivera2004}, GJ 436: \\citealt{Butler2004}, GJ 581: \\citealt{Bonfils2005}, GJ 674: \\citealt{Bonfils2007}, GJ 317: \\citealt{Johnson2007}, GJ 849: \\citealt{Butler2006}, and GJ 176: \\citealt{Endl2007})  from radial velocity searches, but a search for planetary companions to UCDs will likely require observations in the infrared. High-precision optical radial velocities of early-M dwarfs have already been reported in the literature \\citep{Endl2003}, and the techniques for obtaining radial velocities of UCDs in the NIR are approaching the level required for confirming massive, short-period companions \\citep{Blake2007a}. A transit search offers the exciting possibility of detecting close-in companions to UCDs and, because of the relatively small size of the host UCDs, enables, in principle, the detection of terrestrial planets.  The possibility of detecting terrestrial planetary companions to UCDs and M dwarfs through observations of transits has been explored by \\citet{Gould2003}, \\citet{Blake2003}, \\citet{Caballero2003}, \\citet{Snellen2005}, \\citet{Plavchen2005,Plavchen2007}, and \\citet{Nutzman2007}. If photometric precision similar to that realized by transit searches at optical wavelengths can be achieved in observations of UCDs, rocky companions as small as Earth would be detectable.  Searching for transits of UCDs presents many challenges. These objects are intrinsically faint and therefore very rare in the shallow, wide-field surveys typically used to find transiting planets.  Most current transit searches exploit a multiplexing advantage by observing many thousands of stars simultaneously (i.e. \\citealt{Bakos2004}), but the low density of UCDs on the sky requires observing them individually.  While photometric precision sufficient for the detection of transits is regularly achieved by surveys operating in the optical regime, the UCDs are much more luminous in the NIR than in the optical. Therefore, observations must be carried out in a wavelength regime where observational techniques are not as well developed and it may not be possible to design suitable wide-field instruments. Several authors have  discussed these challenges and the prospects for high precision, NIR, time-series photometry. In addition to the results presented by \\citet{Koen2005} and \\citet{Bailer-Jones2003}, it has been demonstrated by \\citet{Snellen2005} that modern ground-based NIR detectors are capable of producing 0.1\\% differential photometry of bright point sources under certain circumstances.  It is important to ascertain the level of intrinsic variability of the UCDs themselves.  While large photometric surveys, such as microlensing searches and transiting planet surveys, have provided a wealth of information concerning the photometric properties of larger and bluer stars, relatively little is known about the time-dependent behavior of UCDs. For a transit search, we are primarily concerned with variability on timescales comparable to the duration of the transit events. Their complex, dynamical atmospheres, in conjunction with  their observed rapid rotation, could result in observable photometric variability due to spots. Many late M dwarfs are thought to have strong magnetic fields and are found to flare. These flares, which can be large in amplitude in the U band (i.e. \\citealt{Rockenfeller2006}), are found to strongly decrease in amplitude with wavelength. There is little evidence in the literature for similar flares on L dwarfs. There is a growing body of literature concerning ``weather'' and the intrinsic variability of UCDs at both optical and NIR  wavelengths \\citep{Gelino2002, Bailer-Jones2003,Koen2005}. Observations spanning the optical to the mid-infrared have produced conflicting accounts of the intrinsic variability of L dwarfs. While it seems clear that these objects are not found to be so variable as to preclude the detection of transiting planets, it is not clear at what level  they do exhibit variability. Recent work by \\citet{Calderon2006} demonstrated that in the mid-infrared at least one L dwarf is photometrically stable at the level of 3 mmag.  Here, we present the results of a pilot study designed to evaluate the feasibility of a NIR, targeted search for transiting terrestrial companions to UCDs. In Section 2 we outline the basic properties of the types of planetary systems that we could hope to detect. In Section 3 we describe the robotic observing system used to gather these observations and the process of data reduction and photometry. In Section 4 we describe some of the problems unique to differential photometry of cool stars in the NIR. In Section 5 we summarize the overall quality of the NIR observations of UCDs and present the results of simulations designed to evaluate the feasibility of detecting terrestrial planets in with a survey like ours. ",
        "conclusions": "We present NIR observations of 13 UCDs that are photometrically stable over several months. We demonstrate that UCDs are viable targets for searches for transiting extrasolar planets and that an intensive observing campaign with a single telescope spanning several years could detect terrestrial companions potentially located in the habitable zones of their UCD hosts. Similar planets could also be found with a campaign spanning less than six months with a world-wide network of of small aperture telescopes (such as envisioned by Las Cumbres Observatory). The current photometric precision ($\\approx 1\\%$) is sufficient for the detection of companions as small as $3M_{\\earth}$ and, with future improvements, it may be possible to detect companions as small as Earth.  A UCD transit search is currently one of the most promising ways to search from the ground for terrestrial planets in the habitable zones of their hosts. Since bright UCDs are relatively rare, such a search needs to be carried out one object at a time. We have simulated our transit detection efficiency and found that with a long time baseline our observing strategy can effectively recover events with short periods. To obtain the required phase coverage to detect planets in short period orbits it is necessary to observe each target over many months. We estimate that a four year program to observe 60 UCDs could place a significant upper limit on close-in, super-Earth to Neptune size companions to UCDs. While analysis of current transit searches has led to the conclusion that fewer than 1 in 300 main sequence stars has a close-in massive companion \\citep{Gould2006a}, we know virtually nothing about the occurrence of planetary companions to UCDs or the rate of super-Earth companions to any type of star. Observations of planet-forming disks around UCDs lead us to believe that it is likely possible to form planets around these low mass hosts. Future observations of these small stars, both by searching for transits in the NIR and gathering increasingly precise NIR radial velocity measurements, will hopefully reveal an exciting new class of planets which are susceptible to study by direct imaging and provide important tests of models of terrestrial planets and their formation."
    },
    "0806/0806.2551_arXiv.txt": {
        "abstract": "We present a detailed numerical study of the Gough \\& McIntyre model for the solar tachocline. This model explains the uniformity of the rotation profile observed in the bulk of the radiative zone by the presence of a large-scale primordial magnetic field confined below the tachocline by flows originating from within the convection zone. We attribute the failure of previous numerical attempts at reproducing even qualitatively Gough \\& McIntyre's idea to the use of boundary conditions which inappropriately model the radiative--convective interface. We emphasize the key role of flows downwelling from the convection zone in confining the assumed internal field. We carefully select the range of parameters used in the simulations to guarantee a faithful representation of the hierarchy of expected lengthscales. We then present, for the first time, a fully nonlinear and self-consistent numerical solution of the Gough \\& McIntyre model which qualitatively satisfies the following set of observational constraints: (i) the quenching of the large-scale differential rotation below the tachocline -- including in the polar regions -- as seen by helioseismology (ii) the confinement of the large-scale meridional flows to the uppermost layers of the radiative zone as required by observed light element abundances and suggested by helioseismic sound-speed data. ",
        "introduction": "The presence of the tachocline, a thin shear layer located at the interface between the radiative and convective regions of the Sun, was established two decades ago (Christensen-Dalsgaard \\& Schou, 1988; Kosovichev, 1988; Brown {\\it et al.} 1989; Dziembowski {\\it et al.} 1989) but its {\\it modus operandi} still remains mysterious.  Anisotropic turbulent stresses associated with rotationally constrained eddies are thought to maintain the differential rotation profile observed within the convection zone: \\begin{equation} \\Omega_{\\rm cz}(\\theta,r) \\simeq \\Omega_{\\rm eq}(r) (1-a_2(r) \\cos^2\\theta - a_4(r) \\cos^4\\theta) \\mbox{   ,   } \\label{eq:ocz} \\end{equation} where for example at $r = 0.75\\rsun$ $\\Omega_{\\rm eq}/2\\pi = 463 {\\rm nHz}$, $a_2 = 0.17$ and $a_4  = 0.08$ (Schou {\\it et al.} 1998; Gough 2007). However, as shown by Spiegel \\& Zahn (1992) (SZ92 hereafter), the mere reduction in the amplitude of these stresses naturally expected to occur across the radiative--convective interface (at $r_{\\rm cz} = 0.713 \\rsun$) cannot explain the transition to near-uniform rotation below. It was later argued by Gough \\& McIntyre (1998) (GM98 hereafter) that only {\\it long-range} stresses can explain the suppression of the rotational shear in the bulk of the radiative zone and in addition maintain the angular velocity of the interior (observed to be $\\Omega_{\\rm rz}/2\\pi \\simeq 430$nHz) close to that of the surface despite the global spin-down induced by magnetic-braking.  Two competing theories for these presumed long-range stresses have been investigated: purely hydrodynamic stresses in the form of gravity waves (see the review by Zahn, 2007) and hydromagnetic stresses (see the review by Garaud, 2007). This paper focuses on the latter mechanism only, although the dynamical balance in the Sun could arguably involve a combination of the two.  It has long been known that even a very small primordial magnetic field embedded in the radiative zone could in principle impose uniform rotation (Ferraro, 1937; Mestel, 1953; Mestel \\& Weiss, 1987). Ferraro's isorotation law, \\begin{equation} \\bB \\cdot \\grad \\Omega = 0 \\mbox{   ,   } \\end{equation} valid in the limit of negligible dissipation and for steady-state, axisymmetric flows, is usually stated as {\\it the angular velocity must be constant on magnetic field lines}. Thus in its simplest form, Ferraro's law predicts the possibility of uniform rotation for the radiative zone provided the magnetic field is entirely confined beneath the radiative--convective interface (R\\\"udiger \\& Kitchatinov 1997). On the other hand, any field line directly connected with the convection zone promotes the propagation of the rotational shear into the radiative zone (MacGregor \\& Charbonneau 1999), inducing what will be referred to from here on as a ``differentially rotating Ferraro state''. Hence, field confinement is the key to the existence of a tachocline.  GM98 were the first to address the question of how the presumed primordial field could indeed be confined within the radiative zone, and proposed that meridional flows driven by Coriolis forces in the convection zone and burrowing downward would interact nonlinearly with the underlying magnetic field lines,  bending them towards the horizontal in the tachocline region, thus effectively suppressing direct radial Alfv\\'enic transport between the  convection zone and the radiative zone. Conveniently, the same meridional flows can also be held responsible for mixing light elements such as  Li and Be between the convection zone and their respective nuclear-burning  regions, reducing He settling (Elliott \\& Gough, 1999) as well as providing weak but sufficient angular-momentum  transport to adjust the mean rotation rates of the convective and radiative zones continuously throughout the spin-down phase. The seminal boundary-layer  analysis of the dynamics of the tachocline presented by GM98 validated the plausibility of this theory, although many of their simplifying scaling assumptions remain to be confirmed through the direct numerical solution of the governing equations.  This paper first briefly reviews existing work and then presents new results on the laminar tachocline dynamics according to GM98. We begin by discussing past attempts at implementing their model numerically in Section \\ref{sec:prevfail}. In particular, we argue that the failure of previous numerical studies in reproducing field confinement can be explained by the selection of inappropriate boundary conditions. A new numerical algorithm is then presented in Section \\ref{sec:model}, and used in Section \\ref{sec:numexp} to revisit the idea proposed by GM98 with much more success. We discuss future prospects in Section \\ref{sec:disconcl}. ",
        "conclusions": "\\label{sec:prevfail}  GM98 argue that the radiative interior should be divided into three dynamically distinct zones including (from the base of the convection zone downward) (1) a more-or-less magnetic-free region in thermal-wind balance, well-ventilated by meridional flows originating from the convection zone, which can be thought of as the bulk of the tachocline, (2) a very thin magnetic advection-diffusion layer where the tachocline flows and the underlying field interact nonlinearly to confine one another and (3) a magnetically dominated, near-uniformly rotating interior.  The nonlinear nature and geometric complexity of the problem precludes analytical solutions, and all existing studies since the original work of GM98 have been numerical (see the review by Garaud, 2007). Among these, only two include all the nonlinear terms required (i.e. the advection terms in the momentum and the magnetic induction equations, and the Lorentz force in the momentum equation) to represent the nonlinear magnetohydrodynamics of the tachocline ``correctly'': Garaud (2002) -- hereafter Paper I -- and Brun \\& Zahn (2006) -- hereafter BZ06. Surprisingly, neither have been able to find evidence for the kind of dynamical balance proposed by GM98, and the time is therefore ripe to take a step back and discuss why.  \\subsection{``Failure'' of previous numerical models}  Paper I and BZ06 are two studies which model magnetohydrodynamic perturbations induced in the solar radiative zone by the differentially rotating convection zone and by an assumed primordial magnetic field. Both papers focus on the issues of field confinement and the suppression of differential rotation below the tachocline.  Paper I presents steady-state, axially symmetric solutions of an incompressible and unstratified version of the GM98 model. BZ06 on the other hand use a time-dependent, three-dimensional algorithm based on the ASH code (Glatzmaier 1984; Clune {\\it et al.} 1999; Miesch {\\it et al.} 2000; Brun {\\it et al.} 2004), and solve the complete set of anelastic MHD equations.  Both studies otherwise consider a similar computational domain, namely a spherical shell which spans the region between the base of the convection zone (at $r = r_{\\rm out}\\simeq r_{\\rm cz}$) and an inner sphere (at $r = r_{\\rm in}$). The boundary conditions are also essentially similar: the ``outer'' boundary, which effectively models the radiative--convective interface, is in both cases assumed to be impermeable and rotating differentially with an angular velocity profile given by equation (\\ref{eq:ocz}); the magnetic field is matched onto a potential field.  In Paper I, the numerical solutions only show confinement of the magnetic field lines in the equatorial regions for low enough diffusivities,  occasionally at mid-latitudes for lower field strength, but never in the  polar regions (see Figure 11 of Paper I for example). It is commonly argued that the failure of Paper I to find fully-confined magnetic field solutions stems entirely from the simplified nature of the model equations (incompressible, unstratified): indeed, within these assumptions the meridional flows are generated only by Ekman-Hartmann pumping on the boundaries, and their amplitude scales with the diffusivities in a way which always maintains the magnetic Reynolds number below unity. As a result, these flows are unable to confine the magnetic field and the differential rotation imposed at the upper boundary of the domain persists within a large part of the radiative zone.  The three-dimensional, time-dependent solutions of BZ06 naturally depend on the initial magnetic field configuration selected. Various cases with an initial field more-or-less deeply confined are discussed. Against expectations, BZ06 find that regardless of the initial conditions, the field lines always spread out and eventually overlap with the convection zone, permitting the propagation of the differential rotation into the radiative interior. Interestingly, the transient state of the system -- prior to any field line connecting with the convection zone -- qualitatively looks in many ways similar to the well-known GM98 picture, although it is clearly not a steady state. The ``tachocline'' thus formed slowly narrows with time until field lines spread into the convection zone, at which point a differentially rotating Ferraro state is rapidly established. One is forced to conclude in one of three ways: (1) for unspecified primordial reasons, the initial magnetic field was more deeply embedded in the radiative zone and the currently observed tachocline is merely a transient phase; (2) the tachocline is indeed in a steady-state and some of the assumptions made in the way have been wrong; (3) the parameter regime studied by BZ06 (where diffusion coefficients are artificially increased by many orders of magnitude to satisfy numerical constraints) does not appropriately reproduce the solar interior dynamics.  One should be uncomfortable in selecting the first option, as it would place very strong and unlikely constraints on the initial field conditions to provide just the right structure for today's tachocline. One could naively tend to favour the third option, but BZ06 {\\it a priori} took care to select a range of parameters for which all expected boundary layer thicknesses are small enough and in the same hierarchical order as those of the model proposed by GM98. The plot thickens while we are left with the uneasy task of reconsidering the key features of either the numerical models or of the GM98 model (both, perhaps).  \\subsection{The source of the problem} \\label{subsec:sourceprob}  At this point it is worth discussing one of the more delicate aspects of the GM98 model, namely the exact mechanism by which these pivotal meridional flows are thought to be generated. This point has been  a source of confusion and debate, but is clearly crucial to a better  understanding of the tachocline dynamics. In the original work of GM98, the principal clue to the nature of the flows can be found in the sentence:  {\\it ``Turbulent stresses in the convection zone induce (through Coriolis  effects) a meridional circulation, causing the gas from the convection zone  to burrow downwards ...''}. The flows considered by GM98 do not originate from within the tachocline and can therefore  not be appropriately modelled  by a numerical scheme in which the radiative--convective interface is assumed to be impermeable (as in Paper I and BZ06). While this conclusion seems obvious in hindsight, the physics of the problem are actually rather subtle and deserve clarification. In what follows, we discuss the issue in more detail and present a unified view of the results of previous works on the subject.  Spiegel \\& Zahn were the first to study the dynamics of the newly-discovered tachocline (SZ92). They considered a non-magnetic radiative zone only, and imposed a latitudinally-varying rotation profile at the radiative--convective interface. They performed two distinct calculations. The first looked at the time-dependent evolution of the angular velocity profile in the radiative zone under such forcing, assuming isotropic viscous stresses. The second sought steady-state ``tachocline'' solutions assuming anisotropic turbulent stresses. Since the latter did not specifically address the question of the meridional flows, we focus here on the results of the time-dependent calculation, which can be interpreted in the following way. The differential rotation imposed by the convection zone to the top of the radiative zone inevitably induces some degree of shear along the rotation axis if the radiative zone is originally in a state of uniform rotation. This generates a thermal gradient in the latitudinal direction by way of the thermal-wind equation (see GM98, equation (1) for example). Meridional flows are then required to balance this thermal gradient when the system is in thermal equilibrium. These flows burrow into the radiative zone, advecting angular momentum thereby helping the propagation of the shear further down. The results of SZ92 imply that the system continues evolving in this fashion until the radiative zone has achieved complete thermal and dynamical equilibrium.  The characteristic amplitude of the time-dependent flows associated with this thermo-dynamical relaxation process is, by way of the assumptions listed above, only dependent on the (evolving) local differential rotation, stratification and thermal conductivity. Their turnover timescale is calculated to be of the order of the local Eddington-Sweet timescale, which is short in the initial relaxation stages and steadily increases as the system evolves. Crucially, this result was derived independently of any boundary condition on the meridional flows. It is therefore correct to think of the induced transient flows as being driven by baroclinic stresses from {\\it within} the tachocline rather than downwelling from the convection zone. In fact, they exist {\\it even if} the radiative--convective interface is assumed to be impermeable. This fact is probably at the origin of the impermeable boundary conditions selected by Paper I and by BZ06, in spite the obvious contradiction with GM98's intent.  However, it is vital to remember that this transient phase and its associated flows both end once the system achieves thermal and dynamical equilibrium. The only flows remaining in the final steady-state are driven, within tiny boundary layers, by the thermo(-magneto)-viscous stresses required to match the bulk equilibrium solutions to the applied boundary conditions. Thus, as expected from any elliptic problem, the nature of the boundary conditions selected entirely controls the steady-state solutions. This is clearly illustrated by the work of Garaud \\& Brummell (2008) (GB08 hereafter), which complements that of SZ92 by calculating the spatial properties as well as characteristic amplitudes of {\\it steady-state} meridional flows in the radiative zone, as induced by various kinds of forcing applied at the radiative--convective interface.  GB08 showed that in the non-magnetic case, the induced steady-state flows can always be viewed -- at least in the linear sense, and for solar parameters -- as the sum of two ``modes'': a viscously-dominated ``Ekman mode'', which very rapidly decays away from the interface (on an Ekman lengthscale), and a ``thermo-viscous mode'' which can essentially span the entire radiative zone when not hindered by a magnetic field. Since the Ekman mode decays too rapidly to have any effect on the tachocline dynamics, the flows discussed by GM98, which are thought to ventilate the bulk of the tachocline, should be identified with the thermo-viscous mode. This view is perfectly consistent with GM98's analysis, since the thermo-viscous mode is in thermal-wind balance (GB08).  The respective amplitudes with which the Ekman and thermo-viscous modes are driven depend on the nature, amplitude and spatial structure of the forcing. GB08 showed that when the interface is impermeable, as assumed in Paper I and by BZ06, then the amplitude of the thermo-viscous mode is negligible for microscopic solar values of the diffusivities. As a result, the magnetic Reynolds number of these tachocline flows is much smaller than unity, which straightforwardly explains why field confinement eluded these two previous studies.  When flows are pumped directly through the interface by stresses within the convection zone (i.e. the radial component of the flows is non-zero at the radiative--convective interface), GB08 find that the amplitude of the thermo-viscous mode can be much higher, in which case the tachocline flows may be expected to confine the field in a scenario qualitatively similar to the one proposed by GM98. The quantitative detailed analysis must be done numerically; this is the purpose of the present study.   \\label{sec:disconcl}  \\subsection{Summary of the results}  Despite the difficulties encountered when attempting to find numerical solutions for asymptotically low values of the diffusivities, the nonlinear dynamics which emerge from our simulations are closer to what one may expect from the GM98 model than any other simulation performed to date.  More precisely we do observe (see Section \\ref{subsec:explor}) the quenching of the large-scale differential rotation by the primordial magnetic field, even in the polar regions. We also observe the partial confinement of the field by the meridional flows to the radiative interior (see Section \\ref{subsec:fieldconf}), with the reduction of the radial field strength (in some regions) by more than 70\\%. Finally, we observe the concurrent confinement of the meridional flows to the upper layers of the radiative zone ($r > 0.67\\rsun$) by the magnetic field. We have not yet been able to reduce the diffusivities down sufficiently far to observe a truly segregated structure where the bulk of the tachocline flows are completely magnetic-free. However, there is evidence in the observed rotation profile (see Figure \\ref{fig:bestsim}) for a transition between regions where angular-momentum transport is dominated by the meridional flows, and regions where it is dominated by the magnetic field. The ``magnetic diffusion layer'' studied by GM98, which is thought to control this transition, appears to have a rather complex geometrical structure which prevents a more detailed study of the GM98 scalings.  The calculated value of the interior angular velocity $\\Omega_{\\rm in}$ (see Section \\ref{subsec:omvsf}) does not match the observed value. However, since even the lowest-diffusivity simulations presented here are only just beginning to enter the asymptotic parameter regime described in Section \\ref{subsubsec:paramsreg} we do not view the poor match with the observations as an intrinsic problem with the GM98 model (yet) but rather as evidence that more work should be done to decrease the diffusivities even further -- a challenging task.  \\subsection{The stability of the solutions?}  As described in Section \\ref{subsec:explor}, the convergence of the solutions becomes rather difficult when the values of the diffusivities are decreased below a certain threshold (in the fiducial model for $\\tilde{f} < 8 \\times 10^9$). The typical reasons for this convergence failure were listed and discussed in Section \\ref{subsubsec:convsol}: intrinsic linear/nonlinear instabilities or insufficient latitudinal resolution. In this particular case, the various scenarios are equally plausible, and difficult to disentangle without a supporting fully nonlinear 3D calculation. The field configuration near the inner core could be subject to instabilities (as seen in the simulations of BZ06). The Ekman(-Hartmann) layer near the outer boundary could also be subject to instabilities, or could be developing latitudinal structures which are too fine to resolve (see Figure \\ref{fig:el}). We have in fact tentatively identified the emergence of strong jets in the angular velocity profile (see Figure \\ref{fig:bestsim} for example) as the most likely source of instabilities (and convergence failure), but the reader should be cautionned that this statement is only speculative.  As mentioned in Section \\ref{subsec:explor}, whether the equilibrium governed by the GM98 model is stable or unstable cannot directly be inferred from the stability of the numerical solutions (since they operate in a different parameter regime) but it is clearly a fundamental question. Should our diagnostic concerning the stability of the jets prove to be correct, then a possible path for future investigation (using this relaxation method) may be to continue searching for solutions, starting where we left off, and slowly lowering simultaneously the amplitude of the imposed flow (and field) with the diffusivities to limit the strength of the jets generated. This could be seen as blindly navigating the stable regions of parameter space while avoiding instability reefs. But what could ideally be derived from such an exercise, should we be able to acquire sufficient data, are scaling relationships between the jet strength, the diffusivities and the imposed flow strength which may then be used to infer tentative information on the stability of the GM98 solution itself.  \\subsection{The role of an Ekman layer?} \\label{subsec:ekmanrole}  \\begin{figure} \\epsfig{file=f12.eps,width=8cm} \\caption{Angular velocity profile on and just below the radiative--convective interface (a) at 0.7$\\rsun$ (b) at 0.6999$\\rsun$, (c) at 0.6995$\\rsun$ and (d) at 0.699$\\rsun$ in the fiducial model for $\\tilde{f}=8\\times 10^9$. Note how the angular velocity profile in (and therefore also below) the Ekman layer is very different from the one imposed by the convection zone.} \\label{fig:el} \\end{figure}  Our simulations have also revealed a fundamental difference with the original GM98 model: the presence and role of an Ekman mode\\footnote{Since the bulk of the GM98 tachocline is presumed to be magnetic-free, one should indeed consider an Ekman mode rather than an Ekman-Hartmann mode}. While it is clear from Figure \\ref{fig:bestsim} that the Ekman flows themselves play no role in confining the field (contrary to the claims made by Kitchatinov \\& R\\\"udiger 2006) our simulations reveal that the role of the Ekman layer is still far from trivial.  Indeed, the differential rotation profile imposed at the radiative--convective interface is quite different from the differential rotation profile transmitted by the Ekman layer. This is illustrated in Figure \\ref{fig:el}, which shows the evolution of the angular velocity profile with depth across the Ekman layer. Thus while the GM98 model correctly describes the non-viscous tachocline dynamics {\\it below} the Ekman layer, the Ekman mode could in fact influence the system by modifying the angular velocity profile ``seen'' by the bulk of the tachocline\\footnote{adding another layer to the sandwich...}.  The extent to which this effect influences the tachocline is unclear. Firstly, other processes which also transport angular momentum (convective overshoot, gravity waves, small-scale and large-scale magnetic stresses associated with the dynamo field) are known to take place in the close vicinity of the radiative--convective interface. These processes were not included in GM98's analysis and cannot be modelled with the present numerical algorithm, but could equally affect the tachocline dynamics. Secondly, Ekman layers (laminar or turbulent) must {\\it by definition} be present in any rotating system which exhibits very rapid changes in the imposed stresses. However, whether the Ekman mode actually plays an important role in the tachocline dynamics depends on the relative importance of the bulk thermal-wind stresses and the combination of all the other rapidly varying turbulent stresses (see Figure \\ref{fig:bls}). Indeed, if the base of the convection zone is already essentially in thermal-wind balance (see Miesch, 2005 for instance), then the amplitude of the Ekman mode and its effect on the angular velocity profile will be small (GB08). On the other hand, if this is not the case then significant Ekman flows are expected (as in the simulations shown in the present work for example), with the aforementionned consequences.  \\begin{figure} \\epsfig{file=f13.eps,width=8cm} \\caption{\\small A pictorial representation of the expected tachocline flows in two extreme cases. In both pictures, flows are downwelling from the convection zone into the radiative zone. As in the GM08 model, the bulk of the tachocline is ventilated by the thermo-viscous mode, which interacts with the magnetic field in a thin advection-diffusion layer. On the left, and as in the simulations presented in this paper, the convection zone is not dominated by thermal-wind balance, and a significant portion of the flows downwelling from the convection zone rapidly return within a thin Ekman layer. The angular velocity profile seen by the tachocline differs from the one observed in the convection zone. On the right, a hypothetical situation where the convection zone is essentially in thermal-wind balance. In that case, the Ekman mode is negligible and the GM98 model directly applies.} \\label{fig:bls} \\end{figure}  \\subsection{Prospects}  The present work has emphasized the necessity of flows downwelling from the solar convection zone as a means to confine the internal primordial field and guarantee the uniform rotation of the radiative interior as originally proposed by GM98. The nature of the thermal and dynamical balance governing the convection zone itself therefore also controls the dynamics of the tachocline through the spatial variation and amplitude of the flows crossing the radiative--convective interface. Meanwhile, steady progress in modelling the convection zone has emphasized its dependence on the thermal stratification of the tachocline (Rempel, 2005; Miesch, Brun \\& Toomre, 2006). One can only conclude that the dynamics of the convective and radiative regions are intrinsically and fundamentally coupled through this internal layer called the tachocline, and that the only sensible way forward from this point on is the construction of whole-Sun models including both regions."
    },
    "0806/0806.0762_arXiv.txt": {
        "abstract": "{} {We try to identify the origins  of field O-stars in the nearest 2 to 3~kpc around the Sun using the best presently available kinematic data on O-stars and on young open clusters. We investigate the question if the present-day data are consistent with the assumption that O-stars have formed in groups (clusters, associations),  or in isolation.} {We apply the epicycle theory for back-tracing the orbits of O-type stars and of candidate parent open clusters.} {From the 370 O-stars in the ``Galactic O star catalog v 2.0'' (GOSV2) we have investigated 93 stars classified as  $field$, and found the origin for 73 of them in 48 open clusters younger than 30 Myrs. Only for 32 stars or about 9\\% of all O-stars from this catalogue, the question of their origin in groups is not solved; some of them may have originated in isolation or may have disintegrated the group in which they formed. Fifty percent of the young  open clusters (age $<$ 30 Myr) in the ``Catalogue of Open Cluster Data'' (COCD) have O-stars as members, or have ejected at least one O-star in the first 10 Myrs of their life, or both. During this period the average mass loss from open clusters by ejecting O-stars is found to be  3 to 5~$M_\\odot$ per Myr. We prove that $\\zeta$ Pup had its origin in the open cluster Trumpler~10  which it left about 2.5 Myrs ago, and that its present-day distance is 300 pc (compared to 440~pc before). The revised distance implies a significant revision of the stellar parameters (a radius of 14~$R_\\odot$, a mass of 22.5~$M_\\odot$, and a luminosity of log $L/L_\\odot$ of 5.74) i.e, $\\zeta$ Pup is closer, less massive, and less luminous than previously thought. Our findings provide independent estimates of the present-day distances and absolute magnitudes of field O-stars.} {} ",
        "introduction": "~\\label{intro}  Do all O-stars form in groups (clusters, associations) as is commonly believed or is the formation of O-stars in isolation possible? This long-standing question can only be answered, when the birth-places of all O-stars will be discovered. A review of the situation is given in the introduction by  \\citet{1987ApJS...64..545G} and recently in \\citet{2007ARA&A..45..481Z}. \\citet{1987ApJS...64..545G} compilied a catalogue of 195 O-stars which he used to infer the first solid statistics about runaway and field O-stars. Recently, a new catalogue of Galactic O-stars (GOSV1 version 1, \\citet{2004ApJS..151..103M}; GOSV2 version 2, \\citet{2007astro.ph..3005S}) was published. Comprising 370 O-stars, the catalogue allows to re-address the statistics of O-star birth-places. In particular, the GOSV2 catalogue contains a subset of 105 O-stars called {\\em field}, which simply means that they cannot be identified as present or former members of recognised groups. Such ''isolated O-stars'' are of essential interest to decide the question if ''isolated'' massive star formation is possible or not.  Because of the relatively short lifetime (a few million years) near the main sequence, the orbit of an O-type star in the wider solar neighbourhood can, in principle, be followed all the way back in time to the onset of its hydrogen-burning stage. This means that the location of its parent star forming cloud can be determined. Investigating the area around these parent clouds, one may find other young objects there, e.g. young star clusters or OB-associations.   During the last decade, after the results of the Hipparcos mission became available, nearby OB-associations have been investigated in considerable detail \\citep[][]{1999AJ....117..354D}. However, not all OB-stars have been found living in associations, some are far way from presently-known stellar groups on the sky.  Using the data from ESA's Hipparcos mission, \\citet{2001A&A...365...49H} back-traced the orbits of 56  OB-type runaway stars and nine compact objects with distances less than 700 pc. They found that at least 21 objects of their sample could be linked back to nearby associations and young open clusters. The authors state that the remaining objects may have originated from distances farther away than 700 pc, where the knowledge of parent groups is poor.  Another line of argument has been followed in two papers by \\citet{2004A&A...425..937D,2005A&A...437..247D}. In their first paper they investigate the origin of 43 O-type {\\em field} stars from the O-star catalogue by \\citet{1987ApJS...64..545G}. The authors search the area around these stars for stellar groups in the near-infrared which are possibly hidden in the optical. In their second paper  \\citep{2005A&A...437..247D} they investigate the same sample kinematically. They excluded as {\\em field} stars the runaway stars, i.e. those with spatial velocities above the limit of 40~km~s$^{-1}$ set by \\citet{1961BAN....15..265B} and those at distances larger than 250~pc from the Galactic plane. Combining the results of both papers they claim that not more than 4 $\\pm$ 2\\% of all O-stars in Gies' catalogue can be called genuine {\\em field} stars.  The argument by \\citet{2004A&A...425..937D,2005A&A...437..247D} is a rather indirect one, they are excluding stars as  {\\em field} without being able to retrace their origin. This is exactly the point where we start our present study. Only if one succeeds to retrace an O-star to a parent group within its past lifetime one can say with certainty that this O-star has originated in a group. Proving or disproving this point is not an easy task given our incomplete knowledge of possible birth-places in the wider neighbourhood of the Sun and the uncertainties of the six-dimensional phase space coordinates (position and motion) of candidate stars and candidate clusters and/or associations.  In this paper we are testing the hypothesis that O-stars, the origin of which is hitherto unknown, may have been ejected from young open clusters (or their protoclusters) during or after the star formation period in the (parental) cluster. For this purpose, we follow the path of stars and clusters back in time in the Galactic potential. In the next section we present the underlying observations, then we describe the method and its  application. Section~\\ref{res} is a presentation and a discussion of the results. In Section~\\ref{indstars} we consider a few selected examples of stars with the adopted solutions, whereas in Section~\\ref{nosolu} we briefly discuss the stars for which we did not find an acceptable solution. A summary concludes the paper. ",
        "conclusions": "~\\label{conclus}  In this paper we have followed the dynamical history of O-stars that left the groups where they originated in. We leave open the physical mechanisms which are behind these events.  For 73 out of 93 O-stars considered, we found acceptable solutions indicating that the present-day data are consistent with the assumption that these O-stars were ejected from young open clusters or protoclusters during the past 10 Myr. The GOSV2 catalogue counts 370 O-stars, for 265 of which the origin is given in that catalogue. We were able to add 73 more cases to the list. For 32 stars (or 9$\\%$) we could not prove the origin in groups.  In this paper we have dealt with aspects of the early phases in the life      of open clusters. O-stars are best suited as tracers of this early-phase evolution because of their short life-time. In our sample of 161 young  open clusters (age $<$ 30 Myr) from the COCD there are 55 (or $\\approx$ 35\\%) that have O-stars as members \\citep{2005A&A...438.1163K,2005A&A...440..403K}, 23 of these have already lost one or more O-stars in their history. Another 24 (or 15\\%) of the COCD clusters had relations to O-stars in the past 10 Myrs, but do not contain O-stars at present. For 82 (51\\%) young clusters we cannot prove a relation to presently living O-stars. Either their most massive member is a main sequence star of spectral type later than O, or it is a former O-star which has already evolved. Of   the 47 clusters that have lost at least one O-star, we find 14 that are so young that O-star and cluster should already have separated in the protocluster phase.  Summing up the statistics above, the following picture emerges. Fifty percent of the clusters being able to survive the infant-mortality phase are so massive that they contain or contained O-stars. These O-stars have not been able to destroy the cluster. This, in parts, answers the question asked by \\citet{2003ARA&A..41...57L}: Do the progenitors of bound open clusters ever contain O-stars? Yes, they did.  The fact, that we could not trace back 9\\% of all O-stars from the GOSV2, does not necessarily mean that ``isolated O-star formation'' is possible. Their known astrophysical data (distance, velocities) may be incorrect or our list of possible host candidates may be incomplete. On the other hand, we can interpret our result - no solution for 20 out of 93 stars - as follows: there is an upper bound of slightly more than 20\\% of O-stars which could have destroyed their family of brothers and sisters with which they may have formed together.  It has been shown by \\citet{Pisk2008b} that classical (gravitationally bound) open clusters in the Milky Way evolve due to stellar and dynamical evolution as well as due to external perturbations. They are losing stellar mass during their live-time at an average rate of 3 to 14~$M_\\odot$/Myr. In this paper we determined the mass loss rate of young open clusters due to O-stars alone to be  3 to 5~$M_\\odot$/Myr in the first few million years of their existence.  As a by-product, we find new distances and absolute magnitudes for 73 O-stars. These indicate that the calibration of absolute magnitudes of O-stars should be revised. Their absolute magnitudes are systematically fainter by about 0.3 to 0.8 mag compared to the calibrations by  \\citet{1992A&AS...94..211G}, \\citet{1972AJ.....77..312W}, or \\citet{schmidtkaler}. This would be consistent with the re-calibration of the absolute magnitudes of early B-type stars by \\citet{1999ASPC..167..263K} using Hipparcos trigonometric parallaxes.  We have also shown that $\\zeta$ Pup, the closest O-star from the Sun, left the young open cluster Trumpler 10 some  2.5 Myrs ago. Its present-day distance from the Sun of 300 pc is compatible with the new Hipparcos distance from \\citet{2007hnrr.book.....V}. This implies a radius of 14 $R_\\odot$, a mass of 22.5 $M_\\odot$, and a luminosity of log$L/L_\\odot$ of 5.74 for $\\zeta$ Pup, i.e. the values are considerably smaller than assumed before."
    },
    "0806/0806.4970_arXiv.txt": {
        "abstract": "To evaluate the effect of turbulent heating in the thermal balance of interstellar clouds, we develop an extension of the log-Poisson intermittency model to supersonic turbulence. The model depends on a parameter, $d$, interpreted as the dimension of the most dissipative structures. By comparing the model with the probability distribution of the turbulent dissipation rate in a simulation of supersonic and super-Alfv\\'{e}nic turbulence, we find a best-fit value of $d=1.64$. We apply this intermittency model to the computation of the mass-weighted probability distribution of the gas temperature of molecular clouds, high-mass star-forming cores, and cold diffuse HI clouds. Our main results are: i) The mean gas temperature in molecular clouds can be explained as the effect of turbulent heating alone, while cosmic ray heating may dominate only in regions where the turbulent heating is low; ii) The mean gas temperature in high-mass star-forming cores with typical FWHM of $\\sim$ 6 km s$^{-1}$ (corresponding to a 1D rms velocity of 2.5 km s$^{-1}$) may be completely controlled by turbulent heating, which predicts a mean value of approximately 36~K,  two to three times larger than the mean gas temperature in the absence of turbulent heating; iii) The intermittency of the turbulent heating can generate enough hot regions in cold diffuse HI clouds to explain the observed CH$^+$ abundance, if the rms velocity on a scale of 1~pc is at least 3~km~s$^{-1}$, in agreement with previous results based on incompressible turbulence. Because of its importance in the thermal balance of molecular clouds and high-mass star-forming cores, the process of turbulent heating may be central in setting the characteristic stellar mass and in regulating molecular chemical reactions. ",
        "introduction": "The dissipation of turbulent kinetic energy provides a potentially important heating source in a variety of Galactic astrophysical environments, such as the solar wind (e.g., Matthaeus et al. 1999), interstellar clouds (e.g., Falgarone and Puget 1995), and the warm ionized medium (e.g., Minter and Balser 1997). The effect of turbulent heating has also been studied in extragalactic environments, such as in the context of the broad-line regions of quasars (Bottorff \\& Ferland 2002), and in intracluster cooling flows (Dennis \\& Chandran 2005).  Astrophysical turbulence is often highly supersonic and magnetized. Energy decay in supersonic turbulence was thought to be very rapid due to shocks (e.g., Goldreich and Kwan 1974), while the presence of strong magnetic fields was believed to suppress the compressible modes and thus reduce the dissipation rate (Arons \\& Max 1975). Recent numerical simulations offered a better understanding of energy dissipation in supersonic MHD turbulence (Stone et al. 1998, Mac Low et al. 1998, Padoan and Norlund 1999). As in incompressible hydrodynamic turbulence, where the kinetic energy decays in a turnover time of the largest eddies, the dissipation timescale in supersonic MHD turbulence is of order the flow crossing timescale (Stone et al. 1998) or, equivalently, the dynamical timescale at the driving scale (Mac Low 1999).  In typical molecular clouds the dynamical timescale is rather short, of order a million year, suggesting the need of continuous energy injection to support the observed turbulence. This result also implies a large turbulent heating rate. An estimate of the average dissipation rate from this timescale shows that it can be several times larger than the cosmic-ray heating rate, and thus may be the primary heating source in molecular clouds. Stone et al. (1998) also argued that the average turbulent heating rate can be comparable to the photoelectric heating in HI clouds with large velocity dispersions (however, according to Wolfire et al. 2003, the overall turbulent heating in the neutral medium may not be sufficient to produce the observed CII luminosity). In this paper, we provide a general theoretical formulation of the problem of turbulent heating, and investigate its effect on various types of interstellar clouds.  Turbulent dissipation is characterized by its strong spatial roughness (see Figures~\\ref{f1} and \\ref{f2}). Extreme dissipation events appear in the smallest structures occupying a tiny volume or mass fraction, while a significant fraction of the flow experiences essentially no dissipation. This implies a broad probability distribution of the dissipation rate, which must be taken into account for a consistent investigation of turbulent heating. The extended tail of this distribution at large dissipation rates, corresponding to highly dissipative structures, is responsible for the anomalous scaling of the high-order velocity structure functions, referred to as intermittency in turbulence theory (Frisch 1995).  Intermittency has been extensively studied in incompressible turbulence. The intermittent model by She and Leveque (1994), which considers a hierarchy of dissipation rates of different levels and relates them to the fractal dimension of the most intermittent dissipative structures, has been very successful in reproducing the measured scaling exponents of structure functions in incompressible turbulence. It has been shown that the model is equivalent to a log-Poisson distribution of the dissipation rate (Dubrulle 1994, She and Waymire 1995). The probability distribution of the dissipation rate in supersonic turbulence has not been studied yet. Although the She and Leveque model with a fractal dimension of 2, corresponding to shocks, agrees well with the structure functions computed from numerical simulations of highly compressible turbulence (Boldyrev et al. 2002, Padoan et al. 2004), it remains to be confirmed whether the distribution of the dissipation rate in supersonic turbulence is consistent with a log-Poisson process. This theoretical concern and the wide application of an intermittency model of turbulent heating to various astrophysical environments are the primary motivations of the present work.  An important effect of the intermittency of the turbulent dissipation is the generation of small regions with very large heating rate, and thus very high temperature. Falgarone and Puget (1995) were probably the first to recognize the importance of this effect in cold HI clouds. Observed molecules, such as CH$^+$, suggest the existence of hot regions in cold HI clouds, because their production requires temperatures much higher than the average. Adopting experimental results from incompressible turbulent flows, they found that strong local turbulent heating in HI clouds could produce a sufficient fraction of hot regions to explain the observed abundance of CH$^+$ molecules. In this work we address the validity of their result in the case of an intermittency model for supersonic turbulence, more appropriate for cold HI clouds.  In \\S 2, we study the intermittent energy dissipation in supersonic turbulence. We show that the log-Poisson intermittency model gives probability distributions for the dissipation rate in excellent agreement with those at the resolved scales of numerical simulations of supersonic turbulence. In \\S 3, we discuss heating and cooling processes in the interstellar medium and give the energy balance equations. We apply the log-Poisson intermittency model to investigate turbulent heating in molecular clouds, high-mass star-forming cores, and cold HI clouds in \\S 4. Discussions and conclusions are given in \\S 5. ",
        "conclusions": "We have studied the energy dissipation and the heating in supersonic turbulence. The turbulent dissipation is characterized by strongly intermittent fluctuations. A significant fraction of the kinetic energy is viscously dissipated in the finest, most intermittent structures, giving rise to a broad tail in the probability distribution function of the dissipation rate. To study the turbulent heating in interstellar clouds, a theoretical model is needed for the probability distribution of the dissipation rate at the dissipation scale, $\\eta$. We have generalized the log-Poisson model, originally proposed for incompressible turbulence by She and Leveque (1994), to supersonic turbulence. Because the dissipation scale, $\\eta$, cannot be resolved by current numerical simulations, we have used results from resolved inertial-range scales in our numerical simulation of supersonic and super-Alfv\\'{e}nic turbulence as a guideline for the sub-grid scales. We have found that the log-Poisson model, with a fractal dimension $d=1.64$ for the most intermittent dissipative structures, gives an excellent fit to the mass-weighted probability distribution of the dissipation rate at resolved scales in the simulation. Extrapolating the model to the corresponding Reynolds numbers, we have studied the turbulent heating in molecular clouds, high-mass star-forming cores, and cold diffuse neutral clouds. Here we summarize our results. \\begin{enumerate} \\item In typical molecular clouds, the average turbulent heating rate exceeds the cosmic ray heating rate by a factor of 3-4. Fluctuations in the heating rate give a lower mean temperature than expected from the average heating rate; temperature fluctuations make cooling more efficient because in molecular clouds the cooling rate increases sensitively with temperature. Taking intermittency into account, the turbulent heating alone gives a mean temperature of approximately 8.5~K, close to the characteristic temperature of molecular clouds. This would suggest that cosmic rays are not even needed to explain the thermal balance in molecular clouds. However, due to the intermittent fluctuations in the turbulent heating rate, a significant mass fraction of the molecular gas is not heated by the turbulence. Cosmic ray heating would dominate in these regions. Assuming a cosmic ray heating rate of $0.8 \\times 10^{-27} n \\hspace{1mm} {\\rm ergs} \\hspace{1mm} {\\rm cm}^{-3} \\hspace{1mm} {\\rm s}^{-1}$, the turbulent heating increases the average temperature by a few degrees, from 9.5~K to 13~K. We also find that turbulent heating plays a similar role in the outer regions of the dark cloud cores in the Orion molecular cloud complex.  \\item Turbulent heating provides an important energy source for the molecular gas in high-mass star-forming cores. Assuming spherical symmetry, in the inner regions of these cores (within $\\sim 0.01$~pc from the center), stellar sources heat the dust grains to a relatively high temperature. Due to the high density of these inner regions ($\\sim 10^6$~cm$^{-3}$), the gas and dust are thermally coupled, and the gas temperature is close to the dust temperature even in the absence of turbulent heating. The turbulent heating increases the gas temperature only by a few K. On the other hand, in the outer regions ($\\sim 0.1$~pc from the center), where most of the core mass resides, the turbulent heating results in a considerable increase in the gas temperature. The low density of these regions ($10^4$-$10^5$~cm$^{-3}$) makes energy exchange between gas and dusts inefficient and, without turbulent heating, the energy transfer from the dust heats the gas only to 10-20~K. Inclusion of turbulent heating increases the temperature to approximately 36~K. Because turbulent heating causes a large temperature increase in most of the core mass, it may have important implications for the dynamical evolution of the cores and for their star-formation process, and it may also be probed by future observations.  \\item The intermittent turbulent heating in diffuse HI clouds can give rise to regions much hotter than the average temperature. These warm regions have been used to explain the existence of molecules such as CH$^+$ in HI clouds, whose production needs high temperatures. Using the log-Poisson intermittency model for supersonic turbulence, we find that a turbulent rms velocity of 3~km~s$^{-1}$ at 1~pc is sufficient to account for the observed abundance of these molecules, which extends the earlier result, based on incompressible turbulence, by Falgarone and Puget (1995).  \\end{enumerate}  We point out that thermal conduction is neglected in our calculations. Conduction tends to transport thermal energy to fill in regions not significantly heated by dissipation, and thus may, to some degree, erase the fluctuations in the heating rate. As discussed in \\S~4, more intermittent fluctuations in the heating rate give a smaller average temperature in molecular clouds. Therefore, if thermal conduction were included, the average temperature in molecular clouds and in the outer regions of high-mass star-forming cores would be even larger, making turbulent heating even more important in these two cases (\\S~4.1 and \\S~4.2). The situation is different in cold diffuse HI clouds (\\S~4.3). If the fluctuations in the heating rate are less intermittent there, the average temperature becomes smaller and, more importantly, the tail of the temperature probability distribution, needed for CH$^{+}$ production, would be less extended. We estimate whether and how much thermal conduction would change our results by calculating the conduction length scale, $l_c \\simeq \\sqrt{\\kappa t_c}$, during a cooling time scale, $t_c$, where $\\kappa$ is the thermal conduction coefficient (approximately equal to the kinematic viscosity, $\\nu$). This is the scale over which thermal conduction can homogenize before the heat from turbulent dissipation is radiated away. We find that in molecular clouds and in HI clouds $l_c$ is smaller than (but comparable to) the dissipation length scale, $\\eta$. Therefore, the heat generated in the most intermittent structures cannot be transported far from these structures by thermal conduction. This justifies our choice of computing the temperature distribution by using the dissipation rate distribution evaluated at $\\eta$. In the outer regions of high-mass star-forming cores, $l_c$ is about an order of magnitude larger than $\\eta$ (the latter is very small because of the very large Reynolds number). Using the distribution of the dissipation rate at $l_c$ instead of $\\eta$ (assuming thermal conduction homogenizes the temperature over a size of $l_c$) gives an average temperature of 40~K in the outer regions of these cores, a little higher than from the distribution of the dissipation rate evaluated at $\\eta$. However, the conduction process in the presence of turbulent motions is more complex than described by the above estimate, and a detailed study of its effects is beyond the scope of this paper.  Although so far neglected, the process of turbulent heating may play an important role in the process of star formation. The mean temperature in molecular clouds defines the mean Jeans mass, which may control the peak of the stellar mass distribution. Because we have found that the mean temperature in molecular clouds and in high-mass star-forming cores may be controlled by turbulent heating, the characteristic stellar mass may be affected by turbulent heating as well. For example, the larger gas temperature predicted in high-mass star-forming cores may partly offset their large density, resulting in almost the same characteristic stellar mass, with respect to molecular clouds following Larson's relations. Finally, because the process of turbulent heating results in broad gas temperature distributions, it may be crucial in many molecular chemical reactions, besides those responsible for the formation of CH$^+$ molecules."
    },
    "0806/0806.1880_arXiv.txt": {
        "abstract": "We have studied the correlation between 2357 Chandra X--ray point sources in a $40 \\times 40$ parsec field and $\\sim$20,000 infrared sources we observed in the corresponding subset of our $2\\arcdeg \\times 1.4\\arcdeg$ {\\it Spitzer}/IRAC Galactic Center Survey at 3.6--8.0 $\\micron$, using various spatial and X--ray hardness thresholds. The correlation was determined for source separations of less than $0\\farcs5$, 1$''$ or 2$''$. Only the soft X--ray sources show any correlation with infrared point sources on these scales, and that correlation is very weak. The upper limit on hard X--ray sources that have infrared counterparts is $<1.7\\%$ (3$\\sigma$). However, because of the confusion limit of the IR catalog, we only detect IR sources with absolute magnitudes $\\lesssim 1$. As a result, a stronger correlation with fainter sources cannot be ruled out. Only one compact infrared source, IRS~13, coincides with any of the dozen prominent X--ray emission features in the $3 \\times 3$ parsec region centered on Sgr A*, and the diffuse X--ray and infrared emission around Sgr A* seems to be anti-correlated on a few-arcsecond scale. We compare our results with previous identifications of near--infrared companions to Chandra X--ray sources. ",
        "introduction": "X-ray surveys of the Galactic center with {\\it Chandra X-Ray Observatory} (Wang et al.\\ 2002; Muno et al.\\ 2003) have provided a deep sampling of the population of X--ray point sources that shows a large increase in source density toward the Galactic Center. These X--ray sources have been modeled as a population mix of various sorts of X-ray binaries, Wolf-Rayet stars, nearby X-ray active stars in the foreground, and background AGN (Pfahl et al.\\ 2002; Belczynski \\& Taam 2004; Ebisawa et al.\\ 2005; Ruiter et al.\\ 2006; Liu \\& Li 2006; Muno et al.\\ 2006).  However, the high extinction toward the Galactic center prohibits the detection of visible light emitted by the stellar components in the expected binaries.  Near--IR searches have also had little success in detecting IR counterparts of the X--ray point sources. Several OB stars and Wolf-Rayet stars have been identified with X-ray sources by near-IR spectroscopy (Muno et al.\\ 2006, Mikles et al.\\ 2006, Mauerhan et al.\\ 2007), but the paucity of near-IR detections sets limits that suggest only a small fraction of the X--ray sources can be high mass X--ray binaries (HMXBs; Laycock et al.\\ 2005; Bandopadhyay et al.\\ 2006).  Our {\\it Spitzer Space Telescope} IRAC survey of the Galactic Center (Stolovy et al.\\ 2006, S. Stolovy et al.\\ 2008 in preparation) provides a new opportunity to search for IR counterparts to the X--ray sources. IRAC observations cover four broad bands at 3.6, 4.5, 5.8 and 8 $\\micron$. The survey imaged a $2.0\\arcdeg \\times 1.4\\arcdeg$ ($280 \\times 200$ parsec at 8.0 kpc) region of the Galactic Center with a nominal resolution of $\\sim2''$ (Figure 1). Since our observations are at longer wavelengths than ground--based near--IR ($J, H, K$) observations, extinction should be less of a hinderance to the detection of stellar companions. Furthermore, at the longest IRAC wavelengths (5.8 and especially 8 $\\micron$), IRAC is sensitive to circumstellar dust emission which may cause significant extinction at shorter wavelengths. Thus, comparison of the X--ray and IR point source catalogs may reveal stellar companions which are at an evolutionary stage where they produce large quantities of dust, or are simply too heavily attenuated by the line of sight extinction at shorter wavelengths.  We calculated the correlation between 2357 hard and soft Chandra X--ray sources identified and catalogued by Muno et al.\\ (2003) and the $\\sim$20,000  {\\it Spitzer}/IRAC infrared point sources that lie within a $40 \\times 40$ parsec ($20 \\times 20$ arcmin) field at the Galactic Center (Figure 2). The IR sources are a small subset of our full catalog (Ram\\'irez et al.\\ 2008) which has a mean confusion limit of $[3.6] = 12.4$ mag. We divide the Chandra sources by their hardness because the high column density of gas towards the GC, $N_H$ $\\sim$ 5 $\\times$ 10$^{22}$ cm$^{-2}$, absorbs all soft X-rays.  Thus soft X-ray sources must be in the foreground towards the GC; hard X-ray sources can be at the 8 kpc distance of the GC or can be foreground/background sources. ",
        "conclusions": "We have analyzed the possible correlations between the largest number of candidate sources to date: 2357 X--ray sources (of which 1809 are hard X--ray sources most likely located at the Galactic Center) and $\\sim$20,000 {\\it Spitzer}/IRAC infrared point sources. Source confusion limits our correlations to only bright infrared sources with absolute magnitudes $\\lesssim 1$ if located near the Galactic Center. The lack of any significant correlation between hard X--ray sources and 3.6 -- 8 $\\micron$ infrared point sources suggests that there is no unique population of sources that are bright at both X--ray and 3.6 -- 8 $\\micron$ wavelengths. Based on this study, we can set the upper limit on the fraction of all hard X--ray sources that can be bright at both X--ray and 3.6 -- 8 $\\micron$ wavelengths to be $<1.7\\%$ (3$\\sigma$)."
    },
    "0806/0806.3557_arXiv.txt": {
        "abstract": "{We present ESO/VLT spectra in the 2.9--4.1 $\\mu$m range for a large sample of infrared stars in the Small Magellanic Cloud (SMC), mainly carbon stars, massive oxygen-rich Asymptotic Giant Branch (AGB) stars, and red supergiants. Strong emission from Polycyclic Aromatic Hyrdrocarbons (PAHs) is detected in the spectrum of the post-AGB object MSX\\,SMC\\,29. Water ice is detected in at least one Young Stellar Object, IRAS\\,01042$-$7215, for the first time in the SMC. The strength and shapes of the molecular bands detected in the evolved stars are compared with similar data for stars in the Large Magellanic Cloud (LMC). Absorption from acetylene in carbon stars is found to be equally strong in the SMC as in the LMC, but the LMC stars show stronger dust emission in their infrared colours and veiling of the molecular bands. This suggests that a critical link exists in the formation of dust from the molecular atmosphere in carbon stars which scales with the initial metallicity. Nucleation seeds based on a secondary element such as titanium or silicon provide a plausible explanation. In oxygen-rich stars, both the nucleation seeds and molecular condensates depend on secondary elements (in particular titanium, silicon, and/or aluminium), which explains the observed lower molecular abundances and lower dust content in the SMC stars. Emission from silicon monoxide seen in some oxygen-rich AGB stars and red supergiants in the SMC suggests that these metal-poor stars are able to drive strong pulsation shocks through their molecular layers. Data for pulsating dusty AGB stars and supergiants in the LMC are used to show that pulsation is likely the critical factor in driving mass loss, as long as dust forms, rather than the stellar luminosity. Finally, we suggest that the reduced dust production and consequently slower winds of metal-poor AGB stars and red supergiants are more likely to result in chemical inhomogeneities and small-scale structure in the interstellar medium. ",
        "introduction": "We owe our existence to the nuclear processing of light elements into heavy elements inside of stars, and their subsequent dispersal into interstellar space by way of stellar mass loss. One of the main contributors of carbon and nitrogen, and arguably the most important ``factory'' of cosmic dust, Asymptotic Giant Branch (AGB) stars represent the final stages of evolution of intermediate-mass stars ($M_{\\rm ZAMS}\\la$1 to $\\sim$8 M$_\\odot$) during which they lose up to $\\ga$80\\% of their mass at rates between $\\dot{M}\\sim10^{-6}$ to $10^{-4}$ M$_\\odot$ yr$^{-1}$. AGB stars thus chemically enrich the local interstellar medium (ISM) on timescales ranging from less than the dynamical timescale of a galaxy to as long as the age of the Universe. Massive stars ($M_{\\rm ZAMS}\\ga$8 M$_\\odot$) may become red supergiants (RSGs) and experience similar dusty mass loss too; this greatly affects the properties of the progenitor and circumstellar environment of the subsequent core-collapse supernova, and makes them an important source of dust produced in starbursts observed at cosmological distances.  Mass loss from red (super)giants happens when strong radial pulsations elevate the stellar atmosphere to a height where the temperature is sufficiently low, but the density is still high, for dust formation to occur. Radiation pressure from these luminous stars ($L\\sim5\\times10^3-5\\times10^5$ L$_\\odot$) is believed to drive away the dust, dragging the gas along with it (e.g., Bowen 1988; Fleischer, Gauger \\& Sedlmayr 1992; H\\\"ofner, Feuchtinger \\& Dorfi 1995), although the threshold luminosity for this to occur in oxygen-rich environments that produce relatively transparent grains (Ferrarotti \\& Gail 2006, and references therein) may be uncomfortably high to explain the mass loss from most oxygen-rich AGB stars (Woitke 2006). At the height of their mass loss dust-enshrouded AGB stars --- and in extreme cases even RSGs --- vanish at optical wavelengths but shine brightly in the infrared (IR).  An unsolved problem of red (super)giant mass loss is how the simple molecules in the stellar photosphere grow into larger assemblies that form the cores for dust growth (scenarios have been proposed by, e.g., Gail \\& Sedlmayr 1988), and how efficient these processes are at the low metallicity that is characteristic for the early Universe. Little is known about the molecular abundances within the dust-formation zone and the dust-to-gas ratio in the stellar wind. The transformation of some oxygen-rich AGB stars into carbon stars is known to have extremely important consequences: the molecular and dust formation chemistry will change from oxygen-dominated to carbon-dominated, giving rise to a vastly different array of molecules and dust particles. Hence these two types of AGB star enrich the ISM with very different material. On the other hand, to solve the opacity problem to drive oxygen-rich winds, H\\\"ofner \\& Andersen (2007) proposed that even oxygen-rich giants may form some carbonaceous grains.  Much of what we know about the mass loss from AGB stars and RSGs has been based on IR observations of dust-enshrouded AGB stars in the Large Magellanic Cloud (LMC, at 50 kpc) and Small Magellanic Cloud (SMC, at 60 kpc), where accurate luminosities and mass-loss rates can be obtained (e.g., Wood et al.\\ 1992; Zijlstra et al.\\ 2006; van Loon et al.\\ 1999b, 2005a; van Loon 2000; Matsuura et al.\\ 2005, 2006; Marshall et al.\\ 2004; Sloan et al.\\ 2006; Groenewegen et al.\\ 2007; Lagadec et al.\\ 2007). The Magellanic environments are metal-poor compared to the Sun; their ISM and star clusters (except the oldest ones) have metallicities of typically [Fe/H]$\\sim40$\\% in the LMC and $\\sim15$\\% in the SMC (e.g., Westerlund 1997; van Loon, Marshall \\& Zijlstra 2005).  The above studies suggest that metal-poor carbon stars are surprisingly abundant in C$_2$ and C$_2$H$_2$, probably because they reach higher C/O ratios than their Galactic solar-metallicity equivalents. This may be due to the oxygen-poor photospheres of metal-poor carbon stars, rather than a larger production of primary carbon. This could have important consequences for the condensation of carbonaceous dust grains in low-metallicity environments (cf.\\ Mattsson et al.\\ 2008) --- which might however depend on seeds sensitive to initial metallicity such as TiC$_2$ (Bernatowicz et al.\\ 1991). It is as yet unknown how the molecular abundances depend on metallicity in oxygen-rich stars, but here too the nucleation seeds for dust condensation may be restricted to species such as TiO and ZrO, which are then coated, first by aluminium-oxides and then by silicates (Vollmer et al.\\ 2006; Nittler et al.\\ 2008); this reliance on elements which are not synthesized inside the stars themselves (Ti, Al, Si) is consistent with the observed initial-metallicity dependence of oxygen-rich outflows (van Loon 2006; Marshall et al.\\ 2004).  Until now, 3--4 $\\mu$m spectra of stars in the SMC were limited to two carbon stars published by van Loon, Zijlstra \\& Groenewegen (1999a), and another carbon star and one post-AGB object published by Matsuura et al.\\ (2005). Here we present ESO/VLT 3--4 $\\mu$m spectra of a large sample of dust-enshrouded carbon stars, oxygen-rich AGB stars and red supergiants in the SMC, an analysis of their molecular atmospheres and a comparison with previous LMC samples (Matsuura et al.\\ 2005; van Loon et al.\\ 2006). We also present spectra of two R\\,CrB type stars, a post-AGB object showing emission from Polycyclic Aromatic Hydrocarbons (PAHs), and the first 3--4 $\\mu$m spectra of Young Stellar Objects (YSOs) in the SMC showing water ice absorption and hydrogen recombination lines. ",
        "conclusions": "We presented 2.9--4.1 $\\mu$m spectra, obtained with ISAAC at the ESO/VLT, of a sample of dusty carbon stars, oxygen-rich AGB stars and red supergiants in the SMC. Strong absorption bands mainly due to acetylene are seen in the spectra of carbon stars, whilst many of the oxygen-rich stars too show strong molecular absorption. OH and/or SiO is detected in some oxygen-rich AGB stars and red supergiants, with the SiO sometimes seen in emission.  We also presented a spectrum of the post-AGB object MSX\\,SMC\\,29, which displays very strong PAH emission. Absorption from aliphatic bonds may have been detected, which would point at the PAHs being released through photo-destruction of the grains in the relic AGB envelope. The spectra of two R\\,CrB-type stars show no molecular features but only a dust emission continuum. We suggest that these may be ``proto-post-AGB'' objects.  The properties of the molecular bands in the SMC sample were compared with similar data in the LMC. The molecular absorption in SMC carbon stars is as strong (but not stronger) as that in LMC carbon stars, although there appears to be more cold molecular gas in the LMC carbon stars. Less conclusive results are obtained for the oxygen-rich stars. Veiling by dust emission is clearly more important in the LMC than in the SMC, suggesting that dust formation is less efficient at lower metallicity. The fact that this is seen also in the carbon stars suggests that the formation of dust in the molecular atmosphere relies on an intermediary agent, e.g., a nucleation seed based on a secondary element such as titanium or silicon. The lower dust content of oxygen-rich stars may also be caused by lower abundances of secondary elements, in the nucleation seed or the condensable molecular material.  The pulsation properties of a sample of AGB stars and red supergiants in the LMC were used to show that the mass-loss rate is likely determined more critically by the star's ability to levitate the molecular atmosphere through pulsation, than by the star's ability to drive the elevated molecular layers away through radiation pressure on the dust (unless dust does not form at all).  The slower winds of less dusty metal-poor AGB stars and red supergiants were discussed in the context of the enrichment and mixing of the mass loss with the ISM. Chemical inhomogeneities and small-scale structure in the dusty ISM are suggested to be more likely in metal-poor environments.  As a by-product of this programme, we obtained spectra of several Young Stellar Objects. These include the first detections of water ice in such objects in the SMC."
    },
    "0806/0806.4231_arXiv.txt": {
        "abstract": "The existence of a gradient in the Faraday rotation measure (RM) of the quasar 3C 273 jet is confirmed by follow-up observations. A gradient transverse to the jet axis is seen for more than 20 mas in projected distance. Taking account of the viewing angle, we estimate it to be more than 100 pc. Comparing to the distribution of the RM in 1995, we detect a time variation of it at the same distance from the core over 7 yr. We discuss the origin of the Faraday rotation based on this rapid time variation. We rule out foreground media such as a narrow-line region, and suggest a helical magnetic field in the sheath region as the origin of this gradient of the  RM. ",
        "introduction": "A gradient of the Faraday rotation measure (RM) across a jet is growing evidence for the existence of a toroidal or helical magnetic field associated with the jet. The first evidence for such a gradient of the RM across a jet was found by VLBA polarimetry toward the VLBI jet of a well-known quasar, 3C 273 (Asada et al. 2002, hereafter A02). Following this report, the same kind of gradient of the RM was reported for several jets of BL Lac objects (Gabuzda et al. 2004), and the gradient of the RM across the 3C 273 jet itself was also confirmed by several observations (Zavala \\& Taylor 2005; Attridge et al. 2005). The role of a toroidal or helical magnetic field has been discussed for the launching and propagating mechanisms of jets based on magnetohydrodynamics from the theoretical point of view (e.g., Meier et al. 2001 and references therein), and it has been suggested that the presence of a toroidal or helical magnetic field could be observed as a gradient of the RM across the jet (Blandford 1993). Recently, it has also been shown that the toroidal magnetic field in a jet's rest frame would be observed as a toroidal magnetic field in the observer frame with a compression of the pitch angle (Lyutikov et al. 2005). In this paper we report on our follow-up observation, which confirms our initial results and indicates a time variation. Throughout this paper, we use a Hubble constant of H$_{0}$ = 100 km s$^{-1}$ Mpc$^{-1}$ and a deceleration parameter of q$_{0}$ = 0.5 in order to keep consistency to the previous papers (e.g., A02).  An angular resolution of one milli-arcsecond (mas) corresponds to a linear resolution of 1.86 pc. ",
        "conclusions": "In order to confirm the RM gradient across the 3C 273 jet, we made a follow-up observation using multifrequency VLBA polarimetry. The systematic gradient across the jet is confirmed for more than 100 pc along the jet, and the trend of the RM gradient is consistent with that revealed by previous observations. Since the amounts of the Faraday rotation exceed 90$^{\\circ}$, the origin of the Faraday rotation should be in the foreground of the emitting jet. On the other hand, we detected a time variation in the distribution of the RM in comparison to that in 1995, and this rapid time variation rules out the possibility that a foreground magnetized cloud independent of the jet, such as a narrow-line region, is responsible for the origin of the Faraday rotation. Therefore, the sheath around the ultra-relativistic jet is likely to be the origin."
    },
    "0806/0806.4007_arXiv.txt": {
        "abstract": "We report our new code (named SACRA) for numerical relativity simulations in which an adaptive mesh refinement algorithm is implemented. In this code, the Einstein equations are solved in the BSSN formalism with a fourth-order finite differencing, and the hydrodynamic equations are solved by a third-order high-resolution central scheme.  The fourth-order Runge-Kutta scheme is adopted for integration in time. To test the code, simulations for coalescence of black hole-black hole (BH-BH), neutron star-neutron star (NS-NS), and black hole-neutron star (BH-NS) binaries are performed, and also, properties of BHs formed after the merger and gravitational waveforms are compared among those three cases. For the simulations of BH-BH binaries, we adopt the same initial conditions as those by Buonanno et al.  [Phys. Rev.  D {\\bf 75}, 124018 (2007)] and compare numerical results. We find reasonable agreement except for a slight disagreement possibly associated with the difference in choice of gauge conditions and numerical schemes. For an NS-NS binary, we performed simulations employing both SACRA and Shibata's previous code, and find reasonable agreement between two numerical results for the final outcome and qualitative property of gravitational waveforms.  We also find that the convergence is relatively slow for numerical results of NS-NS binaries, and again realize that longterm numerical simulations with several resolutions and grid settings are required for validating the results.  For a BH-NS binary, we compare numerical results with our previous ones, and find that gravitational waveforms and properties of the BH formed after the merger agree well with those of our previous ones, although the disk mass formed after the merger is less than 0.1\\% of the total rest mass, which disagrees with the previous result. We also report numerical results of a longterm simulation (with $\\sim 4$ orbits) for a BH-NS binary for the first time. All these numerical results show behavior of convergence, and extrapolated numerical results for time spent in the inspiral phase agree with post-Newtonian predictions in a reasonable accuracy. These facts validate the results by SACRA. ",
        "introduction": "\\label{sec:intro}  Coalescence of binary compact objects such as binaries of two neutron stars (NS-NS), black hole and neutron star (BH-NS), and two black holes (BH-BH) is the most promising source for kilometer-size laser-interferometric gravitational wave detectors such as LIGO, VIRGO, and LCGT. To detect gravitational waves and to analyze the gravitational wave signals for extracting physical information of the sources, it is necessary to prepare theoretical templates of gravitational waves from the coalescing compact binaries. Motivated by this fact, significant effort has been paid in the past two decades. For theoretically computing gravitational waveforms in a relatively early inspiral phase, post-Newtonian approximations are the robust approach \\cite{Blanchet}. On the other hand, for studying the last inspiral and merger phases of the coalescing binaries in which general relativistic effects are significantly strong and any approximation breaks down, numerical relativity is the unique approach.  In the past decade, in particular in the past three years, a wide variety of general relativistic simulations have been performed for the coalescence of NS-NS binaries \\cite{NSNS1,NSNS2,NSNS3,NSNS4,NSNS9,NSNS5,NSNS6,NSNS7,NSNS8} and BH-BH binaries \\cite{BHBH1,BHBH2,BHBH3,BHBH4,BHBH5,BHBH6,BHBH7,BHBH8,BHBH9,BHBH10,BHBH11,BHBH12,BHBH13,BHBH14,BHBH15,BHBH155,BHBH16} (see also early-stage results for merger of BH-NS binaries \\cite{BHNS1,BHNS2,BHNS3}).  Since 1999, a variety of simulations have been performed for the inspiral and merger of NS-NS binaries after the first success of Shibata and Ury\\=u \\cite{NSNS1}. Shibata, Ury\\=u, and Taniguchi have then performed simulations focusing mainly on the merger process and the final fate \\cite{NSNS2}. Their simulations were done for a variety of equations of state (EOSs) as well as for a wide range of mass of two NSs. They have clarified that the final outcome of the merger (formation of a BH or a hypermassive neutron star; hereafter HMNS) depends strongly on the total mass of the system and on the chosen EOSs. In the latest paper \\cite{NSNS3}, they clarified that with stiff EOSs such as Akmal-Pandharipande-Ravenhall one \\cite{APR}, a BH is not promptly formed even for a system of the total mass $\\sim 2.8M_{\\odot}$, but an HMNS is a likely outcome. They also indicated that the formed HMNSs have an elliptical shape because of their rapid rotation, and hence, quasiperiodic gravitational waves of frequency $\\sim 3$--4 kHz will be emitted for a long time (for $\\sim 100$ cycles) in the absence of dissipative mechanisms except for gravitational wave emission.  The integrated effective amplitude of such gravitational waves may be large enough to be detected by advanced laserinterferometric gravitational wave detectors \\cite{NSNS2,NSNSS}. In the last couple of years, longterm simulations for the inspiral of NS-NS binaries have been also done. In particular, in the latest simulations, 3--5 inspiral orbits are stably followed \\cite{NSNS5,NSNS6,NSNS7,NSNS8}, and also, the computations are continued until the system settles down approximately to a stationary state even in the case that a BH is the final outcome. Preliminary simulations for merger of magnetized NSs have been also performed recently \\cite{NSNS6,NSNS9} (although it is not clear whether or not many of crucial magnetohydrodynamic instabilities are resolved in these simulations). However, in most of these works, very simple $\\Gamma$-law EOSs is adopted for modeling the NSs, and hence, realistic simulations with a variety of realistic EOSs have not been done yet.  The last three years have also witnessed great progress in simulations of BH-BH binaries, starting with the first stable simulation of orbiting and merging BHs by Pretorius \\cite{BHBH1} and development of the moving puncture approach \\cite{BHBH2,BHBH3} in 2005. Since then, a large number of simulations have been done on the late inspiral and merger of BH-BH binaries \\cite{BHBH1,BHBH2,BHBH3,BHBH4,BHBH5,BHBH6,BHBH7,BHBH8,BHBH9,BHBH10,BHBH11,BHBH12,BHBH13,BHBH14,BHBH15,BHBH155,BHBH16}. These works have clarified that the merger waveforms are universally characterized by a quasi-normal mode ring-down.  They have also shown that a large kick velocity is excited at the merger in the cases that the masses of two BHs are not equal and/or the spin and orbital angular momentum vectors misalign.  The latest works with a high accuracy \\cite{BHBH12,BHBH13,BHBH14,BHBH15,BHBH155,BHBH16} compare the numerical gravitational waveforms with post-Newtonian ones and assess the accuracy of the post-Newtonian waveforms \\cite{Blanchet}.  In particular, the numerical simulation of Ref.~\\cite{BHBH16} presents highly accurate gravitational waves, which assess the accuracy of the post-Newtonian gravitational waves with a level much beyond the previous analysis. They clarify that the so-called Taylor T4 post-Newtonian gravitational waveforms are very accurate at least up to the last two orbits before the merger for the equal-mass, nonspinning BH-BH binaries. This work shows a monumental achievement of numerical relativity because it demonstrates that numerical relativity could provide inspiral waveforms for BH-BH binaries more accurate than the post-Newtonian waveforms.  However, simulations for coalescing compact binaries have been performed only for a restricted parameter space. Because the ultimate goal is to prepare a template family which covers gravitational waveforms for almost all the possible parameters for binary compact objects, the present status is regarded as a preliminary one from the view point of gravitational wave astronomy. For example, for BH-BH binaries, the simulations have been primarily performed for the case that the spin vector of BHs aligns with the orbital angular momentum vector and the magnitude of the BH spin is not extremely large. The simulations for BH-BH binaries of unequal-mass and misaligned spin have been also performed only for the restricted cases. For NS-NS binaries, the simulations have been also primarily performed for the case that masses of two NSs are equal, and the cases of unequal-mass have been investigated in a small mass range.  Moreover, the simulations have been performed adapting a few EOSs, mostly a simple $\\Gamma$-law EOS. Because the EOS of NSs is still unknown, it is necessary to perform simulations choosing a wide variety of EOSs.  To perform a number of simulations for various parameters of compact objects, an efficient scheme for the numerical simulation is necessary. For the two-body problem considered here, adaptive mesh-refinement (AMR) algorithm is well-suited for this purpose \\cite{BO}.  The reason is described as follows: In the two-body problem, there are three characteristic length scales; the radius of compact objects, $R$, the orbital separation, $r$, and the gravitational wave length, $\\lambda \\approx \\pi(r^3/M)^{1/2}$ where $M$ is the total mass of the system. We have to accurately resolve these three scales. These scales obey the relation $R < r < \\lambda$, and typically, $R \\ll \\lambda$. Thus, an issue to be resolved in this problem is to assign an appropriate resolution for each scale of significantly different magnitude. To resolve each compact object accurately, the grid spacing $\\Delta x$ in its vicinity has to be much smaller than $R$ ($R/\\Delta x$ should be larger than $\\sim 20$). On the other hand, gravitational waves have to be extracted from the geometric variables in the wave zone. This implies that the size of the computational region should be larger than $\\lambda$. By simply using a uniform grid, the required grid number in one direction is $N_g=2 \\lambda/\\Delta x$ where the factor 2 comes from the fact that there are plus and minus directions in each axis. Because of the facts $r \\agt 2R$ and $R >M$, the required value of $N_g$ is larger than several hundreds. To follow the binary inspiral from $r \\sim 5R$, $N_g$ has to be larger than $10^3$. Even by supercomputers currently available for the general users, it approximately takes at least a month to perform a simulation of such a huge grid number. This implies that it is not feasible to perform a number of simulations for a wide variety of the parameters.  In the AMR algorithms, one can change the grid spacing and the grid structure arbitrarily for different scales, preserving the required grid-resolution for each scale. To accurately resolve each star in a binary, we need to take $N_g \\sim 2 R/\\Delta x \\sim 100$ to cover the region in the vicinity of the compact stars. However for other region, we do not have to take such a small grid spacing. In particular, we can save the grid number in the distance zone. To follow the propagation of gravitational waves in the wave zone, the required grid spacing is $\\sim 0.05$--$0.1\\lambda$ which is larger than $\\Delta x$ by an order of magnitude. Thus, by choosing such a large grid spacing (and correspondingly, a large time step) in the wave zone, we can significantly save the grid number for covering the large computational region as well as computational costs.  Due to this reason, the AMR algorithms are employed by many numerical relativity groups now (e.g., \\cite{NSNS5,NSNS6,BHBH1,BHBH6}), which have provided a variety of numerical results recently.  Motivated by the facts mentioned above, we have developed a new code in which an AMR algorithm is implemented, named SACRA ({\\bf S}imul{\\bf A}tor for {\\bf C}ompact objects in {\\bf R}elativistic {\\bf A}strophysics) \\footnote{SACRA is named after Sakura in Japanese (cherry blossom in English)}. This code can evolve not only BH-BH binaries but also NS-NS and BH-NS binaries with a variety of EOSs. In SACRA, the Einstein equations are solved in a similar AMR technique to that adopted in Ref.~\\cite{BHBH6}. Namely, we adopt a fourth-order finite differencing scheme for spatial derivatives and a fourth-order Runge-Kutta scheme for integration forward in time. For the AMR algorithm, six buffer zones are prepared at the refinement boundaries and for the interpolation at the refinement boundaries, fifth-order Lagrangian interpolation scheme in space and second-order Lagrangian interpolation scheme in time are adopted. For simplicity, the size and the grid spacing of computational domain for each refinement level are fixed, although the computational domain can move with the compact objects. We find that this scheme is so stable that we do not have to introduce the Kreiss-Oliger-type dissipation which is often necessary in some AMR codes. For solving the hydrodynamic equations, we adopt a high-resolution central scheme proposed by Kurganov and Tadmor \\cite{KT} with a third-order interpolation for reconstructing the fluid flux at cell interfaces.  For implementing the AMR algorithm, six buffer zones are also prepared as in the gravitational field. Fifth-order and second-order Lagrangian interpolations are basically adopted in space and in time, respectively, although a limiter function is applied in the time interpolation for a region where fluid variables vary steeply. We also find that with this scheme, a stable longterm evolution is feasible for NS-NS and BH-NS binaries.  The paper is organized as follows. In Sec.~\\ref{sec:form}, we briefly describe the basic equations, the gauge conditions, the methods for extracting gravitational waves, and the quantities used in the analysis for the numerical results.  We describe an AMR scheme which we employ in SACRA in Sec.~\\ref{sec:AMR}.  In Sec.~\\ref{sec:res}, numerical results for the simulation of BH-BH, NS-NS, and BH-NS binaries are presented separately. The simulations were performed for a variety of grid resolutions and grid structures.  Convergence of numerical results shows validity of our code. Section \\ref{sec:summary} is devoted to a summary. Throughout this paper, we adopt the geometrical units in which $G=c=1$ where $G$ and $c$ are the gravitational constant and the speed of light.  Latin and Greek indices denote spatial components ($x, y, z$) and space-time components ($t, x, y, z$), respectively: $r \\equiv \\sqrt{x^2+y^2+z^2}$. $\\delta_{ij}(=\\delta^{ij})$ denotes the Kronecker delta. ",
        "conclusions": "\\label{sec:summary}  We have reported our new numerical relativity code, named SACRA, in which an AMR algorithm is implemented. In this code, the Einstein evolution equations are solved in the BSSN formalisms with a fourth-order spatial finite-differencing scheme, the hydrodynamic equations are solved by a third-order high-resolution central scheme, and the time integration is done in the fourth-order Runge-Kutta scheme. Both $F_i$-type and $\\tilde \\Gamma^i$-type BSSN formalisms are implemented. In both cases, $W =e^{-2\\phi}$ is evolved instead of evolving $\\phi$. This enables us to adopt grid-center-grid coordinates.  \\subsection{Technical Points and Issues for the Future}  To check feasibility of SACRA, we performed simulations for coalescence of BH-BH, NS-NS, and BH-NS binaries. All the simulations were performed on personal computers using at most 5 GBytes memory. The required CPU time is at most 1 month even for the best-resolved runs. For simulating BH-BH binaries, we employed the same initial conditions as those adopted by Buonanno et al. \\cite{BHBH12}. Our results agree with theirs in a reasonable manner except for a slight disagreement possibly associated with the difference in choices of gauge conditions and numerical scheme. We also show that our code can follow inspiraling BH-BH binaries at least for about 4.5 orbits even in the absence of dissipation term such as Kreiss-Oliger-type dissipation term. This implies that even in the AMR code, the dissipation term is not always necessary, if appropriate schemes for interpolation and extrapolation are employed for the procedures at refinement boundaries.  Our numerical results for BH-BH binaries indicate that for accurately evolving final 2.5 orbits before the merger, relatively small number of grid points is sufficient. The orbit of the BH is computed accurately and, as a result, gravitational waveforms are computed with a small phase error. In the present simulations, the used memory is at most 3 GBytes, and a personal computer of 4 GBytes memory is sufficient for accurate evolution of the final phase of BH-BH binaries.  By contrast, numerical results, in particular the merger time, depend strongly on the grid resolution, grid structure, and gauge condition for evolving $\\approx 4.5$ inspiral orbits. The estimated phase error in gravitational waveforms for such cases is about $40m_0$ even in the finest-resolution simulation in this paper. For obtaining convergent results within the phase error of, say, $10m_0$ for the whole evolution, the grid resolution has to be finer by a factor of $\\sim 2$. However, we note the following: Numerical results for the final state of the BH formed after merger do not depend on the grid resolution as strongly as the merger time and gravitational wave phase.  It should be noted that for determining the final state of the BH within 1\\% error, it is not necessary to take a high-grid resolution. We find that the present choice is appropriate.   We find that the merger time and gravitational wave phase could depend on spatial gauge conditions.  The reason for this is explained as follows: The physical grid spacing and grid structure depend on the spatial gauge condition, in particular, around BHs. Thus, the magnitude of numerical dissipation also depends on the spatial gauge and may be reduced for a simulation performed with an appropriate choice for the spatial gauge, even if the same grid structure is employed. Therefore, an appropriate choice of the spatial gauge condition may reduce computational costs, and a careful choice is required.  We also show that our code can evolve NS-NS and BH-NS binaries. Numerical results obtained by SACRA agree with those in the previous simulations, if we resolve the NSs and BHs by approximately the same accuracy. However, the computational cost is at most 5\\% of the previous uni-grid simulations and robustness of the AMR scheme is confirmed. Simulations with much better accuracy than those in the previous simulations can be performed by less computational costs. Because we performed the simulations for a wide range of grid resolutions, we can also estimate the magnitude of the phase error of gravitational waveforms in the present and previous numerical results \\cite{BHNS2} in an inexpensive computational cost.  We followed inspiral phase of BH-NS binaries for a long time ($\\sim 4$ orbits) for the first time. In the best-resolved simulation, the inspiral orbit up to the onset of tidal disruption is followed for about 3.7 orbits.  Subsequent merger and ringdown phases are also computed well for producing gravitational waveforms. However, we find that the prepared quasicircular initial condition has a large eccentricity, and the inspiral orbit is highly eccentric for the first $\\sim 2$ orbits, although the eccentricity for a few orbits just before the merger is reduced by the emission of gravitational waves. To perform a realistic simulation for the inspiral phase with small eccentricity, it is necessary to improve the initial condition (see, e.g., Ref.~\\cite{FKPT} for a method). This is an issue for the future.  We compare the duration spent in the inspiral phase obtained by numerical simulations with that predicted by the Taylor T4 formalism for BH-BH and BH-NS binaries.  For longterm runs with the merger time $\\agt 500M_0$, the merger time determined by extrapolation of the numerical results agree with the prediction by the Taylor T4 formalism within an error of $\\sim 10\\%$. This makes us reconfirm that the Taylor T4 formalism provides a good semi-analytical estimate for the time spent in the inspiral phase.  We also find that the Taylor T4 formalism always provides an overestimated value of the merger time for NS-NS and BH-NS binaries. The reason for this overestimation is that in this formalism, tidal effects of NSs, which accelerate the infalling process to merger, are not included. Nevertheless, the error is not extremely large because tidal effects play a crucial role only for close orbits. Therefore, for validating a numerical result, it is useful to compare the merger time with the result derived by the Taylor T4 formalism.  We find that the convergence of the merger time for NS-NS binaries is relatively slow. For this case, the evolution of NSs in the late inspiral phase depends on the effects of tidal deformation of each NS, which in general shortens the merger time.  Thus, to accurately determine the orbital evolution, the tidal deformation of each NS has to be followed accurately in hydrodynamics. The degree of tidal deformation is in general larger near the surface of the NS because the tidal force is approximately proportional to the distance from the center of each NS. In our AMR scheme, the grid resolution around the surface region is not as high as that in the central region. Consequently, the tidal deformation is not followed as accurately as that in the central region. A simple way to overcome this problem is to resolve the surface region as accurately as the central region, i.e., to cover each NS in the finest level. However, doing this in our present scheme is computationally expensive because we have to choose a large value of $N$ for the finest level. There may be a better grid structure to overcome this problem, e.g. to change the cube size in each refinement level. Improving our AMR scheme is an issue in the next step. A completely alternative possibility is to employ a different hydrodynamic scheme which is less dissipative. Improving this scheme is also an issue in the future.  We note that the final state of the BH and surrounding disk after merger of NS-NS binaries do not depend on the grid resolution as strongly as the merger time.  This property is the same as that in the case of BH-BH binaries. Thus, for studying the final state, the present choice of the grid resolution is acceptable.   We check whether or not the conservation relations of energy and angular momentum denoted by Eqs. (\\ref{cons}) and (\\ref{consJ}) or by Eqs. (\\ref{conE}) and (\\ref{conJ}) hold. The energy conservation holds within $\\sim 1\\%$ error irrespective of the binary components for the best-resolved run. The error of angular momentum conservation is larger: The error is $\\sim 3\\%$--5\\%. The resulting total energy and angular momentum of BHs are always smaller than the values predicted by the conservation relations, and hence, numerical dissipation is the most likely source of the error. The error size for the angular momentum conservation may not be negligible, in particular, for studying disk formation around the BH formed after merger. In our results, the disk mass is likely to be underestimated. Indeed, the result for the disk mass in this paper does not agree with the previous result of a BH-NS binary \\cite{BHNS2}.  In the previous result, the angular momentum conservation holds in a much better manner. Thus, the small disk mass in the present results might be partly due to the spurious loss of angular momentum \\cite{footend}.   Another possible drawback in our present AMR scheme is that we might not be able to accurately follow material that spreads around the BH after tidal disruption of the NS.  The reason is that a large fraction of material escapes from the finest level soon after the onset of tidal disruption.  The motion of such material orbiting the BH is located at relatively coarser levels and hence it may not be followed accurately. The material, which forms a spiral arm around the BH, subsequently falls into the BH in a short time scale in the present result. This may be in part due to the fact that its angular momentum is spuriously dissipated. In the present simulations, we found that the resulting mass of accretion disk is much smaller than $10^{-3}M_*$ for $q \\approx 0.33$ and $M_{\\rm NS}/R_{\\rm NS}=0.145$. This result totally disagrees with our previous results \\cite{BHNS2,footBHNS} as mentioned above. Note that the evolution of binaries up to tidal disruption agrees well indicating that the grid structure is appropriate at least up to the onset of tidal disruption.  This suggests that the grid structure in our AMR code might not be well-suited only for following the material orbiting the BH of a distant orbital separation. To improve this situation, it may be necessary to prepare a fine grid which covers a larger region around the BH.  To perform such a simulation, it will be necessary to change the grid structure, e.g. to increase the grid number for the finer levels while fixing that for the coarser levels. Such improvement of our current AMR scheme is an issue in the next step.  \\subsection{Comparison of Numerical Results for Three Types of Binaries}  We performed simulations for three types of binaries. Because of the presence of strong equivalence principle, the orbital evolution and gravitational waveforms in the inspiral phase with a large orbital separation depend very weakly on the components of the binaries.  By contrast, the final outcome and gravitational waveforms in the merger phase depend strongly on the components. As already found in the previous studies (e.g., \\cite{BHBH12}), we found that after merger of slowly spinning two equal-mass BHs, a rotating BH with spin $\\approx 0.7$ is formed.  However, the magnitude of the spin parameter is much higher for a BH formed after merger of NS-NS binaries: The present results show that the spin is $\\sim 0.8$--0.85. This disagreement comes primarily from the difference in amplitude of gravitational waves emitted in the final merger phase. In the case of BH-BH binaries, the BHs can have a closer orbit than the NSs because the BHs are more compact. As a result, gravitational waves are significantly emitted in the final inspiral orbit.  In addition, the quasinormal mode oscillation of fundamental $l=m=2$ mode is excited significantly in the merger phase. Indeed, the gravitational wave amplitude is as high as that emitted at the last inspiral orbit (cf. Fig. \\ref{FIG4}). By these gravitational wave emissions, the angular momentum is significantly dissipated in the final phase. By contrast, in the case of NS-NS binaries, the merger sets in at a relatively distant orbit because NSs are not as compact as BHs, and moreover, the quasinormal mode is not excited as significantly as in the case of BH-BH binaries because of smaller degree of nonaxisymmetric deformation of the spacetime curvature at the merger.  Because of the difference in amplitude of ringdown gravitational waveforms, the property of gravitational waveforms in the final merger phase depends strongly on the binary components. As mentioned above, the amplitude of ringdown gravitational waves is as high as that in the last inspiral phase for the merger of BH-BH binaries.  By contrast, the amplitude is $\\sim 10\\%$ as high as that in the last inspiral phase for the merger of NS-NS binaries. Thus, the wave amplitude quickly decreases in this case.  We also study the merger of BH-NS binaries. In the present paper, we focus on the case that the NS is tidally disrupted before it is swallowed by the companion BH.  In this case, the quasinormal mode is not significantly excited as in the case of NS-NS binaries, and hence, the amplitude of ringdown gravitational waves is also much smaller than that in the last inspiral orbit. However, this may not be always the case. If the mass ratio, $q(=M_{\\rm NS}/M_{\\rm BH})$, is small enough, the NS will not be tidally disrupted before swallowing by the BH. In such case, a quasinormal mode may be excited significantly at a moment that the NS falls into the BH. This topic should be investigated in the future work.  As summarized in this section, gravitational waveforms at merger phase depend strongly on the binary components. This makes us reconfirm that gravitational waves at merger phase will carry information about the properties of binary components. As reviewed in Sec. I, a number of simulations have been performed in the past decade. However, there are a huge parameter space for which numerical study has not been done yet, in particular for NS-NS and BH-NS binaries. Obviously, further study is required. Our new code SACRA will be able to make a contribution to this purpose."
    },
    "0806/0806.2001_arXiv.txt": {
        "abstract": "In this paper, we review the Billion Galaxy Survey that will be carried out at radio--optical wavelengths to micro--nanoJansky levels with the telescopes of the next decades. These are the Low-Frequency Array, the Square Kilometer Array and the Large Synoptic Survey Telescope as survey telescopes, and the Thirty Meter class Telescopes for high spectral resolution+AO, and the James Webb Space Telescope (JWST) for high spatial resolution near--mid IR follow-up. With these facilities, we will be addressing fundamental questions like how galaxies assemble with super-massive black-holes inside from the epoch of First Light until the present, how these objects started and finished the reionization of the universe, and how the processes of star-formation, stellar evolution, and metal enrichment of the IGM proceeded over cosmic time.  We also summarize the high-resolution science that has been done thus far on high redshift galaxies with the Hubble Space Telescope (HST). Faint galaxies have steadily decreasing sizes at fainter fluxes and higher redshifts, reflecting the hierarchical formation of galaxies over cosmic time. HST has imaged this process in great structural detail to z\\cle 6. We show that ultradeep radio-optical surveys may slowly approach the natural confusion limit, where objects start to unavoidably overlap because of their own sizes, which only SKA can remedy with HI redshifts for individual sub-clumps. Finally, we summarize how the 6.5 meter James Webb Space Telescope (JWST) will measure first light, reionization, and galaxy assembly in the near--mid-IR. ",
        "introduction": "For this review paper, I was asked to review the ``Billion Galaxy Survey'' that will be carried out at radio--optical wavelengths to micro--nanoJansky levels with the telescopes of the next decades. These facilities are, for instance, the Low-Frequency Array (LOFAR, R\\\"ottgering et al. 2005), the Square Kilometer Array (SKA; Schilizzi\\ etal 2004), and the Large Synoptic Survey Telescope (LSST; Tyson 2007) --- which will be used as survey telescopes at radio optical wavelengths --- and the Thirty Meter Telescopes (TMT; \\eg Nelson \\etal 2006) --- including \\eg the Giant Segmented Mirror Telescope, Giant Magellan Telescope, plus the EU Extremely Large Telescope ---which will be used for high spectral resolution+ adaptive optics follow-up, as well as the James Webb Space Telescope (Mather and Stockman, 2000), which will provide high spatial resolution near--mid-IR imaging and low-resolution spectroscopy. With a combination of these facilities available in the next decade, we will be addressing fundamental questions like:  \\n \\bul (1) How do HI clouds at z\\cge 6 assemble over cosmic time into the giant spiral and elliptical galaxies seen today?  \\n \\bul (2) How and why did the (dwarf dominated) galaxy luminosity function (LF) and mass function evolve with epoch?  \\n \\bul (3) In the context of the galaxy formation--AGN paradigm, how did supermassive black hole (SMBH) growth keep up with the process of galaxy assembly?  \\n \\bul (4) How does the central accretion disk feed the SMBH, and how are radio jets and lobes in radio galaxies and quasars produced as a result?  \\n \\bul (5) How did AGN feedback control the bright-end evolution of the galaxy LF, and how did SN feedback shape the faint-end of the LF from z\\cge 6 to z=0?  These are some of critical science drivers for radio and optical telescopes of the next decade. Before we give a preview of possible answers to these questions, we will first briefly consider how the Hubble Space Telescope (HST) has revolutionized the topic of galaxy assembly in the last decade. One of the remarkable discoveries by HST was that the numerous faint blue galaxies are in majority late-type (Abraham \\etal\\ 1996, Glazebrook \\etal\\ 1995, Driver \\etal\\ 1995) and small (Odewahn \\etal\\ 1996, Pascarelle \\etal\\ 1996) star-forming objects. These are the building blocks of the giant galaxies seen today. By measuring their distribution over rest-frame type versus redshift, HST has shown that galaxies of all Hubble types formed over a wide range of cosmic time, but with a notable transition around redshifts    \\ve  \\noindent\\begin{minipage}[b]{0.74\\txw} \\psfig{file=windhorst.fig1.eps,angle=0,width=0.73\\txw,height=0.575\\txh} \\end{minipage} \\begin{minipage}[b]{0.2525\\txw} \\noindent{\\footnotesize {\\bf Fig. 1} \\ Galaxy sizes vs. B$_{Vega}$ or \\JAB-mag from the RC3 to the HUDF limit. Short dashed lines indicate survey limits for the HDF (black), HUDF (red), and JWST (orange): the point-source sensitivity is horizontal and the SB-sensitivity has slope=+5 mag/dex. Broken long-dashed pink lines indicate the natural confusion limit, below which objects begin to overlap due to their own sizes. Red and green lines indicate the expectations at faint fluxes of the {\\it non-evolving} median size for RC3 elliptical and spiral galaxies, respectively (Odewahn \\etal\\ 1996). Orange and black squares indicate hierarchical size simulations (Kawata \\etal 2003). Note that most galaxies at \\JAB\\cge 28 mag are expected to be smaller than the HST and JWST diffraction limits (\\ie r$_{hl}$\\cle 0\\arcspt 1).\\\\  \\n \\\\  } \\end{minipage}   \\bn z$\\simeq$0.5--1.0 (\\eg Driver \\etal\\ 1998, Elmegreen \\etal\\ 2007). This was done through HST programs like the Medium-Deep Survey (Griffiths \\etal\\ 1994), GOODS (Giavalisco \\etal\\ 2004), GEMS (Rix \\etal\\ 2005), and COSMOS (Scoville \\etal\\ 2007). Subgalactic units rapidly merged from the end of reionization to grow bigger units at lower redshifts (Pascarelle \\etal\\ 1996). Merger products start to settle as galaxies with giant bulges or large disks around redshifts z$\\simeq$1 (Lilly \\etal\\ 1998, 2007). These evolved mostly passively since then, resulting in giant galaxies today, possibly because the epoch-dependent merger rate was tempered at z\\cle 1 by the extra expansion induced by $\\Lambda$ (Cohen \\etal\\ 2003, Ryan \\etal\\ 2007). To avoid caveats from the morphological K-correction (Giavalisco \\etal\\ 1996, Windhorst \\etal\\ 2002), galaxy structural classification needs to done at rest-frame wavelengths longwards of the Balmer break at high redshifts (Taylor-Mager \\etal\\ 2007). JWST will make such studies possible with 0\\arcspt 1--0\\arcspt 2 FWHM resolution at observed near--IR wavelengths (1--5 \\mum), corresponding to the restframe optical--near-IR at the median redshift of faint galaxies (\\zmed$\\simeq$1--2; Mobasher \\etal\\ 2007). ",
        "conclusions": "HST has led the study of galaxy assembly, showing that galaxies form hierarchically through repeated mergers with sizes growing steadily over time as r$_{\\rm hl}$(z) $\\propto$ r$_{\\rm hl}$(0)$\\cdot$(1+z)$^{\\rm -s}$ and s$\\simeq$ 1. The Hubble sequence thus gradually emerges at z\\cle 1--2, when the epoch-dependent merger rate starts to wind down. The global onset of Pop II-star dominated dwarf galaxies ended the process of reionization at z$\\simeq$6. High resolution rest-frame UV--optical imaging of high redshift galaxies is best done from space, because faint galaxies are small (\\rhl\\cle 0\\arcspt 15), while the ground-based sky is too bright and the PSF not stable enough to obtain good high-resolution images at faint fluxes (AB\\cge 27 mag). JWST will extend these studies into the epoch of reionization and First Light, and trace galaxy SED's in the restframe-optical for z\\cle 20.  Such surveys will slowly approach the natural confusion limit (Fig. 1), where objects at the faintest (nJy) fluxes start to unavoidably blend with their neighbors, not because of instrumental resolution, but because of their own intrinsic sizes. SKA will be critical to help disentangle the non-negligible fraction of such objects at $\\mu$Jy--nJy levels, and provide unique HI-redshifts for each component. For this, it will need to have beam-sizes as small as 0\\arcspt 3 FWHM (see Fig. 1). In order to get the next radio facilities it wants, the science community will need to unify behind current \\& future radio facilities, such as SKA and LOFAR --- building on proto-types that are currently being build (see this Vol.) --- and define the essential synergy of SKA with other future facilities, such as the JWST, LSST, and the TMT's.     \\begin{theacknowledgments} This work was supported by HST grants from STScI, which is operated by AURA for NASA under contract NAS 5-26555, and by NASA JWST grant NAG 5-12460. Other JWST studies are at:\\ www.asu.edu/clas/hst/www/jwst/ . \\end{theacknowledgments}"
    },
    "0806/0806.0232_arXiv.txt": {
        "abstract": "A bosonic dark matter model is examined in detail via high-resolution simulations. These bosons have particle mass of order $10^{-22}eV$ and are non-interacting. If they do exist and can account for structure formation, these bosons must be condensed into the Bose-Einstein state and described by a coherent wave function.  This matter, also known as \\emph{Fuzzy Dark Matter} \\citep{hu00}, is speculated to be able, first, to eliminate the sub-galactic halos to solve the problem of over-abundance of dwarf galaxies, and, second, to produce flat halo cores in galaxies suggested by some observations.  We investigate this model with simulations up to $1024^3$ resolution in an 1 $h^{-1}Mpc$ box that maintains the background matter density $\\Omega_m=0.3$ and $\\Omega_\\Lambda=0.7$. Our results show that the extremely light bosonic dark matter (ELBDM) can indeed eliminate low-mass halos through the suppression of short-wavelength fluctuations, as predicted by the linear perturbation theory. But to the contrary of expectation, our simulations yield singular cores in the collapsed halos, where the halo density profile is similar, but not identical, to the NFW profile \\citep{nfw97}. Such a profile arises regardless of whether the halo forms through accretion or merger.  In addition, the virialized halos exhibit anisotropic turbulence inside a well-defined virial boundary.  Much like the velocity dispersion of standard dark matter particles, turbulence is dominated by the random radial flow in most part of the halos and becomes isotropic toward the halo cores.  Consequently the three-dimensional collapsed halo mass distribution can deviate from spherical symmetry, as the cold dark matter halo does. ",
        "introduction": "Observations of low surface brightness galaxies and dwarf galaxies indicate that the cores of galactic halos have shallow density profiles \\citep{dal97,swt00} instead of central cusps predicted by cold dark matter (CDM) \\citep{nfw97}. In addition, the number density of dwarf galaxies in Local Group turns out to be an order of magnitude fewer than that produced by CDM simulations \\citep{kkvp99}.  These two features cast doubt on the validity of standard CDM.  There have been at least three different solutions proposed to resolve these problems: (1) warm dark matter, (2) collisional dark matter and (3) fuzzy dark matter.  Warm dark matter can suppress small-scale structures by free streaming.  It seems to be able to both solve the over-abundance problem of dwarf galaxies and the singular core problem. In this model the flat core is embedded within a radius a couple of percents of the virial radius \\citep{jing01,colins08}, and the core smoothly connects to the NFW profile \\citep{nfw97} outside. However, this modification may generally adversely affect structures of somewhat larger scales \\citep{hu00}, despite that fine tuning of the thermal velocity of dark matter particles may still be able to have the larger scale structures consistent with observations \\citep{abazajian06, viel08}.  For collisional dark matter, the halo core can be flattened and dwarf galaxies destroyed, and N-body simulations confirm this conjecture \\citep{spst00}.  But simulations also show that very frequent collisions can yield even more singular cores than the standard collisionless CDM does \\citep{yoshida00}. This opposite behavior is indicative of the requirement of fine tuning for collisional parameters.  The third solution to the problem is to treat dark matter as an \\emph{extremely light bosonic dark matter} (ELBDM) or \\emph{fuzzy dark matter} \\citep{pr90,sin94,hu00}. Axion has been thought to be a candidate of light bosonic dark matter. But for the light dark matter to erase the singular galactic core and suppresses low-mass halos, the particle mass must be far smaller than axion (m$\\sim$ $10^{-22}$ eV), so low that the uncertainty principle operates on the astronomical length scale.  Much like axions, the ELBDM is in a Bose-Einstein condense state produced in the early universe. These extremely light particles share the common ground state and is described by a single coherent wave function. Its de-Broglie wavelength is comparable to or even somewhat smaller than the Jean's length \\citep{dw97}, where the quantum fluctuation provides effective pressure against self-gravity.  Several previous works have pondered on such an idea or its variance \\citep{sin94,hu00,slopez03}, in which the wave mechanics is described by the Schr\\\"{o}dinger-Poisson equation with Newtonian gravity or by the Klein-Gordon equation with gravity.  The Schr\\\"{o}dinger-Poisson system addresses the scale-free regime of quantum mechanics, where the Jean's length is a dynamical running parameter.  On the other hand, the Klein-Gordon system makes use of the Compton wavelength as a natural length scale to create the flat core in a halo.  Widrow \\& Kaiser conducted simulations for the two-dimensional Schr\\\"{o}dinger-Poisson system to approximate the standard collisionless cold dark matter \\citep{wk93}.  In the 2D case, the $1/r$ gravitational potential is replaced by $log(r)$, and the 2D force law in their simulation becomes of longer range than it actually is in 3D.  Due to the lack of 3D numerical simulations, some authors resort to spherical symmetry \\citep{sin94,slopez03} or even 1D \\citep{hu00} to study this problem. These simplifications may not capture what actually results in a 3D system with realistic initial conditions. In particular, the existence of a flattened core has been derived or inferred from these previous works of 1D system or with spherical symmetry.  In this paper we report high-resolution fully 3D simulations for this problem. Surprisingly, our simulations reveal that the singular cores of bound objects remain to exist even when the core size is much smaller than the Jean's length.  In Sec.2, we provide an explanation for the possible existence of the Bose-Einstein state for the extremely low mass bosons under investigation here. We then discuss two different representations of ELBDM and the evolution of linear perturbations for the two representations. In Sec.3, the numerical scheme and initial condition are described. We present the simulation results in Sec.4. In Sec.5, we look into the physics of collapsed cores with detailed analyses from different perspectives.  Finally the conclusion is given in Sec.6. In the Appendix, we present results of 1D and 2D simulations and demonstrate singular cores do not occur in 1D and 2D cases. ",
        "conclusions": "As far as we know of, this work presents the first result for the study of Bose-Einstein condensate under self-gravity via high resolution ($1024^{3}$ grids) simulation. \\citet{hu00} conjectured that if the dark matter is ELBDM, it can solve the long standing problem of far too many low-mass halos present in the standard CDM simulations, and also explains the existence of flat cores in some galaxies. In this work, we confirm that low mass halos are indeed suppressed by quantum stress even when the small-scale fluctuations are abundant in the initial power spectrum. This result is a consequence of long-time linear suppression of the small-scale modes. We also find, from our simulations of different grid resolutions, that collapsed halos develop singular cores regardless of the halo formation processes. All these runs produce convergent density profiles. Our $1024^{3}$ highest resolution run gives singular density profiles similar to what standard CDM simulations produce.  In retrospect this singular-core result may not be too surprising, as it arises from an almost scale-free Schr\\\"{o}dinger-Poisson system. This system is not exactly scale free because there exists a Jeans length for fluctuations that are small in amplitude.  However, when the local density much exceeds the background density, the latter becomes locally ill-defined, the Jeans length no longer has any physical significance, and the Schr\\\"{o}dinger-Poisson system becomes locally scale free. Being locally scale-free, the system develops singularities within a finite time.  By contrast, the conservation of phase-space density in classical particle dynamics precludes the space density of standard dark matter particles from developing any singularity \\citep{chiu97,dal01}, and explains the existence of a flat core in the warm dark matter model.  Note that such a phase-space constraint does not exist for nonlinear wave dynamics; one example of this nature is a system described by the nonlinear Schr\\\"{o}dinger equation with attractive self-interaction \\citep{sulem99}.  Most recent observations of rotation curve in low-surface-brightness galaxies indicate inconclusive results, as far as the existence of singular halo core is concerned. Some galaxies are claimed not to possess singular halo cores, and some are if non-circular motion is taken into consideration. Among those that do, many possess concentration parameters inconsistent with the constraint given by $\\Lambda$CDM cosmology \\citep{swt03,zaku06,kuz06,kuz08}. Given the present status of observations, if galaxies indeed contain singular cores, ELBDM will likely be the only viable candidate for the dark matter that, on one hand, permits the galactic-scale, NFW-like halo cores, and on the other hand suppresses the sub-galactic low-mass halos."
    },
    "0806/0806.0468_arXiv.txt": {
        "abstract": "These lectures take a look at how observations with adaptive optics (AO) are beginning to influence our understanding of active galactic nuclei (AGN). By focussing on a few specific topics, the aim is to highlight the different ways in which enhanced spatial resolution from AO can aid the scientific analysis of AGN data. After presenting some background about how AO works, I will describe a few recent observations made with AO of QSO host galaxies, the Galactic Center, and nearby AGN, and show how they have contributed to our knowledge of these enigmatic objects. ",
        "introduction": "\\label{dav:intro}  Adaptive Optics (AO) serves two purposes. It is an enabling technology for even more complex instrumentation such as IR/optical interferometry. But it is also an important technique in its own right. At near-infrared wavelengths (1--5\\,$\\mu$m) it allows one to improve the spatial resolution of ground-based telescopes by an order of magnitude from 0.5--1$^{\\prime\\prime}$ to the diffraction limit, which for an 8-m class telescope in the H-band (1.6\\,$\\mu$m) is about 40\\,mas. AO has applications in many areas of astronomy because it not only increases a telescope's sensitivity for unresolved sources, but it also reduces crowding problems in dense fields, and sharpens our view of the morphologies and kinematics of extended objects. In these lectures, the impact that adaptive optics is having on studies of Active Galactic Nuclei (AGN) is discussed with particular reference to three topics that span scales from $<1$\\,pc in the innermost region of our Galactic Center, through 10\\,pc scales in nearby AGN, to kpc scales in QSO host galaxies at high redshift. By studying these different aspects and combining what we can learn from different spatial scales, we can gain a more holistic view of the structures comprising AGN and the physical processes governing how they are fuelled. ",
        "conclusions": ""
    },
    "0806/0806.0374_arXiv.txt": {
        "abstract": "{The internal dynamics of ultra-compact dwarf galaxies (UCDs) has attracted increasing attention, with most of the UCDs studied to date located in the Virgo cluster.} % {Our aim is to perform a comprehensive census of the internal dynamics of UCDs in the Fornax cluster, and to shed light on the nature of the interface between star clusters and galaxies. } {We obtained high-resolution spectra of 23 Fornax UCDs with $-10.4>M_V>-13.5$ mag ($10^6<{\\rm M/M_{\\sun}<10^8}$), using FLAMES/Giraffe at the VLT. This is the largest homogeneous data set of UCD internal dynamics assembled to date. We derive dynamical M/L ratios for 15 UCDs covered by HST imaging.}% {In the M$_V$-$\\sigma$ plane, UCDs with $M_V<-12$ mag are consistent with the extrapolated Faber-Jackson relation for luminous elliptical galaxies, while most of the fainter UCDs are closer to the extrapolated globular cluster (GC) relation. At a given metallicity, Fornax UCDs have, on average, M/L ratios lower by 30-40\\% than Virgo UCDs, suggesting possible differences in age or dark matter content between Fornax and Virgo UCDs.  For our sample of Fornax UCDs we find no significant correlation between M/L ratio and mass.  We combine our data with available M/L ratio measurements of compact stellar systems with $10^4<{\\rm M/M_{\\sun}}<10^8$M, and normalise all M/L estimates to solar metallicity.  We find that UCDs (M$\\gtrsim$2$\\times 10^6$M$_{\\sun}$) have M/L ratios twice as large as GCs (M$\\lesssim$2$\\times 10^6$M$_{\\sun}$). We argue that dynamical evolution has probably had only a small effect on the current M/L ratios of objects in the combined sample, implying that stellar population models tend to under-predict dynamical M/L ratios of UCDs and over-predict those of GCs. Considering the scaling relations of stellar spheroids, we find that UCDs align well along the 'Fundamental Manifold'.  UCDs can be considered the small-scale end of the galaxy sequence in this context. The alignment for UCDs is especially clear for $r_e\\gtrsim 7$pc, which corresponds to dynamical relaxation times that exceed a Hubble time. In contrast, globular clusters exhibit a broader scatter and do not appear to align along the manifold.} {We argue that UCDs are the smallest dynamically un-relaxed stellar systems, with M$\\gtrsim$2$\\times 10^6$M$_{\\sun}$ and 7$\\lesssim$${\\rm r_e/pc}$$\\lesssim$100. Future studies should aim at explaining the elevated M/L ratios of UCDs and the environmental dependence of their properties.} ",
        "introduction": "\\label{intro} In recent years, significant effort has been devoted to studying the internal dynamics of extragalactic compact stellar systems in the mass regime of massive globular clusters and ultra-compact dwarf galaxies ($10^6<M/M_{\\sun}<10^8$) (Drinkwater et al. 2003, Martini \\& Ho 2004, Ha\\c{s}egan et al. 2005, Maraston et al. 2004, Rejkuba et al. 2007, Evstigneeva et al. 2007, Hilker et al. 2007). A compilation of the available data is presented in Mieske \\& Kroupa (2008) and Dabringhausen, Hilker \\& Kroupa (2008).  A striking outcome of these studies is the finding that the dynamical M/L ratios of massive compact stellar systems are, on average, about two times larger than those of normal globular clusters of comparable metallicity. Several objects have M/L ratios at the limit of, or even beyond, the range predicted by stellar population models assuming canonical IMFs (Ha\\c{s}egan et al. 2005).  Possible explanations for these high M/L ratios include extreme stellar mass functions (Mieske \\& Kroupa 2008; Dabringhausen, Hilker \\& Kroupa 2008) or densely packed dark matter (Goerdt et al. 2008).  The occurence of objects with high M/L ratios is observed to start at $\\sim 2\\times$10$^6$M$_{\\sun}$ (Ha\\c{s}egan et al.  2005, Rejkuba et al.  2007, Mieske \\& Kroupa 2008), coinciding with ``breaks\" in physical size (Ha\\c{s}egan et al. 2005, Mieske et al.  2006, Kissler-Patig et al.  2006) and stellar content (Mieske et al. 2006).  These breaks are consistent with the hypothesis that at $\\sim 2\\times$10$^6$M$_{\\sun}$ ($M_V \\simeq -11$ mag) we observe the transition between simple globular clusters and more complex systems, the UCDs.  Three main formation scenarios have been suggested for UCDs: 1. UCDs are stellar super clusters formed in the tidal arms of violent gas-rich galaxy mergers (Fellhauer \\& Kroupa 2002, 2005). 2. UCDs are tidally stripped compact remnants of nucleated dwarf galaxies (Bassino et al. 1994, Hilker et al. 1999, Bekki et al.  2003, Goerdt et al. 2008). 3. UCDs are genuine compact dwarf galaxies formed from small-scale peaks in the primordial dark matter power spectrum (Drinkwater et al. 2004). In the first case, UCDs are not expected to contain any dark matter. In the second and third case, UCDs would be related to cosmological low-mass dark matter halos and may contain dark matter (Goerdt et al.  2008). The comparably large M/L ratios of 5-10 found for some Virgo UCDs may point towards a cosmological origin. These M/L ratios are similar to the values found for some of the more luminous Local Group dSphs like Sculptor and LeoI (Gilmore et al. 2007), although note that they are still 1-2 orders of magnitude below the M/L values found for the ultra-faint dSph candidates (Gilmore et al. 2007, Simon \\& Geha 2007).  In this paper, we aim to study whether the high M/L ratios are a fundamental trend equally common to all UCDs, or whether environmental variations of the trend exist.  The latter may be expected if competing formation channels dominate in different environments (e.g. Mieske et al. 2006). For both the Virgo cluster (Ha\\c{s}egan et al. 2005, Evstigneeva et al.  2007) and the Centaurus A group (Rejkuba et al. 2007), more than 10 objects in the UCD mass range have measured M/L ratios. For the Fornax cluster, only five sources have high resolution spectroscopy available (Hilker et al.  2007), making it difficult to judge whether differences exist between Fornax and Virgo UCDs. With a comprehensive sample of measurements for Fornax, it will be possible to analyse whether Fornax UCDs extend to such high M/L ratios as Virgo UCDs.  In what follows, we present new measurements of the internal kinematics of 23 compact objects in Fornax (obtained with FLAMES/GIRAFFE at the VLT). We analyse how their M/L ratios relate to predictions from stellar population models, and investigate how they fit into the trend of increasing M/L with mass among compact stellar systems. We also examine how UCDs and globular clusters fit into the broader context of larger and more luminous stellar systems, focusing on the fundamental manifold of stellar spheroids (Zaritsky et al. 2006a, 2006b, 2008) that is an extension of the fundamental plane concept.  The paper is structured as follows: in Sect.~\\ref{data} we present the new spectroscopic data of Fornax UCDs used for the present study. Section~\\ref{reduction} describes the data reduction, including the modelling of the mass distribution. In Sect.~\\ref{results}, the results for the Fornax UCDs are presented and discussed. In Sect.~\\ref{discussion} we combine the Fornax data with other literature results on dynamical M/L ratios of compact stellar systems, and investigate UCDs and GCs in the context of the fundamental manifold. The paper finishes with Summary and Conclusions in Sect.~\\ref{conclusions}.  Throughout this paper we assume a distance modulus to Fornax of (m-M)=31.4 mag (Freedman et al.  2001). ",
        "conclusions": "\\label{conclusions} In this paper we have analysed the internal dynamics of 23 ultra-compact dwarf galaxies in the Fornax cluster. The analysis is based on high-resolution spectroscopy obtained with the FLAMES spectrograph at the VLT. Our targets cover an approximate mass range of $10^6<M<10^8$ M$_{\\sun}$ and a luminosity range $-10.4<M_V<-13.5$ mag, overlapping the bright end of the globular cluster luminosity function. We also compare the dynamical properties of UCDs and GCs, and put them into the context of fundamental scaling relations defined for larger galaxies. We obtain the following results:  \\begin{enumerate}  \\item In the M$_V$-$\\sigma$ plane, we find that UCDs with $M_V<-12$ mag are consistent with the extrapolation of the Faber-Jackson relation for luminous elliptical galaxies. For $M_V>-12$ mag, most objects are located closer to the extrapolation to brighter luminosities of the globular cluster M$_V$-$\\sigma$ relation.   \\item We derive dynamical M/L ratios for those 15 of the 23 UCDs for which HST archival imaging is available, taking into account aperture effects in the spectroscopy (Hilker et al.  2007). Three out of the 15 UCDs have dynamical M/L ratios too high to be explained by canonical stellar populations, but we do not find Fornax UCDs with M/L ratios as extreme as found for some Virgo UCDs (Ha\\c{s}egan et al. 2005). At a given metallicity, Fornax UCDs have on average 30 to 40\\% lower M/L ratios than Virgo UCDs.  \\item We normalise the dynamical M/L ratios of the 15 Fornax UCDs to solar metallicity, using predictions from stellar population models (Bruzual \\& Charlot 2003, Maraston 2005). We find no significant correlation between normalised M/L ratio and mass or relaxation time for our Fornax UCD sample. We do not find a dependence of normalised M/L ratio on projected clustercentric distance.   \\item We add our new measurements for 15 Fornax UCDs to the available data on M/L ratios of compact stellar systems in the broader mass range $10^4<M<10^8$M$_{\\sun}$. We include Galactic globular clusters and UCDs in Virgo, CenA and Fornax. We re-analyse dynamical mass estimates of UCDs in the Virgo cluster (Ha\\c{s}egan et al. 2005) and the CenA group (Rejkuba et al.  2007), using our modelling algorithm (Hilker et al.  2007) to correct for aperture effects in the spectroscopy. The corrections for the global velocity dispersion $\\sigma$ are of the order of 5-10\\%. We also provide previously unpublished M and M/L estimates for 5 CenA compact objects.  \\item We find a clear break in the distribution of normalised M/L ratios at a characteristic mass of $\\simeq$2$\\times 10^6$M$_{\\sun}$, which roughly corresponds to a relaxation time of one Hubble time (see also Dabringhausen, Hilker \\& Kroupa 2008, Mieske \\& Kroupa 2008).  Objects more massive than this limit have normalised M/L ratios twice as large as objects less massive than this limit. In this context we suggest to separate UCDs from GCs by a mass limit of M$\\simeq 2\\times 10^6$M$_{\\sun}$ (see also Ha\\c{s}egan et al. 2005). On average, the M/L ratios of UCDs are 40\\% above the expectations for a 13 Gyr stellar population with canonical IMF, while for GCs they are 40\\% below these expectations.  We find that the M/L ratios estimates of GCs are probably only weakly biased ($\\sim$5\\%) by their dynamical evolution, indicating that stellar population models indeed over-predict M/L ratios for compact stellar systems like GCs.  \\item UCDs extend the 'Fundamental Manifold' in its original formulation (Zaritsky et al.  2006a) by more than a decade in $r_e$ down to $r_e \\sim 5-7$ pc. In this formulation, neither the faintest dwarf spheroidals ($M_V>-9$ mag) nor GCs lie on the manifold. When using also GCs, UCDs and dSphs to define the shape of the FM, UCDs with $r_e \\gtrsim 7$pc and dwarf spheroidals align along the manifold, while GCs and smaller UCDs do not. This characteristic scale of $r_e$$\\simeq$7pc also marks the transition between compact stellar systems with relaxation times below and above a Hubble time.  \\end{enumerate}  \\noindent We suggest a defintion of UCDs as those compact stellar systems with M$\\ge$2$\\times 10^6$M$_{\\sun}$ and 7$\\lesssim$${\\rm r_e/pc}$$\\lesssim$100.  As such, UCDs are the smallest dynamically un-relaxed stellar systems. From their position in the 'Fundamental Manifold' they can be considered the small-scale end of the galaxy sequence.  A key question about UCDs is whether they are of 'cosmological' origin, hence related to compact low-mass dark matter halos.  Their elevated M/L ratios can be interpreted as marking the on-set of dark matter domination in small stellar systems.  However, dark matter can hardly be detected directly, such that observational efforts need to be directed towards verifying/excluding alternative scenarios, such as a variation of the IMF in UCDs (Mieske \\& Kroupa 2008). In parallel, theoretical studies regarding the dynamical evolution of compact stellar systems embedded in dark matter halos are needed for the mass-size regime of UCDs.   In this paper it has been found that Fornax UCDs have 30-40\\% lower dynamical M/L ratios than Virgo UCDs.  A possible explanation for this is that only Virgo UCDs have significant fractions of dark matter. This may be explained by  the dominance of different UCD formation channels in Virgo and Fornax (Mieske et al. 2006). A simple way to test the possibility of different dark matter fractions is to determine the luminosity weighted ages of Fornax and Virgo UCDs.  Younger ages in Fornax UCDs of $\\sim$7 Gyrs would naturally explain the M/L ratio differences and imply similar dark matter fractions as in Virgo.  Together with efforts to constrain the IMF shape in UCDs, such an observational study is the next logical step in understanding the puzzling nature of UCDs."
    },
    "0806/0806.3694_arXiv.txt": {
        "abstract": "{ Gigahertz-Peaked Spectrum (GPS) sources are probably the precursors of local radio galaxies. Existing GPS source samples are small ($<200$). } { It is necessary to extend the availabe sample of the Gigahertz-Peaked Spectrum (GPS) and High Frequency Peaker (HFP) sources in order to study their nature with greater details and higher statistical significance. } { A sample of 214 radio sources, which were extracted from the SPECFIND catalog and show an inverted radio spectrum, were observed quasi-simultaneously at 4.85, 10.45, and 32~GHz with the 100-m Effelsberg radio telescope. Using the VLBA calibrator survey (VCS) we have investigated the parsec-scale morphology of the sources. } { About 45\\% of the sources in our sample are classified as GPS or HFP candidates. We add 65 new GPS/HFP candidates to existing samples. We confirm the expected tendency that HFP are more compact on milliarcsecond scale than the 'classical' GPS sources, which peak at lower frequencies. } { The data mining of the SPECFIND database represents a promising tool for the discovery of new GPS/HFP sources. } ",
        "introduction": "}  GHz-Peaked Spectrum (GPS) sources are powerful ($\\log P_{1.4~{\\rm GHz}} > 25$~W\\,Hz$^{-1}$) and compact ($< 1$~kpc) extragalactic radio sources, which show a convex radio spectrum peaking between 500~MHz and 10~GHz in the observer's frame (for a review see O'Dea~1998). The physical mechanism responsible for the turnover of the spectrum is still unclear with two competing models proposed: the synchrotron self-absorption caused by dense plasma within the source or the free-free absorption caused by a screen external to the source.  High Frquency Peakers (HFP) are radio sources defined via their convex spectrum peaking at frequencies above 5~GHz (Dallacasa et al. 2000). From the anti-correlation found between the turnover frequency and size (O'Dea \\& Baum 1997) HFPs are expected to be smaller and therefore younger radio sources than GPS sources.  GPS sources are associated with either quasars or galaxies. Despite the similar shape of their radio spectrum, these two classes of GPS sources are often considered to be different. Torniainen et al. (2005), who studied the long term variability of 35 inverted-spectrum sources, concluded that genuine quasar-type GPS sources are rare. They found a large number of highly variable blazar sources that can have a convex spectrum peaking at high frequencies (up to $\\sim$100~GHz) during flaring events occurring in the radio jets.  The nature of GPS sources is still under debate. Two possible scenarios have been put forth: (i) the `frustration' scenario, according to which the small size and the inverted spectrum are caused by a dense environment that prohibits the source from growing larger (e.g., Gopal-Krishna \\& Wiita 1991), (ii) the `youth' scenario, suggesting that the GPS sources represent the young precursors of compact steep spectrum (CSS) sources and extended radio sources (e.g., Mutel \\& Phillips 1988, Fanti et al. 1990, 1995).  There is now a wide consensus that, at least the symmetric GPS sources, represent the early evolutionary stage of the extended radio source population. At this stage the radio emitting region grows and expands within the interstellar medium before plunging into the intergalactic medium to form an FR~II radio source (Fanti et al.~1995, Readhead et al.~1996, Begelman~1996, Snellen et al.~2000a).  The detection of kpc-scale emission associated with a few GPS sources seems to be inconsistent with a recent origin of the radio activity. Such an extended emission is interpreted as a sign of a past nuclear activity. In this case the GPS source, i.e. the galactic nucleus, is at the beginning of a new cycle of activity. Since extended radio emission around GPS sources is a rare phenomenon, Stanghellini et al. (2005) conclude that the time scale between subsequent phases of activity is in general longer than the radiative lifetime of the radio emission from the earlier activity ($\\sim 10^{8}$~yr).  The currently existing GPS sample is limited ($<200$ objects, see, e.g., the recent compilation by Labiano et al.~2007). In order to conduct statistical studies of GPS objects and test whether the number of GPS sources at intermediate redshifts is in accordance with that of local FR~II radio galaxies, it is necessary that the sample of available GPS sources is extended. In particular, this could help to test whether GPS sources are the precursors of the local FR~II radio galaxies.  In this paper, we present results of a search for new GPS/HFP candidates. Since the identification of a radio source as a genuine young GPS source depends on the detailed knowledge of its spectrum, flux density variability, VLBI structure and source identification, we will use in the following the words GPS or HFP source as a synonym for GPS candidate or HFP candidate. In order to measure radio-spectra over a range of high frequencies, which are not affected by non-contemporanous measurements, we observed a sample of 214 objects north of declination $-25^\\circ$ that show an inverted radio spectrum, with quasi-simultaneous flux density measurements at 4.85~GHz, 10.45~GHz, and 32~GHz using the Effelsberg 100-m telescope of the Max-Planck-Institut f\\\"ur Radioastronomie (MPIfR). In Sect.~\\ref{sec:sample} we review the sample selection procedure implemented on the basis of the SPECFIND database, observations and primary data reduction are described in Sect.~\\ref{sec:observations}. We present results of observations in Sect.~\\ref{sec:results}, sources extention and variablility characteristics in Sect.~\\ref{sec:ext_var}, cross-identification and spectral fitting in Sect.~\\ref{sec:fitting}, milliarsecond-scale compactness and morphology in Sect.~\\ref{sec:vlbi}, comparison with other available GPS/HFP samples in Sect.~\\ref{sec:discussion}. We summarize our results in Sect.~\\ref{sec:summary}. ",
        "conclusions": "}  In order to investigate if our objects are already identified as GPS/HFP sources, we cross-identified our sample with existing GPS/HFP samples of Snellen et al. (1998), Stanghellini et al. (1998), Marecki et al. (1999), Dallacasa et al. (2000), Fanti et al. (2001) and Labiano et al. (2007).  Table~\\ref{tab:gpscross} shows the results of this cross-identification.  The sample designations can also be found in column~6 of Table~\\ref{tab:table1} and on top of each plot in Fig.~\\ref{fig:spectra}.  We find 32 previously known GPS/HFP sources in our sample, out of which 26 are classified as GPS/HFP candidates by us. Dallacasa et al. (2000) made quasi-simultaneous multi-frequency observations of HFP sources at the VLA.  We have plotted their flux densities as crosses on the spectra of Fig.~\\ref{fig:spectra}. For the majority of the sources there is a good agreement with the Effelsberg and SPECFIND flux densities.  \\begin{table} \\caption{Cross-identification with existing GPS/HFP samples.} \\label{tab:gpscross} \\[ \\begin{array}{lcc} \\hline\\hline {\\rm Article} & {\\rm Designation}  & {\\rm Number\\ of\\ objects} \\\\ \\hline {\\rm Snellen\\ et\\ al.~(1998)} & {\\rm B} & 3 \\\\ % {\\rm Stanghellini\\ et\\ al.~(1998)} & {\\rm S} & 7 \\\\ % {\\rm Fanti\\ et\\ al.~(2001)} & {\\rm F} & 0 \\\\ % {\\rm Dallacasa\\ et\\ al.~(2000)} & {\\rm D} & 15 \\\\ % {\\rm Marecki\\ et\\ al.~(1999)} & {\\rm M} & 9 \\\\ % {\\rm Labiano\\ et\\ al.~(2007)} & {\\rm L} & 14 \\\\ \\hline \\end{array} \\] \\end{table}  }  A sample of 214 radio sources with inverted spectra and declination $\\delta > -25^\\circ$ was extracted from the SPECFIND catalog (Vollmer et al. 2005a). This catalog contains cross-identifications of radio sources from surveys at different frequencies and combines them into one single radio spectrum per object. To obtain quasi-simultaneous radio spectra, we observed those sources with the MPIfR Effelsberg 100-m radio telescope at 4.85, 10.45, and 32~GHz. All the observed sources were detected at 4.85~GHz, while 209 and 180 sources were detected at 10.45~GHz and 32~GHz, respectively. We expect the fraction of variable or confusing sources in our complete sample to be of $\\sim 30$\\%. On the basis of the performed analysis of continuum radio spectra, we have identified 38 GPS and 53 HFP candidates out of which 65 were previously unknown. An inspection of VCS data shows that 24 out of 53 GPS/HFP sources with available VCS data are resolved at 2.3 and 8.6~GHz. We have confirmed the expected tendency for HFP objects to be highly compact and the GPS ones to show a significantly lower level of compactness at the milliarcsecond scale. This independently supports robustness of our source classifications presented here.   We have a success of $\\sim$45\\% for finding GPS/HFP candidates from a selected SPECFIND catalog sample. Once the SPECFIND catalog is upgraded with the inclusion of more radio catalogs, this method comprises a promising way for future identification of new GPS/HFP sources."
    },
    "0806/0806.1144_arXiv.txt": {
        "abstract": "{GRID-launcher-1.0 was built within the VO-Tech framework, as a software interface between the UK-ASTROGRID and a generic GRID infrastructures in order to allow any ASTROGRID user to launch on the GRID computing intensive tasks from the  ASTROGRID Workbench or Desktop. Even though of general application, so far the Grid-Launcher has been tested on a few selected softwares (VONeural-MLP, VONeural-SVM, Sextractor and SWARP) and on the SCOPE-GRID. ",
        "introduction": "The main goal of the International Virtual Observatory \\cite{IVOA} infrastructure is to provide the community at large with an easy and user friendly access to astronomical data, software tools, and computing facilities. While the first two parts of the task, namely the federation and fusion of heterogeneous data archives and the implementation of flexible data reduction and data analysis tools have been widely addressed and, at least in their fundamental aspects, solved, the possibility to access large distributed computing facilities to perform computing intensive tasks has not yet been satisfactorily answered. The main reasons for this delay being mainly two: the need to match the GRID security requirements and the lack of homogeneity in the definitions of the storage space. The first issue can be easily explained: most users of a specific Virtual Organization (VO) do not possess the personal certificates which are requested to access the GRID or, even when they do have a personal certificate, the computing GRID which they need does not recognize their own certification authority. The second issue arises instead from the fact that the data usually reside locally or in a remote storage space which is not seen by the selected GRID as a storage element (SE).  The lack of an user friendly access to the GRID is a main obstacles against the use of computing intensive data mining methods and tools such as, for instance, Support Vector Machines \\cite{cavuoti} or Probabilistic Principal Surfaces \\cite{dabrusco} on massive data sets (MDS). Therefore, in the framework of the VONeural project, which aims at implementing a package of data mining routines capable to work in a distributed computing environment and on large data sets of high dimensionality, we have implemented and tested a general purpose interface between the UK-ASTROGRID \\cite{astrogrid} and the GRID which solves part of the above quoted problems. We wish to stress that the followed approach is quite general and that GRID-Launcher can be easily adapted to other Virtual Organizations and to any GRID. For testing we used the GRID-SCOPE \\cite{scope} which is part of the recently funded GRID infrastructure for Southern Italy. ",
        "conclusions": ""
    },
    "0806/0806.3225_arXiv.txt": {
        "abstract": "We show that within the inverse seesaw mechanism for generating neutrino masses minimal supergravity is more likely to have a sneutrino as the lightest superparticle than the conventional neutralino. We also demonstrate that such schemes naturally reconcile the small neutrino masses with the correct relic sneutrino dark matter abundance and accessible direct detection rates in nuclear recoil experiments. ",
        "introduction": "Over the last fifteen years we have had solid experimental evidence for neutrino masses and oscillations~\\cite{Maltoni:2004ei}, providing the first evidence for physics beyond the Standard Model. On the other hand, cosmological studies clearly show that a large fraction of the mass of the Universe in dark and must be non--baryonic.  The generation of neutrino masses may provide new insight on the nature of the dark matter~\\cite{Berezinsky:1993fm}. In this Letter we show that in a minimal supergravity (mSUGRA) scheme where the smallness of neutrino masses is accounted for within the inverse seesaw mechanism the lightest supersymmetric particle is likely to be represented by the corresponding neutrino superpartner (sneutrino), instead of the lightest neutralino. This opens a new window for the mSUGRA scenario. Here we consider the implications of the model for the dark matter issue. We demonstrate that such a model naturally reconciles the small neutrino masses with the correct relic abundance of sneutrino dark matter and experimentally accessible direct detection rates. ",
        "conclusions": ""
    },
    "0806/0806.4543_arXiv.txt": {
        "abstract": "We compute the effect of the galactic absorption on AGN emission in a cosmological context by including a physical model for AGN feeding and feedback in a semi-analytic model of galaxy formation. This is based on galaxy interactions as triggers for AGN accretion, and on expanding blast waves as a mechanism to propagate outwards the AGN energy injected into the interstellar medium at the center of galaxies.  We first test our model against the observed number density of AGNs with different intrinsic luminosity as a function of redshift. The model yields a ''downsizing'' behavior in close agreement with the observed one for $z\\lesssim  2$. At higher redshifts, the model predicts an overall abundance of AGNs (including Compton-thick sources) larger than the observed Compton-thin sources by a factor $\\approx 2$ for $z\\gtrsim 2$ and $L_X\\leq 10^{44}$ erg/s. Thus, we expect that at such luminosities and redshifts about $1/2$ of the total AGN population is contributed by Compton-thick sources.  We then investigate the dependence of the absorbing column density $N_H$ associated to cold galactic gas (and responsible for the Compton-thin component of the overall obscuration) on the AGN luminosity and redshift. We find that the absorbed fraction of AGNs with $N_H\\geq 10^{22}$ cm$^{-2}$ decreases with luminosity for  $z\\leq 1$; in addition, the total (integrated over luminosity) absorbed fraction increases with redshift up to $z\\approx 2$, and saturates to the value $\\approx 0.8$ at higher redshifts. Finally, we predict the luminosity dependence of the absorbed fraction of AGNs with $L_X\\leq 3\\,10^{44}$ erg/s to weaken with increasing redshift.  We compare our results with recent observations, and discuss their implications in the context of cosmological models of galaxy formation. ",
        "introduction": "The accretion which built up the supermassive black holes (SMBHs) now hosted in many local galaxies is widely thought to be associated with a sequence of output episodes observed as Active Galactic Nuclei (AGNs). Since only a minority ($\\sim 10^{-2}$, see Richstone et al.  1998) of local galaxies host a currently active AGN, the corresponding  lifetimes are estimated to be close to $\\tau\\sim 10^8$ yrs. Several observations indicate the accretion episodes to be fundamentally related to the galaxy growth; one such indication is provided by the narrow scatter in the observed correlation of the SMBH mass $M_{BH}$ with its host galaxy mass $M$, when the former is in the range $10^7\\leq M_{BH}/M_{\\odot}\\leq 10^9$ (Ferrarese \\& Merritt 2000, Gebhardt et al. 2000). The emissions of AGNs thus may conceivably constitute a probe for the history of accretion and growth of SMBHs, and for its interplay with the galaxy building process.  The co-evolution of galaxies and AGNs and their so called ``downsizing'' (faster evolution for more luminous objects) depends also on feedback between nuclear and other galactic activities. In fact, the density of the high luminosity QSOs is peaked at high redshift and declines strongly toward us; similarly, massive galaxies are characterized by a star formation history peaked at high redshifts. Luminous AGNs are efficient in \"sterilizing\" their massive host galaxies by heating the interstellar matter through winds, shocks, and high energy radiation, see Granato et al. (2004); Murray, Quataert, Thompson (2005); Hopkins et al. (2006); Bower et al. (2006); Menci et al. (2006). Intriguingly, the latter authors found that the bimodal color distribution of galaxies at z$\\gtrsim 1.5$ can only be explained if AGN feedback is considered.  In this picture an AGN phase precedes the phase when a galaxy is caught in a passive state with red optical-UV colors, most of the star-formation having been inhibited by the AGN activity. Indeed, Pozzi et al. (2007) using Spitzer photometry found that that a sample of optically obscured QSO at z=1--2 are mainly hosted by red passive galaxies, suggesting a later stage in their evolution.  On the other hand, at low redshift many weak AGNs have been found in star-forming galaxies (Salim et al. 2007). In these cases feedback from less powerful AGNs (the so-called ''radio mode'') is probably acting to self-regulate accretion and star-formation, and cold gas is left available for both processes for a much longer time (Croton et al. 2006). The same cold gas can intercept the line of sight to the nucleus. Indeed, Compton-thin absorbers (with column densities $N_H\\leq 10^{24}$ cm$^{-2}$) may well be located in the galactic disk (Malkan, Gorjian, Tam 1998; Matt 2000; see also Ballantyne, Everett, Murray 2006). Therefore a natural expectation in this scenario is the fraction of obscured AGNs to be large at low AGN luminosities. It is well known since the pioneering work done with the {\\it Einstein} satellite (Lawrence \\& Elvis 1982) and with optically and radio-selected AGNs (Lawrence 1991) that this fraction strongly decreases with increasing AGN luminosity (see Ueda et al. 2003; La Franca et al. 2005; Gilli, Comastri, Hasinger 2007; Triester, Krolik, Dullemond 2008; Hasinger 2008).  A widely shared view holds that the luminosity dependence of the obscured fraction is related to the energy fed back by the AGNs onto the surrounding gas that constitutes the interstellar medium (ISM). Given that the AGN emission is proportional to the fraction of such a gas available for accretion, a \\textit{positive} correlation between luminosity and absorption would be expected instead in the absence of an energy feedback depleting the ISM after the onset of the AGN activity.  A mounting body of observations cogently indicates that strong nuclear feedback is present in galaxies hosting AGNs (see for a review Elvis 2006 and references therein).  On small (sub-pc) scales, the observed X-ray absorption lines indicate the presence of outward winds with velocities up to some $ 10^4$ km/s (Weymann 1981; Turnshek et al. 1988; Creenshaw et al. 2003; Chartas et al. 2002; Pounds et al. 2003, 2006; Risaliti et al. 2005b).  These likely originate from the acceleration of disk outflows due to the AGN radiation field (Proga 2007 and references therein). On larger scales, broad absorption lines in about 10\\% of optically luminous QSOs indicate fast outflows (up to 30,000 km/s).  Massive (10-50 $M_{\\odot}/yr$) flows of neutral gas with speed $\\sim1000$ km/s are observed through 21-cm absorption of radio-loud AGNs (see Morganti, Tadhunter, Oosterloo 2005), indicating that AGNs have a major effect on the circumnuclear gas in the central kiloparsec region around  AGNs. On even larger scales of some $10^{2}$ kpc, the presence of AGN-induced outflows is revealed by X-ray observations of the intra-cluster medium (see McNamara \\& Nulsen 2007 for a review) showing cavities and expanding shocks with Mach numbers ranging from $\\approx 1.5$ to $\\sim 8$. How the outflows produced in the innermost regions of the active galaxies are transported outwards to affect such large scales is still matter of investigation; buoyant bubbles (see Reynolds, Heinz \\& Begelman 2001, Churazov 2001) and expanding blast waves (see Cavaliere, Lapi \\& Menci 2002; Lapi, Cavaliere \\& Menci 2005) constitute viable mechanisms for such a transport.  Nuclear obscuration is directly linked to AGN feedback, since the same gas (and dust) which feed the AGN output may well be responsible for its obscuration. Therefore, modelling AGN obscuration is essential to connect the observed AGN properties to the accretion history of SMBHs over cosmological time. This is a significant theoretical challenge as it requires not only connecting the AGN evolution to the galaxy formation and growth, but also implementing in a model a self-consistent description of the AGN feedback on the galactic gas. Indeed, very few attempts have been carried out in this direction so far. An attempt to include AGN absorption into an {\\it ab initio} galaxy formation model has been proposed by Nulsen \\& Fabian (2000), by relating both the SMBH fueling and the AGN absorption to the cooling flows associated with the hot gas pervading the growing dark matter haloes, but this assumption did not lead to a full description of the statistical distribution of QSO luminosity.  Here we develop our semi-analytic model of hierachical galaxy formation and AGN evolution (see Menci 2006) to self-consistently include the absorption of AGNs, with the aim of investigating the cosmic evolution of the latter and its dependence on AGN properties like luminosity and redshifts. The model is suited to our scope as it includes a detailed treatment of the feedback on the interstellar gas. Our aim is to investigate the dependence of AGN {\\it absorption} on luminosity $L$ and redshift $z$ arising in hierarchical galaxy formation scenarios that include an effective description of the evolution of AGNs and of their feedback.  The plan of the paper is as follows: in Sect. 2 we describe our model for galaxy formation and the associated AGN evolution. Sect. 3 is focussed on describing our treatment of the AGN feedback onto the interstellar gas and  its effects on the AGN absorption. In Sect. 4 we test our feedback-inclusive model for AGN evolution by comparing its outcomes with the redshift distribution of the number density of AGNs for different X-ray luminosities. Our results on the luminosity and redshift dependence of the absorbed fraction of AGNs are shown and discussed in Sect. 5. Sect. 6 is devoted to summarize our conclusions. ",
        "conclusions": "We have computed the effects of the galactic absorption on AGN emission in a cosmological context, by including a physical model for AGN fueling and feedback into a semi-analytic model of galaxy formation in the concordance cosmology. The model is based on galaxy interactions as triggers for AGN accretion and on expanding blast waves as a mechanism to propagate outwards the AGN energy injected into the interstellar medium at the center of galaxies.  We have shown that an {\\it inverse} dependence of AGN absorption on luminosity (fig. 3) and a {\\it direct} dependence on redshift (fig. 4) is a natural outcome in such a context. The former arises from the faster expansion of blast waves induced by feedback of energetic, luminous AGNs onto the galactic interstellar gas; the rapid sweeping of gas in the inner regions, where the density was initially higher, results in the fast formation of a gas-depleted region with size larger for higher AGN intrinsic luminosities. Qualitatively similar conclusions were reached by Hopkins et al. (2005b) on the basis of dedicated hydrodynamical N-body simulations. On the other hand, the redshift dependence of the AGN absorption is due to the larger amount of cold galactic gas available at high redshift, when the higher densities allowed for fast cooling occurring in galactic haloes. The quantitative {\\it predictions} of our model are consistent with existing observations concerning the fraction of absorbed ($N_H\\geq 10^{22}$ cm$^{-2}$) AGNs as a function of their luminosity and redshift. Our model specifically predicts that for AGNs with $L_X\\leq 3\\,10^{44}$ erg/s the luminosity dependence of the absorbed fraction  weakens with increasing redshift (see fig. 5), while for the brightest objects with $L_X\\gtrsim 3\\,10^{44}$ the absorbed fraction quickly decreases with luminosity for $z\\gtrsim 2.5$ (see fig. 5).  Note that, after averaging over the line of sight, unobscured or mildly obscured AGNs correspond to late stages of the feedback action; in particular, for a given orientation of the line of sight, the observed column density depends on the time elapsed since the start of the blast wave expansion. The faster expansion characterizing the blast wave of luminous AGNs thus corresponds to a {\\it larger probability} to observe them when the blast has already swept out the central regions of the galaxy IGM. This picture constitutes an {\\it extension} of the unified picture for AGNs (see Antonucci 1993) beyond the canonical scheme based on the single orientation parameter, since the absorption properties now depend on the combination of {\\it orientation} and {\\it time} needed to sweep the central regions of the galaxy disk. Note that our picture has also straightforward implications on the connection between star formation and obscuration properties (proposed , since from our results we on average expect larger star formation in heavily obscured objects, while luminous, mildly absorbed AGNs should be generally associated to galaxies with lesser star formation (or  in transition to a passive state). The relevance of evolutionary effects of the kind modeled here in determining the absorption properties of AGNs is supported by several observational works. Stevens et al.  (2005) and Page et al. (2004) find that X-ray obscured QSOs have much higher submillimeter detection rates than X-ray unobscured QSOs, suggesting strong star formation on-going in the host of obscured AGN only.  Sajina et al. (2007) and Martinez-Sansigre et al. (2008) report Spitzer IRS spectra dominated by AGN continuum but showing PAHs features in emission, typical of starforming galaxies, in samples of ULIRGs and radio selected obscured QSOs at z$\\sim2$. Lacy et al.  (2007) find evidence for dust-obscured star formation in type-2 QSOs. From the analysis of the X-ray background, Baalntyne ey al. (2006) argue that the AGN obscuration is connected with the star formation in the host galaxy. Finally, Martinez-Sansigre et al. (2005, 2008) found little or no Lyman-$\\alpha$ emission in a sample of z$>1.7$ obscured QSOs, suggesting large scale (kpc) dust distribution. All these findings are in general agreement with the evolutionary picture, although some disagreeing works argue for a dominance of geometrical effect  (see Triester et al. 2008). A geometrical effect may be constituted by the gravitational bending of the interstellar gas due to the BH, as described in Lamastra et al. (2006); however, except for very massive BHs ($M_{BH}\\sim 10^9\\,M_{\\odot}$), this will affect mainly regions below our resolution scale of 50 $pc$. Another class of geometrical models, based on the luminosity dependence of the sublimation radius of the BH accretion disk,  is that commonly referred to as the ''receeding torus'' picture (Lawrence 1991). These however provide an appreciable luminosity dependence only at very high AGN luminosities $L_x\\gtrsim 10^{44-45}$ erg/s and for dust located close to the sublimation radius ($1-10$ pc). Dust located in galaxy disk can hardly be affected directly by the AGN radiation.  In our model we did not try to model the processes leading to Compton-thick absorption with $N_H\\geq 10^{24}$ cm$^{-2}$, generally thought to be caused by gas directly associated with the central regions of the AGNs. On the other hand, our model provides a hint concerning luminosity and redshift dependence to be expected as for the abundance of Compton-thick sources. Inspection of fig. 2, where we compare the predicted global (including Compton-thick sources) density of AGNs with different {\\it intrinsic} luminosities with data corrected for absorption (but not including Compton-thick sources), shows that our model predicts a number of low/intermediate-luminosity AGNs ($L_X\\leq 10^{44}$ erg/s$^{-2}$) larger than the observed Compton-thin sources by a factor around 2 at $z\\gtrsim 2$. While it is possible that our model overestimates the AGN fueling in this range of $L_X$ and $z$ (e.g., due to its specific modeling of interaction-driven destabilization for cold gas in galactic disks), the above excess could support the view that at such luminosities and redshifts a fraction around 1/2 of the total AGNs are Compton-thick. A complementary process which may explain our overprediction  of low-luminosity AGNs at $z\\gtrsim 2$ is suppression of BH growth in small mass galactic haloes (DM masses $M\\leq 10^9\\,M_{\\odot}$). This may be provided by gravitational-rocket effect on the BHs due to the recoil following the emission of gravitational waves during the coalescence of BH binaries following galaxy mergers (see Madau \\& Quataert 2004). Such a recoil may produce BH velocities of order $10^2$ km/s, sufficient to unbind the hole from galaxies with DM velocity dispersion $\\lesssim 50$ km/s. In this respect, the observational selection of Compton-thick sources by combining mid-infrared to near-infrared and optical photometry of galaxies (Fiore et al. 2008a,b) will provide crucial constrains on the relative role of obscuration and BH depletion in low-mass galactic halos.  Our results are also relevant for constraining the physical mechanisms of AGN feedback onto the interstellar gas, a key issue in recent developments of cosmological galaxy formation models. Indeed, the recent realizations of such models include AGN feedback as the key process to suppress gas cooling in massive galactic haloes. Actually, two kinds of feedback are at present implemented in this context: on the one hand, models based on galaxy interactions as triggers of AGN activity and feedback relate the latter to the bright, accretion phase onto supermassive black holes, i.e., to the same phase corresponding to AGN activity (see Menci et al. 2003, 2007; Di Matteo et al. 2005; Hopkins et al. 2005a); on the other hand, other authors associate the AGN feedback only to a quiescent (so called ''radio'') phase of accretion, characterized by very low accretion rates ($\\lesssim 10^{-2}\\,M_{\\odot}$/yr) and not observable as radiative AGNs (Bower et al. 2006, Croton et al. 2006). Since the latter mode would result into a feedback activity continuing down to low redshift, we expect such models to provide a much milder dependence of AGN absorption on redshift;  so observational results on the high-redshift absorption of AGN emission will constitute an effective test for models of AGN feedback.  Finally, we note that the feedback related to the active AGN phase described here is effective to decrease the galactic gas and the associated absorption mainly at low redshift $z\\lesssim 2$ and for bright ($L_X\\gtrsim 10^{44}$ erg/s) AGNs (see, e.g., Cavaliere \\& Menci 2007). This implies that at high redshifts the effective cooling and the continuous replenishing of galactic gas due to fequent merging events will override the depletion due to the AGN feedback, so the latter can not suppress the early growth of supermassive BH. Indeed, the predicted number density of early ($z\\geq 4$), bright ($L_x\\geq 10^{45.5}$ erg/s) AGNs shown in fig. 2 is consistent with that observed for bright ($M_i\\leq -27.5$) optical QSOs up to $z\\approx 6$ (Hopkins et al. 2006). Note however that such a result does not include the BH spin-up which may occur during the growth due to accretion of gas endowed with angular momentum (for different modeling of such a process see, e.g., Volonteri et al. 2005; King, Pringle \\& Hofman 2007), which in turn may yield larger radiative efficiencies up to values $\\eta\\approx 0.3$ for a dominant fraction of BHs; as noted by the above authors, the lower mass accretion rate related to larger radiation efficiencies may delay the mass assembly of massive BH at early epochs when included in a cosmological model. We shall investigate these issues in a following paper."
    },
    "0806/0806.1366_arXiv.txt": {
        "abstract": "A component of dark energy has been recently proposed to explain the current acceleration of the Universe. Unless some unknown symmetry in Nature prevents or suppresses it, such a field may interact with the pressureless component of dark matter, giving rise to the so-called models of coupled quintessence. In this paper we propose a new cosmological scenario where radiation and baryons are conserved, while the dark energy component is decaying into cold dark matter (CDM). The dilution of CDM particles, attenuated with respect to the usual $a^{-3}$ scaling due to the interacting process, is characterized by a positive parameter $\\epsilon$, whereas the dark energy satisfies the equation of state $p_x=\\omega \\rho_x$ ($\\omega < 0$). We carry out a joint statistical analysis involving recent observations from type Ia supernovae, baryon acoustic oscillation peak, and Cosmic Microwave Background shift parameter to check the observational viability of the coupled quintessence scenario here proposed. ",
        "introduction": "According to Einstein's general theory of relativity, the dynamic properties of a given spacetime are determined by its total energy content. In the cosmological context, for instance, this amounts to saying that to understand the spacetime structure of the Universe one needs to identify the relevant sources of energy and their contributions to the total energy momentum tensor. Matter fields (e.g., baryonic matter and radiation), are obvious sources of energy. Nevertheless, according to current observations, two other components, namely, dark matter and dark energy, whose origin and nature are completely unknown, are governing the late time dynamic properties of the Universe. Although fundamental to our understanding of the Universe, several important questions involving these dark components and their roles in the dynamics of the Universe remain unanswered (see, e.g., \\cite{review} for some recent reviews).  Among these questions, the possibility of interaction in the dark sector (dark matter-dark energy), which gave origin to the so-called models of coupled quintessence, has been largely explored in the literature \\cite{cq,cq2}. These scenarios are based on the premise that, unless some special and unknown symmetry in Nature prevents or suppresses a non-minimal coupling between these components (which has not been found -- see, e.g., \\cite{carroll} for a discussion), such interaction is in principle possible and, although no observational piece of evidence has so far been unambiguously presented, a weak coupling still below detection cannot be completely excluded.  >From the observational viewpoint, these models are capable of explaining the current cosmic acceleration, as well as other recent observational results \\cite{cq}. From the theoretical point of view, however, critiques to these scenarios do exist and are mainly related to the fact that in order to establish a model and study their observational and theoretical predictions, one needs first to specify a phenomenological coupling between the cosmic components.  In this concern, an interesting step towards a realistic interaction law was given recently by Wang \\& Meng in Ref. \\cite{wm} (see also \\cite{alclim05}) in the context of models with vacuum decay,  a class of coupled quintessence in which the dark energy equation of state (EoS) is $w = -1$. Actually, in certain sense, one may say that coupled dark energy or quintessence  models are the natural inheritors of the so-called time-varying $\\Lambda(t)$-cosmologies \\cite{Brons1,Brons2,list,list2}. However, instead of the traditional approach, Refs. \\cite{wm,alclim05} deduced a new interaction law from a simple argument about the effect of the dark energy on the cold dark matter (CDM) expansion rate. The resulting expression is a very general law that has many of the previous phenomenological approaches as a particular case.  \\begin{figure*} \\centerline{\\epsfig{figure=qz.eps,width=2.4truein,height=2.8truein} \\epsfig{figure=qz8.eps,width=2.4truein,height=2.8truein} \\epsfig{figure=qzw.eps,width=2.4truein,height=2.8truein} \\hskip 0.1in} \\caption{$q(z)$ in the scenario of coupled quintessence. {\\bf{a)}} Deceleration parameter as a function of redshift for the phantom case in which $\\omega=-1.2$ and selected values of $\\epsilon$. {\\bf{b)}} The same as in the previous Panel for the quintessence case $\\omega=-0.8$. As discussed in the text, the effect of a positive $\\epsilon$ parameter is to decrease the value of $q(z)$, therefore increasing the value of the transition redshift $z_t$. $q(z)$ versus redshift for the specific value of $\\epsilon=0.1$ and selected values of EoS parameter.} \\label{fig:qzw} \\end{figure*}  In this paper, we extend the arguments of Refs. \\cite{wm,alclim05} to a dark energy/dark matter interaction, where the dark energy component is described by an equation of state $p_x = \\omega \\rho_x$ ($w < 0$), and explore theoretical and observational consequences of a new scenario of coupled quintessence. Differently from other interacting quintessence models, we do not consider interaction between the dark sector and the baryonic content of the Universe. We also emphasize that this process of interaction is completely different from the physical point of view from unification scenarios of the dark sector, an idea that has been widely discussed in the recent literature \\cite{chap}.  We have organized this paper as follows. In Sec. II the interaction law and the basic field equations of the model are presented. The influence of the dark energy-dark matter coupling on the epoch of cosmic acceleration is also discussed. In order to test the observational viability of the model, Sec. III presents a statistical analysis involving the most recent type Ia supernovae (SNe Ia) data \\cite{davis,wv07,rnew,Astier06,Ries07}, observations of the baryon acoustic oscillation (BAO) peak (measured from the correlation function of luminous red galaxies) \\cite{bao} and the current estimate of the Cosmic Microwave Background (CMB) shift parameter from WMAP-5 \\cite{Sperg07}. In Sec. IV we end this paper by summarizing our main results. ",
        "conclusions": "The description of the model discussed in the previous Section clearly shows that it comprises a multitude of cosmological solutions. In a model with such a wealth of different possibilities constraints on the parameter space arising from current observational data are likely to rule out many of the possible scenarios (combinations of $\\epsilon$, $w$ and $\\Omega_{dm}$) for the evolution of the Universe. In this Section we investigate such observational constraints by placing cosmological bounds on the parametric spaces $\\epsilon-w$ and $\\epsilon-\\Omega_{dm}$ from statistical analyses involving a large set of cosmological observations. To this end we use the most recent distance measurements to SNe Ia \\cite{davis} and the current estimates of the baryon acoustic oscillations found in the SDSS data \\cite{bao}, as well as, the shift parameter from WMAP observations \\cite{Sperg07}. In our analysis we fix $\\Omega_b=0.0416$ also from WMAP results, a value in good agreement with the constraints derived from primordial nucleosynthesis \\cite{nucleo}. Now, concerning the Hubble parameter, it should be  recalled that the estimates of $H_0$ through different methods fall on the range 62-74 km/s/Mpc with an uncertainty of about $10\\%$ \\cite{H0}.  In what follows, we consider the Hubble Space Telescope ($HST$) Key Project final result, i.e., $h = 0.71\\pm0.08$ \\cite{freedman}, as a Gaussian prior on the Hubble parameter.  \\subsection{SNe Ia}  The predicted distance modulus for a supernova at redshift $z$, given a set of parameters $\\mathbf{s}$, is \\begin{equation} \\label{mag} \\mu_p(z|\\mathbf{s}) = m - M = 5\\mbox{log} d_L + 25, \\end{equation} where $m$ and $M$ are, respectively, the apparent and absolute magnitudes, the complete set of parameters is $\\mathbf{s} \\equiv (H_o, \\Omega_{dm}, \\epsilon, w)$ and $d_L$ stands for the luminosity distance (in units of megaparsecs).    \\begin{figure*} \\centerline{\\epsfig{file=figcont.eps,width=8.5cm,height=6.5cm} \\epsfig{file=figcontwm.eps,width=8.5cm,height=6.5cm}\\hskip 0.1in} \\caption{The results of our statistical analyses. {\\bf a)} Confidence contours at 68.3\\%, 95.4\\% and 99.7\\% in the plane $\\Omega_m$ - $\\epsilon$ from a joint analysis involving SNe Ia + BAO + CMB shift parameter + $H_0$. As discussed in the text, this analysis constrains $\\epsilon$ to values very close to zero ($\\simeq 0.09$ at 3$\\sigma$). {\\bf b)} Same as Panel {\\bf a} for the plane $\\omega$ - $\\epsilon$.} \\label{fig:projhw} \\end{figure*}  We estimate the best fit to the set of parameters $\\mathbf{s}$ by using a $\\chi^{2}$ statistics, with \\begin{equation}\\label{xhi2def} \\chi^{2} = \\sum_{i=1}^{N}{\\frac{\\left[\\mu_p^{i}(z|\\mathbf{s}) - \\mu_o^{i}(z|\\mathbf{s})\\right]^{2}}{\\sigma_i^{2}}}, \\end{equation} where $\\mu_p^{i}(z|\\mathbf{s})$ is given by Eq. (\\ref{mag}), $\\mu_o^{i}(z|\\mathbf{s})$ is the extinction corrected distance modulus for a given SNe Ia at $z_i$, and $\\sigma_i$ is the uncertainty in the individual distance moduli.  In our analysis, we use a combined sample with $N = 192$ SNs also used by Davis {\\it et al.} (2007) \\cite{davis}. This sample consists of the best quality light-curves SNs of Wood-Vasey {\\it et al.} (2007) \\cite{wv07}, which are 60 ESSENCE supernovae \\cite{wv07}, 57 SNLS supernovae \\cite{Astier06}, and 45 nearby supernovae. We also include, as in \\cite{davis}, 30 new released SNe Ia, classified as ``gold'' supernovae by Riess {\\it et al.} (2007) \\cite{Ries07}.  \\subsection{BAO}  The Baryon Acoustic Oscillations (BAO) given by the acoustic oscillations of baryons in the primordial plasma, leave a signature on the correlation function of galaxies as observed by Eisenstein {\\it et al.} (2005) \\cite{bao}. This signature furnishes a standard rule which can be used to constrain the following quantity \\cite{bao}: \\begin{eqnarray} {\\cal{A}} \\equiv \\frac{\\Omega_m^{1/2}}{ {{\\cal{H}}(z_{\\rm{*}})}^{1/3}}\\left[\\frac{1}{z_{\\rm{*}}} \\Gamma(z_*)\\right]^{2/3}  = 0.469 \\pm 0.017, % \\label{A} \\end{eqnarray} where ${\\cal{H}}$ is given by Eq. (\\ref{friedmann}), $z_*=0.35$ is a typical redshift of the SDSS sample, and $\\Gamma(z_*)$ is the dimensionless comoving distance to the redshift $z_*$. As has been shown in Ref. \\cite{baode}, this quantity can be used for models which do not have a large contribution of dark energy at early times.  \\subsection{CMB shift parameter}  A useful quantity to characterize the position of the CMB power spectrum first peak is the shift parameter, which is given, for a flat Universe, by \\cite{efstathiou}: \\begin{equation} {\\cal R}=\\sqrt{\\Omega_m}\\int_0^{z_r}\\frac{dz}{{\\cal H}(z)} = 1.71 \\pm 0.03\\;, \\end{equation} where $z_r = 1089$ is the recombination redshift and the value for ${\\cal R}$ above is calculated from the MCMC of the WMAP 3-yr in the standard flat $\\Lambda$CDM model \\cite{ElgMult}.  As mentioned above, we also include a Gaussian prior on $h$, as given by the final results of HST Key Project \\cite{freedman}. Thus, in our statistical analysis we minimize the following quantity: \\begin{eqnarray} \\label{chi22} \\chi^2 &=& \\sum_{i=1}^{192}\\left(\\frac{\\mu_{obs,i}-\\mu_{th,i}}{\\sigma_{\\mu,i}}\\right)^2 + \\left(\\frac{{\\cal A}-0.469}{0.017}\\right)^2 \\nonumber \\\\ &+& \\left(\\frac{{\\cal R}-1.71}{0.03}\\right)^2 + \\left(\\frac{h-0.72}{0.08}\\right)^2\\;. \\label{chi2} \\end{eqnarray}    \\subsection{Results}  In Figure (3) we show the main results of our statistical analyses. As usual, the total likelihood is written as ${\\cal L}\\propto e^{-\\chi^2/2}$, where $\\chi^2$ is given by Eq. (\\ref{chi22}) . By marginalizing ${\\cal L}$ over the EoS parameter $\\omega$, we can quantify how much the plane ($\\epsilon$-$\\Omega_m$) can be constrained by the data. The contour levels for this analysis are shown on Figure (\\ref{fig:projhw}a). At 68.3\\%, 95.4\\% and 99.7\\% c.l., we have found, respectively, $$ \\Omega_m=0.269^{+0.028 +0.047 +0.066}_{-0.026 -0.042 -0.058} $$ and $$ \\epsilon=0.000^{+0.027 +0.057 +0.088}_{-0.000 -0.000 -0.000}\\;, $$ with the relative $\\chi^2/\\nu\\simeq1.03$, where $\\nu$ is the number of degrees of freedom. These results are much more constraining than those obtained in Ref. \\cite{alclim05} for the case $w = -1$, as we have used more recent CMB and SNe Ia data. While in the above reference the bounds on the interacting parameter were $\\epsilon=0.06\\pm0.10$ at $95.4\\%$ c.l., we have found, at the same level, $\\epsilon=0.000^{+0.057}_{-0.000}$, which clearly constrains this parameter to values very close to the standard non-interacting case ($\\epsilon=0$).   In Figure (\\ref{fig:projhw}b) we show the plane ($\\omega-\\epsilon$) when the total likelihood is marginalized over the density parameter $\\Omega_m$. For this analysis, we have found $$ \\omega=-1.006^{+0.117+0.188+0.258}_{-0.119-0.205-0.296}\\;, $$ whereas the bounds for $\\epsilon$ are very similar to those found in the previous analysis (Fig. 3a). Clearly, the standard $\\Lambda$CDM is preferred by this analysis, although much space is left for an EoS distinct from $-1$. The so-called phantom models ($\\omega<-1$) are slightly more favoured by this analysis than quintessence ($\\omega>-1$) scenarios."
    },
    "0806/0806.0683_arXiv.txt": {
        "abstract": "In this paper we consider $\\phi^2$ scalar field potential as a candidate to dark matter. If it is an ultralight boson particle, it condensates like a Bose-Einstein system at very early times and forms the basic structure of the Universe. Real scalar fields collapse in equilibrium configurations which oscillate in space-time (oscillatons).The cosmological behavior of the field equations are solved using the dynamical system formalism. We use the current cosmological parameters as constraints for the present value of the scalar field and reproduce the cosmological predictions of the standard $\\Lambda$CDM model with this model. Therefore, scalar field dark matter seems to be a good alternative to cold dark matter nature. ",
        "introduction": "Scalar fields are one of the most interesting and most mysterious fields in theoretical physics. Fundamental scalar fields are needed in all unification's theories, however, there are not experimental evidence of its existence. From the standard model of particles which needs the Higgs boson, until the superstring theory which contains the dilaton, passing throught the Kaluza-Klein and the Brans-Dicke theories or throught the inflationary model, scalar fields are necessary fields. Doubtless, if they exist, they have some features which make them very special.  The \\textit{Scalar Field Dark Matter} (SFDM) model paradigm has been constructed step by step. One of the first suggestions that a (complex) scalar field could contribute to structure formation of the Universe was given by \\citet{press} and \\citet{madsen}. Nevertheless, complex scalar fields were used before as matter candidates as boson stars by \\citet{ruffini} (for a recent introduction to boson stars, see for example \\cite{paco06}). One of the first candidates to be scalar field dark matter is the axion, one of the solutions to the strong-CP problem in QCD (see an excellent review in \\cite{Kolb}). Essentially, the axion is a scalar field with mass restricted by observations to $\\sim10^{-5}$eV, which has its origin at $10^{-30}$ seconds after the big bang, when the energy of the Universe was $10^{12}$GeV. This candidate is till now one of the most accepted candidates for the nature of dark matter, if its abundance is about $10^9$ particles per cubic centimetre.  The first in suggesting that a dark halo could be a Bose-Einstein condensate were \\citet{sang-jin} and \\citet{ji} who used the weak field limit to show that a Bose-Einstein Condensate (BEC) with several nodes can fit the rotation galaxy curves with a very good accuracy. Further investigations on this direction were performed by \\citet{joe-weon}, where they incorporated $\\phi^4$ interactions to the scalar field potential and used the Gross-Pitaevskii equation instead of the Schr\\\"odinger one \\citep{joe}. Nevertheless, \\citet{seidel,seidela} showed that when the whole BEC is in the ground state, many nodes in Einstein-Klein-Gordon fields are unstable, since they evolve into the 0-node solution after a while (for a clear explanation to this point see also \\citep{fcoluis}). Thus, the static solutions  given by \\citet{sang-jin, ji, joe-weon} are expected to be unstable.  Later on, \\citet{peebles99} proposed that a scalar field driven by inflation can behave as a perfect fluid and can have interesting observational consequences in structure formation. Besides that, they performed a sound waves analysis of this hypothesis giving some qualitative ideas for the evolution of these fields and called it fluid dark matter \\citep{peea, peeb}. Independently and in an opposite way, \\citet{matosguzman99} proposed a scalar field coming from some unify theory can condensate and collapse to form haloes of galaxies. Very early, this scalar field behaves as a perfect fluid, however its ultralight mass causes that the bosons condensate at very high temperature and collapse in a very different way as the fluid dark matter of \\citet{peebles99} did. They were able to fit reasonably rotation curves of some galaxies using an exact solution of the Einstein equations with an exponential potential \\citep{matosguzman99, gusmanmatos00, bernal}. The first cosmological study of the SFDM was performed in \\citet{matosurena00a, matosurena00b} where a $cosh$ scalar field potential was used. The cosmology reproduces all features of the $\\Lambda$ Cold Dark Matter ($\\Lambda$CDM) model in the linear regime of perturbations.  On the other hand, \\citet{julien} and \\citet{arbey} used a complex scalar field with a quartic potential $m^2\\phi\\phi^{\\dag}+\\lambda(\\phi\\phi^{\\dag})^2$ and solved perturbations equations (weak field limit approximation) to fit the rotational curves of dwarf galaxies with a very good accuracy, provided that $m^4/\\lambda\\sim50-75$ eV$^4$.  The importance of scalar fields in the dark sector has been increased, for instance, several authors have investigated the unification of dark matter and dark energy in a single scalar field \\citep{pad,arbey,bert}. Recently \\citet{liddleurena06,liddleurena08} proposed that the landscape of superstring theory can provide the Universe with a $\\phi^2+\\Lambda$ scalar field potential. Such scalar field can inflate the Universe during its early epoch, after that, the scalar field can decay into dark matter. The constant $\\Lambda$ can be interpreted as the cosmological one. This model could explain all unknown components of the Universe in a simple way. Another interesting model in order to explain the scalar fields unification, dark sector and inflation, is using a complex scalar field protected by an internal symmetry \\citep{perez}.  In the present work the main idea is that if scalar fields are fundamental, they live as unified fields in some very early moment at the origin of the Universe. As the Universe expands, the scalar fields cool together with the rest of the particles until they decouple from the rest of the matter. After that, only the expansion of the Universe will keep cooling the scalar fields. If the scalar field fluctuation is under the critical temperature of condensation, the object will collapse as a BEC. After inflation, primordial fluctuations cause that the scalar fields collapse and form haloes of galaxies and galaxy clusters. The cooling of scalar fields continue till the fluctuation separates from the expansion of the Universe.  In this work we study the most simple model of SFDM, using a $\\phi^2$ scalar field potential. In sections \\ref{sec:estadistica} and \\ref{sec:einstein} we review the statistic of a boson gas to condensate and form a BEC, focusing in the necessary features for the BEC to form a halo of a galaxy and integrate the Einstein equations with a BEC matter. In section \\ref{sec:cosmologia} we transform the Einstein field equations into a dynamical system, then we numerically integrate them and look for the atractor points. We give some conditions on how these field equations can give the right behavior to reproduce the Universe we observed. Finally, in section \\ref{sec:conclusiones} we conclude that this SFDM model could explain the dark matter of the Universe. ",
        "conclusions": "\\label{sec:conclusiones}  SFDM has provided to be an alternative model for the dark matter nature of the Universe. We have shown that the scalar field with a ultralight mass condensates very early in the Universe and generically form BEC's with a density profile which is very similar as that of the CDM model, but with a almost flat central density profile, as it seems to be in LSB and dwarf galaxies. This fact can be a crucial difference between both models. If the flat central density is no confirmed in galaxies, we can rule out the SFDM model, but if this observation is confirmed, this can be a point in favor of the SFDM model. We also show that the $1/2m^2\\phi^2$ potential and the $V_0[\\cosh(\\kappa\\lambda\\phi)-1]$ model are in fact the same. They have the same predictions, a control variable which determines the behavior of the model, given naturally the right expected cosmology and the same cosmology as the CDM model. This implies that the differences between both models, the CDM and SFDM ones, is in the non linear regime of perturbations. In this way they form galaxies and galaxy clusters, specially in the center of galaxies where the SFDM model predicts a flat density profile. If the existence of supersymmetry is confirmed, the DM supersymmetric particles would be observed by detectors and they would have the right mass, DM density and coupling constant, therefore the SFDM model can be ruled out. However, if these observations are not confirmed, the SFDM is an excellent alternative candidate to be the nature of the DM of the Universe."
    },
    "0806/0806.2709_arXiv.txt": {
        "abstract": "Using the high resolution spectra secured with the Nasmyth Echelle Spectrograph NES of the 6 meter telescope we analysed the hydrogen-deficient binary star KS\\,Per. The atmospheric parameters derived are: $T_{\\rm eff}$\\,=\\,9500$\\pm$300\\,K, $\\log g$\\,=\\,2.0$\\pm$0.5, and $\\xi_{\\rm t}$\\,=\\,9.5$\\pm0.5$\\,km\\,s$^{-1}$. The hydrogen deficiency is H/He\\,=\\,3$\\cdot10^{-5}$, iron abundance is reduced by 0.8\\,dex, nitrogen abundance is very high [N/Fe]\\,=\\,1.4, but carbon and oxygen abundances are low. The star luminosity is $\\log L/L_{\\odot}$\\,=\\,3.3. A complex absorption and emission structure of the Na\\,I D doublet was revealed. We suggest that the emission component forms in the circumbinary gaseous envelope. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0806/0806.1546_arXiv.txt": {
        "abstract": "\\noindent We show in this paper that it is possible to attain very high, {\\it including observable}, values for the level of non-gaussianity $f_{NL}$ associated with the bispectrum $B_\\zeta$ of the primordial curvature perturbation $\\zeta$, in a subclass of small-field {\\it slow-roll} models of inflation with canonical kinetic terms. Such a result is obtained by taking care of loop corrections both in the spectrum $P_\\zeta$ and the bispectrum $B_\\zeta$. Sizeable values for $f_{NL}$ arise even if $\\zeta$ is generated during inflation. Five issues are considered when constraining the available parameter space:  1. we must ensure that we are in a perturbative regime so that the $\\zeta$ series expansion, and its truncation, are valid. 2. we must apply the correct condition for the (possible) loop dominance in $B_\\zeta$ and/or $P_\\zeta$. 3. we must satisfy the spectrum normalisation condition.  4. we must satisfy the spectral tilt constraint. 5. we must have enough inflation to solve the horizon problem. ",
        "introduction": "Since COBE \\cite{cobe} discovered and mapped the anisotropies in the temperature of the cosmic microwave background radiation \\cite{smooth}, many balloon and satellite experiments have refined the measurements of such anisotropies, reaching up to now an amazing combined precision. The COBE sequel has continued with the WMAP satellite \\cite{wmap} which has been able to measure the temperature angular power spectrum up to the third peak with unprecedent precision \\cite{hinshaw}, and increase the level of sensitivity to primordial non-gaussianity in the bispectrum by two orders of magnitude compared to COBE \\cite{komatsu2,komatsu1}.  The next-to-WMAP satellite, PLANCK \\cite{planck}, whose launch is programmed for October 2008, is expected to precisely measure the temperature angular power spectrum up to the eighth peak \\cite{planck1}, and improve the level of sensitivity to primordial non-gaussianity in the bispectrum by one order of magnitude compared to WMAP \\cite{komatsu}.  Because of the progressive improvement in the accuracy of the satellite measurements % described above, it is pertinent to study cosmological inflationary models that generate significant (and observable) levels of non-gaussianity. An interesting way to address the problem involves the $\\delta\\textit{N}$ formalism \\cite{starobinsky,sasaki2,lyth4}, which can be employed to give the levels of non-gaussianity $f_{NL}$ \\cite{lyth2} and $\\tau_{NL}$ \\cite{boubekeur1,alabidi2} in the bispectrum $B_\\zeta$ and trispectrum $T_\\zeta$ of the primordial curvature perturbation $\\zeta$ respectively. Such non-gaussianity levels are given, for slow-roll inflationary models, in terms of the local evolution of the universe under consideration, as well as of the $n$-point correlators, evaluated a few Hubble times after horizon exit, of the perturbations $\\delta\\phi_{i}$ in the scalar fields that determine the dynamics of such a universe during inflation.  In the $\\delta\\textit{N}$ formalism for slow-roll inflationary models, the primordial curvature perturbation $\\zeta(\\textbf{x},t)$ is written as a Taylor series in the scalar field perturbations $\\delta\\phi_{i}(\\textbf{x},t_\\star)$, evaluated a few Hubble times after horizon exit, \\begin{eqnarray} \\zeta(t,\\textbf{x})&=&\\sum_{i}N_{i}(t)\\delta\\phi_{i}(\\textbf{x},t_\\star) - \\sum_{i}N_{i}(t) \\langle\\delta\\phi_{i}(\\textbf{x},t_\\star)\\rangle + \\nonumber \\\\ &&+\\frac{1}{2}\\sum_{ij}N_{ij}(t)\\delta\\phi_{i}(\\textbf{x},t_\\star)\\delta\\phi_{j}(\\textbf{x},t_\\star)-\\frac{1}{2}\\sum_{ij}N_{ij}(t) \\langle\\delta\\phi_{i}(\\textbf{x},t_\\star)\\delta\\phi_{j}(\\textbf{x},t_\\star)\\rangle + \\nonumber \\\\ &&+\\frac{1}{3!}\\sum_{ijk}N_{ijk}(t)\\delta\\phi_{i}(\\textbf{x},t_\\star)\\delta\\phi_{j}(\\textbf{x},t_\\star)\\delta\\phi_{k}(\\textbf{x},t_\\star) - \\frac{1}{3!}\\sum_{ijk}N_{ijk}(t)\\langle\\delta\\phi_{i}(\\textbf{x},t_\\star)\\delta\\phi_{j}(\\textbf{x},t_\\star)\\delta\\phi_{k}(\\textbf{x},t_\\star)\\rangle + \\nonumber \\\\ &&+...\\;, \\end{eqnarray} where the brackets mean spatial averages, $N$ is the amount of inflation (or number of e-folds) from a bit later (in Hubble times) than the time when the cosmologically relevant scales exit the horizon and until the time at which one wishes to calculate $\\zeta$, and $N_{i}\\equiv\\frac{\\partial N}{\\partial\\phi_{i}}$, $N_{ij}\\equiv\\frac{\\partial^{2}N}{\\partial\\phi_{i}\\partial\\phi_{j}}$, and so on. It is in this way that the correlation functions of $\\zeta$ (for instance, $\\langle\\zeta_{\\bf k_{1}}\\zeta_{\\bf k_{2}}\\zeta_{\\bf k_{3}}\\rangle$) can be obtained in terms of series, as often happens in Quantum Field Theory where the probability amplitude is a series whose possible truncation at any desired order is determined by the coupling constants of the theory. A highly relevant question is that of whether the series for $\\delta N$ converges in cosmological perturbation theory and whether it is possible in addition to find some quantities that determine the possible truncation of the series, which in this sense would be analogous to the coupling constants in Quantum Field Theory.  In general such  quantities % will depend on the specific inflationary model; the series then cannot be simply truncated at some order until one is sure that it does indeed converge, and besides, one has to be careful not to forget any term that may be leading in the series even if it is of higher order in the coupling constant. This issue has not been investigated in the present literature, and generally the series has been truncated to second- or third-order neglecting in addition terms that could be the leading ones \\cite{sasaki2,lyth2,boubekeur1,alabidi2,zaballa,alabidi1,vernizzi,battefeld,yokoyama1,yokoyama2,seery3,byrnes2}.  The most studied and popular inflationary models nowadays are those of the slow-roll variety with canonical kinetic terms \\cite{liddle,lyth5,lyth6}, because of their simplicity and because they easily satisfy the spectral index requirements for the generation of large-scale structures. One of the usual predictions from inflation and the theory of cosmological perturbations is that the levels of non-gaussianity in the primordial perturbations are expected to be unobservably small when considering this class of models \\cite{zaballa,vernizzi,battefeld,yokoyama1,seery3,maldacena,seery7,seery5,li,seery4}\\footnote{One possible exception is the two-field slow-roll model analyzed in Ref. \\cite{alabidi1} (see also Refs. \\cite{bernardeu1,bernardeu2}) where {\\it observable, of order one, values for} $f_{NL}$ are generated for a reduced window parameter associated with the initial field values when taking into account only the tree-level terms in both $P_\\zeta$ and $B_\\zeta$. However, such a result seems to be incompatible with the general expectation, proved in Ref. \\cite{vernizzi}, of $f_{NL}$ being of order the slow-roll parameters, and {\\it in consequence unobservable}, for two-field slow-roll models with separable potential when considering only the tree-level terms both in $P_\\zeta$ and $B_\\zeta$.  The origin of the discrepancy could be understood by conjecturing that the trajectory in field space, for the models in Refs. \\cite{alabidi1,bernardeu1,bernardeu2}, seems to be sharply curved, being quite near a saddle point; such a condition is required, according to Ref. \\cite{vernizzi}, to generate $f_{NL} \\sim \\mathcal{O}(1)$. \\label{laila}}. This fact leads us to analyze the cosmological perturbations in the framework of first-order cosmological perturbation theory.  Non-gaussian characteristics are then suppressed since the non-linearities in the inflaton potential and in the metric perturbations are not taken into account.  The non-gaussian characteristics are actually present and they are made explicit if second-order \\cite{lyth3} or higher-order corrections are considered.  The whole literature that encompasses the slow-roll inflationary models with canonical kinetic terms reports that the non-gaussianity level $f_{NL}$ is expected to be very small, being of the order of the slow-roll parameters $\\epsilon_i$ and $\\eta_i$, ($\\epsilon_i, |\\eta_i| \\ll 1$) \\cite{vernizzi,battefeld,yokoyama1,maldacena,seery7}. These works have not taken into account either the convergence of the series for $\\zeta$ nor the possibility that loop corrections dominate over the tree level ones in the $n$-point correlators. Our main result in this paper is the recognition of the possible convergence of the $\\zeta$ series, and the existence of some ``coupling constants'' that determine the possible truncation of the $\\zeta$ series at any desired order. When this situation is encountered in a subclass of small-field {\\it slow-roll} inflationary models with canonical kinetic terms, the one-loop corrections may dominate the series when calculating either the spectrum $P_\\zeta$, or the bispectrum $B_\\zeta$. This in turn {\\it may generate sizeable and observable levels of non-gaussianity} in total contrast with the general claims found in the present literature.  The layout of the paper is the following:  in Section \\ref{descriptors} we consider the quantities that describe the statistical properties encoded in any probability distribution function; theoretical explanations as well as observational constraints for $\\zeta$ are given.  Section \\ref{dNf} is devoted to the issue of the $\\zeta$ series convergence and loop corrections in the framework of the $\\delta N$ formalism, as well as to the presentation of the current knowledge about primordial non-gaussianity in slow-roll inflationary models.  A particular subclass of small-field slow-roll inflationary models is the subject of Section \\ref{model} as it is this subclass of models that generate significant levels of non-gaussianity.  The available parameter space for this subclass of models is constrained in Section \\ref{rest} by taking into account some observational requirements such as the COBE normalisation, the scalar spectral tilt, and the minimal amount of inflation. Another requirement of methodological nature,  the possible tree-level or one-loop dominance in $P_\\zeta$ and/or $B_\\zeta$, is considered in this section.  The level of non-gaussianity $f_{NL}$ in the bispectrum $B_\\zeta$ is calculated in Section \\ref{endcal} for models where $\\zeta$ is % generated during inflation; a comparison with the current literature is made. Section \\ref{seccou} is devoted to central issues in the consistency of the approach followed such as satisfying necessary conditions for the convergence of the $\\zeta$ series and working in a perturbative regime. Finally in Section \\ref{concl} we conclude. The professional reader who is already familiarized with the present ideas on the cosmological non-gaussianity may skip Sections \\ref{descriptors} and \\ref{dNf}, leaping directly to the new material starting from Section \\ref{model}. As regards the level of non-gaussianity $\\tau_{NL}$ in the trispectrum $T_\\zeta$, it will be studied following the sequence of ideas presented above in a companion paper \\cite{cogollo}. ",
        "conclusions": "\\label{concl} Observational cosmology is in its golden age: current satellite and balloon experiments are working extremely well \\cite{wmap,hinshaw}, dramatically improving the quality of data \\cite{komatsu1}.  Moreover, foreseen experiments \\cite{planck,planck1} will take the field to a state of unprecedent precission where theoretical models will be subjected to the most demanding tests. Given such a state of affairs, it is essential to study the higher order statistical descriptors for cosmological quantities such as the primordial curvature perturbation $\\zeta$, which give us information about the non-gaussianity in their corresponding probability distribution functions.  $\\zeta$ and its associated non-gaussianity depend on the specific inflationary model that describes the dynamics of the early Universe, the slow-roll class of inflationary models with canonical kinetic terms being the most popular and studied to date. Inflationary models of the slow-roll variety predict very well the spectral index in the spectrum $P_\\zeta$ of $\\zeta$ but, if the kinetic terms are canonical, they seem to generate unobservable levels of non-gaussianity in the bispectrum $B_\\zeta$ and the trispectrum $T_\\zeta$ of $\\zeta$ making them impossible to test against the astonishing forthcoming data. Where does this conclusion come from? The answer relies on careful calculations of the levels of non-gaussianity $f_{NL}$ and $\\tau_{NL}$ by making use of the $\\delta N$ formalism \\cite{vernizzi,battefeld,yokoyama1,seery3}. In this framework, $\\zeta$ is given in terms of the perturbation $\\delta N$ in the amount of expansion from the time the cosmologically relevant scales exit the horizon until the time at which one wishes to calculate $\\zeta$.  Due to the functional dependence of the amount of expansion, $\\zeta$ is usually Taylor-expanded (see Eq. (\\ref{Nexp})) and truncated up to some desired order so that $f_{NL}$ and $\\tau_{NL}$ are easily calculated (see for instance Eq. (\\ref{fdnf})). Two key questions arise when noting that it is impossible to extract general and useful information from the $\\zeta$ series expansion in Eq. (\\ref{Nexp}) until one chooses a definite inflationary model and calculates explicitly the $N$ derivatives. First of all, when writing a general expression for $f_{NL}$ or $\\tau_{NL}$ in terms of the $N$ derivatives, how do we know that such an expression is correct if the series convergence has not been examined? Moreover, if the convergence radius of the $\\zeta$ series is already known, why is each term is the $\\zeta$ series supposed to be smaller than the previous one so that cutting the series at any desired order is thought to be enough to keep the leading terms? Nobody seems to have formulated these questions before and, by following a naive line of thinking, $f_{NL}$ and $\\tau_{NL}$ were calculated for slow-roll inflationary models with canonical kinetic terms without checking the $\\zeta$ series convergence and keeping only the presumably leading tree-level terms \\cite{vernizzi,battefeld,yokoyama1,seery3,maldacena,seery7}.  These two questions have been addressed in this paper by paying attention to a particular quadratic small-field slow-roll model of inflation with two components and canonical kinetic terms (see Eq. (\\ref{pot})). Although the non-diagrammatic approach followed in Section \\ref{seccou} to find the necessary condition for the convergence of the $\\zeta$ series in our model might not be applicable to all the cases, we have been able to show that not being careful enough when choosing the right available parameter space could make the $\\zeta$ series, and therefore the calculation of $f_{NL}$ and $\\tau_{NL}$ from the truncated series (e.g. Eq. (\\ref{fdnf})), meaningless. We also have been able to show in our model that the one-loop terms in the spectrum $P_\\zeta$ and/or the bispectrum $B_\\zeta$ of $\\zeta$ could be bigger or lower than the corresponding tree-level terms, but are always much bigger than the corresponding terms whose order is higher than the one-loop order. If both $P_\\zeta$ and $B_\\zeta$ are dominated by the one-loop terms, {\\it a huge} $f_{NL}$ {\\it is generated} which overwhelms the observational constraint, ruling out the model {\\it by an excess} and not by a shortfall. If $B_\\zeta$ is still dominated by the one-loop correction but $P_\\zeta$ is now dominated by the tree-level term, {\\it sizeable and observable values for} $f_{NL}$ {\\it are generated}, so they can be tested against present and forthcoming observational data. Finally, if both $P_\\zeta$ and $B_\\zeta$ are dominated by the tree-level terms, $f_{NL}$ {\\it is slow-roll suppressed} as was originally predicted in Refs. \\cite{vernizzi,battefeld,yokoyama1}.  What these results teach us is that the issue of the $\\zeta$ series convergence and loop corrections is essential for making correct predictions about the statistical descriptors of $\\zeta$ in the framework of the $\\delta N$ formalism, and promising for finding high levels of non-gaussianity that can be compared with observations. In fact, now that we have learned the lesson, the level of non-gaussianity $\\tau_{NL}$ for the same slow-roll model studied here will be the subject of a companion paper \\cite{cogollo}.  \\bigskip  \\subsection*"
    },
    "0806/0806.1770_arXiv.txt": {
        "abstract": "The Australian National University's SkyMapper telescope is amongst the first of a new generation of dedicated wide-field survey telescopes. Featuring a 5.7 deg$^{2}$ field-of-view Cassegrain imager and 268 Mega-pixel CCD array, its primary goal will be to undertake the Southern Sky Survey: a six color ($uvgriz$), six-epoch digital record of the entire southern sky. The survey will provide photometry for objects between 8th and 23rd magnitude with a global photometric accuracy of 0.03 magnitudes and astrometry to 50 mas. In this contribution we introduce the SkyMapper facility, the survey data products and outline a variety of case-studies in stellar astrophysics for which SkyMapper will have high impact. ",
        "introduction": "It is now possible for CCD mosaic imagers on small telescopes to achieve areal coverage that was once solely the domain of photographic Schmidt plates. Digital surveys offer better photometric and astrometric precision and calibrations, in part due to the excellent linearity and uniformity of modern CCD detectors. Several groups around the world are now actively pursuing digital, multi-colour surveys of the sky. The SkyMapper telescope and Southern Sky Survey fill an urgent demand for high \\emph{etendue} multi-colour optical survey work in the southern hemisphere that will be unrivaled until the commissioning of the Large Synoptic Survey Telescope (LSST) in Chile in 2015. ",
        "conclusions": ""
    },
    "0806/0806.2773_arXiv.txt": {
        "abstract": "Chemically peculiar A stars (Ap) are extreme examples of the interaction of atomic element diffusion processes with magnetic fields in stellar atmospheres. The rapidly oscillating Ap stars provide a means for studying these processes in 3D and are at the same time important for studying the pulsation excitation mechanism in A stars. As part of the first comprehensive, uniform, high resolution spectroscopic survey of Ap stars, which we are conducting in the southern hemisphere with the Michigan Spectral Catalogues as the basis of target selection, we report here the discovery of 17 new magnetic Ap stars having spectroscopically resolved Zeeman components from which we derive magnetic field moduli in the range $3-30$\\,kG. Among these are 1) the current second-strongest known magnetic A star, 2) a double-lined Ap binary with a magnetic component and 3) an A star with particularly peculiar and variable abundances. Polarimetry of these stars is needed to constrain their field geometries and to determine their rotation periods. We have also obtained an additional measurement of the magnetic field of the Ap star HD\\,92499. ",
        "introduction": "The study of chemically peculiar A stars with magnetic fields has ramifications for several other branches of astrophysics.  The interaction among strong magnetic fields, atomic diffusion and energy transfer in (or above) the upper atmospheres of non-degenerate stars has direct implications for observed stellar abundances and for the instability of stellar pulsations in $\\beta$\\,Cephei and sdB stars. The magnetic field lines guide the diffusing elements, so that some ions will be concentrated where the field lines are vertical while others will group where the lines are horizontal (e.g. \\citealt{michaud81}).  The very different magnetic field strengths, orientations and geometries, as well as the great variety of abundance distributions for Ap stars are therefore particularly informative.  Magnetic fields affect spectroscopic lines through, e.g., broadening, intensification and Zeeman splitting of the intrinsic lines \\citep{mathys89}. Several cool Ap stars show light and spectral variability that follow the stellar rotation, and such stars are known as $\\alpha^2$ Canum Venaticorum ($\\alpha^2$\\,CVn) variables. \\citet{wolfetal71} demonstrated that the observed light, spectrum and magnetic variations are related. In particular, they suggested that photometric variability may result from a redistribution of flux either by variations in line blanketing or opacity, such as enhanced absorption at rare earth maximum. The light variability may vary with that of the rare earths which appear to be concentrated predominantly near the region of the strongest magnetic field or that of the negative pole. Thus can light variability be linked to magnetic field variability which, for oblique rotators, follows the stellar rotation.  Magnetic activity cycles similar to the solar 11-y cycle are observed in F--M stars \\citep{baliunas95} and are related to chromospheric activity. However, chromospheres have not been unambiguously detected in Ap stars. Finally, magnetic fields drive winds through Alfv\\'en waves or other magneto-hydrodynamic  waves \\citep{linsky04}.   The chemically peculiar B, A and F stars (henceforth Ap stars) have globally organised, typically dipole or quadrupole fields that with an axis obliquely inclined to the rotation axis of the star so that, for dipolar fields,  one or both magnetic poles come into, or out of, view with rotation. The full geometry of the field may hence be visible for stars with favourably oriented rotation and magnetic axes, which in turn makes it possible to use the Zeeman Doppler imaging technique to create magnetic field maps. Among the Ap stars is Babcock's star = HD\\,215441 \\citep{babcock60}, the non-degenerate star with the strongest field known, 34.4\\,kG. The only other published cases with extremely strong fields are: HD\\,137509, $\\langle B\\rangle=29$\\,kG \\citep{kochukhov06}, HD\\,154708, $\\langle B\\rangle= 24.5$\\,kG \\citep{hubrigetal05} and HD\\,178892, $\\langle B\\rangle=18.0$\\,kG \\citep{ryabchikova06}.  The atmospheres of Ap stars are complex to interpret. As a consequence of vertical abundance stratification, and non-standard temperature gradients, even theoretical modelling of hydrogen line profiles, such as the core-wing anomaly in the H$\\alpha$ lines \\citep{cowleyetal01}, is not yet completely successful \\citep{kochukhovetal02}. However, the observational and theoretical efforts in understanding the Ap atmospheric processes may eventually prove `worth the candle' thanks to their general applicability; the Ap stars are the most extreme examples of magnetic fields and atomic diffusion in non-degenerate stars, processes common in most other stars. Examples of studies of atmospheric abundance stratification in a cool Ap star subgroup, the rapidly oscillating Ap (roAp) stars are given by \\cite{wadeetal01,ryabchikovaetal02,ryabchikovaetal05}. These studies are in agreement with each other, in general: Fe is concentrated by gravitational settling in the observable layer between $-1<\\log\\tau_{\\rm {5000}}<0$ and Pr and Nd are concentrated by radiative levitation above $\\log\\tau_{\\rm {5000}}=-5$. The Pr and Nd forming layers are above the line forming layer of the narrow core of the H$\\alpha$ line which in standard A star models is in the range $-4<\\log\\tau_{\\rm {5000}}<-2$.  Magnetic field measurements are mostly based on difficult, indirect measurements, such as magnetic broadening of spectral lines and spectropolarimetric observations of longitudinal magnetic fields or the line-of-sight field component. The mean longitudinal magnetic field (or, the longitudinal field) $\\langle B_{\\rm z}\\rangle$ is a weighted average over the visible stellar disk of the component of the magnetic vector along the line of sight \\citep{mathys89}, and is typically at least 3 times weaker than the mean magnetic field modulus $\\left<B\\right>$ (see Sect.\\,\\ref{sec:mag} for definition of this). The first large survey for magnetic fields in non-degenerate stars was done by \\citet{babcock58}. It has since been followed by many other studies of individual stars or groups of stars (e.g., \\citealt{borraetal79,bohlenderetal93}). Recent searches for magnetic stars were carried out by \\citet{kudryavtsevetal06} who discovered $\\sim 70$ magnetic stars, and by \\citet{hubrigetal06} who discovered 57 magnetic Ap stars. The total number of firmly established magnetic chemically peculiar main-sequence stars is currently $\\sim 350$ (Romanyuk, 2008, in prep.) For a significant fraction of these magnetic stars the variation in the longitudinal magnetic field has been measured as a function of rotation period (see, e.g., \\citealt{mathys91}).  Such curves provide valuable information about the global geometrical structure of magnetic fields. Better and more reliable measurements are necessary to obtain a complete understanding of the many phenomena related to magnetic fields. For this purpose, magnetic Ap stars with resolved magnetically split lines provide extremely favourable conditions, since one can determine in a straightforward, mostly approximation-free, model-independent manner, and with particularly good precision, the mean magnetic field modulus \\citep{mathysetal97}. However, in spite of many studies (e.g., \\citealt{mathysetal97}; \\citealt{hubrigetal07}), only 51 such stars were known \\citep{hubrigetal07}, prior to this work. Our discovery of 17 new such stars therefore significantly increases the number known.  At present, 40 roAp stars are known (see, e.g., \\citealt{kurtzetal06b,gonzalezetal08}), although several surveys have searched for rapid pulsation in Ap stars, such as \\citet{nelsonetal93, martinezetal94, handleretal99, ashokaetal00, weissetal00, dorokhovaetal05}. For a (non-exhaustive) list of spectroscopic studies of roAp stars, see \\citet{kurtzetal06a}. The roAp stars are thought to be characterised by, e.g., the H$\\alpha$\\ {\\em core-wing anomaly} and by an {\\em ionization disequilibrium} for Nd\\,\\textsc{ii} and Nd\\,\\textsc{iii} and for Pr\\,\\textsc{ii} and Pr\\,\\textsc{iii} \\citep{ryabchikovaetal04}, the latter caused by stratification of those two elements to levels above $\\log \\tau_{{5000}} \\sim -5$ by atomic diffusion. Other typical spectral characteristics are strong magnetic fields, wing-nib anomaly in \\ion{Ca}{ii}\\,K \\citep{cowleyetal06}, and strong abundances of rare earth elements and relatively slow rotation. It is therefore possible to identify promising roAp star candidates with a single, high signal-to-noise ratio, high-resolution spectrum before spending time on a large telescope with fast spectroscopy to search for rapid pulsations.  In 2006 we therefore began a systematic survey of cool Ap stars in the photometric `Cape cool Ap star catalogue' \\citep{martinez93} to identify roAp candidates based generally on a single high-resolution spectrum of each star. The Cape catalogue gives Str\\\"omgren $uvby$ and $\\beta$ photometry for over 500 (almost all) of the SrCrEu subclass of cool Ap stars listed in the Michigan spectral catalogues, volumes $1-4$ \\citep{houk78,houk82,houketal75,houketal88}. Perhaps surprisingly -- given the importance of the Ap stars to stellar astrophysics, and the long history of their study -- there is no large-scale, uniform spectroscopic survey such as the one we are now conducting. We expect the data set to be a rich source of discoveries, and to be the basis of uniform statistical analyses of many of the astrophysically interesting characteristics of the class.  Our first observing season is finished and of 140 stars, we identified dozens of stars with magnetically intensified, broadened or even resolved lines. Of these, 17 are new detections with magnetically resolved or partially resolved lines, in particular for the Zeeman doublet \\ion{Fe}{ii}\\,6149.258\\,\\AA\\ (Figs\\,\\ref{fig:fe6149a}--\\ref{fig:fe6149b}). As pointed out by \\citet{mathysetal97}, this line is particularly important as a diagnostic for a magnetic field as the splitting provides a direct measure of the mean magnetic field, and furthermore because iron is usually rather homogeneously distributed over the surface of Ap stars.  This is the discovery paper for the new Ap stars with magnetically resolved lines, and we provide magnetic field strengths, projected rotation velocities and selected relative abundance estimates. Of particular interest, we discovered a new star with an extremely large magnetic field; a highly peculiar magnetic star; and an uncommon magnetic Ap star in a relatively close binary system. We also re-observed the recently discovered magnetic Ap star HD\\,92499 \\citep{hubrigetal07} to check for stability of its magnetic field strength, abundances and radial velocity.  In the following sections, we describe the selection and observation of the targets along with the data reduction in Sect.\\,2. Then follows the data analyses, including estimation of physical parameters and the magnetic measurements in Sect.\\,3.  Finally we discuss the results in Sect.\\,4. ",
        "conclusions": "As part of a systematic search for new roAp candidates using FEROS on the ESO 2.2-m telescope, we have discovered 17 new magnetic stars with magnetically resolved lines and with these lines measured their mean magnetic field moduli directly.  For 11 stars, spectra were obtained about 30\\,d later at higher resolution with UVES on the VLT. These were used to confirm the discoveries and check the stability of the measured magnetic field strengths and radial velocities.  A new double-lined spectroscopic binary, HD\\,135728AB, was discovered with two similar components, for one of which, the more slowly rotating component, a magnetic field was detected. It is possible that the primary (the faster rotating and more massive star, as deemed from its significantly smaller radial velocity difference between our two spectra) is either an Am or Ap star -- i.e. either magnetic or non-magnetic; as yet we cannot tell.  It is likely that the secondary is a spectrum variable (the primary may be also) so that the rotation periods of the stars can be determined independently and compared to the orbital period, thus showing whether either or both of the stars are synchronously rotating. Since both stars show overabundances typical of Ap and Am stars, a first guess might be that one is Am (non-magnetic) and the other Ap (magnetic). If that is so, how can two, rather similar stars in a close binary with both in the Ap-Am domain end up with one strongly magnetic and the other not? Or, on the other hand, if the primary {\\it is} magnetic, but has a weaker field, then the question is similar, but not so extreme: how do the two components end up with different field strengths. There are Am-Am SB2 systems known -- WW\\,Aurigae is a famous eclipsing example where the stars are very similar in mass, but not identical in abundances. SB2 systems with a magnetic Ap star are very rare and HD\\,135728AB may be particularly promising for illuminating the magnetic field origin question. Examples of other such cases are: HD\\,59435 \\citep{wadeetal99}, HD\\,55719 \\citep{bonsack76} and HD\\,98088 \\citep{abtetal68,hensberge74,bychkovetal05}. These binaries have the following characteristics: magnetic periods of 5.8 -- 1360\\,d, orbital periods of 5.9 -- 1386.1\\,d, magnetic fields of 1.45 -- 8\\,kG and companion-to-Ap star mass ratios of 0.75 -- 1.33. In a recent study of the very young binary HD\\,72106, \\citet{folsometal07} established a dipole field of 1.3\\,kG for the primary star which rotates fast enough to permit these authors to produce 2-D surface abundance maps.  It is possible that two other systems in our sample are binary: HD\\,55540 had a significant radial velocity  change, while HD\\,121661 needs confirmation of a marginally significant change. Neither of these systems, if confirmed, are of the importance of HD\\,135728AB and further observations are planned to determine the orbital period and use spectral disentanglement to study the abundances of both its components in detail.  The most important discovery, HD\\,75049, has the second-largest known magnetic field of any Ap star and was found to be highly variable in magnetic field strength and fine-structure over a time scale of 30\\,d. The extremely strong magnetic field may not only rival, but even surpass the strength of Babcock's star (HD\\,215441), 34.4\\,kG. Follow-up is in progress of this very interesting object.  HD\\,96237  was shown to exhibit extreme abundance variations, possibly related to a photometric variability of 22\\,d that may be the rotation period. Extreme abundance variations with stellar rotation are known from more slowly rotating stars, such as HR\\,465 (HD\\,9996; \\citealt{prestonetal70}) for which Eu, Cr, Ca, Sr vary up to a factor of 3 in line strength (Cr and Eu in antiphase). A magnetic field was detected and measured. We compared the spectra of HD\\,96237 with those of the arguably most peculiar Ap star, HD\\,101065, and demonstrated comparable levels of peculiar abundances. HD\\,96237 is remarkable by also exhibiting fast abundance variations. Follow-up studies are currently in progress of this important star.  From light curves in the ASAS database, the stars HD\\,75049, HD\\,88701, HD\\,96237, HD\\,110274, HD\\,121661, HD\\,117290 and HD\\,92499 were shown to be $\\alpha^2$\\,CVn variables. Two periods were detected for HD\\,75049, while HD\\,88701 exhibits a clear double wave. HD\\,117290 and HD\\,92499 show, as the only cases, variability longer than the time span covered by the photometry (3--5 years). There are some implications of the rotation periods found from the ASAS data and the measured rotational velocities, since these constrain the stellar radii and thus luminosities. An example we discussed is HD\\,88701 for which the rotation period and \\vsini\\ implied an unexpectedly large radius. However, uncertainties in \\vsini\\ measurements, sensitive to line-blending and magnetic broadening and the possibility of an ASAS period only being half of the rotation period (in case of double-wave light curves) leaves some uncertainties. Furthermore, astrometric luminosities suggest that at least half the stars are near the terminal end of the main-sequence, so that larger radii are to be expected in comparison with younger Ap stars.  Included among our original sample 140 cool Ap stars is HD\\,92499, a known magnetic star with Zeeman splitting \\citep{hubrigetal07}. Our new spectra showed a constant magnetic field modulus and radial velocity of the star. As our target selection is relatively unbiased among more than 500 cool Ap stars, the fraction of such stars with magnetically resolved lines appears to be 13 per cent based on the first part of our Ap star survey.  Abundance estimates were used to identify Nd and Pr ionization disequilibrium anomalies in abundances of ions in the two first ionized states. HD\\,44226, HD\\,96237 and HD\\,143487 showed significant abundance anomalies ($\\sim1$\\,dex) and are in addition to HD\\,92499  excellent roAp candidates. With 32\\,min time-series spectroscopy of HD\\,143487, we demonstrated a low-amplitude candidate period of 2\\,mHz that, however, could not be confirmed by individual lines.  We emphasise that the lack of accurate {\\it Hipparcos} \\citep{hip} parallaxes $[\\sigma(\\pi)/\\pi<0.2]$ for the newly detected magnetic stars presented in this paper means that absolute magnitudes are considerably more difficult to determine because of the peculiar spectra of these stars. We note that for a subset of the stars having parallaxes, their relatively clustered location in the H-R diagram indicates that many of these are stars near the end of their main-sequence lifetimes. This new sample of stars with directly measured magnetic fields will aid studies of magnetic field interaction with stellar atmospheres. Polarimetric measurements are needed to establish geometry of the detected fields and to determine or confirm the rotation periods of the stars. We are currently obtaining high time resolution spectroscopy with UVES for 16 of the 18 stars in this study to search for rapid pulsations. The exceptions are HD\\,75049, which is too hot to be a roAp star, and HD\\,135728AB, which is an SB2 system."
    },
    "0806/0806.0295_arXiv.txt": {
        "abstract": "{The International Gamma-Ray Astrophyiscs Laboratory (INTEGRAL) is discovering a large number of new hard X-ray sources, many of them being HMXBs. The identification and spectral characterization of their optical/infrared counterparts is a necessary step to undertake detailed study of these systems. In particular, the determination of the spectral type is crucial in the case of the new class of Supergiant Fast X-ray Transients (SFXTs), which show X-ray properties common to other objects.} {Our goal is to perform spectral analysis and classification of proposed counterparts to HMXBs in order to characterize the system they belong to.} {We used the ESO/NTT SofI spectrograph to observe proposed IR counterparts to HMXBs, obtaining $K_{s}$ medium resolution spectra ($R = 1320$) with a S/N $\\gtrsim$ 100. We classified them through comparison with published atlases.} {We were able to spectrally classify the six sources. This allowed us to ascribe one of them to the new class of SFXTs and confirm the membership of two sources to this class. We confirmed the spectral classification, derived from optical spectroscopy, of a known system, 4U 1907-09, showing for the first time its infrared spectrum. The spectral classification was also used to estimate the distance of the sources. We compared the extinction as derived from X-ray data with effective interstellar extinction obtained from our data, discussing the absorption component due to the circumstellar environment, which we observed in four systems; in particular, intrinsic absorption seems to emerge as a typical feature of the entire class of SFXTs. } {} ",
        "introduction": "\\begin{table*}[!ht] \\caption[]{NTT journal of observations. We report in the fourth column the net accumulated exposure time. Column five gives the obtained signal-to-noise ratio. The references list in the last column relates to the identification of the optical/infrared counterpart.} \\label{table:logobs} \\centering \\begin{tabular}{lcccccl} \\hline \\hline \\noalign{\\smallskip} Source  & K mag &  Start time (UT) &  Exp. time (s) & S/N & IR Counterpart & Reference\\\\ \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} IGR J16207--5129 & 9.1  & 2006-07-14 23:19   & 600 &  100 &  2MASS J16204627-5130060& Tomsick et al. (2006)\\\\ IGR J16465--4507 & 9.8 & 2007-05-26 05:03 & 240  & 100 &  2MASS J16463526-4507045 & Zurita and Walter (2004)\\\\ IGR J16479--4514 & 9.8 & 2007-05-26 05:11 & 240   & 100 & 2MASS J16480656-4512068 & Kennea et al. (2005)\\\\ AX J1841.0--0536 & 8.9 & 2006-07-14 03:57 & 600  & 180 &  2MASS J18410043-0535465& Halpern et al. (2004)\\\\ 4U 1907+097 & 8.8 & 2006-07-15 07:37 & 600 & 130 & 2MASS J19093804+0949473 & Schwartz et al. (1980)\\\\ IGR J19140+0951 & 7.1 & 2006-07-14 04:46 & 360 & 130  & 2MASS J19140422+0952577 &  in't Zand et al. (2004)\\\\ \\noalign{\\smallskip} \\hline \\end{tabular} \\end{table*}  High Mass X-Ray Binaries (HMXBs) are systems composed of an early-type massive star and an accreting compact object. All sub-groups of HMXBs involve OB type stars and are commonly found in the galactic plane and in the Magellanic Clouds, among their OB progenitors.\\  The majority of the known systems are Be/X-ray Binaries (BeXRBs), consisting of a neutron star accreting matter from the circumstellar equatorial disc of a Be star. Most of them are transient, exhibiting short and bright outbursts ($L_{X}\\sim 10^{36} - 10^{37}$ erg s$^{-1}$ in the case of Type I outbursts, generally close to the periastron passage of the neutron star; $L_{X} \\ge 10^{37}$ erg s$^{-1}$ in the case of Type II outbursts).  In the second major class of HMXBs, the Supergiant X-ray Binaries (SGXRBs), the counterpart is an early supergiant star, feeding the compact object with its radially outflowing stellar wind. As a result, the SGXRBs are, generally, persistent systems ($L_{X}\\sim 10^{36}$ erg s$^{-1}$). \\\\  The five-year INTEGRAL data possibly reveal a different scenario. A recent subgroup has been proposed by \\citet{negue06}, named Supergiant Fast X-ray Transients (SFXTs): these objects, associated with a supergiant companion, occasionally undergo a short period of X-ray activity lasting less then a day, typically a few hours \\citep{sguera05}, with a very different behavior from those observed in other X-ray binaries. These outbursts show very sharp rises, reaching the peak of the flare in $\\lesssim$1 hour. The decay is generally of a complex kind, with two or three further flares. The physical reason for these fast outbursts is still unknown, although theoretical speculations would connect them to some sort of discrete mass ejection from the supergiant donor \\citep{gol03} or to wind variability \\citep{int05}, or to the possible presence of a second, equatorial, wind component \\citep{sid07} . Due to high interstellar absorption and to the transient nature of these sources, SFXTs are difficult to detect, and in most cases, the sources had not been detected by previous missions. To date, six objects have been firmly characterized as SFXTs, but many other systems are likely candidates, and their number has grown rapidly since the launch of INTEGRAL \\citep{wink03}, so that they could actually constitute a major class of X-ray binaries.  Up to now, the INTEGRAL survey of the Galactic Plane and central regions has revealed the existence of more than 200 sources \\citep{bird07,bodag07} in the energy range 20--100 keV, with a position accuracy of $2'-3'$, depending on count rate, position in the FOV and exposure. A large fraction of the newly discovered sources are found to be heavily obscured, displaying much larger column densities ($N_{H}\\gtrsim 10^{23}$ cm$^{-2}$) than would be expected along the line of sight \\citep[see][]{kuulkers05}. These sources were missed by previous high-energy missions, whose onboard instruments were sensitive to a softer energy range. Moreover, optical counterparts to these obscured sources are poorly observable due to the high interstellar extinction, with $A_{V}$ in excess of up to $\\sim20$ mag.\\\\  In this context, the recent availability of infrared spectroscopy has emerged as a strong tool to characterize these systems and, together with high-energy data, reveal the HMXB sub-class they belong to. This results in the identification of the mass transfer process of the system, with information about the intrinsic physics of the X-ray binary. The need for low energy data is particularly urgent in the case of SFXTs, which show X-ray properties common to other objects (such as RS CVs binaries and Low Mass X-ray binaries) and thus crucially require the spectral classification of their counterpart in order to be properly discerned.\\\\   In this paper we present spectral analysis and classification of six HMXBs discovered (or re-descovered) by INTEGRAL. The selected IGR sources are the following: IGR J16207--5129, IGR J16465--4507, IGR J16479--4514, AX J1841.0--0536 and IGR J19140+0951. We also included the well known system 4U 1907+09 since the spectral classification of its counterpart has been a matter of debate in the past, and no infrared spectra have been published up to now. The first three sources are located in the direction of the Norma-arm tangent region, the fourth in the Scutum-arm tangent region, the fifth and the last one in the Sagittarius arm tangent. In the next section we describe the observations and data reduction; in Section \\ref{results} we report the obtained spectra, analyze their features and propose a classification; we calculate the interstellar hydrogen column density and estimate the distance to each source; in section \\ref{discussion} we discuss our results, before concluding.  Preliminary results of our data analysis for AX J1841.0--0536 and IGR J19140+0951 were published in \\citet{nespoli07}. ",
        "conclusions": "From near-infrared spectroscopy of the six HMXBs we have found that:  \\begin{itemize} \\item[-] the proposed optical counterparts were confirmed and the spectral classification of the sources provided; \\item[-] one source, IGR J16479--4514, was added to the SFXTs and the confirmation of IGR J16465--4507 and AX J1841.0--0536 as members of the class was proven with infrared data; \\item[-] the comparison between $N_{H}$ obtained from X-ray data and interstellar extinction from our data showed for four systems (IGR J16465--4507, IGR J16479--4514, AX J1841.0--0536 and IGR J19140+0951) the presence of an absorbing envelope, strictly confined around the compact object; \\item[-] all the three identified SFXTs are intrinsically absorbed, suggesting that this might be a characteristic of the class; \\item[-] the distance estimation, compatible with the location of the sources in the respective galactic arms, is a possible confirmation of the spectral classification provided here. \\end{itemize}"
    },
    "0806/0806.2545_arXiv.txt": {
        "abstract": "We discuss a \"compact source\" model of very high energy (VHE) emission from blazars in which the variability time is determined by the blazar central engine. In this model electron or proton acceleration close to the supermassive black hole is followed by the development of electromagnetic cascade in a radiatively inefficient accretion flow. Assuming such a model for the TeV blazar \\pks, we show that the variability properties of the TeV \\gr\\ signal observed during a bright flare from this source, such as the minimal variability time scale and the recurrence period of the sub-flares, constrain the mass and the angular momentum of the supermassive black hole. ",
        "introduction": "Recent observation of fast variability of TeV \\gr\\ emission from several TeV blazars \\citep{aharonian07,albert07} challenges the conventional model in which the TeV \\gr s are supposed to be produced at large distances from the blazar central engine, the supermassive black hole. Within this conventional model the \\gr\\ emitting blobs are assumed to travel to the parsec-scale distances along the AGN jet before emitting in the TeV energy band; it is believed that otherwise the \\gr\\ emission would be strongly absorbed in the accretion flow (see e.g. \\citet{1996MNRAS.280...67G}). Since the radiative cooling time of the TeV emitting electrons is typically shorter than the time of propagation from the central engine to the TeV emission region, it is usually assumed that the TeV emitting particles are produced via shock acceleration locally in the emission region, rather than in the AGN central engine.  At the same time, the observed short variability time scales $t_{\\rm var}\\sim$ a few minutes indicate that the TeV emission comes from very compact regions having the size in the comoving frame $\\Delta x'\\lesssim \\delta(1+z)^{-1}ct_{\\rm var}$, where $\\delta$ is the bulk Doppler factor and $z$ is the source redshift. This implies that in the static frame the longitudinal size of the TeV emitting region is \\begin{equation} \\label{rlab} \\Delta x=\\frac{\\Delta x'}{\\Gamma}\\lesssim (1+z)^{-1}ct_{\\rm var}\\simeq 6\\times 10^{12}(1+z)^{-1}\\left[\\frac{t_{\\rm var}}{200\\mbox{ s}}\\right]\\mbox{cm}\\;. \\end{equation} where we assume that the bulk Lorentz factor $\\Gamma\\sim \\delta$. This is comparable to the minimal possible scale set up by the gravitational radius of the central supermassive black hole \\be \\label{Rg} R_g=GM_{\\rm BH}/c^2\\simeq 1.5\\times 10^{12}\\left[M_{\\rm BH}/10^{7}M_\\odot\\right]\\mbox{ cm}\\;. \\ee Even if one assumes that the TeV emitting plasma blobs are produced close to the black hole, which would explain their initially small size, it is not clear how the blobs propagating downstream the relativistic jet can retain this size up to large distances,  unless they have unreasonably large bulk Lorentz factors. This problem has recently led to a suggestion \\citep{begelman07} that the TeV flares may be not triggered by the black hole but rather are results of enhanced emission intrinsic to the jet.  However even in that case, to explain the observed rapid variability, the TeV \\gr\\ emitting blobs have to travel with quite large bulk Lorentz factors, $\\Gamma\\gg 1$. Such Lorentz factors are in contradiction with the radio observations of the motion of hot spots in the parsec-scale  jets.  Moderate apparent speeds of the blobs of the parsec-scale jets, revealed by radio observations, combined with an estimate of the number of parent objects of TeV blazars, would give  much smaller values of the bulk Lorentz factors, $\\Gamma\\sim 1$ \\citep{henry06}. For example, to explain the fast variability of the July 2006 TeV flare of \\pks\\ \\citep{aharonian07} the bulk Lorentz factor required by the mechanism of TeV emission in the parsec-scale jet should be $\\Gamma>50$~\\citep{aharonian07,begelman07}.  At the same time, the direct observations  of the apparent velocity of hot spots in \\pks\\ jet at the projected distance $(1\\div 2)$ parsecs give the value $v_{\\rm app}=(0.9\\pm 0.3)c$  \\citep{Piner:2008ju}. Assuming that the viewing angle is not too small, $\\theta\\gtrsim 1^\\circ$, this yields $\\Gamma\\lesssim 10$ at the distance of a few tens of parsecs from the central engine.  Both the problems of the fast variability and of the small observed Lorentz factors at parsec distances could be naturally resolved if the site of the VHE \\gr\\ production is located closer to the AGN central engine, at significantly sub-parsec distances. If the VHE emitting region is moving relativistically toward the observer with a bulk Lorenz factor $\\Gamma$, the variability time scale limits the distance $R$ of the \\gr\\ production site from the central engine (see e.g. \\citet{celotti}), \\begin{equation} R\\sim \\Delta x\\,\\Gamma^2\\le 1.5\\times 10^{16}(1+z)^{-1}\\left[\\frac{t_{\\rm var}}{200\\mbox{ s}}\\right]\\left[\\frac{\\Gamma}{50}\\right]^2\\mbox{ cm.} \\end{equation} An immediate difficulty is, however, that at such distances the accretion flow onto the black hole can be opaque to the \\gr s (see e.g. \\citet{blandford95}).  The problem of opacity of the compact source does not arise in the case of low-luminosity AGNs that accrete at significantly sub-Eddington rates \\citep{celotti}. In these sources the accretion flow is described within the framework of the radiatively inefficient accretion flow (RIAF) models \\citep{rees82,narayan94,narayan95,narayan02} in which most of the gravitational energy extracted from the accreted matter is converted into internal energy, rather than into radiation. The possibility of escape of the VHE \\gr s from the vicinity of the AGN central engine is best illustrated by the nearby low-luminosity radio galaxy M87, which was recently found to be a source of the variable TeV \\gr\\ emission \\citep{aharonian03,aharonian06,albert08}, most probably coming from a compact source \\citep{neronov07,aharonian08a}.  In the compact source model the VHE \\gr\\ emission is triggered by high-energy particles accelerated close to the black hole via one of the possible mechanisms (see e.g. \\citet{lovelace76,Lesch1992,kardashev95,Bednarek:1998jq,neronov02, neronov02a,neronov04, Rieger:2007tt,neronov08}). In this case the spectral and timing characteristics of the VHE emission are directly linked to the physics of the processes taking place close to the supermassive black hole, which naturally explains the variability of the signal on the shortest possible time scale.  Within the AGN unification scheme, the TeV blazars (high-energy peaked BL Lacs) are assumed to be the beamed versions of the low-luminosity radio galaxies similar to M87 \\citep{browne83,giroletti04,giroletti06}.  Since the only difference between the TeV blazars and the low-luminosity radio galaxies is their orientation with respect to the line of sight, the compact source model can be applicable also in the case of TeV blazars.  In what follows we adopt this point of view and develop a compact source model of high-energy activity of TeV blazars. We demonstrate that within such compact source model the characteristics of the fast-variable VHE emission can be used to constrain the parameters of the AGN central engine, in particular, the density and luminosity of the accretion flow, the black hole mass and spin. We illustrate this possibility on the particular example of the bright \\pks\\ flare detected by the HESS telescope in July 2006 \\citep{aharonian07}. This flare consists of a number of well-pronounced sub-flares which exhibit quasi-periodic recurrence. We show that the rise time and the recurrence period of the sub-flares can be directly related to the light-crossing time and to the period of rotation over the last stable orbit around a $M_{\\rm BH}\\sim 10^7M_\\odot$ black hole\\footnote{This mass estimate is different from the value $\\sim 10^9 M_\\odot$ quoted by \\citet{aharonian07}. We will comment on this discrepancy in Sec.~\\ref{sec:3.3}.}.  The paper is organized as follows.  In Sec.~\\ref{IR} we discuss the qualitative features of the model, including particle acceleration and propagation through the RIAF environment. The possibility of a new interpretation of the observational data within such a model is demonstrated in Sec.~\\ref{TIMING} where we find the constraints on the black hole mass and angular momentum imposed by the timing analysis of the bright TeV flare of \\pks. We summarize our results in Sec.~\\ref{CONCL}. ",
        "conclusions": "\\label{CONCL}  In this paper we have proposed that the recently observed fast variability of the VHE emission from blazars can be naturally accommodated within the framework of \"compact source\" model. In this model particles responsible for the observed VHE emission are accelerated close to the central supermassive black hole, rather than at large distances downstream the AGN jet.  We have analyzed the problem of escape of the VHE \\gr s from the vicinity of the central engine and demonstrated that the region around the central engine is transparent for TeV \\gr s if the accretion flow in the TeV blazars is radiatively inefficient. If the luminosity of the accretion flow is as low as $L_{\\rm acc}\\lesssim 10^{40}\\mbox{erg/s}$, the TeV \\gr\\ emission can come directly from the immediate neighborhood of the central black hole. Alternatively, for brighter accretion flows, the TeV \\gr s may be produced at some distance from the black hole in a proton-initiated electromagnetic cascade developing in the accretion flow.  The possibility that the properties of the VHE \\gr\\ emission are directly linked to the properties of the central engine of the AGN, if confirmed by future observations, provides a new tool to study the physical conditions in the direct vicinity of the supermassive black hole. In particular, the VHE signal can be used to constrain the parameters of the accretion flow and of the black hole itself, such as its mass and spin.  We demonstrated this possibility on the example of the bright TeV flare of the blazar \\pks. Within the proposed scenario, the characteristic time scales, found in the timing analysis of this flare, are directly related to the parameters of the supermassive black hole. The minimal variability time scale of the signal is identified with the black hole light-crossing time. This sets the bound on the black hole mass and its rotation moment shown in Fig.~\\ref{fig:period_aM}. We also observed that the signal exhibits quasi-periodic oscillations. Identifying the recurrence time of these oscillations with the period of rotation around the black hole we obtained a relation between the black hole mass and its rotation moment.  A detailed modeling based of the framework proposed in this paper should involve calculation of particle acceleration and propagation in the vicinity of the black hole through the environment created by the accretion flow. This can be done assuming particular (numerical) models of RIAF and particle acceleration. We leave this for future work.   \\subsection*"
    },
    "0806/0806.0861.txt": {
        "abstract": " ",
        "introduction": "instead of showing variations of cosmic function with parameters, can compile plots that reflect the %current understanding and that provide a baseline for the discussions  %Part III %Chapters 9--10: survey results and applications %evaluates the status and compiles first results obtained for thesis work  %Part IV %Chapters 11: outlook and future developments %golden phase will continue, provide a sketch of what is coming  %-->3 pages  %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%    %--------------------- KAPITEL 2 -----------------------------------------------  \\chapter{Galaxy Clusters as Astrophysical Laboratories} \\label{c2_cluster_theory}  \\noindent  This chapter will review selected aspects of the physical foundation of galaxy clusters. The main focus will be placed on the X-ray properties of the intracluster medium (ICM) and optical characteristics of the (distant) cluster galaxy population. The discussions on  X-ray emission mechanisms, ICM structure and formation, and cool cores provide the background for understanding the high-redshift cluster detection procedure. X-ray scaling relations are critical for the mass calibration, whereas the \\reds of cluster galaxies is used for the initial redshift estimation. The sections on galaxy formation, evolution, and environmental effects set the stage for the high-redshift cluster studies in Chaps.\\,\\ref{c10_HizClusterStudies}\\,\\&\\,\\ref{c10b_science_outlook}. More details can be found in the classic work of Sarazin \\cite*{Sarazin1986a}, the review of Voit \\cite*{Voit2004a}, or Biviano \\cite*{Biviano2000a} for a historical treatise of clusters.   %--> ~10 pages  %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% %................... 2.1 ....................................................... ",
        "conclusions": "%\\noindent %In this chapter, the  Z--H \\reds method was calibrated and applied to  % data for the confirmation of distant cluster candidates %allowed a number of additional galaxy evolution studies that will help to {\\bf reconcile the observed cluster galaxy properties %with the standard paradigm of hierarchical galaxy formation}:   %\\begin{enumerate} \\begin{itemize} % ? klarer machen  \\item In Sect.\\,\\ref{s10_formation_epoch}, the Z--H \\reds method was calibrated using ten spectroscopically confirmed clusters. It was shown that the ($z_{\\mathrm{f}}\\!=\\!5$, $Z\\!=\\!Z_{\\sun}$)-passive evolution model is fully consistent with all observations and was therefore applied as reference model for the photometric redshift estimation of the full cluster candidate sample.   \\item  The observed Z--H  \\reds color evolution over the redshift baseline $0.2\\!\\la\\!z\\!\\la\\!1.5$ was used to simultaneously constrain the average metallicity and the formation epoch of \\reds galaxies yielding  ${z_{\\mathrm{f}}\\!=\\!4.2\\!\\pm\\!1.1}$ and     $Z\\!=\\!(1.2\\!\\pm\\!0.4)\\,Z_{\\sun}$. These early results, based on only ten reference objects, confirm the well-established old age of the stellar populations of \\reds galaxies. Using additional spectroscopically confirmed high-redshift clusters of the ongoing XDCP follow-up program, the  Z--H  \\reds method has the potential to put stringent constraints on the early-type galaxy formation epoch in the near future.   %, similar to the findings of numerous cluster galaxy studies before. %    was used to  constrain the  {\\bf formation epoch of  \\reds galaxies\\/}  to $\\mathbf{z_{\\mathrm{\\bf f}}\\!=\\!4.2\\!\\pm\\!1.1}$. %The observed Z--H \\reds colors are found to be fully consistent with predictions for passive galaxy evolution models, similar to %the findings of numerous cluster galaxy studies before.  % ? confirmation with one  a posteriori measure redshift % established tool for redshift estimation % confirm preliminary method to select BCGs   %    \\item In Sect.\\,\\ref{s9_cl15_analysis}, the large-scale structure environment of a newly spectroscopically confirmed X-ray %cluster at $z\\!=\\!0.95$ was analyzed, finding a probable associated second X-ray cluster at a projected distance of 3\\,Mpc. %The main cluster exhibits %strong indications for galaxy transformation processes at the cluster outskirts, manifested by an observed color transition of %the galaxies.   \\item In Sect.\\,\\ref{s9_cl15_analysis}, details of the newly discovered X-ray luminous cluster of galaxies         XMMU\\,J0104.4-0630 at redshift $z\\!=\\!0.947\\!\\pm\\!0.005$ were  presented. The compact, intermediate mass cluster is found to be in an evolved state and hosts a strong central radio source. It was shown that the cluster shows a pronounced stratification of galaxy populations. Whereas the  spatial distribution of the red-sequence population of early type galaxies coincides well with the X-ray emission,  a significantly bluer population dominates beyond 1--2 core radii from the center, suggesting a cluster environment-driven effect of differential galaxy evolution, \\eg \\  delayed star formation quenching in the outskirts.   %  which could be a sign of ongoing cooling core activity. %The systems is consistent?? with local Lx-T,.... scaling relation.  % ? BCG consistent \\item The Z--H color of the cluster \\reds was shown to be in good agreement with the  predictions of the reference model and confirmed the applicability of the method for reliable photometric redshift estimates.   \\item A second  X-ray selected cluster, XMMU\\,J0104.1-0635, 6.4\\arcmin \\ to the South-West could be photometrically identified at $z\\!\\simeq\\!0.95$. It was speculated that this object is part of the large-scale structure environment of the main cluster, which could include additional optically selected galaxy overdensities.   %       \\item The cluster shows a pronounced stratification of galaxy populations. The spatial distribution of the red-sequence %population (4 spectroscopic members) of early type galaxies coincides well with the X-ray emission, whereas a significantly bluer %population (3 spectroscopic members) dominates beyond 1--2 core radii from the center, suggesting a cluster environment-driven %effect of differential galaxy evolution, e.g.  delayed star formation quenching in the outskirts.     % ? BCG consistent, confirms that weak sources can be found    % and no clear indications for a color stratisfication of the galaxy populations. %\\item It was shown that all optically color-selected 3--5 sigma peaks are X-ray underluminous/undetected galaxy %overdensities (XUGOs). Geometric alignment is %suggestive that some of these systems are associated with the %large-scale-structure %environment of XMMU\\,J0104.4-0630, which would imply a significant blueing %of the red sequence color in the filaments. %      \\item Further spectroscopic LSS studies in this field have the potential to reconstruct the cosmic web structure at %$z\\!\\sim\\!1$ on the 10\\,Mpc scale and investigate environmental effects of galaxy evolution at a lookback time of 7.5\\,Gyr.   \\end{itemize} %\\end{enumerate}   %\\clearpage    %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%  %--------------------- KAPITEL 11 -----------------------------------------------  \\chapter{Preliminary Studies and Science Outlook} \\label{c10b_science_outlook}  \\noindent The available Z--H follow-up observations aimed primarily at the cluster {\\em identification\\/} and {\\em confirmation\\/} of the principal strategy in terms of limiting depth, photometric precision, and  the choice of filter bands. However, the application of the calibrated Z--H-redshift relation to the full distant cluster candidate sample allows a preliminary study of the luminosity evolution of the {\\em brightest\\/} cluster galaxies in the  redshift range $0.2\\!\\la\\!z\\!\\la\\!1.5$, presented  in Sect.\\,\\ref{s10_bcg_assembly}. Section\\,\\ref{s10_redsequ_population} provides a brief outlook on upcoming investigations of the {\\em faint\\/} end of the cluster red-sequence. The chapter ends with  a gallery of selected, newly discovered high-redshift systems awaiting spectroscopic confirmation. %in Sect.\\,\\ref{s10_highz_gallery}.   %The results presented here should hence motivate additional in-depth studies with designated programs to confirm the observed %effects and improve upon  the achieved accuracy.     % ? Application to photometric studies % preliminary study   %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% %..................... 11.1 .....................................................  \\label{s10b_outlook_conclusions}  \\begin{itemize}  % apply Z--H technique, qualitative results  \\item In Sect.\\,\\ref{s10_bcg_assembly}, the photometric Z--H \\reds redshift estimates were used for a preliminary study of the brightest cluster galaxy evolution by constructing the H-band Hubble diagram for BCGs out to $z\\!\\simeq\\!1.5$.   \\item It was shown, that the   observed X-ray selected cluster  BCGs   are consistent with being `standard candles'  of absolute H-band magnitude $M_{\\mathrm{H}}\\!\\simeq\\!-26.3$\\,mag  out to $z\\!\\simeq\\!1.5$.   \\item The increased redshift baseline allowed to  exclude that the luminosity of BCGs evolves passively at  the 3\\,$\\sigma$ confidence level. Although  a number of possible  selection biases  were identified, which need to be carefully quantified in future studies, none of these  effects could fully account for the observed non-passive evolution component.   \\item The comparison to the expected apparent magnitude evolution of a single model galaxy with passively evolving stellar populations %passive galaxy models revealed the  tentative result that BCGs have at least doubled their mass between redshifts of $z\\!\\simeq\\!1.5$ and  $z\\!\\simeq\\!0.2$, whereas their stellar populations evolve passively.  % observed doubling in total stellar BCG mass  as %an approximate lower limit for the mass growth between redshifts of $z\\!\\simeq\\!1.5$ and  $z\\!\\simeq\\!0.2$.  \\item It was speculated  that the tentatively identified active BCG assembly phase around $z\\!\\sim\\!1$ is accompanied by an increased observed merger rate at the cluster center, followed by an epoch where radially infalling groups of galaxies seem to be important for the final stages of the BCG evolution.  %show that old stellar populations are not a contradiction to the recent BCG assembly \\item The preliminary  results suggest that the well-established old age of the stellar populations of BCGs and the expected late mass assembly can be observationally reconciled.  Predictions for the BCG evolution from  the latest simulations and semi-analytic models are in qualitative agreement with our tentative findings.  %Indications for galaxy transformation processes at work embedded of the probable LSS  \\item Upcoming deep multi-band photometric studies of XDCP clusters at high redshift will allow a systematic investigation of the phenomenon of a \\reds truncation and the related `downsizing' scenario.  \\item A gallery of newly discovered X-ray luminous  high-redshift galaxy cluster candidates was presented. These and other  systems are currently awaiting their spectroscopic redshift confirmation.   % ?? to dos for real study % redshifts, test BCG selection  % ?? next steps: spectroscopy of a number of ho-z cluster, deeper photometry, refined BCG selection, extension of sample % high resolution imaging of BCG merger candidates  \\item The refinement of the  early results of Chap.\\,\\ref{c10_HizClusterStudies} and the preliminary BCG evolution study presented in Sect.\\,\\ref{s10_bcg_assembly} will provide  observational contributions towards an  emerging  consistent picture of galaxy evolution in clusters.  %    In conclusion, this work provided some  observational contributions towards an  emerging  consistent picture %of galaxy evolution in clusters, where the stellar populations of the dominant early-type galaxies are very old, but the major %mass assembly of the most massive objects occurs at a more recent epoch, in qualitatively good agreement with hierarchical %formation models.  %This work could establish observationally that the   \\end{itemize}         %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%  %--------------------- KAPITEL 12 -----------------------------------------------  \\chapter{The Future of Galaxy Cluster Surveys} \\label{c11_Outlook}  \\noindent Before closing this work, a brief outlook of some of the expected major developments in distant galaxy cluster studies over the next decade(s) is presented.    %--> ~8 pages  %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% %................... 11.1 .......................................................   \\noindent When speaking of the {\\em Golden Age for Astronomy} \\ and the dawning %starting {\\em Era of  Precision Cosmology}, one can, without exaggeration, also conceive the next 10--15 years as the {\\em New Age of Galaxy Cluster Surveys}. Here we have only touched on the upcoming developments in X-rays and for the SZ effect. Major survey projects with a strong focus on galaxy clusters are additionally ongoing or in preparation in the optical, near-infrared, and mid-infrared wavebands. Just to name a few, the Dark Energy Survey (DES), the Kilo Degree Survey (KIDS), the Panoramic Survey Telescope and Rapid Response System (PanSTARRS) are optical experiments with expected cluster sample sizes of at least a few thousand systems. At longer wavelength the Visible \\& Infrared Survey Telescope for Astronomy\\footnote{A list of VISTA surveys can be found at  \\url{http://www.eso.org/sci/observing/policies/PublicSurveys/sciencePublicSurveys.html#VISTA}.} (VISTA) or the Spitzer mid-infrared fields\\footnote{A list of Spitzer surveys is available at \\url{http://ssc.spitzer.caltech.edu/legacy/all.html}.} aim at the compilation of hundreds of objects out to high redshift.  The common aim of all projects is the identification of systems out to redshifts of \\zga1, owing to the recognition of galaxy clusters as prime tracers of cosmic evolution. As a survey in an advanced state, the XMM-Newton Distant Cluster Project is in a good position to take on %an explorer role and  smooth the way % pave?? a pathfinder role to largely unexplored redshift regimes over the next few years. In conclusion, the future for distant galaxy cluster studies looks bright. Only the tip of the distant cluster iceberg has been found, the main body is yet to be unveiled.   %now the time has come to unveil the main body...  % as the prime importance of galaxy clusters as tracers of cosmic evolution has be %Major cluster surveys in preparation in basically all wavelength from SZE, MIR, NIR, optical, X-ray  % final statements %future of X-ray surveys looks bright %lots of work ahead  %the era of \\zg1 \\ cluster astrophysics and distant cluster cosmology has just begun. %multi-wavelength %Future for distant cluster studies looks bright...      %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%  %%%%%%%%%%%%%%%%%%%%%%% ANHANG %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% \\begin{appendix}   %%%%%%%%%%%%%%%%%%%%%% Appendix A %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% % >>>>>>>>>>noch Auskommentieren, momentan rausgenommen, da 2 extra Seiten %%\\chapter{XMM-Newton Survey Fields}  %%%%%%%%%%%%%%%%%%%%%% Appendix A %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% \\chapter{XDCP Survey Field List}  %RA  DEC     GL  GB  OBSID   EXPTIME   SPT\\_COVERAGE  (PN\\_CLEAN    MOS1\\_CLEAN  MOS2\\_CLEAN)  EFF\\_CLEAN   PI  TARGETS   %SURVEY\\_STATUS   N\\_H??  %RA & DEC &       GL   &       GB   & NH [E21/cm2] & OBSID       &    EXPTIME  & EFF_EXPTIME & STATUS &              TARGET &      %PI  &       OBSSTART     \\\\    %\\begin{table}[h]    % h = here ; positioning %\\begin{center}  %\\begin{figure}[t] %\\begin{center} %\\parbox{0.45\\textwidth}{ %\\includegraphics[angle=0,clip,width=0.45\\textwidth]{Noras2_cluster_z0.23.eps} %} %\\hfill %\\parbox{0.45\\textwidth}{ %\\caption[NORAS2 Galaxy Cluster]{The NORAS2 cluster of galaxies RXC\\,J0834.9+5534 at $z\\!=\\!0.239$ as a typical example of a %galaxies of early type (red), late type (blue), and Seyfert type\\,1 (yellow).} \\label{f4_SampleCl_NORAS2} %} %\\end{center} %\\end{figure}  %\\ref{tA_field_list}  \\begin{table}[!h] %\\begin{tabular}{|c|c|c|c|c|c|c|} %\\end{tabular} \\caption[XDCP Survey Fields]{XDCP survey field summary. For each of the 546 processed XMM-Newton observations the central equatorial (RA, DEC) and galactic coordinates (GL, GB) are given followed by the corresponding galactic hydrogen column density N$_{\\mathrm{H}}$ in units of $10^{21}$\\,cm$^{-2}$. The last five columns contain the XMM field identifier (OBSID), the nominal exposure time (EXT), the effective clean exposure time (CLT), the processing status (STAT), and the target name of the observation (TARGET). 469 fields have been successfully processed and analyzed (STAT=ok), 48 were highly flared (STAT=flared), and 29 were discarded for other reasons (STAT=discard). } \\label{tA_field_list} \\end{table}    \\begin{longtable}{|c|c|c|c|c|c|c|c|c|l|}  %\\parbox{0.9\\textwidth}{ %\\caption[XDCP Survey Fields]{XDCP survey field summary. For each of the 546 processed XMM-Newton observations the central %equatorial (RA, DEC) and galactic coordinates (GL, GB) are given followed by the corresponding galactic hydrogen column density %N$_{\\mathrm{H}}$ in units of [$10^{21}$\\,cm$^{-2}$]. The last five columns contain the XMM field identifier (OBSID), the nominal %exposure time (EXT), the effective clean exposure time (CLT), the processing status (STAT), and the target name of the %observation (TARGET). 469 fields have been successfully processed and analyzed (STAT=ok), 48 were highly flared (STAT=flared), %and 29 were discarded for other reasons (STAT=discard). } \\\\ %}  \\hline  % \\,cm$^{-2}$ {\\footnotesize\\bf RA} & {\\footnotesize\\bf DEC} &   {\\footnotesize\\bf GL}   &   {\\footnotesize\\bf GB}   & {\\footnotesize\\bf N$_{\\mathbf{H}}$ }  & {\\footnotesize\\bf OBSID}  &    {\\footnotesize\\bf EXT}  & {\\footnotesize\\bf CLT} & {\\footnotesize\\bf STAT} &   {\\footnotesize\\bf TARGET} \\\\ %&   {\\footnotesize\\bf PI}  &   {\\footnotesize\\bf DATE}     \\\\ %[E21/cm2] \\footnotesize h m s & \\footnotesize  d m s & \\footnotesize d.d   & \\footnotesize d.d & \\footnotesize $\\times 10^{21}$ & & \\footnotesize  sec & \\footnotesize sec & & \\\\  \\hline\\hline  \\footnotesize 00 00 29.2 & \\footnotesize  -25 07 20 & \\footnotesize  40.0  & \\footnotesize  -78.3 & \\footnotesize  0.157 & \\footnotesize  0125310101  & \\footnotesize  46065 & \\footnotesize  16360 & \\footnotesize  ok & \\footnotesize  Abell 2690 \\\\ %\\footnotesize  Jansen, Fred & \\footnotesize   2000-06-01 \\\\ \\footnotesize 00 02 51.2 & \\footnotesize  -29 59 11 & \\footnotesize  15.0  & \\footnotesize  -78.9 & \\footnotesize  0.131 & \\footnotesize  0041750101  & \\footnotesize  52270 & \\footnotesize  46458 & \\footnotesize  ok & \\footnotesize  BLANCO1 \\\\ %\\footnotesize  Micela, Giuseppina & \\footnotesize   2002-06-16 \\\\ \\footnotesize 00 03 09.1 & \\footnotesize  -35 57 00 & \\footnotesize  349.2 & \\footnotesize  -76.4 & \\footnotesize  0.113 & \\footnotesize  0145020201  & \\footnotesize  54258 & \\footnotesize  47354 & \\footnotesize  ok & \\footnotesize  A 2717 \\\\ %\\footnotesize  ARNAUD, Monique & \\footnotesize   2002-12-27 \\\\ \\footnotesize 00 03 26.1 & \\footnotesize  -26 02 28 & \\footnotesize  36.0  & \\footnotesize  -79.2 & \\footnotesize  0.163 & \\footnotesize  0103060301  & \\footnotesize  50537 & \\footnotesize  32902 & \\footnotesize  ok & \\footnotesize  Q0000-263 \\\\ %\\footnotesize  Aschenbach, Bernd & \\footnotesize   2002-06-25 \\\\ \\footnotesize 00 00 23.6 & \\footnotesize  +02 04 39 & \\footnotesize  98.0  & \\footnotesize  -58.3 & \\footnotesize  0.300 & \\footnotesize  0201900101  & \\footnotesize  29907 & \\footnotesize  22485 & \\footnotesize  ok & \\footnotesize  RXCJ0003.8+0203 \\\\ %\\footnotesize  Boehringer, Hans & \\footnotesize   2004-06-24 \\\\ \\footnotesize 00 10 30.5 & \\footnotesize  +10 58 43 & \\footnotesize  106.9 & \\footnotesize  -50.6 & \\footnotesize  0.576 & \\footnotesize  0127110201  & \\footnotesize  16311 & \\footnotesize  2625 & \\footnotesize  ok & \\footnotesize  IIIZW2 \\\\ %\\footnotesize  Mason, Keith & \\footnotesize   2000-07-04 \\\\ \\footnotesize 00 11 03.5 & \\footnotesize  -12 07 18 & \\footnotesize  88.7  & \\footnotesize  -72.2 & \\footnotesize  0.257 & \\footnotesize  0150480501  & \\footnotesize  22199 & \\footnotesize  10960 & \\footnotesize  ok & \\footnotesize  NGC34 \\\\ %\\footnotesize  Maiolino, Roberto & \\footnotesize   2002-12-22 \\\\ \\footnotesize 00 11 25.1 & \\footnotesize  -11 28 42 & \\footnotesize  90.0  & \\footnotesize  -71.7 & \\footnotesize  0.268 & \\footnotesize  0111970901  & \\footnotesize  12857 & \\footnotesize  9673 & \\footnotesize  ok & \\footnotesize  WW Cet \\\\ %\\footnotesize  Mason, Keith & \\footnotesize   2001-12-06 \\\\ \\footnotesize 00 14 19.0 & \\footnotesize  -30 22 22 & \\footnotesize  9.0   & \\footnotesize  -81.2 & \\footnotesize  0.165 & \\footnotesize  0042340101  & \\footnotesize  18463 & \\footnotesize  12311 & \\footnotesize  ok & \\footnotesize  RXCJ0014.3-3022 \\\\ %\\footnotesize  Boehringer, Hans & \\footnotesize   2001-05-31 \\\\ \\footnotesize 00 14 34.0 & \\footnotesize  -39 10 56 & \\footnotesize  333.1 & \\footnotesize  -75.7 & \\footnotesize  0.176 & \\footnotesize  0028740101  & \\footnotesize  31669 & \\footnotesize  25236 & \\footnotesize  ok & \\footnotesize  NGC 55 (West) \\\\ %\\footnotesize  Warwick, Robert & \\footnotesize   2001-11-15 \\\\ \\footnotesize 00 15 47.2 & \\footnotesize  -39 15 34 & \\footnotesize  332.1 & \\footnotesize  -75.7 & \\footnotesize  0.176 & \\footnotesize  0028740201  & \\footnotesize  33767 & \\footnotesize  29338 & \\footnotesize  ok & \\footnotesize  NGC 55 (East) \\\\ %\\footnotesize  Warwick, Robert & \\footnotesize   2001-11-14 \\\\ \\footnotesize 00 18 33.6 & \\footnotesize  +16 26 07 & \\footnotesize  111.6 & \\footnotesize  -45.7 & \\footnotesize  0.407 & \\footnotesize  0111000101  & \\footnotesize  37949 & \\footnotesize  26769 & \\footnotesize  ok & \\footnotesize  CL 0016+16 \\\\ %\\footnotesize  Watson, Michael & \\footnotesize   2000-12-30 \\\\ \\footnotesize 00 20 45.6 & \\footnotesize  -25 41 43 & \\footnotesize  42.9  & \\footnotesize  -82.9 & \\footnotesize  0.226 & \\footnotesize  0201900301  & \\footnotesize  31007 & \\footnotesize  12852 & \\footnotesize  ok & \\footnotesize  RXCJ0020.7-2542 \\\\ %\\footnotesize  Boehringer, Hans & \\footnotesize   2004-05-26 \\\\ \\footnotesize 00 21 18.5 & \\footnotesize  -48 39 13 & \\footnotesize  316.1 & \\footnotesize  -67.6 & \\footnotesize  0.262 & \\footnotesize  0152330101  & \\footnotesize  47961 & \\footnotesize  21255 & \\footnotesize  ok & \\footnotesize  SCG 0018-4854 \\\\ %\\footnotesize  Iovino, Angela & \\footnotesize   2002-12-25 \\\\ \\footnotesize 00 26 07.5 & \\footnotesize  +10 41 11 & \\footnotesize  112.8 & \\footnotesize  -51.6 & \\footnotesize  0.498 & \\footnotesize  0001930101 & \\footnotesize  26609 & \\footnotesize  14168 & \\footnotesize  ok & \\footnotesize  IRAS F00235+10 \\\\ %\\footnotesize  Wilman, Richard & \\footnotesize   2001-01-11 \\\\ \\footnotesize 00 26 36.6 & \\footnotesize  +17 09 36 & \\footnotesize  114.4 & \\footnotesize  -45.3 & \\footnotesize  0.431 & \\footnotesize  0050140201  & \\footnotesize  55009 & \\footnotesize  41795 & \\footnotesize  ok & \\footnotesize  Cl0024+17 \\\\ %\\footnotesize  Boehringer, Hans & \\footnotesize   2001-01-07 \\\\ \\footnotesize 00 30 26.7 & \\footnotesize  +04 51 49 & \\footnotesize  113.1 & \\footnotesize  -57.6 & \\footnotesize  0.292 & \\footnotesize  0112320101  & \\footnotesize  31229 & \\footnotesize  8492 & \\footnotesize  ok & \\footnotesize  PSR J0030+0451 \\\\ %\\footnotesize  Aschenbach, Bernd & \\footnotesize   2001-06-20 \\\\ \\footnotesize 00 33 53.0 & \\footnotesize  -43 18 00 & \\footnotesize  314.2 & \\footnotesize  -73.4 & \\footnotesize  0.293 & \\footnotesize  0148961201  & \\footnotesize  19863 & \\footnotesize  15344 & \\footnotesize  ok & \\footnotesize  ELAIS-S1-C \\\\ %\\footnotesize  Fiore, Fabrizio & \\footnotesize   2003-06-06 \\\\ \\footnotesize 00 33 52.9 & \\footnotesize  -43 18 00 & \\footnotesize  314.2 & \\footnotesize  -73.4 & \\footnotesize  0.293 & \\footnotesize  0148961301  & \\footnotesize  48419 & \\footnotesize  32683 & \\footnotesize  ok & \\footnotesize  ELAIS-S1-C \\\\ %\\footnotesize  Fiore, Fabrizio & \\footnotesize   2003-06-06 \\\\ \\footnotesize 00 33 53.6 & \\footnotesize  -12 07 48 & \\footnotesize  106.7 & \\footnotesize  -74.4 & \\footnotesize  0.242 & \\footnotesize  0125920201  & \\footnotesize  50809 & \\footnotesize  7920 & \\footnotesize  ok & \\footnotesize  G158-100 \\\\ %\\footnotesize  Jansen, Fred & \\footnotesize   2000-06-05 \\\\ \\footnotesize 00 33 52.8 & \\footnotesize  -43 18 00 & \\footnotesize  314.2 & \\footnotesize  -73.4 & \\footnotesize  0.293 & \\footnotesize  0148961401  & \\footnotesize  25039 & \\footnotesize  3379 & \\footnotesize  flared & \\footnotesize  ELAIS-S1-C \\\\ %\\footnotesize  Fiore, Fabrizio & \\footnotesize   2003-06-07 \\\\ \\footnotesize 00 34 16.1 & \\footnotesize  -21 25 19 & \\footnotesize  87.5  & \\footnotesize  -83.0 & \\footnotesize  0.156 & \\footnotesize  0044350101  & \\footnotesize  22870 & \\footnotesize  11870 & \\footnotesize  ok & \\footnotesize  IRAS00317-2142 \\\\ %\\footnotesize  Georgantopou, Ioannis & \\footnotesize   2002-05-29 \\\\ \\footnotesize 00 34 24.6 & \\footnotesize  -05 33 40 & \\footnotesize  111.5 & \\footnotesize  -68.0 & \\footnotesize  0.432 & \\footnotesize  0031740101  & \\footnotesize  35738 & \\footnotesize  7057 & \\footnotesize  ok & \\footnotesize  PSR J0034-0534 \\\\ %\\footnotesize  Becker, Werner & \\footnotesize   2002-06-19 \\\\ \\footnotesize 00 35 51.3 & \\footnotesize  -43 17 55 & \\footnotesize  312.9 & \\footnotesize  -73.5 & \\footnotesize  0.263 & \\footnotesize  0148960101  & \\footnotesize  104416 & \\footnotesize  47233 & \\footnotesize  ok & \\footnotesize  ELAIS-S1-A \\\\ %\\footnotesize  Fiore, Fabrizio & \\footnotesize   2003-05-14 \\\\ \\footnotesize 00 37 07.0 & \\footnotesize  +09 09 07 & \\footnotesize  116.9 & \\footnotesize  -53.5 & \\footnotesize  0.478 & \\footnotesize  0084230201  & \\footnotesize  30207 & \\footnotesize  22255 & \\footnotesize  ok & \\footnotesize  Abell 68 \\\\ %\\footnotesize  KNEIB, Jean-Paul & \\footnotesize   2001-12-15 \\\\ \\footnotesize 00 41 42.2 & \\footnotesize  -09 22 26 & \\footnotesize  115.0 & \\footnotesize  -72.0 & \\footnotesize  0.308 & \\footnotesize  0065140101  & \\footnotesize  13253 & \\footnotesize  9026 & \\footnotesize  discard & \\footnotesize  Abell 85 \\\\ %\\footnotesize  Durret, Florence & \\footnotesize   2002-01-07 \\\\ \\footnotesize 00 42 30.8 & \\footnotesize  -09 41 25 & \\footnotesize  115.6 & \\footnotesize  -72.4 & \\footnotesize  0.308 & \\footnotesize  0065140201  & \\footnotesize  13248 & \\footnotesize  10324 & \\footnotesize  ok & \\footnotesize  Abell 85 \\\\ %\\footnotesize  Durret, Florence & \\footnotesize   2002-01-07 \\\\ \\footnotesize 00 43 20.1 & \\footnotesize  +00 51 05 & \\footnotesize  118.6 & \\footnotesize  -61.9 & \\footnotesize  0.245 & \\footnotesize  0090070201  & \\footnotesize  21249 & \\footnotesize  17844 & \\footnotesize  ok & \\footnotesize  UM 269 \\\\ %\\footnotesize  Reeves, James & \\footnotesize   2002-01-05 \\\\ \\footnotesize 00 43 24.8 & \\footnotesize  -20 37 31 & \\footnotesize  106.7 & \\footnotesize  -83.2 & \\footnotesize  0.154 & \\footnotesize  0042340201  & \\footnotesize  15006 & \\footnotesize  8511 & \\footnotesize  ok & \\footnotesize  RXCJ0043.4-2037 \\\\ %\\footnotesize  Boehringer, Hans & \\footnotesize   2002-01-04 \\\\ \\footnotesize 00 43 35.7 & \\footnotesize  -17 59 25 & \\footnotesize  111.3 & \\footnotesize  -80.6 & \\footnotesize  0.170 & \\footnotesize  0112880601  & \\footnotesize  13250 & \\footnotesize  9422 & \\footnotesize  ok & \\footnotesize  HR 188 \\\\ %\\footnotesize  Aschenbach, Bernd & \\footnotesize   2002-01-08 \\\\ \\footnotesize 00 46 51.5 & \\footnotesize  -20 43 15 & \\footnotesize  113.4 & \\footnotesize  -83.5 & \\footnotesize  0.157 & \\footnotesize  0110990301  & \\footnotesize  16889 & \\footnotesize  2953 & \\footnotesize  flared & \\footnotesize  NGC247 \\\\ %\\footnotesize  Watson, Michael & \\footnotesize   2001-07-08 \\\\ \\footnotesize 00 47 23.9 & \\footnotesize  -25 15 25 & \\footnotesize  96.8  & \\footnotesize  -87.9 & \\footnotesize  0.149 & \\footnotesize  0110900101  & \\footnotesize  33732 & \\footnotesize  21384 & \\footnotesize  ok & \\footnotesize  NGC 253 NW \\\\ %\\footnotesize  Watson, Michael & \\footnotesize   2000-12-14 \\\\ \\footnotesize 00 47 36.0 & \\footnotesize  -25 17 52 & \\footnotesize  97.5  & \\footnotesize  -87.9 & \\footnotesize  0.128 & \\footnotesize  0125960201  & \\footnotesize  17537 & \\footnotesize  5496 & \\footnotesize  ok & \\footnotesize  NGC253 \\\\ %\\footnotesize  Jansen, Fred & \\footnotesize   2000-06-04 \\\\ \\footnotesize 00 47 36.2 & \\footnotesize  -25 17 52 & \\footnotesize  97.5  & \\footnotesize  -87.9 & \\footnotesize  0.128 & \\footnotesize  0125960101  & \\footnotesize  60809 & \\footnotesize  31583 & \\footnotesize  ok & \\footnotesize  NGC253 \\\\ %\\footnotesize  Jansen, Fred & \\footnotesize   2000-06-03 \\\\ \\footnotesize 00 47 39.3 & \\footnotesize  -25 17 08 & \\footnotesize  98.0  & \\footnotesize  -87.9 & \\footnotesize  0.128 & \\footnotesize  0152020101  & \\footnotesize  140799 & \\footnotesize  63718 & \\footnotesize  ok & \\footnotesize  NGC 253 \\\\ %\\footnotesize  Pietsch, Wolfgang & \\footnotesize   2003-06-20 \\\\ \\footnotesize 00 49 59.2 & \\footnotesize  -52 08 58 & \\footnotesize  303.4 & \\footnotesize  -64.9 & \\footnotesize  0.340 & \\footnotesize  0133120301  & \\footnotesize  12022 & \\footnotesize  8107 & \\footnotesize  ok & \\footnotesize  BPM 16274 \\\\ %\\footnotesize  Jansen, Fred & \\footnotesize   2000-12-12 \\\\ \\footnotesize 00 50 01.3 & \\footnotesize  -52 08 11 & \\footnotesize  303.4 & \\footnotesize  -64.9 & \\footnotesize  0.340 & \\footnotesize  0123920101  & \\footnotesize  57464 & \\footnotesize  8833 & \\footnotesize  ok & \\footnotesize  BPM 16274 \\\\ %\\footnotesize  Jansen, Fred & \\footnotesize   2000-05-18 \\\\ \\footnotesize 00 50 03.0 & \\footnotesize  -52 08 19 & \\footnotesize  303.4 & \\footnotesize  -64.9 & \\footnotesize  0.340 & \\footnotesize  0125320401  & \\footnotesize  33728 & \\footnotesize  17775 & \\footnotesize  ok & \\footnotesize  BPM16274 \\\\ %\\footnotesize  Jansen, Fred & \\footnotesize   2000-11-24 \\\\ \\footnotesize 00 50 03.3 & \\footnotesize  -52 08 17 & \\footnotesize  303.4 & \\footnotesize  -64.9 & \\footnotesize  0.340 & \\footnotesize  0125320701  & \\footnotesize  45951 & \\footnotesize  10346 & \\footnotesize  ok & \\footnotesize  BPM16274 \\\\ %\\footnotesize  Jansen, Fred & \\footnotesize   2001-05-13 \\\\ \\footnotesize 00 50 07.1 & \\footnotesize  -52 09 41 & \\footnotesize  303.4 & \\footnotesize  -64.9 & \\footnotesize  0.340 & \\footnotesize  0133120401  & \\footnotesize  13707 & \\footnotesize  8611 & \\footnotesize  ok & \\footnotesize  BPM 16274 \\\\ %\\footnotesize  Jansen, Fred & \\footnotesize   2000-12-13 \\\\ \\footnotesize 00 54 55.5 & \\footnotesize  -37 41 09 & \\footnotesize  299.1 & \\footnotesize  -79.4 & \\footnotesize  0.361 & \\footnotesize  0112800101  & \\footnotesize  46711 & \\footnotesize  38155 & \\footnotesize  ok & \\footnotesize  NGC 300 \\\\ %\\footnotesize  Turner, Martin & \\footnotesize   2001-01-02 \\\\ \\footnotesize 00 54 55.5 & \\footnotesize  -37 41 08 & \\footnotesize  299.1 & \\footnotesize  -79.4 & \\footnotesize  0.361 & \\footnotesize  0112800201  & \\footnotesize  36909 & \\footnotesize  27667 & \\footnotesize  ok & \\footnotesize  NGC 300 \\\\ %\\footnotesize  Turner, Martin & \\footnotesize   2000-12-27 \\\\ \\footnotesize 00 56 15.0 & \\footnotesize  -01 18 39 & \\footnotesize  125.6 & \\footnotesize  -64.1 & \\footnotesize  0.310 & \\footnotesize  0012440101  & \\footnotesize  34992 & \\footnotesize  25332 & \\footnotesize  discard & \\footnotesize  LBQS0053-0134 \\\\ %& \\footnotesize  Elvis, Martin & \\footnotesize   2001-01-15 \\\\ \\footnotesize 00 58 37.0 & \\footnotesize  -36 06 02 & \\footnotesize  293.7 & \\footnotesize  -80.8 & \\footnotesize  0.194 & \\footnotesize  0102040701  & \\footnotesize  21536 & \\footnotesize  7855 & \\footnotesize  ok & \\footnotesize  Q 0056-363 \\\\ %\\footnotesize  Jansen, Fred & \\footnotesize   2000-07-05 \\\\ \\footnotesize 00 58 48.8 & \\footnotesize  -28 25 00 & \\footnotesize  251.6 & \\footnotesize  -87.9 & \\footnotesize  0.195 & \\footnotesize  0111282301  & \\footnotesize  14075 & \\footnotesize  650 & \\footnotesize  flared & \\footnotesize  SGP-9 \\\\ %\\footnotesize  Watson, Michael & \\footnotesize   2002-12-01 \\\\ \\footnotesize 00 58 53.8 & \\footnotesize  -27 59 00 & \\footnotesize  240.7 & \\footnotesize  -88.1 & \\footnotesize  0.195 & \\footnotesize  0111280601  & \\footnotesize  10944 & \\footnotesize  7947 & \\footnotesize  ok & \\footnotesize  SGP-6 \\\\ %\\footnotesize  Watson, Michael & \\footnotesize   2002-06-19 \\\\ \\footnotesize 01 02 41.4 & \\footnotesize  -21 52 51 & \\footnotesize  149.5 & \\footnotesize  -84.1 & \\footnotesize  0.160 & \\footnotesize  0144310101  & \\footnotesize  33917 & \\footnotesize  17992 & \\footnotesize  ok & \\footnotesize  Abell 133 \\\\ %\\footnotesize  Sarazin, Craig & \\footnotesize   2002-12-23 \\\\ \\footnotesize 01 03 59.4 & \\footnotesize  -06 41 53 & \\footnotesize  131.8 & \\footnotesize  -69.3 & \\footnotesize  0.506 & \\footnotesize  0112650501  & \\footnotesize  25035 & \\footnotesize  16432 & \\footnotesize  ok & \\footnotesize  G133-69 Pos\\_2 \\\\ %\\footnotesize  Aschenbach, Bernd & \\footnotesize   2000-07-03 \\\\ \\footnotesize 01 04 24.6 & \\footnotesize  -06 24 10 & \\footnotesize  131.9 & \\footnotesize  -69.0 & \\footnotesize  0.508 & \\footnotesize  0112650401  & \\footnotesize  27145 & \\footnotesize  18405 & \\footnotesize  ok & \\footnotesize  G133-69 Pos\\_1 \\\\ %\\footnotesize  Aschenbach, Bernd & \\footnotesize   2000-12-18 \\\\ \\footnotesize 01 06 22.2 & \\footnotesize  +00 49 11 & \\footnotesize  130.8 & \\footnotesize  -61.8 & \\footnotesize  0.316 & \\footnotesize  0150870201  & \\footnotesize  28815 & \\footnotesize  715 & \\footnotesize  flared & \\footnotesize  BRI0103+0032 \\\\ %\\footnotesize  Bechtold, Jill & \\footnotesize   2003-07-15 \\\\ \\footnotesize 01 06 55.9 & \\footnotesize  -80 17 22 & \\footnotesize  302.1 & \\footnotesize  -36.8 & \\footnotesize  0.664 & \\footnotesize  0145540101  & \\footnotesize  31315 & \\footnotesize  22348 & \\footnotesize  ok & \\footnotesize  IRASF01063-8034 \\\\ %\\footnotesize  Greenhill, Lincoln & \\footnotesize   2003-04-21 \\\\ \\footnotesize 01 07 45.8 & \\footnotesize  -17 30 14 & \\footnotesize  145.1 & \\footnotesize  -79.6 & \\footnotesize  0.140 & \\footnotesize  0025540101  & \\footnotesize  12572 & \\footnotesize  3361 & \\footnotesize  flared & \\footnotesize  IC1623 \\\\ %\\footnotesize  Iwasawa, Kazushi & \\footnotesize   2001-06-26 \\\\ \\footnotesize 01 12 53.4 & \\footnotesize  -45 32 04 & \\footnotesize  291.2 & \\footnotesize  -71.1 & \\footnotesize  0.210 & \\footnotesize  0067170101  & \\footnotesize  47870 & \\footnotesize  39004 & \\footnotesize  ok & \\footnotesize  Phoenix field \\\\ %\\footnotesize  Georgakakis, Antonis & \\footnotesize   2002-05-27 \\\\ \\footnotesize 01 13 49.9 & \\footnotesize  -14 50 40 & \\footnotesize  147.0 & \\footnotesize  -76.6 & \\footnotesize  0.163 & \\footnotesize  0147920101  & \\footnotesize  26915 & \\footnotesize  20479 & \\footnotesize  ok & \\footnotesize  MRK 1152 \\\\ %\\footnotesize  Schartel, Norbert & \\footnotesize   2003-06-15 \\\\ \\footnotesize 01 18 56.5 & \\footnotesize  -00 59 16 & \\footnotesize  138.2 & \\footnotesize  -63.0 & \\footnotesize  0.371 & \\footnotesize  0153170101  & \\footnotesize  22204 & \\footnotesize  8626 & \\footnotesize  ok & \\footnotesize  MS 0116.3-0115 \\\\ %\\footnotesize  Lewis, Aaron & \\footnotesize   2003-07-12 \\\\ \\footnotesize 01 20 20.8 & \\footnotesize  -44 07 51 & \\footnotesize  285.9 & \\footnotesize  -72.0 & \\footnotesize  0.241 & \\footnotesize  0103860901  & \\footnotesize  22608 & \\footnotesize  18679 & \\footnotesize  ok & \\footnotesize  ESO 244- G 017 \\\\ %\\footnotesize  Aschenbach, Bernd & \\footnotesize   2000-12-18 \\\\ \\footnotesize 01 20 39.0 & \\footnotesize  -10 56 27 & \\footnotesize  147.3 & \\footnotesize  -72.4 & \\footnotesize  0.353 & \\footnotesize  0113040801  & \\footnotesize  10768 & \\footnotesize  3539 & \\footnotesize  flared & \\footnotesize  C2001 \\\\ %\\footnotesize  Jansen, Fred & \\footnotesize   2001-06-27 \\\\ \\footnotesize 01 22 43.0 & \\footnotesize  -11 14 24 & \\footnotesize  149.2 & \\footnotesize  -72.4 & \\footnotesize  0.353 & \\footnotesize  0113040701  & \\footnotesize  10045 & \\footnotesize  2455 & \\footnotesize  flared & \\footnotesize  C2001 \\\\ %\\footnotesize  Jansen, Fred & \\footnotesize   2001-06-27 \\\\ \\footnotesize 01 23 46.1 & \\footnotesize  -58 48 19 & \\footnotesize  295.0 & \\footnotesize  -57.8 & \\footnotesize  0.282 & \\footnotesize  0101040201  & \\footnotesize  33006 & \\footnotesize  26928 & \\footnotesize  ok & \\footnotesize  Fairall 9 \\\\ %\\footnotesize  Jansen, Fred & \\footnotesize   2000-07-06 \\\\ \\footnotesize 01 24 35.7 & \\footnotesize  +03 47 20 & \\footnotesize  138.7 & \\footnotesize  -58.0 & \\footnotesize  0.337 & \\footnotesize  0025541601  & \\footnotesize  13203 & \\footnotesize  10287 & \\footnotesize  ok & \\footnotesize  NGC520 \\\\ %\\footnotesize  Iwasawa, Kazushi & \\footnotesize   2002-01-01 \\\\ \\footnotesize 01 25 33.4 & \\footnotesize  +01 45 35 & \\footnotesize  140.1 & \\footnotesize  -59.9 & \\footnotesize  0.308 & \\footnotesize  0109860101  & \\footnotesize  41694 & \\footnotesize  33368 & \\footnotesize  ok & \\footnotesize  A 189 \\\\ %\\footnotesize  Kaastra, Jelle & \\footnotesize   2001-01-01 \\\\ \\footnotesize 01 25 47.3 & \\footnotesize  -01 23 55 & \\footnotesize  142.1 & \\footnotesize  -62.9 & \\footnotesize  0.356 & \\footnotesize  0136340101  & \\footnotesize  22515 & \\footnotesize  16496 & \\footnotesize  ok & \\footnotesize  Abell 194 \\\\ %\\footnotesize  Furuzawa, Akihiro & \\footnotesize   2002-12-24 \\\\ \\footnotesize 01 31 54.3 & \\footnotesize  -13 36 58 & \\footnotesize  159.9 & \\footnotesize  -73.5 & \\footnotesize  0.156 & \\footnotesize  0084230301  & \\footnotesize  24709 & \\footnotesize  14485 & \\footnotesize  ok & \\footnotesize  Abell 209 \\\\ %\\footnotesize  KNEIB, Jean-Paul & \\footnotesize   2001-01-16 \\\\ \\footnotesize 01 33 01.1 & \\footnotesize  -40 06 27 & \\footnotesize  272.0 & \\footnotesize  -74.4 & \\footnotesize  0.201 & \\footnotesize  0112630201  & \\footnotesize  37870 & \\footnotesize  21297 & \\footnotesize  ok & \\footnotesize  QSO 0130-403 \\\\ %\\footnotesize  Turner, Martin & \\footnotesize   2001-06-04 \\\\ \\footnotesize 01 36 20.8 & \\footnotesize  +15 44 24 & \\footnotesize  138.5 & \\footnotesize  -45.7 & \\footnotesize  0.483 & \\footnotesize  0154350201  & \\footnotesize  24914 & \\footnotesize  22087 & \\footnotesize  ok & \\footnotesize  SN2002ap \\\\ %\\footnotesize  Jansen, Fred & \\footnotesize   2003-01-07 \\\\ \\footnotesize 01 36 24.3 & \\footnotesize  +15 45 02 & \\footnotesize  138.5 & \\footnotesize  -45.7 & \\footnotesize  0.483 & \\footnotesize  0154350101  & \\footnotesize  37259 & \\footnotesize  26902 & \\footnotesize  ok & \\footnotesize  SN2002ap \\\\ %\\footnotesize  Jansen, Fred & \\footnotesize   2002-02-02 \\\\ \\footnotesize 01 39 52.9 & \\footnotesize  +06 18 31 & \\footnotesize  144.0 & \\footnotesize  -54.5 & \\footnotesize  0.371 & \\footnotesize  0110890901  & \\footnotesize  27412 & \\footnotesize  14774 & \\footnotesize  ok & \\footnotesize  PHL 1092 \\\\ %\\footnotesize  Watson, Michael & \\footnotesize   2003-01-19 \\\\ \\footnotesize 01 40 10.5 & \\footnotesize  -67 48 45 & \\footnotesize  296.0 & \\footnotesize  -48.6 & \\footnotesize  0.235 & \\footnotesize  0032140401  & \\footnotesize  12561 & \\footnotesize  5279 & \\footnotesize  ok & \\footnotesize  1saxj0140.2-6748 \\\\ %\\footnotesize  Fiore, Fabrizio & \\footnotesize   2001-05-03 \\\\ \\footnotesize 01 40 55.2 & \\footnotesize  -67 54 28 & \\footnotesize  295.9 & \\footnotesize  -48.5 & \\footnotesize  0.235 & \\footnotesize  0148440101  & \\footnotesize  29559 & \\footnotesize  26028 & \\footnotesize  ok & \\footnotesize  BL Hyi \\\\ %\\footnotesize  TERADA, Yukikatsu & \\footnotesize   2002-12-17 \\\\ \\footnotesize 01 43 01.9 & \\footnotesize  +13 38 22 & \\footnotesize  141.6 & \\footnotesize  -47.3 & \\footnotesize  0.476 & \\footnotesize  0093641001  & \\footnotesize  12209 & \\footnotesize  8034 & \\footnotesize  ok & \\footnotesize  NGC 660 \\\\ %\\footnotesize  Bauer, Franz & \\footnotesize   2001-01-07 \\\\ \\footnotesize 01 46 14.8 & \\footnotesize  -39 22 05 & \\footnotesize  263.4 & \\footnotesize  -73.2 & \\footnotesize  0.186 & \\footnotesize  0090070101  & \\footnotesize  33020 & \\footnotesize  29584 & \\footnotesize  ok & \\footnotesize  Q 0144-3938 \\\\ %\\footnotesize  Reeves, James & \\footnotesize   2002-06-18 \\\\ \\footnotesize 01 48 59.9 & \\footnotesize  +05 55 10 & \\footnotesize  147.8 & \\footnotesize  -54.1 & \\footnotesize  0.445 & \\footnotesize  0112551501  & \\footnotesize  21446 & \\footnotesize  14748 & \\footnotesize  ok & \\footnotesize  ngc676 \\\\ %\\footnotesize  Turner, Martin & \\footnotesize   2002-07-14 \\\\ \\footnotesize 01 52 42.1 & \\footnotesize  +01 00 30 & \\footnotesize  153.0 & \\footnotesize  -58.2 & \\footnotesize  0.284 & \\footnotesize  0084230401  & \\footnotesize  30199 & \\footnotesize  13453 & \\footnotesize  ok & \\footnotesize  Abell 267 \\\\ %\\footnotesize  KNEIB, Jean-Paul & \\footnotesize   2002-01-02 \\\\ \\footnotesize 01 52 40.0 & \\footnotesize  -13 58 52 & \\footnotesize  173.3 & \\footnotesize  -70.5 & \\footnotesize  0.147 & \\footnotesize  0109540101  & \\footnotesize  53917 & \\footnotesize  42618 & \\footnotesize  ok & \\footnotesize  WARP J0152.7-13 \\\\ %\\footnotesize  Mason, Keith & \\footnotesize   2002-12-24 \\\\ \\footnotesize 01 53 03.9 & \\footnotesize  -13 43 25 & \\footnotesize  173.0 & \\footnotesize  -70.3 & \\footnotesize  0.169 & \\footnotesize  0112300101  & \\footnotesize  47478 & \\footnotesize  18978 & \\footnotesize  ok & \\footnotesize  ngc 720 \\\\ %\\footnotesize  Turner, Martin & \\footnotesize   2002-07-13 \\\\ \\footnotesize 01 53 37.2 & \\footnotesize  -49 36 53 & \\footnotesize  279.1 & \\footnotesize  -64.5 & \\footnotesize  0.241 & \\footnotesize  0092970201  & \\footnotesize  14368 & \\footnotesize  10001 & \\footnotesize  ok & \\footnotesize  ESO 197-G010 \\\\ %\\footnotesize  Goudfrooij, Paul & \\footnotesize   2001-05-30 \\\\ \\footnotesize 01 55 46.2 & \\footnotesize  +00 28 58 & \\footnotesize  154.7 & \\footnotesize  -58.3 & \\footnotesize  0.284 & \\footnotesize  0145450201  & \\footnotesize  12618 & \\footnotesize  2216 & \\footnotesize  flared & \\footnotesize  SDSS015543+00 \\\\ %\\footnotesize  Szkody, Paula & \\footnotesize   2003-07-04 \\\\ \\footnotesize 01 59 49.6 & \\footnotesize  +00 23 52 & \\footnotesize  156.5 & \\footnotesize  -57.9 & \\footnotesize  0.251 & \\footnotesize  0101640201  & \\footnotesize  14778 & \\footnotesize  4169 & \\footnotesize  flared & \\footnotesize  Mrk 1014 \\\\ %\\footnotesize  Aschenbach, Bernd & \\footnotesize   2000-07-30 \\\\ \\footnotesize 02 07 39.4 & \\footnotesize  +02 08 10 & \\footnotesize  158.0 & \\footnotesize  -55.4 & \\footnotesize  0.311 & \\footnotesize  0052140301  & \\footnotesize  35006 & \\footnotesize  26292 & \\footnotesize  ok & \\footnotesize  HCG 15 \\\\ %\\footnotesize  Lloyd-Davies, Edward & \\footnotesize   2002-01-11 \\\\ \\footnotesize 02 07 52.8 & \\footnotesize  +02 43 46 & \\footnotesize  157.6 & \\footnotesize  -54.9 & \\footnotesize  0.350 & \\footnotesize  0048740101  & \\footnotesize  47340 & \\footnotesize  9723 & \\footnotesize  ok & \\footnotesize  NAB 0205+024 \\\\ %\\footnotesize  Boller, Thomas & \\footnotesize   2002-07-23 \\\\ \\footnotesize 02 09 51.5 & \\footnotesize  -63 18 38 & \\footnotesize  288.8 & \\footnotesize  -51.6 & \\footnotesize  0.294 & \\footnotesize  0111970401  & \\footnotesize  15197 & \\footnotesize  6033 & \\footnotesize  ok & \\footnotesize  WX Hyi \\\\ %\\footnotesize  Mason, Keith & \\footnotesize   2002-01-08 \\\\ \\footnotesize 02 15 13.1 & \\footnotesize  -73 59 08 & \\footnotesize  295.3 & \\footnotesize  -41.8 & \\footnotesize  0.339 & \\footnotesize  0049150101  & \\footnotesize  29614 & \\footnotesize  5694 & \\footnotesize  ok & \\footnotesize  Magellanic Bridge \\\\ %\\footnotesize  Snowden, Steve & \\footnotesize   2002-04-01 \\\\ \\footnotesize 02 15 19.6 & \\footnotesize  -73 59 50 & \\footnotesize  295.3 & \\footnotesize  -41.8 & \\footnotesize  0.339 & \\footnotesize  0099840101  & \\footnotesize  14256 & \\footnotesize  9602 & \\footnotesize  ok & \\footnotesize  Magellanic Bridge \\\\ %\\footnotesize  Mushotzky, Richard & \\footnotesize   2000-07-05 \\\\ \\footnotesize 02 15 27.2 & \\footnotesize  -74 01 09 & \\footnotesize  295.3 & \\footnotesize  -41.8 & \\footnotesize  0.339 & \\footnotesize  0049150201  & \\footnotesize  28609 & \\footnotesize  4147 & \\footnotesize flared &\\footnotesize  Magellanic Bridge\\\\ %& \\footnotesize  Snowden, Steve & \\footnotesize   2002-04-03 \\\\ \\footnotesize 02 15 35.6 & \\footnotesize  -73 59 57 & \\footnotesize  295.2 & \\footnotesize  -41.8 & \\footnotesize  0.339 & \\footnotesize  0049150301  & \\footnotesize  31426 & \\footnotesize  15332 & \\footnotesize  ok & \\footnotesize  Magellanic Bridge \\\\ %\\footnotesize  Snowden, Steve & \\footnotesize   2002-06-16 \\\\ \\footnotesize 02 19 23.7 & \\footnotesize  -02 57 50 & \\footnotesize  167.7 & \\footnotesize  -57.9 & \\footnotesize  0.244 & \\footnotesize  0148500201  & \\footnotesize  12218 & \\footnotesize  10172 & \\footnotesize  ok & \\footnotesize  Mira \\\\ %\\footnotesize  Kastner, Joel & \\footnotesize   2003-07-24 \\\\ \\footnotesize 02 22 00.2 & \\footnotesize  -04 50 12 & \\footnotesize  171.0 & \\footnotesize  -58.9 & \\footnotesize  0.255 & \\footnotesize  0112680801  & \\footnotesize  15617 & \\footnotesize  10372 & \\footnotesize  ok & \\footnotesize  MLS-4 \\\\ %\\footnotesize  Turner, Martin & \\footnotesize   2002-02-01 \\\\ \\footnotesize 02 24 03.2 & \\footnotesize  -04 29 12 & \\footnotesize  171.3 & \\footnotesize  -58.3 & \\footnotesize  0.251 & \\footnotesize  0112680501  & \\footnotesize  23645 & \\footnotesize  17817 & \\footnotesize  ok & \\footnotesize  MLS-8 \\\\ %\\footnotesize  Turner, Martin & \\footnotesize   2002-07-25 \\\\ \\footnotesize 02 24 36.9 & \\footnotesize  -04 10 51 & \\footnotesize  171.1 & \\footnotesize  -58.0 & \\footnotesize  0.232 & \\footnotesize  0112680301  & \\footnotesize  23411 & \\footnotesize  19205 & \\footnotesize  ok & \\footnotesize  MLS-3 \\\\ %\\footnotesize  Turner, Martin & \\footnotesize   2003-01-19 \\\\ \\footnotesize 02 25 08.0 & \\footnotesize  +18 47 37 & \\footnotesize  151.7 & \\footnotesize  -38.7 & \\footnotesize  1.042 & \\footnotesize  0150180101  & \\footnotesize  22208 & \\footnotesize  19599 & \\footnotesize  ok & \\footnotesize  RBS 315 \\\\ %\\footnotesize  Lamer, Georg & \\footnotesize   2003-07-25 \\\\ \\footnotesize 02 25 13.3 & \\footnotesize  -29 27 30 & \\footnotesize  224.9 & \\footnotesize  -69.2 & \\footnotesize  0.170 & \\footnotesize  0201900701  & \\footnotesize  29481 & \\footnotesize  0 & \\footnotesize  flared & \\footnotesize  RXCJ0225.1-2928 \\\\ %\\footnotesize  Boehringer, Hans & \\footnotesize   2004-07-06 \\\\ \\footnotesize 02 25 20.5 & \\footnotesize  -04 30 10 & \\footnotesize  171.7 & \\footnotesize  -58.1 & \\footnotesize  0.232 & \\footnotesize  0112681001  & \\footnotesize  41946 & \\footnotesize  19995 & \\footnotesize  ok & \\footnotesize  MLS-7 \\\\ %\\footnotesize  Turner, Martin & \\footnotesize   2002-01-31 \\\\ \\footnotesize 02 26 02.8 & \\footnotesize  -04 09 10 & \\footnotesize  171.5 & \\footnotesize  -57.7 & \\footnotesize  0.232 & \\footnotesize  0112680201  & \\footnotesize  21648 & \\footnotesize  8058 & \\footnotesize  ok & \\footnotesize  MLS-2 \\\\ %\\footnotesize  Turner, Martin & \\footnotesize   2002-07-14 \\\\ \\footnotesize 02 26 43.0 & \\footnotesize  -04 29 13 & \\footnotesize  172.2 & \\footnotesize  -57.9 & \\footnotesize  0.232 & \\footnotesize  0112681301  & \\footnotesize  40381 & \\footnotesize  13322 & \\footnotesize  ok & \\footnotesize  MLS-6 \\\\ %\\footnotesize  Turner, Martin & \\footnotesize   2002-07-26 \\\\ \\footnotesize 02 27 20.6 & \\footnotesize  -04 10 11 & \\footnotesize  172.0 & \\footnotesize  -57.5 & \\footnotesize  0.273 & \\footnotesize  0112680101  & \\footnotesize  30351 & \\footnotesize  22689 & \\footnotesize  ok & \\footnotesize  MLS-1 \\\\ %\\footnotesize  Turner, Martin & \\footnotesize   2002-01-29 \\\\ \\footnotesize 02 28 00.4 & \\footnotesize  -04 30 10 & \\footnotesize  172.7 & \\footnotesize  -57.7 & \\footnotesize  0.273 & \\footnotesize  0112680401  & \\footnotesize  25050 & \\footnotesize  20187 & \\footnotesize  ok & \\footnotesize  MLS-5 \\\\ %\\footnotesize  Turner, Martin & \\footnotesize   2002-02-03 \\\\ \\footnotesize 02 28 56.1 & \\footnotesize  -10 10 25 & \\footnotesize  181.1 & \\footnotesize  -61.4 & \\footnotesize  0.260 & \\footnotesize  0112860101  & \\footnotesize  10035 & \\footnotesize  4697 & \\footnotesize  ok & \\footnotesize  Q0226-1024 \\\\ %\\footnotesize  Griffiths, Richard & \\footnotesize   2000-07-31 \\\\ \\footnotesize 02 32 21.0 & \\footnotesize  -44 19 45 & \\footnotesize  259.9 & \\footnotesize  -63.4 & \\footnotesize  0.261 & \\footnotesize  0042340301  & \\footnotesize  13645 & \\footnotesize  9640 & \\footnotesize  ok & \\footnotesize  RXCJ0232.2-4420 \\\\ %\\footnotesize  Boehringer, Hans & \\footnotesize   2002-07-11 \\\\ \\footnotesize 02 34 22.6 & \\footnotesize  -43 47 43 & \\footnotesize  258.4 & \\footnotesize  -63.4 & \\footnotesize  0.261 & \\footnotesize  0148790101  & \\footnotesize  48479 & \\footnotesize  19176 & \\footnotesize  ok & \\footnotesize  CC Eri \\\\ %\\footnotesize  Kay, Hilary & \\footnotesize   2003-08-08 \\\\ \\footnotesize 02 34 41.4 & \\footnotesize  -08 46 35 & \\footnotesize  180.8 & \\footnotesize  -59.4 & \\footnotesize  0.283 & \\footnotesize  0150470601  & \\footnotesize  57917 & \\footnotesize  12453 & \\footnotesize  ok & \\footnotesize  NGC 985 \\\\ %\\footnotesize  Santos-Lleo, Maria & \\footnotesize   2003-07-16 \\\\ \\footnotesize 02 36 11.5 & \\footnotesize  -52 19 16 & \\footnotesize  272.2 & \\footnotesize  -58.1 & \\footnotesize  0.308 & \\footnotesize  0098810101  & \\footnotesize  24995 & \\footnotesize  20509 & \\footnotesize  ok & \\footnotesize  WW Hor \\\\ %\\footnotesize  Mason, Keith & \\footnotesize   2000-12-04 \\\\ \\footnotesize 02 38 19.4 & \\footnotesize  -52 11 38 & \\footnotesize  271.6 & \\footnotesize  -57.9 & \\footnotesize  0.307 & \\footnotesize  0067190101  & \\footnotesize  34146 & \\footnotesize  18519 & \\footnotesize  ok & \\footnotesize  ESO 198-G24 \\\\ %\\footnotesize  PORQUET, Delphine & \\footnotesize   2001-01-25 \\\\ \\footnotesize 02 38 38.5 & \\footnotesize  +16 36 49 & \\footnotesize  156.7 & \\footnotesize  -39.1 & \\footnotesize  0.885 & \\footnotesize  0110990101  & \\footnotesize  19851 & \\footnotesize  16422 & \\footnotesize  ok & \\footnotesize  AO0235+164 \\\\ %\\footnotesize  Watson, Michael & \\footnotesize   2002-02-10 \\\\ \\footnotesize 02 41 04.4 & \\footnotesize  -08 15 12 & \\footnotesize  182.0 & \\footnotesize  -57.9 & \\footnotesize  0.308 & \\footnotesize  0093630101  & \\footnotesize  16463 & \\footnotesize  12911 & \\footnotesize  ok & \\footnotesize  NGC 1052 \\\\ %\\footnotesize  Weaver, Kimberly & \\footnotesize   2001-08-15 \\\\ \\footnotesize 02 42 40.9 & \\footnotesize  -00 00 47 & \\footnotesize  172.1 & \\footnotesize  -51.9 & \\footnotesize  0.355 & \\footnotesize  0111200101  & \\footnotesize  42258 & \\footnotesize  33615 & \\footnotesize  ok & \\footnotesize  NGC 1068 \\\\ %\\footnotesize  Kaastra, Jelle & \\footnotesize   2000-07-29 \\\\ \\footnotesize 02 42 41.0 & \\footnotesize  -00 00 46 & \\footnotesize  172.1 & \\footnotesize  -51.9 & \\footnotesize  0.355 & \\footnotesize  0111200201  & \\footnotesize  46429 & \\footnotesize  28657 & \\footnotesize  ok & \\footnotesize  NGC 1068 \\\\ %\\footnotesize  Kaastra, Jelle & \\footnotesize   2000-07-30 \\\\ \\footnotesize 02 48 06.6 & \\footnotesize  -03 30 57 & \\footnotesize  177.7 & \\footnotesize  -53.4 & \\footnotesize  0.412 & \\footnotesize  0084230501  & \\footnotesize  33645 & \\footnotesize  25182 & \\footnotesize  ok & \\footnotesize  Abell 383 \\\\ %\\footnotesize  KNEIB, Jean-Paul & \\footnotesize   2002-08-17 \\\\ \\footnotesize 02 49 33.7 & \\footnotesize  -31 11 24 & \\footnotesize  229.0 & \\footnotesize  -63.9 & \\footnotesize  0.227 & \\footnotesize  0146510401  & \\footnotesize  38914 & \\footnotesize  29923 & \\footnotesize  ok & \\footnotesize  IC1860 Group \\\\ %\\footnotesize  Mulchaey, John & \\footnotesize   2003-02-04 \\\\ \\footnotesize 02 52 54.5 & \\footnotesize  -01 15 39 & \\footnotesize  176.4 & \\footnotesize  -51.0 & \\footnotesize  0.525 & \\footnotesize  0151490101  & \\footnotesize  47204 & \\footnotesize  19428 & \\footnotesize  ok & \\footnotesize  NGC 1132 \\\\ %\\footnotesize  Mathews, William & \\footnotesize   2003-07-16 \\\\ \\footnotesize 02 56 05.6 & \\footnotesize  +19 22 48 & \\footnotesize  159.2 & \\footnotesize  -34.5 & \\footnotesize  1.065 & \\footnotesize  0112510301  & \\footnotesize  30242 & \\footnotesize  20763 & \\footnotesize  ok & \\footnotesize  MBM 12 Pos\\_2 \\\\ %\\footnotesize  Aschenbach, Bernd & \\footnotesize   2001-02-05 \\\\ \\footnotesize 02 56 04.5 & \\footnotesize  +19 28 50 & \\footnotesize  159.2 & \\footnotesize  -34.4 & \\footnotesize  1.096 & \\footnotesize  0110660101  & \\footnotesize  24196 & \\footnotesize  10395 & \\footnotesize  ok & \\footnotesize  MBM 12 \\\\ %\\footnotesize  Jansen, Fred & \\footnotesize   2000-08-26 \\\\ \\footnotesize 02 56 33.4 & \\footnotesize  +00 06 00 & \\footnotesize  175.9 & \\footnotesize  -49.4 & \\footnotesize  0.665 & \\footnotesize  0056020301  & \\footnotesize  25094 & \\footnotesize  8676 & \\footnotesize  ok & \\footnotesize  RX J0256.5+0006 \\\\ %\\footnotesize  ARNAUD, Monique & \\footnotesize   2001-02-15 \\\\ \\footnotesize 02 57 52.5 & \\footnotesize  +13 02 08 & \\footnotesize  164.3 & \\footnotesize  -39.4 & \\footnotesize  1.058 & \\footnotesize  0112260101  & \\footnotesize  15409 & \\footnotesize  9110 & \\footnotesize  ok & \\footnotesize  A399 \\\\ %& \\footnotesize  %Turner, Martin & \\footnotesize   2000-08-18 \\\\ \\footnotesize 02 58 58.4 & \\footnotesize  +13 33 47 & \\footnotesize  164.2 & \\footnotesize  -38.8 & \\footnotesize  1.019 & \\footnotesize  0112260301  & \\footnotesize  13806 & \\footnotesize  9988 & \\footnotesize  ok & \\footnotesize  A399 \\\\ %& \\footnotesize  %Turner, Martin & \\footnotesize   2002-02-04 \\\\ \\footnotesize 02 59 51.4 & \\footnotesize  +13 46 47 & \\footnotesize  164.2 & \\footnotesize  -38.5 & \\footnotesize  1.019 & \\footnotesize  0112260401  & \\footnotesize  12755 & \\footnotesize  8971 & \\footnotesize  ok & \\footnotesize  A399 \\\\ %& \\footnotesize  %Turner, Martin & \\footnotesize   2002-02-04 \\\\ \\footnotesize 03 02 39.1 & \\footnotesize  +00 07 31 & \\footnotesize  177.5 & \\footnotesize  -48.3 & \\footnotesize  0.733 & \\footnotesize  0041170101  & \\footnotesize  51724 & \\footnotesize  40890 & \\footnotesize  ok & \\footnotesize  CFRS3h \\\\ %\\footnotesize  Gear, Walter & \\footnotesize   2001-02-17 \\\\ \\footnotesize 03 05 22.2 & \\footnotesize  +17 29 11 & \\footnotesize  162.7 & \\footnotesize  -34.8 & \\footnotesize  1.106 & \\footnotesize  0112190101  & \\footnotesize  13647 & \\footnotesize  9886 & \\footnotesize  ok & \\footnotesize  MS0302.5+1717 \\\\ %\\footnotesize  Mason, Keith & \\footnotesize   2002-08-23 \\\\ \\footnotesize 03 06 37.6 & \\footnotesize  +00 00 28 & \\footnotesize  178.6 & \\footnotesize  -47.6 & \\footnotesize  0.668 & \\footnotesize  0142610101  & \\footnotesize  73918 & \\footnotesize  32406 & \\footnotesize  ok & \\footnotesize  S2F1a \\\\ %\\footnotesize  Ceca, Roberto & \\footnotesize   2003-02-12 \\\\ \\footnotesize 03 07 04.4 & \\footnotesize  -28 40 29 & \\footnotesize  223.9 & \\footnotesize  -60.0 & \\footnotesize  0.136 & \\footnotesize  0042340501  & \\footnotesize  14767 & \\footnotesize  10108 & \\footnotesize  ok & \\footnotesize  RXCJ0307.0-2840 \\\\ %\\footnotesize  Boehringer, Hans & \\footnotesize   2001-02-16 \\\\ \\footnotesize 03 11 55.8 & \\footnotesize  -76 51 52 & \\footnotesize  293.4 & \\footnotesize  -37.5 & \\footnotesize  0.748 & \\footnotesize  0122520201  & \\footnotesize  46099 & \\footnotesize  23563 & \\footnotesize  ok & \\footnotesize  PKS0312-770 \\\\ %\\footnotesize  Jansen, Fred & \\footnotesize   2000-03-31 \\\\ \\footnotesize 03 17 56.7 & \\footnotesize  -44 14 10 & \\footnotesize  252.9 & \\footnotesize  -56.0 & \\footnotesize  0.253 & \\footnotesize  0105660101  & \\footnotesize  24470 & \\footnotesize  19260 & \\footnotesize  ok & \\footnotesize  A3112 \\\\ %\\footnotesize  Jansen, Fred & \\footnotesize   2000-12-24 \\\\ \\footnotesize 03 18 14.9 & \\footnotesize  -66 29 48 & \\footnotesize  283.3 & \\footnotesize  -44.6 & \\footnotesize  0.440 & \\footnotesize  0106860101  & \\footnotesize  42416 & \\footnotesize  15312 & \\footnotesize  ok & \\footnotesize  NGC 1313 \\\\ %\\footnotesize  Aschenbach, Bernd & \\footnotesize   2000-10-17 \\\\ \\footnotesize 03 19 59.4 & \\footnotesize  +11 14 10 & \\footnotesize  171.1 & \\footnotesize  -37.3 & \\footnotesize  1.671 & \\footnotesize  0110661101  & \\footnotesize  12796 & \\footnotesize  8587 & \\footnotesize  ok & \\footnotesize  MBM 16 \\\\ %\\footnotesize  Jansen, Fred & \\footnotesize   2000-08-20 \\\\ \\footnotesize 03 20 07.6 & \\footnotesize  -62 26 33 & \\footnotesize  278.6 & \\footnotesize  -47.0 & \\footnotesize  0.422 & \\footnotesize  0084960101  & \\footnotesize  27842 & \\footnotesize  6428 & \\footnotesize  ok & \\footnotesize  Galactic Halo \\\\ %\\footnotesize  Shelton, Robin & \\footnotesize   2002-05-03 \\\\ \\footnotesize 03 23 21.1 & \\footnotesize  -49 30 26 & \\footnotesize  260.9 & \\footnotesize  -53.3 & \\footnotesize  0.172 & \\footnotesize  0140190101  & \\footnotesize  29614 & \\footnotesize  25722 & \\footnotesize  ok & \\footnotesize  RX J0323-4931 \\\\ %\\footnotesize  Walter, Roland & \\footnotesize   2003-08-16 \\\\ \\footnotesize 03 23 50.5 & \\footnotesize  -37 17 11 & \\footnotesize  240.2 & \\footnotesize  -56.4 & \\footnotesize  0.167 & \\footnotesize  0091770101  & \\footnotesize  60362 & \\footnotesize  27252 & \\footnotesize  ok & \\footnotesize  NGC 1316 \\\\ %\\footnotesize  Iyomoto, Naoko & \\footnotesize   2002-02-04 \\\\ \\footnotesize 03 32 53.2 & \\footnotesize  -09 28 23 & \\footnotesize  195.8 & \\footnotesize  -48.0 & \\footnotesize  0.413 & \\footnotesize  0112880501  & \\footnotesize  13420 & \\footnotesize  11301 & \\footnotesize  ok & \\footnotesize  HR 1084 \\\\ %\\footnotesize  Aschenbach, Bernd & \\footnotesize   2003-01-19 \\\\ \\footnotesize 03 32 52.8 & \\footnotesize  -63 27 53 & \\footnotesize  278.6 & \\footnotesize  -45.3 & \\footnotesize  0.504 & \\footnotesize  0084960201  & \\footnotesize  12846 & \\footnotesize  10014 & \\footnotesize  ok & \\footnotesize  Galactic Halo \\\\ %\\footnotesize  Shelton, Robin & \\footnotesize   2002-05-03 \\\\ \\footnotesize 03 33 30.3 & \\footnotesize  -36 10 22 & \\footnotesize  238.0 & \\footnotesize  -54.6 & \\footnotesize  0.135 & \\footnotesize  0151370101  & \\footnotesize  19417 & \\footnotesize  15229 & \\footnotesize  ok & \\footnotesize  NGC 1365-X1 \\\\ %\\footnotesize  Fabbiano, Giuseppina & \\footnotesize   2003-01-17 \\\\ \\footnotesize 03 33 38.8 & \\footnotesize  -36 08 55 & \\footnotesize  237.9 & \\footnotesize  -54.5 & \\footnotesize  0.135 & \\footnotesize  0151370701  & \\footnotesize  11417 & \\footnotesize  6839 & \\footnotesize  ok & \\footnotesize  NGC 1365-X1 \\\\ %\\footnotesize  Fabbiano, Giuseppina & \\footnotesize   2003-08-13 \\\\ \\footnotesize 03 36 41.1 & \\footnotesize  +00 36 23 & \\footnotesize  184.8 & \\footnotesize  -41.5 & \\footnotesize  0.805 & \\footnotesize  0116150601  & \\footnotesize  64197 & \\footnotesize  22191 & \\footnotesize  ok & \\footnotesize  HR1099 \\\\ %\\footnotesize  Schartel, Norbert & \\footnotesize   2000-01-30 \\\\ \\footnotesize 03 36 41.3 & \\footnotesize  +00 36 24 & \\footnotesize  184.8 & \\footnotesize  -41.5 & \\footnotesize  0.805 & \\footnotesize  0116340601  & \\footnotesize  28610 & \\footnotesize  17959 & \\footnotesize  discard & \\footnotesize  HR1099 \\\\ %\\footnotesize  Schartel, Norbert & \\footnotesize   2000-02-03 \\\\ \\footnotesize 03 36 43.7 & \\footnotesize  -36 01 03 & \\footnotesize  237.6 & \\footnotesize  -53.9 & \\footnotesize  0.139 & \\footnotesize  0140950201  & \\footnotesize  17417 & \\footnotesize  14528 & \\footnotesize  ok & \\footnotesize  NGC1386 \\\\ %\\footnotesize  Guainazzi, Matteo & \\footnotesize   2002-12-30 \\\\ \\footnotesize 03 36 47.2 & \\footnotesize  +00 35 21 & \\footnotesize  184.9 & \\footnotesize  -41.5 & \\footnotesize  0.805 & \\footnotesize  0134540601  & \\footnotesize  35740 & \\footnotesize  28264 & \\footnotesize  ok & \\footnotesize  HR1099 \\\\ %\\footnotesize  Jansen, Fred & \\footnotesize   2002-08-23 \\\\ \\footnotesize 03 36 46.9 & \\footnotesize  +00 35 27 & \\footnotesize  184.9 & \\footnotesize  -41.5 & \\footnotesize  0.805 & \\footnotesize  0129350201  & \\footnotesize  31310 & \\footnotesize  8460 & \\footnotesize  flared & \\footnotesize  HR1099 \\\\ %\\footnotesize  Jansen, Fred & \\footnotesize   2000-08-28 \\\\ \\footnotesize 03 36 47.0 & \\footnotesize  +00 35 20 & \\footnotesize  184.9 & \\footnotesize  -41.5 & \\footnotesize  0.805 & \\footnotesize  0134540101  & \\footnotesize  46325 & \\footnotesize  13660 & \\footnotesize  discard & \\footnotesize  HR1099 \\\\ %\\footnotesize  Jansen, Fred & \\footnotesize   2001-02-23 \\\\ \\footnotesize 03 36 47.2 & \\footnotesize  +00 35 19 & \\footnotesize  184.9 & \\footnotesize  -41.5 & \\footnotesize  0.805 & \\footnotesize  0134540801  & \\footnotesize  51927 & \\footnotesize  27 & \\footnotesize  discard & \\footnotesize  HR1099 \\\\ %\\footnotesize  Jansen, Fred & \\footnotesize   2004-08-13 \\\\ \\footnotesize 03 37 44.3 & \\footnotesize  -25 22 17 & \\footnotesize  219.8 & \\footnotesize  -52.8 & \\footnotesize  0.155 & \\footnotesize  0107860401  & \\footnotesize  63138 & \\footnotesize  26483 & \\footnotesize  ok & \\footnotesize  SHARC4 \\\\ %\\footnotesize  Jansen, Fred & \\footnotesize   2001-08-19 \\\\ \\footnotesize 03 38 28.3 & \\footnotesize  -35 26 52 & \\footnotesize  236.7 & \\footnotesize  -53.6 & \\footnotesize  0.145 & \\footnotesize  0012830101  & \\footnotesize  29364 & \\footnotesize  6170 & \\footnotesize  ok & \\footnotesize  NGC 1399 \\\\ %\\footnotesize  Buote, David & \\footnotesize   2001-06-28 \\\\ \\footnotesize 03 38 29.8 & \\footnotesize  +00 21 44 & \\footnotesize  185.5 & \\footnotesize  -41.3 & \\footnotesize  0.735 & \\footnotesize  0036540101  & \\footnotesize  23004 & \\footnotesize 17982 & \\footnotesize ok & \\footnotesize SDSS033829+00 \\\\ %& \\footnotesize  Brandt, William & \\footnotesize   2002-02-22 \\\\ \\footnotesize 03 38 36.3 & \\footnotesize  +09 57 50 & \\footnotesize  176.2 & \\footnotesize  -35.0 & \\footnotesize  1.864 & \\footnotesize  0109870101  & \\footnotesize  30612 & \\footnotesize  9118 & \\footnotesize  discard & \\footnotesize  2A 0335+096 \\\\ %\\footnotesize  Kaastra, Jelle & \\footnotesize   2001-02-09 \\\\ \\footnotesize 03 38 40.7 & \\footnotesize  +09 58 17 & \\footnotesize  176.2 & \\footnotesize  -35.0 & \\footnotesize  1.864 & \\footnotesize  0147800201  & \\footnotesize  140323 & \\footnotesize  99140 & \\footnotesize discard & \\footnotesize 2A 0335+096 \\\\ %& \\footnotesize  Kaastra, Jelle & \\footnotesize   2003-08-05 \\\\ \\footnotesize 03 39 35.7 & \\footnotesize  -35 25 55 & \\footnotesize  236.6 & \\footnotesize  -53.4 & \\footnotesize  0.148 & \\footnotesize  0055140101  & \\footnotesize  51410 & \\footnotesize  39957 & \\footnotesize  ok & \\footnotesize  LP 944-20 \\\\ %\\footnotesize  Martin, Eduardo & \\footnotesize   2001-01-08 \\\\ \\footnotesize 03 41 30.2 & \\footnotesize  -44 52 34 & \\footnotesize  252.0 & \\footnotesize  -51.8 & \\footnotesize  0.166 & \\footnotesize  0045940301  & \\footnotesize  31893 & \\footnotesize  22907 & \\footnotesize  ok & \\footnotesize  AXJ0341.4-4453 \\\\ %\\footnotesize  Georgantopou, Ioannis & \\footnotesize   2002-08-26 \\\\ \\footnotesize 03 45 46.0 & \\footnotesize  -41 12 39 & \\footnotesize  246.0 & \\footnotesize  -51.7 & \\footnotesize  0.188 & \\footnotesize  0201900801  & \\footnotesize  26912 & \\footnotesize  12926 & \\footnotesize  ok & \\footnotesize  RXCJ0345.7-4112 \\\\ %\\footnotesize  Boehringer, Hans & \\footnotesize   2004-03-05 \\\\ \\footnotesize 03 53 43.6 & \\footnotesize  -00 04 33 & \\footnotesize  188.9 & \\footnotesize  -38.5 & \\footnotesize  1.183 & \\footnotesize  0134920901  & \\footnotesize  20467 & \\footnotesize  9150 & \\footnotesize  ok & \\footnotesize  SA95-42 \\\\ %\\footnotesize  Jansen, Fred & \\footnotesize   2001-02-06 \\\\ \\footnotesize 03 53 50.5 & \\footnotesize  -10 24 35 & \\footnotesize  200.6 & \\footnotesize  -44.0 & \\footnotesize  0.428 & \\footnotesize  0093160201  & \\footnotesize  28014 & \\footnotesize  13589 & \\footnotesize  ok & \\footnotesize  BR 0351-1034 \\\\ %\\footnotesize  Mathur, Smita & \\footnotesize   2002-08-23 \\\\ \\footnotesize 03 54 35.6 & \\footnotesize  -20 30 52 & \\footnotesize  214.1 & \\footnotesize  -47.8 & \\footnotesize  0.394 & \\footnotesize  0150910101  & \\footnotesize  18920 & \\footnotesize  7796 & \\footnotesize  ok & \\footnotesize  NGC 1482 \\\\ %\\footnotesize  Strickland, David & \\footnotesize   2003-02-14 \\\\ \\footnotesize 03 57 25.3 & \\footnotesize  +01 11 41 & \\footnotesize  188.3 & \\footnotesize  -37.0 & \\footnotesize  1.301 & \\footnotesize  0094790201  & \\footnotesize  21892 & \\footnotesize  19221 & \\footnotesize  ok & \\footnotesize  Hawaii 167 \\\\ %\\footnotesize  Watson, Michael & \\footnotesize   2002-08-22 \\\\ \\footnotesize 03 58 50.8 & \\footnotesize  +10 25 21 & \\footnotesize  179.8 & \\footnotesize  -31.0 & \\footnotesize  1.452 & \\footnotesize  0064600301  & \\footnotesize  11914 & \\footnotesize  6857 & \\footnotesize  ok & \\footnotesize  3C 98 \\\\ %\\footnotesize  Makishima, Kazuo & \\footnotesize   2003-02-05 \\\\ \\footnotesize 03 58 57.9 & \\footnotesize  +10 26 37 & \\footnotesize  179.8 & \\footnotesize  -31.0 & \\footnotesize  1.452 & \\footnotesize  0064600101  & \\footnotesize  28904 & \\footnotesize  9336 & \\footnotesize  ok & \\footnotesize  3C 98 \\\\ %\\footnotesize  Makishima, Kazuo & \\footnotesize   2002-09-07 \\\\ \\footnotesize 04 12 35.5 & \\footnotesize  +10 15 34 & \\footnotesize  182.4 & \\footnotesize  -28.5 & \\footnotesize  1.432 & \\footnotesize  0112231601  & \\footnotesize  38757 & \\footnotesize  17000 & \\footnotesize  ok & \\footnotesize  A478 off-set \\\\ %\\footnotesize  Turner, Martin & \\footnotesize   2002-03-08 \\\\ \\footnotesize 04 13 30.2 & \\footnotesize  +10 28 05 & \\footnotesize  182.4 & \\footnotesize  -28.2 & \\footnotesize  1.527 & \\footnotesize  0109880101  & \\footnotesize  126748 & \\footnotesize  53202 & \\footnotesize  ok & \\footnotesize  A 478 \\\\ %\\footnotesize  Kaastra, Jelle & \\footnotesize   2002-02-17 \\\\ \\footnotesize 04 19 34.3 & \\footnotesize  +02 23 54 & \\footnotesize  190.9 & \\footnotesize  -31.8 & \\footnotesize  1.159 & \\footnotesize  0152150101  & \\footnotesize  30912 & \\footnotesize  19525 & \\footnotesize  ok & \\footnotesize  NGC1550 \\\\ %\\footnotesize  Takahashi, Isao & \\footnotesize   2003-02-22 \\\\ \\footnotesize 04 19 33.2 & \\footnotesize  +14 32 50 & \\footnotesize  179.9 & \\footnotesize  -24.5 & \\footnotesize  1.470 & \\footnotesize  0141400301  & \\footnotesize  25218 & \\footnotesize  8165 & \\footnotesize  ok & \\footnotesize  LH0416+14 \\\\ %\\footnotesize  Stelzer, Beate & \\footnotesize   2003-02-18 \\\\ \\footnotesize 04 19 47.6 & \\footnotesize  +15 37 33 & \\footnotesize  179.0 & \\footnotesize  -23.8 & \\footnotesize  1.764 & \\footnotesize  0024140101  & \\footnotesize  63578 & \\footnotesize  45054 & \\footnotesize  ok & \\footnotesize  GammaTau \\\\ %\\footnotesize  Brickhouse, Nancy & \\footnotesize   2001-03-14 \\\\ \\footnotesize 04 20 01.1 & \\footnotesize  -50 23 54 & \\footnotesize  258.1 & \\footnotesize  -44.3 & \\footnotesize  0.196 & \\footnotesize  0141750101  & \\footnotesize  22201 & \\footnotesize  19269 & \\footnotesize  ok & \\footnotesize  RXJ0420.0-5022 \\\\ %\\footnotesize  Haberl, Frank & \\footnotesize   2002-12-30 \\\\ \\footnotesize 04 20 01.0 & \\footnotesize  -50 23 55 & \\footnotesize  258.1 & \\footnotesize  -44.3 & \\footnotesize  0.196 & \\footnotesize  0141751001  & \\footnotesize  21914 & \\footnotesize  16225 & \\footnotesize  ok & \\footnotesize  RXJ0420.0-5022 \\\\ %\\footnotesize  Haberl, Frank & \\footnotesize   2003-01-01 \\\\ \\footnotesize 04 20 00.7 & \\footnotesize  -50 23 56 & \\footnotesize  258.1 & \\footnotesize  -44.3 & \\footnotesize  0.196 & \\footnotesize  0141751101  & \\footnotesize  22420 & \\footnotesize  19204 & \\footnotesize  ok & \\footnotesize  RXJ0420.0-5022 \\\\ %\\footnotesize  Haberl, Frank & \\footnotesize   2003-01-19 \\\\ \\footnotesize 04 20 05.5 & \\footnotesize  -50 21 47 & \\footnotesize  258.1 & \\footnotesize  -44.3 & \\footnotesize  0.196 & \\footnotesize  0141751201  & \\footnotesize  21919 & \\footnotesize  19404 & \\footnotesize  ok & \\footnotesize  RXJ0420.0-5022 \\\\ %\\footnotesize  Haberl, Frank & \\footnotesize   2003-07-26 \\\\ \\footnotesize 04 22 12.6 & \\footnotesize  -38 45 55 & \\footnotesize  241.7 & \\footnotesize  -44.9 & \\footnotesize  0.190 & \\footnotesize  0104860101  & \\footnotesize  22865 & \\footnotesize  8417 & \\footnotesize  ok & \\footnotesize  [HB89] 0420-388 \\\\ %\\footnotesize  Mason, Keith & \\footnotesize   2003-01-15 \\\\ \\footnotesize 04 24 12.7 & \\footnotesize  +14 45 35 & \\footnotesize  180.5 & \\footnotesize  -23.5 & \\footnotesize  1.670 & \\footnotesize  0101440601  & \\footnotesize  48757 & \\footnotesize  33085 & \\footnotesize  ok & \\footnotesize  VB50 \\\\ %\\footnotesize  Pallavicini, Roberto & \\footnotesize   2000-09-09 \\\\ \\footnotesize 04 24 12.9 & \\footnotesize  +14 45 28 & \\footnotesize  180.5 & \\footnotesize  -23.5 & \\footnotesize  1.670 & \\footnotesize  0101441501  & \\footnotesize  57513 & \\footnotesize  32175 & \\footnotesize  ok & \\footnotesize  VB50 \\\\ %\\footnotesize  Pallavicini, Roberto & \\footnotesize   2002-03-06 \\\\ \\footnotesize 04 26 11.2 & \\footnotesize  +16 55 48 & \\footnotesize  179.0 & \\footnotesize  -21.8 & \\footnotesize  1.915 & \\footnotesize  0056020401  & \\footnotesize  23659 & \\footnotesize  8990 & \\footnotesize  ok & \\footnotesize  RX J0426.1+1655 \\\\ %\\footnotesize  ARNAUD, Monique & \\footnotesize   2002-09-06 \\\\ \\footnotesize 04 28 35.0 & \\footnotesize  +15 57 45 & \\footnotesize  180.2 & \\footnotesize  -21.9 & \\footnotesize  1.711 & \\footnotesize  0101440501  & \\footnotesize  47811 & \\footnotesize  37399 & \\footnotesize  ok & \\footnotesize  VB71 \\\\ %\\footnotesize  Pallavicini, Roberto & \\footnotesize   2000-09-03 \\\\ \\footnotesize 04 31 38.8 & \\footnotesize  +18 10 15 & \\footnotesize  178.9 & \\footnotesize  -20.0 & \\footnotesize  1.792 & \\footnotesize  0109060301  & \\footnotesize  57818 & \\footnotesize  48640 & \\footnotesize  ok & \\footnotesize  L1551 cloud \\\\ %\\footnotesize  Favata, Fabio & \\footnotesize   2000-09-10 \\\\ \\footnotesize 04 33 39.3 & \\footnotesize  -13 15 45 & \\footnotesize  209.5 & \\footnotesize  -36.4 & \\footnotesize  0.568 & \\footnotesize  0135120201  & \\footnotesize  30867 & \\footnotesize  17308 & \\footnotesize  discard & \\footnotesize  A 496 \\\\ %\\footnotesize  Bleeker, Johan & \\footnotesize   2001-02-01 \\\\ \\footnotesize 04 33 56.5 & \\footnotesize  -08 35 24 & \\footnotesize  204.4 & \\footnotesize  -34.3 & \\footnotesize  0.579 & \\footnotesize  0150480201  & \\footnotesize  23929 & \\footnotesize  17124 & \\footnotesize  ok & \\footnotesize  NGC1614 \\\\ %\\footnotesize  Maiolino, Roberto & \\footnotesize   2003-02-14 \\\\ \\footnotesize 04 37 14.1 & \\footnotesize  +00 44 13 & \\footnotesize  195.3 & \\footnotesize  -29.0 & \\footnotesize  0.819 & \\footnotesize  0042341601  & \\footnotesize  18238 & \\footnotesize  2955 & \\footnotesize  flared & \\footnotesize  RXCJ0437.1+0043 \\\\ %\\footnotesize  Boehringer, Hans & \\footnotesize   2002-09-07 \\\\ \\footnotesize 04 37 22.0 & \\footnotesize  -47 15 30 & \\footnotesize  253.4 & \\footnotesize  -41.9 & \\footnotesize  0.169 & \\footnotesize  0112320201  & \\footnotesize  69412 & \\footnotesize  24942 & \\footnotesize  ok & \\footnotesize  PSR J0437-4715 \\\\ %\\footnotesize  Aschenbach, Bernd & \\footnotesize   2002-10-09 \\\\ \\footnotesize 04 40 16.9 & \\footnotesize  -43 33 25 & \\footnotesize  248.4 & \\footnotesize  -41.5 & \\footnotesize  0.147 & \\footnotesize  0104860201  & \\footnotesize  12868 & \\footnotesize  10136 & \\footnotesize  ok & \\footnotesize  [HB89] 0438-436 \\\\ %\\footnotesize  Mason, Keith & \\footnotesize   2002-04-06 \\\\ \\footnotesize 04 47 52.1 & \\footnotesize  -06 27 17 & \\footnotesize  204.0 & \\footnotesize  -30.3 & \\footnotesize  0.547 & \\footnotesize  0014740601  & \\footnotesize  28887 & \\footnotesize  6505 & \\footnotesize  ok & \\footnotesize  a0447-0627 \\\\ %\\footnotesize  Ceca, Roberto & \\footnotesize   2002-09-08 \\\\ \\footnotesize 04 48 12.6 & \\footnotesize  -20 29 08 & \\footnotesize  219.5 & \\footnotesize  -35.9 & \\footnotesize  0.321 & \\footnotesize  0142240101  & \\footnotesize  15208 & \\footnotesize  7895 & \\footnotesize  ok & \\footnotesize  Abell 514 \\\\ %\\footnotesize  Schindler, Sabine & \\footnotesize   2003-02-07 \\\\ \\footnotesize 04 48 30.8 & \\footnotesize  -20 35 42 & \\footnotesize  219.6 & \\footnotesize  -35.8 & \\footnotesize  0.321 & \\footnotesize  0142240201  & \\footnotesize  15203 & \\footnotesize  7860 & \\footnotesize  ok & \\footnotesize  Abell 514 \\\\ %\\footnotesize  Schindler, Sabine & \\footnotesize   2003-03-16 \\\\ \\footnotesize 04 52 34.4 & \\footnotesize  -29 53 01 & \\footnotesize  231.1 & \\footnotesize  -37.5 & \\footnotesize  0.180 & \\footnotesize  0153100101  & \\footnotesize  15925 & \\footnotesize  9924 & \\footnotesize  ok & \\footnotesize  IRAS 04505-2958 \\\\ %\\footnotesize  Anabuki, Naohisa & \\footnotesize   2003-09-09 \\\\ \\footnotesize 04 54 10.2 & \\footnotesize  -53 22 40 & \\footnotesize  261.1 & \\footnotesize  -38.7 & \\footnotesize  0.350 & \\footnotesize  0148650101  & \\footnotesize  58916 & \\footnotesize  39557 & \\footnotesize  ok & \\footnotesize  NGC 1705 \\\\ %\\footnotesize  Heckman, Timothy & \\footnotesize   2003-01-31 \\\\ \\footnotesize 04 54 44.8 & \\footnotesize  -18 08 30 & \\footnotesize  217.4 & \\footnotesize  -33.6 & \\footnotesize  0.425 & \\footnotesize  0140210101  & \\footnotesize  38947 & \\footnotesize  29209 & \\footnotesize  ok & \\footnotesize  RXCJ0454-1808 \\\\ %\\footnotesize  Boehringer, Hans & \\footnotesize   2003-03-26 \\\\ \\footnotesize 04 59 34.8 & \\footnotesize  +01 47 17 & \\footnotesize  197.6 & \\footnotesize  -23.7 & \\footnotesize  0.747 & \\footnotesize  0112880401  & \\footnotesize  20600 & \\footnotesize  16321 & \\footnotesize  ok & \\footnotesize  Gl 182 \\\\ %\\footnotesize  Aschenbach, Bernd & \\footnotesize   2000-09-03 \\\\ \\footnotesize 05 00 55.6 & \\footnotesize  -38 40 12 & \\footnotesize  242.4 & \\footnotesize  -37.3 & \\footnotesize  0.307 & \\footnotesize  0083151201  & \\footnotesize  13516 & \\footnotesize  8438 & \\footnotesize  ok & \\footnotesize  A3301 \\\\ %\\footnotesize  David, Laurence & \\footnotesize   2002-10-07 \\\\ \\footnotesize 05 01 05.9 & \\footnotesize  -24 25 18 & \\footnotesize  225.2 & \\footnotesize  -34.3 & \\footnotesize  0.237 & \\footnotesize  0110870101  & \\footnotesize  25867 & \\footnotesize  21755 & \\footnotesize  ok & \\footnotesize  CL0500-24 \\\\ %\\footnotesize  Watson, Michael & \\footnotesize   2002-03-28 \\\\ \\footnotesize 05 05 19.4 & \\footnotesize  -28 48 53 & \\footnotesize  230.6 & \\footnotesize  -34.6 & \\footnotesize  0.154 & \\footnotesize  0111160201  & \\footnotesize  49616 & \\footnotesize  26291 & \\footnotesize  ok & \\footnotesize  SHARC-2 \\\\ %\\footnotesize  Watson, Michael & \\footnotesize   2001-09-02 \\\\ \\footnotesize 05 07 42.0 & \\footnotesize  -37 31 00 & \\footnotesize  241.2 & \\footnotesize  -35.9 & \\footnotesize  0.278 & \\footnotesize  0110980801  & \\footnotesize  43606 & \\footnotesize  33560 & \\footnotesize  ok & \\footnotesize  NGC 1808 \\\\ %\\footnotesize  Jansen, Fred & \\footnotesize   2002-04-07 \\\\ \\footnotesize 05 12 55.9 & \\footnotesize  -16 12 25 & \\footnotesize  217.2 & \\footnotesize  -28.9 & \\footnotesize  0.600 & \\footnotesize  0143370101  & \\footnotesize  47265 & \\footnotesize  37618 & \\footnotesize  ok & \\footnotesize  Mu Lep \\\\ %\\footnotesize  Behar, Ehud & \\footnotesize   2003-03-24 \\\\ \\footnotesize 05 14 06.2 & \\footnotesize  -40 02 43 & \\footnotesize  244.5 & \\footnotesize  -35.0 & \\footnotesize  0.361 & \\footnotesize  0151750101  & \\footnotesize  23914 & \\footnotesize  8177 & \\footnotesize  ok & \\footnotesize  4U 0513-40 \\\\ %\\footnotesize  Chakrabarty, Deepto & \\footnotesize   2003-04-01 \\\\ \\footnotesize 05 15 41.5 & \\footnotesize  +01 04 28 & \\footnotesize  200.4 & \\footnotesize  -20.6 & \\footnotesize  1.091 & \\footnotesize  0010620101  & \\footnotesize  31641 & \\footnotesize  20228 & \\footnotesize  ok & \\footnotesize  V1309 Ori \\\\ %\\footnotesize  Reinsch, Klaus & \\footnotesize   2001-03-18 \\\\ \\footnotesize 05 16 43.8 & \\footnotesize  -54 30 26 & \\footnotesize  262.2 & \\footnotesize  -35.3 & \\footnotesize  0.625 & \\footnotesize  0042340701  & \\footnotesize  15114 & \\footnotesize  6620 & \\footnotesize  ok & \\footnotesize  RXCJ0516.7-5430 \\\\ %\\footnotesize  Boehringer, Hans & \\footnotesize   2001-10-17 \\\\ \\footnotesize 05 19 49.8 & \\footnotesize  -45 46 59 & \\footnotesize  251.6 & \\footnotesize  -34.6 & \\footnotesize  0.449 & \\footnotesize  0090050701  & \\footnotesize  20521 & \\footnotesize  8647 & \\footnotesize  ok & \\footnotesize  PICTOR A \\\\ %\\footnotesize  Grandi, Paola & \\footnotesize   2001-03-17 \\\\ \\footnotesize 05 21 01.6 & \\footnotesize  -25 21 58 & \\footnotesize  227.8 & \\footnotesize  -30.2 & \\footnotesize  0.193 & \\footnotesize  0085640101  & \\footnotesize  12423 & \\footnotesize  7693 & \\footnotesize  ok & \\footnotesize  IRAS05189-2524 \\\\ %\\footnotesize  Heckman, Timothy & \\footnotesize   2001-03-17 \\\\ \\footnotesize 05 23 03.3 & \\footnotesize  -36 27 22 & \\footnotesize  240.6 & \\footnotesize  -32.6 & \\footnotesize  0.329 & \\footnotesize  0065760201  & \\footnotesize  31919 & \\footnotesize  28588 & \\footnotesize  ok & \\footnotesize  PKS0521-365 \\\\ %\\footnotesize  Chiappetti, Lucio & \\footnotesize   2002-10-10 \\\\ \\footnotesize 05 25 01.7 & \\footnotesize  -33 43 47 & \\footnotesize  237.5 & \\footnotesize  -31.7 & \\footnotesize  0.218 & \\footnotesize  0149500801  & \\footnotesize  13706 & \\footnotesize  6745 & \\footnotesize  ok & \\footnotesize  PMN-525-3343 \\\\ %\\footnotesize  Fabian, Andrew & \\footnotesize   2003-03-06 \\\\ \\footnotesize 05 25 01.7 & \\footnotesize  -33 43 46 & \\footnotesize  237.5 & \\footnotesize  -31.7 & \\footnotesize  0.218 & \\footnotesize  0149500901  & \\footnotesize  12205 & \\footnotesize  5727 & \\footnotesize  ok & \\footnotesize  PMN-525-3343 \\\\ %\\footnotesize  Fabian, Andrew & \\footnotesize   2003-03-16 \\\\ \\footnotesize 05 25 01.5 & \\footnotesize  -33 43 46 & \\footnotesize  237.5 & \\footnotesize  -31.7 & \\footnotesize  0.218 & \\footnotesize  0149500601  & \\footnotesize  12206 & \\footnotesize  8994 & \\footnotesize  ok & \\footnotesize  PMN-525-3343 \\\\ %\\footnotesize  Fabian, Andrew & \\footnotesize   2003-02-14 \\\\ \\footnotesize 05 25 01.7 & \\footnotesize  -33 43 45 & \\footnotesize  237.5 & \\footnotesize  -31.7 & \\footnotesize  0.218 & \\footnotesize  0149500701  & \\footnotesize  12206 & \\footnotesize  8534 & \\footnotesize  ok & \\footnotesize  PMN-525-3343 \\\\ %\\footnotesize  Fabian, Andrew & \\footnotesize   2003-02-24 \\\\ \\footnotesize 05 25 06.2 & \\footnotesize  -33 42 51 & \\footnotesize  237.5 & \\footnotesize  -31.6 & \\footnotesize  0.218 & \\footnotesize  0050150301  & \\footnotesize  28528 & \\footnotesize  17343 & \\footnotesize  ok & \\footnotesize  PMN 0525-3343 \\\\ %\\footnotesize  Fabian, Andrew & \\footnotesize   2001-09-15 \\\\ \\footnotesize 05 25 06.7 & \\footnotesize  -33 43 12 & \\footnotesize  237.5 & \\footnotesize  -31.6 & \\footnotesize  0.218 & \\footnotesize  0050150101  & \\footnotesize  21870 & \\footnotesize  8102 & \\footnotesize  ok & \\footnotesize  PMN 0525-3343 \\\\ %\\footnotesize  Fabian, Andrew & \\footnotesize   2001-02-11 \\\\ \\footnotesize 05 25 08.9 & \\footnotesize  -33 42 09 & \\footnotesize  237.5 & \\footnotesize  -31.6 & \\footnotesize  0.218 & \\footnotesize  0149501201  & \\footnotesize  12419 & \\footnotesize  5663 & \\footnotesize  ok & \\footnotesize  PMN-525-3343 \\\\ %\\footnotesize  Fabian, Andrew & \\footnotesize   2003-08-09 \\\\ \\footnotesize 05 28 56.0 & \\footnotesize  -39 27 36 & \\footnotesize  244.3 & \\footnotesize  -32.1 & \\footnotesize  0.210 & \\footnotesize  0042340801  & \\footnotesize  14112 & \\footnotesize  5543 & \\footnotesize  ok & \\footnotesize  RXCJ0528.9-3927 \\\\ %\\footnotesize  Boehringer, Hans & \\footnotesize   2001-09-15 \\\\ \\footnotesize 05 33 01.9 & \\footnotesize  -37 01 21 & \\footnotesize  241.7 & \\footnotesize  -30.8 & \\footnotesize  0.293 & \\footnotesize  0042341801  & \\footnotesize  12740 & \\footnotesize  8292 & \\footnotesize  ok & \\footnotesize  RXCJ0532.9-3701 \\\\ %\\footnotesize  Boehringer, Hans & \\footnotesize   2002-10-07 \\\\ \\footnotesize 05 33 44.0 & \\footnotesize  -24 10 49 & \\footnotesize  227.6 & \\footnotesize  -27.1 & \\footnotesize  0.233 & \\footnotesize  0068940101  & \\footnotesize  129506 & \\footnotesize  21632 & \\footnotesize  ok & \\footnotesize  CL1 \\\\ %\\footnotesize  ARNAUD, Monique & \\footnotesize   2001-10-02 \\\\ \\footnotesize 05 33 43.7 & \\footnotesize  -24 10 49 & \\footnotesize  227.6 & \\footnotesize  -27.1 & \\footnotesize  0.233 & \\footnotesize  0068940601  & \\footnotesize  35178 & \\footnotesize  11860 & \\footnotesize  ok & \\footnotesize  CL1 \\\\ %& \\footnotesize  %ARNAUD, Monique & \\footnotesize   2001-10-07 \\\\ \\footnotesize 05 40 12.0 & \\footnotesize  -40 54 59 & \\footnotesize  246.5 & \\footnotesize  -30.2 & \\footnotesize  0.353 & \\footnotesize  0149420101  & \\footnotesize  18217 & \\footnotesize  14790 & \\footnotesize  ok & \\footnotesize  ESO03060170 \\\\ %\\footnotesize  Forman, William & \\footnotesize   2002-10-11 \\\\ \\footnotesize 05 47 38.6 & \\footnotesize  -31 52 42 & \\footnotesize  236.9 & \\footnotesize  -26.6 & \\footnotesize  0.195 & \\footnotesize  0201900901  & \\footnotesize  25320 & \\footnotesize  20193 & \\footnotesize  ok & \\footnotesize  RXCJ0547.6-3152 \\\\ %\\footnotesize  Boehringer, Hans & \\footnotesize   2004-03-07 \\\\ \\footnotesize 05 50 40.8 & \\footnotesize  -32 16 09 & \\footnotesize  237.5 & \\footnotesize  -26.1 & \\footnotesize  0.220 & \\footnotesize  0111830201  & \\footnotesize  54576 & \\footnotesize  15715 & \\footnotesize  ok & \\footnotesize  PKS 0548-32 \\\\ %\\footnotesize  Kaastra, Jelle & \\footnotesize   2001-10-04 \\\\ \\footnotesize 05 59 46.1 & \\footnotesize  -50 26 59 & \\footnotesize  257.9 & \\footnotesize  -28.5 & \\footnotesize  0.484 & \\footnotesize  0125110101  & \\footnotesize  57265 & \\footnotesize  8000 & \\footnotesize  ok & \\footnotesize  PKS0558-504 \\\\ %\\footnotesize  Jansen, Fred & \\footnotesize   2000-05-24 \\\\ \\footnotesize 05 59 47.2 & \\footnotesize  -50 26 52 & \\footnotesize  257.9 & \\footnotesize  -28.5 & \\footnotesize  0.484 & \\footnotesize  0129360201  & \\footnotesize  26409 & \\footnotesize  9031 & \\footnotesize  ok & \\footnotesize  PKS0558-504 \\\\ %& \\footnotesize  Jansen, Fred & \\footnotesize   2000-10-10 \\\\ \\footnotesize 05 59 46.6 & \\footnotesize  -50 27 07 & \\footnotesize  257.9 & \\footnotesize  -28.5 & \\footnotesize  0.484 & \\footnotesize  0119100201  & \\footnotesize  45437 & \\footnotesize  4340 & \\footnotesize  flared & \\footnotesize  PKS0558 \\\\ %\\footnotesize  Schartel, Norbert & \\footnotesize   2000-03-02 \\\\ \\footnotesize 05 59 48.9 & \\footnotesize  -50 26 37 & \\footnotesize  257.9 & \\footnotesize  -28.5 & \\footnotesize  0.484 & \\footnotesize  0137550201  & \\footnotesize  14968 & \\footnotesize  11568 & \\footnotesize  ok & \\footnotesize  PKS0558-504 \\\\ %\\footnotesize  Jansen, Fred & \\footnotesize   2001-06-26 \\\\ \\footnotesize 05 59 48.4 & \\footnotesize  -50 27 10 & \\footnotesize  257.9 & \\footnotesize  -28.5 & \\footnotesize  0.484 & \\footnotesize  0137550601  & \\footnotesize  14785 & \\footnotesize  4640 & \\footnotesize  ok & \\footnotesize  PKS0558-504 \\\\ %\\footnotesize  Jansen, Fred & \\footnotesize   2001-10-19 \\\\ \\footnotesize 06 02 02.8 & \\footnotesize  -39 57 37 & \\footnotesize  246.4 & \\footnotesize  -26.0 & \\footnotesize  0.501 & \\footnotesize  0151900101  & \\footnotesize  47206 & \\footnotesize  19753 & \\footnotesize  ok & \\footnotesize  A3376 \\\\ %\\footnotesize  Markevitch, Maxim & \\footnotesize   2003-04-01 \\\\ \\footnotesize 06 15 24.3 & \\footnotesize  -33 24 56 & \\footnotesize  240.5 & \\footnotesize  -21.5 & \\footnotesize  0.325 & \\footnotesize  0092360301  & \\footnotesize  15513 & \\footnotesize  6363 & \\footnotesize  ok & \\footnotesize  3EG 0616-3310 \\\\ %\\footnotesize  Caraveo, Patrizia & \\footnotesize   2001-10-29 \\\\ \\footnotesize 06 15 24.1 & \\footnotesize  -32 55 00 & \\footnotesize  240.0 & \\footnotesize  -21.3 & \\footnotesize  0.410 & \\footnotesize  0092360401  & \\footnotesize  15514 & \\footnotesize  8972 & \\footnotesize  ok & \\footnotesize  3EG 0616-3310 \\\\ %\\footnotesize  Caraveo, Patrizia & \\footnotesize   2001-10-29 \\\\ \\footnotesize 06 16 47.1 & \\footnotesize  -47 48 12 & \\footnotesize  255.6 & \\footnotesize  -25.3 & \\footnotesize  0.481 & \\footnotesize  0201901101  & \\footnotesize  44676 & \\footnotesize  2156 & \\footnotesize  flared & \\footnotesize  RXCJ0616.8-4748 \\\\ %\\footnotesize  Boehringer, Hans & \\footnotesize   2004-04-27 \\\\ \\footnotesize 06 17 47.1 & \\footnotesize  -32 54 57 & \\footnotesize  240.1 & \\footnotesize  -20.9 & \\footnotesize  0.410 & \\footnotesize  0092360101  & \\footnotesize  12512 & \\footnotesize  8503 & \\footnotesize  ok & \\footnotesize  3EG 0616-3310 \\\\ %\\footnotesize  Caraveo, Patrizia & \\footnotesize   2001-10-29 \\\\ \\footnotesize 06 17 47.1 & \\footnotesize  -33 24 55 & \\footnotesize  240.6 & \\footnotesize  -21.0 & \\footnotesize  0.384 & \\footnotesize  0092360201  & \\footnotesize  15068 & \\footnotesize  10551 & \\footnotesize  ok & \\footnotesize  3EG 0616-3310 \\\\ %\\footnotesize  Caraveo, Patrizia & \\footnotesize   2001-10-19 \\\\ \\footnotesize 06 45 21.9 & \\footnotesize  -54 12 57 & \\footnotesize  263.6 & \\footnotesize  -22.5 & \\footnotesize  0.657 & \\footnotesize  0201901201  & \\footnotesize  26397 & \\footnotesize  10965 & \\footnotesize  ok & \\footnotesize  RXCJ0645.4-5413 \\\\ %\\footnotesize  Boehringer, Hans & \\footnotesize   2004-05-07 \\\\ \\footnotesize 06 45 25.6 & \\footnotesize  -54 12 11 & \\footnotesize  263.6 & \\footnotesize  -22.5 & \\footnotesize  0.657 & \\footnotesize  0201903401  & \\footnotesize  23236 & \\footnotesize  11474 & \\footnotesize  ok & \\footnotesize  RXCJ0645.4-5413 \\\\ %\\footnotesize  Boehringer, Hans & \\footnotesize   2004-06-12 \\\\ \\footnotesize 06 58 16.8 & \\footnotesize  -55 57 23 & \\footnotesize  266.0 & \\footnotesize  -21.2 & \\footnotesize  0.634 & \\footnotesize  0112980201  & \\footnotesize  46756 & \\footnotesize  21497 & \\footnotesize  ok & \\footnotesize  RXJ0658-55 \\\\ %\\footnotesize  Turner, Martin & \\footnotesize   2000-10-22 \\\\ \\footnotesize 07 51 08.8 & \\footnotesize  +18 07 50 & \\footnotesize  202.7 & \\footnotesize  21.0  & \\footnotesize  0.502 & \\footnotesize  0111100301  & \\footnotesize  38499 & \\footnotesize  9271 & \\footnotesize  ok & \\footnotesize  PSR J0751+18 \\\\ %\\footnotesize  Watson, Michael & \\footnotesize   2000-10-01 \\\\ \\footnotesize 08 06 27.5 & \\footnotesize  +15 27 49 & \\footnotesize  206.9 & \\footnotesize  23.4  & \\footnotesize  0.248 & \\footnotesize  0150800101  & \\footnotesize  28719 & \\footnotesize  17091 & \\footnotesize  ok & \\footnotesize  RXJ0806.3+1527 \\\\ %\\footnotesize  Israel, GianLuca & \\footnotesize   2002-11-02 \\\\ \\footnotesize 08 39 57.5 & \\footnotesize  +19 32 39 & \\footnotesize  206.0 & \\footnotesize  32.3  & \\footnotesize  0.309 & \\footnotesize  0101440401  & \\footnotesize  48305 & \\footnotesize  39011 & \\footnotesize  ok & \\footnotesize  Praesepe \\\\ %\\footnotesize  Pallavicini, Roberto & \\footnotesize   2000-11-08 \\\\ \\footnotesize 08 40 52.2 & \\footnotesize  +13 12 39 & \\footnotesize  212.9 & \\footnotesize  30.1  & \\footnotesize  0.422 & \\footnotesize  0147670301  & \\footnotesize  30914 & \\footnotesize  8263 & \\footnotesize  ok & \\footnotesize  3C 207 \\\\ %\\footnotesize  Fiore, Fabrizio & \\footnotesize   2003-10-16 \\\\ \\footnotesize 08 51 26.7 & \\footnotesize  +11 47 08 & \\footnotesize  215.7 & \\footnotesize  31.9  & \\footnotesize  0.377 & \\footnotesize  0109461001  & \\footnotesize  10160 & \\footnotesize  7267 & \\footnotesize  ok & \\footnotesize  EU Cnc \\\\ %\\footnotesize  Mason, Keith & \\footnotesize   2001-10-21 \\\\ \\footnotesize 09 08 13.8 & \\footnotesize  -09 36 35 & \\footnotesize  239.1 & \\footnotesize  24.6  & \\footnotesize  0.421 & \\footnotesize  0136740201  & \\footnotesize  18366 & \\footnotesize  5363 & \\footnotesize  ok & \\footnotesize  A754\\_f2 \\\\ %\\footnotesize  Turner, Martin & \\footnotesize   2001-05-14 \\\\ \\footnotesize 09 08 45.6 & \\footnotesize  -09 46 03 & \\footnotesize  239.3 & \\footnotesize  24.6  & \\footnotesize  0.421 & \\footnotesize  0112950401  & \\footnotesize  16810 & \\footnotesize  9772 & \\footnotesize  ok & \\footnotesize  A754\\_f4 \\\\ %\\footnotesize  Turner, Martin & \\footnotesize   2002-05-06 \\\\ \\footnotesize 09 08 55.4 & \\footnotesize  -09 32 04 & \\footnotesize  239.1 & \\footnotesize  24.8  & \\footnotesize  0.421 & \\footnotesize  0112950301  & \\footnotesize  14771 & \\footnotesize  9836 & \\footnotesize  ok & \\footnotesize  A754\\_f3 \\\\ %\\footnotesize  Turner, Martin & \\footnotesize   2002-05-07 \\\\ \\footnotesize 09 09 26.7 & \\footnotesize  -09 40 22 & \\footnotesize  239.3 & \\footnotesize  24.8  & \\footnotesize  0.459 & \\footnotesize  0136740101  & \\footnotesize  18370 & \\footnotesize  11654 & \\footnotesize  ok & \\footnotesize  A754\\_f1 \\\\ %\\footnotesize  Turner, Martin & \\footnotesize   2001-05-14 \\\\ \\footnotesize 09 11 27.6 & \\footnotesize  +05 51 05 & \\footnotesize  224.6 & \\footnotesize  33.6  & \\footnotesize  0.391 & \\footnotesize  0083240201  & \\footnotesize  20814 & \\footnotesize  11284 & \\footnotesize  discard & \\footnotesize  RX J0911.4+0551 \\\\ %& \\footnotesize  Hjorth, Jens & \\footnotesize   2001-11-02 \\\\ \\footnotesize 09 18 00.9 & \\footnotesize  -12 05 11 & \\footnotesize  242.9 & \\footnotesize  25.0  & \\footnotesize  0.486 & \\footnotesize  0109980301  & \\footnotesize  32995 & \\footnotesize  14662 & \\footnotesize  ok & \\footnotesize  Hydra A cluster \\\\ %\\footnotesize  Bleeker, Johan & \\footnotesize   2000-12-08 \\\\ \\footnotesize 09 43 44.0 & \\footnotesize  +16 44 34 & \\footnotesize  216.3 & \\footnotesize  45.5  & \\footnotesize  0.282 & \\footnotesize  0046940401  & \\footnotesize  15560 & \\footnotesize  11234 & \\footnotesize  ok & \\footnotesize  WARPJ0943.7+16 \\\\ %\\footnotesize  Ponman, Trevor & \\footnotesize   2001-11-02 \\\\ \\footnotesize 09 43 43.8 & \\footnotesize  +16 44 05 & \\footnotesize  216.3 & \\footnotesize  45.5  & \\footnotesize  0.282 & \\footnotesize  0046940201  & \\footnotesize  33116 & \\footnotesize  3149 & \\footnotesize  flared & \\footnotesize  WARPJ0943.7+16 \\\\ %& \\footnotesize  Ponman, Trevor & \\footnotesize   2001-05-07 \\\\ \\footnotesize 09 45 41.8 & \\footnotesize  -14 19 40 & \\footnotesize  249.7 & \\footnotesize  28.7  & \\footnotesize  0.517 & \\footnotesize  0147920301  & \\footnotesize  28917 & \\footnotesize  19551 & \\footnotesize  ok & \\footnotesize  NGC 2992 \\\\ %\\footnotesize  Schartel, Norbert & \\footnotesize   2003-05-19 \\\\ \\footnotesize 09 52 08.6 & \\footnotesize  -01 48 07 & \\footnotesize  239.4 & \\footnotesize  38.0  & \\footnotesize  0.384 & \\footnotesize  0065790101  & \\footnotesize  10111 & \\footnotesize  5800 & \\footnotesize  ok & \\footnotesize  RXJ 095208.7-014 \\\\ %\\footnotesize  Vikhlinin, Alexey & \\footnotesize   2001-11-12 \\\\ \\footnotesize 09 53 04.9 & \\footnotesize  +07 55 24 & \\footnotesize  228.8 & \\footnotesize  43.6  & \\footnotesize  0.301 & \\footnotesize  0103260801  & \\footnotesize  84291 & \\footnotesize  15600 & \\footnotesize  ok & \\footnotesize  PSR B0950+08 \\\\ %\\footnotesize  Jansen, Fred & \\footnotesize   2002-05-09 \\\\ \\footnotesize 09 53 14.3 & \\footnotesize  -15 58 43 & \\footnotesize  252.5 & \\footnotesize  28.9  & \\footnotesize  0.527 & \\footnotesize  0140210201  & \\footnotesize  38915 & \\footnotesize  33996 & \\footnotesize  ok & \\footnotesize  RXCJ0953-1558 \\\\ %\\footnotesize  Boehringer, Hans & \\footnotesize   2002-12-20 \\\\ \\footnotesize 09 53 36.5 & \\footnotesize  +01 34 33 & \\footnotesize  236.1 & \\footnotesize  40.3  & \\footnotesize  0.359 & \\footnotesize  0070940401  & \\footnotesize  26760 & \\footnotesize  7633 & \\footnotesize  ok & \\footnotesize  NGC 3044 \\\\ %\\footnotesize  Dettmar, Ralf-Juergen & \\footnotesize   2002-05-11 \\\\ \\footnotesize 09 53 41.0 & \\footnotesize  +01 35 03 & \\footnotesize  236.1 & \\footnotesize  40.3  & \\footnotesize  0.359 & \\footnotesize  0070940101  & \\footnotesize  11428 & \\footnotesize  3713 & \\footnotesize  flared & \\footnotesize  NGC 3044 \\\\ %\\footnotesize  Dettmar, Ralf-Juergen & \\footnotesize   2001-11-25 \\\\ \\footnotesize 09 54 48.8 & \\footnotesize  -20 23 03 & \\footnotesize  256.2 & \\footnotesize  26.0  & \\footnotesize  0.389 & \\footnotesize  0141170101  & \\footnotesize  61227 & \\footnotesize  30122 & \\footnotesize  ok & \\footnotesize  EIS0954-2023 \\\\ %\\footnotesize  ARNAUD, Monique & \\footnotesize   2003-05-09 \\\\ \\footnotesize 09 54 56.7 & \\footnotesize  +17 43 19 & \\footnotesize  216.4 & \\footnotesize  48.3  & \\footnotesize  0.339 & \\footnotesize  0112850101  & \\footnotesize  33376 & \\footnotesize  14111 & \\footnotesize  ok & \\footnotesize  0952+179 \\\\ %\\footnotesize  Griffiths, Richard & \\footnotesize   2001-05-11 \\\\ \\footnotesize 09 56 11.7 & \\footnotesize  -10 03 09 & \\footnotesize  248.0 & \\footnotesize  33.5  & \\footnotesize  0.507 & \\footnotesize  0148170101  & \\footnotesize  102815 & \\footnotesize  64021 & \\footnotesize  ok & \\footnotesize  Abell 901 \\\\ %\\footnotesize  Gray, Meghan & \\footnotesize   2003-05-07 \\\\ \\footnotesize 09 58 17.6 & \\footnotesize  -11 03 29 & \\footnotesize  249.3 & \\footnotesize  33.2  & \\footnotesize  0.506 & \\footnotesize  0201903501  & \\footnotesize  18285 & \\footnotesize  6509 & \\footnotesize  ok & \\footnotesize  RXCJ0958.3-1103 \\\\ %\\footnotesize  Boehringer, Hans & \\footnotesize   2004-06-17 \\\\ \\footnotesize 10 00 11.9 & \\footnotesize  -19 38 08 & \\footnotesize  256.7 & \\footnotesize  27.4  & \\footnotesize  0.449 & \\footnotesize  0041180301  & \\footnotesize  22358 & \\footnotesize  18131 & \\footnotesize  ok & \\footnotesize  NGC 3091 group \\\\ %\\footnotesize  Ponman, Trevor & \\footnotesize   2001-11-22 \\\\ \\footnotesize 10 04 15.7 & \\footnotesize  +05 12 53 & \\footnotesize  234.1 & \\footnotesize  44.6  & \\footnotesize  0.233 & \\footnotesize  0150610101  & \\footnotesize  23624 & \\footnotesize  8447 & \\footnotesize  ok & \\footnotesize  PG 1001+054 \\\\ %\\footnotesize  Schartel, Norbert & \\footnotesize   2003-05-04 \\\\ \\footnotesize 10 07 21.7 & \\footnotesize  +12 48 35 & \\footnotesize  225.1 & \\footnotesize  49.1  & \\footnotesize  0.369 & \\footnotesize  0140550601  & \\footnotesize  22205 & \\footnotesize  19384 & \\footnotesize  ok & \\footnotesize  PG 1004+130 \\\\ %\\footnotesize  Brandt, William & \\footnotesize   2003-05-05 \\\\ \\footnotesize 10 11 05.7 & \\footnotesize  -04 41 20 & \\footnotesize  246.1 & \\footnotesize  39.8  & \\footnotesize  0.391 & \\footnotesize  0026340201  & \\footnotesize  21618 & \\footnotesize  6073 & \\footnotesize  ok & \\footnotesize  Sextans A \\\\ %\\footnotesize  Fabian, Walter & \\footnotesize   2001-11-27 \\\\ \\footnotesize 10 19 59.8 & \\footnotesize  +08 13 28 & \\footnotesize  233.5 & \\footnotesize  49.5  & \\footnotesize  0.329 & \\footnotesize  0093640301  & \\footnotesize  24263 & \\footnotesize  3845 & \\footnotesize  flared & \\footnotesize  IRAS 10173+0828 \\\\ %\\footnotesize  Bauer, Franz & \\footnotesize   2001-05-26 \\\\ \\footnotesize 10 21 37.4 & \\footnotesize  +13 06 40 & \\footnotesize  227.2 & \\footnotesize  52.3  & \\footnotesize  0.379 & \\footnotesize  0146990101  & \\footnotesize  21917 & \\footnotesize  16357 & \\footnotesize  ok & \\footnotesize  IRAS 10190+1322 \\\\ %\\footnotesize  Risaliti, Guido & \\footnotesize   2003-05-05 \\\\ \\footnotesize 10 23 30.8 & \\footnotesize  +19 51 56 & \\footnotesize  216.9 & \\footnotesize  55.4  & \\footnotesize  0.222 & \\footnotesize  0101040301  & \\footnotesize  40495 & \\footnotesize  32296 & \\footnotesize  ok & \\footnotesize  NGC 3227 \\\\ %\\footnotesize  Jansen, Fred & \\footnotesize   2000-11-29 \\\\ \\footnotesize 10 23 40.0 & \\footnotesize  +04 11 27 & \\footnotesize  239.3 & \\footnotesize  47.9  & \\footnotesize  0.287 & \\footnotesize  0108670101  & \\footnotesize  56459 & \\footnotesize  48492 & \\footnotesize  ok & \\footnotesize  ZW 3146 \\\\ %\\footnotesize  Mushotzky, Richard & \\footnotesize   2000-12-06 \\\\ \\footnotesize 10 28 34.2 & \\footnotesize  -08 44 41 & \\footnotesize  253.9 & \\footnotesize  40.1  & \\footnotesize  0.459 & \\footnotesize  0153290101  & \\footnotesize  43368 & \\footnotesize  15915 & \\footnotesize  ok & \\footnotesize  RXJ1028.6-0844 \\\\ %\\footnotesize  Yuan, Weimin & \\footnotesize   2003-06-13 \\\\ \\footnotesize 10 30 22.4 & \\footnotesize  +05 24 45 & \\footnotesize  239.4 & \\footnotesize  50.0  & \\footnotesize  0.339 & \\footnotesize  0148560501  & \\footnotesize  103876 & \\footnotesize  60819 & \\footnotesize  ok & \\footnotesize  SDSS 1030+05 \\\\ %\\footnotesize  Farrah, Duncan & \\footnotesize   2003-05-23 \\\\ \\footnotesize 10 36 28.2 & \\footnotesize  -03 43 17 & \\footnotesize  251.1 & \\footnotesize  45.1  & \\footnotesize  0.471 & \\footnotesize  0150870401  & \\footnotesize  31418 & \\footnotesize  27773 & \\footnotesize  ok & \\footnotesize  BRI1033-0327 \\\\ %\\footnotesize  Bechtold, Jill & \\footnotesize   2002-12-20 \\\\ \\footnotesize 10 41 12.8 & \\footnotesize  +06 10 10 & \\footnotesize  241.0 & \\footnotesize  52.6  & \\footnotesize  0.270 & \\footnotesize  0151390101  & \\footnotesize  59915 & \\footnotesize  40999 & \\footnotesize  ok & \\footnotesize  4C 06.41 \\\\ %\\footnotesize  Weaver, Kimberly & \\footnotesize   2003-05-19 \\\\ \\footnotesize 10 44 32.9 & \\footnotesize  -01 25 13 & \\footnotesize  250.8 & \\footnotesize  48.2  & \\footnotesize  0.432 & \\footnotesize  0125300101  & \\footnotesize  62310 & \\footnotesize  29778 & \\footnotesize ok & \\footnotesize J104433.04-0125\\\\ %& \\footnotesize  Jansen, Fred & \\footnotesize   2000-05-28 \\\\ \\footnotesize 10 50 26.2 & \\footnotesize  -12 50 39 & \\footnotesize  262.7 & \\footnotesize  40.4  & \\footnotesize  0.501 & \\footnotesize  0146510301  & \\footnotesize  40921 & \\footnotesize  20572 & \\footnotesize  ok & \\footnotesize  NGC3411 Group \\\\ %\\footnotesize  Mulchaey, John & \\footnotesize   2002-12-21 \\\\ \\footnotesize 10 56 59.8 & \\footnotesize  -03 37 38 & \\footnotesize  256.5 & \\footnotesize  48.6  & \\footnotesize  0.379 & \\footnotesize  0094800101  & \\footnotesize  41021 & \\footnotesize  22223 & \\footnotesize  ok & \\footnotesize  MS1054.4-0321 \\\\ %\\footnotesize  Watson, Michael & \\footnotesize   2001-06-21 \\\\ \\footnotesize 11 01 52.6 & \\footnotesize  -34 42 37 & \\footnotesize  278.6 & \\footnotesize  22.9  & \\footnotesize  0.473 & \\footnotesize  0112880201  & \\footnotesize  29776 & \\footnotesize  25669 & \\footnotesize  ok & \\footnotesize  TW Hya \\\\ %\\footnotesize  Aschenbach, Bernd & \\footnotesize   2001-07-09 \\\\ \\footnotesize 11 06 34.2 & \\footnotesize  -18 21 44 & \\footnotesize  270.8 & \\footnotesize  37.8  & \\footnotesize  0.461 & \\footnotesize  0112630101  & \\footnotesize  36428 & \\footnotesize  10464 & \\footnotesize  ok & \\footnotesize  HE 1104-1805 \\\\ %\\footnotesize  Turner, Martin & \\footnotesize   2001-06-14 \\\\ \\footnotesize 11 13 17.2 & \\footnotesize  -26 45 06 & \\footnotesize  277.2 & \\footnotesize  31.1  & \\footnotesize  0.561 & \\footnotesize  0071340201  & \\footnotesize  22234 & \\footnotesize  10958 & \\footnotesize  ok & \\footnotesize  NGC 3585 \\\\ %\\footnotesize  Ponman, Trevor & \\footnotesize   2001-12-28 \\\\ \\footnotesize 11 15 18.1 & \\footnotesize  -21 56 38 & \\footnotesize  275.2 & \\footnotesize  35.7  & \\footnotesize  0.408 & \\footnotesize  0094380101  & \\footnotesize  33673 & \\footnotesize  26861 & \\footnotesize  ok & \\footnotesize  GRB011211 \\\\ %\\footnotesize  Perlman, Eric & \\footnotesize   2001-12-12 \\\\ \\footnotesize 11 17 15.8 & \\footnotesize  +17 57 42 & \\footnotesize  230.8 & \\footnotesize  66.4  & \\footnotesize  0.143 & \\footnotesize  0099030101  & \\footnotesize  23795 & \\footnotesize  13319 & \\footnotesize  ok & \\footnotesize  DP Leo \\\\ %\\footnotesize  Mason, Keith & \\footnotesize   2000-11-22 \\\\ \\footnotesize 11 18 17.3 & \\footnotesize  +07 46 11 & \\footnotesize  249.8 & \\footnotesize  60.6  & \\footnotesize  0.336 & \\footnotesize  0082340101  & \\footnotesize  63358 & \\footnotesize  49529 & \\footnotesize  ok & \\footnotesize  PG1115+080 \\\\ %\\footnotesize  CHARTAS, GEORGE & \\footnotesize   2001-11-26 \\\\ \\footnotesize 11 18 55.1 & \\footnotesize  +13 05 19 & \\footnotesize  241.3 & \\footnotesize  64.2  & \\footnotesize  0.227 & \\footnotesize  0082140301  & \\footnotesize  33519 & \\footnotesize  28398 & \\footnotesize  ok & \\footnotesize  NGC3623 \\\\ %\\footnotesize  Irwin, Jimmy & \\footnotesize   2002-05-22 \\\\ \\footnotesize 11 20 14.9 & \\footnotesize  +12 59 24 & \\footnotesize  241.9 & \\footnotesize  64.4  & \\footnotesize  0.264 & \\footnotesize  0093641101  & \\footnotesize  11363 & \\footnotesize  6979 & \\footnotesize  ok & \\footnotesize  NGC 3627 \\\\ %\\footnotesize  Bauer, Franz & \\footnotesize   2001-05-26 \\\\ \\footnotesize 11 20 17.1 & \\footnotesize  +13 35 31 & \\footnotesize  240.8 & \\footnotesize  64.7  & \\footnotesize  0.196 & \\footnotesize  0110980101  & \\footnotesize  86739 & \\footnotesize  41222 & \\footnotesize  ok & \\footnotesize  NGC 3628 \\\\ %\\footnotesize  Jansen, Fred & \\footnotesize   2000-11-28 \\\\ \\footnotesize 11 23 09.2 & \\footnotesize  +05 30 29 & \\footnotesize  254.7 & \\footnotesize  59.8  & \\footnotesize  0.433 & \\footnotesize  0083000301  & \\footnotesize  33361 & \\footnotesize  24867 & \\footnotesize  ok & \\footnotesize  3C 257 \\\\ %\\footnotesize  O'Brien, Paul & \\footnotesize   2001-12-16 \\\\ \\footnotesize 11 23 16.1 & \\footnotesize  +01 37 42 & \\footnotesize  259.5 & \\footnotesize  56.8  & \\footnotesize  0.423 & \\footnotesize  0145750101  & \\footnotesize  42208 & \\footnotesize  21840 & \\footnotesize  ok & \\footnotesize  Q1120+0195 \\\\ %\\footnotesize  CHARTAS, GEORGE & \\footnotesize   2003-06-24 \\\\ \\footnotesize 11 30 02.3 & \\footnotesize  -14 49 25 & \\footnotesize  275.2 & \\footnotesize  43.6  & \\footnotesize  0.406 & \\footnotesize  0112850201  & \\footnotesize  30089 & \\footnotesize  13489 & \\footnotesize  ok & \\footnotesize  1127-145 \\\\ %\\footnotesize  Griffiths, Richard & \\footnotesize   2002-07-01 \\\\ \\footnotesize 11 31 56.5 & \\footnotesize  -19 55 42 & \\footnotesize  278.5 & \\footnotesize  39.1  & \\footnotesize  0.451 & \\footnotesize  0042341001  & \\footnotesize  15025 & \\footnotesize  10122 & \\footnotesize  ok & \\footnotesize  RXCJ1131.9-1955 \\\\ %\\footnotesize  Boehringer, Hans & \\footnotesize   2001-07-01 \\\\ \\footnotesize 11 32 00.5 & \\footnotesize  -34 36 30 & \\footnotesize  284.8 & \\footnotesize  25.4  & \\footnotesize  0.504 & \\footnotesize  0112880101  & \\footnotesize  29921 & \\footnotesize  26857 & \\footnotesize  ok & \\footnotesize  CoD-33\\_7795 \\\\ %\\footnotesize  Aschenbach, Bernd & \\footnotesize   2003-01-09 \\\\ \\footnotesize 11 38 22.5 & \\footnotesize  +03 22 05 & \\footnotesize  263.4 & \\footnotesize  60.5  & \\footnotesize  0.235 & \\footnotesize  0111970701  & \\footnotesize  12866 & \\footnotesize  10213 & \\footnotesize  ok & \\footnotesize  T Leo \\\\ %\\footnotesize  Mason, Keith & \\footnotesize   2002-06-01 \\\\ \\footnotesize 11 39 01.7 & \\footnotesize  -37 44 08 & \\footnotesize  287.4 & \\footnotesize  22.9  & \\footnotesize  0.941 & \\footnotesize  0112210201  & \\footnotesize  137818 & \\footnotesize  16000 & \\footnotesize  ok & \\footnotesize  NGC3783 \\\\ %\\footnotesize  Kaastra, Jelle & \\footnotesize   2001-12-19 \\\\ \\footnotesize 11 39 01.7 & \\footnotesize  -37 44 08 & \\footnotesize  287.4 & \\footnotesize  22.9  & \\footnotesize  0.941 & \\footnotesize  0112210501  & \\footnotesize  137815 & \\footnotesize  20000 & \\footnotesize  ok & \\footnotesize  NGC3783 \\\\ %\\footnotesize  Kaastra, Jelle & \\footnotesize   2001-12-21 \\\\ \\footnotesize 11 39 02.3 & \\footnotesize  -37 44 13 & \\footnotesize  287.4 & \\footnotesize  22.9  & \\footnotesize  0.941 & \\footnotesize  0112210101  & \\footnotesize  40412 & \\footnotesize  13509 & \\footnotesize  ok & \\footnotesize  NCG3783 \\\\ %\\footnotesize  Kaastra, Jelle & \\footnotesize   2000-12-29 \\\\ \\footnotesize 11 39 06.9 & \\footnotesize  +17 07 15 & \\footnotesize  240.1 & \\footnotesize  70.5  & \\footnotesize  0.203 & \\footnotesize  0066950201  & \\footnotesize  13383 & \\footnotesize  4620 & \\footnotesize  ok & \\footnotesize  UGC 6614 \\\\ %\\footnotesize  Mao, Shude & \\footnotesize   2002-06-13 \\\\ \\footnotesize 11 41 07.0 & \\footnotesize  -01 43 09 & \\footnotesize  269.7 & \\footnotesize  56.5  & \\footnotesize  0.245 & \\footnotesize  0151230101  & \\footnotesize  10919 & \\footnotesize  6607 & \\footnotesize  ok & \\footnotesize  UN J1141-0143 \\\\ %\\footnotesize  Hall, Patrick & \\footnotesize   2003-07-06 \\\\ \\footnotesize 11 41 19.7 & \\footnotesize  -12 16 20 & \\footnotesize  277.3 & \\footnotesize  47.0  & \\footnotesize  0.330 & \\footnotesize  0201901601  & \\footnotesize  37932 & \\footnotesize  24496 & \\footnotesize  ok & \\footnotesize  RXCJ1141.4-1216 \\\\ %\\footnotesize  Boehringer, Hans & \\footnotesize   2004-07-09 \\\\ \\footnotesize 11 44 39.9 & \\footnotesize  +19 47 00 & \\footnotesize  235.0 & \\footnotesize  73.0  & \\footnotesize  0.255 & \\footnotesize  0061740101  & \\footnotesize  33366 & \\footnotesize  26714 & \\footnotesize  discard & \\footnotesize  A1367 \\\\ %\\footnotesize  Forman, William & \\footnotesize   2001-05-27 \\\\ \\footnotesize 11 50 41.5 & \\footnotesize  +01 45 50 & \\footnotesize  270.5 & \\footnotesize  60.7  & \\footnotesize  0.215 & \\footnotesize  0044740201  & \\footnotesize  49072 & \\footnotesize  43910 & \\footnotesize  ok & \\footnotesize  Beta Vir \\\\ %\\footnotesize  Schmitt, Juergen & \\footnotesize   2001-06-06 \\\\ \\footnotesize 11 51 02.4 & \\footnotesize  -28 48 06 & \\footnotesize  287.2 & \\footnotesize  32.2  & \\footnotesize  0.627 & \\footnotesize  0027340101  & \\footnotesize  45533 & \\footnotesize  29905 & \\footnotesize  ok & \\footnotesize  NGC 3923 \\\\ %\\footnotesize  Buote, David & \\footnotesize   2002-01-04 \\\\ \\footnotesize 12 00 48.0 & \\footnotesize  -03 27 39 & \\footnotesize  279.2 & \\footnotesize  57.0  & \\footnotesize  0.243 & \\footnotesize  0056020701  & \\footnotesize  30114 & \\footnotesize  24816 & \\footnotesize  ok & \\footnotesize  RX J1200.8-0328 \\\\ %\\footnotesize  ARNAUD, Monique & \\footnotesize   2001-12-12 \\\\ \\footnotesize 12 01 53.1 & \\footnotesize  -18 51 42 & \\footnotesize  286.9 & \\footnotesize  42.4  & \\footnotesize  0.400 & \\footnotesize  0085220101  & \\footnotesize  24913 & \\footnotesize  15885 & \\footnotesize  ok & \\footnotesize  The Antennae \\\\ %\\footnotesize  Fabbiano, Giuseppina & \\footnotesize   2002-01-19 \\\\ \\footnotesize 12 01 52.6 & \\footnotesize  -18 51 41 & \\footnotesize  286.9 & \\footnotesize  42.4  & \\footnotesize  0.400 & \\footnotesize  0085220201  & \\footnotesize  52777 & \\footnotesize  34943 & \\footnotesize  ok & \\footnotesize  The Antennae \\\\ %\\footnotesize  Fabbiano, Giuseppina & \\footnotesize   2002-01-09 \\\\ \\footnotesize 12 04 27.4 & \\footnotesize  +01 54 30 & \\footnotesize  276.9 & \\footnotesize  62.3  & \\footnotesize  0.186 & \\footnotesize  0093060101  & \\footnotesize  16004 & \\footnotesize  11301 & \\footnotesize  ok & \\footnotesize  MKW4 \\\\ %\\footnotesize  Vrtilek, Jan & \\footnotesize   2001-12-21 \\\\ \\footnotesize 12 12 15.8 & \\footnotesize  +13 12 33 & \\footnotesize  267.6 & \\footnotesize  73.3  & \\footnotesize  0.249 & \\footnotesize  0112550501  & \\footnotesize  23614 & \\footnotesize  17974 & \\footnotesize  ok & \\footnotesize  ngc4168 \\\\ %\\footnotesize  Turner, Martin & \\footnotesize   2001-12-04 \\\\ \\footnotesize 12 13 46.0 & \\footnotesize  +02 48 53 & \\footnotesize  280.9 & \\footnotesize  64.0  & \\footnotesize  0.174 & \\footnotesize  0081340801  & \\footnotesize  23206 & \\footnotesize  19206 & \\footnotesize  ok & \\footnotesize  IRAS12112+03 \\\\ %\\footnotesize  Franceschini, Alberto & \\footnotesize   2001-12-30 \\\\ \\footnotesize 12 16 40.5 & \\footnotesize  -12 01 17 & \\footnotesize  289.6 & \\footnotesize  49.9  & \\footnotesize  0.385 & \\footnotesize  0143210801  & \\footnotesize  33916 & \\footnotesize  22276 & \\footnotesize  ok & \\footnotesize  cl1216-12 \\\\ %\\footnotesize  Zaritsky, Dennis & \\footnotesize   2003-07-06 \\\\ \\footnotesize 12 18 45.0 & \\footnotesize  +14 24 53 & \\footnotesize  270.3 & \\footnotesize  75.1  & \\footnotesize  0.271 & \\footnotesize  0147610101  & \\footnotesize  50920 & \\footnotesize  15073 & \\footnotesize  ok & \\footnotesize  NGC 4254 \\\\ %\\footnotesize  Ehle, Matthias & \\footnotesize   2003-06-29 \\\\ \\footnotesize 12 19 23.2 & \\footnotesize  +05 49 47 & \\footnotesize  281.8 & \\footnotesize  67.3  & \\footnotesize  0.163 & \\footnotesize  0056340101  & \\footnotesize  33360 & \\footnotesize  23324 & \\footnotesize  ok & \\footnotesize  NGC 4261 \\\\ %\\footnotesize  Sambruna, Rita & \\footnotesize   2001-12-16 \\\\ \\footnotesize 12 20 13.5 & \\footnotesize  +06 41 15 & \\footnotesize  281.5 & \\footnotesize  68.2  & \\footnotesize  0.158 & \\footnotesize  0105070101  & \\footnotesize  11482 & \\footnotesize  5419 & \\footnotesize  ok & \\footnotesize  1WGA J1220.3+06 \\\\ %\\footnotesize  Mason, Keith & \\footnotesize   2002-07-05 \\\\ \\footnotesize 12 22 54.9 & \\footnotesize  +15 49 19 & \\footnotesize  271.1 & \\footnotesize  76.8  & \\footnotesize  0.232 & \\footnotesize  0106860201  & \\footnotesize  36633 & \\footnotesize  18017 & \\footnotesize  ok & \\footnotesize  NGC 4321 \\\\ %\\footnotesize  Aschenbach, Bernd & \\footnotesize   2001-12-29 \\\\ \\footnotesize 12 23 06.8 & \\footnotesize  +10 37 22 & \\footnotesize  279.5 & \\footnotesize  72.1  & \\footnotesize  0.217 & \\footnotesize  0108860101  & \\footnotesize  22271 & \\footnotesize  17203 & \\footnotesize  ok & \\footnotesize  NGC 4325 \\\\ %\\footnotesize  Mushotzky, Richard & \\footnotesize   2000-12-24 \\\\ \\footnotesize 12 25 42.4 & \\footnotesize  +12 39 33 & \\footnotesize  279.0 & \\footnotesize  74.3  & \\footnotesize  0.263 & \\footnotesize  0110930301  & \\footnotesize  18760 & \\footnotesize  433 & \\footnotesize  flared & \\footnotesize  NGC 4388 \\\\ %\\footnotesize  Watson, Michael & \\footnotesize   2002-07-07 \\\\ \\footnotesize 12 25 51.5 & \\footnotesize  +12 39 43 & \\footnotesize  279.1 & \\footnotesize  74.3  & \\footnotesize  0.263 & \\footnotesize  0110930701  & \\footnotesize  11915 & \\footnotesize  8140 & \\footnotesize  ok & \\footnotesize  NGC 4388 \\\\ %\\footnotesize  Watson, Michael & \\footnotesize   2002-12-12 \\\\ \\footnotesize 12 26 07.3 & \\footnotesize  +12 56 38 & \\footnotesize  279.0 & \\footnotesize  74.6  & \\footnotesize  0.264 & \\footnotesize  0108260201  & \\footnotesize  85636 & \\footnotesize  55860 & \\footnotesize  discard & \\footnotesize  M86 \\\\ %\\footnotesize  Aschenbach, Bernd & \\footnotesize   2002-07-02 \\\\ \\footnotesize 12 27 18.9 & \\footnotesize  +01 29 11 & \\footnotesize  289.2 & \\footnotesize  63.7  & \\footnotesize  0.181 & \\footnotesize  0110990201  & \\footnotesize  29679 & \\footnotesize  9818 & \\footnotesize  ok & \\footnotesize  HI1225+01 \\\\ %\\footnotesize  Watson, Michael & \\footnotesize   2001-06-23 \\\\ \\footnotesize 12 28 13.7 & \\footnotesize  -15 46 59 & \\footnotesize  294.7 & \\footnotesize  46.7  & \\footnotesize  0.378 & \\footnotesize  0002740301  & \\footnotesize  12234 & \\footnotesize  6450 & \\footnotesize  ok & \\footnotesize  Denis-J1228 \\\\ %\\footnotesize  Neuhaeuser, Ralph & \\footnotesize   2001-12-29 \\\\ \\footnotesize 12 28 44.2 & \\footnotesize  +02 07 20 & \\footnotesize  289.7 & \\footnotesize  64.4  & \\footnotesize  0.179 & \\footnotesize  0126700401  & \\footnotesize  27251 & \\footnotesize  2696 & \\footnotesize  flared & \\footnotesize  3C 273off+7min \\\\ %\\footnotesize  Jansen, Fred & \\footnotesize   2000-06-15 \\\\ \\footnotesize 12 29 03.3 & \\footnotesize  +02 01 42 & \\footnotesize  289.9 & \\footnotesize  64.3  & \\footnotesize  0.179 & \\footnotesize  0126700201  & \\footnotesize  26140 & \\footnotesize  8678 & \\footnotesize  ok & \\footnotesize  3C 273off-1.5min \\\\ %\\footnotesize  Jansen, Fred & \\footnotesize   2000-06-13 \\\\ \\footnotesize 12 29 30.2 & \\footnotesize  +01 59 32 & \\footnotesize  290.2 & \\footnotesize  64.3  & \\footnotesize  0.179 & \\footnotesize  0126700101  & \\footnotesize  27248 & \\footnotesize  135 & \\footnotesize  flared & \\footnotesize  3C 273off-7min \\\\ %\\footnotesize  Jansen, Fred & \\footnotesize   2000-06-13 \\\\ \\footnotesize 12 29 41.4 & \\footnotesize  +07 59 45 & \\footnotesize  286.8 & \\footnotesize  70.1  & \\footnotesize  0.159 & \\footnotesize  0112550601  & \\footnotesize  24622 & \\footnotesize  11715 & \\footnotesize  discard & \\footnotesize  ngc4472 \\\\ %\\footnotesize  Turner, Martin & \\footnotesize   2002-06-05 \\\\ \\footnotesize 12 29 57.5 & \\footnotesize  +13 38 11 & \\footnotesize  281.4 & \\footnotesize  75.6  & \\footnotesize  0.280 & \\footnotesize  0112552101  & \\footnotesize  14126 & \\footnotesize  10301 & \\footnotesize  ok & \\footnotesize  ngc4477 \\\\ %\\footnotesize  Turner, Martin & \\footnotesize   2002-06-08 \\\\ \\footnotesize 12 30 06.3 & \\footnotesize  +13 38 53 & \\footnotesize  281.5 & \\footnotesize  75.6  & \\footnotesize  0.258 & \\footnotesize  0112550701  & \\footnotesize  13488 & \\footnotesize  4252 & \\footnotesize  flared & \\footnotesize  ngc4477 \\\\ %\\footnotesize  Turner, Martin & \\footnotesize   2001-12-05 \\\\ \\footnotesize 12 30 45.5 & \\footnotesize  +13 43 32 & \\footnotesize  282.0 & \\footnotesize  75.7  & \\footnotesize  0.258 & \\footnotesize  0106060401  & \\footnotesize  13275 & \\footnotesize  8099 & \\footnotesize  ok & \\footnotesize  Virgo 4 \\\\ %\\footnotesize  Aschenbach, Bernd & \\footnotesize   2002-07-04 \\\\ \\footnotesize 12 30 45.4 & \\footnotesize  +14 43 22 & \\footnotesize  280.6 & \\footnotesize  76.6  & \\footnotesize  0.247 & \\footnotesize  0106060701  & \\footnotesize  14375 & \\footnotesize  10637 & \\footnotesize  ok & \\footnotesize  Virgo 7 \\\\ %\\footnotesize  Aschenbach, Bernd & \\footnotesize   2002-07-05 \\\\ \\footnotesize 12 30 45.3 & \\footnotesize  +15 23 29 & \\footnotesize  279.6 & \\footnotesize  77.3  & \\footnotesize  0.247 & \\footnotesize  0106060901  & \\footnotesize  15622 & \\footnotesize  8970 & \\footnotesize  ok & \\footnotesize  Virgo 9 \\\\ %\\footnotesize  Aschenbach, Bernd & \\footnotesize   2002-06-10 \\\\ \\footnotesize 12 30 45.3 & \\footnotesize  +15 43 30 & \\footnotesize  279.0 & \\footnotesize  77.6  & \\footnotesize  0.230 & \\footnotesize  0106061001  & \\footnotesize  15720 & \\footnotesize  6231 & \\footnotesize  ok & \\footnotesize  Virgo 10 \\\\ %\\footnotesize  Aschenbach, Bernd & \\footnotesize   2002-06-06 \\\\ \\footnotesize 12 30 45.2 & \\footnotesize  +16 03 33 & \\footnotesize  278.4 & \\footnotesize  77.9  & \\footnotesize  0.230 & \\footnotesize  0106061101  & \\footnotesize  16624 & \\footnotesize  7389 & \\footnotesize  ok & \\footnotesize  Virgo 11 \\\\ %\\footnotesize  Aschenbach, Bernd & \\footnotesize   2002-06-09 \\\\ \\footnotesize 12 30 45.2 & \\footnotesize  +14 03 18 & \\footnotesize  281.6 & \\footnotesize  76.0  & \\footnotesize  0.258 & \\footnotesize  0106060501  & \\footnotesize  17738 & \\footnotesize  12421 & \\footnotesize  ok & \\footnotesize  Virgo 5 \\\\ %\\footnotesize  Aschenbach, Bernd & \\footnotesize   2002-07-06 \\\\ \\footnotesize 12 30 45.4 & \\footnotesize  +14 23 22 & \\footnotesize  281.1 & \\footnotesize  76.3  & \\footnotesize  0.258 & \\footnotesize  0106060601  & \\footnotesize  14876 & \\footnotesize  8756 & \\footnotesize  ok & \\footnotesize  Virgo 6 \\\\ %\\footnotesize  Aschenbach, Bernd & \\footnotesize   2002-07-08 \\\\ \\footnotesize 12 30 45.2 & \\footnotesize  +16 23 31 & \\footnotesize  277.8 & \\footnotesize  78.2  & \\footnotesize  0.230 & \\footnotesize  0106061201  & \\footnotesize  18441 & \\footnotesize  11858 & \\footnotesize  ok & \\footnotesize  Virgo 12 \\\\ %\\footnotesize  Aschenbach, Bernd & \\footnotesize   2002-06-09 \\\\ \\footnotesize 12 30 45.4 & \\footnotesize  +17 53 28 & \\footnotesize  274.5 & \\footnotesize  79.6  & \\footnotesize  0.251 & \\footnotesize  0106061301  & \\footnotesize  16624 & \\footnotesize  12174 & \\footnotesize  ok & \\footnotesize  Virgo 13 \\\\ %\\footnotesize  Aschenbach, Bernd & \\footnotesize   2002-06-10 \\\\ \\footnotesize 12 30 45.3 & \\footnotesize  +15 03 29 & \\footnotesize  280.1 & \\footnotesize  77.0  & \\footnotesize  0.247 & \\footnotesize  0106060801  & \\footnotesize  13057 & \\footnotesize  1660 & \\footnotesize  flared & \\footnotesize  Virgo 8 \\\\ %\\footnotesize  Aschenbach, Bernd & \\footnotesize   2002-06-06 \\\\ \\footnotesize 12 30 49.1 & \\footnotesize  +12 23 22 & \\footnotesize  283.7 & \\footnotesize  74.4  & \\footnotesize  0.250 & \\footnotesize  0114120101  & \\footnotesize  60109 & \\footnotesize  30143 & \\footnotesize  discard & \\footnotesize  M87 \\\\ %\\footnotesize  Jansen, Fred & \\footnotesize   2000-06-19 \\\\ \\footnotesize 12 30 49.9 & \\footnotesize  +12 43 17 & \\footnotesize  283.4 & \\footnotesize  74.8  & \\footnotesize  0.259 & \\footnotesize  0106060101  & \\footnotesize  10020 & \\footnotesize  4160 & \\footnotesize  flared & \\footnotesize  Virgo 1 \\\\ %\\footnotesize  Aschenbach, Bernd & \\footnotesize   2001-07-12 \\\\ \\footnotesize 12 30 53.9 & \\footnotesize  +11 00 03 & \\footnotesize  285.2 & \\footnotesize  73.1  & \\footnotesize  0.209 & \\footnotesize  0145800101  & \\footnotesize  107002 & \\footnotesize  50080 & \\footnotesize  discard & \\footnotesize  LBQS 1228+1116 \\\\ %& \\footnotesize  Fujimoto, Ryuichi & \\footnotesize   2003-07-14 \\\\ \\footnotesize 12 31 58.3 & \\footnotesize  +14 25 19 & \\footnotesize  282.3 & \\footnotesize  76.5  & \\footnotesize  0.247 & \\footnotesize  0112550801  & \\footnotesize  14163 & \\footnotesize  6669 & \\footnotesize  ok & \\footnotesize  ngc4501 \\\\ %\\footnotesize  Turner, Martin & \\footnotesize   2001-12-04 \\\\ \\footnotesize 12 35 34.9 & \\footnotesize  -39 54 50 & \\footnotesize  299.6 & \\footnotesize  22.8  & \\footnotesize  0.719 & \\footnotesize  0006220201  & \\footnotesize  46207 & \\footnotesize  35856 & \\footnotesize  ok & \\footnotesize  NGC 4507 \\\\ %\\footnotesize  Turner, Tracey & \\footnotesize   2001-01-05 \\\\ \\footnotesize 12 35 39.7 & \\footnotesize  +12 33 17 & \\footnotesize  287.9 & \\footnotesize  74.9  & \\footnotesize  0.259 & \\footnotesize  0141570101  & \\footnotesize  44838 & \\footnotesize  20074 & \\footnotesize  ok & \\footnotesize  NGC 4552 \\\\ %\\footnotesize  Sarazin, Craig & \\footnotesize   2003-07-10 \\\\ \\footnotesize 12 36 39.4 & \\footnotesize  -33 54 03 & \\footnotesize  299.4 & \\footnotesize  28.8  & \\footnotesize  0.561 & \\footnotesize  0201901701  & \\footnotesize  26147 & \\footnotesize  5973 & \\footnotesize  ok & \\footnotesize  RXCJ1236.7-3354 \\\\ %\\footnotesize  Boehringer, Hans & \\footnotesize   2004-07-28 \\\\ \\footnotesize 12 37 38.8 & \\footnotesize  +11 49 03 & \\footnotesize  290.3 & \\footnotesize  74.3  & \\footnotesize  0.261 & \\footnotesize  0112840101  & \\footnotesize  23669 & \\footnotesize  15453 & \\footnotesize  ok & \\footnotesize  NGC4579 \\\\ %\\footnotesize  Griffiths, Richard & \\footnotesize   2003-06-12 \\\\ \\footnotesize 12 39 38.9 & \\footnotesize  -05 20 43 & \\footnotesize  297.4 & \\footnotesize  57.4  & \\footnotesize  0.228 & \\footnotesize  0109970101  & \\footnotesize  28062 & \\footnotesize  0 & \\footnotesize  flared & \\footnotesize  NGC 4593 \\\\ %\\footnotesize  Kaastra, Jelle & \\footnotesize   2000-07-02 \\\\ \\footnotesize 12 39 59.3 & \\footnotesize  -11 37 11 & \\footnotesize  298.4 & \\footnotesize  51.1  & \\footnotesize  0.385 & \\footnotesize  0084030101  & \\footnotesize  43456 & \\footnotesize  23727 & \\footnotesize  ok & \\footnotesize  M104 \\\\ %\\footnotesize  Fabbiano, Giuseppina & \\footnotesize   2001-12-28 \\\\ \\footnotesize 12 40 20.4 & \\footnotesize  -83 14 57 & \\footnotesize  302.5 & \\footnotesize  -20.3 & \\footnotesize  0.963 & \\footnotesize  0092360801  & \\footnotesize  16835 & \\footnotesize  11397 & \\footnotesize  ok & \\footnotesize  3EG 1249-8330 \\\\ %\\footnotesize  Caraveo, Patrizia & \\footnotesize   2001-03-30 \\\\ \\footnotesize 12 40 21.2 & \\footnotesize  -83 44 57 & \\footnotesize  302.6 & \\footnotesize  -20.8 & \\footnotesize  0.860 & \\footnotesize  0092360701  & \\footnotesize  13236 & \\footnotesize  2275 & \\footnotesize  flared & \\footnotesize  3EG 1249-8330 \\\\ %\\footnotesize  Caraveo, Patrizia & \\footnotesize   2001-03-30 \\\\ \\footnotesize 12 41 13.5 & \\footnotesize  +18 34 32 & \\footnotesize  287.0 & \\footnotesize  81.1  & \\footnotesize  0.170 & \\footnotesize  0149900301  & \\footnotesize  17412 & \\footnotesize  14070 & \\footnotesize  ok & \\footnotesize  Abell 1589 \\\\ %\\footnotesize  Forman, William & \\footnotesize   2003-06-18 \\\\ \\footnotesize 12 42 38.3 & \\footnotesize  -11 19 31 & \\footnotesize  299.4 & \\footnotesize  51.4  & \\footnotesize  0.374 & \\footnotesize  0136950201  & \\footnotesize  30265 & \\footnotesize  26157 & \\footnotesize  ok & \\footnotesize  RXJ1242 \\\\ %\\footnotesize  Jansen, Fred & \\footnotesize   2001-06-22 \\\\ \\footnotesize 12 42 48.2 & \\footnotesize  +02 41 21 & \\footnotesize  297.7 & \\footnotesize  65.4  & \\footnotesize  0.175 & \\footnotesize  0111190701  & \\footnotesize  64406 & \\footnotesize  53805 & \\footnotesize  ok & \\footnotesize  NGC4636 \\\\ %\\footnotesize  Kaastra, Jelle & \\footnotesize   2001-01-06 \\\\ \\footnotesize 12 42 51.8 & \\footnotesize  +13 15 40 & \\footnotesize  294.2 & \\footnotesize  75.9  & \\footnotesize  0.236 & \\footnotesize  0112551001  & \\footnotesize  15112 & \\footnotesize  10990 & \\footnotesize  ok & \\footnotesize  ngc4639 \\\\ %\\footnotesize  Turner, Martin & \\footnotesize   2001-12-16 \\\\ \\footnotesize 12 43 39.8 & \\footnotesize  +11 33 15 & \\footnotesize  295.8 & \\footnotesize  74.3  & \\footnotesize  0.194 & \\footnotesize  0021540201  & \\footnotesize  54210 & \\footnotesize  43902 & \\footnotesize  ok & \\footnotesize  NGC 4649 \\\\ %\\footnotesize  Sarazin, Craig & \\footnotesize   2001-01-03 \\\\ \\footnotesize 12 45 04.5 & \\footnotesize  -00 27 41 & \\footnotesize  299.5 & \\footnotesize  62.3  & \\footnotesize  0.174 & \\footnotesize  0110980201  & \\footnotesize  58237 & \\footnotesize  50432 & \\footnotesize  ok & \\footnotesize  NGC 4666 \\\\ %\\footnotesize  Jansen, Fred & \\footnotesize   2002-06-28 \\\\ \\footnotesize 12 48 22.0 & \\footnotesize  +08 29 24 & \\footnotesize  300.5 & \\footnotesize  71.3  & \\footnotesize  0.192 & \\footnotesize  0112551101  & \\footnotesize  16912 & \\footnotesize  10529 & \\footnotesize  ok & \\footnotesize  ngc4698 \\\\ %\\footnotesize  Turner, Martin & \\footnotesize   2001-12-17 \\\\ \\footnotesize 12 48 48.9 & \\footnotesize  -41 18 42 & \\footnotesize  302.4 & \\footnotesize  21.5  & \\footnotesize  0.825 & \\footnotesize  0046340101  & \\footnotesize  47807 & \\footnotesize  30632 & \\footnotesize  discard & \\footnotesize  Centaurus cluster \\\\ %& \\footnotesize  Ikebe, Yasushi & \\footnotesize   2002-01-03 \\\\ \\footnotesize 12 49 13.8 & \\footnotesize  -05 59 34 & \\footnotesize  301.9 & \\footnotesize  56.8  & \\footnotesize  0.212 & \\footnotesize  0060370201  & \\footnotesize  41273 & \\footnotesize  30143 & \\footnotesize  ok & \\footnotesize  Q1246-057 \\\\ %\\footnotesize  Mathur, Smita & \\footnotesize   2001-07-12 \\\\ \\footnotesize 12 52 24.2 & \\footnotesize  -29 15 02 & \\footnotesize  303.1 & \\footnotesize  33.6  & \\footnotesize  0.610 & \\footnotesize  0111020201  & \\footnotesize  36207 & \\footnotesize  6496 & \\footnotesize  ok & \\footnotesize  EX Hydrae \\\\ %\\footnotesize  Kaastra, Jelle & \\footnotesize   2000-07-02 \\\\ \\footnotesize 12 52 24.2 & \\footnotesize  -29 15 02 & \\footnotesize  303.1 & \\footnotesize  33.6  & \\footnotesize  0.610 & \\footnotesize  0111020101  & \\footnotesize  57802 & \\footnotesize  0 & \\footnotesize  flared & \\footnotesize  EX Hydrae \\\\ %\\footnotesize  Kaastra, Jelle & \\footnotesize   2000-07-01 \\\\ \\footnotesize 12 52 59.3 & \\footnotesize  -29 27 13 & \\footnotesize  303.3 & \\footnotesize  33.4  & \\footnotesize  0.600 & \\footnotesize  0057740301  & \\footnotesize  68621 & \\footnotesize  60049 & \\footnotesize  ok & \\footnotesize  c1252.9-2927 \\\\ %\\footnotesize  Rosati, Piero & \\footnotesize   2003-01-03 \\\\ \\footnotesize 12 52 59.3 & \\footnotesize  -29 27 03 & \\footnotesize  303.3 & \\footnotesize  33.4  & \\footnotesize  0.600 & \\footnotesize  0057740401  & \\footnotesize  68626 & \\footnotesize  60993 & \\footnotesize  ok & \\footnotesize  c1252.9-2927 \\\\ %\\footnotesize  Rosati, Piero & \\footnotesize   2003-01-11 \\\\ \\footnotesize 12 53 10.5 & \\footnotesize  -09 11 56 & \\footnotesize  303.6 & \\footnotesize  53.6  & \\footnotesize  0.293 & \\footnotesize  0112270701  & \\footnotesize  12866 & \\footnotesize  9766 & \\footnotesize  ok & \\footnotesize  HCG62 \\\\ %\\footnotesize  Turner, Martin & \\footnotesize   2003-01-15 \\\\ \\footnotesize 12 53 37.6 & \\footnotesize  +15 42 24 & \\footnotesize  305.5 & \\footnotesize  78.5  & \\footnotesize  0.191 & \\footnotesize  0082990101  & \\footnotesize  50900 & \\footnotesize  29195 & \\footnotesize  ok & \\footnotesize  3C 277.2 \\\\ %\\footnotesize  Gianfranco, Brunetti & \\footnotesize   2002-12-11 \\\\ \\footnotesize 12 54 00.0 & \\footnotesize  +10 11 25 & \\footnotesize  305.0 & \\footnotesize  73.0  & \\footnotesize  0.162 & \\footnotesize  0001930301  & \\footnotesize  25335 & \\footnotesize  18403 & \\footnotesize  ok & \\footnotesize  IRAS F12514+10 \\\\ %\\footnotesize  Wilman, Richard & \\footnotesize   2001-12-28 \\\\ \\footnotesize 12 54 33.6 & \\footnotesize  -29 08 17 & \\footnotesize  303.7 & \\footnotesize  33.7  & \\footnotesize  0.610 & \\footnotesize  0030140101  & \\footnotesize  17248 & \\footnotesize  13787 & \\footnotesize  discard & \\footnotesize  A3528 \\\\ %\\footnotesize  Ettori, Stefano & \\footnotesize   2001-12-24 \\\\ \\footnotesize 12 57 16.8 & \\footnotesize  -30 22 11 & \\footnotesize  304.4 & \\footnotesize  32.4  & \\footnotesize  0.596 & \\footnotesize  0030140301  & \\footnotesize  16895 & \\footnotesize  11650 & \\footnotesize  ok & \\footnotesize  A3532 \\\\ %\\footnotesize  Ettori, Stefano & \\footnotesize   2002-07-03 \\\\ \\footnotesize 12 57 28.2 & \\footnotesize  -17 20 29 & \\footnotesize  304.9 & \\footnotesize  45.5  & \\footnotesize  0.465 & \\footnotesize  0010420201  & \\footnotesize  22806 & \\footnotesize  12570 & \\footnotesize  discard & \\footnotesize  Abell 1644 \\\\ %\\footnotesize  Sarazin, Craig & \\footnotesize   2001-01-08 \\\\ \\footnotesize 12 58 00.5 & \\footnotesize  -83 14 57 & \\footnotesize  303.1 & \\footnotesize  -20.3 & \\footnotesize  0.710 & \\footnotesize  0092360501  & \\footnotesize  15383 & \\footnotesize  7728 & \\footnotesize  ok & \\footnotesize  3EG 1249-8330 \\\\ %\\footnotesize  Caraveo, Patrizia & \\footnotesize   2001-03-23 \\\\ \\footnotesize 12 57 59.9 & \\footnotesize  -83 44 59 & \\footnotesize  303.1 & \\footnotesize  -20.8 & \\footnotesize  0.689 & \\footnotesize  0092360601  & \\footnotesize  15380 & \\footnotesize  8747 & \\footnotesize  ok & \\footnotesize  3EG 1249-8330 \\\\ %\\footnotesize  Caraveo, Patrizia & \\footnotesize   2001-03-23 \\\\ \\footnotesize 12 58 48.0 & \\footnotesize  -01 44 44 & \\footnotesize  306.7 & \\footnotesize  61.0  & \\footnotesize  0.154 & \\footnotesize  0093200101  & \\footnotesize  43240 & \\footnotesize  33950 & \\footnotesize  ok & \\footnotesize  A1650 \\\\ %\\footnotesize  Yamashita, Koujun & \\footnotesize   2001-12-30 \\\\ \\footnotesize 13 02 46.1 & \\footnotesize  -02 30 29 & \\footnotesize  308.6 & \\footnotesize  60.2  & \\footnotesize  0.171 & \\footnotesize  0201901801  & \\footnotesize  29149 & \\footnotesize  20164 & \\footnotesize  ok & \\footnotesize  RXCJ1302.8-0230 \\\\ %\\footnotesize  Boehringer, Hans & \\footnotesize   2004-06-23 \\\\ \\footnotesize 13 05 13.1 & \\footnotesize  -10 20 50 & \\footnotesize  308.4 & \\footnotesize  52.3  & \\footnotesize  0.322 & \\footnotesize  0032141201  & \\footnotesize  16003 & \\footnotesize  12223 & \\footnotesize  ok & \\footnotesize  1saxj1305.2-1020 \\\\ %\\footnotesize  Fiore, Fabrizio & \\footnotesize   2002-01-03 \\\\ \\footnotesize 13 05 43.9 & \\footnotesize  +18 01 06 & \\footnotesize  323.5 & \\footnotesize  80.3  & \\footnotesize  0.198 & \\footnotesize  0017940101  & \\footnotesize  56505 & \\footnotesize  40187 & \\footnotesize  ok & \\footnotesize  GP Com \\\\ %\\footnotesize  Mauche, Christopher & \\footnotesize   2001-01-04 \\\\ \\footnotesize 13 11 29.3 & \\footnotesize  -01 20 18 & \\footnotesize  313.3 & \\footnotesize  61.1  & \\footnotesize  0.180 & \\footnotesize  0093030101  & \\footnotesize  39806 & \\footnotesize  32201 & \\footnotesize  ok & \\footnotesize  Abell 1689 \\\\ %\\footnotesize  Hughes, John & \\footnotesize   2001-12-24 \\\\ \\footnotesize 13 15 23.8 & \\footnotesize  -16 22 55 & \\footnotesize  311.2 & \\footnotesize  46.1  & \\footnotesize  0.491 & \\footnotesize  0037950101  & \\footnotesize  24109 & \\footnotesize  18674 & \\footnotesize  ok & \\footnotesize  NGC 5044 \\\\ %\\footnotesize  Mathews, William & \\footnotesize   2001-01-13 \\\\ \\footnotesize 13 15 28.6 & \\footnotesize  -16 45 29 & \\footnotesize  311.1 & \\footnotesize  45.7  & \\footnotesize  0.491 & \\footnotesize  0152360101  & \\footnotesize  39116 & \\footnotesize  34508 & \\footnotesize  ok & \\footnotesize  NGC 5044 \\\\ %\\footnotesize  Lewis, Aaron & \\footnotesize   2003-01-09 \\\\ \\footnotesize 13 19 16.5 & \\footnotesize  -14 50 42 & \\footnotesize  312.9 & \\footnotesize  47.4  & \\footnotesize  0.491 & \\footnotesize  0110980601  & \\footnotesize  53647 & \\footnotesize  27626 & \\footnotesize  ok & \\footnotesize  NGC 5073 \\\\ %\\footnotesize  Jansen, Fred & \\footnotesize   2002-07-06 \\\\ \\footnotesize 13 24 05.1 & \\footnotesize  +13 58 28 & \\footnotesize  334.6 & \\footnotesize  74.8  & \\footnotesize  0.175 & \\footnotesize  0108860701  & \\footnotesize  26193 & \\footnotesize  14219 & \\footnotesize  ok & \\footnotesize  NGC 5129 \\\\ %\\footnotesize  Mushotzky, Richard & \\footnotesize   2002-07-08 \\\\ \\footnotesize 13 25 19.9 & \\footnotesize  -38 24 43 & \\footnotesize  310.1 & \\footnotesize  23.9  & \\footnotesize  0.481 & \\footnotesize  0110890101  & \\footnotesize  64019 & \\footnotesize  52553 & \\footnotesize  ok & \\footnotesize  IRAS 13224-3809 \\\\ %\\footnotesize  Watson, Michael & \\footnotesize   2002-01-19 \\\\ \\footnotesize 13 27 57.2 & \\footnotesize  -31 30 19 & \\footnotesize  311.9 & \\footnotesize  30.7  & \\footnotesize  0.363 & \\footnotesize  0107260101  & \\footnotesize  44615 & \\footnotesize  37883 & \\footnotesize  discard & \\footnotesize  A3558 \\\\ %\\footnotesize  Aschenbach, Bernd & \\footnotesize   2002-01-22 \\\\ \\footnotesize 13 29 21.7 & \\footnotesize  +11 46 24 & \\footnotesize  334.8 & \\footnotesize  72.2  & \\footnotesize  0.195 & \\footnotesize  0041180801  & \\footnotesize  24948 & \\footnotesize  16594 & \\footnotesize  ok & \\footnotesize  NGC 5171 Group \\\\ %\\footnotesize  Ponman, Trevor & \\footnotesize   2001-12-31 \\\\ \\footnotesize 13 30 53.8 & \\footnotesize  -01 50 15 & \\footnotesize  322.6 & \\footnotesize  59.5  & \\footnotesize  0.232 & \\footnotesize  0112240301  & \\footnotesize  34779 & \\footnotesize  26166 & \\footnotesize  ok & \\footnotesize  A1750 \\\\ %\\footnotesize  Turner, Martin & \\footnotesize   2001-07-29 \\\\ \\footnotesize 13 31 45.6 & \\footnotesize  -31 48 02 & \\footnotesize  312.8 & \\footnotesize  30.2  & \\footnotesize  0.398 & \\footnotesize  0105261401  & \\footnotesize  23338 & \\footnotesize  10413 & \\footnotesize  ok & \\footnotesize  A3562\\_f2 \\\\ %\\footnotesize  Aschenbach, Bernd & \\footnotesize   2002-07-27 \\\\ \\footnotesize 13 32 24.9 & \\footnotesize  -31 37 27 & \\footnotesize  313.0 & \\footnotesize  30.4  & \\footnotesize  0.391 & \\footnotesize  0105261601  & \\footnotesize  23847 & \\footnotesize  17065 & \\footnotesize  ok & \\footnotesize  A3562\\_f4 \\\\ %\\footnotesize  Aschenbach, Bernd & \\footnotesize   2002-08-02 \\\\ \\footnotesize 13 32 55.3 & \\footnotesize  -31 52 05 & \\footnotesize  313.1 & \\footnotesize  30.1  & \\footnotesize  0.398 & \\footnotesize  0105261701  & \\footnotesize  20911 & \\footnotesize  17323 & \\footnotesize  ok & \\footnotesize  A3562\\_f5 \\\\ %\\footnotesize  Aschenbach, Bernd & \\footnotesize   2003-01-21 \\\\ \\footnotesize 13 33 05.4 & \\footnotesize  -31 41 25 & \\footnotesize  313.1 & \\footnotesize  30.3  & \\footnotesize  0.391 & \\footnotesize  0105261301  & \\footnotesize  47167 & \\footnotesize  36328 & \\footnotesize  ok & \\footnotesize  A3562\\_f1 \\\\ %\\footnotesize  Aschenbach, Bernd & \\footnotesize   2003-01-13 \\\\ \\footnotesize 13 33 54.7 & \\footnotesize  -31 30 26 & \\footnotesize  313.4 & \\footnotesize  30.4  & \\footnotesize  0.391 & \\footnotesize  0105261501  & \\footnotesize  22648 & \\footnotesize  18919 & \\footnotesize  ok & \\footnotesize  A3562\\_f3 \\\\ %\\footnotesize  Aschenbach, Bernd & \\footnotesize   2002-07-31 \\\\ \\footnotesize 13 34 25.2 & \\footnotesize  -31 44 20 & \\footnotesize  313.5 & \\footnotesize  30.2  & \\footnotesize  0.391 & \\footnotesize  0105261801  & \\footnotesize  21847 & \\footnotesize  8385 & \\footnotesize  ok & \\footnotesize  A3562\\_f6 \\\\ %\\footnotesize  Aschenbach, Bernd & \\footnotesize   2003-01-27 \\\\ \\footnotesize 13 35 53.9 & \\footnotesize  -34 17 49 & \\footnotesize  313.2 & \\footnotesize  27.6  & \\footnotesize  0.409 & \\footnotesize  0111570101  & \\footnotesize  46453 & \\footnotesize  8000 & \\footnotesize  ok & \\footnotesize  MCG-6-30-15 \\\\ %\\footnotesize  Kaastra, Jelle & \\footnotesize   2000-07-11 \\\\ \\footnotesize 13 35 53.9 & \\footnotesize  -34 17 47 & \\footnotesize  313.2 & \\footnotesize  27.6  & \\footnotesize  0.409 & \\footnotesize  0111570201  & \\footnotesize  66197 & \\footnotesize  8000 & \\footnotesize  ok & \\footnotesize  MCG-6-30-15 \\\\ %\\footnotesize  Kaastra, Jelle & \\footnotesize   2000-07-12 \\\\ \\footnotesize 13 37 05.1 & \\footnotesize  -29 51 46 & \\footnotesize  314.6 & \\footnotesize  31.9  & \\footnotesize  0.397 & \\footnotesize  0110910201  & \\footnotesize  30638 & \\footnotesize  20121 & \\footnotesize  ok & \\footnotesize  M 83 \\\\ %\\footnotesize  Watson, Michael & \\footnotesize   2003-01-27 \\\\ \\footnotesize 13 37 44.4 & \\footnotesize  -12 57 20 & \\footnotesize  320.0 & \\footnotesize  48.3  & \\footnotesize  0.481 & \\footnotesize  0147670201  & \\footnotesize  13917 & \\footnotesize  11814 & \\footnotesize  ok & \\footnotesize  PKS B1334-127 \\\\ %\\footnotesize  Fiore, Fabrizio & \\footnotesize   2003-01-31 \\\\ \\footnotesize 13 38 11.5 & \\footnotesize  +04 32 24 & \\footnotesize  331.2 & \\footnotesize  64.8  & \\footnotesize  0.197 & \\footnotesize  0152940101  & \\footnotesize  67313 & \\footnotesize  30686 & \\footnotesize  ok & \\footnotesize  NGC 5252 \\\\ %\\footnotesize  Dadina, Mauro & \\footnotesize   2003-07-18 \\\\ \\footnotesize 13 39 56.2 & \\footnotesize  -31 38 39 & \\footnotesize  314.8 & \\footnotesize  30.1  & \\footnotesize  0.390 & \\footnotesize  0035940301  & \\footnotesize  47216 & \\footnotesize  29616 & \\footnotesize  ok & \\footnotesize  NGC 5253 \\\\ %\\footnotesize  Heckman, Timothy & \\footnotesize   2001-08-09 \\\\ \\footnotesize 13 41 19.7 & \\footnotesize  +00 23 53 & \\footnotesize  329.1 & \\footnotesize  60.7  & \\footnotesize  0.186 & \\footnotesize  0111281001  & \\footnotesize  10377 & \\footnotesize  7333 & \\footnotesize  ok & \\footnotesize  F864-1 \\\\ %\\footnotesize  Watson, Michael & \\footnotesize   2002-07-12 \\\\ \\footnotesize 13 41 19.6 & \\footnotesize  -00 00 07 & \\footnotesize  328.8 & \\footnotesize  60.3  & \\footnotesize  0.191 & \\footnotesize  0111281301  & \\footnotesize  14541 & \\footnotesize  2342 & \\footnotesize  flared & \\footnotesize  F864-4 \\\\ %\\footnotesize  Watson, Michael & \\footnotesize   2002-07-27 \\\\ \\footnotesize 13 47 15.5 & \\footnotesize  +17 27 21 & \\footnotesize  358.9 & \\footnotesize  73.8  & \\footnotesize  0.181 & \\footnotesize  0144570101  & \\footnotesize  66431 & \\footnotesize  39983 & \\footnotesize  ok & \\footnotesize  Tau Boo \\\\ %\\footnotesize  Maggio, Antonio & \\footnotesize   2003-06-24 \\\\ \\footnotesize 13 47 22.7 & \\footnotesize  -33 09 54 & \\footnotesize  316.2 & \\footnotesize  28.2  & \\footnotesize  0.430 & \\footnotesize  0147040101  & \\footnotesize  31428 & \\footnotesize  11078 & \\footnotesize  discard & \\footnotesize  Abell 3571 \\\\ %\\footnotesize  De Grandi, Sabrina & \\footnotesize   2003-08-10 \\\\ \\footnotesize 13 47 30.7 & \\footnotesize  -11 45 11 & \\footnotesize  324.0 & \\footnotesize  48.8  & \\footnotesize  0.492 & \\footnotesize  0112960101  & \\footnotesize  38392 & \\footnotesize  30423 & \\footnotesize  ok & \\footnotesize  RXJ1347-1145 \\\\ %\\footnotesize  Turner, Martin & \\footnotesize   2002-08-01 \\\\ \\footnotesize 13 47 30.1 & \\footnotesize  -32 51 52 & \\footnotesize  316.3 & \\footnotesize  28.5  & \\footnotesize  0.393 & \\footnotesize  0086950201  & \\footnotesize  33642 & \\footnotesize  23354 & \\footnotesize  discard & \\footnotesize  A3571 \\\\ %\\footnotesize  Dupke, Renato & \\footnotesize   2002-07-29 \\\\ \\footnotesize 13 49 19.0 & \\footnotesize  -30 18 30 & \\footnotesize  317.4 & \\footnotesize  30.9  & \\footnotesize  0.443 & \\footnotesize  0101040401  & \\footnotesize  13925 & \\footnotesize  9262 & \\footnotesize  discard & \\footnotesize  IC 4329a \\\\ %\\footnotesize  Jansen, Fred & \\footnotesize   2001-01-31 \\\\ \\footnotesize 13 50 43.9 & \\footnotesize  -33 43 05 & \\footnotesize  316.8 & \\footnotesize  27.5  & \\footnotesize  0.480 & \\footnotesize  0201901901  & \\footnotesize  26515 & \\footnotesize  15442 & \\footnotesize  ok & \\footnotesize  RXCJ1350.7-3343 \\\\ %\\footnotesize  Boehringer, Hans & \\footnotesize   2004-02-16 \\\\ \\footnotesize 13 54 12.6 & \\footnotesize  -02 21 52 & \\footnotesize  332.5 & \\footnotesize  56.8  & \\footnotesize  0.353 & \\footnotesize  0112250101  & \\footnotesize  33646 & \\footnotesize  22728 & \\footnotesize  ok & \\footnotesize  RXJ1354.3-0222 \\\\ %\\footnotesize  Turner, Martin & \\footnotesize   2002-07-20 \\\\ \\footnotesize 13 56 03.0 & \\footnotesize  +18 22 33 & \\footnotesize  5.8   & \\footnotesize  72.7  & \\footnotesize  0.208 & \\footnotesize  0094401201  & \\footnotesize  26850 & \\footnotesize  22981 & \\footnotesize  ok & \\footnotesize  Mrk 463 \\\\ %\\footnotesize  Sanders, David & \\footnotesize   2001-12-22 \\\\ \\footnotesize 14 01 01.8 & \\footnotesize  +02 52 40 & \\footnotesize  340.3 & \\footnotesize  60.5  & \\footnotesize  0.224 & \\footnotesize  0098010101  & \\footnotesize  61353 & \\footnotesize  27531 & \\footnotesize  ok & \\footnotesize  A 1835 \\\\ %\\footnotesize  Kaastra, Jelle & \\footnotesize   2000-06-28 \\\\ \\footnotesize 14 01 02.0 & \\footnotesize  +02 52 33 & \\footnotesize  340.3 & \\footnotesize  60.5  & \\footnotesize  0.224 & \\footnotesize  0147330201  & \\footnotesize  106844 & \\footnotesize  38487 & \\footnotesize  discard & \\footnotesize  Abell 1835 \\\\ %\\footnotesize  Fabian, Andrew & \\footnotesize   2003-07-30 \\\\ \\footnotesize 14 01 34.6 & \\footnotesize  -11 07 35 & \\footnotesize  329.2 & \\footnotesize  48.1  & \\footnotesize  0.472 & \\footnotesize  0109910101  & \\footnotesize  100219 & \\footnotesize  42837 & \\footnotesize  ok & \\footnotesize  A 1837 \\\\ %\\footnotesize  Kaastra, Jelle & \\footnotesize   2001-01-12 \\\\ \\footnotesize 14 04 16.6 & \\footnotesize  -34 02 05 & \\footnotesize  319.7 & \\footnotesize  26.4  & \\footnotesize  0.549 & \\footnotesize  0058940301  & \\footnotesize  20049 & \\footnotesize  15753 & \\footnotesize  ok & \\footnotesize  1401-33 \\\\ %\\footnotesize  Harrison, Fiona & \\footnotesize   2002-02-14 \\\\ \\footnotesize 14 08 51.9 & \\footnotesize  -07 52 38 & \\footnotesize  333.8 & \\footnotesize  50.2  & \\footnotesize  0.273 & \\footnotesize  0151590101  & \\footnotesize  32318 & \\footnotesize  10405 & \\footnotesize  ok & \\footnotesize  pks1406-076 \\\\ %\\footnotesize  McEnery, Julie & \\footnotesize   2003-07-05 \\\\ \\footnotesize 14 08 51.9 & \\footnotesize  -07 52 43 & \\footnotesize  333.8 & \\footnotesize  50.2  & \\footnotesize  0.273 & \\footnotesize  0151590201  & \\footnotesize  14299 & \\footnotesize  4257 & \\footnotesize  flared & \\footnotesize  pks1406-076 \\\\ %\\footnotesize  McEnery, Julie & \\footnotesize   2003-08-10 \\\\ \\footnotesize 14 14 52.4 & \\footnotesize  -00 23 45 & \\footnotesize  342.4 & \\footnotesize  55.9  & \\footnotesize  0.322 & \\footnotesize  0145480101  & \\footnotesize  23567 & \\footnotesize  12539 & \\footnotesize  ok & \\footnotesize  Abell 1882 \\\\ %\\footnotesize  Miller, Chris & \\footnotesize   2003-02-08 \\\\ \\footnotesize 14 15 41.3 & \\footnotesize  +11 29 27 & \\footnotesize  358.6 & \\footnotesize  64.7  & \\footnotesize  0.180 & \\footnotesize  0112251301  & \\footnotesize  29642 & \\footnotesize  25672 & \\footnotesize  ok & \\footnotesize  H1413+117 \\\\ %\\footnotesize  Turner, Martin & \\footnotesize   2002-08-03 \\\\ \\footnotesize 14 15 45.9 & \\footnotesize  +11 29 28 & \\footnotesize  358.7 & \\footnotesize  64.7  & \\footnotesize  0.180 & \\footnotesize  0112250301  & \\footnotesize  26643 & \\footnotesize  21461 & \\footnotesize  ok & \\footnotesize  H1413+117 \\\\ %\\footnotesize  Turner, Martin & \\footnotesize   2001-07-30 \\\\ \\footnotesize 14 29 06.5 & \\footnotesize  +01 17 01 & \\footnotesize  349.2 & \\footnotesize  55.1  & \\footnotesize  0.285 & \\footnotesize  0102040501  & \\footnotesize  17951 & \\footnotesize  2069 & \\footnotesize  flared & \\footnotesize  MARK 1383 \\\\ %\\footnotesize  Jansen, Fred & \\footnotesize   2000-07-28 \\\\ \\footnotesize 14 37 33.7 & \\footnotesize  -15 00 31 & \\footnotesize  337.5 & \\footnotesize  40.6  & \\footnotesize  0.778 & \\footnotesize  0081341401  & \\footnotesize  21892 & \\footnotesize  14225 & \\footnotesize  ok & \\footnotesize  IRAS14348-14 \\\\ %\\footnotesize  Franceschini, Alberto & \\footnotesize   2002-07-30 \\\\ \\footnotesize 14 49 23.7 & \\footnotesize  -10 10 23 & \\footnotesize  344.4 & \\footnotesize  42.9  & \\footnotesize  0.788 & \\footnotesize  0149620201  & \\footnotesize  72555 & \\footnotesize  17408 & \\footnotesize  ok & \\footnotesize  SN 1995N \\\\ %\\footnotesize  Zampieri, Luca & \\footnotesize   2003-07-28 \\\\ \\footnotesize 14 49 23.4 & \\footnotesize  +08 59 53 & \\footnotesize  5.4   & \\footnotesize  56.7  & \\footnotesize  0.202 & \\footnotesize  0148520101  & \\footnotesize  32854 & \\footnotesize  20727 & \\footnotesize  ok & \\footnotesize  Daddi field \\\\ %\\footnotesize  Andrea, Comastri & \\footnotesize   2003-07-22 \\\\ \\footnotesize 14 49 23.4 & \\footnotesize  +08 59 52 & \\footnotesize  5.4   & \\footnotesize  56.7  & \\footnotesize  0.202 & \\footnotesize  0148520301  & \\footnotesize  32664 & \\footnotesize  26914 & \\footnotesize  ok & \\footnotesize  Daddi field \\\\ %\\footnotesize  Andrea, Comastri & \\footnotesize   2003-08-17 \\\\ \\footnotesize 14 49 28.7 & \\footnotesize  +08 59 47 & \\footnotesize  5.4   & \\footnotesize  56.7  & \\footnotesize  0.202 & \\footnotesize  0057560301  & \\footnotesize  43147 & \\footnotesize  34519 & \\footnotesize  ok & \\footnotesize  ERO field \\\\ %\\footnotesize  Andrea, Comastri & \\footnotesize   2001-08-09 \\\\ \\footnotesize 14 52 43.7 & \\footnotesize  +16 45 00 & \\footnotesize  18.9  & \\footnotesize  60.1  & \\footnotesize  0.184 & \\footnotesize  0091140201  & \\footnotesize  32749 & \\footnotesize  26980 & \\footnotesize  ok & \\footnotesize  A1983 \\\\ %\\footnotesize  ARNAUD, Monique & \\footnotesize   2002-02-15 \\\\ \\footnotesize 14 53 53.2 & \\footnotesize  +03 32 34 & \\footnotesize  359.4 & \\footnotesize  52.4  & \\footnotesize  0.362 & \\footnotesize  0150350101  & \\footnotesize  47159 & \\footnotesize  19737 & \\footnotesize  ok & \\footnotesize  NGC5775 \\\\ %\\footnotesize  Dettmar, Ralf-Juergen & \\footnotesize   2003-07-28 \\\\ \\footnotesize 14 54 36.3 & \\footnotesize  +18 38 53 & \\footnotesize  22.8  & \\footnotesize  60.4  & \\footnotesize  0.239 & \\footnotesize  0145020101  & \\footnotesize  42037 & \\footnotesize  27935 & \\footnotesize  ok & \\footnotesize  A 1991 \\\\ %\\footnotesize  ARNAUD, Monique & \\footnotesize   2003-02-16 \\\\ \\footnotesize 14 58 21.8 & \\footnotesize  -31 40 17 & \\footnotesize  332.2 & \\footnotesize  23.8  & \\footnotesize  0.958 & \\footnotesize  0067750101  & \\footnotesize  53559 & \\footnotesize  20664 & \\footnotesize  ok & \\footnotesize  Cen X-4 \\\\ %\\footnotesize  Rutledge, Robert & \\footnotesize   2001-08-21 \\\\ \\footnotesize 15 06 29.5 & \\footnotesize  +01 36 24 & \\footnotesize  0.4   & \\footnotesize  48.7  & \\footnotesize  0.425 & \\footnotesize  0021540101  & \\footnotesize  30178 & \\footnotesize  13115 & \\footnotesize  ok & \\footnotesize  NGC 5846 \\\\ %\\footnotesize  Sarazin, Craig & \\footnotesize   2001-01-25 \\\\ \\footnotesize 15 10 51.6 & \\footnotesize  +05 44 22 & \\footnotesize  6.4   & \\footnotesize  50.5  & \\footnotesize  0.307 & \\footnotesize  0111270201  & \\footnotesize  22222 & \\footnotesize  10337 & \\footnotesize  discard & \\footnotesize  A2029 \\\\ %\\footnotesize  Watson, Michael & \\footnotesize   2002-08-25 \\\\ \\footnotesize 15 16 14.6 & \\footnotesize  +00 05 46 & \\footnotesize  1.1   & \\footnotesize  45.9  & \\footnotesize  0.459 & \\footnotesize  0201902001  & \\footnotesize  28413 & \\footnotesize  23397 & \\footnotesize  ok & \\footnotesize  RXCJ1516.3+0005 \\\\ %\\footnotesize  Boehringer, Hans & \\footnotesize   2004-07-23 \\\\ \\footnotesize 15 16 29.4 & \\footnotesize  -00 57 06 & \\footnotesize  0.0   & \\footnotesize  45.2  & \\footnotesize  0.548 & \\footnotesize  0201902101  & \\footnotesize  30416 & \\footnotesize  25126 & \\footnotesize  ok & \\footnotesize  RXCJ1516.5-0056 \\\\ %\\footnotesize  Boehringer, Hans & \\footnotesize   2004-08-03 \\\\ \\footnotesize 15 16 35.5 & \\footnotesize  +00 14 41 & \\footnotesize  1.3   & \\footnotesize  45.9  & \\footnotesize  0.459 & \\footnotesize  0103860601  & \\footnotesize  13646 & \\footnotesize  10581 & \\footnotesize  ok & \\footnotesize  CGCG 21-63 \\\\ %\\footnotesize  Aschenbach, Bernd & \\footnotesize   2002-08-02 \\\\ \\footnotesize 15 16 43.9 & \\footnotesize  +07 01 12 & \\footnotesize  9.4   & \\footnotesize  50.1  & \\footnotesize  0.290 & \\footnotesize  0109920101  & \\footnotesize  31597 & \\footnotesize  24994 & \\footnotesize  ok & \\footnotesize  A 2052 \\\\ %\\footnotesize  Kaastra, Jelle & \\footnotesize   2000-08-21 \\\\ \\footnotesize 15 17 00.9 & \\footnotesize  -16 08 25 & \\footnotesize  346.1 & \\footnotesize  34.1  & \\footnotesize  0.840 & \\footnotesize  0164570401  & \\footnotesize  53218 & \\footnotesize  33847 & \\footnotesize  ok & \\footnotesize  GRB040827 \\\\ %\\footnotesize  Schartel, Norbert & \\footnotesize   2004-08-28 \\\\ \\footnotesize 15 21 49.6 & \\footnotesize  +07 41 18 & \\footnotesize  11.3  & \\footnotesize  49.4  & \\footnotesize  0.315 & \\footnotesize  0109930101  & \\footnotesize  55413 & \\footnotesize  32636 & \\footnotesize  ok & \\footnotesize  MKW 3s \\\\ %\\footnotesize  Kaastra, Jelle & \\footnotesize   2000-08-23 \\\\ \\footnotesize 15 32 22.7 & \\footnotesize  -08 32 07 & \\footnotesize  356.2 & \\footnotesize  37.1  & \\footnotesize  0.868 & \\footnotesize  0100240701  & \\footnotesize  29906 & \\footnotesize  14965 & \\footnotesize  ok & \\footnotesize  UZ LIB \\\\ %\\footnotesize  Jansen, Fred & \\footnotesize   2000-08-19 \\\\ \\footnotesize 15 32 23.0 & \\footnotesize  -08 31 56 & \\footnotesize  356.2 & \\footnotesize  37.1  & \\footnotesize  0.868 & \\footnotesize  0100240801  & \\footnotesize  30204 & \\footnotesize  22437 & \\footnotesize  ok & \\footnotesize  UZ LIB \\\\ %\\footnotesize  Jansen, Fred & \\footnotesize   2001-02-01 \\\\ \\footnotesize 15 32 29.2 & \\footnotesize  +04 40 42 & \\footnotesize  9.8   & \\footnotesize  45.5  & \\footnotesize  0.400 & \\footnotesize  0091140401  & \\footnotesize  45414 & \\footnotesize  24947 & \\footnotesize  ok & \\footnotesize  MKW9 \\\\ %\\footnotesize  ARNAUD, Monique & \\footnotesize   2001-08-21 \\\\ \\footnotesize 15 35 02.0 & \\footnotesize  +01 20 39 & \\footnotesize  6.6   & \\footnotesize  43.0  & \\footnotesize  0.466 & \\footnotesize  0112190401  & \\footnotesize  15111 & \\footnotesize  3593 & \\footnotesize  flared & \\footnotesize  MS1532.5+0130 \\\\ %\\footnotesize  Mason, Keith & \\footnotesize   2001-08-31 \\\\ \\footnotesize 15 40 48.3 & \\footnotesize  -24 43 07 & \\footnotesize  344.9 & \\footnotesize  24.0  & \\footnotesize  1.347 & \\footnotesize  0152630301  & \\footnotesize  13911 & \\footnotesize  5242 & \\footnotesize  ok & \\footnotesize  SXRB\\_3 \\\\ %\\footnotesize  Willingale, Richard & \\footnotesize   2003-08-11 \\\\ \\footnotesize 15 50 37.0 & \\footnotesize  -03 53 27 & \\footnotesize  4.3   & \\footnotesize  36.8  & \\footnotesize  0.974 & \\footnotesize  0138951501  & \\footnotesize  12412 & \\footnotesize  2336 & \\footnotesize  flared & \\footnotesize  IRAS15480-0344 \\\\ %\\footnotesize  Guainazzi, Matteo & \\footnotesize   2003-07-31 \\\\ \\footnotesize 15 54 34.7 & \\footnotesize  -25 14 54 & \\footnotesize  347.1 & \\footnotesize  21.5  & \\footnotesize  1.224 & \\footnotesize  0142630301  & \\footnotesize  22210 & \\footnotesize  15005 & \\footnotesize  ok & \\footnotesize  3 Sco \\\\ %\\footnotesize  Drake, Stephen & \\footnotesize   2003-09-06 \\\\ \\footnotesize 15 56 25.2 & \\footnotesize  -23 38 00 & \\footnotesize  348.6 & \\footnotesize  22.4  & \\footnotesize  1.159 & \\footnotesize  0112380101  & \\footnotesize  48309 & \\footnotesize  33833 & \\footnotesize  ok & \\footnotesize  SCO-CEN PT1 \\\\ %\\footnotesize  Turner, Martin & \\footnotesize   2000-08-27 \\\\ \\footnotesize 16 01 14.2 & \\footnotesize  +08 44 22 & \\footnotesize  19.8  & \\footnotesize  41.6  & \\footnotesize  0.376 & \\footnotesize  0147580201  & \\footnotesize  13305 & \\footnotesize  3462 & \\footnotesize flared &\\footnotesize 1AXGJ160118+08 \\\\ %& \\footnotesize  Uneda, Yoshihiro & \\footnotesize   2003-08-09 \\\\ \\footnotesize 16 05 02.3 & \\footnotesize  +17 44 50 & \\footnotesize  31.5  & \\footnotesize  44.5  & \\footnotesize  0.336 & \\footnotesize  0147210301  & \\footnotesize  19008 & \\footnotesize  6312 & \\footnotesize  ok & \\footnotesize  Abell 2151 \\\\ %\\footnotesize  Ponman, Trevor & \\footnotesize   2003-08-11 \\\\ \\footnotesize 16 05 02.4 & \\footnotesize  +17 44 49 & \\footnotesize  31.5  & \\footnotesize  44.5  & \\footnotesize  0.336 & \\footnotesize  0147210101  & \\footnotesize  29850 & \\footnotesize  4682 & \\footnotesize  flared & \\footnotesize  Abell 2151 \\\\ %\\footnotesize  Ponman, Trevor & \\footnotesize   2003-08-09 \\\\ \\footnotesize 16 05 06.4 & \\footnotesize  +17 45 08 & \\footnotesize  31.5  & \\footnotesize  44.5  & \\footnotesize  0.336 & \\footnotesize  0147210201  & \\footnotesize  29598 & \\footnotesize  1204 & \\footnotesize  flared & \\footnotesize  Abell 2151 \\\\ %\\footnotesize  Ponman, Trevor & \\footnotesize   2003-03-11 \\\\ \\footnotesize 16 07 13.2 & \\footnotesize  +08 04 58 & \\footnotesize  20.0  & \\footnotesize  40.0  & \\footnotesize  0.395 & \\footnotesize  0067340601  & \\footnotesize  15391 & \\footnotesize  10580 & \\footnotesize  ok & \\footnotesize  Field VI \\\\ %\\footnotesize  Willingale, Richard & \\footnotesize   2001-03-04 \\\\ \\footnotesize 16 13 06.9 & \\footnotesize  -83 42 17 & \\footnotesize  308.2 & \\footnotesize  -23.0 & \\footnotesize  0.898 & \\footnotesize  0008820401  & \\footnotesize  54463 & \\footnotesize  35628 & \\footnotesize  ok & \\footnotesize  GRB 020321 \\\\ %\\footnotesize  Rauw, Gregor & \\footnotesize   2002-03-22 \\\\ \\footnotesize 16 14 34.2 & \\footnotesize  -06 09 14 & \\footnotesize  6.5   & \\footnotesize  30.7  & \\footnotesize  1.184 & \\footnotesize  0112230901  & \\footnotesize  30809 & \\footnotesize  7591 & \\footnotesize  ok & \\footnotesize  A2163 of3 \\\\ %\\footnotesize  Turner, Martin & \\footnotesize   2000-08-31 \\\\ \\footnotesize 16 15 46.1 & \\footnotesize  -06 09 09 & \\footnotesize  6.7   & \\footnotesize  30.4  & \\footnotesize  1.227 & \\footnotesize  0112230601  & \\footnotesize  16364 & \\footnotesize  8421 & \\footnotesize  ok & \\footnotesize  A2163 \\\\ %\\footnotesize  Turner, Martin & \\footnotesize   2000-08-29 \\\\ \\footnotesize 16 15 45.9 & \\footnotesize  -06 09 11 & \\footnotesize  6.7   & \\footnotesize  30.4  & \\footnotesize  1.227 & \\footnotesize  0112231501  & \\footnotesize  13263 & \\footnotesize  5502 & \\footnotesize  ok & \\footnotesize  A2163 \\\\ %\\footnotesize  Turner, Martin & \\footnotesize   2001-08-29 \\\\ \\footnotesize 16 15 46.1 & \\footnotesize  -05 51 16 & \\footnotesize  7.0   & \\footnotesize  30.6  & \\footnotesize  1.184 & \\footnotesize  0112230801  & \\footnotesize  29400 & \\footnotesize  15264 & \\footnotesize  ok & \\footnotesize  A2163 of2 \\\\ %\\footnotesize  Turner, Martin & \\footnotesize   2000-08-31 \\\\ \\footnotesize 16 15 46.0 & \\footnotesize  -06 27 13 & \\footnotesize  6.4   & \\footnotesize  30.2  & \\footnotesize  1.227 & \\footnotesize  0112231001  & \\footnotesize  29695 & \\footnotesize  21671 & \\footnotesize  ok & \\footnotesize  A2163 of4 \\\\ %\\footnotesize  Turner, Martin & \\footnotesize   2000-09-08 \\\\ \\footnotesize 16 16 28.6 & \\footnotesize  +12 12 24 & \\footnotesize  26.1  & \\footnotesize  39.8  & \\footnotesize  0.457 & \\footnotesize  0103460801  & \\footnotesize  14448 & \\footnotesize  10925 & \\footnotesize  ok & \\footnotesize  McNaught-Hartley \\\\ %\\footnotesize  Aschenbach, Bernd & \\footnotesize   2001-01-29 \\\\ \\footnotesize 16 16 47.8 & \\footnotesize  -06 09 12 & \\footnotesize  6.9   & \\footnotesize  30.2  & \\footnotesize  1.227 & \\footnotesize  0112230701  & \\footnotesize  30909 & \\footnotesize  11901 & \\footnotesize  ok & \\footnotesize  A2163 of1 \\\\ %\\footnotesize  Turner, Martin & \\footnotesize   2000-08-29 \\\\ \\footnotesize 16 17 02.2 & \\footnotesize  +12 24 23 & \\footnotesize  26.4  & \\footnotesize  39.8  & \\footnotesize  0.457 & \\footnotesize  0103460901  & \\footnotesize  15700 & \\footnotesize  11978 & \\footnotesize  ok & \\footnotesize  McNaught-Hartley \\\\ %\\footnotesize  Aschenbach, Bernd & \\footnotesize   2001-01-29 \\\\ \\footnotesize 16 17 37.0 & \\footnotesize  +12 37 00 & \\footnotesize  26.7  & \\footnotesize  39.7  & \\footnotesize  0.427 & \\footnotesize  0103461001  & \\footnotesize  13100 & \\footnotesize  9580 & \\footnotesize  ok & \\footnotesize  McNaught-Hartley \\\\ %\\footnotesize  Aschenbach, Bernd & \\footnotesize   2001-01-30 \\\\ \\footnotesize 16 21 34.3 & \\footnotesize  -15 42 30 & \\footnotesize  359.3 & \\footnotesize  23.4  & \\footnotesize  1.499 & \\footnotesize  0153950201  & \\footnotesize  25040 & \\footnotesize 20409 & \\footnotesize discard & \\footnotesize Sco X-1 large \\\\ %offset & \\footnotesize  Jansen, Fred & \\footnotesize   2002-02-18 \\\\ \\footnotesize 16 21 44.6 & \\footnotesize  -02 16 31 & \\footnotesize  11.4  & \\footnotesize  31.4  & \\footnotesize  0.847 & \\footnotesize  0164560801  & \\footnotesize  26462 & \\footnotesize  6683 & \\footnotesize  ok & \\footnotesize  SN 2004dk \\\\ %\\footnotesize  Schartel, Norbert & \\footnotesize   2004-08-12 \\\\ \\footnotesize 16 32 46.2 & \\footnotesize  +05 34 35 & \\footnotesize  21.0  & \\footnotesize  33.2  & \\footnotesize  0.594 & \\footnotesize  0112230301  & \\footnotesize  22964 & \\footnotesize  16543 & \\footnotesize  ok & \\footnotesize  A2204 \\\\ %\\footnotesize  Turner, Martin & \\footnotesize   2001-09-12 \\\\ \\footnotesize 16 42 16.7 & \\footnotesize  +03 11 04 & \\footnotesize  20.0  & \\footnotesize  30.0  & \\footnotesize  0.535 & \\footnotesize  0067340501  & \\footnotesize  15278 & \\footnotesize  10211 & \\footnotesize  ok & \\footnotesize  Field V \\\\ %\\footnotesize  Willingale, Richard & \\footnotesize   2001-02-28 \\\\ \\footnotesize 16 52 58.4 & \\footnotesize  +02 24 18 & \\footnotesize  20.7  & \\footnotesize  27.2  & \\footnotesize  0.549 & \\footnotesize  0101640601  & \\footnotesize  19003 & \\footnotesize  7532 & \\footnotesize  ok & \\footnotesize  NGC 6240 \\\\ %\\footnotesize  Aschenbach, Bernd & \\footnotesize   2002-03-13 \\\\ \\footnotesize 16 52 58.9 & \\footnotesize  +02 24 00 & \\footnotesize  20.7  & \\footnotesize  27.2  & \\footnotesize  0.549 & \\footnotesize  0147420401  & \\footnotesize  14115 & \\footnotesize  7003 & \\footnotesize  ok & \\footnotesize  NGC 6240 \\\\ %\\footnotesize  Netzer, Hagai & \\footnotesize   2003-08-13 \\\\ \\footnotesize 16 52 58.8 & \\footnotesize  +02 23 53 & \\footnotesize  20.7  & \\footnotesize  27.2  & \\footnotesize  0.549 & \\footnotesize  0101640101  & \\footnotesize  30111 & \\footnotesize  10396 & \\footnotesize  ok & \\footnotesize  NGC 6240 \\\\ %\\footnotesize  Aschenbach, Bernd & \\footnotesize   2000-09-22 \\\\ \\footnotesize 16 52 58.8 & \\footnotesize  +02 24 07 & \\footnotesize  20.7  & \\footnotesize  27.2  & \\footnotesize  0.549 & \\footnotesize  0147420201  & \\footnotesize  31638 & \\footnotesize  4342 & \\footnotesize  flared & \\footnotesize  NGC 6240 \\\\ %\\footnotesize  Netzer, Hagai & \\footnotesize   2003-03-15 \\\\ \\footnotesize 16 52 58.8 & \\footnotesize  +02 24 05 & \\footnotesize  20.7  & \\footnotesize  27.2  & \\footnotesize  0.549 & \\footnotesize  0147420301  & \\footnotesize  28077 & \\footnotesize  1663 & \\footnotesize  flared & \\footnotesize  NGC 6240 \\\\ %\\footnotesize  Netzer, Hagai & \\footnotesize   2003-03-19 \\\\ \\footnotesize 16 52 58.7 & \\footnotesize  +02 23 56 & \\footnotesize  20.7  & \\footnotesize  27.2  & \\footnotesize  0.549 & \\footnotesize  0147420501  & \\footnotesize  31219 & \\footnotesize  1542 & \\footnotesize  flared & \\footnotesize  NGC 6240 \\\\ %\\footnotesize  Netzer, Hagai & \\footnotesize   2003-08-21 \\\\ \\footnotesize 16 54 00.2 & \\footnotesize  +14 17 47 & \\footnotesize  33.1  & \\footnotesize  32.3  & \\footnotesize  0.551 & \\footnotesize  0113070101  & \\footnotesize  24532 & \\footnotesize  15159 & \\footnotesize  ok & \\footnotesize  nps\\_pos1 \\\\ %\\footnotesize  Aschenbach, Bernd & \\footnotesize   2001-09-14 \\\\ \\footnotesize 17 26 05.4 & \\footnotesize  +02 19 41 & \\footnotesize  25.0  & \\footnotesize  20.0  & \\footnotesize  0.842 & \\footnotesize  0067340401  & \\footnotesize  15280 & \\footnotesize  10474 & \\footnotesize  ok & \\footnotesize  Field IV \\\\ %\\footnotesize  Willingale, Richard & \\footnotesize   2001-02-28 \\\\ \\footnotesize 18 19 44.6 & \\footnotesize  -63 45 31 & \\footnotesize  330.8 & \\footnotesize  -20.7 & \\footnotesize  0.775 & \\footnotesize  0146340601  & \\footnotesize  18216 & \\footnotesize  5958 & \\footnotesize  ok & \\footnotesize  PKS 1814-637 \\\\ %\\footnotesize  Salvati, Marco & \\footnotesize   2003-04-16 \\\\ \\footnotesize 18 19 43.7 & \\footnotesize  -63 45 14 & \\footnotesize  330.8 & \\footnotesize  -20.7 & \\footnotesize  0.775 & \\footnotesize  0146340501  & \\footnotesize  33855 & \\footnotesize  2365 & \\footnotesize  flared & \\footnotesize  PKS 1814-637 \\\\ %\\footnotesize  Salvati, Marco & \\footnotesize   2003-03-23 \\\\ \\footnotesize 18 47 02.3 & \\footnotesize  -78 31 54 & \\footnotesize  315.7 & \\footnotesize  -26.3 & \\footnotesize  1.008 & \\footnotesize  0147920401  & \\footnotesize  10851 & \\footnotesize  261 & \\footnotesize  flared & \\footnotesize  H1846-786 \\\\ %\\footnotesize  Schartel, Norbert & \\footnotesize   2003-04-14 \\\\ \\footnotesize 19 21 13.9 & \\footnotesize  -58 40 17 & \\footnotesize  338.1 & \\footnotesize  -26.7 & \\footnotesize  0.495 & \\footnotesize  0101040501  & \\footnotesize  55489 & \\footnotesize  14130 & \\footnotesize  ok & \\footnotesize  ESO 141-G55 \\\\ %\\footnotesize  Jansen, Fred & \\footnotesize   2001-10-10 \\\\ \\footnotesize 19 31 19.4 & \\footnotesize  -72 39 09 & \\footnotesize  322.5 & \\footnotesize  -28.7 & \\footnotesize  0.597 & \\footnotesize  0081341001  & \\footnotesize  23322 & \\footnotesize  13874 & \\footnotesize  ok & \\footnotesize  IRAS19254-72 \\\\ %\\footnotesize  Franceschini, Alberto & \\footnotesize   2001-03-30 \\\\ \\footnotesize 19 35 13.0 & \\footnotesize  -52 51 43 & \\footnotesize  344.9 & \\footnotesize  -27.7 & \\footnotesize  0.371 & \\footnotesize  0051610101  & \\footnotesize  22624 & \\footnotesize  14980 & \\footnotesize  ok & \\footnotesize  grb980425 \\\\ %\\footnotesize  Pian, Elena & \\footnotesize   2002-03-28 \\\\ \\footnotesize 19 39 59.3 & \\footnotesize  -30 57 50 & \\footnotesize  8.7   & \\footnotesize  -23.2 & \\footnotesize  0.858 & \\footnotesize  0111300101  & \\footnotesize  26709 & \\footnotesize  20646 & \\footnotesize  ok & \\footnotesize  M 55 \\\\ %\\footnotesize  Watson, Michael & \\footnotesize   2001-10-18 \\\\ \\footnotesize 20 03 23.7 & \\footnotesize  -32 51 32 & \\footnotesize  8.4   & \\footnotesize  -28.5 & \\footnotesize  0.749 & \\footnotesize  0104860601  & \\footnotesize  30089 & \\footnotesize  14036 & \\footnotesize  ok & \\footnotesize  [HB89] 2000-330 \\\\ %\\footnotesize  Mason, Keith & \\footnotesize   2002-04-14 \\\\ \\footnotesize 20 07 47.8 & \\footnotesize  -11 09 16 & \\footnotesize  31.3  & \\footnotesize  -21.9 & \\footnotesize  0.675 & \\footnotesize  0044350501  & \\footnotesize  14921 & \\footnotesize  10438 & \\footnotesize  ok & \\footnotesize  IRAS20051-1117 \\\\ %\\footnotesize  Georgantopou, Ioannis & \\footnotesize   2002-10-21 \\\\ \\footnotesize 20 07 50.9 & \\footnotesize  -11 08 20 & \\footnotesize  31.3  & \\footnotesize  -21.9 & \\footnotesize  0.675 & \\footnotesize  0044350201  & \\footnotesize  21638 & \\footnotesize  4395 & \\footnotesize  flared & \\footnotesize  IRAS20051-1117 \\\\ %\\footnotesize  Georgantopou, Ioannis & \\footnotesize   2002-04-20 \\\\ \\footnotesize 20 10 57.8 & \\footnotesize  -56 48 12 & \\footnotesize  340.8 & \\footnotesize  -33.1 & \\footnotesize  0.459 & \\footnotesize  0105260501  & \\footnotesize  19503 & \\footnotesize  11975 & \\footnotesize  discard & \\footnotesize  A3667\\_f5 \\\\ %\\footnotesize  Aschenbach, Bernd & \\footnotesize   2000-10-04 \\\\ \\footnotesize 20 11 09.7 & \\footnotesize  -56 38 40 & \\footnotesize  341.0 & \\footnotesize  -33.2 & \\footnotesize  0.459 & \\footnotesize  0105260201  & \\footnotesize  20146 & \\footnotesize  16130 & \\footnotesize  ok & \\footnotesize  A3667\\_f2 \\\\ %\\footnotesize  Aschenbach, Bernd & \\footnotesize   2000-10-13 \\\\ \\footnotesize 20 11 21.8 & \\footnotesize  -57 25 34 & \\footnotesize  340.1 & \\footnotesize  -33.2 & \\footnotesize  0.333 & \\footnotesize  0042341101  & \\footnotesize  13916 & \\footnotesize  3615 & \\footnotesize  flared & \\footnotesize  RXCJ2011.3-5725 \\\\ %\\footnotesize  Boehringer, Hans & \\footnotesize   2001-04-21 \\\\ \\footnotesize 20 11 54.3 & \\footnotesize  -56 54 54 & \\footnotesize  340.7 & \\footnotesize  -33.3 & \\footnotesize  0.459 & \\footnotesize  0105260601  & \\footnotesize  26467 & \\footnotesize  18472 & \\footnotesize  discard & \\footnotesize  A3667\\_f6 \\\\ %\\footnotesize  Aschenbach, Bernd & \\footnotesize   2000-10-02 \\\\ \\footnotesize 20 12 13.7 & \\footnotesize  -56 37 25 & \\footnotesize  341.1 & \\footnotesize  -33.3 & \\footnotesize  0.459 & \\footnotesize  0105260401  & \\footnotesize  17469 & \\footnotesize  13205 & \\footnotesize  ok & \\footnotesize  A3667\\_f4 \\\\ %\\footnotesize  Aschenbach, Bernd & \\footnotesize   2000-10-03 \\\\ \\footnotesize 20 13 04.4 & \\footnotesize  -56 53 11 & \\footnotesize  340.7 & \\footnotesize  -33.4 & \\footnotesize  0.459 & \\footnotesize  0105260101  & \\footnotesize  21303 & \\footnotesize  7504 & \\footnotesize  discard & \\footnotesize  A3667\\_f1 \\\\ %\\footnotesize  Aschenbach, Bernd & \\footnotesize   2000-09-21 \\\\ \\footnotesize 20 13 08.2 & \\footnotesize  -56 43 17 & \\footnotesize  340.9 & \\footnotesize  -33.4 & \\footnotesize  0.459 & \\footnotesize  0105260301  & \\footnotesize  17871 & \\footnotesize  13421 & \\footnotesize  ok & \\footnotesize  A3667\\_f3 \\\\ %\\footnotesize  Aschenbach, Bernd & \\footnotesize   2000-10-04 \\\\ \\footnotesize 20 13 29.0 & \\footnotesize  -41 47 22 & \\footnotesize  358.7 & \\footnotesize  -32.4 & \\footnotesize  0.498 & \\footnotesize  0081340501  & \\footnotesize  21913 & \\footnotesize  10696 & \\footnotesize  ok & \\footnotesize  IRAS20100-41 \\\\ %\\footnotesize  Franceschini, Alberto & \\footnotesize   2001-04-21 \\\\ \\footnotesize 20 13 59.7 & \\footnotesize  -87 57 32 & \\footnotesize  305.0 & \\footnotesize  -27.8 & \\footnotesize  1.141 & \\footnotesize  0083210701  & \\footnotesize  10779 & \\footnotesize  5646 & \\footnotesize  ok & \\footnotesize  1SAXJ2011.9-8758 \\\\ %\\footnotesize  Perri, Matteo & \\footnotesize   2002-05-06 \\\\ \\footnotesize 20 18 18.1 & \\footnotesize  -70 51 15 & \\footnotesize  324.1 & \\footnotesize  -32.5 & \\footnotesize  0.497 & \\footnotesize  0022340101  & \\footnotesize  32858 & \\footnotesize  18375 & \\footnotesize  ok & \\footnotesize  Pavo group \\\\ %\\footnotesize  Davis, David & \\footnotesize   2002-04-01 \\\\ \\footnotesize 20 40 10.0 & \\footnotesize  -00 52 23 & \\footnotesize  45.2  & \\footnotesize  -24.4 & \\footnotesize  0.703 & \\footnotesize  0111180201  & \\footnotesize  17270 & \\footnotesize  10630 & \\footnotesize  ok & \\footnotesize  AE Aqr \\\\ %\\footnotesize  Watson, Michael & \\footnotesize   2001-11-08 \\\\ \\footnotesize 20 46 20.5 & \\footnotesize  -02 48 54 & \\footnotesize  44.2  & \\footnotesize  -26.7 & \\footnotesize  0.586 & \\footnotesize  0112600501  & \\footnotesize  11360 & \\footnotesize  8305 & \\footnotesize  ok & \\footnotesize  Mrk 896 \\\\ %\\footnotesize  Turner, Martin & \\footnotesize   2001-11-16 \\\\ \\footnotesize 20 48 14.0 & \\footnotesize  -17 49 51 & \\footnotesize  28.7  & \\footnotesize  -33.5 & \\footnotesize  0.478 & \\footnotesize  0201902401  & \\footnotesize  28257 & \\footnotesize  21534 & \\footnotesize  ok & \\footnotesize  RXCJ2048.1-1750 \\\\ %\\footnotesize  Boehringer, Hans & \\footnotesize   2004-05-14 \\\\ \\footnotesize 20 54 18.0 & \\footnotesize  -15 55 36 & \\footnotesize  31.5  & \\footnotesize  -34.2 & \\footnotesize  0.436 & \\footnotesize  0083210101  & \\footnotesize  11462 & \\footnotesize  7154 & \\footnotesize  ok & \\footnotesize  1SAXJ2054.3-1556 \\\\ %\\footnotesize  Perri, Matteo & \\footnotesize   2001-05-03 \\\\ \\footnotesize 20 56 22.3 & \\footnotesize  -04 37 59 & \\footnotesize  43.8  & \\footnotesize  -29.7 & \\footnotesize  0.502 & \\footnotesize  0112190601  & \\footnotesize  17511 & \\footnotesize  12272 & \\footnotesize  ok & \\footnotesize  MS2053.7-0449 \\\\ %\\footnotesize  Mason, Keith & \\footnotesize   2001-11-14 \\\\ \\footnotesize 20 58 26.1 & \\footnotesize  -42 38 54 & \\footnotesize  358.6 & \\footnotesize  -40.7 & \\footnotesize  0.382 & \\footnotesize  0081340401  & \\footnotesize  21863 & \\footnotesize  12027 & \\footnotesize  discard & \\footnotesize  IRAS20551-42 \\\\ %& \\footnotesize  Franceschini, Alberto & \\footnotesize   2001-04-21 \\\\ \\footnotesize 21 04 10.3 & \\footnotesize  -11 21 39 & \\footnotesize  37.7  & \\footnotesize  -34.5 & \\footnotesize  0.466 & \\footnotesize  0041150101  & \\footnotesize  40565 & \\footnotesize  30058 & \\footnotesize  ok & \\footnotesize  NGC 7009 \\\\ %\\footnotesize  Chu, You-Hua & \\footnotesize   2001-04-30 \\\\ \\footnotesize 21 04 39.5 & \\footnotesize  -12 20 03 & \\footnotesize  36.7  & \\footnotesize  -35.0 & \\footnotesize  0.465 & \\footnotesize  0038540301  & \\footnotesize  17217 & \\footnotesize  12791 & \\footnotesize  ok & \\footnotesize  NGC7010 \\\\ %\\footnotesize  Zabludoff, Ann & \\footnotesize   2001-05-01 \\\\ \\footnotesize 21 14 55.7 & \\footnotesize  +06 08 31 & \\footnotesize  57.0  & \\footnotesize  -28.0 & \\footnotesize  0.663 & \\footnotesize  0150610201  & \\footnotesize  16914 & \\footnotesize  7408 & \\footnotesize  ok & \\footnotesize  PG 2112+059 \\\\ %\\footnotesize  Schartel, Norbert & \\footnotesize   2003-05-14 \\\\ \\footnotesize 21 24 43.5 & \\footnotesize  -33 58 34 & \\footnotesize  10.9  & \\footnotesize  -45.4 & \\footnotesize  0.572 & \\footnotesize  0112320601  & \\footnotesize  70315 & \\footnotesize  20215 & \\footnotesize  ok & \\footnotesize  PSR J2124-3358 \\\\ %\\footnotesize  Aschenbach, Bernd & \\footnotesize   2002-04-15 \\\\ \\footnotesize 21 27 37.2 & \\footnotesize  -44 48 41 & \\footnotesize  355.3 & \\footnotesize  -45.9 & \\footnotesize  0.345 & \\footnotesize  0088020201  & \\footnotesize  32370 & \\footnotesize  10164 & \\footnotesize  ok & \\footnotesize  IRAS F21243-4501 \\\\ %\\footnotesize  DiCocco, Guido & \\footnotesize   2001-04-29 \\\\ \\footnotesize 21 29 11.5 & \\footnotesize  -15 38 33 & \\footnotesize  35.9  & \\footnotesize  -41.8 & \\footnotesize  0.493 & \\footnotesize  0103060101  & \\footnotesize  23564 & \\footnotesize  15955 & \\footnotesize  ok & \\footnotesize  PKS 2126-158 \\\\ %\\footnotesize  Aschenbach, Bernd & \\footnotesize   2001-05-01 \\\\ \\footnotesize 21 29 34.7 & \\footnotesize  +00 04 53 & \\footnotesize  53.6  & \\footnotesize  -34.4 & \\footnotesize  0.422 & \\footnotesize  0093030201  & \\footnotesize  58916 & \\footnotesize  34048 & \\footnotesize  ok & \\footnotesize  RX J2129.6+0005 \\\\ %\\footnotesize  Hughes, John & \\footnotesize   2002-10-29 \\\\ \\footnotesize 21 32 29.9 & \\footnotesize  +10 09 18 & \\footnotesize  63.6  & \\footnotesize  -29.0 & \\footnotesize  0.465 & \\footnotesize  0150470701  & \\footnotesize  37917 & \\footnotesize  12325 & \\footnotesize  ok & \\footnotesize  UGC 11763 \\\\ %\\footnotesize  Santos-Lleo, Maria & \\footnotesize   2003-05-17 \\\\ \\footnotesize 21 37 45.1 & \\footnotesize  -14 32 51 & \\footnotesize  38.4  & \\footnotesize  -43.3 & \\footnotesize  0.477 & \\footnotesize  0092850201  & \\footnotesize  59850 & \\footnotesize  28012 & \\footnotesize  ok & \\footnotesize  PKS 2135-147 \\\\ %\\footnotesize  Canizares, Claude & \\footnotesize   2001-04-29 \\\\ \\footnotesize 21 40 14.7 & \\footnotesize  -23 39 31 & \\footnotesize  26.4  & \\footnotesize  -46.9 & \\footnotesize  0.356 & \\footnotesize  0008830101  & \\footnotesize  22769 & \\footnotesize  9500 & \\footnotesize  ok & \\footnotesize  ms2137-23 \\\\ %\\footnotesize  Allen, Steve & \\footnotesize   2001-04-29 \\\\ \\footnotesize 21 51 55.1 & \\footnotesize  -30 27 42 & \\footnotesize  17.0  & \\footnotesize  -50.7 & \\footnotesize  0.191 & \\footnotesize  0103060401  & \\footnotesize  25028 & \\footnotesize  19943 & \\footnotesize  ok & \\footnotesize  PKS 2149-306 \\\\ %\\footnotesize  Aschenbach, Bernd & \\footnotesize   2001-05-01 \\\\ \\footnotesize 21 52 01.3 & \\footnotesize  -27 31 59 & \\footnotesize  21.6  & \\footnotesize  -50.3 & \\footnotesize  0.236 & \\footnotesize  0062940401  & \\footnotesize  39421 & \\footnotesize  17220 & \\footnotesize  ok & \\footnotesize  HE2149-2745 \\\\ %\\footnotesize  CHARTAS, GEORGE & \\footnotesize   2001-11-08 \\\\ \\footnotesize 21 53 36.5 & \\footnotesize  +17 41 56 & \\footnotesize  73.9  & \\footnotesize  -27.8 & \\footnotesize  0.669 & \\footnotesize  0111270101  & \\footnotesize  23263 & \\footnotesize  9766 & \\footnotesize  ok & \\footnotesize  A2390 \\\\ %\\footnotesize  Watson, Michael & \\footnotesize   2001-06-19 \\\\ \\footnotesize 22 01 49.7 & \\footnotesize  -59 57 41 & \\footnotesize  332.2 & \\footnotesize  -46.3 & \\footnotesize  0.284 & \\footnotesize  0149670101  & \\footnotesize  25437 & \\footnotesize  19965 & \\footnotesize  ok & \\footnotesize  Abell 3827 \\\\ %\\footnotesize  Gomez, Percy & \\footnotesize   2002-11-16 \\\\ \\footnotesize 22 02 01.8 & \\footnotesize  -31 52 13 & \\footnotesize  15.1  & \\footnotesize  -53.0 & \\footnotesize  0.165 & \\footnotesize  0147920601  & \\footnotesize  16858 & \\footnotesize  11831 & \\footnotesize  ok & \\footnotesize  NGC 7172 \\\\ %\\footnotesize  Schartel, Norbert & \\footnotesize   2002-11-19 \\\\ \\footnotesize 22 03 10.6 & \\footnotesize  +18 53 03 & \\footnotesize  76.7  & \\footnotesize  -28.5 & \\footnotesize  0.572 & \\footnotesize  0130920101  & \\footnotesize  20217 & \\footnotesize  13922 & \\footnotesize  ok & \\footnotesize  HD 209458 \\\\ %\\footnotesize  Jansen, Fred & \\footnotesize   2000-11-30 \\\\ \\footnotesize 22 04 30.7 & \\footnotesize  -64 44 31 & \\footnotesize  326.2 & \\footnotesize  -44.0 & \\footnotesize  0.291 & \\footnotesize  0143250101  & \\footnotesize  22216 & \\footnotesize  17825 & \\footnotesize  ok & \\footnotesize  HD 209295 \\\\ %\\footnotesize  James, David & \\footnotesize   2002-11-16 \\\\ \\footnotesize 22 05 09.3 & \\footnotesize  -01 55 11 & \\footnotesize  58.0  & \\footnotesize  -42.9 & \\footnotesize  0.607 & \\footnotesize  0012440301  & \\footnotesize  35366 & \\footnotesize  24897 & \\footnotesize  ok & \\footnotesize  PB5062 \\\\ %\\footnotesize  Elvis, Martin & \\footnotesize   2001-05-24 \\\\ \\footnotesize 22 09 15.9 & \\footnotesize  -47 10 00 & \\footnotesize  349.5 & \\footnotesize  -52.5 & \\footnotesize  0.194 & \\footnotesize  0111810101  & \\footnotesize  49610 & \\footnotesize  18189 & \\footnotesize  ok & \\footnotesize  NGC7213 \\\\ %\\footnotesize  Kaastra, Jelle & \\footnotesize   2001-05-29 \\\\ \\footnotesize 22 15 31.8 & \\footnotesize  -17 44 09 & \\footnotesize  39.2  & \\footnotesize  -52.9 & \\footnotesize  0.228 & \\footnotesize  0106660501  & \\footnotesize  11568 & \\footnotesize  6085 & \\footnotesize  ok & \\footnotesize  LBQS 2212-1759 \\\\ %\\footnotesize  Clavel, Jean & \\footnotesize   2001-11-17 \\\\ \\footnotesize 22 15 32.0 & \\footnotesize  -17 44 06 & \\footnotesize  39.2  & \\footnotesize  -52.9 & \\footnotesize  0.228 & \\footnotesize  0106660101  & \\footnotesize  60508 & \\footnotesize  51704 & \\footnotesize  ok & \\footnotesize  LBQS 2212-1759 \\\\ %\\footnotesize  Clavel, Jean & \\footnotesize   2000-11-18 \\\\ \\footnotesize 22 15 31.9 & \\footnotesize  -17 44 09 & \\footnotesize  39.2  & \\footnotesize  -52.9 & \\footnotesize  0.228 & \\footnotesize  0106660201  & \\footnotesize  53769 & \\footnotesize  37107 & \\footnotesize  ok & \\footnotesize  LBQS 2212-1759 \\\\ %\\footnotesize  Clavel, Jean & \\footnotesize   2000-11-19 \\\\ \\footnotesize 22 15 31.3 & \\footnotesize  -17 44 08 & \\footnotesize  39.2  & \\footnotesize  -52.9 & \\footnotesize  0.228 & \\footnotesize  0106660601  & \\footnotesize  110168 & \\footnotesize  82222 & \\footnotesize  ok & \\footnotesize  LBQS 2212-1759 \\\\ %\\footnotesize  Clavel, Jean & \\footnotesize   2001-11-19 \\\\ \\footnotesize 22 15 31.9 & \\footnotesize  -17 44 08 & \\footnotesize  39.2  & \\footnotesize  -52.9 & \\footnotesize  0.228 & \\footnotesize  0106660401  & \\footnotesize  35114 & \\footnotesize  12805 & \\footnotesize  ok & \\footnotesize  LBQS 2212-1759 \\\\ %\\footnotesize  Clavel, Jean & \\footnotesize   2001-11-17 \\\\ \\footnotesize 22 17 04.5 & \\footnotesize  -16 38 35 & \\footnotesize  41.1  & \\footnotesize  -52.8 & \\footnotesize  0.261 & \\footnotesize  0148060201  & \\footnotesize  12800 & \\footnotesize  2657 & \\footnotesize flared & \\footnotesize BPS\\,CS\\,22892-05 \\\\ %& \\footnotesize  Schlegel, Eric & \\footnotesize   2003-05-15 \\\\ \\footnotesize 22 17 15.4 & \\footnotesize  +14 15 03 & \\footnotesize  76.0  & \\footnotesize  -34.2 & \\footnotesize  0.493 & \\footnotesize  0103660301  & \\footnotesize  47294 & \\footnotesize  10301 & \\footnotesize  ok & \\footnotesize  Mark 304 \\\\ %\\footnotesize  Aschenbach, Bernd & \\footnotesize   2002-05-13 \\\\ \\footnotesize 22 17 31.1 & \\footnotesize  +00 14 15 & \\footnotesize  63.0  & \\footnotesize  -44.0 & \\footnotesize  0.459 & \\footnotesize  0094310101  & \\footnotesize  68821 & \\footnotesize  58820 & \\footnotesize  ok & \\footnotesize  SSA22 \\\\ %\\footnotesize  Almaini, Omar & \\footnotesize   2002-11-18 \\\\ \\footnotesize 22 17 31.1 & \\footnotesize  +00 14 12 & \\footnotesize  63.0  & \\footnotesize  -44.0 & \\footnotesize  0.459 & \\footnotesize  0094310201  & \\footnotesize  81270 & \\footnotesize  60699 & \\footnotesize  ok & \\footnotesize  SSA22 \\\\ %\\footnotesize  Almaini, Omar & \\footnotesize   2002-12-16 \\\\ \\footnotesize 22 19 18.4 & \\footnotesize  +12 08 05 & \\footnotesize  74.7  & \\footnotesize  -36.1 & \\footnotesize  0.536 & \\footnotesize  0103861201  & \\footnotesize  13014 & \\footnotesize  8526 & \\footnotesize  ok & \\footnotesize  IIZw 177 \\\\ %\\footnotesize  Aschenbach, Bernd & \\footnotesize   2001-06-07 \\\\ \\footnotesize 22 20 45.8 & \\footnotesize  -24 40 57 & \\footnotesize  28.4  & \\footnotesize  -56.1 & \\footnotesize  0.199 & \\footnotesize  0049340201  & \\footnotesize  28359 & \\footnotesize  22176 & \\footnotesize  ok & \\footnotesize  NGC 7252 \\\\ %\\footnotesize  Ponman, Trevor & \\footnotesize   2001-11-14 \\\\ \\footnotesize 22 28 29.7 & \\footnotesize  -05 18 46 & \\footnotesize  59.1  & \\footnotesize  -49.6 & \\footnotesize  0.514 & \\footnotesize  0100440101  & \\footnotesize  46681 & \\footnotesize  40544 & \\footnotesize  ok & \\footnotesize  PHL 5200 \\\\ %\\footnotesize  Aschenbach, Bernd & \\footnotesize   2001-05-28 \\\\ \\footnotesize 22 29 44.9 & \\footnotesize  -20 50 16 & \\footnotesize  36.1  & \\footnotesize  -57.1 & \\footnotesize  0.268 & \\footnotesize  0125911001  & \\footnotesize  16692 & \\footnotesize  12301 & \\footnotesize  ok & \\footnotesize  NGC 7293 \\\\ %\\footnotesize  Jansen, Fred & \\footnotesize   2002-11-26 \\\\ \\footnotesize 22 35 45.7 & \\footnotesize  -26 02 55 & \\footnotesize  27.1  & \\footnotesize  -59.7 & \\footnotesize  0.147 & \\footnotesize  0111790101  & \\footnotesize  44663 & \\footnotesize  13638 & \\footnotesize  ok & \\footnotesize  NGC7314 \\\\ %\\footnotesize  Kaastra, Jelle & \\footnotesize   2001-05-02 \\\\ \\footnotesize 22 36 10.8 & \\footnotesize  +13 44 45 & \\footnotesize  79.9  & \\footnotesize  -37.5 & \\footnotesize  0.474 & \\footnotesize  0153220601  & \\footnotesize  12915 & \\footnotesize  5617 & \\footnotesize  ok & \\footnotesize  PG 2233+134 \\\\ %\\footnotesize  Matsumoto, Chiho & \\footnotesize   2003-05-28 \\\\ \\footnotesize 22 36 55.4 & \\footnotesize  -22 13 04 & \\footnotesize  34.6  & \\footnotesize  -59.1 & \\footnotesize  0.216 & \\footnotesize  0103860201  & \\footnotesize  11864 & \\footnotesize  5675 & \\footnotesize  ok & \\footnotesize  ESO 602- G 031 \\\\ %\\footnotesize  Aschenbach, Bernd & \\footnotesize   2001-05-25 \\\\ \\footnotesize 22 37 00.5 & \\footnotesize  -15 16 18 & \\footnotesize  46.8  & \\footnotesize  -56.6 & \\footnotesize  0.400 & \\footnotesize  0056021601  & \\footnotesize  25111 & \\footnotesize  15633 & \\footnotesize  ok & \\footnotesize  RX J2237.0-1516 \\\\ %\\footnotesize  ARNAUD, Monique & \\footnotesize   2001-11-28 \\\\ \\footnotesize 22 39 56.5 & \\footnotesize  -05 51 30 & \\footnotesize  61.1  & \\footnotesize  -52.1 & \\footnotesize  0.402 & \\footnotesize  0149410401  & \\footnotesize  34475 & \\footnotesize  19256 & \\footnotesize  ok & \\footnotesize  BR 2237-0607 \\\\ %\\footnotesize  Mathur, Smita & \\footnotesize   2003-05-17 \\\\ \\footnotesize 22 40 32.9 & \\footnotesize  +03 22 16 & \\footnotesize  71.8  & \\footnotesize  -46.1 & \\footnotesize  0.524 & \\footnotesize  0110960101  & \\footnotesize  42870 & \\footnotesize  12551 & \\footnotesize  ok & \\footnotesize  Q2237+0305 \\\\ %\\footnotesize  Watson, Michael & \\footnotesize   2002-05-29 \\\\ \\footnotesize 22 41 59.1 & \\footnotesize  -44 04 00 & \\footnotesize  351.9 & \\footnotesize  -59.0 & \\footnotesize  0.182 & \\footnotesize  0153220101  & \\footnotesize  10919 & \\footnotesize  3344 & \\footnotesize  flared & \\footnotesize  RX J2241.8-4405 \\\\ %\\footnotesize  Matsumoto, Chiho & \\footnotesize   2003-05-17 \\\\ \\footnotesize 22 44 19.2 & \\footnotesize  -72 42 47 & \\footnotesize  314.8 & \\footnotesize  -41.3 & \\footnotesize  0.409 & \\footnotesize  0150050101  & \\footnotesize  59915 & \\footnotesize  20302 & \\footnotesize  ok & \\footnotesize  Galactic Halo \\\\ %\\footnotesize  Shelton, Robin & \\footnotesize   2002-11-27 \\\\ \\footnotesize 22 49 39.9 & \\footnotesize  -27 06 51 & \\footnotesize  25.8  & \\footnotesize  -62.9 & \\footnotesize  0.178 & \\footnotesize  0111970101  & \\footnotesize  13358 & \\footnotesize  10382 & \\footnotesize  ok & \\footnotesize  TY PsA \\\\ %\\footnotesize  Mason, Keith & \\footnotesize   2001-11-28 \\\\ \\footnotesize 22 49 50.0 & \\footnotesize  -64 23 08 & \\footnotesize  322.0 & \\footnotesize  -47.9 & \\footnotesize  0.280 & \\footnotesize  0112240101  & \\footnotesize  35013 & \\footnotesize  25781 & \\footnotesize  ok & \\footnotesize  A3921 \\\\ %\\footnotesize  Turner, Martin & \\footnotesize   2000-10-26 \\\\ \\footnotesize 22 51 48.7 & \\footnotesize  -17 52 16 & \\footnotesize  45.1  & \\footnotesize  -60.9 & \\footnotesize  0.269 & \\footnotesize  0081340901  & \\footnotesize  23365 & \\footnotesize  19561 & \\footnotesize  ok & \\footnotesize  IRAS22491-18 \\\\ %\\footnotesize  Franceschini, Alberto & \\footnotesize   2001-05-25 \\\\ \\footnotesize 22 57 09.8 & \\footnotesize  -36 27 33 & \\footnotesize  4.6   & \\footnotesize  -64.1 & \\footnotesize  0.115 & \\footnotesize  0135980201  & \\footnotesize  32903 & \\footnotesize  25360 & \\footnotesize  ok & \\footnotesize  IC1459 \\\\ %\\footnotesize  Fukazawa, Yasushi & \\footnotesize   2002-05-01 \\\\ \\footnotesize 23 02 28.9 & \\footnotesize  +08 36 41 & \\footnotesize  82.6  & \\footnotesize  -45.5 & \\footnotesize  0.487 & \\footnotesize  0032140201  & \\footnotesize  13275 & \\footnotesize  6836 & \\footnotesize  ok & \\footnotesize  1saxJ2302.5+0836 \\\\ %\\footnotesize  Fiore, Fabrizio & \\footnotesize   2001-06-10 \\\\ \\footnotesize 23 03 15.9 & \\footnotesize  +08 52 22 & \\footnotesize  83.0  & \\footnotesize  -45.4 & \\footnotesize  0.529 & \\footnotesize  0112170301  & \\footnotesize  24590 & \\footnotesize  8989 & \\footnotesize  ok & \\footnotesize  NGC7469 \\\\ %\\footnotesize  Kaastra, Jelle & \\footnotesize   2000-12-26 \\\\ \\footnotesize 23 04 45.6 & \\footnotesize  +03 11 34 & \\footnotesize  78.4  & \\footnotesize  -50.2 & \\footnotesize  0.527 & \\footnotesize  0033541001  & \\footnotesize  13307 & \\footnotesize  10070 & \\footnotesize  ok & \\footnotesize  pg 2302+029 \\\\ %\\footnotesize  Andrea, Comastri & \\footnotesize   2001-11-30 \\\\ \\footnotesize 23 04 56.1 & \\footnotesize  +12 19 30 & \\footnotesize  86.2  & \\footnotesize  -42.8 & \\footnotesize  0.519 & \\footnotesize  0025541001  & \\footnotesize  13263 & \\footnotesize  6410 & \\footnotesize  ok & \\footnotesize  NGC7479 \\\\ %\\footnotesize  Iwasawa, Kazushi & \\footnotesize   2001-06-19 \\\\ \\footnotesize 23 08 22.4 & \\footnotesize  -02 11 19 & \\footnotesize  73.8  & \\footnotesize  -54.9 & \\footnotesize  0.435 & \\footnotesize  0042341201  & \\footnotesize  12596 & \\footnotesize  1651 & \\footnotesize  flared & \\footnotesize  RXCJ2308.3-0211 \\\\ %\\footnotesize  Boehringer, Hans & \\footnotesize   2001-06-10 \\\\ \\footnotesize 23 13 58.5 & \\footnotesize  -42 43 26 & \\footnotesize  348.3 & \\footnotesize  -64.8 & \\footnotesize  0.185 & \\footnotesize  0123900101  & \\footnotesize  66664 & \\footnotesize  29998 & \\footnotesize  ok & \\footnotesize  A S 1101 \\\\ %\\footnotesize  Kaastra, Jelle & \\footnotesize   2000-05-11 \\\\ \\footnotesize 23 13 59.1 & \\footnotesize  -42 43 43 & \\footnotesize  348.3 & \\footnotesize  -64.8 & \\footnotesize  0.185 & \\footnotesize  0147800101  & \\footnotesize  126314 & \\footnotesize  83040 & \\footnotesize  ok & \\footnotesize  Sersic 159-03 \\\\ %\\footnotesize  Kaastra, Jelle & \\footnotesize   2002-11-21 \\\\ \\footnotesize 23 15 42.4 & \\footnotesize  -59 04 11 & \\footnotesize  323.7 & \\footnotesize  -54.0 & \\footnotesize  0.243 & \\footnotesize  0081340301  & \\footnotesize  23413 & \\footnotesize  8782 & \\footnotesize  ok & \\footnotesize  IRAS23128-59 \\\\ %\\footnotesize  Franceschini, Alberto & \\footnotesize   2002-11-19 \\\\ \\footnotesize 23 16 10.7 & \\footnotesize  -42 34 57 & \\footnotesize  348.1 & \\footnotesize  -65.2 & \\footnotesize  0.185 & \\footnotesize  0093640701  & \\footnotesize  20760 & \\footnotesize  7748 & \\footnotesize  ok & \\footnotesize  NGC 7552 \\\\ %\\footnotesize  Bauer, Franz & \\footnotesize   2001-05-30 \\\\ \\footnotesize 23 18 15.2 & \\footnotesize  +00 16 14 & \\footnotesize  79.8  & \\footnotesize  -54.6 & \\footnotesize  0.411 & \\footnotesize  0066950301  & \\footnotesize  12208 & \\footnotesize  2387 & \\footnotesize  flared & \\footnotesize  NGC 7589 \\\\ %\\footnotesize  Mao, Shude & \\footnotesize   2001-06-04 \\\\ \\footnotesize 23 18 16.4 & \\footnotesize  +00 15 50 & \\footnotesize  79.8  & \\footnotesize  -54.6 & \\footnotesize  0.411 & \\footnotesize  0066950401  & \\footnotesize  13324 & \\footnotesize  4524 & \\footnotesize  ok & \\footnotesize  NGC 7589 \\\\ %\\footnotesize  Mao, Shude & \\footnotesize   2001-11-28 \\\\ \\footnotesize 23 18 22.2 & \\footnotesize  -42 22 05 & \\footnotesize  348.0 & \\footnotesize  -65.6 & \\footnotesize  0.196 & \\footnotesize  0112310201  & \\footnotesize  23362 & \\footnotesize  18843 & \\footnotesize  ok & \\footnotesize  NGC 7582 \\\\ %\\footnotesize  Turner, Martin & \\footnotesize   2001-05-25 \\\\ \\footnotesize 23 19 43.4 & \\footnotesize  -73 12 43 & \\footnotesize  311.6 & \\footnotesize  -42.3 & \\footnotesize  0.191 & \\footnotesize  0201903201  & \\footnotesize  31029 & \\footnotesize  0 & \\footnotesize  flared & \\footnotesize  RXCJ2319.6-7313 \\\\ %\\footnotesize  Boehringer, Hans & \\footnotesize   2004-04-19 \\\\ \\footnotesize 23 19 50.1 & \\footnotesize  -73 12 55 & \\footnotesize  311.6 & \\footnotesize  -42.3 & \\footnotesize  0.191 & \\footnotesize  0201903301  & \\footnotesize  12616 & \\footnotesize  7975 & \\footnotesize  ok & \\footnotesize  RXCJ2319.6-7313 \\\\ %\\footnotesize  Boehringer, Hans & \\footnotesize   2004-05-15 \\\\ \\footnotesize 23 25 17.9 & \\footnotesize  -12 06 23 & \\footnotesize  65.3  & \\footnotesize  -64.8 & \\footnotesize  0.250 & \\footnotesize  0147330101  & \\footnotesize  120513 & \\footnotesize  52958 & \\footnotesize  ok & \\footnotesize  Abell 2597 \\\\ %\\footnotesize  Fabian, Andrew & \\footnotesize   2003-06-28 \\\\ \\footnotesize 23 25 21.4 & \\footnotesize  -12 07 20 & \\footnotesize  65.3  & \\footnotesize  -64.8 & \\footnotesize  0.250 & \\footnotesize  0108460201  & \\footnotesize  21073 & \\footnotesize  12680 & \\footnotesize  ok & \\footnotesize  Abell 2597 \\\\ %\\footnotesize  Mushotzky, Richard & \\footnotesize   2000-11-30 \\\\ \\footnotesize 23 31 49.2 & \\footnotesize  +19 56 29 & \\footnotesize  98.5  & \\footnotesize  -39.1 & \\footnotesize  0.422 & \\footnotesize  0112880301  & \\footnotesize  15772 & \\footnotesize  11809 & \\footnotesize  ok & \\footnotesize  EQ Peg \\\\ %\\footnotesize  Aschenbach, Bernd & \\footnotesize   2000-07-09 \\\\ \\footnotesize 23 31 52.8 & \\footnotesize  +19 38 10 & \\footnotesize  98.4  & \\footnotesize  -39.4 & \\footnotesize  0.422 & \\footnotesize  0032140701  & \\footnotesize  12311 & \\footnotesize  7348 & \\footnotesize  ok & \\footnotesize  1saxj2331.9+19 \\\\ %\\footnotesize  Fiore, Fabrizio & \\footnotesize   2001-07-05 \\\\ \\footnotesize 23 31 57.8 & \\footnotesize  +19 44 16 & \\footnotesize  98.5  & \\footnotesize  -39.3 & \\footnotesize  0.422 & \\footnotesize  0147580301  & \\footnotesize  15371 & \\footnotesize  9780 & \\footnotesize  ok & \\footnotesize  1AXGJ233200+19 \\\\ %\\footnotesize  Uneda, Yoshihiro & \\footnotesize   2002-12-28 \\\\ \\footnotesize 23 33 40.6 & \\footnotesize  -15 17 12 & \\footnotesize  62.2  & \\footnotesize  -68.4 & \\footnotesize  0.196 & \\footnotesize  0093550401  & \\footnotesize  25209 & \\footnotesize  19527 & \\footnotesize  ok & \\footnotesize  LEDA 913766 \\\\ %\\footnotesize  Charles, Phil & \\footnotesize   2000-12-23 \\\\ \\footnotesize 23 36 12.1 & \\footnotesize  +02 08 25 & \\footnotesize  88.1  & \\footnotesize  -55.5 & \\footnotesize  0.493 & \\footnotesize  0112522601  & \\footnotesize  17140 & \\footnotesize  13884 & \\footnotesize  ok & \\footnotesize  NGC 7714 \\\\ %\\footnotesize  Watson, Michael & \\footnotesize   2002-12-08 \\\\ \\footnotesize 23 36 18.1 & \\footnotesize  +02 10 07 & \\footnotesize  88.2  & \\footnotesize  -55.5 & \\footnotesize  0.493 & \\footnotesize  0112521301  & \\footnotesize  22232 & \\footnotesize  13123 & \\footnotesize  ok & \\footnotesize  NGC 7714 \\\\ %\\footnotesize  Watson, Michael & \\footnotesize   2002-06-07 \\\\ \\footnotesize 23 37 41.6 & \\footnotesize  +00 16 24 & \\footnotesize  87.0  & \\footnotesize  -57.3 & \\footnotesize  0.385 & \\footnotesize  0042341301  & \\footnotesize  14115 & \\footnotesize  9841 & \\footnotesize  ok & \\footnotesize  RXCJ2337.6+0016 \\\\ %\\footnotesize  Boehringer, Hans & \\footnotesize   2001-12-06 \\\\ \\footnotesize 23 37 50.0 & \\footnotesize  -56 24 43 & \\footnotesize  322.1 & \\footnotesize  -57.8 & \\footnotesize  0.204 & \\footnotesize  0083210201  & \\footnotesize  10458 & \\footnotesize  7655 & \\footnotesize  ok & \\footnotesize  1SAXJ2337.8-56 \\\\ %\\footnotesize  Perri, Matteo & \\footnotesize   2001-11-29 \\\\ \\footnotesize 23 39 54.9 & \\footnotesize  -12 17 38 & \\footnotesize  70.9  & \\footnotesize  -67.6 & \\footnotesize  0.306 & \\footnotesize  0055990301  & \\footnotesize  14965 & \\footnotesize  10380 & \\footnotesize  ok & \\footnotesize  Arp222 \\\\ %\\footnotesize  Read, Andrew & \\footnotesize   2001-11-28 \\\\ \\footnotesize 23 40 25.9 & \\footnotesize  -11 43 12 & \\footnotesize  72.2  & \\footnotesize  -67.3 & \\footnotesize  0.281 & \\footnotesize  0152200101  & \\footnotesize  53961 & \\footnotesize  45458 & \\footnotesize  ok & \\footnotesize  Abell 2638 \\\\ %\\footnotesize  Ponman, Trevor & \\footnotesize   2002-12-22 \\\\ \\footnotesize 23 47 20.2 & \\footnotesize  -02 18 52 & \\footnotesize  88.4  & \\footnotesize  -60.8 & \\footnotesize  0.362 & \\footnotesize  0152860101  & \\footnotesize  33656 & \\footnotesize  24541 & \\footnotesize  ok & \\footnotesize  HCG 97 \\\\ %\\footnotesize  Dupke, Renato & \\footnotesize   2002-12-29 \\\\ \\footnotesize 23 47 22.3 & \\footnotesize  +00 53 03 & \\footnotesize  91.5  & \\footnotesize  -58.0 & \\footnotesize  0.362 & \\footnotesize  0147580401  & \\footnotesize  15369 & \\footnotesize  13043 & \\footnotesize  ok & \\footnotesize  1AXGJ234725+00 \\\\ %\\footnotesize  Uneda, Yoshihiro & \\footnotesize   2002-12-28 \\\\ \\footnotesize 23 51 42.3 & \\footnotesize  -26 04 13 & \\footnotesize  34.0  & \\footnotesize  -76.6 & \\footnotesize  0.166 & \\footnotesize  0148990101  & \\footnotesize  31161 & \\footnotesize  20013 & \\footnotesize  ok & \\footnotesize  Abell 2667 \\\\ %\\footnotesize  Allen, Steve & \\footnotesize   2003-06-21 \\\\ \\footnotesize 23 54 08.5 & \\footnotesize  -10 23 50 & \\footnotesize  81.3  & \\footnotesize  -68.5 & \\footnotesize  0.292 & \\footnotesize  0108460301  & \\footnotesize  33298 & \\footnotesize  13607 & \\footnotesize  ok & \\footnotesize  Abell 2670 \\\\ %\\footnotesize  Mushotzky, Richard & \\footnotesize   2001-06-20 \\\\ \\footnotesize 23 56 55.8 & \\footnotesize  -24 24 39 & \\footnotesize  42.3  & \\footnotesize  -77.4 & \\footnotesize  0.173 & \\footnotesize  0125910301  & \\footnotesize  11216 & \\footnotesize  8211 & \\footnotesize  ok & \\footnotesize  HD 224317 \\\\ %\\footnotesize  Jansen, Fred & \\footnotesize   2002-06-17 \\\\ \\footnotesize 23 57 02.9 & \\footnotesize  -34 45 39 & \\footnotesize  356.3 & \\footnotesize  -76.0 & \\footnotesize  0.110 & \\footnotesize  0109950101  & \\footnotesize  30767 & \\footnotesize  18525 & \\footnotesize  ok & \\footnotesize  A 4059 \\\\ %\\footnotesize  Kaastra, Jelle & \\footnotesize   2000-11-24 \\\\ \\footnotesize 23 57 02.9 & \\footnotesize  -34 45 39 & \\footnotesize  356.3 & \\footnotesize  -76.0 & \\footnotesize  0.110 & \\footnotesize  0109950201  & \\footnotesize  25967 & \\footnotesize  19251 & \\footnotesize  discard & \\footnotesize  A 4059 \\\\ %\\footnotesize  Kaastra, Jelle & \\footnotesize   2000-11-24 \\\\ \\footnotesize 23 57 27.0 & \\footnotesize  -30 27 29 & \\footnotesize  14.0  & \\footnotesize  -77.7 & \\footnotesize  0.137 & \\footnotesize  0103861501  & \\footnotesize  16336 & \\footnotesize  9388 & \\footnotesize  flared & \\footnotesize  AM2354-304S \\\\ %\\footnotesize  Aschenbach, Bernd & \\footnotesize   2001-05-25 \\\\    \\hline \\end{longtable}  %\\caption[Spectroscopically Confirmed \\zg1 X-ray Clusters]{Spectroscopically confirmed \\zg1 X-ray luminous galaxy clusters in June 2007. ++++++. References: 1\\,=\\,Stanford \\etal~\\cite*{Stanford2006a}, 2\\,=\\,Mullis \\etal~\\cite*{Mullis2005a} , 3\\,=\\,  4\\,=\\,  5\\,=\\,  6\\,=\\,  7\\,=\\,  8\\,=\\,  9\\,=\\,  10\\,=\\, } %\\label{tA_field_list} %\\end{center} %\\end{table}      %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% %%%%%%%%%%%%%%%%%%%%%% Appendix B %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% \\chapter{List of Technical Terms and Acronyms}  \\begin{description}  %\\item[:]  \\item[2MASS:] Two Micron All-Sky Survey  \\item[AB Magnitudes:] Magnitude system which uses a fiducial source of constant flux $S_{\\nu}^{\\mathrm{AB}}\\!=\\!2.89\\!\\times\\!10^{-21}$\\,erg \\,cm$^{-2}$Hz$^{-1}$ as reference standard \\cite{Schneider2006a}. This value is chosen in a way that the V-band magnitudes are equivalent to the Vega system $m_V^{\\mathrm{AB}}=m_V^{\\mathrm{Vega}}$. Redder filter bands have a positive offset to the Vega system (fainter magnitudes), bluer bands have negative offsets (brighter magnitudes).   \\item[ACS:] Advanced Camera for Surveys on board the Hubble Space Telescope  \\item[ACT:] Atacama Cosmology Telescope  \\item[AGB:] Asymptotic Giant Branch  \\item[AGN:]  Active Galactic Nuclei  \\item[APEX:] Atacama Pathfinder EXperiment  \\item[Baffle:] Light shield placed in the optical path to block stray light and reduce the thermal background radiation.  \\item[BAOs:] Baryon Acoustic Oscillations  \\item[BCG:]  Brightest Cluster Galaxy   \\item[Bias:] An offset voltage applied to all pixels in an array detector.  \\item[CA:] Calar Alto Observatory in Southern Spain  \\item[CAL:] Calibration Access Layer   \\item[CCD:] Charge-Coupled Device  %\\item[COMBO-17:] %    Classifying Objects by Medium-Band Observations with 17 filters  \\item[CDM:] Cold Dark Matter  \\item[CFHT:] Canada-France-Hawaii Telescope  \\item[Chandra:] American-led X-ray observatory launched in 1999 with an on-axis spatial resolution of 0.5\\arcsec.  \\item[CMB:] Cosmic Microwave Background  \\item[COSMOS:] Cosmic Evolution Survey   %\\item[Command Line Parameter:] %    Parameter that is passed to a MIDAS routine with a function call %    to allow a high degree of flexibility in the program  \\item[Cosmic:] Cosmic ray event in an image. Signal in isolated pixels caused by a charged cosmic particle hit.  \\item[CTIO:] Cerro Tololo Inter-American Observatory  \\item[CXB:]  Cosmic X-ray Background  \\item[DDT:] Director's Discretionary Time  \\item[DEC:] Declination  \\item[DES:] Dark Energy Survey  \\item[Descriptor:] Data buffer with a name that contains information about an associated image.  \\item[DETF:] Dark Energy Task Force  \\item[DET\\,ML:]  Detection Maximum Likelihood  \\item[Dithering:] Common observation technique in infrared astronomy. Multiple images of an object are taken with slight telescope offsets in between the images. By overlaying several images and using a median process, the local sky for the data reduction can be extracted from the science images.  \\item[ECF:]  Energy Conversion Factor  \\item[EMSS:] Einstein Medium Sensitivity Survey  \\item[EMMI:] The ESO Multi Mode Instrument at the NTT telescope in La Silla.  \\item[EoS:] Equation-of-State relating density and pressure  \\item[EPIC:] European Photon Imaging Cameras  on XMM-Newton  \\item[eROSITA:] extended R\\\"ontgen Survey with an Imaging Telescope Array  \\item[ESA:] European Space Agency  \\item[ESO:] European Southern Observatory, with its headquarters in Garching, Germany.  \\item[EXT\\,ML:]  Extent Maximum Likelihood  \\item[FITS:] Flexible Image Transport System. Standard format for astronomical images.  \\item[Flatfield:] Image of a uniformly-illuminated area used for the correction of sensitivity variations across the detector.  \\item[FORS\\,2:] The FOcal Reducer and low dispersion Spectrograph for the Very Large Telescope. %The {\\bf FO}cal {\\bf R}educer and low dispersion {\\bf S}pectrograph for the Very Large Telescope.  \\item[FoV:] Field-of-View. Sky area that can be covered with one image of a particular instrument.  \\item[FP:]   False Positive  \\item[FPN:] Fixed Pattern Noise. Pixel-to-pixel sensitivity variations that are corrected by dividing through a normalized flatfield.  \\item[Fringing:] Interference patterns produced in CCD layers by strong monochromatic atmospheric emission lines at long optical wavelengths.  \\item[FWHM:] Full-Width Half-Maximum,  a measure of the width of an object in an image. The FWHM is a well-defined number obtained by fitting a Gaussian curve of the form $f(x)\\!=\\!1 / (\\sigma \\sqrt{2 \\pi}) \\times exp(-\\frac{x^2}{2 \\sigma ^2})$ to the intensity profile of an object. The standard deviation $\\sigma$ and the full width at half the peak intensity of the profile only differ by a constant factor of $\\sigma\\!=\\!0.424 \\cdot$FWHM \\cite{stoecker1998}.  \\item[Gain:] Conversion factor from detected electrons to digital counts.  \\item[GDDS:] Gemini Deep Deep Survey  \\item[GISSEL:] Galaxy Isochrone Synthesis Spectral Evolution Library  \\item[GMOS:] Gemini Multi-Object Spectrograph at the Gemini South 8\\,m telescope.  \\item[GNIRS:] Gemini NIR Spectrograph at the Gemini South 8\\,m telescope.  \\item[GR:] General Relativity  \\item[GROND:]  Gamma-Ray Burst Optical Near-IR Detector  \\item[GTI:]  Good Time Intervals  \\item[HEW:] Half Energy Width   \\item[HgCdTe:] Mercury-cadmium-telluride  \\item[HIFLUGCS:] HIghest X-ray FLUx Galaxy Cluster Sample  \\item[HIROCS:] Heidelberg InfraRed / Optical Cluster Survey  \\item[HR:]  Hardness Ratio  \\item[HR:] Hertzsprung-Russell diagram  \\item[HST:] Hubble Space Telescope  \\item[ICM:] Intracluster Medium  \\item[Image Catalog:] List of image names that can be used within the MIDAS environment to apply a set of commands to all images in the list.  %\\item[InSb:] %    Indium antimonide  \\item[IMF:] Stellar Initial Mass Function  \\item[IRAC:] InfraRed Array Camera on the Spitzer observatory   \\item[IRAF:] Image Reduction and Analysis Facility. Multi-purpose astronomical software for image reduction and analysis.   \\item[I/O:] Input/Output  \\item[IR:] InfraRed  \\item[ISAAC:] Infrared Spectrometer And Array Camera at the VLT  \\item[Keyword:] Variable of a specified data type within the MIDAS environment.  \\item[LF:] Luminosity Function  \\item[LH:] Lockman Hole  \\item[LM:] Limiting Magnitude. Apparent magnitude of objects with a signal-to-noise ratio~of~5.  \\item[LSS:] Large-Scale Structure  \\item[M67:] Open star cluster in the constellation Cancer at $\\alpha\\!=\\!8^h51^m$ and $\\delta\\!=\\!11.\\degr 82$.  %\\item[M51:] %    Spiral galaxy in the constellation Canes Venatici at $\\alpha\\!=\\!12^h17^m$ %    and $\\delta\\!=\\!47.\\degr 20$.  \\item[Median:] Central value in an ordered sequence of numbers.  \\item[MEKAL:] Plasma emission code from MEwe, KAastra, and Liedahl  \\item[MH:]  Mexican Hat wavelet  \\item[MIDAS:] Munich Image Data Analysis System. Astronomical software package developed and maintained by the European Southern Observatory.  \\item[ML:]  Maximum Likelihood  \\item[MOS:] Metal Oxid Semi-conductor instrument on XMM-Newton  \\item[NED:]  NASA Extragalactic Data Base  \\item[NEP:] North Ecliptic Pole survey  %\\item[Multiplexer:] %    Device that combines many signals into a small number of signals.  %\\item[NGST:] %    Next Generation Space Telescope. Now named the James Webb Space %    Telescope (JWST).  \\item[NICMOS:] Near Infrared Camera and Multi-Object Spectrometer on the Hubble Space Telescope  \\item[NIR:] Near-InfraRed. In this thesis, usually used for the spectral region 1--2.5\\,\\microns, unless otherwise noted.  \\item[NORAS:] NOrthern ROSAT All-Sky survey   \\item[NTT:] The 3.5\\,m New Technology Telescope at La Silla Observatory in Chile.   \\item[OBSID:] Unique identification number for each XMM observation.  \\item[ODF:] Observation Data File   \\item[OMEGA2000:] The NIR wide-field imager at the 3.5\\,m Calar Alto telescope.  \\item[OoT:]  Out-of-Time events  \\item[PanSTARRS:] Panoramic Survey Telescope \\& Rapid Response System  \\item[PI:] Principal Investigator  \\item[Pipeline:] Automated data reduction software requiring only minimal user interaction.  \\item[Pixel Coordinates:] Coordinate system based on the pixel row and column numbers on the detector.  \\item[PN:] Imaging camera on XMM-Newton   \\item[Pointing:] One principal telescope position. The sum of all images with a specific target in the field of view.  \\item[PSF:] Point-Spread Function. Intensity distribution of a point-like light source in an astronomical image.  \\item[PSPC:] Position Sensitive Proportional Counter, main instrument of ROSAT.  \\item[QE:] Quantum Efficiency.  \\item[QSO:] Quasi-Stellar Object  \\item[Seeing:] ``Blurring'' of point-like astronomical objects caused by atmospheric turbulence. Without adaptive optics, the seeing limits the highest possible resolution in ground based observations. Typical seeing values for the Calar Alto observatory are somewhat better than 1\\,arcsec. %%    The median seeing in the years 2001 and 2002 %%    was 0.82\\,arcsec\\footnote{See %%    \\url{http://www.caha.es/CAHA/MISC/seeing.html} for %%    further details.}.  \\item[RA:]  Right Ascension  \\item[RASS:] ROSAT All-Sky Survey  \\item[RCS:] Red-Sequence Cluster Survey  \\item[RDCS:] ROSAT Deep Cluster Survey  \\item[REFLEX:] ROSAT-ESO Flux-Limited X-ray cluster survey  \\item[RGS:] Reflecting Grating Spectrometer on XMM-Newton  \\item[ROSAT:] ROentgen SATellite, German X-ray survey mission from 1990 to 1999.  \\item[RWM:] Robertson-Walker-Metric  \\item[SAM:] Semi-Analytic Model  \\item[SAS:] XMM-Newton Science Analysis Software  \\item[SDSS:] Sloan Digital Sky Survey  \\item[SED:] Spectral Energy Distribution  \\item[SEP:] South Ecliptic Pole  \\item[SFH:] Star Formation History  \\item[SFR:] Star Formation Rate  \\item[SGP:] South Galactic Pole  \\item[SMF:]  Spectral Matched Filter Scheme  \\item[SNR:] Signal-to-Noise-Ratio.  \\item[SPIE:] Society of Photo-Optical Instrumentation Engineers  \\item[Spitzer:] Mid-infrared space telescope launched in 2003.  \\item[SPT:] South Pole Telescope  \\item[SRT:] Special Relativity Theory  \\item[SSP:] Simple Stellar Population models. Galaxy evolution models assuming a single formation redshift of the stellar population       and only passive evolution afterwards.  \\item[STD:]  Standard Scheme  \\item[SZE:] Sunyaev-Zeldovich Effect  \\item[UT:] Universal time.  \\item[VADER:] Very Ambitious Dark Energy Research mission  \\item[Vega Magnitudes:] Magnitude system calibrated to the A0 star Vega for which all filter bands have by definition a magnitude of 0. Since Vega's   spectrum has decreasing flux towards the red optical and NIR filters, these bands have magnitude values which are smaller (\\ie \\ brighter) than the corresponding AB magnitudes. The Vega system is still the standard photometric system for NIR observations.  \\item[Vignetting:] Obscuration of parts of the primary mirror as seen by a detector pixel.  \\item[VISTA:] Visible and Infrared Survey Telescope for Astronomy in Chile  \\item[VLT:] Very Large Telescope in Chile  \\item[VST:] VLT Survey Telescope  \\item[WARPS:] Wide Angle ROSAT Pointed Survey  \\item[WMAP:] Wilkinson Microwave Anisotropy Probe  \\item[World Coordinates:] Coordinate system that is based on absolute positions on the celestial sphere.  \\item[XCS:] XMM Cluster Survey  \\item[XDCP:] XMM-Newton Distant Cluster Project  \\item[XLF:] X-ray Luminosity Function  \\item[XMM-LSS:] XMM Large-Scale Structure survey  \\item[XMM-Newton:] X-ray Spectroscopy Multi-Mirror Mission  \\item[XSA:] XMM-Newton Science Archive  \\end{description}      \\end{appendix}     %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%   %------------------------ BIBLIOGRAPHIE ----------------------------------------  \\newpage  \\thispagestyle{empty}"
    },
    "0806/0806.0630_arXiv.txt": {
        "abstract": "The general relativistic precession rate of periastra in close-in exoplanets can be orders of magnitude larger than the magnitude of the same effect for Mercury.  The realization that some of the close-in exoplanets have significant eccentricities raises the possibility that this precession might be detectable. We explore in this work the observability of the periastra precession using radial velocity and transit light curve observations. Our analysis is independent of the source of precession, which can also have significant contributions due to additional planets and tidal deformations. We find that precession of the periastra of the magnitude expected from general relativity can be detectable in timescales of $\\lesssim 10$ years with current observational capabilities by measuring the change in the primary transit duration or in the time difference between primary and secondary transits. Radial velocity curves alone would be able to detect this precession for super-massive, close-in exoplanets orbiting inactive stars if they have $\\sim 100$ datapoints at each of two epochs separated by $\\sim 20$ years. We show that the contribution to the precession by tidal deformations may dominate the total precession in cases where the relativistic precession is detectable. Studies of transit durations with {\\it Kepler} might need to take into account effects arising from the general relativistic and tidal induced precession of periastra for systems containing close-in, eccentric exoplanets. Such studies may be able to detect additional planets with masses comparable to that of Earth by detecting secular variations in the transit duration induced by the changing longitude of periastron. ",
        "introduction": "Following the discovery of an extra-solar planet around the solar type star 51 Pegasi \\citep{Mayor1995a} there has been rapid progress in the detection and characterization of extra-solar planetary systems. The very early discoveries have shattered our view on planetary systems, as certain systems exhibited short periods (51 Peg), high eccentricities \\citep[e.g.\\ 70 Virginis b,][]{Marcy1996}, and massive planetary companions \\citep[e.g.\\ Tau Boo b,][]{Butler1997a}.  Interestingly, systems with all these properties combined (i.e.~massive planets with short periods, small semi-major axes, high eccentricities) have been also discovered \\citep[e.g., HAT-P-2b, XO-3b;][]{Bakos2007a,Johns-Krull2008a}. The high eccentricities are somewhat surprising, as hot Jupiters with short periods are generally expected to be  circularized in timescales shorter than the lifetime of the system if the parameter $Q$, inversely proportional to the planet's tidal dissipation rate, is assumed to be similar to that inferred for Jupiter \\citep{GoldreichSoter1966,Rasio1996}.  By virtue of their small semi-major axes and high eccentricities, the longitude of periastron $\\omega$ of some of the newly discovered systems are expected to precess due to General Relativistic (GR) effects at rates of degrees per century. This is orders of magnitude larger than the same effect observed in Mercury in our Solar System ($43\\arcsec$/century), which offered one of the cornerstone tests of GR. Furthermore, the massive, close-in eccentric planets induce significant reflex motion of the host star, therefore enhancing the detectability of the precession directly via radial velocities.  In this work we explore the observability of the precession of the longitude of periastron with the magnitude expected from GR in exoplanets using radial velocity and transit timing observations. We also consider in this work the periastra precession due to planetary perturbers and tidal deformations, which can have contributions comparable or greater than that of GR. Previous works \\citep{Miralda-Escude2002a,Heyl2007a} have explored some aspects of the work presented here in the context of using timing observations to detect terrestrial mass planets. We refer the reader to independent work by \\cite{Pal2008b} that also explores the measurable effects of the periastra precession induced by GR. ",
        "conclusions": "In this work we have studied the observability of the precession of periastra caused by general relativity in exoplanets. We additionally consider the precession caused by tidal deformations and planetary perturbers, which can produce a precession of comparable or greater magnitude. We consider radial velocities and transit light curve observations and conclude that for some methods precessions of the magnitude expected from GR will be detectable in timescales of $\\sim 10$ years or less for some close-in, eccentric systems.  In more detail, we find that:  \\begin{enumerate}  \\item For transiting systems, precession of periastra of the magnitude expected from GR will manifest itself through detectable changes in the duration of primary transit (\\S\\ref{subsec:D}) or through the change in the time between primary and secondary transits (\\S\\ref{subsec:Deltat}) in timescales of $\\lesssim 10$ years.  The two methods are most effective at different values of the true anomaly.  A determination of the primary transit duration time to $\\sim$ a few seconds and that of the secondary to $\\sim$ a minute will lead to measurable effects. The effects of GR and tidal deformations might need to be included in the analysis of {\\it Kepler} data for eccentric, close-in systems. The transit duration of near-grazing systems will be particulary sensitive to changes in $\\omega$.   \\item Radial velocity observations alone would be able to detect changes in the longitude of periastron of the magnitude expected from GR effects only for eccentric super-massive ($K \\sim 1000$ m sec$^{-1}$) exoplanets orbiting close to a host star with a low-level of stellar jitter (\\S\\ref{subsec:rv}). For the detection to be statistically significant, on the order of $100$ precise radial velocity observations are needed at each of two epochs separated by $\\sim 20$ years.  \\item Measurements of the change over time of the period between primary transits is not currently a method that will lead to a detection of changes in $\\omega$ of the magnitude expected from GR (\\S\\ref{subsec:dPdt}). Previous works have shown that measuring the small difference between the radial velocity period and that of transits are not sensitive enough to lead to detectable changes due to GR.  \\end{enumerate}   In order to contrast any detected change in the transit duration (\\S\\ref{subsec:D}) or the time between primary and secondary (\\S\\ref{subsec:Deltat}) to the predictions of a given mechanism one needs to know the eccentricity and longitude of periastron of the systems, for which radial velocities are needed (although not necessarily of the precision required to directly detect changes in $\\omega$ with them\\footnote{The eccentricity can be constrained using transit information alone, see \\citet{Ford2008a}}). Conversely, photometric monitoring of primary transits are useful in order to elucidate the nature of a detected change in $\\omega$ by probing for the presence of transit time variations. The presence of the latter would imply that at least part of any observed changes in $\\omega$ could have been produced by additional planetary companions (\\S\\ref{sec:distinguish}).  Precession of periastra caused by planetary perturbers and the effects of tidal deformations can be of comparable magnitude to that caused by GR (\\S\\ref{ssec:add_prec}). The effects of tidal deformations on the precession of periastra in particular may be of the same magnitude or dominate the total $\\dot{\\omega}$ in the regime where the GR effects are detectable (\\S\\ref{ssec:rel_mag}). While this limits the ability to directly extract the precession due to GR given the uncertainty in the expected precession from tides, it might allow to study the tidally induced precession by considering the residual precession after subtracting the effects of GR. The latter possibility is particularly attractive in systems where the tidally induced precession may dominate the signal. We note that even without considering the confusing effects of tidal contributions to the precession, a measurement of \\dotw as described in this work would not be competitive in terms of precision with binary pulsar studies \\citep[see, e.g.,][for a review ]{Will2006a} and would therefore not offer new tests of GR.  The upcoming {\\it Kepler} mission expects to find a large number of massive planets transiting close to their host stars \\citep{Borucki2003a}, some of which will certainly have significant eccentricities. Furthermore, systems observed by {\\it Kepler} will be extensively monitored for variations in their transiting time periods in order to search for terrestrial-mass planets using transit-time variations. We have shown that modeling of the transit time durations and further characterization of close-in, eccentric systems might need to take into account the effects of GR and tidal deformations as they will become detectable on timescales comparable to the 4-year lifetime of the mission, and certainly on follow-up studies after the mission ends. We have also shown that planetary companions with super-Earth masses may be detectable by {\\it Kepler} by the change in transit durations they induce (\\S~\\ref{ssec:tpl}). Additionally, well sampled radial velocity curves spanning $\\gtrsim 5$ years may also be able to detect companions with super-earth massses by measuring a change in $\\omega$ over the expected GR value for the case of super-massive, close-in systems orbiting inactive stars (\\S~\\ref{ssec:tpl})."
    },
    "0806/0806.2129_arXiv.txt": {
        "abstract": "We report an upper limit on the flux of relativistic monopoles based on the nonobservation of in-ice showers by the Radio Ice Cherenkov Experiment (RICE) at the South Pole. We obtain a 95\\% C.L. limit of order $10^{-18}(\\text{cm}^2\\text{ s sr})^{-1}$ for intermediate-mass monopoles of $10^7\\leq\\gamma\\leq10^{12}$ at the anticipated energy $E_\\text{tot} =10^{16}$ GeV. This bound is over an order of magnitude stronger than all previously published experimental limits for this range of boost parameters $\\gamma$ and exceeds 2 orders of magnitude improvement over most of the range.  We review the physics of radio detection, describe a Monte Carlo simulation including continuous and stochastic energy losses, and compare to previous experimental limits. ",
        "introduction": "Magnetic monopoles are hypothetical particles carrying a net magnetic charge. Dirac\\cite{Dirac31} calculated the quantum unit of magnetic charge while showing that existence of magnetic charge leads to charge quantization. The relationship between the fundamental magnetic charge $g$ and the quantum of electric charge $e$ takes on the especially simple form $g=\\frac{e}{2\\alpha}$ in Gaussian units, where $\\alpha$ is the fine structure constant.   Magnetic monopoles occur in  grand unified theories (GUTs)\\cite{thooft}.  Most GUTs predict a symmetry-breaking phase transition at an energy-scale $\\sim10^{16}$ GeV\\cite{Vilenkin94}.  Such a phase transition can lead to localized topological defects in the form of magnetic monopoles.  An order-of-magnitude estimate of one magnetic monopole per cubic Hubble radius gives a magnetic monopole number density at the GUT time of $n_M(t_{GUT})\\approx 10^{82}$ m$^{-3}$, which would lead to magnetic monopoles dominating the Universe today\\cite{Ryden03}, in conflict with observation.  Meanwhile, some inflationary models predict dilution of GUT monopole density to hopelessly undetectable levels\\cite{Guth81}.  No monopoles have been verifiably detected to date\\cite{Eidelman04}.  Reports of magnetic monopole detections\\cite{Price75, Cabrera82, Caplin86} have been challenged, often by the original authors themselves\\cite{Price78, Huber90}.  Alternate explanations range from ionized nuclei\\cite{Price78} to hardware malfunctions\\cite{Caplin86} to new physics other than monopoles\\cite{Price78}.  In this paper we present limits on relativistic monopole fluxes determined from five years of data collection by the RICE experiment\\cite{Rice1}.  As noted by Wick {\\it et al.}\\cite{Wick03} the RICE array is particularly well-suited to ultrarelativistic monopole detection because of a combination of large effective volume and favorable scaling with energy.  Our limits for fluxes over the range of monopole Lorentz boost parameters $ 10^8 \\lesssim \\gamma \\lesssim 10^{12}$ are about 100 times more restrictive than the tightest published limits from AMANDA\\cite{Wissing07}, Baikal\\cite{Aynutdinov05}, and MACRO\\cite{Ambrosio02}, and more than 1000 times more restrictive than the original Parker bound.  This paper first reviews monopole properties and previous limits in Sec. \\ref{sec:monoreview}.  We then provide an overview of RICE in Sec. \\ref{RICEsection}.  Sections \\ref{sec:eloss}, \\ref{sec:MC}, and \\ref{sec:response} discuss monopole energy loss, our Monte Carlo simulation thereof, and RICE's response.  Results are summarized in Sec. \\ref{sec:results}. ",
        "conclusions": "-\\frac{dE}{dx} = \\alpha + \\beta E. \\end{equation} The $\\alpha$ term is the energy loss per distance (units: g/cm$^2$) due to ionization of the medium.  The $\\beta$ term\\cite{notebeta} is the sum of three terms reflecting bremsstrahlung, pair production, and photonuclear effect energy losses.  Each energy loss mechanism is calculated separately.  Defining $y$ to be the fraction of its energy lost by the particle in a single interaction with the medium, each of the three terms $\\beta_i$ is found by convolving $y$ with the partial interaction cross section with respect to $y$ \\eqref{inty}: \\begin{equation} \\label{inty} \\beta_i(E) = \\frac{N}{A} \\int_{y_{\\text{min}_i}}^{y_{\\text{max}_i}} y \\frac{d \\sigma_i(y,E)}{dy} dy \\end{equation} Here, $N$ is Avogadro's number and $A$ is the average atomic mass number of the medium.  The full formulae for $\\alpha$ and the $y_{\\text{min}_i}$, $y_{\\text{max}_i}$, and ${d \\sigma_i}/{dy}$ needed to calculate each $\\beta_i$ are given in\\cite{Dutta01}.  In general, the expressions are functions of particle mass and energy and of various properties of the medium, although $\\alpha$ and the individual $\\beta_i$'s are only weakly energy dependent.  \\subsection{Discrete Loss Processes}  Although energy loss due to ionization can be treated as smooth and continuous with little loss of accuracy, we explicitly model the stochastic fluctuation in pair production and photonuclear energy losses.  Combining Eq. \\eqref{losssummary} with \\eqref{inty} and replacing the integral in \\eqref{inty} with the corresponding Riemann sum gives the result \\eqref{Rsum}, where $\\Delta E_i$ is the energy loss via process $i$ (brem., pair, or photonuclear) over a small distance $\\Delta x$: \\begin{equation} \\label{Rsum} \\Delta E_i \\approx \\sum_{ \\substack{j \\\\ y_j=y_{\\text{min}_i}} }^{y_j=y_{\\text{max}_i}}\\left(\\frac{N}{A}\\right)(y_j E)\\left( \\Delta x \\frac{d\\sigma_i}{dy_j}\\right)\\Delta y \\end{equation} Recasting the energy loss equation this way effectively sorts the total energy loss into an arbitrary number of bins, each of which spans a length $\\Delta y$ of the possible $y$ values.  Since $y_j$ is the fractional energy loss in a single interaction within bin $j$ and $E$ is the total energy of the particle, $(y_j E)$ is the energy loss for a single interaction in the $j^{\\text{th}}$ bin.  Each term of the Riemann sum represents an energy loss, so if $(y_j E)$ is the energy lost in a single interaction, the remaining multiplicative terms in the summand give the expectation number of interactions in the $j$th bin $\\langle n_{ij} \\rangle$: \\begin{equation} \\label{expectation} \\langle n_{ij} \\rangle = \\frac{N}{A}\\Delta x \\frac{d\\sigma_i}{dy_j}\\Delta y \\end{equation} Therefore, accurately modeling the stochastic variation in bremsstrahlung, pair production, and the photonuclear effect is equivalent to replacing $\\langle n_{ij} \\rangle$ by a random number drawn from a Poisson distribution of expectation value $\\langle n_{ij} \\rangle$ when evaluating the energy loss expressions \\eqref{Rsum}.  \\subsection{Generalization to Monopoles} \\label{sec:model}  Only a few changes are needed to convert  the stochastic model of muon energy loss to a model of magnetic monopole energy loss.  First, the muon mass must be replaced by the magnetic monopole mass.  Because bremsstrahlung falls off by inverse powers  of particle mass, the bremsstrahlung energy loss contribution is negligible for even light magnetic monopoles and will be subsequently disregarded\\cite{Wick03}.  It should be noted that at large masses ($\\gtrapprox$1 TeV), $\\beta_\\text{pair production}$ can become difficult to calculate numerically due to rounding error.  However, pair production energy loss approaches an asymptotic limit with increasing particle mass and varies with mass by only a few percent for masses above $\\sim 100$ MeV.  Next, due to Dirac's quantization condition a magnetic monopole of 1 Dirac charge will lose energy equivalent to an electric charge of $1/(2\\alpha)$ times the proton charge\\cite{Jackson62}.  Accounting for this large effective charge only requires multiplying the expectation number of interactions by $1/(2\\alpha)^2 \\approx 4700$.  This procedure for modelling magnetic monopole energy loss in matter has assumed the particle to be a simple Dirac monopole, that is, a point source of magnetic charge with no further internal structure.  ``Actual'' magnetic monopoles may contain internal color fields and lose additional energy through hadronic interactions beyond the photonuclear effect\\cite{Rubakov}. However calculations of such processes are highly model-dependent\\cite{Wick03} and not further considered in this analysis.  Figure \\ref{energyloss} shows the average monopole energy loss in ice and ``standard rock'' (A=22, $\\rho=2.65 \\text{~g/cm}^3$\\cite{Dutta01}), along with the same results for muons.  While the muon mass is fixed, the monopole rest mass is constrained to vary inversely with gamma such that total energy is fixed at a reference energy of $10^{16}$ GeV.   Figure \\ref{monobyparts} shows various contributions to the energy loss of a  $10^{16}$ GeV monopole. Much of the difference between muons and monopole energy losses is due to the monopole's large effective charge.   The curves indicate average energy loss due to the three principal mechanisms, while the points show actual stochastic energy loss (as averaged over a 50 m interval).  The photonuclear effect is the dominant energy loss mechanism at $\\gamma>10^4$, while ionization energy losses dominate below this value\\cite{noteerror}.  Because the photonuclear mechanism results in hadronic showers generated by nuclear recoils, we may ignore LPM\\cite{LPM} effects.  There are substantial uncertainties in extrapolations of photonuclear losses to ultrahigh energies.  The primary unknown is the hadronic contribution of real photon-nucleon scattering. Consider an extrapolation based on the Froissart bound setting in at about 50 GeV\\cite{Froissart}.  The energy dependence of the photon-nucleon cross section $\\sigma_{\\gamma N}$ of such a model is \\begin{eqnarray} \\sigma_{\\gamma N}(E_{\\gamma}) = 114.3+1.67 \\,{\\rm ln}^2 (0.0213 E/{\\rm GeV})\\, \\mu{\\rm b} \\label{bb-photonuclear} \\end{eqnarray}  This 1981 model of Bezrukov and Bugaev\\cite{Bezrukov81} predates the discovery at HERA of parton distributions at small$-x$ that causes cross sections to grow like fractional powers.  For comparison, the 1998-2001 post-HERA photonuclear cross section of Donnachie and Landshoff \\cite{DL} is within 10\\% of Eqn. \\eqref{bb-photonuclear} at $E_{\\gamma}=10^{6}$ GeV, while being about 16 times larger at $10^{11}$ GeV.  As shown in Fig. 3 of Ref. \\cite{Dutta01}, the more complicated cross sections developed by extrapolating structure functions track the simple expression of Eq. \\ref{bb-photonuclear} very well. Using small and slow-growing cross section models such as these is in some sense conservative.  It results in dim showers that are less likely to trigger.  However, small cross sections also tend to develop fewer showers failing the time-over-threshold cut discussed below, and fewer showers are stopped by the Earth.  No one cross section model is in all cases the ``most conservative.''  \\begin{figure}\\centering \\includegraphics[width=0.5\\textwidth]{ice_rock_eloss_revise_color} \\caption{(Color online) Muon energy loss in ice and standard rock compared to monopole energy loss in ice and standard rock, as a function of boost parameter $\\gamma$.  Monopole results are shown for total energy $10^{16}$ GeV (the energy assumed in this analysis) and, for comparison, $10^{14}$ GeV.  The figure shows that, over the kinematic range of interest, monopole energy loss depends strongly on $\\gamma$, but for a given $\\gamma$ it is only weakly mass dependent.} \\label{energyloss} \\end{figure}  \\begin{figure} \\centering \\includegraphics[width=0.5\\textwidth]{monobyparts_constE_revise_color} \\caption{(Color online) Total energy loss versus $\\gamma$ for $10^{16}$ GeV monopoles, showing stochastic variation over 50m intervals.  Lines show average contributions from different processes.  Also shown, again for comparison, is the average total energy loss for $10^{14}$ GeV monopoles.  In the latter case, stochastic variation and energy loss contributions from the various processes are almost identical to the former case.} \\label{monobyparts} \\end{figure}    From the nonobservation of highly ionizing shower ``trails'' we have derived the monopole flux upper limits shown in Fig. \\ref{figflux}, which are on the order of $10^{-18} (\\text{cm}^2\\text{ s sr})^{-1}$.  Previously, AMANDA\\cite{Wissing07}, Baikal\\cite{Aynutdinov05}, and MACRO\\cite{Ambrosio02} determined monopole flux limits on the order of $10^{-16}(\\text{cm}^2\\text{ s sr})^{-1}$ for $\\beta$ greater than $0.8$, $0.8$, and $4\\times10^{-5}$, respectively.  Although the results of this study cover a much narrower range of $\\beta$ values than previous works, it is the range that is of the greatest interest for IMM searches.  Within much of this kinematic range ($E=10^{16}$ GeV; $\\gamma\\geq10^{8}$), monopole flux limits from RICE are stronger than the limits from any previous astrophysical monopole search by more than an order of magnitude.  \\medskip  {\\bf Acknowledgments} We acknowledge helpful conversations with Alfred Goldhaber, Tom Weiler, and Stuart Wick.  This research was supported by the University of Kansas and by the National Science Foundation under grants No. OPP-0338219 and No. PHY-0243935.  JPR is supported by DOE-HEP Grant No. DE-FG02-04ER14308."
    },
    "0806/0806.3204_arXiv.txt": {
        "abstract": "{}{We analyse absorption characteristics and physical conditions of extraplanar intermediate- and high-velocity gas to study the distribution of the neutral and weakly ionised Milky Way halo gas and its relevance for the evolution of the Milky Way and other spiral galaxies.} {We combine optical absorption line measurements of \\ion{Ca}{ii}/\\ion{Na}{i} and 21\\,cm emission line observations of \\ion{H}{i} along 103 extragalactic lines of sight towards quasars (QSOs) and active galactic nuclei (AGN). The archival optical spectra were obtained with the Ultraviolet and Visual Echelle Spectrograph (UVES) at the ESO Very Large Telescope, while the 21\\,cm \\ion{H}{i} observations were carried  out using the 100-m radio telescope at Effelsberg.} {The analysis of the UVES spectra shows that single and multi-component \\ion{Ca}{ii}/\\ion{Na}{i} absorbers at intermediate and high velocities are present in about 35 percent of the sight lines, indicating the presence of neutral extraplanar gas structures. In some cases the \\ion{Ca}{ii}/\\ion{Na}{i} absorption is connected with \\ion{H}{i} 21\\,cm intermediate- or high-velocity gas with \\ion{H}{i} column densities in the range of $10^{18}$ to $10^{20}\\,\\mathrm{cm}^{-2}$ (i.e., the classical IVCs and HVCs), while other \\ion{Ca}{ii}/\\ion{Na}{i} absorbers show no associated \\ion{H}{i} emission. The observed \\ion{H}{i} line widths vary from $\\Delta v_\\mathrm{FWHM}=3.2$\\,km\\,s$^{-1}$ to $32.0$\\,km\\,s$^{-1}$ indicating a range of upper gas temperature limits of 250\\,K up to about 22500\\,K.} {Our study suggests that the Milky Way halo is filled with a large number of neutral gaseous structures whose high column density tail represents the population of common \\ion{H}{i} high-velocity clouds seen in 21\\,cm surveys. The \\ion{Ca}{ii} column density distribution follows a power-law $f(N)=CN^{\\beta}$ with a slope of $\\beta \\approx -1.6$, thus comparable to the distribution found for intervening metal-line systems toward QSOs. Many of the statistical and physical properties of the \\ion{Ca}{ii} absorbers resemble those of strong ($W_\\mathrm{\\lambda 2796}>0.3\\,\\AA{}$) \\ion{Mg}{ii} absorbing systems observed in the circumgalactic environment of other galaxies, suggesting that both absorber populations may be closely related.} ",
        "introduction": "Spiral galaxies are surrounded by large gaseous halos. That the disk of the Milky Way has a hot envelope was first proposed by \\citet{spitzer56}. Spitzer considered such a ``Galactic Corona'' to explain spectroscopic observations that have been made earlier by \\citet{adams49} and \\citet{muench52}. In recent years, great instrumental progress has been made to measure extraplanar gas structures around other galaxies, as well. It became clear, that the gaseous halos represent the interface between the condensed galactic discs and the surrounding intergalactic medium (IGM) \\citep[e.g.,][ and references therein]{savage_massa87, majewski04, fraternalietal_07}.  The properties of this medium around galaxies presumably are determined by both, the accretion of metal-poor gaseous matter from intergalactic space onto the galactic disc, as well as the outflow of metal-enriched gas caused by star formation activity within the galaxy \\citep[e.g.,][ and references therein]{sembach_wakker_savage_richter_etal03, fraternaliandbinney06}.   Gas in the halos and intergalactic environment of galaxies leaves its imprint in the spectra of distant quasars (QSOs) in the form of hydrogen- and metal-line absorption \\citep[for a recent review, see][]{richterp06}. Therefore, QSO absorption spectroscopy has become a powerful method to study the physical properties, the kinematics, and the spatial distribution of gas in the halos of galaxies over a large range of column densities at low and high redshifts. The analysis of intervening \\ion{Mg}{ii} and \\ion{C}{iv} absorption line systems \\citep[e.g.,][]{charltonetal00, dingetal03, masieroetal05, boucheetal06} and their relation to galactic structures suggest rather complex absorption characteristics of these ions indicating the multi-phase nature of gas in the outskirts of galaxies with density and temperature ranges spanning several order of magnitudes. Stronger intervening low-ion absorbers (e.g., strong \\ion{Mg}{ii} systems) preferentially arise at low impact parameters ($< 35\\,h^{-1}$\\,kpc) of intervening galaxies, while weaker \\ion{Mg}{ii} systems and high-ion absorbers (e.g., \\ion{C}{iv} systems) apparently are often located at larger distances up to $\\sim 100\\,h^{-1}$\\,kpc  (Churchill et al.\\,1999; Milutinovic et al.\\,2006). Due to the lack of additional information the exact nature and origin of the various circumgalactic absorber populations is not yet fully understood. Most likely, gas outflow and infall processes both contribute to the complex absorption pattern observed.  Our own Galaxy also is surrounded by large amounts of neutral and ionised gas \\citep{richterp06}. Most prominent are the so-called intermediate- and high-velocity clouds \\citep[IVCs, HVCs,][] {mulleroortraimond63} which represent clouds of neutral atomic hydrogen seen in 21\\,cm emission at radial velocities inconsistent with a simple model of Galactic disk rotation. Most important for our understanding of the nature of IVCs and HVCs and their role for the evolution of the Milky Way is the determination of accurate metal abundances and distances of these clouds. Metallicity measurements are particularly important to learn about the origin of IVCs/HVCs. The metal abundances of some IVCs/HVCs have been determined by absorption line measurements along several lines of sight \\citep[see][]{wakker01, richterp06}. The results show that the metallicities are varying between $\\sim 0.1$ and $\\sim 1.0$ solar. This wide range of metallicities suggests that many IVCs and HVCs cannot have a common origin. In fact, it is know widely accepted that various different processes contribute to the neutral gas flow in the Milky Way halo including the Galactic fountain \\citep{shapirofield76, bregman80, shapiro_benjamin91}, the accretion of gas from surrounding satellite galaxies \\citep[e.g., Magellanic Stream, ][]{mathewson74}, and the infall of metal-poor gas from the intergalactic medium (e.g., Wakker et al.\\,1999).  To determine the total mass of the gas that is falling toward the Milky Way disk in form of IVCs and HVCs a reliable distance estimate of these clouds is required. Measuring the distances of IVCs/HVCs is very difficult, however. The most reliable method to infer a distance bracket requires high-resolution spectra of stars with known distances in which IVCs/HVCs appear in absorption. The problem is the limited number of suitable background stars. The distance estimates of IVCs and HVCs around the Milky Way \\citep[e.g.,][]{Sembachetal91, vanWoerden99, wakker01, thom2006, wakker_york_howketal07, wakkeryorkwilhelmetal08} indicate that most IVCs are relatively nearby objects with distances of $d< 2$\\,kpc while the majority of the HVCs are more distant clouds, located in the halo of the Milky Way with distances of $5<d<50$\\,kpc. These numbers imply that the typical mass-accretion rate from IVCs and HVC are on the order of one solar mass per year.  While there are a large number of recent absorption studies on the nature of IVCs and HVCs and their role for the evolution of the Milky Way, relatively little effort has been made to investigate the connection between the Galactic population of IVCs and HVCs and the distribution and nature of intervening metal-absorption systems from galaxy halos seen in QSO spectra. In fact, almost all recent absorption studies of IVCs and HVCs were carried out in the FUV to study in detail metal abundances and ionisation conditions of  halo clouds using the many available transitions of low and high ions in the ultraviolet regime (e.g., Richter et al.\\,2001). These studies were designed as follow-up absorption observations of known IVCs and HVCs, thus providing an 21\\,cm {\\it emission-selected} data set. However, to statistically compare the absorption characteristics of the extraplanar Galactic halo structures with the properties of intervening metal-absorption systems towards QSOs one requires an {\\it absorption-selected} data set of IVCs and HVCs. Since in the UV band there are currently only a very limited number ($<50$) of high-quality spectra available such a statistical comparison can be done best in the optical regime where a large number of high-quality spectra of low- and high-redshift QSOs are available.  In this paper we discuss low-column density extraplanar structures which are most likely located in the environment of the Milky Way detected in optical \\ion{Ca}{ii} and \\ion{Na}{i} absorption towards quasars along 103 sight lines through the Milky Way halo. Our study allows us to directly compare the observed absorption column-density distribution of gas in the Milky Way halo with the overall column-density distribution of intervening metal absorbers toward QSOs and AGN. Moreover, our study enables us to identify neutral and ionised absorption structures at low gas column densities and small angular extent that remain unseen in the large 21\\,cm IVC and HVC all-sky surveys but that possibly have a considerable absorption cross section \\citep[see][]{richterwestmeierbruens05}. We supplement our absorption-line data with new \\ion{H}{i} 21\\,cm observations to investigate the relation between intermediate- and high-velocity \\ion{Ca}{ii} absorption features and halo 21\\,cm emission. The study presented here discusses the first results of our analysis of the physical and statistical properties of the detected absorption and emission features. Additional absorption-line observations as well as follow-up \\ion{H}{i} synthesis observations of some of the absorbers are under way and will be presented in a subsequent paper (Ben Bekhti et al., 2008; in preparation).  Our paper is organised as follows. In Section\\,\\ref{data} we describe the data acquisition and data reduction. In Section\\,\\ref{results} the results of the analysis of the UVES and the Effelsberg data of 103 sight lines in the direction of several quasars are presented. A statistical investigation (column density distribution functions, distribution of $b$-values, deviation velocities, etc.)  of the data is presented in Section\\,\\ref{stat.properties}. In Section\\,\\ref{disscussion} we discuss the results regarding multiple absorption lines and corresponding emission line measurements, the determination of metallicities and a probable association of the gas with known IVC and HVC complexes. In Section\\,\\ref{summary} we summarise our results and discuss future observations that will be required to determine the metallicities of the gas, important to understand the nature and origin of these intermediate- and high-velocity absorption features. ",
        "conclusions": "\\label{disscussion}  The fact that among 103~random lines of sight through the halo we observe 35 optical extraplanar absorption systems with at least one intermediate- or high-velocity component shows that the Milky Way halo contains a large number of neutral gas structures that give rise to intermediate- and high-velocity \\ion{Ca}{ii} and \\ion{Na}{i} absorption.  \\subsection{Area filling factors}  In the following we want to determine the area filling factors of our absorbers observed with UVES and Effelsberg. The area filling factor which we define is the fraction of sight lines which show \\ion{Ca}{ii} absorption or \\ion{H}{i} absorption to the total number of observed sight lines. We get a total \\ion{Ca}{ii} area filling factor of about $30\\%$ and a significantly larger \\ion{H}{i} area filling factor for the Effelsberg data of about $68\\%$. Both factors are defined for the same lower \\ion{H}{i} column density regime as described above. The discrepancy between the total absorption and emission area filling factors could possibly be explained by the effect that the neutral gas structures that give rise to the \\ion{Ca}{ii} absorption only partly fill the $9\\arcmin$ beam of the Effelsberg telescope. This could be an indication that the angular extent of the absorbers is small. The beam sizes of the different instruments (UVES, Effelsberg, LAB) also must have an effect on the column density distributions. Especially possible small-scale, high column density structures get affected by beam smearing, leading to a bias towards smaller column densities. The slope of the column-density distribution from 21\\,cm data therefore is expected to steepen with stronger beam-smearing. Thus, the different slopes of the \\ion{H}{i} column-density distributions from the UVES data and the Effelsberg data can be explained naturally by the limited beam size of the radio observations.  We also determined the area filling factor for our intermediate- and high-velocity \\ion{Ca}{ii} absorbers separately. We distinguish between IVCs and HVCs by using the deviation velocity of the clouds as a separation criterion (see Section\\,\\ref{Deviation velocities}). The \\ion{Ca}{ii} area filling factor is about $16\\%$ for the intermediate-velocity and about $17\\%$ for the high-velocity gas for converted \\ion{H}{i} column densities above $\\log(N_\\mathrm{HI}/\\mathrm{cm}^{-2})\\approx16$. Using our Effelsberg observations, we calculate an \\ion{H}{i} area filling factor of about $37\\%$ (IVCs) and $16\\%$ (HVCs) for densities above $\\log(N_\\mathrm{HI}/\\mathrm{cm}^{-2})\\approx18$.  The analysed UVES and Effelsberg spectra show that for several cases the \\ion{Ca}{ii} and \\ion{Na}{i} absorbers have no intermediate- or high-velocity counterpart in the 21\\,cm data obtained with Effelsberg. This implies that the \\ion{H}{i} column densities of the \\ion{Ca}{ii} and \\ion{Na}{i} absorption components at these velocities fall below the detection limit of the Effelsberg 21\\,cm observations and also of current large surveys like the LAB survey. Another possibility for this non-detection is a small diameter of the clouds, so that beam-smearing effects will again become important.  \\subsection{Possible association with known IVC or HVC complexes} The directions of 19 of the intermediate- and high-velocity clouds as well as their velocities indicate a possible association with known large and extended IVC or HVC complexes (Tables\\,\\ref{tab_27_quasars} and \\ref{tab_20_quasars}). Since we do not know the distances to the several absorbing systems the only criteria for the determination of a possible association with IVC/HVC complexes are the positions and the velocities.  Eight absorbing systems, and therefore the majority, are associated with the extensive Magellanic Stream (MS). In all eight cases the sight lines passes only the outer regions of the MS. In the case of QSO\\,J0139$-$0824 and QSO\\,J0105$-$1846 it is not clear if they are possibly parts of the MS or the Anti-Center shell (AC shell) because they are located in the intermediate range. Only for the sight line towards QSO~J0003$-$2323 we have corresponding \\ion{H}{i} observations obtained with the 100-m telescope at Effelsberg.  Four systems are most likely associated with the intermediate-velocity Spur (IV-Spur). The sight lines toward QSO\\,J1211$+$1030 and QSO\\,B1331+170 pass the inner, denser part of the IV-Spur whereas the sight lines in the direction of QSO\\,B0952$+$179 and QSO\\,B2314$-$409 are located in the outer, clumpy structures of the complex. For three of these sight lines we have additional \\ion{H}{i} observations but only in one case (QSO\\,J1211$+$1030) we found a corresponding emission line.  The position and the velocity of the absorbing systems in the direction of QSO\\,B1212$+$3326 indicate that these systems are possibly part of the inner region of the intermediate-velocity Arch (IV-Arch). For these sight line we have obtained additional \\ion{H}{i} data. Only one component has a counterpart in \\ion{H}{i}.  One absorption component in the direction of QSO\\,J1039$-$2719 is most likely associated with the HVC complex WD. We re-observed this sight line with the 100-m telescope at Effelsberg but we find no corresponding \\ion{H}{i} emission lines.  The sight line towards QSO\\,B0458$-$0203 passes the very outer regions of the Anti-Center shell. The \\ion{H}{i} data reveal a corresponding emission line.  The velocity and the coordinates of the absorption component with the high positive LSR velocity in the direction of QSO\\,B1101$-$26 is possibly associated with the Leading Arm (LA) of the Magellanic System. The sight line passes the outer region of the LA. The Effelsberg data show no corresponding \\ion{H}{i} emission line.  The absorbing systems towards QSO\\,1448$-$232 are possibly associated with the HVC complex L \\citep{richterwestmeierbruens05}.  The two absorbing components in the direction of QSO\\,J2155$-$0922 are possibly part of the HVC complex GCN. The Effelsberg data show no corresponding \\ion{H}{i} emission line.  \\subsection{Comparison with \\ion{Mg}{ii} systems and column density distribution functions}  It is evident that low-column density intermediate- and high-velocity \\ion{Ca}{ii} absorbers have a large area filling factor because of the high \\ion{Ca}{ii} and \\ion{Na}{i} detection rate. If these absorbers represent a typical phenomenon of spiral galaxies like our Milky Way or M31 and if they are located in the halo of these galaxies, they should produce \\ion{H}{i} Lyman-Limit and \\ion{Mg}{ii} absorption along sight lines which pass through the halo gas of other galaxies. There are two different populations of \\ion{Mg}{ii} absorbers, the weak ($W_\\mathrm{\\lambda 2796} \\leq0.3\\,\\AA{}$) and the strong ($W_\\mathrm{\\lambda 2796}>0.3\\,\\AA{}$) \\ion{Mg}{ii} systems. While the weak \\ion{Mg}{ii} absorbers are thought to be located in the outskirts of galaxies \\citep[e.g.,][]{churchill99}, the strong \\ion{Mg}{ii} systems are most likely located in the halo \\citep[e.g.,][]{petitjeanbergeron90, charltonandchurchill98, dingcharltonchurchillpalma03} or even in the discs of galaxies \\citep[Damped Lyman Alpha systems, DLAs,][]{raoetal06}, and they possibly represent the analogues of IVCs/HVCs \\citep[e.g.,][]{savageetal00}. In fact, many of the properties of the high-velocity \\ion{Ca}{ii} and \\ion{Na}{i} absorption lines resemble those of the strong \\ion{Mg}{ii} absorption line systems observed in the circumgalactic environment of other galaxies \\citep{dingetal03, boucheetal06, prochteretal06}. The high-velocity \\ion{Ca}{ii} and \\ion{Na}{i} absorbers thus may possibly represent the Galactic counterparts of strong \\ion{Mg}{ii} systems at low redshift.  In Section\\,\\ref{stat.properties} we have shown that absorption systems at intermediate and high velocities with a single or a double absorption component structure are more common than systems with more than two absorption components. The analysis of strong \\ion{Mg}{ii} absorption systems by \\citet{prochteretal06} shows that there is a tendency for two or more \\ion{Mg}{ii} absorption lines (Table\\,\\ref{tab_7_components}). The more complex absorption structure for \\ion{Mg}{ii} systems can be explained by longer sight lines through the halo of the host galaxy compared to the typically shorter sight lines through the halo of the Milky Way. Our statistical analysis shows that there is no obvious difference in the observed parameters of single-component and multi-component absorption systems. As described by \\citet{dingcharltonchurchill05}, multiple cloud components in absorption systems can be divided into two sub-classes. The ``kinematically spread'' absorbers show one or more dominant \\ion{Ca}{ii} absorbers and several weaker ones spread over a wide velocity range, as seen towards QSO 0109-3518 and QSO B1101-26 (Figures\\,\\ref{fig_he0151_J222756_q0109_q0002} and \\ref{fig_QSOB1212_QSOB1101_PKS1448-232_LBQS1229-0207}, online version). The ``kinematically compact'' subclass is characterised by multiple absorption components with comparable equivalent widths blended together and spread over less than about $100\\,\\mathrm{km\\,s}^{-1}$. An example of the latter sub-class is the absorber towards QSO J0003-2323, as displayed in Figure\\,\\ref{fig_he0001-2340_he0151_he1341_J092913}.  Important information about the distribution of neutral and weakly ionised extraplanar gas in the Milky Way is provided by column density distribution functions $f(N)$ of \\ion{Ca}{ii} and \\ion{H}{i} that we have derived in Section\\,\\ref{stat.properties}. We now can compare the properties of these distribution functions with results from QSO absorption-line studies and extragalactic \\ion{H}{i} surveys. The slope of the \\ion{Ca}{ii} column density distribution of Milky Way absorbers turns out to be $\\beta=-1.6 \\pm 0.3$. This is identical to the slope of $\\beta=-1.59 \\pm 0.05$ derived for strong \\ion{Mg}{ii} absorbers at intermediate reshifts ($z=0.4-1.2$), as presented by \\citet{churchillvogtcharlton03}. The equality of the slopes suggests that (despite somewhat different ionisation and dust-depletion properties) \\ion{Ca}{ii} and strong \\ion{Mg}{ii} absorbers probe similar gaseous structures that are located in the environment of galaxies.  Concerning neutral hydrogen, \\citet{petitjean93} have investigated the \\ion{H}{i} column density distribution function of intergalactic QSO absorption line systems at high redshift (mean redshift of $z \\approx 2.8$) with the help of high spectral resolution data. Their \\ion{H}{i} data span a range from $\\log (N_\\mathrm{HI}/\\mathrm{cm}^{-2})=12$ to $\\log (N_\\mathrm{HI}/\\mathrm{cm}^{-2})=22$. After they could show that a single power law with a slope of $\\beta=-1.49$ provides only a poor fit to the data they divided the sample into two subsamples (lower and higher than $10^{16}$\\,cm$^{-2}$). They found a slope of $\\beta=-1.83$ for the $\\log (N_\\mathrm{HI}/\\mathrm{cm}^{-2}) < 16$ sample and $\\beta=-1.32$ for the $\\log (N_\\mathrm{HI}/\\mathrm{cm}^{-2}) > 16$ sample. \\citet{petitjean93} point out that the overall \\ion{H}{i} distribution is more complex than had been thought and that even two power laws only poorly fit the data. In a later study, \\citet{kimcarswelletal02} indeed show that the slope of the \\ion{H}{i} column density distribution varies between $\\sim -1.4$ and $\\sim -2.0$, depending on the mean absorber redshift and the column-density interval used. Unfortunately, very little is known about the \\ion{H}{i} absorber distribution in the for us interesting range $\\log (N_\\mathrm{HI}/\\mathrm{cm}^{-2})=15\\ldots19$ and $z=0\\ldots0.5$ due to the lack of low-redshift QSO data in the UV band.  Although we have substantial uncertainties in the conversion of \\ion{Ca}{ii} into \\ion{H}{i} as discussed above, it is interesting that the slope of the \\ion{H}{i} column density distribution of $\\beta=-1.3 \\pm 0.1$, as indirectly derived from our \\ion{Ca}{ii} data for the $\\log (N_\\mathrm{HI}/\\mathrm{cm}^{-2})$ range between 19 and 22, is in general agreement with the statistics of low- and high-redshift QSO absorption line data. This implies that a significant fraction of high-column density (log $N>16$) intervening QSO absorbers are closely related to galaxies in a way similar as the intermediate- and high-velocity \\ion{Ca}{ii} absorbers are connected to the Milky Way. It appears that the column density distribution of extraplanar neutral gas structures (i.e., HVCs and their extragalactic analogues) is roughly universal at low and high redshift. This underlines the Overall importance of the processes that lead to the circulation of neutral gas in the environment of galaxies (e.g., fountain flows, gas accretion, tidal interactions) for the evolution of galaxies."
    },
    "0806/0806.4424_arXiv.txt": {
        "abstract": "The famous extreme solar and particle event of 20 January 2005 is analyzed from two perspectives. Firstly, using multi-spectral data, we study temporal, spectral, and spatial features of the main phase of the flare, when the strongest emissions from microwaves up to 200 MeV gamma-rays were observed. Secondly, we relate our results to a long-standing controversy on the origin of solar energetic particles (SEP) arriving at Earth, \\textit{i.e.}, acceleration in flares, or shocks ahead of coronal mass ejections (CMEs). Our analysis shows that all electromagnetic emissions from microwaves up to 2.22 MeV line gamma-rays during the main flare phase originated within a compact structure located just above sunspot umbrae. In particular, a huge ($\\approx 10^5$ sfu) radio burst with a high frequency maximum at 30 GHz was observed, indicating the presence of a large number of energetic electrons in very strong magnetic fields. Thus, protons and electrons responsible for various flare emissions during its main phase were accelerated within the magnetic field of the active region. The leading, impulsive parts of the ground-level enhancement (GLE), and highest-energy gamma-rays identified with $\\pi^0$-decay emission, are similar and closely correspond in time. The origin of the $\\pi^0$-decay gamma-rays is argued to be the same as that of lower-energy emissions, although this is not proven. On the other hand, we estimate the sky-plane speed of the CME to be 2\\,000\\,--\\,2\\,600 km\\thinspace s$^{-1}$, \\textit{i.e.}, high, but of the same order as preceding non-GLE-related CMEs from the same active region. Hence, the flare itself rather than the CME appears to determine the extreme nature of this event. We therefore conclude that the acceleration, at least, to sub-relativistic energies, of electrons and protons, responsible for both the major flare emissions and the leading spike of SEP/GLE by 07~UT, are likely to have occurred nearly simultaneously within the flare region. However, our analysis does not rule out a probable contribution from particles accelerated in the CME-driven shock for the leading GLE spike, which seemed to dominate at later stages of the SEP event. ",
        "introduction": "The solar eruptive-flare event of 20 January 2005 has attracted great attention from the solar and solar-terrestrial community due to its outstanding characteristics (see, \\textit{e.g.}, http:/\\negthinspace/creme96.nrl.navy.mil/20Jan05/). Despite its occurrence at the deep descending phase of the solar cycle, the event was characterized, in particular, by strong gamma-ray emission with a photon energy up to at least 200 MeV; by one of the strongest microwave bursts with a spectral maximum frequency $\\approx 30$~GHz, by a fast halo coronal mass ejection (CME), and was accompanied by a very hard-spectrum solar energetic particle (SEP) flux near Earth including the second largest ground-level enhancement (GLE) of cosmic ray intensity in observational history.   Many studies are devoted to this extreme event, but they mainly concern the cosmic ray aspect. The central points of these studies are the features and origin of the SEP/GLE. Additionally, a long-standing discussion over the origin of accelerated protons in solar events, \\textit{i.e.}, CME-driven shock-acceleration or flare-acceleration, has been strengthened. In particular, Gopalswamy \\textit{et al.} (2005) estimated that the CME had the largest sky-plane speed exceeding 3000 km\\thinspace s$^{-1}$ and concluded that the GLE had a shock origin. On the other hand, Simnett (2006, 2007) considered the relative timing of various manifestations in this event and several other factors, and stated that the CME was not responsible for the relativistic ion acceleration. Kuznetsov \\textit{et al.} (2005a,b, 2006a,b, 2007) also argued in favor of flare-related proton acceleration in this and other events.  These contradicting opinions reflect a long-standing controversy over different viewpoints of the origin of high-energy protons related to solar events [see Cliver (2000) for a historical review]. Many papers have been published in support of each of these opposing viewpoints. Some researchers insist that all energetic particles arriving at Earth are accelerated exclusively by CME-driven shocks, rather far from flare regions (Reames, 1999; Kahler, 2001). Some others argue that traveling shocks could perhaps contribute to lower-energy proton fluxes ($E < 10$ MeV), while higher-energy ones ($E > 100$ MeV) are flare-accelerated on the Sun (\\textit{e.g.}, Klein and Trottet, 2001; Chertok, 1995; Livshits and Belov, 2004; Li \\textit{et al.}, 2007a,b).  The 20 January 2005 event provides a unique chance to address this problem for the following reasons: (a)~availability of spectral-imaging gamma-ray data from the Reuven Ramaty High Energy Solar Spectroscopic Imager (RHESSI; Lin \\textit{et al.}, 2002), (b)~availability of highest-energy spectral gamma-ray data up to 200~MeV from the non-imaging SOlar Neutron and Gamma experiment (SONG; Kuznetsov \\textit{et al.}, 2004, 2008) on board the CORONAS-F spacecraft (Oraevsky and Sobelman, 2002), (c)~the strongest GLE in over half a century of observations, caused by the unusually prompt arrival of protons at Earth (Bieber \\textit{et al.}, 2005; Belov \\textit{et al.}, 2005; Simnett, 2006; Cliver, 2006), and (d)~availability of many other kinds of observational data, for both the flare region on the Sun, and for the near-Earth space, up to the Earth's surface. In particular, the source region on the Sun was observed in soft X-rays (SXR) by GOES-10 and GOES-12 including the Soft X-ray Imager (SXI; Hill \\textit{et al.}, 2005). The H$\\alpha$ flare was recorded at the IPS Culgoora Solar Observatory in Australia and at the Hiraiso Observatory in Japan. Nobeyama Radio Polarimeters (NoRP; Nakajima \\textit{et al.}, 1985; Torii \\textit{et al.}, 1979) recorded an unusually strong microwave burst at six frequencies: 2, 3.75, 9.4, 17, 35, and 80 GHz. The Transition Region and Coronal Explorer (TRACE; Handy \\textit{et al.}, 1999) observed the flare after its peak in the 1600 \\AA\\ channel. The Extreme-Ultraviolet Imaging Telescope (EIT; Delaboudini\\`ere \\textit{et al.}, 1995) and the Large-Angle Spectrometric Coronagraph (LASCO; Brueckner \\textit{et al.}, 1995) on SOHO also supplied images related to this event.  In this paper, we first analyze in detail the flare itself. We consider its source region on the Sun in different spectral domains, using available imaging data and reconciling time profiles of various emissions as well as spatial and spectral characteristics of their sources. In this way, we endeavor to understand this remarkable flare, its configuration, and the physical conditions in the flare region. Then, based on the results of this analysis, we consider which of the two above-mentioned models of the origin of high-energy protons better matches observations of the 20 January 2005 event. Our methodology is based on joint analysis of multi-spectral data. To minimize the influence of instrumental effects and model-dependent estimations, we verify our results and estimates by using different methods and independent observations as much as possible. Their agreement enhances the reliability of our results and conclusions. Is is important to analyze manifestations of accelerated protons primarily during their main peak (corresponding to the main phase of the highest rate of flare energy release), when the influence of subsequent post-eruptive processes and transport effects on the Sun and/or in interplanetary space, as well as probable contributions from different acceleration mechanisms, which are difficult to distinguish, is minimal.  We start our consideration in Section~\\ref{Observations_Analysis} with observational solar data and their brief analysis, including the development of the active region and pre-flare situation. Then, flare manifestations in gamma-rays, hard and soft X-rays, microwaves, H$\\alpha$, 1600~\\AA, and extreme-ultraviolet (EUV) emissions, as well as the CME are addressed. We conclude that the extreme features of the flare were due to its occurrence in the strong magnetic fields just above sunspot umbrae. Section~\\ref{SEP_GLE} is devoted to relating flare features to the observed SEP/GLE characteristics. In particular, its proton productivity is compared with other events from the same active region. Finally, the temporal parameters of the impulsive parts of the gamma-ray burst and GLE are compared. The same population of energetic particles appears to be responsible for both the $\\pi^0$-decay gamma-ray burst and the leading spike of the GLE. In turn, the $\\pi^0$-decay emission was close, but not identical to, emissions whose sources were located within the solar active region. In Section~\\ref{Discussion}, we discuss the results of our analysis and argue that the $\\pi^0$-decay emission also originated in the flare region and that data available to us appear to favor the flare-related initial acceleration of particles responsible for the leading GLE spike, but do not preclude their origin in a CME-driven shock. Section~\\ref{Summary} summarizes the results of our analyses of the flare and origin of SEP/GLE particles and briefly addresses their implications. ",
        "conclusions": "\\label{Summary}  In considering the extreme proton event of 20 January 2005, we first carried out a comprehensive analysis of the flare, based on multi-spectral observations, and second, investigated the origin of the energetic protons corresponding to the leading SEP/GLE spike. Our conclusions regarding the flare are as follows.  \\begin{enumerate}  \\item Imaging data in H$\\alpha$, 1600~\\AA, soft and hard X-rays, EUV, and 2.22 MeV neutron-capture line show that the sources of these emissions were compact and localized within a volume confined by flare ribbons, which crossed the umbrae of the largest sunspots with strong magnetic fields.  \\item Non-imaging data extend this conclusion to microwave/millimeter, lower-energy, and medium-energy gamma-ray emissions, because the time profiles of the major flare component, caused by accelerated electrons and protons, were similar and closely related, particularly in time.   \\item The $\\pi^0$-decay emission time profile was much like those of lower energy, but their differences prevent us from confidently asserting a common source. Nevertheless, co-temporal structural details in the $\\pi^0$-decay and lower-energy time profiles (their magnitude corresponds to a continuous spectrum) indicate their common nature; the timing of the $\\pi^0$-decay gamma-ray burst corresponds to the highest energy release rate in the flare (Section~\\ref{where_from}). Thus, its flare origin appears to be preferable. However, one should not consider the flare origin to be proven.  \\item The above considerations favor the acceleration of both the electrons and protons, responsible for various emissions (at least, at moderate energies), within the magnetic field of the active region at a moderate height above the sunspots. The manifestations of the electron and proton components correspond to the standard picture of a flare, in which chromospheric ribbons and compact sources of hard emissions map the footpoints of newly formed magnetic loops, along which energetic particles fly downward and bombard the lower atmosphere.  \\item The presence of a large number of accelerated particles in strong magnetic fields is demonstrated by the fact that the flare was accompanied by one of the strongest and hardest-spectrum microwave bursts ever observed, with a peak flux density of almost $10^{5}$ sfu at 30 GHz and a very strong ($5\\times 10^{4}$ sfu) emission at 80 GHz.  \\item Evaluations of the 20 January CME sky-plane speed from various available data result in an estimate of 2000\\,--\\,2600 km\\thinspace s$^{-1}$.  \\end{enumerate}  Regarding the SEP/GLE aspect of this event, we first note that there is no clear observational evidence to state with certainty if it was due to the flare or the CME-driven shock. However, our analysis favors the flare region as the probable site of acceleration of the particles responsible for the leading SEP/GLE spike, rather than the shock ahead of the CME, as supported by the following arguments.  \\begin{enumerate}  \\item The compactness of the gamma-ray sources and similarity of the time profiles of different emissions could hardly be expected if high-energy particles were accelerated in an extended shock front ahead of the receding CME.   \\item The proton event under consideration was extreme, not in its CME speed, but in the parameters of the flare itself, particularly in the huge microwave/millim\\-eter radio burst.   \\item There is a significant correspondence between the hardness of the microwave spectrum and the hardness of the energetic spectrum of the SEP and GLE protons.   \\item The time of emission from the Sun of high-energy protons responsible for the onset of the GLE is close to that of the $\\pi^0$-decay emission.   \\item The gamma-ray burst and the leading GLE spike have similar shapes.  \\item The decay times of gamma-ray emissions computed by G.H.~Share (2007, private communication) from RHESSI data agree with the collision times of protons trapped in the observed flare loop, but not with those stored higher in the corona.  \\end{enumerate}  The available data, together with our results, do not rule out a possible role for the CME-driven shock in the acceleration of particles during the initial part of the GLE. Similarly, the role of flare acceleration cannot be completely ruled out for later stages of the SEP event, when the shock seems to dominate. Advocating ``shock-acceleration only'' or ``flare-acceleration only'' does not seem to be productive (see {\\v S}vestka, 2001; Kallenrode, 2003). Studies of the relative contribution from acceleration in flares and CME-driven shocks appears to be more fruitful. They can vary for dissimilar events, their different stages, and in different energy ranges.  The 20 January 2005 event represents a distinct class of extreme solar flares. They amount to a small percentage of all flares, but probably constitute the majority of events which can be proton-rich under favorable Sun\\,--\\,Earth connections. The expected features of these flares are: occurrence near or above sunspot umbrae in strong magnetic fields; powerful bursts in microwaves; and at long millimeter wavelengths. These events are very dangerous due to a high probability of strong proton fluxes with hard spectra. This fact highlights the importance of measuring strong magnetic fields in solar active regions, as well as patrol observations of the total radio flux at long millimeter wavelengths, for the forecast and diagnosis of major proton events. Currently, only NoRP observations at 35 and 80~GHz are available from $\\approx\\,$22 till $\\approx\\,$08~UT."
    },
    "0806/0806.3497_arXiv.txt": {
        "abstract": "\\noindent If general relativity (GR) describes % the expansion of the % Universe, the observed cosmic acceleration implies the existence of a % `dark energy'. However, % while the Universe is on average % homogeneous on large scales, it is inhomogeneous on smaller scales. While GR governs % the dynamics of the \\emph{in}homogeneous Universe, % the averaged % \\emph{homogeneous} Universe obeys modified Einstein equations. Can such % modifications alone % explain the acceleration? For a simple generic model with realistic initial conditions, we show the answer to be `no'. Averaging effects negligibly influence the cosmological dynamics. % ",
        "introduction": " ",
        "conclusions": ""
    },
    "0806/0806.2458_arXiv.txt": {
        "abstract": "We try to constrain the noncommutativity length scale of the theoretical model given in \\cite{cmbpaper} using the observational data from ACBAR, CBI and five year WMAP. The noncommutativity parameter is not constrained by WMAP data, however ACBAR and CBI data restrict the lower bound of its energy scale to be around $10$ TeV. We also derive an expression for the amount of non-causality coming from spacetime noncommutativity for the fields of primordial scalar perturbations that are space-like separated. The amount of causality violation for these field fluctuations are direction dependent. ",
        "introduction": "In 1992, the Cosmic Background Explorer (COBE) satellite detected anisotropies in the CMB radiation, which led to the conclusion that the early universe was not smooth: there were small density perturbations in the photon-baryon fluid before they decoupled from each other. Quantum corrections to the inflaton field generate perturbations in the metric and these perturbations could have been carried over to the photon-baryon fluid as density perturbations.  We then observe them today in the distribution of large scale structure and anisotropies in the CMB radiation.  Inflation \\cite{Starobinsky79, Starobinsky82, Guth81, Linde82, Albrecht82} stretches a region of Planck size into cosmological scales. So, at the end of inflation, physics at the Planck scale can leave its signature on cosmological scales too. Physics at the Planck scale is better described by models of quantum gravity or string theory. There are indications from considerations of either quantum gravity or string theory that spacetime is noncommutative with a length scale of the order of Planck length. CMB radiation, which consists of photons from the last scattering surface of the early universe can carry the signature of spacetime noncommutativity. With these ideas in mind, in this paper, we look for a constraint on the noncommutativity length scale from the WMAP5 \\cite{WMAP1, WMAP2, WMAP3}, ACBAR \\cite{ACBAR1, ACBAR2, ACBAR3} and CBI \\cite{CBI1, CBI2, CBI3, CBI4, CBI5} observational data.  In a noncommutative spacetime, the commutator of quantum fields at space-like separations does not in general vanish, leading to violation of causality. This type of violation of causality in the context of the fields for the primordial scalar perturbations is also discussed in this paper. It is shown that the expression for the amount of causality violation is direction-dependent.  In \\cite{Sachin}, it was shown that causality violation coming from noncommutative spacetimes leads to violation of Lorentz invariance in certain scattering amplitudes. Measurements of these violations would be another way to put limits on the amount of spacetime noncommutativity.  This paper is a sequel to an earlier work \\cite{cmbpaper}. The latter explains the theoretical basis of the formulae used in this paper. In \\cite{Queiroz1} another approach of noncommutative inflation is considered based on target space noncommutativity of fields \\cite{Queiroz2}. ",
        "conclusions": "The power spectrum becomes direction dependent in the presence of spacetime noncommutativity, indicating a preferred direction in the universe. We tried a best-fit of the theoretical model in \\cite{cmbpaper} with the WMAP data and saw that to improve the bound on $H\\theta$, we need data at higher $l$. (The last data point for WMAP is at $l=839$.) We therefore conclude that the WMAP data do not constrain $H\\theta$. We also see that tighter error bars at these higher $l$ will also help constrain the noncommutativity parameter. The small-scale CMB data like ACBAR and CBI give the CMB power spectrum for larger multipoles and hence may be better suited for the determination of $H\\theta$. ACBAR+CBI data only restrict $H\\theta$ to $H\\theta < 0.01$ Mpc and do not indicate whether the best fit is at $H \\theta = 0$ Mpc or some small non-zero value. However, this restriction corresponds to a lower bound for the energy of $\\theta$ of around $10$ TeV.  Further work is needed before rejecting the initial hypothesis that the other parameters of the $\\Lambda$CDM cosmology are unaffected by noncommutivity. It requires performing a full MCMC study of all seven parameters.  Also, we have shown the the existence and direction-dependence of non-causality coming from spacetime noncommutativity for the fields describing the primordial scalar perturbations when they are space-like separated. We see that the amount of causality violation is maximum when the two vectors, $\\vec{\\theta}^{0}$ and ${\\bf r}={\\bf x}_{1}-{\\bf x}_{2}$, are aligned. Here ${\\bf r}$ is the relative spatial coordinate of the fields at spatial locations ${\\bf x}_{1}$ and ${\\bf x}_{2}$."
    },
    "0806/0806.0494_arXiv.txt": {
        "abstract": "{Deep surveys planned as a Key Science Project of LOFAR provide completely new opportunities for gravitational lens searches. For the first time do large-scale surveys reach the resolution required for a direct selection of lens candidates using morphological criteria. We briefly describe the strategies that we will use to exploit this potential. The long baselines of an international \\emph{E}-LOFAR are essential for this project. }  \\FullConference{From planets to dark energy: the modern radio universe\\\\ October 1-5 2007\\\\ University of Manchester, Manchester, UK}  \\begin{document} ",
        "introduction": "Gravitational lensing is the best method to determine the mass distribution of distant galaxies with high accuracy. With a large sample of lens systems at a range of redshifts, we can study not only the structure but also the evolution of galaxies. We identify three main topics, in which lensing can help to resolve controversial problems. \\begin{itemize}\\parskip0pt\\topsep0pt\\partopsep0pt\\parsep0pt\\itemsep0pt \\item Most galaxies at moderate redshifts seem to have very close to isothermal ($\\rho\\propto r^{-2}$) mass distributions. Why is this the case, and does it also hold for higher redshifts? \\item How does the \\emph{central} density profile of galaxies look like? Do they have cores or cusps, how do masses of central black holes evolve? \\item Are small sub-halos as abundant as predicted by the CDM structure formation scenario? \\end{itemize}  \\begin{figure}[hb] \\hspace*{0.2\\textwidth}% \\includegraphics[height=0.2\\textwidth]{claeskens}% \\hfill% \\includegraphics[height=0.2\\textwidth]{8oclock} \\hspace*{0.2\\textwidth} \\caption{Two examples from the bright end of typical lens systems expected to be found in LOFAR surveys. Left: RXS\\,J1131--1231 \\citep[HST image from][]{claeskens06}. Right: The `8 o'clock arc' \\citep{allam07}. WSRT 5\\,GHz contours superimposed upon SDSS image. } \\label{fig:1131+8uur} \\end{figure}    We can help answering these questions by modelling the mass distributions to fit for the structure of gravitationally lensed sources. In order to obtain information about the complete mass distribution, lensed \\emph{extended} sources like the ones shown in Fig.~\\ref{fig:1131+8uur} should be used and modelled with LensClean \\citep{lc1}. There are several advantages in using radio observations. Most important are the wide range of resolutions that can be achieved with radio arrays and the absence of effects like dust extinction and microlensing. So far we know less than 50 radio lenses, whilst the number of optical lenses has grown well above that and is increasing further. ",
        "conclusions": ""
    },
    "0806/0806.0388_arXiv.txt": {
        "abstract": "In most high energy cosmic ray surface arrays, the primary energy is currently determined from the value of the lateral distribution function at a fixed distance from the shower core, $r_{0}$. The value of $r_{0}$ is mainly related to the geometry of the array and is, therefore, considered as fixed independently of the shower energy or direction. We argue, however, that the dependence of $r_{0}$ on energy and zenith angle is not negligible. Therefore, in the present work we propose a new characteristic distance, which we call $r_{opt}$, specifically determined for each individual shower, with the objective of optimizing the energy reconstruction. This parameter may not only improve the energy determination, but also allow a more reliable reconstruction of the shape and position of rapidly varying spectral features. We show that the use of a specific $r_{opt}$ determined on a shower-to-shower basis, instead of using a fixed characteristic value, is of particular benefit in dealing with the energy reconstruction of events with saturated detectors, which are in general a large fraction of all the events detected by an array as energy increases. Furthermore, the $r_{opt}$ approach has the additional advantage of applying the same unified treatment for all detected events, regardless of whether they have saturated detectors or not. ",
        "introduction": "The energy spectrum of cosmic rays (CR) is observed from energies below 1 GeV up to more than $10^{20}$ eV. CR hit the top of the atmosphere at a rate $\\gtrsim 10^{3}$ per square meter per second. For energies below  $\\sim 10^{2}$ TeV, CR can be detected by direct measurement from high altitude balloons or satellites, albeit with rapidly decreasing statistics since the CR spectrum is very steep, $\\sim E^{-2.7}$. For $E \\gtrsim 10^{2}$ TeV, the point of maximum development of the cascades of daughter particles initiated by collisions with atmospheric nuclei starts reaching the ground at high altitudes. From that point onward, the properties of primary cosmic ray can be determined indirectly from the measurement of extensive air showers (EAS).  Two different experimental approaches have been traditionally used to study the highest energy EAS. The first one consists on the inference of lateral distribution of particles from the discrete sampling of the shower front at ground level by a surface array of detectors. Scintillators (e.g., Volcano Ranch, AGASA, KASCADE) and water Cherenkov detectors (e.g., Haverah Park) have been mainly used for this purpose. The second technique consists on the reconstruction of the longitudinal profile of the shower from the fluorescence light produced by atmospheric Nitrogen as it is excited by charged EAS particles along the atmosphere (e.g., Fly's Eye, HiRes). The latter is considered to be close to a calorimetric measurement of the primary CR particle energy. Extensive reviews on CR theory and experiments can be found in \\cite{Nagano,HillasReview,Unger}.  A special case from the experimental point of view is the Pierre Auger Observatory \\cite{Auger} which pioneers the simultaneous use of both techniques, water Cherenkov detectors and fluorescence telescopes. For these {\\it hybrid} events, systematic errors in their energy estimate are greatly reduced. Unfortunately, fluorescence can only be observed during dark nights and, consequently, this technique can only be applied to $\\sim 13$\\% of the incoming events. Therefore, even if the hybrid technique, when simultaneously available, allows for an independent calibration of a ground detector, high statistics must come from surface arrays with a duty cycle close to 100\\%.  In this work we are interested on the highest energies, $E > 10^{17}$ eV, at which point extragalactic flux is likely to penetrate the Galaxy and start dominating the CR flux. At these energies at least three very important spectral features are located: the second knee, the ankle and the GZK complex (bump plus steep flux suppression) \\cite{GMTancoPeVZeV}. The determination of their exact position and shape is a fundamental experimental problem. Therefore, the most relevant parameter in this spectral region is, arguably, the primary energy of the impinging CR.  The procedure to determine the primary energy in sufrace arrays is a two step process. First, the lateral distribution function (LDF), i.e. the shower particle density or signal  versus distance to the shower axis, is fitted assuming a known functional form. This fit suffers from uncertainties related to the statistical shower fluctuations, the uncertainties in core location and the ignorance of the exact form of the LDF. The normalization constant of the LDF of an extensive air shower is a monotonous (almost linear) increasing function of the energy of the primary cosmic ray. Therefore, Hillas \\cite{Hillas} proposed to use the interpolated signal at some fixed, {\\it characteristic distance\\/} from the shower core, $S(r_{0})$, at which fluctuations in the LDF are minimal. The uncertainty due to the lack of knowledge of the LDF is also minimized by this procedure \\cite{AugerLDF}.  The use of the signal interpolated at $r_{0}$, $S(r_0)$ is widely used as energy estimator by surface detector arrays. AGASA \\cite{AGASALDF,AGASADai}, Yakutsk \\cite{Yakutsk_ropt} and Haverah Park \\cite{HP_ropt}, for example, choose $r_{0}=600\\; m$, while Auger uses 1000 m due to its larger array spacing \\cite{AugerSpectrumPRL}. The characteristic distance $r_{0}$ is mainly, although not completely, determined by the geometry of the array. Thus, the same value of $r_{0}$ is used to estimate the energy for all the showers, regardless of primary energy or incoming direction. In the second step, there are at least two possible approaches to calibrate $S(r_0)$ as a function of primary energy: either via Monte Carlo simulations or, as in the case of Auger, by using the almost calorimetric measurement obtained from the fluorescence observation of high quality hybrid events \\cite{AugerSpectrumPRL}.   As an alternative, but motivated by Hillas' original idea \\cite{Hillas}, in the present work we focus in the shower-to-shower determination of an {\\it optimal} distance to the core, which we name hereafter $r_{opt}$, at which the interpolation of the {\\it reconstructed} signal is the best energy estimator for each individual shower, regardless of whether this point is actually the one that minimizes shower to shower LDF fluctuations.  We perform a detailed study of $r_{opt}$ as function of array spacing, primary energy and the zenith angle of the incoming cosmic ray and demonstrate that, although array geometry is an important underlying factor, the dependence of $r_{opt}$ on the remaining parameters is not negligible. We study the bias associated with both techniques, $r_{0}$ and $r_{opt}$, and show that, if the dynamical range of the detector covers a wide interval of energies, it is much safer to estimate an $r_{opt}$ for the energy reconstruction of each individual event than to fix a single $r_{0}$ for the whole data set. In fact, not only the bias as a function of energy can be kept negligible over at least 2.5 decades in energy, but also the error distribution functions are much better behaved, i,e, without appreciable kurtosis or skewness and very much Gaussian in the mentioned energy range. The latter has a potential impact in the reconstruction accuracy of the energy spectrum. We demonstrate this by applying a fixed $r_{0}$ as well as a shower-to-shower $r_{opt}$, to a simplified version of the actual energy spectrum between $\\sim 1$ and $\\sim 100$ EeV.  A further advantage of the $r_{opt}$ approach is the straightforward treatment of events with saturated detectors. The problem of saturation is very common in all surface experiments, specially when dealing with high energy vertical showers. In fact, at the highest energies inside the designed dynamical range of any experiment, usually events with saturated detectors can account for a large, if not dominant, fraction of all the observed events. Different strategies have been used to deal with them. In some cases saturated detectors are directly discarded from the LDF fit, while in others the saturation value is used as a lower limit to the true signal during the fitting procedure. The Auger Collaboration is developing at present special, more sophisticated algorithms to estimate the signal of a saturated detector \\cite{MarisThesis} in order to more properly account for them in the LDF fit. We show here that it is actually not possible to define a single characteristic $r_0$ distance for both kinds of events. In fact, even if well defined medians values of $r_{opt}$ for events with and without saturated detectors do exist, the dispersions around the median at any energy are so large that both sets cannot be clearly differentiated as to use, for example, just two fixed distances instead of a single one. Nevertheless, using a shower-to-shower $r_{opt}$ distance, the inferred energy is unbiased for events with and without saturated detectors alike. This reconstruction strategy allows for an homogeneous treatment of the data set regardless of the increasing number events with saturated detectors when the energy increases.    In a recent work \\cite{Newton}, Newton and co-workers also estimated an optimal shower-to-shower distance, but used a different algorithm and with a somewhat different scope. They were mainly concerned with demonstrating the existence of a single distance for any given shower at which fluctuations in the LDF are minimum. By assuming that such fluctuations can be well described by the fluctuations of just one parameter, the slope of the LDF, externally fed into their procedure, and using a combination of simulations and semi-analytical analysis, they claim that, regardless of the functional form of the LDF considered, there exist a {\\it convergence} point of the LDFs, at a characteristic distance they call optimal, where shower-to-shower fluctuations are minimal. Their results, combined together for a mix of energies drawn from a flat spectrum, seem to support their claim and lead them to the conclusion that a single fixed distance, depending only on the geometry and spacing of a given array, would be a good choice for the energy determination in the whole energy range of the experiment. Furthermore, it is not clear from their study how to deal with the events with saturated detectors in the later scenario.  Alternatively, in the present work we do not constrain the parameters of the LDF, which are an output of the fit to the simulated data. We introduce instead reliable error estimations for the reconstruction of the core position, as calculated by \\cite{M.C.Medina} for arrays of varying spacing as a function of energy. Furthermore, our final scope is the determination of energy all along the dynamical range of an experiment, and not the study of the manifestation of signal fluctuations in the LDF. Therefore, we study in detail the dependence of $r_{opt}$ and of its distribution function as a function of energy, zenith angle and array spacing. This study is performed for events with and without saturated detectors. We also give a comparative description of error and biases for the fix distance and the $r_{opt}$ distance approaches in that parameter space. In the same line, we further extend our analysis to the reconstruction of a simulated energy spectrum of known shape, and show what the potential effects are of using each technique.  The paper is organized as follows. Section 2 describes our general algorithm. Two different detector arrays are considered, scintillators and water Cherenkov tanks. In Section 3, in order to study the $r_{opt}$ dependencies with array spacing and the energy and incoming direction of primary cosmic ray, water Cherenkov (Auger-like) stations have been used. In Section 4 we deal with the issue of energy determination. In that analysis, (AGASA-like) scintillators are considered. A general discussion and conclusions are given in Section 5. While different detectors are used in Sections 3 and 4, the algorithm to find $r_{opt}$ is the same for both and the results and conclusions of the paper are not affected by the array under consideration. ",
        "conclusions": "The primary CR energy is generally estimated in surface arrays by interpolating the lateral distribution function of particles in the shower front at ground level at a fixed optimum distance $r_{0}$ from the shower core. This parameter is assumed to be predominantly dependent on the detector separation distance for a given layout geometry and, therefore, is considered as a constant for a given array.  In this work we propose an algorithm to evaluate an equivalent, but shower-to-shower optimal distance, which we call $r_{opt}$. We have performed a thorough analysis of the dependence of $r_{opt}$ on energy and zenith angle, and demonstrate that, contrary to reference \\cite{Newton}, these are not negligible factors. In fact, not taking into account an event-specific $r_{opt}$, produce wider error distribution functions that can even affect the reconstruction of a highly structured, rapidly varying spectrum. The shower-to-shower $r_{opt}$ approach, on the other hand, is an unbiassed estimator of the CR primary energy, which produce also narrower, symmetric, almost Gaussian error distribution functions for energy reconstruction. Those properties of $r_{opt}$ can additionally lead to much more reliable spectral reconstruction. The differences emerging from the two procedures, $r_{0}$ vs. $r_{opt}$, when applied to spectral reconstruction may have astrophysical implications, specially in the coming era of improved precision.  An important aspect of the $r_{opt}$ approach is that it has the additional advantage of allowing the same unified treatment for events with and without saturated detectors; something that, in the $r_{0}$ approach is generally not possible, requiring either the selection of events through quality cuts, or the separate reconstruction with different techniques of the two types of events. Since the fraction of events presenting saturation is a rapidly increasing function of energy, the later greatly reduces the effective energy range for spectral reconstruction in almost all practical situations.  For practical application to real experiments, however, a proper calibration curve should be deduced specifically for $r_{opt}$, which would further optimize it as an energy estimator."
    },
    "0806/0806.2805_arXiv.txt": {
        "abstract": "\\vspace*{.5mm}\\noindent A modified-gravity theory is considered with a four-form field strength $F$, a variable gravitational coupling parameter $G(F)$, and a standard matter action. This theory provides a concrete realization of the general vacuum variable $q$ as the four-form amplitude $F$ and allows for a study of its dynamics. The theory gives a flat Friedmann--Robertson--Walker universe with rapid oscillations of the effective vacuum energy density (cosmological ``constant''), whose amplitude drops to zero asymptotically. Extrapolating to the present age of the Universe, the order of magnitude of the average vacuum energy density agrees with the observed near-critical vacuum energy density of the present universe. It may even be that this type of oscillating vacuum energy density constitutes a significant part of the so-called cold dark matter in the standard Friedmann--Robertson--Walker framework. ",
        "introduction": "\\label{sec:introduction}  In a previous article~\\cite{KlinkhamerVolovik2008}, we proposed to characterize a Lorentz-invariant quantum vacuum  by a nonzero conserved relativistic ``charge'' $q$. This approach  allowed us to discuss the \\emph{thermodynamics} of the quantum vacuum, in particular, thermodynamic properties as stability and compressibility. We found that the vacuum energy density appears in two guises.  The \\emph{microscopic} vacuum energy density is characterized by an ultraviolet energy scale, $\\epsilon(q)\\sim E^4_\\text{UV}$. For definiteness, we will take this energy scale $E_\\text{UV}$ to be close to the Planck energy scale $E_\\text{Planck}\\equiv \\sqrt{\\hbar\\, c^5/G_\\text{N}} \\approx 1.22 \\times 10^{19}\\,\\text{GeV}$. The \\emph{macroscopic} vacuum energy density is, however, determined by a particular thermodynamic quantity, \\mbox{$\\widetilde{\\epsilon}_\\text{vac}(q) \\equiv \\epsilon - q\\,d\\epsilon/dq$,} and it is this energy density that contributes to the effective gravitational field equations at low energies. For a self-sustained vacuum in full thermodynamic equilibrium and in the absence of matter, the effective (coarse-grained) vacuum energy density $\\widetilde{\\epsilon}_\\text{vac}(q)$ is automatically nullified (without fine tuning) by the spontaneous adjustment of the vacuum variable $q$ to its equilibrium value $q_0$, so that $\\widetilde{\\epsilon}_\\text{vac}(q_0)=0$. This implies that the effective cosmological constant $\\Lambda$ of a perfect quantum vacuum is strictly zero, which is consistent with the requirement of Lorentz invariance.   The presence of thermal matter makes the vacuum state Lorentz noninvariant and leads to a readjustment of the variable $q$ to a new equilibrium value, $q_0^\\prime = q_0 + \\delta q$, which shifts the effective vacuum energy density away from zero, $\\widetilde{\\epsilon}_\\text{vac}(q_0 + \\delta q) \\ne 0$. The same happens with other types of perturbations that violate Lorentz invariance, such as the existence of a spacetime boundary or an interface. According to this approach, the present value of $\\widetilde{\\epsilon}_\\text{vac}$ is nonzero but small because the universe is close to equilibrium and Lorentz-noninvariant perturbations of the quantum vacuum are small (compared with the ultraviolet scale which sets the microscopic energy density $\\epsilon$).   The situation is different for Lorentz-invariant perturbations of the vacuum, such as the formation of scalar condensates as discussed in Ref.~\\cite{KlinkhamerVolovik2008} or quark/gluon condensates derived from quantum chromodynamics (cf. Ref.~\\cite{Shifman-etal1978}). In this case, the variable $q$ shifts in such a way that it completely compensates the energy density of the perturbation and the effective cosmological constant is again zero in the new Lorentz-invariant equilibrium vacuum.   The possible origin of the conserved vacuum charge $q$ in the perfect Lorentz-invariant quantum vacuum was discussed in Ref.~\\cite{KlinkhamerVolovik2008}  in general terms. But a specific example was also given in terms of  a four-form field strength $F$~\\cite{DuffNieuwenhuizen1980,Aurilia-etal1980,Hawking1984+Duff+Wu, DuncanJensen1989,BoussoPolchinski2000,Aurilia-etal2004}. Here, we use this explicit realization with a four-form field $F$ to study the \\emph{dynamics} of the vacuum energy, which describes the relaxation of the vacuum energy density $\\widetilde{\\epsilon}_\\text{vac}$ (effective cosmological ``constant'') from its  natural Planck-scale value at early times to a naturally small value at late times. In short, the present cosmological constant is small because the Universe happens to be old.\\footnote{An extensive but nonexhaustive list of references to research papers and reviews on the so-called ``cosmological constant problem(s)'' can be found in Ref.~\\cite{KlinkhamerVolovik2008}. A recent review on cosmic ``dark energy''  is given in Ref.~\\cite{Frieman-etal2008}.}  The results of the present article show that, for the type of theory considered, the decay of $\\widetilde{\\epsilon}_\\text{vac}$ is accompanied by rapid oscillations of the vacuum variable $F$ and that the relaxation of $\\widetilde{\\epsilon}_\\text{vac}$ mimics the behavior of cold dark matter (CDM) in a standard Friedmann--Robertson--Walker (FRW) universe. This suggests that part of the inferred CDM may come from dynamic vacuum energy density and may also give a clue to the solution of the so-called coincidence problem \\cite{Frieman-etal2008}, namely, why the approximately constant vacuum energy density is precisely now of the same order as the time-dependent CDM energy density.   These results are obtained by the following steps. In Sec.~\\ref{sec:Gravity-with-F-field}, a modified-gravity theory with a four-form field $F$ is defined in terms of general functions for the microscopic energy density $\\epsilon(F)$ and variable gravitational coupling parameter $G(F)$. In Sec.~\\ref{sec:de-Sitter-expansion}, the dynamics of the corresponding de-Sitter universe without matter is discussed and, in Sec.~\\ref{sec:FRW universe}, the dynamics of a flat FRW universe with matter, using simple \\emph{Ans\\\"{a}tze} for the functions $\\epsilon(F)$ and $G(F)$. In Sec.~\\ref{sec:Equilibrium-approach}, the approach to equilibrium in such a FRW universe is studied in detail and the above mentioned vacuum oscillations are established. In Sec.~\\ref{sec:Conclusion}, the main results are summarized. ",
        "conclusions": "\\label{sec:Conclusion} \\vspace*{0mm}  The considerations of the present article and its predecessor~\\cite{KlinkhamerVolovik2008} by no means solve the cosmological constant problems, but may provide hints. Specifically, the new results are \\begin{itemize} \\vspace*{0mm}\\item[(i)] a  mechanism of vacuum-energy decay, which, starting from a ``natural'' Planck-scale value at very early times, leads to the correct order of magnitude \\eqref{eq:rhoVaverage-now} for the present cosmological constant; \\vspace*{0mm}\\item[(ii)] the realization from result \\eqref{eq:y+h+eV-Asymptotes} that a substantial part of the inferred CDM may come from an oscillating vacuum energy density; \\vspace*{0mm}\\item[(iii)] the important role of oscillations of the vacuum variable $q$ (here, $F$), which drive the vacuum energy density oscillations responsible for the first two results. \\vspace*{-0mm} \\end{itemize} Expanding on the last point, another consequence of $q$ oscillations is that they naturally lead to the creation of hot (ultrarelativistic) matter from the vacuum. This effective mechanism of energy exchange between vacuum and matter deserves further study.  \\vspace*{0mm}"
    },
    "0806/0806.2728_arXiv.txt": {
        "abstract": "High-resolution observations provide evidence about the existence of small-amplitude transverse oscillations in solar filament fine structures. These oscillations are believed to represent fast magnetohydrodynamic (MHD) waves and the disturbances are seen to be damped in short timescales of the order of 1 to 4 periods. In this Letter  we propose that, due to the highly inhomogeneous nature of the filament plasma at the fine structure spatial scale, the phenomenon of resonant absorption is likely to operate in the temporal attenuation of fast MHD oscillations. By considering transverse inhomogeneity in a straight flux tube model we find that, for density inhomogeneities typical of filament threads, the decay times are of a few oscillatory periods only. ",
        "introduction": "Quiescent solar filaments form along the inversion polarity line or between the weak remnants of active regions. Early filament observations (\\citealt{Engvold98}) as well as recent high-resolution H$\\alpha$ observations, obtained with the Swedish Solar Telescope (SST) in La Palma (\\citealt{Lin05}) and the Dutch Open Telescope (DOT), have revealed that their fine structure is composed by many horizontal and thin dark threads. The measured average width of resolved threads is about $0.3$ arc sec ($\\sim$ $210$ km) while the length is between $5$ and $40$ arc sec ($\\sim$ 3500 - 28,000  km). They seem to be partially filled with cold plasma \\citep{Lin05}, typically two orders of magnitude denser than that of the corona, and it is generally assumed that they outline their magnetic flux tubes (\\citealt{Engvold98,Linthesis,Lin05,Martin08}). This idea is strongly supported by the fact that they are inclined with respect to the filament long axis at a similar angle to what has been found for the magnetic field (\\citealt{Leroy80,Bommier94,BL98}).    Small amplitude oscillations have been observed in filaments and it is well established that these periodic changes are of local nature. The detected peak velocity ranges from the noise level (down to 0.1 km s$^{-1}$ in some cases) to 2--3~km s$^{-1}$, although larger values have also been reported. Two-dimensional observations of filaments by \\citet{YiEngvold91} and \\citet{Yi91} revealed that individual threads or groups of threads may oscillate independently with their own periods, which range between 3 and 20 minutes. More recently, \\citet{Linthesis} reports that spatially coherent oscillations are found over slices, with an area of $1.4 \\times 54$ arc sec$^2$, of a polar crown filament, and that among other, a significant periodicity at 26 minutes, strongly damped after 4 periods, appears.  Furthermore, \\citet{Lin07} have shown evidence about traveling waves along a number of filament threads with an average phase velocity of $12$ \\ km s$^{-1}$, a wavelength of $4''$ ($\\sim 2800$ \\ km), and oscillatory periods of the individual threads that vary from $3$ to $9$ minutes. The observed periodic signals are obtained from Doppler velocity measurements and can therefore be associated to the transverse displacement of the fine structures.      The observed small amplitude oscillations  have been interpreted in terms of magnetohydrodynamic (MHD) waves (\\citealt{OB02}) and theoretical models have been developed (see \\citealt{Ballester05,Ballester06}, for recent reviews). \\citet{diaz02} modeled a prominence thread as a straight cylindrical flux tube with a cool region representing the filament material, confined by two symmetric hot regions. They found that the fundamental fast mode is  always confined in the dense part of the flux tube, hence, for an oscillating cylindrical filament thread, it should be difficult to induce oscillations in adjacent threads, unless they are very close.   The time damping of prominence oscillations has been unambiguously determined in some observations. Reliable values for the damping time have been derived, from different Doppler velocity time series, by \\citet{Molowny-Horas99} and \\citet{Terradas02}, in prominences, and by \\citet{Linthesis}, in filaments.  The values thus obtained are usually between 1 and 4 times the corresponding period, and large regions of prominences/filaments display similar damping times. The damping of perturbations is probably a common feature of filament oscillations, hence theoretical mechanisms must be explored and their damping time scales should be compared with those obtained from observations.  Linear non-adiabatic MHD waves have been proposed  as a potential mechanism to explain the observed attenuation time scales (\\citealt{Carbonell04,Terradas01,Terradas05,Soler07,Soler08}).  Using thermal mechanisms only slow waves can be damped in an efficient manner while fast waves remain almost undamped.  Ion-neutral collisions provide a possible mechanism to damp fast waves (as well as Alfv\\'en waves) that is able to reproduce observed damping times for given parameter values, in particular for a quasi-neutral gas (\\citealt{Forteza07}).   Apart from the mentioned non-ideal damping mechanisms, there is another possibility to attenuate fast waves in thin filament threads. The phenomenon of resonant wave damping is well documented for fast kink waves in coronal loops (see e.g.\\ \\citealt{GAA06,goossens08}, for recent reviews) and provides a plausible explanation to quickly damped transverse loop oscillations observed by TRACE (\\citealt{Nakariakov99,Aschwanden99}). In this Letter, we address the resonant damping mechanism in the context of filament thread oscillations and assess its relevance in explaining the observed attenuation time scales. ",
        "conclusions": "\\label{conclusion}  In this Letter we have shown that due to the highly inhomogeneous nature of filaments  at their transverse scales the process of resonant absorption is an efficient damping mechanism for fast MHD oscillations propagating in these structures. The relevance of the mechanism has been assessed in a flux tube model, with the inclusion of transverse inhomogeneity in the Alfv\\'en velocity. Also, the accuracy of analytical estimates, in terms of wavelength, density contrast, and transverse inhomogeneity, has been quantified. For the typical large filament to coronal density contrast the mechanism produces rapid damping in timescales of the order of a few oscillatory periods only. The obtained damping rates are only slightly dependent on the wavelength of perturbations. An important result is that the damping rate becomes independent of density contrast for large values of this parameter. This has two seismological consequences. First, the observational determination of density contrast is less critical than in the low contrast regime. Second, according to seismic inversion results that combine theoretical and observed periods and damping times (\\citealt{Arregui07,GABW08}), high density thread models would be compatible with relatively short transverse inhomogeneity length-scales. Analytical estimates of $l/a\\sim 0.15$  can even be calculated using equation~(\\ref{dampingrate}) for a given observed $\\tau_d/P=4$, taking the limit $c\\rightarrow\\infty$.  We believe our conclusion on the relevance of resonant damping in filament thread oscillations is robust in front of the main simplification adopted for the present study. Observations show that threads are only partially filled with cool and dense plasma and our one-dimensional model misses this property. Although the resonant damping mechanism relies on the Alfv\\'en velocity inhomogeneity in the direction transverse to the magnetic field, an investigation of the damping properties in such a two-dimensional configuration is a necessary issue. Finally, the presence of flows is commonly observed along filament threads and the interplay between resonant damping and flows must also be explored in this context."
    },
    "0806/0806.4803_arXiv.txt": {
        "abstract": "A simple, yet powerful, algorithm for computed tomography of the solar corona is presented and demonstrated using synthetic EUV data. A minimum of three perspectives are required. These may be obtained from \\textit{STEREO}/EUVI plus an instrument near Earth, e.g.\\ \\textit{TRACE} or \\textit{SOHO}/EIT. ",
        "introduction": "Observations by the Extreme Ultraviolet Imager \\citep[EUVI,][]{Wuelser2004} aboard \\textit{STEREO} provide the first simultaneous, stereoscopic image pairs of the solar corona and transition region. Ideally, these data are simple projections through an optically thin corona. However, the 3D distribution of emission is difficult to estimate with only two projections \\citep[see][and references therein]{Gary1998}. This difficulty may be explained as follows. Consider an object  $I(x,y,z)$ on domain $D$. The coordinates $x$ and $y$ need not be orthogonal, but $z$ is orthogonal to $x$ and $y$. Given two projections $f(x,z) = \\int_D I\\,dy$ and $g(y,z)=\\int_D I\\,dx$, a plausible reconstruction of $I$ is \\begin{equation}\\label{eq:separable} I'(x,y,z) = \\frac{f(x,z)\\,g(y,z)}{T(z)}, \\end{equation} where $T(z) = \\int_D I(x,y)\\,dx\\,dy = \\int_D f\\,dx = \\int_D g\\,dy$ is the total emission of a plane of constant $z$. Unfortunately, this solution fails utterly in practice. For each pair of sources in $I$,  $I'$ introduces a pair of ``ghost'' artifacts. These are systematic errors, independent of the noise and apparently unavoidable. Traditional regularization strategies are not fruitful: $I'$ is positive, is as smooth as the observed images $f$ and $g$, and is precisely the maximum entropy solution. More information is therefore required to guide the tomographic reconstruction.  Previously described approaches to the \\textit{STEREO} coronal tomography problem rely on assumptions about the geometry of the coronal plasma distribution. The triangulation method \\citep{Gary1998,Aschwanden2005,Aschwanden2008} assumes loops with circular cross-section, and relies on the identification of the same loops in both images. The magnetic tomography approach \\citep{WiegelmannInhester2006} also assumes loops with circular cross-section, and incorporates magnetic field extrapolations to constrain loop geometries.  These methods are powerful, but they require assumptions about things that one might reasonably hope to learn from the 3D reconstruction.  I propose that EUV images taken from a third perspective---e.g.\\ \\textit{TRACE} or \\textit{SOHO}/EIT---may provide adequate additional constraints for coronal tomography, without any assumptions about loop geometry or magnetic fields. I describe a simple computed tomography algorithm, fast enough to run in real time, and demonstrate its performance using synthetic data with three viewpoints. ",
        "conclusions": "\\label{sec:discussion}  Coronal tomography is possible with as few as three viewpoints, making no prior assumptions about coronal morphology or magnetic fields.  The test cases demonstrate recovery of complex geometry without reference to magnetic field extrapolations or assumptions about loop geometry. Loops that run in an east-west direction, however, provide insufficient depth cues for three instruments confined to the  ecliptic.  The SMART algorithm described here provides a noise-insensitive tomographic reconstruction by finding an optimal balance between goodness of fit and local smoothness.  \\textit{STEREO} will obtain data at large separation angles in Fall 2008. Observations from \\textit{SOHO} and/or \\textit{TRACE} at that time should allow the best opportunity for application of SMART to coronal tomography."
    },
    "0806/0806.4342_arXiv.txt": {
        "abstract": "Thermohaline mixing has recently been proposed to occur in low mass red giants, with large consequences for the chemical yields of low mass stars. We investigate the role of thermohaline mixing during the evolution of stars between 1$\\mso$ and 3$\\mso$, in comparison to other mixing processes acting in these stars. We confirm that thermohaline mixing has the potential to destroy most of the $^3$He which is produced earlier on the main sequence during the red giant stage. In our models we find that this process is working only in stars with initial mass  $M \\simle1.5\\mso$. Moreover, we report that thermohaline mixing can be present during core helium burning and beyond in stars which still have a $^3$He reservoir. While rotational and magnetic mixing is negligible compared to the thermohaline mixing in the relevant layers, the interaction of thermohaline motions with differential rotation and magnetic fields may be essential to establish the time scale of thermohaline mixing in red giants. ",
        "introduction": "Thermohaline mixing is not usually considered as an important mixing process in single stars, since the ashes of thermonuclear fusion consists of heavier nuclei than its fuel, and stars usually burn from the inside out. The condition for thermohaline mixing, however, is that the mean molecular weight ($\\mu$) decreases inward. Recently \\citet[CZ07]{cz07} identified thermohaline mixing as an important mixing process which significantly modifies the surface composition of red giants after the first dredge-up. The work by CZ07 was triggered by the paper of \\citet[EDL06]{edl06}, who found a $\\mu$-inversion in their $1\\mso$ stellar evolution model, occurring after the so-called luminosity bump on the red giant branch (RGB), which is produced after the first dredge-up, when the hydrogen-burning shell reaches the chemically homogeneous part of the envelope. The $\\mu$-inversion is produced by the reaction $^3$He($^3$He,2p)$^4$He (as predicted by \\citet{ulr72}). It does not show up earlier, since the magnitude of the $\\mu$-inversion is small and negligible compared to a stabilizing $\\mu$-stratification.  The mixing process below the convective envelope in models of low-mass stars turns out to be essential for the prediction of the chemical yield of $^3$He (EDL06), and to understand the surface abundances of red giants, in particular the $^{12}$C/$^{13}$C ratio, and the $^7$Li, carbon and nitrogen abundances (CZ07). This may also be important for other occurrences of thermohaline mixing in stars, i.e., in single stars when a $\\mu$-inversion is produced by off-center ignition in semi-degenerate cores, or in stars which accrete chemically enriched matter from a companion in a close binary \\citep[e.g.,][]{sgi+07}. Accreted  metal-rich matter during the phases of planetary formation also leads to thermohaline mixing. The host stars of exoplanets present a metallicity excess compared to stars in which no planets have been detected. This metallicity excess can be reconciled with the overabundances expected in cases of accretion if thermohaline mixing is included in the picture \\citep{vau04}. ",
        "conclusions": "We comfirm the results of EDL06 and CL07: thermohaline mixing in low-mass giants is capable of destroying large quantities of ${^3}$He, as well as decreasing the ratio $\\c1213$. Thermohaline mixing starts when the hydrogen burning shell moves into the chemically homogeneous layers established by the first dredge-up. Our models show further that thermohaline mixing remains important during core helium burning, and can still be relevant during the AGB phase --- including the termally-pulsing AGB stage. This can result in important changes to the surface abundances of low-mass stars. The quantitative discussion is complicated by the fact that the capability of thermohaline mixing to change surface abundances depends on an uncertain efficiency parameter, as well as on the local ${^3}$He abundance. The efficiency parameter $\\ath$ is still a matter of debate and is probably strongly affected by the interaction of thermohaline mixing with rotation and magnetic fields. The local ${^3}$He abundance depends on the previous history of mixing.  Our calculations show that in the relevant layers, thermohaline mixing generally has a higher diffusion coefficient than rotational instabilities and magnetic diffusion. Still, the interaction of thermohaline mixing with magneto-rotational instabilities is important: we expect the speed of the mixing to be strongly affected by the presence of differential rotation and magnetic fields. To better understand the picture it would be desirable to have realistic MHD simulations of thermohaline mixing."
    },
    "0806/0806.4174_arXiv.txt": {
        "abstract": "A new model for the formation of Jovian planets is proposed. We consider planets forming at large distances from a protostar ($\\gtrsim 100$ AU) through direct fragmentation of a gas cloud, by the same formation mechanism as wide stellar and brown dwarf binaries. We model the gravitational evolution of a system of these distant planets and a second population formed in a disk closer to the star. We compute the typical closest approach of these planets to the star (i.e., smallest pericenter) over the course of their evolution. When the planets reach a pericenter within a gaseous disk surroundig the star, dynamical friction from this disk slows down the planet at each plunge, causing its orbit to be gradually circularized and made coplanar with the disk. After the disk dissipates, a large fraction of these planets may be left at orbits small enough to be detected in present radial velocity surveys. A brief analytic derivation of the rate of orbital energy dissipation during these disk crossings is presented. Observational tests of this model are discussed. ",
        "introduction": "Our understanding of planetary systems and brown dwarfs has been revolutionized over the last decade. About 250 exoplanets have been detected so far by the method of radial velocities, with orbital periods up to $\\sim$ 10 years and radial velocity amplitudes down to $K \\gtrsim 3 \\mets$, although the velocity threshold for detection is still higher (K $\\gtrsim 10 \\mets$ ) for the majority of observed stars \\citep{but06}. Several brown dwarfs have been directly imaged at large distances ($\\gtrsim 30 \\au$) from nearby stars \\citep{nak95,giz01,neu04}, and detected as well by radial velocities especially around M dwarfs \\citep{clo03}. A population of planetary mass objects has been identified in young star clusters \\citep{zap00, cab07}. The majority of the exoplanets known so far are of Jovian type because of the sensitivity limit of the observations, but a few terrestrial planets have been identified by radial velocities down to masses of $\\sim 5 M_E$. The microlensing technique has resulted in the detection of several additional planets of a few Earth masses \\citep{bea06,gou06,gau08}, and promises to rapidly increase the detection rate in the future. Planets very close to their stars are additionally being found by transits, where a rapid increase of the number of detections is also likely from the space missions CoRoT \\citep{bor03} and Kepler \\citep{bas05}.  The present set of known planets and brown dwarfs already constitutes a rich data set, which informs us on the final orbital distribution of planets in semimajor axis and eccentricity as a function of the planet mass and the properties of the host star (basically mass and metallicity). The principal challenge being faced in the field of exoplanets today is to understand how the formation mechanisms for planets and other stellar companions have determined their initial masses and orbits, and to discern how the subsequent dynamical and physical evolution of planetary systems has led to the final distribution of orbits that is observed.  Generally, two possible formation mechanisms have been discussed for Jovian exoplanets. The first is the core-accretion model \\citep{ste82, pol96}, where a terrestrial planet is formed first by aggregation of planetessimals within a gas disk, and then the solid core starts accreting hydrogen and helium gas after reaching the minimum mass required for accretion to take place. The second is gravitational instability in a gaseous disk, where a Jovian planet forms directly by gravitational collapse of a gas clump in a disk \\citep{bos97, bos07}. A third mechanism is the direct turbulent fragmentation of the protostellar cloud and gravitational collapse of the fragments before the gas has settled into a disk, a process that is believed to be fundamental for the formation of stars and brown dwarfs and the determination of their Initial Mass Function \\citep{pad99,pad04,bon08}. This third mechanism has not been discussed as much in the literature as a possible origin for planets, with the exception of \\citet{pap01} and \\citet{ter02}, probably because it can only form planets at very large distances from a star ($\\gtrsim 100 \\au$), so it has not been thought that the planets detected by radial velocities could be formed in this way. However, \\citet{raf05} has argued that the second mechanism of gas disk instability can form only massive planets at similarly large distances. In this paper, we propose a mechanism by which planets formed through the third mechanism might migrate inwards from the distances of the widest observed binary systems (probably formed by the same fragmentation mechanism) to the distances at which planets have been found in radial velocity surveys.  A complete theory for the formation of stellar companions should be addressing at the same time the formation of planets, brown dwarfs and binary stars. It must be born in mind that the mass limits chosen to divide planets, brown dwarfs and stars into separate classes of objects (which are the minimum masses required to ignite fusion of deuterium and hydrogen, $13$ and $80$ Jupiter masses, respectively) have no relation with the physics of the formation process. We should therefore expect that, if there are several formation mechanisms, each mechanism may give rise to objects over mass ranges that partly overlap, and that extend over more than one of these three classes of objects. In fact, we already know that the core-accretion mechanism must be operating at some level, because some planets (at least Saturn, Uranus and Neptune, and the transiting planet HD149026b) have a mean density that requires the presence of a large core enriched with heavy elements. At the same time, the presence of isolated planets in young star clusters, as well as the existence of brown dwarfs at very large distances from their primary stars (e.g., the pair of brown dwarfs at $\\sim 1500 \\au$ from the nearby star $\\epsilon$ Indi, \\citet{mcc04} favors the operation of the third mechanism as well (although some of these distant planets could be the result of dynamical ejection from a system of planets formed by core accretion). In addition, the existence of binary stars at semimajor axis $a \\lesssim 30 \\au$ with relatively large mass ratios suggests that some process of gravitational instability in an opaque, disk-like structure at distances comparable to the known exoplanets can result in the formation of a massive companion by direct gravitational collapse. But could the latter two mechanisms of direct gas collapse also be responsible for some of the short-period exoplanets detected by radial velocities?  To address this question, it will be useful to list a series of relevant facts we have already learned from observations: \\begin{itemize} \\item[\\em 1)] Many Jovian planets must have migrated radially, probably through their interaction with a gaseous disk, to explain their presence at orbital semimajor axes smaller than a few AU \\citep{lin96}. These planets cannot form very close to their stars because the solid cores do not grow massive enough to start the process of runaway gas accretion at the high temperatures of the inner disk \\citep{ste82,raf06}. In principle, this migration process could have occurred starting from distances larger than a few AU as long as there is enough gas mass in the disk at large distances. \\item[\\em 2)] The abundance of Jovian planets increases with the metallicity of the star \\citep{fis05}. This must be due to a more efficient formation of planets around metal-rich stars, and not to an enrichment of the envelopes of stars hosting planets \\citep{pin01}. This fact favors the core-accretion model as the mechanism for forming the majority of Jovian planets, because high metallicities would allow solid cores to grow to a larger mass and facilitate the formation of Jovian planets. \\item[\\em 3)] Gas disks around young stars survive for only a few million years \\citep{hai01,sic06,mor07,ale07}. Therefore, the accretion of gas by Jovian planets, as well as any migration process requiring the presence of gas, must take place during this period. Although the possibility of radial migration by planetesimals has been proposed \\citep{mur98}, it seems unlikely that a disk of planetesimals could be massive enough to allow for migration of Jovian planets over large radial intervals. \\item[\\em 4)] Orbits of planets have a mean eccentricity $\\sim 0.3$. Their orbits have therefore been perturbed from the original circular orbits if they formed by the core accretion process, either by the interaction with the gas disk, or as a result of the dynamical interaction among several planets after the disk evaporated. The latter process may generically produce an eccentricity distribution similar to the observed one \\citep{jur08,cha08}. \\item[\\em 5)] There is a brown dwarf desert (i.e., a large reduction in the abundance of brown dwarfs relative to planets and stellar companions) orbiting solar type stars with orbital periods less than $\\sim 10$ years. Brown dwarfs are more common at large distances, and also at short distances around M dwarf stars. \\end{itemize}  The brown dwarf desert, together with the difficulties for forming Jovian planets at small distances by gravitational instability in a gas disk (see Rafikov 2005, 2007), suggest that two mechanisms operate to form objects of different mass at short distances: core accretion, which forms planets, and gravitational instability of a massive disk or the protostellar cloud, which forms binary stars. At the same time, planets and brown dwarfs must be forming by direct gas collapse at large distances, where they seem to be abundant as well. If they were able to migrate inwards from there, a natural outcome would be that the known planet population detected by radial velocities is the result of two different formation processes, as already suggested previously (Black 1997, Mayor \\etal 1998). Radial inwards migration of planets formed by core accretion must occur within a few AU, in order to explain the close-in planets, so it would not be highly surprising that similar migration ocurred at larger distance as well, owing to an extended gas disk.  In Ribas \\& Miralda-Escud\\'e (2007), tentative evidence was found that the host stars of the most massive planets ($M \\gtrsim 4 M_J$) are less metal rich than the host stars of lower mass planets. We suggested that many of the massive planets could have formed by a process of gas collapse that is independent of metallicity, with the opacity limit for fragmentation providing a natural mass lower limit, while lower-mass planets would have formed by core accretion. We further explore this idea in this paper, focusing on the possibility of inward migration for a population of massive planets formed by gas collapse at large distances. Following the earlier suggestion of Papaloizou \\& Terquem (2001) and Terquem \\& Papaloizou (2002), we consider in \\S \\ref{sec_simu} and \\S \\ref{sec_resu} the dynamical evolution of a system of brown dwarfs and high-mass planets formed at large distances, together with a population of lower mass planets formed at small distances in initially circular orbits. We examine the possibility that some of the planets formed at large distances may be relocated into an orbit with a greatly reduced semimajor axis, making them detectable by current radial velocity surveys, simply by interacting gravitationally among themselves and with the population of inner planets formed by core accretion. We shall show that this dynamical evolution can easily increase the eccentricity of some objects to values close to unity, and therefore greatly reduce their pericenter, but it is much more unlikely to reduce the semimajor axis and leave the object in a stable orbit. We propose in \\S \\ref{sec_disk} that adding the interaction of a planet with the gas disk can act as a dissipative mechanism and help reduce the semimajor axis, achieving the desired result of migration of a planet formed far from the disk into an orbit within the gas disk at small radius. These results are discussed in \\S \\ref{sec_disc}. ",
        "conclusions": "\\label{sec_disc}  We propose in this paper a new model for the formation of the most massive Jovian planets based on direct fragmentation of a pre-stellar cloud. This fragmentation process would form planets at distances from the star much larger than the orbital sizes that are detectable in radial velocity surveys, implying that a migration mechanism is required to transport these planets to orbits with semimajor axis much smaller than their initial values at formation. Using orbital evolution simulations we have explored the orbital exchange mechanism, in which a massive planet formed by fragmentation would gradually be transported to a small orbit by gravitational interactions with other smaller planets in the system. The results show that only about 1\\% of the planets born beyond 30~AU reach final states with semi-major axes smaller than about 3 AU.  A better mechanism to transport planets to small orbits is obtained when considering the dynamical friction effect with a disk around the star. In this case, a planet can be transported inwards by decreasing only the pericenter through perturbations with other planets, therefore increasing by a large factor the amount of phase space from which the required migration can be achieved. Our results show that a significant fraction of the outer planets spend several orbits at pericenter distances below 3 AU over their dynamical evolution. The approximate analytical estimate in the previous section confirms that the semimajor axis can be reduced with a relatively small number of disk crossings, placing a planet formed by fragmentation into an orbit of the typical size detected by radial velocities.  This dynamical mechanism for migration does not provide an explanation for the absence of brown dwarfs in orbits with $a \\lesssim 5$ AU around solar-type stars (the so-called brown dwarf desert). As explained in \\S \\ref{sec_resu}, the probability for a planet to reach a small pericenter leading to disk crossings is nearly independent of its mass. In addition, in \\S \\ref{sec_disk} we showed that the number of disk crossings needed to slow down a planet is inversely proportional to the mass (see eq. \\ref{ncross}), suggesting that brown dwarfs should actually be easier to capture than Jovian planets. However, a possible difficulty for capturing objects that are very massive may be that the disk quickly destroyed by the plunges themselves as a result of the induced heating. Brown dwarfs might simply be too massive to undergo the process of orbital grinding before they destroy the gas disk. If it is captured at large radius in a gas disk, a brown dwarf might also not be able to migrate inwards because its angular momentum would easily be larger than that of the disk.  The scenario of direct fragmentation we have proposed may be tested observationally in several ways, and we discuss some of them here: ${\\emph i)}$ As explained in \\S \\ref{sec_resu}, the dynamical evolution of planetary systems that perturbs planets into the high eccentricities required for interaction with the disk should also produce a large number of ejected planets. If the star hosting the planetary system is part of stellar cluster, the ejection velocity of the planet would often be small enough for the ejected planet to remain gravitationally bound to the cluster, implying that the abundance of isolated planetary-mass objects should be comparable to that of stars. Some isolated planets have in fact been identified in young star-forming regions \\citep{zap00,cab07}, and our hypothesis may be tested as the observational determination of the abundances of these objects improves. Note, however, that many planetary-mass objects may be also be born by fragmentation without being initially bound to any star, so the abundance of isolated planets bound in clusters can only yield an upper limit to the number of planets that are initially bound to an individual star. Some of the ejected planets will have a high enough velocity to escape from the cluster, and these may in principle provide a more specific test for our model: they should be discovered around young open clusters and associations moving away from the clusters at the typical ejection velocities ($\\sim 1 \\kms$), with a distribution of distances that would indicate the rate at which they have been ejected from the young planetary systems.  ${\\emph ii)}$ If massive planets are indeed formed by gas fragmentation, they should probably not contain a rocky core (although see \\citet{bos02} for the possibility that a core could form by sedimentation of dust grains to the center), in which case they would be less dense than objects formed by core accretion. It might be possible to test this for planets of known age with measured radii from transits. These will probably have to be, however, transiting planets that are not very close to their host stars, to avoid possible effects from gravitational tides and the stellar radiation that seem to be affecting the radius evolution \\citep{bod01,gui06}.  ${\\emph iii)}$ Formation of planets in a protoplanetary disk will naturally predict that the spin axis of the star and the orbital axis of the planet are in close alignment. However, formation by fragmentation in a random pre-stellar cloud would result in a random orientation of the planet orbit relative to the star spin. When the planet is captured by the disk, the disk plane will be changed for a massive enough planet, when the planet mass is comparable to that of the disk, in which case a large misalignment of the star spin and the planet orbit would be expected. This may be tested for transiting planets, for which the stellar spin orientation can be measured from the Rossiter-McLaughlin effect \\citep{que00,win05}. Interestingly, the most massive transiting planets, such as XO-3 b (11.8~M$_{\\rm J}$) or HD 17156 b (3.1~M$_{\\rm J}$), do indeed tend show a moderate degree of misalignment \\citep{nar08,heb08}."
    },
    "0806/0806.2752_arXiv.txt": {
        "abstract": "In this paper we derive the probability of the radial profiles of spherically symmetric inhomogeneities in order to provide an improved estimation of the number density of primordial black holes (PBHs). We demonstrate that the probability of PBH formation depends sensitively on the radial profile of the initial configuration. We do this by characterising this profile with two parameters chosen heuristically: the amplitude of the inhomogeneity and the second radial derivative, both evaluated at the centre of the configuration. We calculate the joint probability of initial cosmological inhomogeneities as a function of these two parameters and then find a correspondence between these parameters and those used in numerical computations of PBH formation. Finally, we extend our heuristic study to evaluate the probability of PBH formation taking into account for the first time the radial profile of curvature inhomogeneities. ",
        "introduction": "The idea that large amplitude matter overdensities in the universe could have collapsed through self-gravity to form  primordial black holes (PBHs) was first put forward by Zeldovich and Novikov \\cite{zeldovich66}, and then independently by Hawking \\cite{hawking71}, more than three decades ago. This theory suggests that large amplitude inhomogeneities in the very early universe overcome internal pressure forces and collapse to form black holes. A lower threshold for the amplitude of such inhomogeneities $\\delta_{\\textrm{th}} \\equiv (\\delta \\rho / \\rho)_{\\textrm{th}}$, was first provided by Carr \\cite{carr75}, giving $ \\delta_{\\textrm{th}} \\approx 1/3 $  at the time of radiation domination. This value was found by  comparing the Jeans length of the overdensity with the scale of the cosmological horizon at the time of formation.  The mass fraction of the universe turning into PBHs of mass $M$ at their formation time, $\\beta_{\\textrm{PBH}}(M)$, is computed using the probability density function (PDF) of the relevant field of inhomogeneities, which is provided by the cosmological theory. The mass fraction $\\beta_{\\textrm{PBH}}(M)$ is customarily given by the integral of this PDF over the amplitude $\\delta \\equiv \\delta\\rho / \\rho$, with a lower limit equal to $\\delta_{\\textrm{th}}$ \\cite{press74,carr75}.  The probability of PBH formation is a useful tool to constrain the mean amplitude of inhomogeneities on scales which cannot be probed by any other methods. The PBH contribution to the energy density increases with time during the radiation-dominated epoch. For this reason, the  PBHs formed considerably before the end of radiation domination are the most relevant to cosmology \\cite{lidsey94,liddle98,sendouda06,zaballa06,bugaev06,kohri07}. We will focus our study on these kind of PBHs and assume that the background matter at the time of PBH formation is radiation-dominated. To make this cosmological tool more precise, we must improve the calculations of the probability of PBH formation. This demands a more accurate evaluation of the threshold value $\\delta_{\\textrm{th}}$, or the equivalent curvature inhomogeneity $\\R_{\\textrm{th}}$ \\cite{bicknell79,niemeyer97,jedamzik99,shibata99,hawke02,musco04}. In search of these values, it was evident that the process of PBH formation depends on the pressure gradients in the collapsing configuration in addition to the amplitude \\cite{nadezhin78,shibata99,musco07}. It was also found that such pressure gradients can modify the value of $\\delta_{\\textrm{th}}$ significantly. Hence, when calculating the probability of PBH formation, one should consider the shape and radial profile of the initial configurations. These profiles are directly related with the internal pressure gradients. This is the main motivation for the present work. Previous studies, concerning the probability of PBH formation, take the amplitude of perturbations to be the only parameter determining the probability density. In addition to that, we include for the first time a parameter related to the slope of curvature profile at the edge of the configuration \\footnotemark\\footnotetext{In the context of dark matter haloes the question of initial profiles is effectively irrelevant because galaxies are formed from pressureless configurations. The density profiles and shapes of virialized haloes result from the evolution of the initial high peaks and are not linked to the profile of initial configurations that we investigate here (see e.g. \\cite{cooray02} for a review on the profiles of dark matter haloes and \\cite{maccio08} for some recent results on this topic).}.  In this paper, we calculate the probability of finding a curvature configuration with a given radial profile. As follows from \\cite{polnarev81,polnarev82, khlopov83,polnarev85,zabotin87}, PBH formation takes place only from nearly spherical configurations, so in this paper we restrict ourselves to the spherically symmetric case. In this first approximation we describe the radial profiles by introducing two parameters: the central amplitude of the curvature inhomogeneity $\\R(\\br = 0)$ and the second radial derivative at the centre $\\R''(\\br = 0)$, which is chosen mainly to avoid technical difficulties. The introduction of these parameters is a first step towards the full parametrisation of profiles in terms of even derivatives at the centre of configurations, i.e. in terms of $\\R^{(2n)}(0)$ (the odd derivatives $\\R^{(2n + 1)}(0)$   are all zero due to the assumed spherical symmetry). In the future, with the results from more accurate codes simulating the formation of PBHs, we will have at hand a larger number of conditions for the collapse of a curvature profile. An equal number of parameters will be required for the complete description of these profiles and the probability of finding them. In the meantime, only families of curvature profiles described by two parameters are available. We consequently limit ourselves to the two-parametric description of initial curvature profiles.    The central amplitude $\\R(0)$ has been used in previous calculations of gravitational collapse \\cite{shibata99}, and the probability of PBH formation \\cite{green04,zaballa06}. In the present paper we compute the probability to find a given configuration as a function of the two parameters $[\\R(0),\\R''(0)]$. We subsequently illustrate how this two-parametric probability is used to correct the probability of PBH formation. Such an exercise is presented for illustration purposes. The results, based on a non-rigorous but physically meaningful determination of the parameters which describe the initial profiles, show how the corrections to $\\beta_{\\mathrm{PBH}}$ are significant and they will be considered in more detail in future studies of PBH formation.  This paper is organised as follows. In Section \\ref{prob:construction} we calculate the joint probability distribution for $\\R(0)$ and $\\R''(0)$. In Section \\ref{metrics}, we relate these parameters to those used in the most recent numerical computations of PBH formation. In Section \\ref{prob:comparison} we present the total probability of PBH formation, integrating the probability distribution derived in Section \\ref{prob:construction} over the relevant region of parameter space $\\left[\\R(0)\\,,\\,\\R''(0)\\right]$. In Section \\ref{conclusion} we summarise our results and discuss future research in this area. ",
        "conclusions": "\\label{conclusion}  We have developed a method for calculating the two-parametric probability of PBH formation, taking into account  the radial profiles of non-linear curvature cosmological inhomogeneities. This is the fist step towards calculating the $N$-parametric probability, which takes into account the radial profiles more precisely than studies using the amplitude as the only relevant parameter.  Using the results of sophisticated numerical computations, we obtain the values of $\\R''(0)$ that are relevant for PBH formation. Subsequently we have incorporated these values to the total probability of PBH formation. Finally, we have provided an example of the consequences of this probability for the statistics of PBHs.   The results obtained show that, if we restrict ourselves to the PBH formation calculated for parabolic profiles only (as described in Section \\ref{metrics}), then the total probability of PBH formation is orders of magnitude below  previous estimations. On the other hand, with the aid of heuristic arguments we show that a much larger region of parameter-space $[\\R(0),\\R''(0)]$ representing non-parabolic profiles should also be considered in the estimation of the probability of PBH formation (see Fig.~\\ref{fig11}). In this case, the total probability of PBH formation is higher than the single-parameter estimate of previous works. In this case, we have an opportunity to impose new bounds on the power spectrum on the scales relevant for PBH formation. Analysing the uncertainty of our results, mostly due to the heuristic nature of the present study, we have demonstrated how much we still have to understand about the formation and statistics of PBHs. The physical arguments supporting our results should be made rigorous by direct verification with numerical hydrodynamical simulations of PBH formation. This in turn would provide  valuable support for the initial motivation of this work.  The main conclusion of this paper is that the amplitude of initial inhomogeneities is not the only parameter which determines the probability of PBH formation. The ultimate solution of the problem requires a greater set of parameters and a larger range of their values to determine all high curvature configurations that form PBHs, which is a huge task for future research. In the meantime, we have a method to operate with the statistics of all these parameters."
    },
    "0806/0806.2422_arXiv.txt": {
        "abstract": "We study an inflationary scenario with a vector impurity. We show that the universe undergoes anisotropic inflationary expansion due to a preferred direction determined by the vector. Using the slow-roll approximation, we find a formula to determine anisotropy of the inflationary universe. We discuss possible observable predictions of this scenario. In particular, it is stressed that primordial gravitational waves can be induced from curvature perturbations. Hence, even in low scale inflation, a sizable amount of primordial gravitational waves may be produced during inflation. ",
        "introduction": "It is often mentioned that cosmology has entered into a new stage, so-called precision cosmology. Of course, it is referring to developments of observational side.  From theoretical point of view, however, we have not yet exhausted possible phenomenology on the order of a few percent. Clearly, it is important to explore qualitatively new scenarios at the percent level. Here, above all, we would like to point out that an inflationary model with a few percent of anisotropy yields significant consequences.  Indeed, a few percent does not mean the consequent effects are negligible. Rather, it provides the leading component of the primordial gravitational waves in low scale inflationary models which are preferred by recent model construction in string theory~ \\cite{Kallosh:2004yh}.  One may feel that to seek the anisotropic inflation is against for the basics of the inflationary scenario.  Actually, the isotropy is the most robust property of inflation because of the cosmic no-hair theorem on the isotropization of Bianchi universes~\\cite{Wald:1983ky}. However, it is possible to evade the cosmic no-hair theorem  by incorporating the Kalb-Ramond action~\\cite{Kaloper:1991rw} or considering higher curvature theories of gravity~\\cite{Kawai:1998bn,Barrow:2005qv}. In spite of the possibility, no one has attempted to construct any inflation models based on these ideas. The apparent other possibility is to break the Lorentz invariance by introducing a condensation of vector field. The model proposed in~\\cite{Ackerman:2007nb,Dulaney:2008bp} seems to be successful, however, it is known to be metastable~\\cite{Clayton:2001vy,Dulaney:2008ph,Bluhm:2008yt}. A more natural possibility would be to realize a slow-roll phase of vector fields like as inflaton fields in chaotic inflationary scenarios. So far, it has been believed that it is difficult to make the vector field slow-roll without fine tuning~\\cite{Ford:1989me} \\footnote{Recently, a non-standard spinor driven inflation has been proposed~\\cite{Boehmer:2007ut} . It is interesting to see if the natural slow-rolling can be achieved in this case.}. Very recently, the situation has changed by the discovery of slow-roll mechanism for the vector field due to non-minimally coupling~\\cite{Golovnev:2008cf}. Hence, an apparent difficulty to construct the anisotropic inflationary scenario has been resolved. At this point, it is important to realize that both scalar and vector fields are ingredients of fundamental particle physics models. Therefore, it is natural to consider both scalar and vector fields exist during inflation. Of course, the vector fields should be subdominant in the inflationary dynamics in order to reconcile the scenario with current observational data ~\\cite{Pullen:2007tu}. In this sense, the vector field should be regarded as impurity. Nevertheless, the effect of the vector impurity on observables should not be overlooked under the current precision cosmology.  In this paper, we propose an anisotropic inflation model with the vector impurity. To the best of our knowledge, this is the first concrete model which realizes anisotropic inflation, exits successfully to the isotropic standard universe, and provides a framework to discuss interesting phenomenology. We argue that the anisotropic inflation yields the statistical anisotropy in fluctuations. More importantly, as is expected, the primordial gravitational waves could be induced from curvature perturbations through the anisotropic background. Hence, we can expect the correlation between the curvature perturbations and the gravitational waves. In addition to these, we point out that linear polarization of the gravitational waves is created, which should be observed through CMB or direct interferometer observations.  The organization of this paper is as follows. In section II, we present the model for the anisotropic inflation. In section III, we analyze the system numerically and show that the anisotropic inflation is realized successfully. In section IV, using the slow-roll approximation, we obtain degrees of the anisotropy. In section V, we discuss phenomenology of the anisotropic inflation induced by the vector impurity. There, we emphasize that the gravitational waves can be produced through the anisotropy of the spacetime. The final section is devoted to the discussion. ",
        "conclusions": "We have explored the cosmological implications of the vector impurity. It turned out that the vector impurity affects the cosmic no-hair theorem. Consequently, the anisotropic accelerating universe is possible in the presence of the vector impurity. To prove it, we have numerically solved the equations and shown the phase flow from which we see the anisotropic inflation successfully occurs and ends with oscillation. We found the explicit formula to determine the anisotropy of the inflationary universe by using the slow-roll approximation.  We have also discussed possible consequences of the anisotropic inflation on the cosmological fluctuations. Since rotational invariance is violated, the statistical isotropy of CMB temperature fluctuations can not be expected. It is intriguing to seek for the relation to the large scale anomaly discovered in CMB by WMAP~\\cite{Eriksen:2003db,Land:2005ad,Copi:2005ff}. More interestingly, the tensor perturbations could be induced from the curvature perturbations through the anisotropy of the background spacetime. One immediate consequence is the correlation between the curvature perturbations and the tensor perturbations. This correlation should be detected through the analysis of temperature-B-mode correlation in CMB. This new possibility implies, even in the low scale inflation, we can expect the primordial gravitational waves. This is an important result for future observational planning, because there has been worry that string cosmology tend to suggest the low scale inflation. Moreover, because of the anisotropy, there might be linear polarization in the primordial gravitational waves. This polarization can be detected either through the CMB observations or direct interferometer observations. These predictions can be checked by future observations.  To make these predictions more precise, we need to develop the  perturbative analysis~ \\cite{Pereira:2007yy, Pitrou:2008gk, Gumrukcuoglu:2007bx,Tomita:1985me}. These are now under investigation~\\cite{kimura}. The calculation of the perturbations is much more complicated due to the violation of rotational invariance. However, since the anisotropic inflationary universe is smoothly connected to the isotropic radiation dominant phase, the interpretation of the results should be clear. The implications in primordial magnetic fields and the structure formation of the universe~\\cite{Ando:2008zz} should be also studied in future work."
    },
    "0806/0806.2614_arXiv.txt": {
        "abstract": "We calculate the one-photon loop radiative corrections to charged pion Compton scattering, $\\pi^- \\gamma \\to \\pi^- \\gamma $. Ultraviolet and infrared divergencies are both treated in dimensional regularization. Analytical expressions for the ${\\cal O}(\\alpha)$ corrections to the invariant Compton scattering amplitudes, $A(s,u)$ and $B(s,u)$, are presented for 11 classes of contributing one-loop diagrams. Infrared finiteness of the virtual radiative corrections is achieved (in the standard way) by including soft photon radiation below an energy cut-off $\\lambda$, and its relation to the experimental detection threshold is discussed. We find that the radiative corrections are maximal in backward directions, reaching e.g. $-2.4\\%$ for a center-of-mass energy of $\\sqrt{s}=4m_\\pi$ and $\\lambda=5\\,$MeV. Furthermore, we extend our calculation of the radiative corrections by including the leading pion structure effect (at low energies) in form of its electric and magnetic polarizability difference, $\\alpha_\\pi - \\beta_\\pi \\simeq  6\\cdot 10^{-4}\\,$fm$^3$. We find that this structure effect does not change the relative size and angular dependence of the radiative corrections to pion Compton scattering. Our results are particularly relevant for analyzing the COMPASS experiment at CERN which aims at measuring the pion electric and magnetic polarizabilities with high statistics using the Primakoff effect. ",
        "introduction": "Pion Compton scattering, $\\pi^- \\gamma \\to\\pi^- \\gamma$, allows one to extract the electric and magnetic polarizabilities of the (charged) pion. In a classical picture these polarizabilities characterize the deformation response (i.e. induced dipole moments) of a composite system in external electric and magnetic fields. In the proper quantum field theoretical formulation the electric and magnetic polarizabilities, $\\alpha_\\pi$ and $\\beta_\\pi$, are defined as expansion coefficients of the Compton scattering amplitudes at threshold.  However, since pion targets are not directly available, real pion Compton scattering has been approached using different artifices, such as high-energy pion-nucleus bremsstrahlung $ \\pi^- Z\\to\\pi^- Z\\gamma$, radiative pion photoproduction off the proton $\\gamma p \\to \\gamma \\pi^+ n$, and the crossed channel two-photon reaction $\\gamma \\gamma \\to \\pi^+\\pi^-$.  From the theoretical side there is an extraordinary interest in a precise (experimental) determination of the pion polarizabilities. Within the framework of current algebra \\cite{terentev} it has been shown (long ago) that the polarizability difference $\\alpha_\\pi-\\beta_\\pi$ of the charged pion is directly related to the axial-vector-to-vector form factor ratio $h_A/h_V \\simeq 0.44$ measured in the radiative pion decay $\\pi^+\\to e^+ \\nu_e \\gamma$ \\cite{frlez}. At leading (nontrivial) order the result of chiral perturbation theory \\cite{doho}, $\\alpha_\\pi- \\beta_\\pi =\\alpha(\\bar l_6-\\bar l_5)/(24 \\pi^2 f_\\pi^2 m_\\pi) +{\\cal O}(m_\\pi)$, is of course the  same after identifying the combination of low-energy constants as $\\bar l_6-\\bar l_5 = 6h_A/h_V +{\\cal O} (m^2_\\pi)$. Recently, the systematic corrections to this current algebra result have been worked out in refs.\\cite{buergi,gasser} by performing a full two-loop  calculation of pion Compton scattering in chiral perturbation theory. The outcome of that extensive analysis is that altogether the higher order corrections are rather small and the value $\\alpha_\\pi- \\beta_\\pi=(5.7\\pm 1.0)\\cdot 10^{-4}\\,$fm$^3$ \\cite{gasser} for the pion polarizability difference stands now as a firm prediction of the (chiral-invariant) theory. The non-vanishing value $\\alpha_\\pi +\\beta_\\pi=(0.16\\pm 0.1) \\cdot 10^{-4}\\,$fm$^3$ for the pion polarizability sum (obtained also at two-loop order) is presumably too small to cause an observable effect in low-energy pion Compton scattering.  However, the chiral prediction $\\alpha_\\pi- \\beta_\\pi=(5.7\\pm 1.0)\\cdot 10^{-4}\\,$fm$^3$ is in conflict with the existing experimental determinations of $\\alpha_\\pi- \\beta_\\pi=(15.6\\pm 7.8) \\cdot 10^{-4}\\,$fm$^3$ from Serpukhov \\cite{serpukhov} and $\\alpha_\\pi-\\beta_\\pi =(11.6\\pm 3.4)\\cdot 10^{-4}\\,$fm$^3$ from Mainz \\cite{mainz}, which amount to values more than twice as large. These existing experimental determinations of $\\alpha_\\pi-\\beta_\\pi$ certainly raise doubt as to their correctness since they violate the chiral low-energy theorem notably by a factor 2. In that contradictory situation it is promising that the ongoing COMPASS \\cite{compass} experiment at CERN aims at measuring the pion polarizabilities with high statistics using the Primakoff effect. The scattering of high-energy negative pions in the Coulomb field of a heavy nucleus gives (in the region of sufficiently small photon virtualities) access to cross sections for $\\pi^- \\gamma$ reactions. As an alternative, one could also directly analyze the bremsstrahlung process $ \\pi^- Z\\to\\pi^- Z\\gamma$, omitting the whole kinematical extrapolation from virtual to real photons.  In the appendix we will write down the corresponding fivefold differential  cross section $d^5\\sigma/d\\omega d\\Omega_\\gamma d\\Omega_\\pi$ including Born terms, the pion polarizability difference $\\alpha_\\pi-\\beta_\\pi$, and an (equally) important pion-loop correction.  In any case, the effects of the pion's low-energy structure on (real or virtual) Compton scattering observables turn out to be relatively small. For center-of-mass energies $\\sqrt{s} < 4m_\\pi$ (i.e. sufficiently below the prominent $\\rho(770)$-resonance) the differential cross sections $d\\sigma/ d\\Omega_{\\rm cm}$ in backward directions are reduced at most by about $11\\%$ in comparison to the ones of a structureless pion \\cite{picross}. Therefore, a precise knowledge of the pure QED radiative corrections to pion Compton scattering is indispensable if one wants to extract the pion polarizabilities from the data with good accuracy. In certain kinematical regions the effects from the pion's low-energy  structure and the pure QED radiative corrections may become of comparable size. Although calculations of radiative corrections are abundant in the literature, we are not aware of a detailed and utilizable exposition of the radiative corrections to scalar (spin-0) boson Compton scattering. The case of spin-1/2 electron Compton scattering has been treated already in the early days of quantum electrodynamics by Brown and Feynman \\cite{feyn}. The radiative corrections to the Thomson limit, valid close to threshold, have been reported in ref.\\cite{corin}.  In the work of Akhundov et al. \\cite{akhundov} the one-photon loop diagrams to virtual pion Compton scattering have been considered, but no accessible sources to the corresponding analytical expressions (which are necessary for an implementation into data analyses) are given. Since those previous results were mainly presented in numerical form it is difficult to implement them independently into future data analyses. As a consequence of that deficit the radiative corrections to pion Compton scattering are sometimes merely adapted from the known ones for myons by simply replacing the mass: $m_\\mu \\to m_\\pi$. Of course, such a substitutional procedure mistreats profound differences in the couplings of photons to scalar spin-0 bosons and spin-1/2 fermions.  The purpose of the present paper is to fill this gap and to present a detailed calculation of the one-photon loop radiative corrections to pion Compton scattering, $\\pi^- \\gamma \\to \\pi^- \\gamma$. We give closed-form analytical expressions for the ${\\cal O}(\\alpha)$ corrections to the two invariant amplitudes $A(s,u)$ and  $B(s,u)$ as they emerge from 11 classes of contributing one-loop diagrams. Ultraviolet and infrared divergencies are both treated by the method of dimensional regularization which ensures gauge invariance at every step. While the ultraviolet divergences can be absorbed in renormalization constants, the cancellation of infrared divergencies requires the inclusion of soft photon radiation below an experimental detection threshold $\\lambda$. The total finite radiative correction depends then logarithmically on the small energy resolution scale $\\lambda$. We find that the radiative corrections become maximal in backward directions, reaching about $-2.4\\%$ at a center-of-mass energy of $\\sqrt{s} = 4m_\\pi$ for  $\\lambda =5\\,$MeV. With such a size and kinematical signature the radiative corrections are not negligible in comparison to the effects from the pion polarizabilities, which also show up preferentially in the backward directions. Furthermore, we include in our calculation of the radiative corrections also the leading pion structure effect through a two-photon contact-vertex proportional to the pion polarizability difference $\\alpha_\\pi -\\beta_\\pi \\simeq 6 \\cdot 10^{-4}\\,$fm$^3$. We find that these structure effects do not modify the relative size and angular dependence of the radiative  corrections. Our results can be utilized for analyzing the COMPASS experiment at CERN. ",
        "conclusions": ""
    },
    "0806/0806.0575_arXiv.txt": {
        "abstract": "Warm inflationary universe models in the context of a  Chaplygin gas equation  are studied. General conditions required for these models to be realizable are derived and discussed. By using a chaotic potential we develop  models for a dissipation coefficient of the form  $\\Gamma\\propto\\,\\phi^n$, with $n=0$ or $n\\neq 0$. ",
        "introduction": "It is well  know that warm inflation, as opposed to the conventional cool inflation, presents the attractive feature that it avoids the reheating period \\cite{warm}. In these kind of models dissipative effects are important during the inflationary period, so that radiation production occurs concurrently together with the inflationary expansion. If the radiation field is in a highly excited state during inflation, and this has a strong damping effect on the inflaton dynamics, then, it is found a strong regimen of  warm inflation. Also, the dissipating effect arises from a friction term which describes the processes of the scalar field dissipating into a thermal bath via its interaction with other fields. Warm inflation shows how thermal fluctuations during inflation may play a dominant role in producing the initial fluctuations  necessary for Large-Scale Structure (LSS) formation. In these kind of models the density fluctuations arise from thermal rather than quantum fluctuations \\cite{62526}. These fluctuations have their origin in the hot radiation and influence the inflaton through a friction term in the equation of motion of the inflaton scalar field \\cite{1126}. Among the most attractive features of these models, warm inflation end  at the epoch when the universe stops inflating and \"smoothly\" enters in a radiation dominated Big-Bang phase\\cite{warm}. The matter components of the universe are created by the decay of either the remaining inflationary field or the dominant radiation field \\cite{taylorberera}.   On the other hand, the generalized Chaplygin gas    has been proposed as an  alternative model for  describing the present accelerating of the universe. The generalized Chaplygin gas is described by an exotic equation of state of the form \\cite{Bento} \\begin{equation} p_{ch} = - \\frac{A}{\\rho_{ch}^\\beta},\\label{1} \\end{equation} where $\\rho_{ch}$ and $p_{ch}$ are the energy density and pressure of the generalized Chaplygin gas, respectively. $\\beta$ is a constant that lies in  the range $0 <\\beta\\leq 1$, and $A$ is a positive constant. The original  Chaplygin gas corresponds to the case $\\beta = 1$ \\cite{2}. The above equation of state leads to a density evolution in the form \\cite{Bento}  \\begin{equation} \\rho_{ch}=\\left[A+\\frac{B}{a^{3(1+\\beta)}}\\right]^{\\frac{1}{1+\\beta}}, \\label{2} \\end{equation} where $a$ is the scale factor  and $B$ is a positive integration constant.  The Chaplygin gas emerges as a effective fluid of a generalized d-brane in a (d+1, 1) space time, where the   action can be written as a generalized Born-Infeld action \\cite{Bento}. These models have been extensively studied in the literature \\cite{other}. The model parameters were constrained using current cosmological observations, such as, CMB \\cite{CMB} and  supernova of type Ia (SNIa) \\cite{SIa}.    In the  model of  Chaplygin inspired inflation usually the scalar field, which drives inflation, is the standard inflaton field, where the energy density given by Eq.(\\ref{2}), can be extrapolate for  obtaining a successful inflationary period\\cite{Ic}. Recently, tachyon-Chaplygin inflationary universe model was studied in Ref.\\cite{SR}, and the dynamics of the early universe and the initial conditions for inflation in a model with radiation and a Chaplygin gas was considered  in Ref.\\cite{Monerat:2007ud}. The main goal of the present work is to investigate the possible realization of a warm-Chaplygin inspired inflationary model, where the universe is filled with a self-interacting scalar field and a radiation field. We use astronomical data for constraining the parameters appearing in this model. Specifically, the   parameters are constrained from the WMAP observations\\cite{WMAP3,WMAP}.        The outline of the paper is a follows. The next section presents a short review of  the modified  Friedmann equation and the warm-Chaplygin Inflationary phase. Section \\ref{sectpert} deals with the scalar and tensor perturbations.  In Section \\ref{exemple} it is presented a chaotic potential in the high dissipation approximation. Here, we give  explicit expressions for scalar power spectrum and tensor-scalar ratio for our models. Finally, sect.\\ref{conclu} summarizes our findings. We chose units so that $c=\\hbar=1$. ",
        "conclusions": "}  In this paper we have investigated the  warm-Chaplygin inflationary scenario. In the slow-roll approximation we have found a general relationship between the radiation and scalar field energy densities. This has led us to a general criterium for warm-Chaplygin inflation  to occur.     Our  specific models are described by  a chaotic potential and we have consider  different  cases for   the dissipation coefficient, $\\Gamma$. In the first case, we took $\\Gamma=constant=\\Gamma_0 $. Here, we have found  solutions for the inflaton field and the  Hubble parameter. The relation between  the radiation field and the inflaton field energy densities   presents the same characteristic to that corresponding to the   warm inflation case, except that it  depends on the extra parameter  $A$. For the case in which the dissipation coefficient $\\Gamma$ is taken to be a function of the scalar field, i.e. $\\Gamma \\propto \\phi^n$, it was possible to describe an  appropriate warm inflationary universe model for $n=1$ and $n=2$. In these cases, we have obtained  explicit expressions for the corresponding scalar spectrum  and the running of the scalar spectrum indices.  Finally, by using the WMAP five year data\\cite{WMAP} and for the special case where $T\\simeq T_r\\simeq0.24\\times 10^{16}$GeV, $k_*=0.002$Mpc$^{-1}$ and  $n_s\\simeq 0.95$,  we have found new constraints on the parameters $\\alpha_n$, $A$ and $m$."
    },
    "0806/0806.1059_arXiv.txt": {
        "abstract": "We present the discovery of two nearby L dwarfs from our 2MASS proper motion search, which uses multi-epoch 2MASS observations covering $\\sim$4700 square degrees of sky.  2MASS J18212815+1414010 and 2MASS J21481628+4003593 were overlooked by earlier surveys due to their faint optical magnitudes and their proximity to the Galactic Plane (10$^{\\circ}$ $\\le$ $|$b$|$ $\\le$ 15$^{\\circ}$). Assuming that both dwarfs are single, we derive spectrophotometric distances of $\\sim$10 pc, thus increasing the number of known L dwarfs within 10 pc to 10.  In the near-infrared, 2MASS J21481628+4003593 shows a triangular-shaped $H$-band spectrum, strong CO absorption, and a markedly red $J-K_s$ color (2.38~$\\pm$~0.06) for its L6 optical spectral type. 2MASS J18212815+1414010 also shows a triangular-shaped $H$-band spectrum and a slightly red $J-K_s$ color (1.78~$\\pm$~0.05) for its L4.5 optical spectral type. Both objects show strong silicate absorption at 9--11 $\\mu$m. Cumulatively, these features imply an unusually dusty photosphere for both of these objects.  We examine several scenarios to explain the underlying cause for their enhanced dust content and find that a metal-rich atmosphere or a low-surface gravity are consistent with these results.  2MASS J18212815+1414010 may be young (and therefore have a low-surface gravity) based on its low tangential velocity of 10 km s$^{-1}$.  On the other hand, 2MASS J21481628+4003593 has a high tangential velocity of 62 km s$^{-1}$ and is therefore likely old.  Hence, high metallicity and low-surface gravity may lead to similar effects. ",
        "introduction": "Many of the objects known to reside in the Solar Neighborhood were discovered via their high proper motion. Early proper motion surveys by the pioneers of photographic astrometry built the foundation of our knowledge about the local stellar census (e.g., \\citealt{luyten1979, luyten1980, giclas1978, ross1939, wolf1918, barnard1916, innes1915}).  Discoveries of our nearest stellar and substellar neighbors continue to-date, with surveys such as the SIPS survey \\citep{deacon2007, deacon2005}, the LEHPM survey \\citep{pokorny2004}, the SCR survey \\citep{finch2007}, and the SUPERBLINK survey \\citep{lepine2008b, lepine2008a, lepine2002}. Despite their successes, these surveys are not ideally suited to the discovery of nearby brown dwarfs because of their dependence on optical data.  As we have now learned from longer-wavelength surveys such as the Two Micron All Sky Survey (2MASS; \\citealt{skrutskie2006}), the Deep Near Infrared Survey of the southern sky (DENIS; \\citealt{epchtein1997}), and the Sloan Digital Sky Survey (SDSS; \\citealt{york2000}), the Solar Neighborhood is home to many L and T dwarfs -- the coldest brown dwarfs currently known \\citep{kirkpatrick2005}. Most of these cool objects have been selected photometrically using color selections appropriate for L and T dwarfs with near solar metallicity. This method may be missing low-temperature objects with unusual physical properties -- low or high metallicity dwarfs or objects with unusual atmospheric traits --  but a measurement of their high motions will help to identify them as nearby.  Another benefit of proper motion surveys over photometric ones is their ability to find nearby objects against the confusion of the Galactic Plane.  Due to reddening caused by intervening material, extinguished background objects in the Plane can have colors mimicking those of nearby dwarfs in 2MASS, DENIS, and SDSS colors. Proper motions help ferret out these nearby interlopers without relying on color information.  Using data contained in the 2MASS Survey Working Databases, we have searched for proper motion objects by examining the subset of multi-epoch data having time differences in excess of two months. The total coverage satisfying this criterion is approximately 4700 square degrees, or $\\sim$11\\% of the sky. In this paper we present two discoveries from this exclusively near-infrared search. Both objects are optically faint, nearby L dwarfs located within 15$^\\circ$ of the Galactic Plane and hence missed by previous optical proper motion and near-infrared photometric surveys.  We briefly describe our 2MASS proper motion search and measurements from the 2MASS catalog in $\\S$2, along with the observations; we give the analysis in $\\S$3; discussion in $\\S$4; and conclusions in $\\S$5. ",
        "conclusions": "We have presented spectral classification and distance estimates for two nearby and peculiar L dwarfs identified in our $\\sim$4700 sq.\\ degree 2MASS proper motion search. With optical spectral types of L4.5 and L6, we estimate distances of $\\sim$10 pc for both 2MASS J1821+1414 and 2MASS J2148+4003, although 2MASS J2148+4003 may actually be closer. Both objects have overall red spectral energy distributions, weak H$_2$O bands, slightly triangular-shaped $H$-band continua, and silicate absorption at 9--11 $\\mu$m. These features are likely due to unusually thick dust clouds in their atmospheres.  While the underlying cause for this dust is uncertain, we suggest that it may arise from a metal-rich atmosphere or a lower surface gravity, the latter implying a young age. The former cause is more likely in the case of 2MASS J2148+4003, whose large v$_{tan}$ indicates that it is not a young source. Further follow-up is needed to distinguish between these two scenarios and to identify clear diagnostics for characterizing the detailed physical properties of other L dwarfs in the vicinity of the Sun."
    },
    "0806/0806.1090_arXiv.txt": {
        "abstract": "Elliptical galaxies and globular clusters (GCs) have traditionally been regarded as physically distinct entities due to their discontinuous distribution in key scaling diagrams involving size, luminosity and velocity dispersion. Recently this distinctness has been challenged by the discovery of stellar systems with mass intermediate between those of GCs and dwarf ellipticals (such as Ultra Compact Dwarfs and Dwarf Galaxy Transition Objects). Here we examine the relationship between the virial and stellar mass for a range of old stellar systems, from GCs to giant ellipticals, and including such Intermediate Mass Objects (IMOs). Improvements on previous work in this area include the use of (i) near-infrared magnitudes from the 2MASS survey, (ii) aperture corrections to velocity dispersions, (iii) homogeneous half light radii and (iv) accounting for the effects of non-homology in galaxies. We find a virial-to-stellar mass relation that ranges from $\\sim$10$^4$ M$_{\\odot}$ systems (GCs) to $\\sim$10$^{11}$ M$_{\\odot}$ systems (elliptical galaxies). The lack of measured velocity dispersions for dwarf ellipticals with --16 $>$ M$_K$ $>$ --18 ($\\sim$10$^8$ M$_{\\odot}$) currently inhibits our ability to determine how, or indeed if, these galaxies connect continuously with GCs in terms of their virial-to-stellar mass ratios. We find elliptical galaxies to have roughly equal fractions of dark and stellar matter within a virial radius; only in the most massive (greater than 10$^{11}$ M$_{\\odot}$) ellipticals does dark matter dominate the virial mass. Although the IMOs reveal slightly higher virial-to-stellar mass ratios than lower mass GCs, this may simply reflect our limited understanding of their IMF (and hence their stellar mass-to-light ratios) or structural properties. We argue that most of these intermediate mass objects have similar properties to massive GCs, i.e. IMOs are essentially massive star clusters. Only the dwarf spheroidal galaxies exhibit behaviour notably distinct from the other stellar systems examined here, i.e. they display a strongly increasing virial-to-stellar mass ratio (equivalent to higher dark matter fractions) with decreasing stellar mass. The data used in this study is available in electronic format. ",
        "introduction": "The scalings between basic parameters such as the size, luminosity or surface brightness, and line-of-sight velocity dispersion of stellar systems have provided a key tool in which to understand self-gravitating systems. When viewed as a 3 parameter space, they are collectively know as the Fundamental Plane (\\cite{djorgovski87}). Such scaling relations have been used to probe the structural properties, origin, and even to classify objects depending on where they lie in parameter space. The scalings in `$\\kappa$-space' (with axes related to mass, mass-to-light ratio and surface brightness) of dynamically hot galaxies were explored by \\cite{bender92}. These hot systems included elliptical, dwarf spheroidal and the bulges of spiral galaxies. \\cite{burstein97} extended the $\\kappa$-space analysis to include disk galaxies, groups and clusters of galaxies and globular clusters (GCs). More recently, \\cite{zaritsky06} defined the Fundamental Manifold of spheroids revealing a continuity from clusters of galaxies to dwarf ellipticals, and possible extension to dwarf spheroidals.   Like elliptical galaxies, GCs are self-gravitating systems with a strong component of pressure support from the random motions of their stars (i.e. they are dynamically hot) and are dominated by stars of old age (i.e. older than 10 Gyrs). However, they were either excluded from past studies (e.g. \\cite{bender92}; \\cite{zaritsky06}) or found to be distinct entities based on their different scalings and large separation in mass from galaxies (e.g. \\cite{kormendy85}; \\cite{burstein97}).  Only in the last decade have objects of mass intermediate between those of massive GCs and dwarf ellipticals been discovered (\\cite{hilker99}; \\cite{drinkwater00}). These objects have masses of $\\sim$ 10$^{7}$ M$_{\\odot}$ and relatively compact sizes with measured half light radii of $\\le$50 pc. They are usually referred to as Ultra Compact Dwarfs (UCDs) or Dwarf Globular Transition Objects (DGTOs). Although they share many properties with the nuclei of nucleated dwarf galaxies (e.g. \\cite{geha03,cote06,boker08}) they also resemble massive GCs (e.g. \\cite{kisslerpatig06,hasegan05,hilker07,gilmore07}). A number of papers have proposed various possible mechanisms to explain such intermediate mass objects (IMOs; we prefer this terminology as it describes their physical state and not their uncertain origin). These include the remnant nucleus of a stripped dwarf galaxy (\\cite{drinkwater03}; \\cite{bekki03}), the merger of several smaller GCs (\\cite{fellhauer02}; \\cite{bekki04}), a completely new type of galaxy (\\cite{drinkwater00}) or an extension of the GC sequence to higher masses (\\cite{mieske02}). However, each of these possible explanations has difficulties (e.g. \\cite{goerdt08}; \\cite{evstigneeva08}).  A number of authors have examined the scaling relations of IMOs, sometimes including GCs and galaxies in their analysis (\\cite{martini04,hasegan05,evstigneeva07,hilker07}).  A recent work in this fast moving field is that of \\cite{dabringhausen08} (hereafter D08) who include GCs and giant ellipticals, but focus on IMOs and dwarf galaxies. They confirm a transition in globular cluster and IMO properties (to larger sizes, higher stellar densities and higher inferred mass-to-light ratios) at a mass of $\\sim$10$^6$ M$_{\\odot}$. They interpret this as either as evidence for dark matter or a different initial mass function (IMF) in these somewhat higher mass objects. D08 also included dwarf spheroidal (dSph) galaxies in their analysis. These galaxies have similar velocity dispersions to globular clusters but very high inferred mass-to-light ratios. Debate continues whether these high ratios are due to tidal heating or large dark matter halos (see for example \\cite{penarrubia08} and \\cite{metz08}), and how such galaxies fit into the general scaling relations.  After submission of our paper to the journal, the work of \\cite{mieske09} was made publicly available. This work discusses the nature of UCDs focusing on their internal dynamics and re-examining various UCD scaling relations. In the appropriate sections of this paper we comment on the \\cite{mieske09} results. In general, they reach similar conclusions to us.     Here we focus on the relationship between virial and stellar mass for a wide mass range of old, pressure-supported systems. In particular, we examine elliptical galaxies and  globular clusters along with IMOs and dwarf spheroidals. In general, such systems contain little, if any, cold or hot gas and so the stellar mass is a good proxy for the baryonic mass in these systems. They are usually dominated, in mass, by old stellar populations. We also apply several improvements on previous work through the use of:\\\\ (i) near-infrared magnitudes which are a much better tracer of stellar mass than optical light;\\\\ (ii) aperture corrections to the literature velocity dispersions of unresolved GCs and IMOs to reflect central values; \\\\ (iii) half light radii that account for the deviations in galaxy light profiles from the simple R$^{1/4}$ law; and \\\\ (iv) variations to the calculated virial mass for non-homology effects between galaxies.  In Section 2, we describe the physical parameters which we use, while Section 3 lists the data samples (which are given in Table 1 and available fully in electronic form). Section 4 presents the scaling relations of both velocity dispersion and radius with near-infrared luminosity before examining the virial versus stellar mass relation. Finally, in Section 5, we highlight prospects for future work and give our conclusions. ",
        "conclusions": "We have collected various data samples from the literature for old, pressure-supported systems which includes globular clusters, massive star clusters, intermediate mass objects (such as ultra compact dwarfs), dwarf spheroidals, dwarf ellipticals and giant ellipticals, and covers a range in mass from $\\sim$10$^4$ to 10$^{12}$ M$_{\\odot}$. We have applied several improvements on past work that has examined their virial and stellar masses. We have employed aperture corrections to the velocity dispersion measurements of GCs and intermediate mass objects. We have also derived new half light radii for elliptical galaxies based on their sizes and ellipticities from the homogeneous 2MASS survey. Near-infrared magnitudes from the 2MASS survey are converted into total stellar masses using a K-band mass-to-light ratio that depends on the metallicity or colour of the object. Virial masses are calculated taking into account non-homology effects for galaxies.   Although the scalings of velocity dispersion and half light radius with absolute K-band magnitude vary depending on the mass regime probed, these scalings combine to give a virial versus stellar mass relation that shows a remarkable near continuous trend from GCs to ellipticals. We find that the Fundamental Manifold of \\cite{zaritsky06} is not a good representation for GCs. Dwarf and normal elliptical galaxies are found to have virial-to-stellar mass ratios of $\\sim$2:1. This ratio only increases in the very most massive ellipticals, with masses greater than 10$^{11}$ M$_{\\odot}$. Our results are subject to systematic effects from remaining uncertainties, e.g. in the distribution of dark matter, the accuracy of 2MASS total K-band magnitudes, the appropriate IMF to use etc. However, such trends are generally consistent with results from strong lensing studies.    The recently discovered intermediate mass ($\\sim$ 10$^{7}$ M$_{\\odot}$) objects, e.g. Ultra Compact Dwarfs and Dwarf Globular Transition Objects, cover the same parameter space of velocity dispersion, half light radius and mass as massive globular clusters, possible dwarf galaxy nuclei, massive star clusters and the nuclei of dE,N galaxies. To date, these intermediate mass objects do not exceed the maximum mass of known star clusters in galaxies (\\cite{whitmore00}) and they are spatially concentrated near large galaxies (\\cite{wehner07}). All of these facts would support an interpretation that intermediate mass objects are essentially massive star clusters.  Given that there is no evidence in the literature for dark matter in massive GCs, this would also argue against dark matter in intermediate mass objects as they occupy a similar parameter space. However a mystery remains, in that intermediate mass objects (and massive GCs) exhibit higher virial-to-stellar mass ratios, when we apply the same virial coefficient and Chabrier IMF as low mass GCs. Possible solutions to this mystery may include a different virial coefficient due to a longer relaxation timescale in these systems, or that a bottom heavy Salpeter-like IMF is more appropriate in these objects.  In general agreement with D08, we find dwarf spheroidal galaxies to be distinct in terms of their scaling parameters, following neither an obvious extension of the elliptical galaxy or globular cluster relations. Their virial-to-stellar mass ratios reach one thousand.   Although we have probed mass scales from $\\sim$10$^4$ to 10$^{12}$ M$_{\\odot}$, there is a mass regime which remains largely unexplored observationally, i.e. $\\sim$10$^8$ M$_{\\odot}$. This is greater than the most massive star clusters and intermediate mass objects known but less than the mass of dwarf ellipticals for which velocity dispersions are available. It is not clear if the virial versus stellar mass relation will connect smoothly across this mass gap between dwarf ellipticals and globular clusters. Thus an observational campaign to measure central velocity dispersions for a sample of very low mass dwarf ellipticals is needed. Observations are also needed to determine whether the radial velocity dispersion profiles in intermediate mass objects are flat or fall with radius, like their light profiles. Obtaining both high spectral and spatial resolution for such small, low surface brightness systems will be observationally challenging. Near-infrared spectra that could constrain the IMF would also be useful (\\cite{mieske08}).  On the theoretical side, for massive globular clusters and intermediate mass objects we require a detailed understanding of how the virial coefficient varies with the mass and/or type of stellar system, and the expected stellar mass-to-light ratio that includes the effects of multiple stellar populations and dynamical evolutionary processes."
    },
    "0806/0806.3789_arXiv.txt": {
        "abstract": "We present cosmological magnetohydrodynamic simulations of the formation of a galaxy cluster with magnetic energy feedback from an active galactic nuclei (AGN). We demonstrate that X-ray cavities can be produced by the magnetically dominated jet-lobe system that is supported by a central axial current. The cavities are magnetically dominated and their morphology is determined jointedly by the magnetic fields and the background cluster pressure profile.  The expansion and motion of the cavities are driven initially by the Lorentz force of the magnetic fields, and the cavities only become buoyant at late stages ($> 500$ Myr). We find that up to $80\\%-90\\%$ of the injected magnetic energy goes into doing work against the hot cluster medium, heating it, and lifting it in the cluster potential. ",
        "introduction": "The absence of spectral signatures of cooling plasmas at the galaxy cluster centers (Tamura et al. 2001; Peterson et al. 2003) has led to the suggestion that the intra-cluster medium (ICM) in cluster centers must be heated.  Powerful radio jet-lobes emanating from supermassive black holes (SMBHs) in AGNs of clusters are considered to be the promising heating sources (Binney \\& Tabor 1995; Tucker \\& David 1997; see also McNamara \\& Nulsen 2007 for a recent review).  High resolution X-ray images of galaxy clusters by {\\it Chandra} have revealed giant cavities and weak shock fronts in the hot gas (Fabian et al. 2000; McNamara et al. 2000, 2005), which are commonly associated with energetic radio lobes (Blanton et al. 2005; Nulsen et al. 2002) and suggest that magnetic fields play an important role.  Large uncertainties concerning the nature of these cavities, their formation, evolution, and survivability in the ICM still remain.  Numerical simulations of hot, underdense bubbles in galaxy clusters have been performed by a number of authors (e.g., Churazov et al.  2001; Reynolds, Heinz, \\& Begelman 2001; Br\\\"{u}ggen \\& Kaiser 2002; Omma et al. 2004).  It is generally possible to inject a large amount of energy into the ICM via AGNs but it is not exactly clear how the AGN energy can be efficiently utilized (Vernaleo \\& Reynolds 2006; though see Heinz et al.  2006).  One of the most interesting characteristics of the radio bubbles is that they are intact, whereas most hydrodynamic simulations \\citep{Quilis01,Bruggen02,Dalla04} have shown that purely hydrodynamic bubbles will disintegrate in timescales much less than $10^8$ yrs, markedly different from observations.  (we will use X-ray cavities, bubbles, and radio bubbles interchangeably in this {\\it Letter}).  The stabilizing role of magnetic fields has been suggested and studied by a few authors (e.g., Jones \\& De Young 2005).  A rather different class of models has been proposed and studied, in which the AGN energy output is modeled in the magnetically dominated limit (Li et al. 2006; Nakamura, Li, \\& Li 2006; see also the previous work of Blandford 1976, Lovelace 1976, Lynden-Bell 1996, Li et al. 2001). The key feature of this model is to inject simultaneously both the poloidal and the (more dominant) toroidal magnetic fields in a small volume. This is to mimic the possible outcome of an accretion disk dynamo around an SMBH that shears and twists up the poloidal magnetic fields and generates large amounts of toroidal fields with an axial current (as high as $10^{19}$ amperes) flowing along the central axis of this magnetic structure.  The injection of magnetic fields and their associated currents lasts a finite time (mimicking the lifetime of an AGN), after which the magnetic fields and their currents will no longer be injected but will continue to evolve and gradually dissipate away. This global current is essential to maintaining the magnetic structure throughout the lifetime of a magnetic field and its associated current. The system is not force-free initially, and so the fields will self-collimate and expand predominately axially, producing a collimated structure reminiscent of a ``magnetic tower''.  Extensive 3-D magnetohydrodynamic (MHD) simulations based on this model in a static cluster-like background, have demonstrated that such magnetically dominated structures can reproduce some of the global features of the jet-lobe systems, especially in maintaining the integrity of the bubbles.  In this {\\it Letter}, we present cosmological MHD simulations of a cluster formation with the feedback of an AGN, with the aim of understanding the X-ray cavity formation using the magnetically dominated models proposed by Li et al. (2006). This differs from all previous studies \\citep{Bruggen02,Heinz06} in that the formation of the X-ray cavity is studied in a realistic and self-consistent cosmological setting where the dark matter, baryon dynamics and magnetic fields are all evolved self-consistently. In \\S \\ref{sec:method}, we describe our approach and the parameters of the simulations. In \\S \\ref{sec:result}, we present the key results. Conclusions and discussions are given in \\S \\ref{sec:dis}. ",
        "conclusions": "\\label{sec:dis}  With our cosmological MHD simulations, we find that, in the realistic cluster environment where the ICM plasma interacts dynamically with the magnetic jet-lobe, X-ray cavities can naturally form using the magnetically dominated models proposed by Li et al. (2006). The magnetic fields inside the bubbles stabilize the interface instabilities so that bubbles can remain intact. The lifetime of these bubbles can be quite long and only become truly buoyant probably after $\\sim 500$ Myr.  We have performed additional cosmological MHD simulations with radiative cooling and star formation feedback and found that our conclusions for the X-ray cavity formation mechanism do not change. While we have demonstrated the formation of X-ray cavities, much more studies are needed in order to address comprehensively the cooling flow problem at the cluster cores.  The present simulation, with just one AGN, already has important implications for understanding the ICM heating problem. Up to $80\\%-90\\%$ of the injected energy has been dissipated in the surrounding ICM. Further studies are underway with AGNs at different redshifts that have different cluster environments so that we can gain a comprehensive understanding of the overall heating of the ICM by AGNs. These simulations will be presented in future publications.  The morphology of the jet-lobe is dependent on the background density radial profile, which is different for massive clusters (such as the one presented here) and groups or poor clusters. Future work will address this issue."
    },
    "0806/0806.2811_arXiv.txt": {
        "abstract": "The rapidly oscillating Ap (roAp) star 10\\,Aql shows one of the lowest photometric pulsation amplitudes and is characterized by an unusual spectroscopic pulsational behavior compared to other roAp stars. In summer 2006 this star became target of an intense observing campaign, that combined ground-based spectroscopy with space photometry obtained with the MOST satellite. More than 1000 spectra were taken during 7 nights over a time span of 21 days with high-resolution spectrographs at the 8-m ESO VLT and 3.6-m TNG telescopes giving access to radial velocity variations of about 150 lines from different chemical species. A comparison of pulsation signatures in lines formed at different atmospheric heights allowed us to resolve the vertical structure of individual pulsation modes in 10\\,Aql\\ which is the first time for a multiperiodic roAp star. Taking advantage of the clear oscillation patterns seen in a number of rare-earth ions and using the contemporaneous MOST photometry to resolve aliasing in the radial velocity measurements, we improve also the determination of pulsation frequencies. The inferred propagation of pulsation waves in 10\\,Aql is qualitatively similar to other roAp stars: pulsation amplitudes become measurable in the layers where Y and Eu are concentrated, increase in layers where the H$\\alpha$ core is formed, reach a maximum of 200--300\\,\\ms\\ in the layers probed by Ce, Sm, Dy lines and then decrease to 20--50\\,\\ms\\ in the layers where \\nd\\ and \\pr\\ lines are formed. A unique pulsation feature of 10\\,Aql is a second pulsation maximum indicated by Tb\\iii\\ lines which form in the uppermost atmospheric layers and oscillate with amplitudes of up to 350\\,\\ms. The dramatic decline of pulsations in the atmospheric layers probed by the strong \\pr\\ and \\nd\\ lines accounts for the apparent peculiarity of 10\\,Aql when compared to other roAp stars. The phase-amplitude diagrams and bisector measurements of the \\nd\\ 5102\\,\\AA\\ line reveal a rapid change of phase and amplitude with height for all three main pulsation modes, indicating the presence of a pulsation node in the stellar atmosphere. Finally, we report the discovery of a puzzling asymmetry of the strong \\nd\\ lines with their blue wing extending up to $-50$\\,\\kms, which is about 25 times the estimated value of \\vs. ",
        "introduction": "\\label{intro}  More than 30 members among the group of late A magnetic chemically peculiar stars exhibit high-overtone, low-degree, non-radial $p$-mode pulsations with periods in the range of 6--21 minutes \\citep{KM00}. These \\textit{rapidly oscillating Ap} (roAp) stars are characterized by strong global magnetic fields with a polar strengths of the order of 1--10 kG.  The atmospheres of roAp stars are enriched with heavy elements, brought  up from the stellar interior by diffusion. Conspicuous lateral and vertical variations of chemical abundances in Ap stars and the prominence of the magnetic field signatures make them primary targets for detailed investigations of the magnetic topology and magnetically driven formation of structures in stellar atmospheres (\\citealt{K04a}; \\citealt{R04}).  The observed pulsational amplitudes of roAp stars are modulated according to the visible magnetic field structure. This observation led to the oblique pulsator model \\citep{k1982}, where axisymmetric $\\ell=1$ modes are aligned with the magnetic field axis, which itself is oblique to the axis of stellar rotation. Calculations by \\citet{SG04} showed that dipolar modes are significantly distorted by a magnetic field of kG-strength, but retain their axisymmetric character. On the other hand, the structure of pulsational perturbations can include significant non-axisymmetric components in weak-field stars \\citep{BD02}. The indirect pulsation Doppler imaging of oscillations in the prototype roAp stars HR\\,3831 \\citep{K04b} vindicated the oblique pulsator model and for the first time characterized observationally the magnetic distortion of the global {\\it p-}modes.  The roAp stars are key objects for asteroseismology, which presently is the most powerful tool for testing theories of stellar structure and evolution. The classical asteroseismic analysis, utilizing precise frequency measurements, helps to constrain the luminosity and internal chemical composition of pulsating magnetic stars (e.g. Matthews, Kurtz \\& Martinez \\citeyear{MKM99}; Cunha, Fernandes \\& Monteiro \\citeyear{CFM03}). Recent time-resolved spectroscopic observations of roAp stars uitilizing large telescopes  (see review by \\citealt{K07a}) demonstrated the possibility of another type of asteroseismic investigation, which is focused on the upper atmospheric layers. A significant chemical stratification and a short vertical wavelength of pulsation modes lead to a remarkable diversity of the pulsational characteristics observed in individual spectral lines, notably of absorption lines of the rare-earth elements (REE) (e.g., \\citealt{KR01}; Mkrtichian, Hatzes \\& Kanaan \\citeyear{MHK03}; \\citealt{ryab2007a}). The information extracted from the lines formed at different optical depths opens access to different modes and can be combined to yield a vertical tomographic map of the pulsating stellar atmosphere \\citep{ryab2007b}. Thus, applying Doppler imaging techniques and pulsation tomography to line profile variation in roAp stars promises an unprecedented 3-D picture of  magnetoacoustic pulsations \\citep{K05}.  These spectacular results have been obtained despite the fact that most spectroscopic observations of roAp stars were obtained in a snapshot mode and hence cover only a few hours of stellar oscillations. A detailed frequency analysis is highly ambiguous with such limited data, because of the well known aliasing problem. It can only marginally improved by combining observations from several nights, but which rarely are available when large telescopes had to be used. Therefore, despite an outstanding clarity of the pulsation curves often produced with spectroscopy, for most multiperiodic roAp stars they correspond to an unresolved mixture of different modes and thus are exceedingly difficult to interpret. This fundamental problem can be alleviated by simultaneous photometric observations with a high duty cycle, as it was successfully demonstrated by a combination of space photometry of the roAp star HD\\,24712 with contemporaneous VLT/UVES spectroscopy \\citep{ryab2007a}. In this paper we present an analysis of a much more extensive spectroscopic time series for another pulsating Ap star, 10\\,Aql, which was observed simultaneously from ground and space.  10\\,Aql (HR\\,7167, HD\\,176232, HIP\\,93179) is one of the brightest roAp stars. \\citet{RSH00} performed its model atmosphere and abundance analysis using high-resolution spectra. They have estimated \\teff\\,=\\,7550\\,K, \\lgg\\,=\\,4.0, and derived $M$\\,=\\,$2.0\\pm0.2$\\,$M_\\odot$, $R$\\,=\\,$2.5\\pm0.2$\\,$R_\\odot$ from the comparison of the stellar parameters with theoretical evolutionary tracks. \\citet{KB06} have included 10\\,Aql in their study of the evolutionary state of magnetic chemically peculiar stars. Adopting a photometric temperature \\teff\\,=\\,7900\\,K, they found $M$\\,=\\,$1.95\\pm0.04$\\,$M_\\odot$, $\\log L$\\,=\\,$1.32\\pm0.05$ and a fractional age of 64--76\\% of the main sequence lifetime. The abundance pattern of 10\\,Aql is characteristic of a cool Ap star. A notable spectral anomaly is the overabundance of the doubly ionized REEs, \\pr\\ and \\nd.  A longitudinal magnetic field of about 500\\,G was detected in 10\\,Aql by \\citet{B58}, \\citet{P70} and \\citet{RWA05}. Using magnetic Zeeman splitting in unpolarized spectra \\citet{KLR02} provided the first measurement of the mean magnetic field modulus in 10\\,Aql, \\bs\\,=\\,$1.5\\pm0.1$\\,kG, which later was confirmed by Leone, Vacca \\& Stift (\\citeyear{LVS03}),  and he estimated the projected rotational velocity \\vs\\,=\\,2.0\\,\\kms. The sharpness of spectral lines and the absence of measurable variations of the mean longitudinal magnetic field indicate that 10\\,Aql has a long rotational period.  Recently, Ryabchikova, Kochukhov \\& Bagnulo (\\citeyear{RKB08}) have studied the vertical stratification and isotope anomaly of Ca in a number of Ap stars, including 10\\,Aql. These authors found a concentration of Ca in deeper atmospheric layers and invoked a vertical separation of the Ca isotopes to explain the shape of the Ca\\ii\\ infrared triplet lines. A comprehensive chemical stratification analysis of 10\\,Aql is currently underway by Nesvacil et al (in preparation), but already a preliminary result (Nesvacil, Weiss \\& Kochukhov \\citeyear{N08}) shows that Si, Ca, Cr, Fe, and Sr share a qualitatively similar vertical abundance distribution, characterized by a rapid increase of the element abundances below  $\\log\\tau_{5000}\\approx -1$. A re-determination of the atmospheric parameters taking into account individual chemical abundances and stratification supports the spectroscopic \\teff\\,$\\approx$\\,7600\\,K obtained by \\citet{RSH00}.  Photometric variations with a period close to 11.4\\,min and an amplitude below 0.5\\,mmag were discovered in 10\\,Aql by  \\citet{HK88}.  \\citet{HK90}  identified three pulsation modes with periods of 11.06--13.45\\,min, based on 26\\,h of high-speed photometry in the $B$-band. Despite an aliasing problem in their data, \\citet{HK90} were able to determine the large frequency separation, $\\Delta\\nu=50.6$\\,$\\mu$Hz. Belmonte, Martinez Roger \\& Roca Cortes (\\citeyear{BMR91}) collected 47\\,h of time-resolved photometry in the $J$-band, confirming oscillations in 10\\,Aql. Moreover, their study suggested that the amplitude of the infrared photometric variations significantly exceeds variations in the $B$-band, which is unusual for a roAp star (e.g. Martinez, Sekiguchi \\& Hashimoto \\citeyear{MSH94}; \\citealt{MWR96}). However, no oscillations were detected in the follow-up photometric monitoring conducted in the $H$-band (Belmonte, Kreidl \\& Martinez Roger \\citeyear{BKM92}).  Radial velocity variations with previously known oscillation periods were discovered in 10\\,Aql by \\citet{KLR02}. Over two nights they observed the star during a total of 8\\,h in a short wavelength region centred at $\\lambda$\\,6150\\,\\AA. Unlike the majority of other roAp stars, 10\\,Aql showed only a weak variability with an amplitude of $\\approx$\\,30\\,\\ms\\ in the strong Nd\\iii\\ 6145\\,\\AA\\ line, but exhibited amplitudes at the level of 80--130\\,\\ms\\ in the lines of first REE ions, Gd\\ii\\ and Eu\\ii. Aiming to extend the study of pulsational variation of 10\\,Aql, \\citet{HM05} analysed spectra covering 11.6\\,h of time-resolved observations during three nights. Somewhat surprisingly, despite a broad wavelength coverage of their data, \\citet{HM05} were able to confirm the presence of RV variations in only 5 individual spectral lines, two of which remained unidentified. The authors found the highest RV amplitude of 400\\,\\ms\\ in an  unidentified line at $\\lambda$~5471.42\\,\\AA\\ but could not detect pulsations in strong doubly ionized REE lines which show prominent variability in other roAp stars. In addition to its remarkably low-amplitude photometric pulsational variations, 10\\,Aql seems to exhibit an unusual spectroscopic pulsational behaviour, making this star perhaps an intermediate object between high-amplitude roAp stars and apparently constant non-pulsating magnetic Ap stars.  10\\,Aql was chosen as a target for an extensive observing campaign by the Canadian photometric space telescope MOST \\citep{W2003}. The star was observed as a primary target in June-July 2006, and nearly 120000 10\\,s exposures were collected during 31\\,d at a sampling interval of 20\\,s. The frequency analysis by \\citet{huber} revealed three definite periodicities and two candidate frequencies (see Table\\,\\ref{mostf}). Huber et al. showed that the two largest amplitude peaks reported by \\citet{HK90} are 1\\,d$^{-1}$ aliases of the intrinsic stellar frequencies $f_1$ and $f_2$. The third significant MOST frequency, $f_3=1.43$\\,mHz, has not been detected in previous photometric observations of 10\\,Aql. The lack of rotational modulation signal in the MOST data confirms the suspected slow rotation of the star. \\citet{huber} derived an improved value of the large frequency separation, $\\Delta\\nu=50.95$\\,$\\mu$Hz, and have compared the observed frequencies with the predictions of the theoretical models of non-adiabatic, magnetically distorted {\\it p-}modes \\citep{S05}.  \\begin{table}% \\caption{Pulsation frequencies of 10\\,Aql identified in the MOST photometric campaign \\citep{huber}. The numbers in brackets give the estimated error in units of the last significant digit. The two candidate frequencies, $f_i$ and $f_j$, are possibly present in the MOST data but could not be definitely confirmed. \\label{mostf} } \\begin{tabular}{lllc} \\hline id & $\\nu$ (mHz) & $P$ (min) & $A$ (mmag) \\\\ \\hline $f_1$ & 1.44786(3) & 11.5112(2) & 0.17(1) \\\\ $f_2$ & 1.39691(3) & 11.9311(3) & 0.15(1) \\\\ $f_3$ & 1.42709(4) & 11.6788(3) & 0.12(1) \\\\ $f_i$ & 1.3662(1)  & 12.199(1)  & 0.03(2) \\\\ $f_j$ & 1.4686(1)  & 11.349(1)  & 0.03(2) \\\\ \\hline \\end{tabular} \\end{table}  With the aim to improve the frequency analysis, study line profile variations and perform a tomographic analysis of the pulsations in the atmosphere of 10\\,Aql, we have organized a ground-based spectroscopic observing campaign simultaneously with the MOST observations. Using high-resolution spectrographs at 4- and 8-m telescopes,  we have collected a superb observational material with very high signal-to-noise ratio, wide wavelength coverage and a high spectral resolving power. This combination of the data quality and time coverage has been never achieved before for any roAp star.  Preliminary results of our pulsation study of 10\\,Aql were reported by \\citet{SKR07}. In that paper we outlined the frequency analysis procedure and discussed the main pulsation properties of 10\\,Aql, such as the presence of a node-like behaviour in the bisector variations, phases and amplitudes of radial velocity variation for different REE ions. Subsequently, Elkin, Kurtz \\& Mathys (\\citeyear{EKM08}) published an analysis of a short spectroscopic UVES time series of 10\\,Aql. These data, covering only 2 hours, were analysed without taking the multiperiodic character into account. The pulsation frequency of 1.428\\,mHz adopted by Elkin et al. coincides with neither of the two highest-amplitude pulsation modes in our data, but is close to the frequency $f_3$ identifyed by \\citet{huber} and by our spectroscopic analysis (see below).  The present paper is organized as follows. In Sect.\\,\\ref{observ} we describe spectroscopic observations and data reduction. Section\\,\\ref{RV} describes methodology of the line identification and radial velocity measurements. The frequency analysis is outlined in Sect.\\,\\ref{freq}. We establish a relation between photometric and spectroscopic variations of 10\\,Aql in Sect.\\,\\ref{phase}. Line profile variations, indicating a pulsation node in the stellar atmosphere, and the discovery of a remarkable asymmetry in the \\nd\\ lines are presented in Sect.\\,\\ref{puls}. The paper ends with conclusions and discussion in Sect.\\,\\ref{discus}. ",
        "conclusions": "\\label{discus} \\subsection{Reassessment of asteroseismic models}  The additional spectroscopic frequencies could offer a new view on the large frequency separation $\\Delta \\nu$, a crucial factor for asteroseismology which is directly connected to the mean density in the star and describes the separation of consecutive radial overtones for high-order acoustic pulsation. Figure~\\ref{spacing} shows a schematic amplitude spectrum including all intrinsic and candidate frequencies of both MOST photometry and spectroscopy. To first view, there is no apparent equal spacing visible. If we consider the four frequencies with high S/N (solid lines) to estimate the spacing, the only reasonable solution corresponds to $f_{1}-f_{2}=51\\,\\mu$Hz, a value already noted by \\citet{huber}. Lower values (such as 20 or 30\\,$\\mu$Hz) although seem tempting, would however contradict the previously determined temperature and luminosity of 10\\,Aql (see Fig. 4 in \\citealt{huber}). Assuming $\\Delta\\nu=51\\,\\mu$Hz to be correct, Fig.\\,\\ref{spacing} shows the expected position of other radial orders for three different spherical degrees $\\ell$ (vertical dashed, dotted and dashed-dotted lines, respectively) based on the position of the {\\bf four highest} S/N peaks. Evidently, the agreement is not very convincing. Nevertheless, an average deviation of about 5\\,$\\mu$Hz (which is needed to align the observed values to the expected position) is about the order of the frequency shifts of consecutive radial overtones due to the magnetic field perturbation as predicted by theory (see, e.g., \\citealt{cunha2006}). It must be noted, however, that such suggestion can lead to (almost) any desired result and therefore have to be considered with extreme caution.  \\begin{figure}% \\figps{10_Aql_fig4.eps} \\caption{The observed secure (black lines) and tentative frequencies from the photometry (blue lines) and spectroscopy (red lines). The expected position of radial orders for three different spherical degrees $\\ell$ (vertical dashed, dotted and dashed-dotted lines, respectively) is calculated based on a large separation of $f_{1}-f_{2}=51\\,\\mu$Hz.} \\label{spacing} \\end{figure}  Generally, the frequencies we derived from spectroscopy do not contradict to those from MOST photometry. But spectroscopic data provide an additional information useful for modelling. According to \\citet{huber} the best model describing the pulsation features in 10\\,Aql derived from spectroscopy corresponds to a star of 1.95$M_{\\odot}$, \\teff\\,=\\,7730 K and normal solar composition with the boundary reflection layer at $\\log\\tau_{5000}\\sim-4$. The spectroscopic analysis reveals a pulsation node in the Nd\\ii\\ - Nd\\iii\\ line formation region. A Non-LTE analysis of these lines suggested that they are formed  in roAp stars in atmospheric layers above $\\log\\tau_{5000}=-4$ \\citep{MRR05}, thus indicating the position of the node. Indeed, a pulsation node is predicted at $\\log\\tau_{5000}=-4$ in a star of 1.95$M_{\\odot}$, \\teff\\,=\\,7730 K but assuming helium depletion (Saio, private communication), which would better correspond to the typical chemistry of a magnetic peculiar star. Therefore, combining simultaneous photometric and spectroscopic observations, we expect to improve the pulsation model of 10\\,Aql. This analysis will be presented in the next paper.  \\subsection{Interpretation of spectroscopic pulsation measurements}  Here, as in many other recent studies of pulsational RV variations of roAp stars, we interpret the measurements of amplitude and phase in terms of outward propagation of pulsation waves in a chemically stratified stellar atmosphere. Previously, this interpretation of spectroscopic observations of multiperiodic roAp stars was handicapped by short spectroscopic time series which did not allow to resolve frequency patterns of multiperiodic pulsators. Here we overcome this difficulty by combining time-resolved observations acquired during several nights, but with a poor duty cycle, with a continuous photometric monitoring by MOST at the same time.  This strategy allows us to resolve the vertical structure of the three principal modes in 10\\,Aql and to study the propagation of pulsation waves independently for all three frequencies.  We detect a radial node in the region sampled by the formation heights of Nd and Pr for all three frequencies. The maximum amplitudes are observed for the \\dy\\ and Tb\\iii\\ lines that are presumably formed below and above the region of the node. The existence of a node is supported by the bisector analysis of the \\nd\\ 5102~\\AA\\ line.  By attributing the differences in the RV curves of different REE ions to the vertical structure of oscillations we have implicitly assumed that different species sample the horizontal structure of pulsation modes in a similar manner. This might not be the case if the vertical distribution of pulsation amplitude and phase depends strongly on the stellar surface position  due to a spotty element distribution combined with an intrinsic dependency of the mode structures on the orientation and strength of the magnetic field. The importance of the latter effect, arising from a superposition of magnetic and acoustic components of pulsation waves, has been recently emphasized by \\citet{SC07}.  Due to a long rotation period we are unable to constrain the surface abundance distribution of 10\\,Aql in the same way as it could be done for rapidly rotating roAp stars \\citep{K06x}. For this reason it is not straightforward to distinguish between the vertical and horizontal structural effects. At the same time we see only small changes in the line profiles in a comparison of observations obtained in 2001 and 5 years later. Furthermore, the independent evidence for a node provided by phase-amplitude diagrams and bisector analyses of different REE ions suggests that the observed phase variation corresponds rather to a vertical structure and not to a different horizontal sampling of the stellar surface.  \\subsection{Comparison with 33\\,Lib} \\label{33Lib}  A comparison of the pulsation properties of 33\\,Lib and 10\\,Aql is of considerable interest because these are the only roAp stars with clear signatures of a radial node in the upper atmosphere. \\citet{ryab2007b} performed a pulsation analysis of 33\\,Lib using phase-amplitude diagrams, which enable a direct comparison of the vertical pulsation structure in two stars. The top panel of Fig.\\,\\ref{33Lib_ph-rv_all} displays the phase-amplitude diagram for 33\\,Lib, while bisector measurements of \\nd\\ lines are compared for both stars in the bottom panel of Fig.\\,\\ref{33Lib_ph-rv_all}.  A pulsation node is clearly present in the atmospheres of both stars but the phase jump has opposite sign. Taking into account that the REE elements are concentrated in the upper atmospheric layers close to or even above the H$\\alpha$ line core formation zone, we can conclude from Figs.\\,\\ref{rv-ph} and \\ref{33Lib_ph-rv_all}  that the phase changes by 0.5 between the formation zones of Dy\\iii\\ and Tb\\iii\\ in the upper atmosphere of 10\\,Aql, while it jumps by $-0.5$ between Nd\\ii\\ and Nd\\iii\\ in 33\\,Lib.  Interestingly, the bisector measurements give us just the opposite phase variation with depth if one treats the line profile as being produced in a normal stellar atmosphere where the line core is formed higher than the line wings. As we show in the lower panel of Fig.~\\ref{33Lib_ph-rv_all}, the phase decreases from the wing to the core in 10\\,Aql, and increases in 33\\,Lib. At present we do not have a definite interpretation of this discrepant behaviour. One can suspect that the the broad wings of strong REE lines actually originate in the outer atmosphere and hence show pulsational characteristics of these layers.  Although the atmospheric parameters of 10\\,Aql and 33\\,Lib are similar, 33\\,Lib has a much stronger magnetic field, \\bs\\,=\\,5.0~kG, and it has a higher overabundance of REEs (see \\citealt{RNW04}). In addition, it has a shorter main pulsation period, $P\\,=\\,8.27$~min, with the first harmonic exhibiting the highest amplitude close to the position of a radial node. The negative phase jump and the shape of the amplitude-phase diagram (Fig.~\\ref{33Lib_ph-rv_all}, upper panel) may be interpreted as a superposition of standing and running pulsation waves mimicing an inwardly propagating wave, as discussed by \\citet{SC07}. A detailed study of chemical stratification and atmospheric structure of both stars is required for a secure interpretation of pulsation results and subsequent theoretical modelling.  \\begin{figure}% \\figps{33Lib_ph-rv_all.eps}\\\\ \\vspace*{-0.1cm}  \\figps{10Aql-33Lib_bis1.eps} \\caption{Amplitude-phase diagram for 33\\,Lib (upper panel) and pulsation phases of bisector measurements at various normalized flux levels of \\nd\\ lines in 10\\,Aql and 33\\,Lib (lower panel). The phases are given as fractional pulsation period. For demonstration purpose the phases of 33\\,Lib (lower panel) are shifted by +0.6.} \\label{33Lib_ph-rv_all} \\end{figure}  \\subsection{Comparison with previous spectroscopic pulsation studies of 10\\,Aql} \\label{comp}  Non-radial pulsations in 10\\,Aql were investigated in four different spectroscopic observing campaigns,  including the present one. \\citet{KLR02} found that RV amplitudes in Eu\\ii\\ and Gd\\ii\\ lines exceed those determined for Nd\\iii\\ line, which was not typical for a roAp star. Additionally, \\citet{KLR02} detected a change of RV amplitude between two consecutive nights of their observations. These characteristics of pulsational behaviour of 10\\,Aql are confirmed in the present more extensive analysis.  \\citet{HM05} observed the star during 3 nights, clearly detecting pulsational variability in 5 lines. They found the highest pulsation amplitudes of a few hundred \\ms\\ in the three spectral lines of $\\lambda\\lambda$~5373, 5471, 5730\\,\\AA,  which could not be identified. We also find significant amplitudes in these spectral features and tentatively identify these  lines as Dy\\iii, which allowed us to investigate variation of these lines in context of pulsation wave propagation through a stellar atmosphere with chemical stratification.  A recent study by \\citet{EKM08} presented measurements of the pulsation amplitudes and phases based on UVES spectroscopic time-series over two hours, which we included in our present analysis allowing for an increase of our spectroscopic data by 10\\%. \\citet{EKM08} assumed that 10\\,Aql pulsates with a main frequency of 1.428~mHz, which is, in fact, close to one of the three main frequencies unambiguously detected in our spectroscopy and in the MOST photometry,  but with the lowest amplitude. The second frequency detected by \\citet{EKM08} at 1.309~mHz is based on variations of Tb\\iii\\ lines, and is not seen in our more voluminous data. It probably is an artifact caused by, e.g., the limited time span of their spectroscopic observations of this multiperiodic pulsator.  Our study is the first to resolve the frequency spectrum of 10\\,Aql with RV measurements. We find, in agreement with precise space photometry, that pulsations of 10\\,Aql are dominated by \\textit{three} modes of comparable amplitude.  Based on simulteneous spectroscopic and photometric observations, we were able to infer pulsational characteristics of individual modes as a function of height in the atmosphere of 10\\,Aql. We also obtained phase shifts between luminosity and RV variations providing constraints for modelling of roAp pulsation.  Since our time-series analysis is based on a set of three frequencies, the RV amplitudes that we report for individual lines cannot be directly compared with the results of previous studies which assumed monoperiodic variation of the star. However, a good agreement with the results by \\citet{EKM08} is found if we compare relative amplitudes and phase differences of the spectral features in common. The two studies obtained qualitatively similar phase-amplitude diagrams, with the exception of discordant behaviour of weak Er\\iii~5903~\\AA\\ line studied by \\citet{EKM08}. The pulsation characteristics of this line deviate from the generally smooth phase-amplitude behaviour of other REE ions. We omitted this line from our pulsation analysis because it is located in the spectral region full of telluric lines and is blended by 3 atmospheric lines of different intensity. This blending is the most likely the reason for the deviating behaviour of the spectral line Er\\iii~5903~\\AA.  \\subsection{Amplitude modulation in spectroscopy} \\label{concl}  Time-resolved spectroscopic observations of roAp stars seemed to indicate that the pulsation amplitudes of REE lines can be modulated on a time scale of a few hours. This was first noted by \\citet{KR01b} for $\\alpha$\\,Cir and later observed by \\citet{KEM06} for a few other roAp stars. Since in many cases this modulation cannot be linked to known photometric pulsation frequencies, \\citet{KEM06} speculated that the amplitude modulation in spectroscopy suggests the discovery of a new type of pulsational behaviour in the upper atmospheres of roAp stars. However, using relatively short and typically 2-hour long time series resulting in only 100--150 individual measurements, this discovery is doubtful, because such short observational data sets do not allow to resolve the frequency spectrum of multiperiodic roAp stars. In this context it comes not as a surprise that for the only two roAp stars, HD\\,24712 \\citep{ryab2007a} and 10\\,Aql (this paper), for which an extensive spectroscopic (and photometric) monitoring over many nights was performed, no unexplained spectroscopic amplitude modulation could be found. All amplitude modulations in 10\\,Aql can be explained by beating effects of close frequencies. Hence, we recommend extreme caution in the interpretation of amplitude changes seen in short data sets."
    },
    "0806/0806.0952_arXiv.txt": {
        "abstract": "We study the models with radiative neutrino mass generation and explore the relation between the neutrino masses and dark energy. In these models, the pseudo-Nambu-Goldston bosons (pNGBs) arise at two-loop level via the Majorana neutrino masses. In particular, we demonstrate that the potential energy of the pNGB can be the dark energy potential and the observed value of the equation of state (EoS) parameter of the universe, $i.e.$, $w\\simeq -1$, can be realized. ",
        "introduction": "In the standard model (SM) of particle physics, neutrinos are massless. However, various experimental searches indicate that neutrinos have tiny masses ($\\leq O(10^{-2})~\\mathrm{eV}$) \\cite{PDG2006}. It is a challenging and important problem to explain the origin of the small neutrino masses. Various mechanisms could generate neutrino masses~\\cite{Mohapatra:1998rq}, in which the one with radiative neutrino mass generation without right-handed neutrinos by extending the Higgs sector \\cite{RadNu1,RadNu2,RadNu3} is particularly interesting because the neutrino masses are naturally small. It is clear that without the right-handed states the active neutrinos can only have Majorana masses.  On the other hand, recent cosmological observations have confirmed that not only  there existed the inflationary stage in the early universe, but also  at present the expansion of the universe is accelerating~\\cite{WMAP1, SN1,Frieman:2008sn}. Although various scenarios for the late-time acceleration in the expansion of the universe have been proposed, the cosmic acceleration mechanism is still not well understood \\cite{Peebles:2002gy, Padmanabhan:2002ji, Copeland:2006wr, Durrer:2007re, NO-rev}.  In the framework of general relativity, the current accelerating universe is due to the so-called dark energy (or cosmological constant) with its density at the present time is only about $(10^{-3}\\mathrm{eV})^4$, which is much smaller than any known energy scale in particle physics except the neutrino masses. It is interesting to note that the energy scale of dark energy coincides with the neutrino masses as discussed in Ref.~\\cite{Hill:1988vm}.   Recently, it has been suggested~\\cite{Fardon:2003eh,XMZhang,Peccei:2004sz, EV} that the neutrino masses vary as a function of a scalar field, called the ``acceleron'', which drives the universe to its present accelerating phase. Such neutrinos are referred as mass varying neutrinos (MaVaNs). The effects of the MaVaNs on the anisotropy of the cosmic microwave background (CMB) radiation and large scale structure (LSS) have been studied in Ref. \\cite{CE}. Several models with the generation of the MaVaNs through the see-saw mechanism with right-handed neutrinos to account for the baryon asymmetry in the universe have been proposed in Refs.~\\cite{Hill:2006hj, Gu:2007ps}. In these studies, one has to break a global symmetry spontaneously to get a Nambu-Goldston boson (NGB) and introduce a soft symmetry breaking term so that the NGB receives a soft mass via a loop diagram and becomes a pseudo-Nambu-Goldston boson (pNGB). This pNGB corresponds to the acceleron field in the scenario proposed in Ref.~\\cite{Fardon:2003eh}. Some models to explain neutrino masses and dark energy at the TeV scale have been explored in Refs.~\\cite{Gu:2007gy,Bhatt:2007ah}. Moreover, the Majorana neutrino superfluidity and the stability of the neutrino dark energy have been  discussed in Ref.~\\cite{Bhatt:2008hr}.  In this paper, we consider the generation of the small Majorana neutrino masses through the radiative mechanism  without right-handed neutrinos in the framework of the extended Babu-Zee model~\\cite{RadNu2}. Here, we do not introduce a soft breaking term ``by hand'', but induce one from loop diagrams. In other words, we break global symmetries spontaneously in the first place, and then, via loop diagrams, introduce a soft (original) symmetry breaking term followed by a mass term for the pNGB. This pNGB also plays a role of the acceleron field. As a result, we show that the small neutrino masses depend on the pNGB and we argue that the potential energy of the pNGB can be the potential of dark energy. Furthermore, we demonstrate that the observed value of the equation of state (EoS) parameter from the Wilkinson Microwave Anisotropy Probe (WMAP) data on the anisotropy of the CMB radiation can be realized by following the discussion in Refs.~\\cite{Fardon:2003eh,XMZhang,Peccei:2004sz}. ",
        "conclusions": "In summary, we have considered the generation of the small neutrino mass through the radiative mechanism in the extended Babu-Zee models. We have shown that the generated small neutrino masses depend on a pNGB, which can play a role of the acceleron field and the potential energy of the pNGB can be the dark energy potential. In particular, we have demonstrated that the observed value of the EoS parameter from WMAP can be realized."
    },
    "0806/0806.2485_arXiv.txt": {
        "abstract": "{} {Type II AGNs with polarimetric broad emission line provided strong evidence for the orientation-based unified model for AGNs. We want to investigate whether the polarimetric broad emission line in type II AGNs can be used to calculate their central supermassive black hole (SMBH) masses, like that for type I AGNs.} {We collected 12 type II AGNs with polarimetric broad emission line width from the literatures, and calculated their central black hole masses from the polarimetric broad line width and the isotropic \\oiii\\ luminosity. We also calculate the mass from stellar velocity dispersion, $\\sigma_*$, with the $\\mbh-\\sigma_*$ relation.} {We find that: (1) the black hole masses derived from the polarimetric broad line width is averagely larger than that from the $\\mbh- \\sigma_*$ relation by about 0.6 dex, (2) If these type II AGNs follow $\\mbh-\\sigma_*$ relation, we find that the random velocity can't not be omitted and is comparable with the BLRs Keplerian velocity. It is consistent with the scenery of large outflow from the accretion disk suggested by Yong et al. } {} ",
        "introduction": "The standard paradigm for active galactic nuclei (AGNs) posits an accretion disk surrounding a central supermassive black hole (SMBH), along with other components, such as the broad-line regions (BLRs), narrow-line regions (NLRs), jet, and torus (e.g., Antonoucci 1993). The black hole mass ($\\mbh$) is an important parameter for us to understand the nuclear energy mechanics, the formation and evolution of SMBH and the galaxies (e.g., Rees 1984; Tremaine et al. 2002).  In the past decade, one of the most progresses in the study of AGNs is that the masses of SMBHs can be calculated by using the width of the broad emission lines from BLRs (e.g., $\\hb$, $\\ha$, $\\mgii$, $\\civ$) by the reverberation mapping method and several corresponding empirical relations (e.g., Kaspi et al.2000, Bian \\& Zhao 2004, Greene \\& Ho 2006). In the orientation-based unified model for AGNs, the distinction between type I AGNs and type II AGNs depends upon whether the central engine and BLRs are viewed directly (type I) or are obscured by the circumnuclear torus (type II). Because of the absence of broad emission lines in the spectrum of type II AGNs, above methods for the mass calculation are only applicable to type I AGNs. The SMBH mass in the center of type II AGNs generally may be estimated by the $\\mbh- \\sigma_*$ relation (e.g., Kauffmann et al. 2003; Bian \\& Gu 2007).  With spectro-polarimetric observation, some type II AGNs show hidden BLRs (HBLRs) and some do not (e.g., Antonucci $\\&$ Miller 1985; Tran 1995). It is still not clear what kind of physical process is related to the presence of HBLRs in type II AGNs (e.g., Bian \\& Gu 2007). In this paper, we calculate the SMBH mass in type II AGNs with HBLRs by using broad emission lines in their polarimetric spectrum as well as the $\\mbh- \\sigma_*$ relation. In $\\S$ 2, we briefly introduce our sample. $\\S$ 3 introduces the methods to calculate the SMBH masses. $\\S$ 4 is the data analysis. Our results and discussions are given in $\\S$ 5. The last section is our conclusion. All of the cosmological calculations in this paper assume $H_{\\rm 0}$=70$\\kms$ $Mpc^{-1}$, $\\Omega_{M}$=0.3, $\\Omega_{\\wedge }$=0.7. ",
        "conclusions": "The central SMBHs masses for 12 type II AGNs with polarimetric broad lines are estimated from the broad line FWHM and the \\oiii luminosity, as well as the $\\mbh-\\sigma_*$ relation. We find that: (1)the SMBH masses derived from the width of polarimetric broad emission lines is averagely larger than that from the $\\mbh-\\sigma_*$ relation by about 0.6 dex, (2)If these type II AGNs follow $\\mbh-\\sigma_*$ relation, our result suggested that the random velocity can't not be omitted and is comparable with the BLRs Keplerian velocity. It is consistent with the scenery of larger outflow from the accretion disk suggested by Yong et al. (2007).   {\\noindent \\rm {Acknowledgements}. We thank the anonymous referee for helpful suggestions. This work has been supported by the NSFC (No. 10733010), the Science-Technology Key Foundation from Education Department of P. R. China (No. 206053), and the China Postdoctoral Science Foundation (No. 20060400502).}"
    },
    "0806/0806.4740_arXiv.txt": {
        "abstract": "{} % {We study the formation of water and methanol in the dense cloud conditions to find the dependence of its production rate on the binding energies, reaction mechanisms, temperatures, and grain site number. We wish to find the effective grain surface area available for chemical reaction and the effective recombination timescales as functions of grain and gas parameters. } {We used a Monte Carlo simulation to follow the chemical processes occurring on the grain surface. We carried out the simulations on the Olivine grains of different sizes, temperatures, gas phase abundances and different reaction mechanisms. We consider H, O, and CO as the accreting species from the gas phase and allow ten chemical reactions among them on the grains.} {We find that the formation rate of various molecules is strongly dependent on the binding energies. When the binding energies are high, it is very difficult to produce significant amounts of the molecular species. Instead, the grain is found to be full of atomic species. The production rates are found to depend on the number density in the gas phase. When the density is high, the production of various molecules on the grains is small as grain sites are quickly filled up by atomic species. If both the Eley-Rideal and Langmuir-Hinselwood mechanisms are considered, then the production rates are at this maximum and the grains are filled up relatively faster. Thus, if allowed, the Eley-Rideal mechanism can also play a major role and more so when the grain is full of immobile species. We show that the concept of the effective grain surface area, which we introduced in our earlier work, plays a significant role in grain chemistry.} {We compute the abundance of water and methanol and show that the results strongly depend on the density and composition in the gas phase, as well as various grain parameters. In the rate equation, it is generally assumed that the recombination efficiencies are independent of the grain parameters, and the surface coverage. Presently, our computed parameter $\\alpha$ for each product is found to depend on the accretion rate, the grain parameters and the surface coverage of the grain. We compare our results obtained from the rate equation and the one from the effective rate equation, which includes $\\alpha$. A comparison  of our results with the observed abundance shows very good agreement.} ",
        "introduction": "Formation of complex organic molecules in the interstellar medium (ISM) is an active subject of research. Out of a host of organic and inorganic molecules that have been observed, water and methanol are most certainly two very important organic species found both in the gas and solid phases of ISM. The abundance of these species in the various regions of ISM is also different. It was thought for a long time that the water is one of the possible reservoirs of elemental oxygen in the gas phase, but recent Submillimeter Wave Astronomy Satellite (SWAS) observation found surprisingly a low abundance of water. Snell {\\it {et al.}} (2000) find that the abundance of water relative to H$_2$ in Orion and M17 cloud is between $10^{-10}$ to $8 \\times 10^{-10}$. The water abundance in the hot cores range from  $10^{-6}$ to $10^{-4}$ (van Dishoeck \\& Helmich 1996, Helmich {\\it {et al.}} 1996; Boogert \\& Ehrenfreund, 2004). The abundance of water  on grains with respect to the total H column density is typically $10^{-4}$ and is the most abundant component (Tielens et al. 1991). Similarly, interstellar methanol has three types of abundance profile: flat profiles at CH$_3$OH/H$_2$ $\\sim 10^{-9}$ for the coldest sources, profiles with a jump in its abundance from $\\sim 10^{-9}$ to $\\sim 10^{-7}$ for the warmer sources, and flat profiles at a $\\sim$ few $10^{-8}$ for the hot cores (van der Tak {\\it {et al.}} 2000). On the grain surface, the methanol abundance varies from  $5 \\%$ to $30 \\%$ with respect to H$_2$O. In some sources, such as SgrA and Elias 16, the abundance is even less (Gibb {\\it at al.}, 2000). The observed abundance for methanol along the line of sight towards high mass and low-mass proto-stars is between $0.2 - 2 \\times 10^{-5}$ (Gibb et al. 2004; Pontoppidan et al. 2003, 2004). Much higher methanol abundances are found to be associated with the outflows in the regions of low mass star formation, L1157-MM and NGC1333-IRAS2 $2 \\times 10^{-5}$ and $2 \\times 10^{-6}$ (Bachiller \\& Perez Gutierrez, 1997; Bachiller et al. 1998). In other words, either on the grain surfaces or in the hotter region, abundances of theses species are high. This correlation suggests that these species perhaps originate from grains and their productions in the gas phase are inadequate. Therefore, the understanding of the formation of water and methanol on grain surfaces is of primary importance.  The grain surface reactions were first introduced to explain the formation of molecular hydrogen (Hollenbach \\& Salpeter, 1970). Since then it has been used very extensively by several authors (Watson \\& Salpeter 1972ab; Allen \\& Robinson 1975, 1976, 1977; Tielens \\& Hagen 1982; Hasegawa \\& Herbst 1992; Charnley 2001; Stantcheva {\\it et al.} 2002, Green {\\it et al. 2001}; Biham {\\it et al.} 2001; Stantcheva {\\it et al.} 2002). These studies mainly belong to two different categories, the deterministic approach and the stochastic approach. In the deterministic approach, one can completely determine the time evolution of the system once the initial conditions are known. The rate equation method belongs to this category. This method is very extensively used by several authors to study the grain surface chemistry (Hasegawa \\& Herbst 1992; Roberts {\\it et al.} 2002; Acharyya et al. 2005). However, this method is only applicable when there are large numbers of reactants on the grain surface. Given that the interstellar medium is very dilute, very often this criteria is not fulfilled and this method cannot be applied. But this method is computationally faster and can very easily be coupled with the gas phase reactions. In the stochastic approach, fluctuations in the surface abundance due to the statistical nature of the grain is preserved. The Monte Carlo method and the Master equation methods belong to this category. Both these methods are used by several authors (Charnley 2001; Stantcheva {\\it et al.} 2002, Green {\\it et al. 2001}; Biham {\\it et al.} 2001; Stantcheva {\\it et al.} 2002 ). Its major disadvantage is that it takes enormous computational time. Coupling of the Monte Carlo method (for grain surface reactions) and rate equation method (for gas phase reactions) is extremely difficult. Although the Master equation method can be coupled with the rate equations, but it is disadvantageous because one has to solve a large number (ideally infinite) of reactions.  \\begin {figure} \\vskip 0.6cm \\centering{ \\includegraphics[width=6cm]{8422fig1.eps}} \\caption{Cartoon diagram to show the different reaction schemes (C1-C4) that are considered in our calculation. Black and shaded circles are representatives of reactive species and the clear circles are for nonreactive species. A circle around two such species indicates that the new species is created due to Eley-Rideal scheme.} \\label{fig-1} \\end {figure} Recently, Chang {\\it et al. } (2005) have argued that the stochastic methods used so far can also lead to error because the rate of reaction is determined by the rate of hopping (or tunneling) of a hydrogen atom from one site to the nearest neighboring site multiplied by the probability of finding a reactant partner in this site. This is also an average treatment since, on any given grain, the reactant partner is unlikely to lie in the nearest-neighboring site. They used a continuous random work technique to study the formation of molecular hydrogen. Chakrabarti {\\it et al. (2006ab)} used a similar method that keeps track of each individual reactant and their movements and calculated the effective grain surface area involved in the formation of molecular hydrogen in the interstellar clouds. Chakrabarti {\\it {et al.}} (2006, hereafter referred as Paper-I) illustrated that the formation rate per unit grain site itself strongly depends on the nature and size (i.e., no. of sites) of the grains for a given gas phase condition (abundance, temperature, etc.). In Paper-I, this was demonstrated only for H forming H$_2$ molecules.  In the present paper, we carry out a similar analysis where we consider the accretion of H, O, and CO onto the grain surface and show that the formation rates of water and methanol are indeed dependent on the intrinsic and extrinsic parameters of the system. The plan of this paper is the following. In the next section, we discuss various mechanisms of reactions on the surface. In Section 3, we discuss the procedure of our computation. In Section 4, we describe various models used for our simulations. In Section 5, we present our results. Finally, in Section 6, we draw our conclusions. \\begin {figure} \\vskip 0.6cm \\centering{ \\includegraphics[width=7cm]{8422fig2.eps}} \\caption{Evolutions of the number of a few selected species on a grain having $10^4$ sites are shown with respect to the surface coverage of a mono-layer (along X-axis), for various sets of binding energies  (Models 1-3) and for different mechanisms (C1-C4). These simulations were carried out for high abundances of the accreting species (Table 1). Different styles of the curves are marked with the species names on the right. Note that methanol and CO$_2$ are absent in the Model 1 simulation where the binding energy was very high.} \\label{fig-2} \\end {figure} ",
        "conclusions": "In this paper, we studied the formation of water, methanol, and other related species on a grain surface using a Monte-Carlo method. We used three different sets of binding energies, two different gas-phase abundances and considered both the ER and LH mechanisms. We considered all the four possible ways (C1-C4) of the landing of a species on the occupied site of a grain. Besides the formation of these species, we also calculated two parameters $\\gamma$ and $\\alpha$ that represent the recombination efficiencies to show that they are indeed dependent on the number of sites on the grain and the populations of various species on a grain surface.  We found that the formation of various molecules is dependent on the binding energies. We find that, when the higher binding energies are used, it is very difficult to produce a significant amount of the molecular species, instead, the grain is found to be full of atomic species. The formation of these species is also dependent on the gas phase density. We found that, for the high density case, the production of various molecules is small and the grains are filled up relatively quickly by atomic species. We found that, if both the ER and LH mechanisms are considered, then the production is always high and the grain is filled up very quickly. As expected, we find that, when the grain is more or less empty, the LH scheme is most important. The ER scheme starts dominating more and more as the grains are filled up. We have verified that our results from the Monte-Carlo simulation match with that from the effective rate equation.  It is important to see if we can put a constraint on the model parameters from the observational results. We already mentioned in the Introduction that the abundance of water relative to H$_2$ in cold clouds is between $10^{-10}$ and $8 \\times 10^{-10}$ and between $10^{-6}$ to $10^{-4}$ in hot cores. The abundance of methanol with respect to H$_2$ is $ \\sim 10^{-9}$ for coldest clouds and between $\\sim 10^{-9}$ to $10^{-7}$ for warmer clouds and $\\sim$ a few $\\times 10^{-8}$ in hot cores. In the grain surfaces, the solid state methanol abundance should be at the most $30\\%$ with respect to H$_2$O. If we glance at Table 5 and assume that the surface coverage on multiple layered grains is similar to that on a mono-layer, then it is clear that the case of Model 1 and high abundance cases for Models 2 and 3 are not very relevant for the production of methanol. Indeed, the low abundance cases with Models 2 and 3 produce a similar solid state methanol with respect to H$_2$O ($\\sim 43\\%$ - $\\sim 57\\%$). Corresponding high abundance models produce only $4\\% - 6\\%$. Thus the observed abundance of $\\sim 30\\%$ must be from some intermediate region of the cloud. In Fig. 15 we show the ratio of methanol and water as a function of the number density of the gas phase, i.e., accretion rate. We clearly see that to obtain the observed ratio the number density should be around $2.6-8 \\times 10^4$cm$^{-3}$.  If the temperature is high enough, some methanol is released in the gas phase. If $n_l$ layers were formed on the grains and were all assumed to evaporate to the gas phase, the gas phase abundances of methanol would be given by $$ R = \\frac{n_g n_l s S}{n_{H_2}} =3 \\times 10^{-6} f_{-12} n_{10} s_{i,0.3} S_6 . $$ Here, $n_g \\sim n_{H_2} 1.33 \\times 10^{-12}=n_{H_2} f_{-12}$ the number density of grains, $n_{H_{2}}$ is the number density of H$_2$ in the gas phase, $s_{i,0.3}$ the fractional surface coverage of $i$th species (e.g., methanol or water) in units of  $0.3$, $S_6$ the number of sites in the grain in units of $10^6$, and $n_{10}$ in units of $10$ mono-layers. Since time taken for production of a mono-layer is $ \\sim 0.1$ Myr, by the time $10$ mono-layers are produced, the grains are deeply inside in higher abundance region. Here the mono-layer production time is lower but the production efficiency of formation of methanol is lower as well. Thus $R$ computed for methanol above, for $n_l=1$ to $10$ varies from $\\sim 10^{-7}$ ($s_{i,0.3} \\sim 0.03$) to $3 \\times 10^{-5}$ ($s_{i,0.3} \\sim 1$), generally agrees with the observational results. If, for instance, a fraction of grains are sublimated, say, $f_{-12} \\sim 0.1$, the abundance could be even lower. For water, $R\\sim 10^{-6}$ to $6 \\times 10^{-5}$, where we have put $n_{10}=1$ to $10$, $s_{i,0.3}=0.33$ for high abundance, and $s_{i,0.3}=2$ for low abundance. This is also in the observed range.  One of our important findings is that the parameter $\\alpha(t)$ strongly depends on the population of the reactant species on the grain surface. This deviates significantly from unity. This seems to be a very important parameter, because in the usual rate equation we assume that $\\alpha(t)$ is always $1$ (a consequence of the assumption that the recombination is totally a random walk process). This is an overestimation. In Paper-I we also computed this parameter for the H$_2$ molecule. We defined another parameter $\\gamma$ called the `catalytic capacity' and found that it goes down with the increased surface population. This shows that the rate of production indeed increases as the grain is filling up. We also found that the behaviors of $\\alpha(t)$ and $\\gamma$ strongly depends upon the grain temperature.  In our present calculation we restricted ourselves in two ways: (a) We considered only dense clouds where the accretion is such that most of the hydrogen is already in the molecular form in the gas phase. Thus accreting gas composition produces primarily water and methanol as is truly the case for dense clouds. In diffused clouds, on the other hand, the  composition is such that mostly molecular hydrogen is formed and are desorbed into the gas phase. Such a work was presented in Paper-I. (b) We restricted ourselves up to the formation of a mono-layer due to the fact that the complexity of the problem rises with layer number and the unavailability of sufficient data (e.g., the binding energies at different surfaces). In reality, multi-layers would be produced and each layer is expected to have a different abundance due to the freezing out effect. For instance, in the first layers would be dominated by H$_2$O and methanol, etc. In the later stage, accretion species will be dominated by CO, O$_2$, N$_2$, etc. Detailed results are in progress and will be reported elsewhere.  The paper has been greatly improved due to the helpful comments of the anonymous referee who is acknowledged. The work of AD was supported by a RESPOND grant from ISRO."
    },
    "0806/0806.4606_arXiv.txt": {
        "abstract": "One of the most exciting results of the Spitzer era has been the ability to construct longitudinal brightness maps from the infrared phase variations of hot Jupiters. We presented the first such map in \\cite[Knutson et al. (2007)]{Knutson_2007}, described the mapping theory and some important consequences in \\cite[Cowan \\& Agol (2008)]{Cowan_2007} and presented the first multi waveband map in \\cite[Knutson et al. (2008)]{Knutson_2008}. In these proceedings, we begin by putting these maps in historical context, then briefly describe the mapping formalism. We then summarize the differences between the complementary N-Slice and Sinusoidal models and end with some of the more important and surprising lessons to be learned from a careful analytic study of the mapping problem. ",
        "introduction": "Observations of secondary eclipses in exoplanetary systems, starting with HD~209458b \\cite[(Deming et al. 2005)]{Deming_2005} and TrES-1b \\cite[(Charbonneau et al. 2005)]{Charbonneau_2005}, made it possible to estimate the integrated day-side brightness of transiting exoplanets. Constraining the \\emph{global} brightness map of exoplanets, on the other hand, requires observations at various orbital phases, involving more sophisticated calibration of observations, much longer observing campaigns, or both. The first measurements of thermal phase curves for exoplanet systems were reported by \\cite[Harrington et al. (2006)]{Harrington_2006}, which reported a large phase function for $\\upsilon$ Andromeda~b, and \\cite[Cowan, Agol \\& Charbonneau (2007)]{Cowan_2007}, which detected a phase function for HD~179949b, and obtained useful upper limits for HD~209458b and 51~Peg~b. These results proved valuable in constraining the day-night brightness contrast ---and hence the energy recirculation efficiency--- of those planets and indicated that hot Jupiters represent a heterogeneous group. Those first two studies, however, had very incomplete phase coverage (5 epochs for the  \\cite[Harrington et al. 2006]{Harrington_2006} campaign, and 8 epochs for each of the \\cite[Cowan et al. 2007]{Cowan_2007} campaigns). Furthermore, three of the four observed planets were not in transiting systems, and the one transiting system (HD~209458) was deliberately observed outside of transit or secondary eclipse.  The $33$~hours of continuous monitoring of HD~189733b presented in \\cite[Knutson et al. (2007)]{Knutson_2007} differs in three important ways from those first detections of phase variations: 1) The observed system exhibits transits, so the planet's orbital inclination with respect to the celestial plane is known. 2) A secondary eclipse of the planet was observed during the course of the observations, making it possible to quantify not just the relative but the \\emph{absolute} flux of the planet as a function of orbital phase. 3) The continuous observing campaign, the system's relative proximity to the Earth, its favorable contrast ratio, and ingenious corrections for detector systematics conspired to produce the highest S/N light curve of its kind ever measured. Although the observations spanned little more than half an orbit of HD~189733b, the unprecedented quality of the light curve enabled us not only to measure the planet's day/night contrast, but also to generate the first ever brightness map of an extrasolar planet.  \\begin{figure}[htb] \\includegraphics[width=0.5 \\textwidth]{f1.eps} \\includegraphics[width=0.5 \\textwidth]{f2.eps} \\caption{The solid line in the left panel shows the phase function response to a $\\delta$-function in brightness (AKA ``the kernel''). The dashed and and dot-dashed lines represent the contributions from a single slice in 4-slice and 2-slice models, respectively. In the right panel, the solid band shows the 1-$\\sigma$ range for a sample sinusoidal map, while the dotted line shows the equivalent N-slice map and associated error bars. Both maps have 5 parameters and produce indistinguishable light curves.} \\label{slice_kernels} \\end{figure} ",
        "conclusions": ""
    },
    "0806/0806.1419_arXiv.txt": {
        "abstract": "We present a detailed optical study of the ultra-compact X-ray binary \\target. We have used 63\\,hrs of time-resolved optical photometry taken with three  different telescopes (IAC80, NOT and SPM) to search for optical modulations. The power spectra of each dataset reveals sinusoidal modulations with different periods, which are not always present.  The strongest modulation has a period of 51.3\\,mins, a semi-amplitude of 4.6\\,mmags, and is present in  the IAC80 data. The SPM and NOT data show periods of 42\\,mins and 64\\,mins respectively, but with much weaker amplitudes, 2.6\\,mags and 1.3\\,mmags respectively. These modulations arise from either X-ray irradiation of the inner face of the secondary star and/or a superhump modulation from the accretion disc, or quasi-periodic modulations in the accretion disc.  It is unclear whether these  periods/quasi-periodic modulations are related to the orbital period,  however, the strongest period of 51.3\\,mins is close to earlier tentative orbital periods. Further observations taken over a long base-line are encouraged. ",
        "introduction": "Low-mass X-ray binaries (LMXBs) are systems in which a low-mass companion star transfers material onto a neutron star or black hole. Most of the systems have orbital periods of between a few hours to days and contain ordinary hydrogen-rich donor stars. While these systems have  minimum orbital periods around 80 mins, systems with hydrogen-poor or  degenerate donor stars can, however,  systems with hydrogen-poor or degenerate donor stars can evolve to extremely small binary separations with orbital periods  as short as a few minutes \\citep{Nelson86}. Such systems are called ultra-compact X-ray binaries (UCXBs) and have a range in orbital periods from 11 to 50 minutes (see \\citealt{NJ06})  Finding UCXBs is difficult because measuring orbital periods (\\porb) is generally difficult in LMXBs.  The current known sample consists of 27 systems, 8 with known periods, 4 with tentative periods and 15 candidate systems \\citep{Zand07}. The identification of UCXBs is mostly done through measuring \\porb\\ via (1) timing of Doppler-delayed pulses if the accretor is a pulsar; (2)  measuring  periodic X-ray eclipses/dips if the binary inclination is high enough and (3) through the measurement of periodic modulations resulting  from X-ray heating of the inner side of the donor star or from the superhump phenomenon that is predicted for extreme mass ratio systems. There are also two indirect methods to identify UCXBs without measuring \\porb, which depend on the fact that in an UCXB the accretion disk are relatively small; (1) for the same X-ray flux, $M_V$ is about 4 mags fainter for UCXBs than for normal LMXBs and (2) a method which depends on the  critical accretion rate below which a system becomes transient; if the persistent accretion luminosity is 1\\% of Eddington \\citep{Zand07}   \\target\\ is a low luminosity X-ray binary. Thermonuclear type I X-ray bursts were observed by {\\it OSO-8} \\citep{Swank78} and the the compact object is a neutron star. Historically, \\target\\ has been associated with the brightest X-ray bursts.  Based on the detection of a bright burst with {\\it Watch}, and Eddington limit  arguments, \\citet{Brandt92} argue that the distance to \\target\\ is probably  $<$3\\,kpc. Several lines of evidence point to the conclusion that \\target is an UCXB. Based on the comparison of the enhanced neon to oxygen ratio to known  UCXBs, \\citet{Juett01} argue that \\target\\ is also a UCXB.  Further support is provided by  optical spectroscopy, which has revealed carbon and oxygen emission lines, but no evidence for hydrogen or helium  \\citep{Nelemans06}. The optical counterpart to \\target, V1055\\,Ori, is also intrinsically faint ($V$=18.5). The faintness of its persistent X-ray emission and nearby distance suggest a low accretion rate, which is consistent with an orbital period $<$1\\,hr \\citep{Deloye03}.  In this paper we present the results on a long term campaign to find a stable period in \\target, which would most likely represent the orbital period. ",
        "conclusions": "\\label{DISCUSSION}  \\citet{Nelemans04} obtained VLT optical spectra of \\target\\ and identified features of  relatively low ionisation states of carbon and oxygen. This clearly identifies \\target\\ as an UCXB and suggests that the donor in the system in a carbon-oxygen white dwarf. For \\target\\ and 4U\\,1626-67 there are clear indications that the discs are dominated by C and O. \\citet{Werner06} have also obtained VLT spectra and compared them with detailed NLTE models for spectra of UCXBs. Unfortunately, the NTLE models do not sufficiently agree with the observed spectra  for quantitative abundance analysis. Although simple LTE models seem to fit the data better, they also cannot be used for quantitative measurements because NLTE effects mainly due to X-ray irradiation, need to be taken into account.  We have found  several sinusoidal modulation  in the  optical lightcurve of \\target. These modulations most likely arise from either X-ray irradiation of the inner face of the secondary star and/or a superhump modulation from the accretion disc, or quasi-periodic oscillations in the accretion disc.   This is not surprising as UCXBs are known to show orbital modulations as well as QPOs. e.g. strong 15 minute optical/UV quasi-periodic oscillations were previously detected in the 42\\,min UCXB 4U\\,1626--67 \\citep{Chak01}, showing that photometric variability in an UCXB need not only occur near the orbital period.  \\citet{OBrien05} have reported a 50\\,min periodicity in their high-speed optical data taken with ULTRACAM.   Time-resolved VLT spectroscopy also shows some evidience, although marginal, for a 49\\,min periodicity in the  weak absorption line near 4960\\AA\\ \\citep{Nelemans06}. The strongest period we detect is at 51.3\\,mins and is present in the IAC80  data. This may reflect the superhump period, slightly longer than the orbital period. However,  only observations long enough to contain many modulation cycles can distinguish between a periodic and a quasi-periodic modulation and allow a secure measurement of the orbital period."
    },
    "0806/0806.1844_arXiv.txt": {
        "abstract": "We consider two different toy cosmological models based on two fields (one normal scalar and one phantom) realizing the same evolution of the Bang-to-Rip type. One of the fields (pseudoscalar) interacts with the magnetic field breaking the conformal invariance of the latter. The effects of the amplification of cosmic magnetic fields are studied and it is shown that the presence of such effects can discriminate between different cosmological models realizing the same global evolution of the universe. ",
        "introduction": "The discovery of the cosmic acceleration \\cite{cosmic} and the search for dark energy responsible for its origin \\cite{dark} have stimulated the study of different field models driving the cosmological evolution. Such a study usually is called the potential reconstruction \\cite{recon}, because the most typical examples of these models are those with a scalar field, whose potential should be found to provide a given dynamics of the universe. In the flat Friedmann models with a single scalar field, the form of the potential and the time dependence of the scalar field are uniquely determined by the evolution of the Hubble variable (up to a shift of the scalar field). During last years the models with two scalar fields have also become very popular. This is connected with the fact that some observations point out that the relation between the pressure and the energy density could be less than -1 \\cite{obs}. Such equation of state arises if the matter is represented by a scalar field with a negative kinetic term. This field is called ``phantom'' \\cite{phantom}. Moreover, according to some observations \\cite{obs1} the universe undergoes a transition between normal and phantom phase. Such an effect is dubbed ``phantom divide line crossing'' \\cite{divide}. In principle, the effect of phantom divide line crossing can be explained in the model with the only scalar field provided a special form of the potential and initial conditions is chosen \\cite{AndCanKam} or in models with a non-minimally coupled scalar field \\cite{non-minim}. However, the models with two scalar fields, one standard and one phantom, look more ``natural'' for the description of the phantom divide line crossing \\cite{two-field,two-field1,we-two-field}\\footnote{These two fields   may have their origin in spontaneous breaking of primordial symmetries of moduli space with complex potential \\cite{two-field1}.}.  In our preceding paper \\cite {we-two-field} we have studied the procedure of reconstruction of the potential in two-field models. It was shown that there exists a huge variety of potentials and time dependences of the fields realizing the same cosmological evolution. Some concrete examples were considered, corresponding to the evolution beginning with the standard Big Bang singularity and ending in the Big Rip singularity \\cite{Rip}.  One can ask oneself: what is the sense of  studying  different potentials and scalar field dynamics if they imply the same cosmological evolution?    The point is that the scalar and phantom field can interact with other  fields and   influence not only  the global cosmological evolution but also  other observable quantities.  One of the possible effects of the presence of normal and phantom fields could be their influence on the dynamics of cosmic magnetic fields. The problem of the origin and of possible amplification of cosmic magnetic fields is widely discussed in the literature \\cite{magnetic}. In particular, the origin of such fields can be attributed to primordial quantum fluctuations \\cite{quantum} and their further evolution can be influenced by hypothetic interaction with pseudoscalar fields breaking the conformal invariance of the electromagnetic field \\cite{break,Sorbo}. In the present paper we consider the evolution of magnetic fields created as a result of quantum fluctuations, undergoing the inflationary period with unbroken conformal invariance and beginning the interaction with pseudoscalar or pseudophantom fields after exiting the inflation and entering the Big Bang expansion stage, which is a part of the Bang-to-Rip scenario described in the preceding paper \\cite{we-two-field}. We shall use different field realizations of this scenario and shall see how the dynamics of the field with negative parity influences the dynamics of cosmic magnetic fields. To our knowledge the possible influence of the two-field dynamics, (when one of two (pseudo)scalar fields is a phantom one) on the magnetic fields was not yet discussed in the literature.  Speaking of cosmic magentic fields we mean the large-scale galactic, intergalactic or super-cluster magnetic fields of order from $10^{-6} G$ to $10^{-11} G$ with correlation from 100 kpc to several Mpc to the extent that they are originated from scalar and, possibly gauge field fluctuations  after exiting the inflation. Their seeds may well have $10^{-18} - 10^{-27} G$ or less (see \\cite{magnetic}).  The structure of the paper is as follows: in Sec.~2 we recall  the Bang-to-Rip scenario and describe some examples of  different dynamics of scalar and phantom fields; in Sec.~3 we introduce the interaction of the fields (phantom or normal) with an electromagnetic field and write down the corresponding equations of motion; in Sec.~4 we describe the numerical simulations of the evolution of magnetic fields and present the results of these simulations; Sec.~5 is devoted to concluding remarks. ",
        "conclusions": "We have seen that the evolution of the cosmic magnetic  fields interacting with a pseudoscalar (pseudophantom) field is quite sensitive to the concrete form of the dynamics of this field in two-field models where different scalar field dynamics and potentials realize the same cosmological evolution.  We confirm the sensitivity of the evolution of the magnetic field with respect to its helicity given the sign of the coupling constant $\\alpha$ and that the $\\Phi$ is a monotonic function of time (as it is really so in our models). We give also some numerical estimates of the actual magnetic fields up to the factor of amplification of such fields during the inflationary period. The toy model of the Bang-to-Rip evolution studied in this paper, cannot  be regarded as the only responsible for the amplification of cosmic magnetic fields implying their present observable values. It rather complements some other mechanisms acting before. However, the difference  between cosmic magnetic fields arising  in various models (giving the same expansion law after the inflation) is essential. It may provide a discriminating test for such models.  Naturally, the influence of the interaction between a pseudoscalar (phantom)  field and a cosmic magnetic field on the dynamics of the latter depends on the velocity of change with time of the former. The larger is the time derivative of the pseudoscalar field, the more intensive is the growth of the magnetic field (remember that the evolution of the scalar field is  monotonic). The results of numerical calculations illustrated in Sec. 4 confirm these qualitative considerations. Moreover, one can see that there exists a certain range of values of the wave number $k$ (and hence of the corresponding wavelengths of the magnetic field) where the effect is stronger. Indeed, if $k$ is too small the interaction term is small as well and the evolution of the magnetic field is damped. On the other hand if the wave number $k$ is too large the interaction term is small compared to the oscillatory term, proportional as usual to $k^2$ and the evolution has practically oscillatory character. Thus, the study of interplay between the dynamics of global scalar fields providing the cosmological evolution and the magnetic fields looks promising. This interplay has another interesting aspect. The pseudoscalar-electromagnetic field interaction can imply a conversion of the photons into axions. Such an effect can cause the observable dimming of supenovae. While it was shown \\cite{dimming} that this effect cannot mimic the cosmic acceleration, it could nevertheless mimic the dark energy fluid with a phantom equation of state \\cite{dimming1}. Thus, the interrelation amongst an electromagnetic field, a scalar and a phantom field can reveal some surprises.      \\ack This work was partially supported by Grants RFBR  08-02-00923  and  LSS-4899.2008.2. The work of A.A. was  supported by grants FPA2007-66665, 2005SGR00564, 2007PIV10046,  by the Consolider-Ingenio 2010 Program CPAN (CSD2007- 00042) and Program RNP 2.1.1.1112."
    },
    "0806/0806.3888_arXiv.txt": {
        "abstract": "We present a short overview of the properties of faint Galactic X--ray binaries. We place emphasis on current classification scenarios. One of the important parameters for the faint sources is their intrinsic luminosity. In the case of low--mass X--ray binaries it has recently been realised that besides a phase of radius expansion, the duration of type I X--ray bursts can be used as a primer for the source luminosity in some cases. Further, we show that a very low equivalent width of hydrogen and helium emission lines in the optical spectrum alone is not a tell--tale sign for an ultra--compact system. Finally, we list and discuss some unusual sources that could be X--ray binaries. ",
        "introduction": "In recent years many new and often faint X--ray binaries have been discovered.  Sources are found in the deep images made in hard X--rays by INTEGRAL and in images obtained by {\\it Swift}, \\xmm\\ and \\chan. The sensitivity to soft X--rays of the CCD instruments on board the latter satellites as well as their small point--spread function allows for optical or near--infrared identification of the counterpart to the X--ray source which is essential for source classification. Especially for sources in the Galactic plane, crowding often requires the superb \\chan\\ positional accuracy and precludes the use of the existing optical and near--infrared sky survey data such as that of the (S)DSS and 2MASS. Although, in the case of obscured high--mass X--ray binaries as well as the supergiant fast X--ray transients, lower resolution X--ray images together with especially 2MASS often provide a fair assessment of the source type. There are currently several programs providing identifications of the newly discovered sources (e.g.~see various contributions to these proceedings and \\citealt{2005ATel..494....1S}, \\citealt{2005ATel..629....1S}, \\citealt{2007ATel.1072....1T}, \\citealt{2007ATel.1193....1T}, \\citealt{2006ApJ...647.1309T}, \\citealt{2008arXiv0802.0988M}, \\citealt{2008arXiv0802.1774C} to cite just a few). ",
        "conclusions": ""
    },
    "0806/0806.2240_arXiv.txt": {
        "abstract": "Quasi-stationary flows of  gas accreting onto a compact center are analyzed in the framework of general-relativistic radiation hydrodynamics, under assumptions of spherical symmetry and thin gas approximation. Numerical investigation shows that luminosity, redshift and gas abundance are correlated. The  gas  can constitute up to one third of the total mass of brightest low-redshift sources, but its abundance goes down to $1/30$ for  sources with luminosities close to the Eddington limit. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0806/0806.0245_arXiv.txt": {
        "abstract": "We report the X-ray observations of two radio pulsars with drifting subpulses: B0834$+$06 and B0826$-$34 using \\xmm\\. PSR B0834$+$06 was detected with a total of 70 counts from the three EPIC instruments over 50 ks exposure time. Its spectrum was best described as that of a blackbody (BB) with temperature $T_s=(2.0^{+2.0}_{-0.9}) \\times 10^6$~K and bolometric luminosity of $L_b=(8.6^{+14.2}_{-4.4}) \\times 10^{28}$ erg~s$^{-1}$. As it is typical in pulsars with  BB thermal components in their X-ray spectra, the hot spot surface area is much smaller than that of the canonical polar cap, implying a non-dipolar surface magnetic field much stronger than the dipolar component derived from the pulsar spin-down (in this case about 50 times smaller and stronger, respectively). The second pulsar PSR B0826$-$34 was not detected over 50 ks exposure time, giving an upper limit for the bolometric luminosity $L_b \\leq 1.4 \\times 10^{29}$ erg~s$^{-1}$. We use these data as well as the radio emission data concerned with drifting subpulses to test the Partially Screened Gap (PSG) model of the inner accelerator in pulsars. This model predicts a simple and very intuitive relationship between the polar cap thermal X-ray luminosity ($L_b$) and the ``carousel'' period ($P_4$) for drifting subpulses detected in the radio band. The PSG model has been previously successfully confronted with four radio pulsars whose $L_b$ and $P_4$ were both measured: PSR B0943$+$10, PSR B1133$+$16, PSR B0656$+$14, and PSR B0628$-$28. The \\xmm\\ X-ray data of PSR B0834$+$16 reported here are also in agreement with the model prediction, and the upper limit derived from the PSR B0826$-$34 observation does not contradict with such a prediction. We also include two other pulsars PSR B1929$+$10 and B1055$-52$ whose $L_b$ and/or $P_4$ data became available just recently. These pulsars also follow the prediction of the PSG model. The clear prediction of the PSG model is now supported by all pulsars whose $L_b$ and $P_4$ are measured and/or estimated. ",
        "introduction": "More than forty years after the discovery of radio pulsars, the mechanism by which they emit coherent radio beams is still not fully understood. Also, many properties of this radiation remain a mystery, especially the phenomenon of drifting subpulses. This puzzling phenomenon was widely regarded as a powerful tool for the investigation of the pulsar radiation mechanism. Recently, this phenomenon received renewed attention, mostly owing to the newly developed techniques for the analysis of the pulsar radio emission fluctuations (Edwards \\& Stappers 2002,2003; ES02,ES03 henceforth). Using these techniques, Weltevrede et al. (2006 a,b; W06a,b henceforth) presented results of the systematic, unbiased search for the drifting subpulses and/or phase stationary intensity modulations in single pulses of a large sample of pulsars. They found that the fraction of pulsars showing evidence of drifting subpulses is at least 60~\\% and concluded that the conditions for the drifting mechanism to work cannot be very different from the emission mechanism of radio pulsars.  It is therefore likely that the drifting subpulse phenomenon originates from the so-called inner acceleration region right above the polar cap, which powers the pulsar radiation. In the classical model of Ruderman \\& Sutherland (1975; RS75 henceforth) the subpulse-associated spark filaments of plasma circulate in the pure Vacuum Gap (VG hereafter) around the magnetic axis due to well known drift of plasma with non-corotational charge density (see Appendix A for more details). There are few periodicities characteristic for this model, called also the pulsar carousel model: the primary period $P_3$ which can be measured as a distance between the observed subpulse drift bands, the secondary period (apparent when drifting is aliased; Gil \\& Sendyk 2003 for detailed description), and the tertiary period $P_4$ (called also the carousel time\\footnote{designated as $\\hat P_3$ in RS75. Although this symbol is still in use, we advocate to replace it by $P_4$.}, as it is the time interval after which the gap plasma completes one full circulation around the magnetic pole). The carousel model is widely regarded as a natural and qualitative explanation of the drifting subpulse phenomenon. However, its original version published by Ruderman \\& Sutherland (1975; RS75 hereafter) predicts too high a drifting rate of the sparks around the polar cap, as compared with the observations of drifting subpulses (e.g. Deshpande \\& Rankin 1999, 2001; DR99,DR01 henceforth), and too high a heating rate of the polar cap (PC henceforth) surface due to the spark-associated back-flow bombardment, as compared with X-ray observations (e.g. Zhang et al. 2000). Another difficulty of the RS75 model is that recent calculations strongly suggest that the surface binding energy of both ions and electrons are too low to allow the development of a vacuum gap. Indeed, when the surface magnetic field is purely dipolar, then the gap can develop only in magnetars and several highest B-field pulsars (Medin \\& Lai 2007). Another type of inner accelerator model, named space-charge-limited flow (SCLF, Arons \\& Scharlemann 1979; Harding \\& Muslimov 1998), has been discussed in the literature, which assumes that both ions and electrons can be freely striped off the neutron star surface. Although this approximation is valid for most pulsars assuming a pure dipolar field at the polar cap region, a stronger, multipole magnetic field near the polar cap region (which is needed to make a large number of radio pulsars above the radio emission death line, Ruderman \\& Sutherland 1975; Zhang et al. 2000) would introduce a non-negligible binding energy of ions/electrons (Medin \\& Lai 2007), which renders the SCLF approximation no longer valid. Another difficulty of the steady-state SCLF model widely discussed in the literature is that it does not predict the existence of any ``sparks'' that could give rise to the drifting sub-pulses. So, in our opinion, it is not an attractive inner accelerator model to interpret pulsar radio emission.  Motivated by these shortcomings of the otherwise attractive VG model Gil, Melikidze \\& Geppert (2003; G03 henceforth) developed further the idea of the inner acceleration region above the polar cap by including the partial screening caused by the thermionic flow of ions from the PC surface heated by sparks. We call this kind of the inner acceleration region a \"Partially Screened Gap\" (PSG hereafter). The PSG is thermally self-regulated in such a way that the surface temperature is always close to but slightly lower (less than 1 percent) than the critical temperature at which the maximum co-rotational ion outflow occurs and the gap is fully screened (see Appendix and/or G03 for more details). Moreover, if the surface temperature was even few percent lower than the critical temperature, there would be a pure vacuum gap, with all the problems discussed above. Since the actual potential drop in the PSG is much lower than that of the pure VG model (RS75), the intrinsic drift rate and PC heating rate are compatible with measurements of $P_4$ and $L_b$, respectively.  The PSG model can be tested if two observational quantities are known: (i) the circulational period $P_4$ for drifting subpulses observed in radio-emission and (ii) the X-ray bolometric luminosity $L_b$ of thermal BB radiation from the hot polar cap (see Appendix A). Radio pulsars were targeted since beginning of X-ray astronomy for various scientific reasons. Zhang, Sanwal \\& Pavlov (2005; Z05 henceforth) were the first who made an attempt to resolve the mystery of drifting subpulses in radio pulsars by observing them in X-rays. They proposed to detect thermal X-ray photons from the PC heated by sparks of plasma likely to be associated with drifting subpulses observed in radio band. Their choice was the best studied drifting subpulse pulsar B0943+10. Using \\xmm\\  X-ray observatory they detected a weak source coincident with the target pulsar. Due to very small number of counts detected, no unambiguous spectrum could be obtained. However, they were able to fit the BB model to the data, although a power law model was acceptable as well. Within a BB model they inferred a bolometric luminosity $L_b \\sim 5\\times 10^{28}$ erg/s emitted from the hot spot (few MK) with a surface area much smaller (about 60 times) than the conventional polar cap area as defined by the bundle of last closed dipolar field lines. This radio pulsar was well studied by DR99, who described the number of sparks and the circulation time $P_4=37.4 P$ needed for them to complete one full revolution around the pole (where $P$ is the basic pulsar period). These properties as well could not be accounted for by the conventional theory, and some radical modification of RS75 model was required. It appears that PSG model not only resolves all the problems of the RS75 model, but also offers a clean prediction that can be used to test theories of the inner pulsar accelerator. ",
        "conclusions": "Within the partially screened gap (PSG) model of the inner acceleration region in pulsars developed by G03, we derived in Paper I a simple and clean relationship (eq.~[1]) between the thermal X-ray bolometric luminosity $L_b$ from hot PC heated by sparks and the circulation time $P_4$ of the spark-associated drift detected as the subpulse drift in pulsar radio emission. This relationship expresses the well justified assumption (Appendix A) that both the drifting rate and the polar cap heating rate are determined by the same value of electric field within the inner acceleration region. Indeed, the drifting rate described by measurable $P_4$ is determined by the tangent (with respect to surface magnetic field) component of the electric field, while the heating rate described by measurable $L_b$ is determined by its component parallel to the surface magnetic field in the (partially screened) gap. In Paper II we showed that PSRs B0943$+$10, B1133+16, B0628$-$20 and B0654+14, which were the only pulsars with both $L_b$ and $P_4$ known at that time, satisfied equation~(1) quite well (see also Fig.~1 and Table~1). This suggested that the PSG model may indeed be a reasonable description of the inner accelerator region near the polar cap. In this paper we support this view by demonstrating that another two pulsars (B0834+06, B1929+10 and B1055$-$52) also satisfy the equation~(1). Yet another pulsar B0826-34, in which only the upper limit for $L_b$ was obtained, demonstrated a consistency with equation~(1) as well.  Only for a handful of pulsars the circulation (carousel) time was measured or constrained so far. Measurement of $P_4$ by means of modulation spectral analysis requires a strong unevenness in the circulating system, maybe a distinguished group of adjacent sparks or even just a single spark (see also the scenario discussed by Gil \\& Sendyk, 2003; GS03 hereafter). Moreover, this feature should persist considerably longer than the circulation time. Such favorable conditions do not occur frequently in pulsars and therefore direct or indirect measurements of $P_4$ are very rare. In principle, in a clean case one should be able to detect the primary feature $P_3$, reflecting the phase modulation of regularly drifting subpulses, flanked by two symmetrical features corresponding to slower amplitude modulation associated with carousel circulation as well as direct low frequency feature $1/P_4$ (like in the case of PSR B0943+10; DR01, AD01 and GS03). However, results of Paper II clearly showed that $P_4$ can be found also in pulsars without regularly drifting subpulses (and/or in erratic drifting modes). This strongly suggested that no matter the degree of the organization of spark plasma filaments at the polar cap, the slow circumferential plasma drift was always performed at about the same rate in a given pulsar. The problem was how to reveal this motion. Two new methods were discussed or at least mentioned in Paper I. The 2-D phase resolved modulation spectral analysis developed by ES02 and ES03 and implemented by W06a, b was the first one. The second method based on examination of the distribution of nulls in the long sequence of single pulses was recently developed by HR07 and Rankin and Wright (2007; RW07 henceforth). In view of the main results obtained in this paper the latter method deserves some more detailed discussion here.  As discussed in section 3.1 there is a controversy about the actual value of $P_4$ in PSR 0834+06. AD05 reported that the alias-corrected $P_3/P=1.88 \\pm 0.01$ and $P_4/P=15 \\pm 0.8$, implying the number of sparks $N=P_4/P_3=8$. These authors found just one sequence of 64 pulses in which the fluctuation spectrum analysis revealed the low frequency feature at about 1/15=0.067. On the other hand RW07 found the non-aliased primary drift periodicity $P_3/P=2.16 \\pm 0.011$ and the number of sparks $N=15$, implying the tertiary long periodicity $P_4/P=30.24 \\pm 0.15$. This longer cycle with $P_4 \\sim 30 P$ was supported by our measurements of $L_b$ and PSG model expressed by equation~(1). RW07 examined an interaction between nulls and emission in PSR B0834+06. They found that null pulses are not randomly distributed and that the most likely periodicity in their appearance is about $30 P$. Following the previous discovery of HR07 that null pulses and drifting subpulses in PSR B1133+16 are associated with the same long periodicity (about $33 P$) RW07 convincingly argued that short pseudo-nulls (one pulsar period or less) are just a result of irregular distribution of subpulse subbeams/sparks that persist on time scales of at least hundreds of pulsar periods. The short-time pseudo-nulls appear when the line-of-sight cuts through the low-level emission region in the radio beam. Our results on both B1133+16 and B0834+16 strongly support this picture. The interesting question is then why AD05 obtained such a strong feature at $15 ~P$ for a sequence of 64 single pulses from B0834+06. RW07 admitted that they also found in their data some sequences showing $15P$ periodicity, which seemed to be a sub-harmonic of $30P$ cycle. We noticed yet another problem with the result of AD05. In our opinion, these authors have used incorrectly their equations~(2) and (3). In fact, as $\\Delta \\phi$ they used the longitudinal distance between the profile components and in consequence, the azimuthal magnetic angle between the neighboring subbeams was $\\Delta \\theta=50$ degs.  This ignored the subpulses appearing in the saddle of the profile. We believe that they should use $\\Delta \\theta \\sim 25$ degs, and as a result, the number of sparks would be $N=360/25=14$ instead 8. This is consisted with $P_4=N P_3=14\\cdot 2.16=30.24 P$ obtained by RW07 and supported by our results presented in this paper. In summary, we strongly believe that the actual value of $P_4$ in PSR B0834+06 is close to 30 pulsar periods and that $15 P$ corresponds to a first harmonic of the basic cycle. Some evidence of low frequency spectral features at both $0.033~c/P$ and $0.066~c/P$ can be seen in Figure~A19 of W06. Moreover, it seems that 14 sparks inferred by RW07 are more likely than 8 sparks inferred by AD05.  The essence of the PSG pulsar model is the presence of a strong, nondipolar surface magnetic field $B_s$, although it does not appear explicitly in equation (1); see Appendix A for details. The strong value of $B_s$ is necessary for providing enough binding (cohesive energy) to prevent the free flow of iron ions from the surface (Medin \\& Lai 2007; ML07 hereafter), while the small radius of curvature is needed to develop cascading pair production (e.g. Gil \\& Melikidze 2002). The latter phenomenon is essential for both shorting out the gap potential drop and providing a dense electron-positron plasma in the radio emission region (eg. Melikidze \\& Gil, 2000 and Gil, Lyubarski \\& Melikidze, 2004). When the calculations of ML07 are adapted to the PSG model, then one can derive the dependence of the surface magnetic field on the surface temperature $B_s=B_s(T_s=T_i)$; (we will give detailed description of this topic in a separate paper, but see Appendix A for some details). For the condensed Fe surface this relationship is represented by the solid red line in Figure 7 of ML07. We can apply this apparatus to our case of PSR B0834+06, with $L_b=(6.8^{+1.1}_{-1.3})\\times 10^{28}$ erg~$s^{-1}$, $T_s=(2.0^{+2.0}_{-0.9})\\times 10^{6}$ K, and the associated effective surface area of the hot spot $A_{p}=940~$ m$^2$. On the other hand, one can read off from Figure 7 in ML07 the range of values $B_s \\sim (1^{+1.3}_{-0.6})\\times 10^{14}$ G corresponding to $T_s=(2.0^{+2.0}_{-0.9})\\times 10^6$~K. Since the dipolar surface magnetic field and polar cap area are $B_d=3 \\times 10^{12}$ G and $A_{pc}=4.85 \\times 10^{4}$ m$^2$, respectively, we can find the effective surface area $A_p=A_{pc}B_s/B_d=(1.5^{+1.4}_{-0.9})\\times 10^{3}$ m$^2$. This is consistent with our estimate, in which $A_p$ is about 50 smaller than $A_{pc}$. Theoretically, this results naturally from the flux conservation of the open magnetic field lines. As pointed out in Paper II (see also references therein), the small size of the hot spot relative to the canonical polar cap area is a typical property of hot BB thermal radiation detected in a number of pulsars. An extreme case was published just recently by Pavlov, Kargaltsev, Wong et al. (2008; P08 hereafter), who reported on the Chandra detection of a very old (170 Myr) and close to the Earth (0.13 kpc and 0.184$^{+0.01}_{-0.017}$ ~kpc, according to ATNF (Manchester, Hobbs, Teoh et al. 2005) and NE2001 (Cordes \\& Lazio 2003) database, respectively) radio pulsar PSR J0108-1431, with a very weak dipolar surface magnetic field $B_d=2.52 \\times 10^{11}$ ~G and a low spindown $\\dot {E}=5.8 \\times 10^{30}$ erg~$s^{-1}$. During 30 ks exposure they detected 53 counts and found that the spectrum can be described by PL model or BB model equally well. For the latter model they obtained the bolometric luminosity $L_b=1.3 \\times 10^{28} d^{2}_{130}$~erg~$s^{-1}$, $T_s=3.2 \\times 10^{6}$~K and $A_p=50~d^{2}_{130}$~m$^2$, which translates into the hot spot radius as small as 4 meters. This is the smallest hot polar cap ever observed$^8$, with the ratio $b=A_{pc}/A_p=1.77 {\\times 10^3}/d^{2}_{130}$, equal to 1770 or 923 (highest ever obtained) for distances 0.13 and 0.18 kpc, respectively . Accordingly, the actual surface magnetic field $B_s=bB_d$ (see Gil \\& Sendyk 2000 and ML07) is equal to 4.5 or 2.3 $\\times 10^{14}$~G for a distance of 0.13 or 0.18 kpc, respectively. Interestingly, the latter value agrees almost exactly with ML07 (red solid line in their Fig. 7), while the former one implies too high a surface temperature exceeding 5 MK. Thus, the extremely small hot polar cap with $T_s=$3.2 MK results from the fact that the actual surface magnetic field must be about 1000 times stronger than the dipolar component, in order to provide enough cohesive energy to develop PSG in this pulsar. We can therefore say that the case of PSR J0108-1331 supports strongly the PSG pulsar model, the ML07 cohesive energy calculations for the condensed Fe polar cap surface and NE2001 distance to this pulsar (about 0.184 kpc). If one adopts 0.184 kpc as the proper distance to PSR J0108-1331, then the bolometric BB luminosity is $L_b \\sim 2.5 \\times 10^{28}$ erg~$s^{-1}$ and the efficiency $L_b/\\dot {E} \\sim 4.3 \\times 10^{-3}$. With this value the equation (1) predicts the tertiary periodicity $P_4/P \\sim$~12. However, the confirmation of this by means of single pulse radio observations of PSR J0108-1431 seems hopeless with present day possibilities, as the pulsar is also extremely weak in radio band (Tauris, Nicastro, Johnston et al. 1994).  Thus, our PSG model seems to account for the physical phenomena at and above the actual pulsar polar cap quite well. Other available inner acceleration models do not match the observations well. The pure vacuum gap model (Ruderman \\& Sutherland 1975) has $\\eta=1$. Although it also satisfies Eq.(1), it predicts a very high polar cap heating rate, typically $L_b \\sim 10^{-1}-10^{-2} \\dot E$ (Zhang et al. 2000), and therefore a very small $P_4$. The predicted high $L_b$ has been ruled out by the X-ray observations of many old pulsars (ZSP05, TO05, K06 and this paper), and the predicted low $P_4$ is also inconsistent with the radio observations. On the other hand, as discussed in \\S1 the steady-state SCLF model does not predict the existence of the ``sparks'' whose drifts around the polar cap region provide the most natural interpretation of the observed drifting sub-pulse patterns. A modified unsteady SCLF model (which has not been discussed in the literature) may be able to introduce a sparking-like behavior. Based on the similar logic (i.e. the potential drop along the magnetic field line in the gap is equal to the horizontal potential drop across the spark, see Appendix), a similar equation as Eq.(1) can be derived for the SCLF model. However, since this model introduces a very small effective $\\eta$ value ($\\eta \\sim (2\\pi R_*/cP)^{1/2} << 1$, Harding \\& Muslimov 2001), the predicted polar cap heating rate is too low to interpret the observations, typically $L_b \\sim 10^{-4}-10^{-5} \\dot E$ (Harding \\& Muslimov 2002). Also the corresponding drifting velocity is too small so that the predicted $P_4$ is too long as compared with the radio data. The PSG model predicts an intermediate particle inflow rate, and gives the clean prediction (Eq.[1]) which allows $L_b$ to be a moderate value. This is strongly supported by the data.  In order to solve the binding energy problem in the canonical dipolar magnetic field at the neutron star surface, it has been conjectured that drifting subpulse pulsars are bare strange stars (Xu et al. 1999). The simplest model does not allow a hot polar cap because of the high thermal conductivity of the bare strange star surface layer, which is ruled out by the data. Yue et al. (2006) argued that PSR B0943$+$10 may be a low mass quark star ($\\sim 0.02 {\\rm M}_\\odot$). However, pulsar drifting seems to be the most common behavior of radio pulsars (W06a,b), some of which have well measured mass around $1.4 {\\rm M}_\\odot$ (Thorsett \\& Chakrabarty 1999). We regard that the quark star scenario is no longer attractive in view of the latest observations. The cohesive energy calculations of Fe ion chains in ultra-strong magnetic field by ML07 seem to be strongly supported by the X-ray observations discussed in this paper.  Finally, we would like to address a hypotheses put forward by Becker, Kramer \\& Jessner et al. (2006) that in old pulsars ($>10^6$ yrs) the magnetospheric emission dominates over thermal emission, including both cooling radiation and hot polar cap emission component. These authors suggested that the latter radiation component decreases along with the former one, and if so, the hot polar caps in cooling neutron stars could be formed by anisotropic heat flow due to the presence of the magnetic field rather than by particle bombardment. While in young NSs with core temperature $\\simeq 10^8$~K the strong crustal magnetic fields may channel the heat toward the polar cap resulting in $T_s$ of a few MK (Perez-Azorin, Miralles \\& Pons 2006; Geppert, K{\\\"u}ker \\& Page 2006), in pulsars older than $10^6$ years this mechanism is much less efficient and the only viable process that can produce such hot and small polar caps is the back-flow particle bombardment. Almost all pulsars presented and examined in this paper are older than 1 Myr (an exception is 110 kyr PSR B0656+14). For instance, PSR B0834+06 is 3 Myr old and its X-ray emission is dominated by hot BB component (an obvious counter-example arguing against Becker's claim). In PSR B1929+10 (3.1 Myr old) the luminosity of hot BB component is at least comparable with the magnetospheric X-ray radiation (M07). The very old (170 Myr) rotation powered non-recycled pulsar J0108-1431 clearly shows BB radiation from the hot polar cap (P08), probably accompanied by the magnetospheric emission, but no evidence of cooling radiation from the whole surface, as expected for such an old pulsar.  In summary, both the polar cap full cascade (Zhang \\& Harding 2000) and the downward outer gap cascade (Cheng, Gil \\& Zhang 1998) that have been proposed to interpret non-thermal X-ray emission from spindown-powered pulsars are expected to be less significant in pulsars from our sample with respect to the young pulsars. The predicted values of X-ray luminosity in these models are typically lower than that of the polar cap heating in the PSG model (Eq.[1]). In view that other available models of the pulsar inner accelerator (pure vacuum gap model and space-charge-limited flow model) either overpredict or underpredict the polar cap heating level, we conclude that the pulsar inner accelerator is likely partially screened due to a self-regulated sub-Goldreich-Julian flow. Also, the pure vacuum gap model predicts too fast a drifting and the space-charge-limited flow model has no natural explanation for the subpulse drift phenomenon at all. We thus strongly believe that thermal radiation associated with a polar cap heating due to partially screened inner accelerator (PSG) is a common component of pulsar X-ray emission regardless of its age, and this component plays especially significant role in the spectra of old pulsars."
    },
    "0806/0806.0629_arXiv.txt": {
        "abstract": "Transiting exoplanetary systems are surpassingly important among the planetary systems since they provide the widest spectrum of information for both the planet and the host star. If a transiting planet is on an eccentric orbit, the duration of transits $\\TD$ is sensitive to the orientation of the orbital ellipse relative to the line of sight. The precession of the orbit results in a systematic variation in both the duration of individual transit events and the observed period between successive transits, $P_{\\obs}$. The periastron of the ellipse slowly precesses due to general relativity and possibly the presence of other planets in the system. This secular precession can be detected through the long-term change in $P_{\\obs}$ (transit timing variations, TTV) or in $\\TD$ (transit duration variations, TDV). We estimate the corresponding precession measurement precision for repeated future observations of the known eccentric transiting exoplanetary systems (XO-3b, HD~147506b, GJ~436b and HD~17156b) using existing or planned space-borne instruments. The TDV measurement improves the precession detection sensitivity by orders of magnitude over the TTV measurement. We find that TDV measurements over a $\\sim4$\\,year period can typically detect the precession rate to a precision well exceeding the level predicted by general relativity. ",
        "introduction": "Since the discovery of the first transiting extrasolar planet \\citep{charbonneau2000,brown2001}, the number of such systems has increased to more than 30\\footnote{See http://exoplanet.eu for up to date information}. These transiting extrasolar planets (TEPs) provide unique information on the properties of the system. Based on the geometry provided by the transit light curve(s), the inclination, the physical radius and mass, therefore the density and the surface gravity can be derived, in addition to the mass of the planet. Moreover, the time between successive transits can be measured with an exceedingly high accuracy ($\\sim 10^{-6}$~--~$10^{-7}$, relative to the period). The detection of long--term transit timing variations can be used to learn more beyond the properties of the parent-star system \\citep{miralda2002,steffen2007}. They can be indicative of the presence of other planetary companions \\citep[see e.g.][]{holman2005,agol2005,millerricci2008}, co-orbital companions \\citep[Trojans, see][]{ford2007}, or satellites \\citep{simon2007} in the system, could provide information on the oblateness of the host star, or can be used to detect the additional prograde periastron precession predicted by general relativity (GR) \\citep{miralda2002,heyl2007}. Secular variations in the semimajor axis (and therefore in the transit timing) are also predicted on the time scale of stellar life due to the anisotropic light redistribution \\citep[a.k.a. Yarkovski-effect, see][]{fabrycky2008}. Furthermore, \\citet{iorio2006} has shown that TEP observations can in principle also test the gravitoelectric correction of GR by measuring the radial velocity amplitude and transiting periodicity simultaneously, in order to verify that the third Kepler's law requires a semimajor axis dependent correction.  In a pioneer study, \\citet{miralda2002} derived the modification of the observed time period between successive transits $P_{\\obs}$, called transit timing variations (TTVs), caused by the standard periastron precession due to GR \\citep[e.g.][]{misner1973} and the perturbations of other planets if present. Recently, \\citet{heyl2007} have extended these studies and estimated the precision of precession rate measurements for long--term mock observations of eccentric transiting extrasolar planets (ETEPs). Both studies restricted to small eccentricities. At that time, the existence of close eccentric planets was known only through radial velocity measurements, and no ETEPs had been observed. Since their publication, four transiting extrasolar planets have been discovered with significant eccentricity: XO-3\\lowercase{b} \\citep{johnskrull2007}, HD~147506\\lowercase{b} \\citep[a.k.a. HAT-P-2,][]{bakos2007}, GJ~436\\lowercase{b} \\citep{gillon2007,butler2004}, and HD~17156\\lowercase{b} \\citep{fischer2007}. Therefore it is now possible, for the first time, to make specific predictions for future, long--term measurements of periastron precession effects for real exoplanetary systems.  In this paper we determine the precision by which repeated long--term future ETEP observations will be able to detect the periastron precession rate for existing systems. In addition to TTVs, i.e. the slow modulation of $P_{\\obs}$ considered previously \\citep{miralda2002,heyl2007}, the periastron precession also changes the time durations $\\TD$ of individual transits. We examine whether these transit duration variations (TDVs) can be used to improve the sensitivity of periastron precession measurements. We estimate the precession rate measurement precision for long term repeated observations of $P_{\\obs}$ and $\\TD$ for the known ETEPs. Since several of the observed ETEPs have large eccentricities, we derive expressions for both TTVs and TDVs which are applicable for arbitrary eccentricities. We estimate whether future observations of currently known ETEPs will be able to reach the sensitivity necessary to test the prediction of GR, using existing or planned space-borne instruments. We refer the reader to a recent independent study by \\citet{jordan2008}, of precession rates in eccentric transiting extrasolar planets.  The next section of this paper introduces the geometrical description which is the basis of our calculations, and derives the expected transit timings and durations for planets orbiting a star with an arbitrarily large eccentricity. In \\S~\\ref{s:realsystems}, we utilize our results for the confirmed four ETEP systems, and give predictions for future observations of periastron precession with space-borne observations. Our conclusions are discussed in \\S~\\ref{s:summary}. ",
        "conclusions": "\\label{s:summary}  The first four eccentric transiting exoplanetary systems have been discovered during 2007. The precession of an eccentric orbit causes variations both in the transit timings and transit durations. We estimated the significance of measuring the corresponding observable effects compared to the inevitable precession rate of general relativity. We applied these calculations to predict the significance of measuring the effect for the four known eccentric transiting planetary systems. Our calculations show that a space-borne telescope is adequate to detect the change in the transit durations to a high significance better than the GR periastron precession rate within a 3~--~4 year timespan (in a continuous observing mode). The same kind of instruments would need more than a decade to detect the corresponding transit time variations to this sensitivity even for the most optimistic known system.  The CoRoT mission has already found two transiting planets \\citep[see][]{barge2008,alonso2008} and there are two known planets in the planned field-of-view of the Kepler mission \\citep[see][]{odonovan2006,pal2008}. Our results suggest that if an \\emph{eccentric} transiting planet is found in the Kepler or CoRoT field, these missions will be able to measure the periastron precession rate to a very high significance within their mission lifetime or with the support of ground-based observations on a longer time scale. This will provide an independent test of the theory of general relativity and will also be useful for testing for the presence of other planets in these systems."
    },
    "0806/0806.2306_arXiv.txt": {
        "abstract": "Thermal radio and X-ray emission has been traditionally associated with the formation of stars. However, in recent years,  non-thermal radiation from massive star forming regions has been detected.  Synchrotron radio emission and non-thermal X-rays from the outflows of massive young stellar objects (YSOs) provide evidence of the presence of relativistic particles in these sources. In YSOs, the acceleration of particles is likely produced where the thermal jet impacts on the surrounding medium and a shock wave is formed. Thus, particles might be accelerated up to relativistic energies through a Fermi-I type mechanism.  Relativistic electrons and protons can interact with thermal particles and photons, producing then $\\gamma$-rays. These energetic photons could be detected by the new generation of instruments,  making massive YSOs  a new population of $\\gamma$-ray surces.  In the present contribution we briefly describe some massive star forming regions from which non-thermal radio emission has been detected. In addition, we present a general model for high-energy radiation from the massive YSOs embedded in these regions. We take into account both leptonic and hadronic interactions of particles accelerated at the termination points of the collimated outflows ejected from the protostar. ",
        "introduction": "The mechanism of formation of massive stars ($M > 8M_{\\odot}$) remains one of the open questions in astrophysics. Massive stars appear in massive star associations where cloud fragmentation seems to be common. It is known that these stars originate inside giant and massive molecular clouds but the sequence of processes that take place during the formation of the star are mostly unknown. It has been suggested, for example, that the coalescence of various protostars in the same cloud can lead to the emergence of a massive star (Bonnell et al. 1998). Alternatively, a massive star could form by the collapse of the core of a molecular cloud, with associated episodes of mass accretion and ejection, as observed in low-mass stars (Shu et al. 1987). In such a case, the effects of jets propagating through the medium that surrounds the protostar should be detectable.  Until now, the formation of stars has been mostly associated with thermal radio and X-ray emission. However, non-thermal radio emission has been detected in some massive star forming regions. This is a clear evidence that efficient particle acceleration is occurring there, which may have as well a radiative outcome at energies much higher than radio ones.  In the present contribution, based on recent multiwavelength observations and reasonable physical assumptions, we show that massive protostars could produce a significant amount of radiation in the gamma-ray domain, because of the dense and rich medium in which they are formed. ",
        "conclusions": "In this work we show that, if the source is located at few kpc, the high-energy emission may be detected by GLAST and even by forthcoming Cherenkov telescope arrays after long enough exposure.  This opens a new window to the study of star formation and related processes. Also, determinations of the particle spectrum and its high-energy for different sources with a variety of environmental conditions can shed light on the properties of galactic supersonic outflows, and on the  particle acceleration processes occurring at their termination points.  Radio observations already demonstrate that relativistic electrons are produced in some sources. According to the presence of non-thermal emission detected at cm-wavelengths and IR observations of the protostar emission we can suggest several good candidates to be targeted by GLAST. These objects are IRAS~16547$-$4247 (Araudo et al. 2007), the multiple  radio source in Serpens (Rodr\\'iguez et al. 1989, Curiel et al. 1993), HH~80$-$81 (Mart\\'i et al 1993) and W3 (Reid et al. 1995, Wilner et al. 1999).  To conclude, we emphasize that massive YSO with bipolar outflows and non-thermal radio emission can form a new population of gamma-ray sources that could be unveiled by the next generation of $\\gamma$-ray instruments."
    },
    "0806/0806.0135_arXiv.txt": {
        "abstract": "We have reexamined the similarity solution for a self-gravitating isothermal gas sphere and examined implication to star formation in a turbulent cloud.  When parameters are adequately chosen, the similarity solution expresses an accreting isothermal gas sphere bounded by a spherical shock wave.  The mass and radius of the sphere increases in proportion to the time, while the central density decreases in proportion to the inverse square of time. The similarity solution is specified by the accretion rate and the infall velocity.  The accretion rate has an upper limit for a given infall velocity.   When the accretion rate is below the upper limit, there exist a pair of similarity solutions for a given set of the accretion rate and infall velocity.  One of them is confirmed to be unstable against a spherical perturbation.  This means that the gas sphere collapses to initiate star formation only when the accretion rate is larger  than the upper limit.  We have also examined stability of the similarity solution against non-spherical perturbation.  Non-spherical perturbations are found to be damped. ",
        "introduction": "Similarity solutions have contributed very much to our understanding of star formation process.  The classical similarity solution by \\citet{larson69} and \\citet{penston69} elucidated the runaway nature of gravitational collapse. The density increases in proportion to the inverse square of the time during the runaway collapse phase.  We learned from the similarity solutions of \\citet{shu77} and \\citet{hunter77} that the accretion rate of a protostar is of the order of $ c _s {}^3 / G $ where $ G $ and $ c _s $ denote the gravitational constant and the isothermal sound speed of gas.  The similarity solution is also used to evaluate the effects of rotation and magnetic field.  The similarity solutions of \\citet{narita84} and \\citet{saigo98} indicated that the runaway collapse cannot be prevented by rotation if once initiated. Collapse of a rotating magnetized gas cloud is described by the similarity solution of \\citet{krasnopolsky02}. Ambipolar diffusion is taken into account in the the similarity solution of \\citet{adams07}. \\citet{tsai95} extended the similarity solution to include shock wave.  \\citet{shu02} extended the similarity solution involving a shock wave for application to champagne phase of an \\ion{H}{2} region.  \\citet{tsai95} found two classes of similarity solution; the first class describes accretion onto protostar while the second one does failure of star formation.  The central density decreases in proportion to the inverse square of the time, $ \\rho _c \\, \\propto \\, t ^{-2} $, in the second class solution.  Although this solution has not gained much attention thus far, it provides an insight on dynamical compression of a molecular cloud core. If a dense clump of gas is compressed by an external force, the temporal increase in the density may trigger gravitational collapse and star formation.  One can surmise existence of threshold of gravitational collapse.  If the dynamical compression is either weak or short, the clump will bounce back to expansion.  A shock wave will be formed when accreting gas is stoped by the expansion \\citep[see, e.g.] {adams07}. The similarity solution of \\cite{tsai95} demonstrated that a spherical cloud can expand even when it is steadily compressed by a shock wave. On the other hand, the shock compressed gas sphere will collapse owing to its self gravity if the shock is strong and lasts for a long enough period.  In this paper we reexamine the similarity solution of \\cite{tsai95} while keeping its negative implication in mind.  We find the condition for existence of similarity solution describing expansion of a gas sphere. Conversely it will tell us condition for a shock compressed gas sphere to collapse by its self gravity. We also study stability of the similarity solution. The similarity solution denies collapse due to the self gravity only when it is stable.  We review the similarity solution in \\S 2.1 and show the method of linear stability analysis in \\S 2.2. Technical details on the stability analysis are given in Appendix.  Properties of similarity solutions, such as accretion rate and infall velocity, are shown in \\S 3.1. Stability of the similarity solution is given in \\S 3.2 and \\S 3.3. We discuss implications of our analysis in \\S 4. ",
        "conclusions": "We obtained the critical accretion rate above which there exits no similarity solution.  The critical rate can be interpreted as the minimum accretion rate for a high density clump to initiate self-gravitational collapse. The critical accretion rate can be rewritten as \\begin{equation} \\left. \\frac{dM}{dt} \\right| _{\\rm cr} \\; = \\; \\frac{3.6 \\, c _s ^4}{G v} \\, , \\label{critical-phys} \\end{equation} for $ v \\, \\ga \\, 3 c _s$ in the dimensional form.  Equation (\\ref{critical-phys}) gives us an estimate for a converging flow to initiate gravitational collapse.  We shall consider a spherical region of which surface is surrounded by a converging flow. The radius, inflow velocity, and density are assumed to be $ r$, $ v $ and $ \\rho $, respectively.  Then the gravitational collapse will be initiated when the mass accretion rate exceeds the critical, \\begin{equation} 4 \\pi r ^2 \\, \\rho \\, v \\; > \\; \\frac{3.6 \\, c _s ^4}{G \\, v} \\, . \\label{Jeans1} \\end{equation} Equation (\\ref{Jeans1}) can be rewritten as \\begin{equation} \\frac{2r}{\\lambda _{\\rm J}} \\; > \\; \\frac{0.3 \\, c _{\\rm s}}{v} \\, , \\label{Jeans2} \\end{equation} where \\begin{equation} \\lambda _{\\rm J} \\; = \\; \\frac{2 \\pi \\, c _s} {\\displaystyle \\sqrt{4 \\pi G \\rho}} \\; . \\end{equation} Since $ \\lambda _{\\rm J} $ denotes the Jeans length, Equation (\\ref{Jeans2}) means that the effective Jeans length reduces in proportion to the inverse of the Mach number.  The Jeans mass is proportional to the cube of the Jeans length for a given density.  Thus the effective Jeans mass should reduce to \\begin{eqnarray} M _{\\rm J,~eff} & = & M _{\\rm J} \\, \\left( \\frac{\\lambda _{\\rm J,~eff}}{\\lambda _{\\rm J}}\\right) ^3 \\\\ & = & M _{\\rm J} \\, \\left( \\frac{v}{0.3~c_s}\\right) ^{-3} \\, . \\end{eqnarray} This implies that compression of sub Jeans mass clump may result in gravitational collapse in the region of flow convergence.  Note that the effective Jeans mass is several order of magnitude smaller than the classical one when $ v \\, \\ga \\, 3 \\, c _s $.  The compression should continue for a certain timescale for a dynamically compressed clump to collapse by its self gravity. If we evaluate the minimum timescale to be the effective Jeans length divided by the flow velocity, it is shorter than the free-fall timescale by a factor of the Mach number squared, \\begin{equation} \\tau _{\\rm comp} \\; \\simeq \\; \\frac{\\lambda _{\\rm J,~eff}}{v} \\; \\simeq \\; \\tau _{\\rm ff} \\, \\left( \\frac{v}{c _s}\\right) ^{-2} \\, . \\end{equation}  The timescale can be translated into the wavelength of perturbation. A compressed clump can collapse by the self gravity if the wavelength of velocity perturbation is longer than the effective Jeans length. If turbulence contains velocity perturbations of long wavelengths, gravitational collapse due to dynamical compression will take place somewhere in the cloud. In such case we can expect a number of clumps of which masses are much smaller than the classical Jeans mass."
    },
    "0806/0806.4136_arXiv.txt": {
        "abstract": "{High-energy observations  have unveiled peculiar  classes of isolated  neutron stars  which, at  variance with  radio  pulsars, are mostly radio silent and not powered by the star rotation.  Among these objects   are   the    magnetars,   hyper-magnetized   neutron   stars characterized  by transient  X-ray/$\\gamma$-ray emission,  and neutron stars  with  purely  thermal,  and  in most  cases  stationary,  X-ray emission (a.k.a., X-ray dim  isolated neutron stars or XDINSs).  While apparently  dissimilar in  their  high-energy behavior  and age,  both magnetars and XDINSs have  similar periods and unusually high magnetic fields.  This suggests a  tantalizing scenario where the former evolve into  the  latter.}{Discovering  so  far  uninvestigated  similarities between the multi-wavelength properties  of these two classes would be a further step forward to  establish an evolutionary scenario.  A most promising channels is the near infrared (NIR) one, where magnetars are characterized by a distinctive spectral flattening with respect to the extrapolation of the soft  X-ray spectrum.}{We observed the two XDINSs \\zerofour\\ and \\oneight\\ with the  \\madn\\ (\\mad) at the \\vltn\\ (\\vlt), as  part of the  instrument guaranteed  time observations  program, to search for  their NIR  counterparts.}  {Both \\oneight\\  and \\zerofour\\ were  not  detected  down  to   K$_s  \\sim20$  and  K$_s  \\sim  21.5$, respectively.  }   {In order to constrain the  relation between XDINSs and  magnetars  it  would  be  of importance  to  perform  deeper  NIR observations.  A good candidate is  \\oneseven\\ which is the XDINS with the highest inferred magnetic field.} ",
        "introduction": "Over the last two  decades, high-energy observations have unveiled the existence of  peculiar classes of  isolated neutron stars  (INSs) like the anomalous X-ray pulsars  and the soft $\\gamma$-repeaters (AXPs and SGRs),  i.e., the magnetar  candidates (e.g.,  \\citealt{mererev}), and the X-ray  Dim INSs (XDINSs; e.g.,  \\citealt{frank07}).  Despite their vastly  different  observational  manifestations, both  magnetars  and XDINSs  are slow  rotators ($P\\sim  1$--10 s),  with  surface magnetic fields reaching $\\approx  10^{14}$--$10^{15}$~G in the magnetars case, have  persistent  X-ray luminosities  much  larger  than the  inferred rotational  energy losses $L_{\\mathrm  X}\\approx 10$--$100\\,  \\dot E$, and they are usually radio-silent.  The XDINSs X-ray spectra are  purely thermal and well represented by a blackbody ($kT\\approx 50$--100 eV) whose emission radius is consistent with a sizable  fraction of the neutron star  surface.  In most cases, one     or     possibly     more     broad     absorption     features ($E_{\\mathrm{line}}\\approx  0.2$--0.7  keV)  have been  observed  (see Haberl 2007 and references therein), which hint toward magnetic fields of  $\\sim 10^{13}-10^{14}$ ~G,  comparable to  those derived  from the spin  down.  The  absence of  power-law tails  is consistent  with the faint  rotation--powered  magnetospheric  emission expected  from  the inferred spin-down luminosity ($\\dot E \\approx 3-5 \\times 10^{30}$ erg s$^{-1}$).  In  the optical,  only a few  XDINSs have  been identified (see Kaplan  2008 and references therein) and  a candidate counterpart was recently proposed  for \\oneseven\\ (Zane et al.   2008), while that of  \\zerofour\\ was  not confirmed  (Mignani et  al.,  in preparation). Their spectra mostly follow a Rayleigh-Jeans, with fluxes exceeding by a  factor  $\\ge  10$   the  low-energy  extrapolation  of  the  X-rays best-fitting blackbody \\cite[e.g.,][]{kap08}.   Whether this is indeed due  to  emission  from  regions  of the  star  surface  at  different temperature \\cite[e.g.,][]{pons02} is still under debate.  Like in the X-rays, no rotation--powered  magnetospheric emission is detectable in the optical, unless XDINSs are a factor of $\\ge 10^{3}$ more efficient emitters than rotation--powered neutron stars (Zharikov et al.  2006).  Whether there is a  (evolutionary?)  link between magnetars and XDINSs is still hotly debated (see McLaughlin et al. 2006; Popov et al. 2006, for  a  discussion). As  proposed  by  Mignani  et al.   (2007a),  the detection of  the optical/near  infrared (NIR) flattening  observed in the  magnetar  spectra,  ascribed  either to  magnetospheric  emission powered by the  magnetic field or to thermal  emission from a fallback disk, would strengthen  such a link.  For the  XDINSs, such flattening would be easily recognizable in the NIR, where the contribution of the optical   Rayleigh-Jeans  continuum  is   negligible.   So   far,  NIR observations of XDINSs performed with the \\vltn\\ (\\vlt) did not unveil counterparts down to H= 21.5--22.9 (Lo Curto et al. 2007).  We report on new NIR observations of \\oneight\\ and \\zerofour\\ recently performed in  the K$_s$  band with the  \\vlt, the first  presented for these sources.   Observations are described in Sect.  2, while results are presented and discussed in Sects. 3 and 4, respectively. ",
        "conclusions": "Our measured K$_s$-band spectral flux  upper limits are well above the extrapolation of  the optical spectrum  of \\oneight\\ (van  Kerkwijk \\& Kulkarni  2001),  which  would   predict  K$_s  \\sim  26.5$,  and  the Rayleigh-Jeans extrapolation  in the  optical/NIR domain of  the \\xmm\\ spectrum  of \\zerofour\\  (Haberl  et  al.  2004).   We  note that,  for \\oneight, the  K$_s$-band spectral flux upper limit  is shallower than the H-band one  of Lo Curto et al.~(2007), while  for \\zerofour\\ it is slightly  deeper. Thus,  we  are not  yet  able to  constrain the  NIR spectrum of \\oneight\\ and \\zerofour.  A blackbody spectrum produced by the neutron star surface is obviously a  possibility.  It  is clear  that for  the \\oneight\\  and \\zerofour\\ distances of  $\\sim 160$  pc (van Kerkwijk  \\& Kaplan 2007)  and $\\sim 350$ pc (Posselt  et al.  2007), respectively, and  for any reasonable combination of temperature and emitting area, the predicted flux would fall well below our upper  limits.  A blackbody spectrum produced by a fallback disk  (Perna et  al. 2000) could  yield to a  flux compatible with the  present NIR  flux upper limits  of \\oneight\\  and \\zerofour, depending  on the actual  disk size  and accretion  rate (see  also Lo Curto et al.  2007). A power--law spectrum could  be produced by non--thermal emission from the neutron  star magnetosphere.   In Fig.\\ref{Lir} (left),  we plotted the NIR  luminosity ($L_{NIR}$) vs  the rotational energy  loss ($\\dot E$) for  different classes  of INSs: rotation--powered  neutron stars, magnetars, and XDINSs. Our upper limits do not rule out a NIR emission efficiency, $\\eta_{NIR} \\equiv L_{NIR}/\\dot  E$, comparable to that of rotation--powered neutron  stars, or higher.  This  would be possible, for instance,  if the fraction  of energy emitted in  the relativistic wind is  much less than in  rotation--powered neutron stars  or if the NIR emission is powered not by the neutron star rotation but, e.g.,  by its magnetic field.  Although  XDINSs are less luminous than magnetars in absolute terms (Fig.\\ref{Lir},  left), they might have a comparable NIR emission efficiency.  This is shown in Fig.\\ref{Lir} (right) where the NIR emission efficiency is  plotted vs.  the dipole magnetic field $B$ inferred  from the  spin down. We  warn  here   that  NIR  observations  of  the   radio  pulsar  PSR J1119$-$6127 (Mignani et al.  2007b) suggest a NIR emission efficiency comparable to  that of other rotation--powered  neutron stars, despite of its relatively high magnetic field ($B \\sim 4.1 \\times 10^{13}$ G). If the inferred  magnetic field value of PSR  J1119$-$6127 is correct, this  might  argue  against  a  connection between  the  NIR  emission efficiency and the magnetic field.  Much deeper  observations are required  to better constrain  the XDINS optical/NIR  emission  properties and  to  investigate their  possible connection  with   the  magnetars.   In   this  respect,  particularly important would be the study of \\oneseven, possibly the XDINS with the highest  magnetic field,  as  hinted  by the  detection  of the  X-ray absorption feature (Zane et  al.  2005).  Its anomalously high optical emission is  indeed incompatible with  both thermal emission  from the neutron star surface and rotation--powered non--thermal emission (Zane et al. 2008)."
    },
    "0806/0806.1525_arXiv.txt": {
        "abstract": "We study the probability distribution function (PDF) of the mass density in simulations of supersonic turbulence with properties appropriate for molecular clouds.  For this study we use Athena, a new higher-order Godunov code. We find there are surprisingly similar relationships between the mean of the time-averaged PDF and the turbulent Mach number for driven hydrodynamic and strong-field MHD turbulence.  There is, however, a large scatter about these relations, indicating a high level of temporal and spatial variability in the PDF.  Thus, the PDF of the mass density is unlikely to be a good measure of magnetic field strength. We also find that the PDF of decaying MHD turbulence deviates from the mean-Mach relation found in the driven case.  This implies that the instantaneous Mach number alone is not enough to determine the statistical properties of turbulence that is out of equilibrium. The scatter about the mean-Mach relation for driven turbulence, along with the large departure of decaying turbulence PDFs from those of driven turbulence, may illuminate one factor contributing to the large observed cloud-to-cloud variation in the star formation rate per solar mass. ",
        "introduction": "The mechanism that determines the star formation rate (SFR) within a molecular cloud (MC) is not well understood.  Observations show that the SFR per solar mass, as measured by the ratio of CO to IR luminosity, varies by as much as 2-3 orders of magnitude from region to region; instead Gao \\& Solomon (2004) found that HCN emission, which measures molecular gas at much higher density ($\\gtrsim 3 \\times 10^4 \\unit{cm^{-3}}$), is the better indicator of star formation.  This suggests that the amount of mass to reach high density is the key factor in determining the SFR per solar mass within a cloud (McKee \\& Ostriker 2007).  Observed non-thermal line widths in MCs (Falgarone \\& Philips 1990) indicate that supersonic turbulence may be responsible for creating the high density contrasts that lead to clump formation.   Although magnetic fields have been shown not to lengthen appreciably the decay timescale of supersonic turbulence (Stone, Ostriker, \\& Gammie 1998, hereafter S98; Mac Low 1999), they do create anisotropy in the structures within the medium (Vestuto et al. 2003). Observations using Zeeman splitting, such as those described by Crutcher (1999), have found magnetic fields in some clouds strong enough (e.g. $\\beta \\approx 0.04$) that they cannot be neglected. The magnetic field within a MC is typically difficult to detect, motivating the determination of new diagnostics (e.g. Heyer et al. 2008).   The probability distribution function (PDF) of the logarithm of mass density can be used to quantify the amount of material within a turbulent medium that has a given density.  As the density is likely to have a strong impact on star formation, many groups have investigated the properties of such PDFs.  Padoan et al. (1997, hereafter P97), for example, conducted simulations of driven hydrodynamic turbulence, determining a relation between the mean of the PDF and the Mach number.  Most of the mass in the simulated clouds was found to be in only a small fraction of the volume, with the width of the approximately Gaussian PDF increasing with Mach number in a predictable way.  Ostriker, Stone, \\& Gammie (2001) investigated the effect of magnetic field strength on the mean-Mach relation for decaying turbulence, finding that the fast magnetosonic Mach number can be used to predict only a lower limit on the width of the distribution.   Recently, we have undertaken a comprehensive study of the properties of supersonic turbulence (Lemaster \\& Stone 2008, hereafter Paper I) with Athena, a new higher-order Godunov method (Stone et al. 2008).  This study represents one of the first applications of Godunov methods to the study of supersonic MHD turbulence, and therefore represents an important test of previous results. In this letter, we investigate the effects of magnetic field strength and Mach number on the PDF.  We survey a much larger range of Mach number and field strength, and at a higher numerical resolution, than has been used in previous studies.  This provides much better statistics and allows us to constrain the form of the relations more tightly.  We describe the numerical methods used to conduct our simulations in \\S\\ref{sec:methods}, explain the focus of our analysis in \\S\\ref{sec:pdfs}, and present the results in \\S\\ref{sec:results}. Finally, we summarize in \\S\\ref{sec:concl}. ",
        "conclusions": "\\label{sec:concl}   For both supersonic hydrodynamic and strong-field MHD turbulence, we have found a one-to-one correspondence between the mean of the time-averaged PDF and the Mach number.  The mean-Mach pairs used to fit these relations have very small residuals; however, the mean of the PDF at any given Mach number, in both the hydro and MHD cases, is smaller than was found by P97 for the purely hydrodynamic case.  Although there is substantial scatter of the mean-Mach pairs computed from instantaneous sub-volumes, the ensemble average of these values still falls close to the time-averaged global relation.  The scatter puts a small fraction of the instantaneous values in the vicinity of the P97 relation.  Although the relations found for hydro and MHD differ, the scatter about the mean relation of the instantaneous sub-PDF values creates substantial overlap between the two.  Since the MHD relation gives means that are typically smaller than the corresponding hydrodynamic values by only 1$\\sigma$, it will be very difficult to distinguish between the two observationally.  We have also found that PDFs in decaying MHD turbulence differ from those of driven MHD turbulence at the same magnetic $\\beta$.  It would seem that the instantaneous Mach number alone does not adequately describe the statistical state of the turbulent gas when not in equilibrium.  Again, however, there is substantial overlap between the equilibrium and non-equilibrium values, preventing this diagnostic from being used to distinguish between the two.  The scatter we have found about the mean-Mach relation may help explain the large variation in the observed SFR per solar mass in molecular clouds.  Since MCs are likely to be transient entities, relations found from driven turbulence may not even be applicable to real clouds.  If large departures from the mean-Mach relation are in fact linked to the large variation in the SFR per solar mass, this may indicate that turbulence in MCs is indeed decaying rather than forced."
    },
    "0806/0806.4888_arXiv.txt": {
        "abstract": "Large-format (sub)millimeter wavelength imaging arrays are best operated in scanning observing modes rather than traditional position-switched (chopped) modes. The choice of observing mode is critical for isolating source signals from various types of noise interference, especially for ground-based instrumentation operating under a bright atmosphere. Ideal observing strategies can combat $1/f$ noise, resist instrumental defects, sensitively recover emission on large scales, and provide an even field coverage -- all under feasible requirements of telescope movement. This work aims to guide the design of observing patterns that maximize scientific returns. It also compares some of the popular choices of observing modes for (sub)millimeter imaging, such as random, Lissajous, billiard, spiral, On-The-Fly (OTF), DREAM, chopped and stare patterns. Many of the conclusions are also applicable other imaging applications and imaging in one dimension (e.g.~spectroscopic observations). ",
        "introduction": "\\label{sec:intro}  %  The latest generation of imaging arrays for submillimeter and far-infrared applications, both on the ground\\cite{sharc2, gismo, laboca, scuba2, artemis} and in space\\cite{spire, pacs}, produce total-power readouts, providing snapshot views not unlike optical and infrared cameras. This is in contrast to the instruments of the past where data consisted of difference signals from source and a nearby ``off-position''. It is true, that the differential readout scheme provides an effective way of rejecting atmospheric variations (for ground-based instruments), and/or other sources of $1/f$ type noise interference, provided that the differencing happens at a fast enough rate. In practice, however, the noise rejection is rarely perfect (resulting in ``striping'') and the reconstruction of images\\cite{emerson88, emerson95, wright} from differenced data is at once challenging and riddled with problems due to the inadvertent filtering of spatial scales. Nevertheless, differencing remains the only effective way of observing from the ground, with a bright and variable atmospheric foreground, with just a single or a few pixels.  However, the large imaging arrays of today can do better since they collect information at many positions simultaneously, hence no longer needing explicit differencing of ``on'' and ``off'' positions. The noise (e.g.~correlated atmospheric variations, and detector $1/f$ noise) can be effectively separated from the astronomical source signals by capable algorithms\\cite{crush}, provided that the source signals are ``moved'' from pixel-to-pixel during the observation. However, sensitive recovery is only possible for source components, which do not overlap with the predominant noise signals. Some modes of observing are inherently better in isolating the source signals from noise interference and providing sensitivities to more extended spatial scales.  The benefits of faster scanning and that of cross-linking have been widely recognised\\cite{fastscanning, weferling, tegmark, waskett}. Faster scanning moves the source signals into the higher frequencies of the detector signal spectra, where the $1/f$ interference is less. Cross-linking assures that all Fourier components of the source, along all spatial directions, are scanned at the higher speeds. Several patterns are now used or have been proposed for various instruments. However, some of the observing modes are objectively better than others. Earlier attempts at systematic comparisons explored only some of the aspects involved\\cite{tegmark}, or described a framework but did not actually compare observing strategies\\cite{thesis}. This paper aims to provide the most comprehensive evaluation yet of observing strategies for (sub)millimeter-wave imaging.  In the first part of the paper, we outline the criteria that make good observing modes and provide guidelines for designing these. In the second part, we use this understanding to define quantitative measures that are consequently used on a selection of commonly used, or suggested, imaging observing patterns. ",
        "conclusions": "We have used objective criteria, such as resistance to noise measured by phase-space moments, and short-timescale scanning ranges as a indicators of large-scale sensitivities, to compare some of the commonly used observing strategies for astronomical (esp.~submillimeter) imaging.  Accordingly the best observing strategy is to move source signals randomly across the imaging array. Such modes are not easily implemented by scanning telescopes, which typically move along smooth, connected trajectories. For them, Lissajous or {\\em billiard} patterns offer a reasonable compromise for mapping small and/or large fields. Spirals also exhibit formidable qualities, both as standalone patterns and when combined in rasters. However, the more traditional OTF and position-switched (chopped) modes make relatively poor choices because of their more limited abilities to stand up to adverse interference from noise, and because they can be strongly directional in their sensitivities to the larger scales, unless appropriately cross-linked at an angle or several angles. The DREAM pattern, destined for SCUBA-2, proved surprisingly weak in the simulations, both in its noise immunity and sensitivities to large scales. Finally, stare modes are most easily ruined by noise and require significant overheads in observing time when compared to the scanning modes.  Observing strategies that require secondary reflector movement are not always effective. In ground-based submillimeter applications these modes produce strong signals that result from small changes in the primary illumination. Stare modes are similarly handicapped is such applications due to the optical loading changes between dark-frame calibrations and on-source observing, not to mention the significantly longer observing times that stare modes can require to reach the comparable sensitivities to scans.  To retain optimal sensitivities on the large scales, scanning patterns need to be comparable in size to the measured scales. Large fields may be mapped either by scaled versions of the patterns (where applicable), or by mosaicing together smaller fields. The two strategies are similar in their characteristic resistances to noise, which is mainly defined by the type rather than the size of the pattern used."
    },
    "0806/0806.3185_arXiv.txt": {
        "abstract": "% Two relativistic X-ray jets have been detected with the \\textit{Chandra} X-ray observatory from the black hole X-ray transient XTE J1550-564. We report a full analysis of the evolution of the two jets with a gamma-ray burst external shock model. A plausible scenario suggests a cavity outside the central source and the jets first travelled with constant velocity and then are slowed down by the interactions between the jets and the interstellar medium (ISM). The best fitted radius of the cavity is $\\sim$0.36 pc on the eastern side and $\\sim$0.46 pc on the western side, and the densities also show asymmetry, of $\\sim$0.015 cm$^{-3}$ on the east to $\\sim$0.21 cm$^{-3}$ on the west. A large scale low density region is also found in another microquasar system, H 1743-322. These results are consistent with previous suggestions that the environment of microquasars should be rather vacuous, compared to the normal Galactic environment. A generic scenario for microquasar jets is proposed, classifying the observed jets into three main categories, with different jet morphologies (and sizes) corresponding to different scales of vacuous environments surrounding them. ",
        "introduction": "Microquasars are well known miniatures of quasars, with a central black hole (BH), an accretion disk and two relativistic jets very similar to those found in the centers of active galaxies, only on much smaller scales (Mirabel $\\&$ Radr\\'{\\i}guez 1999). Since discovered in 1992, radio jets have been observed in several BH binary systems and some of them showed apparent superluminal features. In the two well known microqusars, GRS 1915+105 (Mirabel $\\&$ Radr\\'{\\i}guez 1999) and GRO J1655-40 (Tingay et al.1995; Hjellming $\\&$ Rupen 1995), relativistic jets with actual velocities greater than 0.9$c$ were observed. In some other systems, small-size ``compact jets\", e.g. Cyg X-1 (Stirling et al. 2001), and large scale diffuse emission, e.g. SS433 (Dubner et al. 1998), were also detected.  XTE J1550-564 was discovered with RXTE during its strong X-ray outburst on September 7, 1998 (Smith 1998). It is believed to be an X-ray binary system at a distance of $\\sim$5.3 kpc, containing a black hole of 10.5$\\pm$1.0 solar masses and a low mass companion star (Orosz et al. 2002). Soon after the discovery of the source, a jet ejection with an apparent velocity greater than 2$c$ was reported (Hannikainen et al. 2001). In the period between 1998 and 2002, several other outbursts occurred but no similar radio and X-ray flares were detected again in these outbursts (Tomsick et al. 2003).  With the help of the \\textit{Chandra} satellite, Corbel et al (2002) found two large scale X-ray jets lying to the east and the west of the central source, which were also in good alignment with the central source. The eastern jet has been detected first in 2000 at a projected distance of $\\sim$21$\\arcsec$ from the central black hole. Two years later, the jet could only be seen marginally in the X-ray image, while a western counterpart became visible at $\\sim$22$\\arcsec$ on the other side. The corresponding radio maps are consistent with the X-ray observations (Corbel et al. 2002).  There are altogether eight two-dimentional imaging observations of XTE J1550-564 in \\textit{Chandra} archive during June 2000 and October 2003 (henceforth observations 1$\\sim$8). Here we report a full analysis of these X-ray data, together with the kinematic and spectral evolution fittings for all these observations. ",
        "conclusions": "A GRB external shock model shows that a large scale cavity exists outside XTE J1550-564. This model has also been applied to another X-ray transient H 1743-322. Chandra X-ray and ATCA radio observations of H 1743-322 from 2003 November to 2004 June revealed the presence of large-scale ($\\sim$0.3 pc) jets with velocity $v/c\\sim0.8$ (Rupen et al. 2004; Corbel et al. 2005). Deceleration is also confirmed in this system. The external shock model describes the data of this source well. A cavity of size $\\sim$0.12 pc most likely exists, but the conclusion is not firm in this case. Even if there is no vacuum cavity, the ISM density is found to be very low($\\sim3\\times10^{-4}$ cm$^{-3}$), compared to the canonical Galactic value.  These studies led us to the suggestion that in microquasars the interactions between the ejecta and the environmental gas play major roles in the jet evolution and the low density of the environment is a necessary requirement for the jet to develop to a long distance. We find that microquasar jets can be classified into roughly three groups: small scale moving jets, large scale moving jets and large scale jet relics. For the first type, the ``small jets\", only radio emissions are detected. The jets are always relatively close to the central source and dissipate very quickly, including GRS 1915+105 (Rodr\\'{\\i}guez \\& Mirabel 1999; Miller-Jones et al. 2007), GRO J1655-40 (Hjellming \\& Rupen 1995), and Cyg X-3 (Marti et al. 2001). The typical spatial scale is 0$\\sim$0.05 pc and the time scale is several tenths of days. No obvious deceleration is observed before the jets become too faint. For the second type, the ``large jets\", both X-ray and radio detections are obtained, at a place far from the central source several years after the outburst. Examples are XTE J1550-564, H1743-322, and GX 339-4 (Gallo et al.2004). The typical jet travelling distance for this type is 0.2$\\sim$0.5 pc from the central engine and deceleration is clearly observed. The last type, the ``large relics\", is a kind of diffuse structures observed in radio, optical and X-ray band, often ring or nebula shaped that are not moving at all. In this class, some well studied sources, Cygnus X-1 (Gallo et al.2005), SS433 (Dubner el al.1998), Circinus X-1 (Stewart et al. 1993), 1E~1740.7-2942 (Mirabel et al. 1992) and GRS 1758-258 (Rodr\\'{\\i}guez et al. 1992) are included. The typical scale for this kind is 1$\\sim$30 pc, an order of magnitude larger than the second type. The estimated lifetime often exceeds one million years, indicating that they are related to previous outbursts.  From these properties, it is reasonable to further suggest a consistent picture involving all the sources together. We make a conjecture that large scale cavities exist in all microquasar systems. The ``small jets\" observed right after the ejection are just travelling through these cavities. Since there are few or none interactions between the jets and the surrounding gas in this region, the jets travel without obvious deceleration. The emission mechanism is synchrotron radiation by particles accelerated in the initial outburst. The emissions of jets decay very quickly and are not detectable after several tenths of days. In some cases (e.g. XTE J1550-564), the cavity has a dense (compared to the cavity interior) boundary at some radius and the interactions between the jets and the boundary gas heat the particles again and thus make the jets detectable again. Those are the ``large jets\". The emission mechanism then is synchrotron radiation by the re-heated particles in the external shocks. Then, after these interactions, the jets lost most of their kinetic energy into the ISM gradually, causing the latter to expand to large scale structures, the ``large relics\", in a comparatively long time (several millions of years).  The creation of the cavities is not clear at this stage. Possible mechanism may involve supernovae explosions, companion star winds or disk winds. Since some of the sources most likely have never had supernovae before and the winds from the companion stars are not strong enough, the accretion disk winds may be the most plausible possibility. However, these assumptions all require further observations to justify.  Microquasars are powerful probes of both the central engine and their surrounding environment. More studies of their jet behaviors may give us information on the ISM gas properties, as well as the ejecta components. It will provide insights of the jet formation process and offer another approach into black hole physics and accretion flow dynamics."
    },
    "0806/0806.4796.txt": {
        "abstract": "We have analyzed 200 Rossi X-ray Timing Explorer observations of the black hole candidate GX 339--4, all from the bright hard state periods between 1996 and 2005.  Purpose of our study is to investigate the radiation  mechanisms in the hard state of  GX 339--4. The broadband 3--200 keV spectra were successfully modeled by a simple analytic model, power--law with an exponential cut-off modified with a smeared edge. The obtained energy cut-off ($E_{\\rm{cut}}$) was distributed over 50--200 keV, and the photon index over 1.4--1.7.  We found a clear anti-correlation ($E_{\\rm{cut}} \\propto L^{-0.70\\pm0.06}$)    between the X-ray luminosity ($L$) in 2--200 keV and $E_{\\rm{cut}}$ , when $L$  is larger than $7 \\times 10^{37}$ erg s$^{-1}$ (assuming a distance of 8 kpc), while $E_{\\rm{cut}}$ is roughly constant at around 200 keV when $L$ is smaller than $7 \\times 10^{37}$ erg s$^{-1}$. This anti-correlation remained unchanged by adopting a more physical thermal Comptonization model, which resulted in the anti-correlation that can be expressed as $kT_{\\rm{e}} \\propto L^{-0.24\\pm0.06}$. These anti-correlations can be quantitatively explained by a picture in which the energy-flow rate from protons to electrons balances with the inverse Compton cooling. ",
        "introduction": "Black hole X-ray binaries show large varieties in their X-ray properties. Several spectral and variability states have been recognized in the past, with varying names and varying definitions (see e.g.\\ McClintock \\& Remillard 2006, Homan \\& Belloni 2005) and it is believed that they correspond to different accretion geometries. In this paper we focus on the spectral properties of the so-called hard state, which is the state observed at low X-ray luminosities (i.e.\\ below a few percent of the Eddington luminosity, $L_{\\rm Edd}$), although it can also be seen during the rising phase of transient outbursts below $\\sim$0.2$L_{\\rm Edd}$. The accretion geometry in some of the X-ray states is relatively well understood -- e.g.\\ the soft state is thought to be governed by an optically thick and geometrically thin accretion disk -- this is not the case for the hard state. X--ray spectra in the hard state are well represented by a power--law with a photon index of 1.4$\\sim$1.7  (Tanaka \\& Shibazaki 1996) and, in some sources, a high energy cut-off at $\\sim$100 keV (Grove et al.\\ 1998). Several models have been proposed to explain the hard state spectra. Broadband X-ray energy spectra of a large number of black hole binaries in the hard state have been successfully modeled by thermal Comptonization model (Dove et al.\\ 1997, Pountanen \\&  Svensson\\ 1996), although location and geometry of the Comptonizing medium is still debated. Other models include advection dominated accretion flows (ADAF) (Esin et al.\\ 1997), synchrotron and/or synchrotron self-Compton radiation from the base of a jet (Markoff, Nowak, \\& Wilms 2005), and external Compton scattering in the jets (Georganopoulos, Aharonian, and Kirk 2002).  \\begin{figure*}[htbp] \\begin{center} \\FigureFile(61mm,50mm){figure1a.ps} \\hspace{0.0cm} \\FigureFile(57mm,50mm){figure1b.ps} \\hspace{-0.2cm} \\FigureFile(57mm,50mm){figure1c.ps} \\end{center} \\caption{RXTE/ASM light curves (upper panels) and hardness ratios HR2 (= 5--12 keV/3--5 keV; lower panels) of the black hole candidate GX 339--4 in the active periods from 1996 to 2005. The arrows labelled $a$ to $h$ indicate the hard state intervals of which the RXTE/PCA observations were analyzed in this paper.} \\label{lc_1996_2005} \\end{figure*}  GX 339--4 was discovered with the X-ray satellite OSO-7 in 1971 (Markert et al.\\ 1973). Since it was similar to Cygnus X--1 in terms of X-ray spectral and variability properties, it was considered a black hole candidate (BHC, Samimi et al.\\ 1979). A recent measurement of the mass function (5.8$\\pm$0.5 $M_\\odot$, Hynes et al. 2004) strengthens this. We thus assume a mass of 5.8 $M_\\odot$ for this  black hole candidate as a secure lower limit. GX 339--4 is one of the best-studied BHCs in X-rays and gamma-rays; it was observed with Ginga/LAC (Ueda, Ebisawa, and Done 1994), CGRO/OSSE (Grabelsky et al.\\ 1995, Smith et al.\\ 1999), ASCA  (Wilms et al.\\ 1998), RXTE (Smith et al.\\ 1999, Belloni et al. 2005), Beppo-SAX  (Corongiu et al.\\ 2003) and INTEGRAL (Belloni et al.\\ 2006, Joinet et al.\\ 2006). GX 339--4 has shown  various spectral states during its pre-RXTE outbursts (Tanaka \\& Shibazaki 1996), but often stayed in the low-luminosity hard state (low/hard state). All the spectra in the low/hard state can be roughly described by a Comptonization model, requiring the addition of a reflection component, soft excess and iron-K lines. Detailed broadband analysis in the 2--1000 keV range was made using Ginga, RXTE and OSSE data from 1991 and 1996 observations by Zdziarski et al.\\  (1998) and Wardzi$\\acute{n}$ski et al.\\ (2002).  Wardzi$\\acute{n}$ski et al (2002) concluded that the four hard state spectra have very similar intrinsic photon indices of $\\cong$ 1.75. Furthermore, they found that the energy cut-off possibly decreases with increasing the luminosity. Zdziarski et al.\\ (2004) studied the long-term behavior by compiling GINGA/ASM, CGRO/BATSE and RXTE/ASM data (spanning 16 years) and also found that the electron temperature depends on the luminosity by showing a positive correlation between the BATSE flux and the photon index in the 70--160 keV range which reflects the breaking energy.  GX 339--4 has also been extensively studied at other wavelengths. The optical counter part was first identified by Doxsey et al.\\ (1979) as a V$\\sim$18 blue star. However, even during the X-ray off state, optical spectroscopic observations with the VLT revealed no spectral features from the  companion star (Shahbaz et al.\\ 2001). Hence, the type and distance of the secondary star are still unknown. There are two papers which give a lower limit on the distance: Maccarone (2003), 7.6 kpc,  and Hynes et al.\\ (2004), 5.6 kpc. A careful study of the distance is presented by Zdziarski et al.\\ (2004) who argue for a most likely distance of 8 kpc. We assume a distance of 8 kpc in this paper. Multi-wavelength observations have shown that the radio emission of GX 339--4 correlates very tightly with the X-ray fluxes over more than two orders of magnitude, suggesting that the jet plasma may also play a role in high energy band (Corbel et al .\\ 2000, Corbel et al.\\ 2003,  Markoff et al.\\ 2003). Corbel et al.\\ (2003) showed that a significant  fraction of the X-ray flux observed in the low/hard state of black hole  candidates may be due to optically thin synchrotron emission from the  compact jet. Also jet emission have been found in the near-infrared (Corbel \\& Fender 2002) and Homan et al. (2005) discovered a tight relation between near-infrared and X-ray fluxes in the hard state of GX 339--4, similar to that of the X-ray/radio correlation.  Many authors have only studied a small number of observations during individual outbursts of GX 339--4.  In order to improve our understanding of the  radiation mechanisms in the hard state of GX 339--4,  we have performed a systematic study of detailed correlations among spectral parameters, using the large archive of public RXTE data for GX 339--4.  In Section 2 we describe our data analysis, and in Section 3 we present our results, focusing on the correlations between spectral parameters such as luminosity and high energy cut-off or electron temperature. Finally in Section 4, we will discuss the origin of the observed correlations. ",
        "conclusions": "\\subsection{Summary of Results and Comparison with Previous Results}  We have analyzed a large number of RXTE spectra of the black hole candidate GX 339--4 when it was in a bright hard state, and have found a clear luminosity dependence of the high energy cut-off in the X-ray spectrum. The 2--200 keV X-ray luminosities ranged from 1.0 $\\times$ 10$^{37}$ to 2.1 $\\times$ 10$^{38}$ erg s$^{-1}$, covering a factor of 20. This corresponds to 1.3--29\\% of Eddington luminosity for a 5.8 solar mass object (assuming a distance of 8 kpc).  This is the first systematic and quantitive spectral study over a wide-luminosity range in the hard state of GX 339--4. We further notice that the anti-correlation between X-ray luminosity and  energy cut-off was only found above a luminosity of 7$\\times10^{37}$ erg s$^{-1}$.  In GX 339--4, such luminosities were seen mainly during the rising phase  of the 2002/2003,  when the source was in a bright hard state. This state was difficult  to observe prior to RXTE, because typical black hole transients pass through this state within a few weeks of the start of an outburst.  RXTE's fast response and flexible scheduling were critical in observing this state in an increasing number of transient BHCs.  A high energy cut-off around 100 keV has been observed in the energy spectra of many BHCs in their hard state (Grove et al.\\ 1998), although it is not observed in the hard state of every source as recent INTEGRAL observations of GRO J1655--40 show (Caballero-Garcia et al.\\ 2007). For GX 339--4 we reconfirmed that it had a high energy cut-off in  almost all of its hard state observations. It is difficult to make  comparisons with previous results (Zdziarski et al.\\ 1998), because the authors  used different models, but we clearly found a varying energy cut-off  in GX 339--4, with values between 40 keV and 200 keV or more.  High energy cut-off energies above 200 keV are beyond the HEXTE bandpass,  so the highest cut-off energies that we found may have some intrinsic uncertainties since a part of them were not constrained. Other high  sensitive observations with INTEGRAL or Suzaku are needed to confirm cut-off energies above 200 keV in GX 339--4.  \\subsection{Anti-correlation between Luminosity and Temperature}  We found that the high energy cut-off or electron temperature strongly depended on the X-ray luminosity when the luminosity was higher than 7$\\times$10$^{37}$ erg s$^{-1}$, i.e. 0.1 $L_{\\rm Edd}$. A possible anti-correlation in GX 339--4 was already  suggested by Wardzi$\\acute{n}$ski et al.\\ (2002) and Zdziarski et al.\\ (2004).  A simple interpretation of the observed anti-correlation, i.e.  that the hot electrons are efficiently cooled by soft photons via inverse-Compton scattering, is  discussed in more detail in the Section 4.4. Similar anti-correlations have already been suggested in GS 2023+238 (Inoue 1994), GRO J0422+32 (Kurfess 1996) and 10 other BHCs including XTE J1550--564, GS 1354--645, and GX 339--4 (Yamaoka et al.\\ 2005).  Hence, this anti-correlation might be not intrinsic to GX 339--4  but universal among BHCs.   When the GX 339--4 luminosity was lower than 7$\\times$10$^{37}$ erg s$^{-1}$,  the source seemed to show different behavior; the high energy cut-off and electron temperature might have a very weak dependence on the luminosity, remaining a nearly constant at 200 keV and 26 keV, respectively. The reason why the high  energy cutoff reaches constant value is beyond the scope of this paper,  but some other cooling mechanisms such as  bremsstrahlungs and synchrotron emissions may have to be taken into account.  The hard state of Cygnus X--1, which has a luminosity of 0.01--0.03 $L_{\\rm Edd}$, belongs to the latter constant-cutoff regime.   This may be the reason why this source has almost the same cut-off in the hard state at any time (Gierlinski et al.\\ 1997).  %It is important to stress that these %two different types of behavior associated with the X-ray luminosity %were seen in two of the three different outbursts (the first %outburst, from 1996/1997, did not enter the high luminosity range). %As one can see in Figure 5, the spectral parameters of the source %always seems to track similar nearly %identical correlations for each of the three outbursts. Hence, these %correlations are probably determined by the properties of the system %itself.  \\subsection{Cooling Time Scale due to Inverse Compton Scattering}  Based on the spectral fit results, the thermal Comptonization model gives a reasonably acceptable fit. In the following, we show whether the Inverse Compton scattering can be really effective emission process using derived physical parameter. %In the following we show that the %cooling time scale due to inverse Compton scattering  is short in %comparison with the viscous time scale. We consider a spherical hot plasma (like an ADAF for example) with a radius of $R$ ($\\sim 10 R_{\\rm{s}} = 171 (\\frac{M}{5.8M_{\\odot}})$ km; $R_{\\rm{s}}$: Schwarzchild Radius) around a black hole. In this situation, protons will be heated up by the gravitational energy release of the accretion flow through a mechanism such as viscous heating. The protons in the plasma lose their energy due to Coulomb collisions with electrons, while the heated electrons cool due to the Inverse Compton scattering.  We define two time scales: $t_{\\rm pe}$ ($\\propto \\frac{ {(\\frac{kT_{\\rm{p}}}{m_{\\rm{p}}} + \\frac{k T_{\\rm{e}}}{m_{\\rm{e}}})}^{\\frac{3}{2}}}{n}$; $n$ is number density of the plasma, $kT_{\\rm p}$ is proton temperature and $m_{\\rm p}$ is the proton mass.) as the time to establish a Maxwellian velocity distributions for protons interacting with electron field particles (for see details Spitzer 1962), and $t_{\\rm comp}$  ($ = \\frac{n k T_{\\rm{e}}}{(\\frac{dE}{dt})_{\\rm comp}}$) as a cooling time scale of the electron due to inverse Compton scattering. The cooling rate per unit volume $(\\frac{dE}{dt})_{\\rm comp}$ is approximately given as $\\frac{4kT_{\\rm{e}}}{m_{\\rm{e}} c^2}$ $U_{\\rm{rad}} n$ ${\\sigma}_{\\rm T}$ where $c$ is the light speed, $U_{\\rm{rad}}$ is the photon flux density, which is given by $U_{\\rm rad} \\cong \\frac{L \\tau}{\\pi R^2}$,  and $\\sigma_{\\rm T}$ is the cross section for Thomson scattering. These two time scales are given by \\begin{equation} t_{\\rm pe} \\sim 2.5 \\times {10}^{-4} \\frac{{(\\frac{k T_{\\rm{e}}}{30 keV})}^{\\frac{3}{2}} (\\frac{R}{10R_{\\rm{s}}})}{{(\\frac{\\tau}{5})}} [{\\rm sec}], \\end{equation} and \\begin{equation} t_{\\rm comp} \\sim 5.7 \\times {10}^{-6} \\frac{{(\\frac{R}{10R_{\\rm{s}}})}^2}{(\\frac{L}{{10}^{37} erg s^{-1}})(\\frac{\\tau}{5})} [{\\rm sec}], \\end{equation} respectively. On the other hand, $t_{\\rm vis}$($ \\sim \\frac{R^2}{\\alpha H^2 \\Omega}$), the viscous time scale in the ADAF, that is, the time scale in which electrons fall to the central black hole, is \\begin{equation} t_{\\rm vis} \\sim 2.6 \\times {10}^{-3}\\frac{{(\\frac{R}{H})}^2 {(\\frac{R}{10 R_{\\rm{s}}})}^{3/2}}{\\alpha} [{\\rm sec}], \\end{equation} where 2$H$ is the thickness of disk, $\\Omega$ is the angular velocity (= $\\sqrt{\\frac{GM}{R^3}}$), and where we assumed $\\frac{R}{H}\\sim 1$ and $\\alpha \\sim 1$ ($\\alpha$ is viscous parameter) for the hard state.  Clearly, for reasonable parameters, $t_{\\rm pe}$ and $t_{\\rm comp}$ are much smaller than $t_{\\rm vis}$, which means that inverse Compton scattering is an efficient cooling process in this situation.  %{\\bf the viscous time scale that is given, is it for disk accretion or is it for an ADAF? }  \\subsection{Qualitative Explanation for Anti-correlation between Luminosity and Electron Temperature}  Next, we try to explain the anti-correlation between $L$ and $kT_{\\rm e}$ quantitavely in the following simple discussion. The proton temperature ($kT_{\\rm p}$) is assumed to be approximately constant at $\\sim \\frac{GM m_{\\rm p}}{R}$, that is, the energy loss rate of protons is much smaller than viscous heating rate through the accretion. These protons will give their energy to electrons through collisions and the energy loss rate per unit volume is given as $\\frac{\\frac{3}{2}nkT_{\\rm p}}{t_{\\rm{pe}}}$ . %if $T_{\\rm{p}}$ is much larger than $T_{\\rm{e}}$, %where $n$ is the number density of the plasma. If $\\frac{k T_{\\rm{e}}}{m_e c^2}$ is larger than $\\frac{k T_p}{m_p c^2}$,  we will get  \\begin{equation} t_{\\rm pe} \\propto \\frac{T_{\\rm{e}}^{\\frac{3}{2}}}{n} \\end{equation}  In a steady state, the heating rate from proton to electron  should balance with the cooling rate due to inverse Compton scattering, i.e. \\begin{equation} \\frac{\\frac{3}{2} nk T_{\\rm{p}}}{t_{\\rm{pe}}} = \\frac{4k T_{\\rm{e}}}{m_{\\rm{e}} c^2} U_{\\rm{rad}} n {\\sigma}_{\\rm T} \\end{equation} Using these equation (4) and (5), we can work out the anti-correlation of $k T_{\\rm{e}} \\propto L^{-\\frac{2}{5}}$ quantitatively. This is attributed to the fact that the radiation mechanism in the hard state is due to inverse Compton scattering.  Zdziarski (1998) suggested that the high energy cut-off depends on the luminosity from a theoretical point of view of thermal Comptonization by a more precise model. The relations  $k T_{\\rm{e}} \\propto L^{-2/7}$ and $k T_{\\rm{e}} \\propto L^{-1/6}$ were derived for the advection dominated and cooling dominated cases, respectively.  In comparison with this prediction, our result, $k T_{\\rm{e}} \\propto L^{-0.24\\pm0.06}$, is close to  these values. Furthermore, the maximum luminosity in the low/hard state  is predicted as $L_{\\rm max} \\sim 0.15 y^{\\frac{3}{5}} \\alpha^{\\frac{7}{5}} L_{\\rm Edd}$ in the model of Zdziarski (1998). The highest luminosity observed in 2002, $\\sim$0.29$L_{\\rm Edd}$, agrees within a factor of 2 with this value.  In recent years the X-ray spectra of the hard state have also been discussed in the framework of jet outflows. In the most recent version of this model (Markoff, Nowak, \\& Wilms 2005) the hard X-ray spectra are dominated by synchrotron (at the low end) and synchrotron self-Compton (SSC) emission (at the high end) from the magnetized plasma at the base of a jet, which shares many properties with other proposed types of Comptonizing coronae.  In the simplest jet SSC interpretation, the observed decrease of the energy cut-off with luminosity means that either the size of the base of the jet is increasing or that the electron temperature is decreasing as the luminosity goes up. Changes are also expected in the synchrotron cut-off, which should decrease with luminosity, because of a higher cooling rate. There is likely an interplay between the synchrotron and the SSC components that needs to be worked out in more detail before the jet model can be applied to our data.  \\subsection{Difference between bright hard state and common low/hard state}  Differences between the bright hard state (i.e.\\ $\\sim$0.2 $L_{Edd}$) and the more common low/hard state have not been widely discussed in the literature.  From an observational point of view, we found that the spectral properties in the bright hard state were not very different from the typical low/hard state except for one point: the high energy cut-off in the bright hard state is considerable lower ($E_{cut}\\approx$50 keV) than in the low hard state ($E_{cut}\\approx$200 keV). The timing properties are also suggested to be almost the same, with a moderate increase of the frequencies with luminosity (Belloni et al. 2005). A clear correlation between the radio and X-ray flux has been established  in the hard state of GX 339--4 (Corbel et al.\\ 2003), but this relation has only been measured up to a luminosity of $\\sim6\\times10^{37}$ erg s$^{-1}$, i.e. just up to the luminosity above which we see a decrease in the energy cut-off.  The low/hard state has been considered to correspond to an accretion-disk solution of the optically thin ADAF (Esin et al.\\ 1997). Observed luminosity ($\\sim$ 0.2L$_{\\rm Edd}$) in the bright hard state can only be explained by this solution (0.4$\\alpha^2 L_{\\rm Edd}$) if we assume a large viscosity parameter $\\alpha \\sim 1$. It is difficult to estimate an actual value of $\\alpha$, but recent three-dimensional MHD simulations suggest that $\\alpha$ takes a value of reasonably 10$^{-2}$ to 10$^{-1}$ (Hawley \\& Krolik 2001).  Furthermore, considering the presence of hysteresis in the state transition, i.e. the fact that the bright hard state is only in the rise phase of the outburst, the bright hard state may  suggest a presence of another solution like luminous hot accretion flow (LHAF: Yuan 2001) rather than an ADAF.  Further investigations including the radio/X-ray correlations are needed for the bright hard state through multi-wavelength observations.  \\subsection{Constancy of the Photon Index Variations}  Another important result of our systematic study is that the slope of photon index remains almost constant at 1.4--1.7 regardless of the source intensity. This value is typical for the hard state of black hole candidates.  From the point of view of thermal Comptonization, the photon index is correlated with Compton $y$ parameter (= $\\frac{kT_e}{m_{\\rm{e}} c^2}$max$(\\tau,\\tau^2)$; Sunyaev \\& Titartuck 1980). Thus, the obtained $y$ parameter distribution is remarkably close to unity for more than an order of  magnitude as seen in Figure \\ref{L_y}. This is consistent with  power-law shape of the spectrum in the hard state. In the thermal  Comptonization model, the $y$ parameter should always be close to be  unity when there is a large number of soft photons (Shapiro, Lightman,  and Eardly 1976). The constancy of the $y$ parameter means that the ratio of Comptonization luminosity to that in the seed soft photons  is constant, indicating a constant geometry of the corona.  \\begin{figure}[htbp] \\begin{center} \\FigureFile(80mm,80mm){figure8.ps} \\end{center} \\caption{Relation between the luminosity($L$) and $y$ parameter obtained by COMPST model. The different symbols correspond to the epoch defined in Figure \\ref{lc_1996_2005}. } \\label{L_y} \\end{figure}  %\\subsection{Comparison with Neutron Stars and Active Galactic Nucleus}  %Weakly-magnetized neutron stars (NSs) in the hard state  are known to %exhibit spectral shape similar to black hole candidates in the hard %state. Barret et al. (2000) suggested that the high energy cut-off of %NSs is systematically lower than that of black holes, due to addtional %cooling by soft photons from the neutron star surface. However, we %have observed a luminosity variation of high energy  cut-off down to %40 keV in BHC candidate GX 339--4, which is comparable to 58$\\pm$5 keV %(Strickman et al.\\ 1996), 49$\\pm$3 keV (Del Sordo et al.\\ 1999), and %97$\\pm$8 keV (Barret et al.\\ 2000) measured during the low/hard state %of the NS binary GS 1826--238  (they used the same cutoff power-law %model.).  Hence, the ranges in the high energy cut-off of BHCs and NSs %have some  overlap.  In fact, Yamaoka et al.\\ (2005) compared the high %energy cut-off in various BHCs with that of the NS binary GS 1826--238 %in the hard state.  The luminosity of GS 1826--238 when the same high %energy cut-off is observed is lower by a factor of 7--8 than in GX %339--4, which can be explained by  the ratio of the luminosity  to the %Eddignton luminosity (L/L$_{\\rm Edd}$).  In addition, a similar energy %cut-off is frequently seen in the gamma-ray spectrum of the Seyfert %galaxies such as NGC 4151 (Zdziarski et al. 1996) and IC 4329A %(Madejski et al. 1995).  Hence, the same accretion mechanisms  in the %low/hard state might work in the wide mass range of the  central %object, as is also suggested on the basis of variability  properties %(McHardy et al.\\ 2006)."
    },
    "0806/0806.0009_arXiv.txt": {
        "abstract": "{The determination of cluster masses is a complex problem that would be aided by information about the cluster shape and orientation (along the line-of-sight).} {It is in this context, that we have developed a scheme for identifying the intrinsic morphology and inclination of a cluster, by looking for the signature of the true cluster characteristics in the inter-comparison of the different deprojected emissivity profiles (that all project to the same X-ray brightness distribution) and by using SZe data when available.} {We deproject the cluster X-ray surface brightness profile under the assumptions four different ellipsoidal geometry and inclination configurations that correspond to four extreme scenarios, by the non-parametric algorithm DOPING. The formalism is tested with model clusters and then is applied to a sample of 24 clusters. While the shape determination is possible by implementing the X-ray brightness alone, the estimation of the inclination with the line-of-sight is usually markedly improved upon by the usage of SZe data that is available for the considered sample.}{We spot 8 prolate systems, 1 oblate and 15 of the clusters in our sample as triaxial. In fact, for systems identified as triaxial, we are able to discern how the three axis ratios compare with each other. This, when compounded by the information about the line-of-sight extent, allows us to constrain the inclination quite tightly.}{} ",
        "introduction": "\\noindent The identification of the three dimensional cluster shape is of significant importance in the pursuit of quantities of cosmological interest. In particular, the misinterpretation of the true cluster geometry can affect the extraction of the Hubble constant from X-ray and SZe information \\citep{zaroubi01} and of the cluster mass. Erroneous cluster masses are often the cause of less reliable understanding of cluster physics and poorer constraints from cluster cosmology. A mandatory exercise that should be undertaken towards the extraction of the correct cluster masses is the identification of the correct geometry and inclination of the cluster. However, the lack of knowledge about these attributes are typically bypassed by resorting to the assumption of sphericity, even when the projected ellipticity of the system is indicated to be non-zero from observations \\citep{fabian81, yoshikawa99, pizzolato03}. This is also the adopted geometry in frequently used packages for determining dark matter distributions in clusters, like JACO \\citep{jaco}.  On the contrary, halos realised in cosmological simulations have been found to be flattened and triaxial in shape; the distribution of ellipsoidal shapes of simulated halos have been discussed by \\cite{frenk_88, dubinski_91, warren_92, cole_lacey, bailin_steinmetz}, among others. Moreover, the formation of large scale structure has a bearing on the intrinsic shape of cluster sized halos \\citep{plionis04}; thus, the 3D cluster morphology could be a device that can be used to constrain cosmological models.  While the statistical analysis of the 3D cluster morphology has been looked into \\citep[see][and references therein]{jing02,kasun05,hayashi07}, the determination of the full 3-D shapes and orientations of individual halos is a harder problem that has attracted relatively less attention \\citep{sereno06}.  Recovery of the full 3-D morphology of a cluster, using measurements of accessible cluster characteristics (such as X-ray emission profiles or 2-D X-ray surface brightness distributions or SZe data), would require expertise in deprojecting such observed information into the full 3-D distribution, under general triaxial geometries. Such can be achieved via parametric fits; in fact, parametric deprojection of the observed optical surface brightness maps of galaxies, under assumptions of sphericity and (more rarely) axisymmetry, has been studied before \\citep{palmer, bendinelli}. However, the fundamental problems with parametric fits are (i) the answer depends on the choice of the parametrisation and (ii) the goodness-of-fit quantifier (such as a $\\chi^2$ measure) can appear spuriously inflated, particularly in the presence of non-homogeneous measurement noise \\citep{bissantz_munk}.  Thus, non-parametric deprojection is a better option. However, deprojection algorithms that promise improved three dimensional cluster mass distributions, by taking the measured ellipticity and the true cluster morphology and orientation into account, are limited in availability. \\cite{zaroubi01} report on the application of a non-parametric deprojection algorithm that assumes axisymmetry, to a set of simulated clusters. However, as the authors state, the applicability of this scheme to the current state of data sets appears difficult.  A comprehensive approach that calls for the amalgamation of two or more mass indicators include the implementation of X-ray observations, SZe information, lensing results and dynamical measurements. Such exercises have been undertaken already \\citep{zaroubi98,reblinsky00,sand02}.  An attempt has been made in the recent past \\citep{sereno06} to decipher the three dimensional morphology and inclination distribution of a sample of clusters from \\cite{reese02}, with the aim of improving the cluster mass estimation (from X-ray measurements) more accurately. This work employs the rudimentary SZe temperature decrement data \\citep{betty05} in the $\\beta$-modelling of the cluster. Thus, this work is susceptible to very large error bars that currently plague SZe data. This is of course topped by the errors that can be introduced by the choice of ellipsoids of revolution that are implemented to model their sample clusters.  In this work, we present a novel, model-independent trick to determine the correct morphology and inclination of a cluster, without resorting to the assumption of axisymmetry. Such a determination is made possible by analysing the {\\it multiple} deprojected emissivity or X-ray luminosity density distributions that are recovered under distinct assumptions about the cluster geometry and inclination and that will all project to the same observed X-ray brightness map. Thus, the deprojection in question needs to be performed by a scheme that is efficient in carrying out deprojection under general geometries, at any assumed inclination.  This is possible with a new inverse deprojection algorithm DOPING (Deprojection of Observed Photometry using an INverse Gambit) that has been reported in \\cite{doping}.  The {\\it inter-comparison of the amplitudes and shapes of the different deprojected emissivity distributions} tell us the correct shape of the cluster (including triaxial ones) while the inclination constraints are improved upon by the implementation of the available cluster elongation information from the SZe data, as reported in the literature. In case of triaxial systems, the SZe data is used to determine the two intrinsic axial ratios.  The paper is organised as follows. Following the first introductory section, we discuss the deprojection algorithm in brief. Section~3 is devoted to the methodology that we use to extract the cluster shapes from the deprojected profiles. The testing of the advanced scheme is dealt within the following section. We then proceed to apply this method to our cluster sample (Section~5), which precedes the results section. The paper is rounded off by recounting some relevant aspects of the presented work. ",
        "conclusions": "\\label{sec:discussions} \\noindent In this work we have presented a simple but novel technique to extract the intrinsic shape and inclination class of galaxy clusters, using X-ray brightness maps alone. The availability of SZe data is then shown to improve the inclination estimate considerably.  The main motivation behind this exercise is the improvement on mass estimates which is expected to refine the constraints that we can place on cosmological constants, from the analysis of cluster data.  Our suggested formalism relies upon the inter-comparison of the density profiles that are deprojected from the measured X-ray surface brightness maps, under 4 different deprojection scenarios that combine the extrema of the inclination scale and the geometry scale, when clusters are treated as figures of revolution. It is found that each type of observed system leaves its unique signature in this inter-comparison amongst the deprojected X-ray luminosity density profiles. The knowledge of the cluster elongation along the LOS (from SZe data) is then used to supplement our shape and inclination determination; in particular, the recovered constraints on the inclination are significantly improved, if we believe the SZe measurements. In case of triaxial systems, the intrinsic axial ratios can also be tracked, using the SZe data.  \\subsection{Differences with Results of \\cite{sereno06}} \\label{sec:diff} \\noindent Our shape determination agrees with the shape advanced by \\cite{sereno06}, for 13 of the 24 sample clusters. This quantification includes those cases that we identify as triaxial and \\cite{sereno06} find compatible with both the prolate and oblate geometries. For another 10 cases, while we find the system to be definitely triaxial, \\citep{sereno06} report the system to be prolate. This might result from the fact that \\cite{sereno06} use the observed values of projected axial ratio and LOS extent, in equations relating these quantities and the intrinsic axial ratio to inclination $i$, for the two geometries of prolateness and oblateness (see Equations~8-11 in Sereno et. al, 2006).  They constrain their choice of geometry by identifying realistic solutions for inclinations, such as $-1\\leq\\cos(i)\\leq 1$. Thus, if for example, they identify an unrealistic solution for $i$ for the oblate case, they declare the system prolate. But strictly speaking, all that they can conclude is that the cluster at hand is {\\it not oblate}. In other words, it is possible that the cluster is triaxial, in such a way that given the high measurement errors, the equation for $i$ in the prolate case also yields acceptable values. The 10 sample clusters that we spot as triaxial, but \\cite{sereno06} call prolate, are similarly 10 non-oblate clusters.  The one system that we identify as prolate and \\cite{sereno06} call oblate i.e. non-prolate, is MS~1358.4+6245. However, when we use the values of 1.325 for the projected axial ratio (ignoring the $\\pm$1$\\%$ error) and 0.91 for $e_{LOS}$ (using the upper bound from the measured range of 0.72$\\pm$0.19), in Equations~8 and 9, we do actually get a realistic inclination of about 44$^\\circ$, under the assumption of prolateness. In other words, the methodology used by \\cite{sereno06}, does not rule out prolateness for this cluster, in line with our inference.  Thus, our methodology is relatively more powerful, since we have greater resolution ability than simply ruling out prolateness/oblateness. Moreover, our implementation of the SZe data allows for the recovery of a much narrower range of inclinations, for the same measurement as used by previous workers.  \\subsection{X-rays vs. Optical $SB$} \\noindent We would like to emphasise that it is very much possible to utilise the optical surface brightness distribution of clusters, if available, to make an estimate of the morphological structure of the cluster with the aid of DOPING.  This would allow the verification and quantification of possible discrepancies between results obtained using X-ray and optical data. This is particularly interesting, given the study presented by \\cite{gottboler_07} which suggests a relatively more spherical central gas component, compared to the dark matter one.  To this aim, we are now applying the DOPING algorithm to both dark matter halos from the Millennium Simulations and to gas+dark matter $SB$ distribution of the same halos (from the Millennium Gas Simulations). Results will be published in Chakrabarty et. al, (2008). This will answer the question of how the shape determination is affected by changes in the nature of the input.  In this connection, it is further important to mention that our sample of real clusters show more eccentric gas distribution in projection, than the simulated ones studied in the hydrodynamics simulations of \\cite{gottboler_07} and, consequently, the intrinsic axial ratios that we recover, are not akin to $q_1$=1, $q_2=$2, as suggested in that work.  \\subsection{Core Sizes} \\noindent We tested our predictions on a suite of model clusters and found that the recovery of shapes and inclinations is independent of the axial ratios that were assigned to these test systems, as well as the core sizes that these test systems were described by. That the core size is not influential in the recovery of the intrinsic shape of the cluster is not surprising since it is the outer part of the brightness profile that is conventionally considered when identifying shapes of clusters. Thus, whether it is Abell 2261 or MS 1358, our formalism is able to do justice to systems at both ends of core sizes in the sample used herein. This brings us to an important point in regard to the applicability of our technique - it can be implemented to estimate the 3-D geometry of clusters at varying redshifts! In fact, the sample that we used to illustrate the efficacy of our formalism is rather eclectic in its redshift coverage, ($z$=0.023 to 0.784).  Crucially, our work has clearly indicated {\\it the dependence of the recovered core size of the deprojected luminosity density profile, on the deprojection scenario} used. Thus, for Abell~665, the deprojected density corresponding to an assumption of oblateness and $i=i_{min}$ implies a core size that is less than half ($\\approx0.06$ arcmin) that obtained by deprojection performed under prolateness and $i=i_{min}$ and inclination combinations ($\\approx$0.15 arcmin).  It is important to stress here that we {\\it do not need to delete the cores from our analysis} due to any conceived inability of our code to deal with the modelling of the very central regions of the clusters. The code used in the work is in fact non-parametric and its functionality is not challenged by local changes in slopes of the observed X-ray brightness profiles or even by local enhancements of observational errors in the brightness distributions. The deprojected profile will of course bear larger error bars when the input brightness profile manifests the same, than otherwise. The fact that the inclusion of the core does not affect the shape determination is brought home by the tests that we have performed by varying the core sizes of toy clusters.   \\subsection{Central Density} \\noindent One ramification of the observed nature of the luminosity-temperature relationship in clusters is the need to invoke some (unknown) pressure that is effective in reducing the density of the intracluster gas in the cores \\citep{voit02}. However, as we have seen in this work, the X-ray luminosity density at the centres of clusters depends critically on the deprojection scenario adopted in the model. Thus, in Abell~1689, we notice that the density recovered under an assumption of prolateness and $i=i_min$, at about 3 arcmin, is about 10 times less than that recovered under the other three deprojection scenarios. Thus, our work indicates the important contribution that deprojection uncertainties can make towards the observed trends in the cluster self-similar relationships.  \\subsection{Effect of Errors of Cluster Elongation} \\noindent The current state of affairs regarding the quality of SZe data is indeed unsatisfactory, albeit improvements are impending. Thus, when we use the SZe data in this work, we make the conscious decision to work with the value of the cluster elongation that is reported at the centre of the error band. Incorporation of the measurement errors would have greatly reduced the quality of the constraints on the inclinations.  \\subsection{The Case of Abell~1995 and RXJ1347} \\noindent The clusters Abell~1995 and RXJ1347 have been found to be triaxial, of Type~III, which implies that these clusters are nearly prolate except that $c < a$; also, we expect $b-a \\ll a$. The density profiles recovered under the assumption of prolateness are flatter for these clusters than the profiles deprojected under oblateness. This comparison is allowed if the ellipsoidal radius in the prolate cases, in general, exceed that in the oblate cases. This is in turn ensured if $a$ is the photometric semi-axis and if the inclination is small. We find that the inclination is indeed constrained to very small angles, given the reported values of $q_p$ and $q_{LOS}$ (rather the medial value of these quantities within the reported measurement error bands). Thus, these two clusters are deciphered to be nearly face-on systems.  \\subsection{The Case of Abell~370} \\noindent The cluster Abell~370 was reported to bear a ``pronounced triaxial morphology'' by \\cite{betty05_370}, on the basis of the X-ray and SZe data. However, \\cite{sereno06} find the prolate solution to be consistent with the observed X-ray and SZe data for this cluster. In contrast to this, we actually find this cluster to be triaxial, with an inclination in excess of 50$^\\circ$.  \\subsection{The Case of CL~0016+609} \\noindent \\cite{hughes98} suggested a ``reasonable triaxial'' morphology for the cluster CL~0016+609, on the basis of the discrepancy between the Hubble constant values that were estimated from the oblate and prolate models. This suggestion is in line with our conclusion of a triaxial shape for this cluster.   \\subsection{Future Work} The identification of the intrinsic shape and inclination of a cluster is a major step in the characterisation of clusters, both in terms of cluster masses and quantification of the contribution of deprojection uncertainties to the observed scatter in the self-similar relationships of clusters. These applications are planned in future contributions.  In fact, it is envisaged that the improved understanding of cluster spatial configurations leads to better constraints on cluster masses. If this is supplemented with information of mass distributions from dynamical considerations (using an algorithm such as CHASSIS - Chakrabarty $\\&$ Saha, 2001), we could potentially place bounds on the distribution of gas that is at hydrostatic equilibrium. Furthermore, using estimation of the mass within a given projected radius from lensing measurements, would constrain the dynamical mass distribution even better, leading to added improvements in cluster characterisation.  An exercise in the determination of improvement in the dynamical mass of a lensing galaxy, using such lensing constraints, is currently underway (Chakrabarty, et. al, 2008)."
    },
    "0806/0806.2018_arXiv.txt": {
        "abstract": "If a primordial magnetic field (PMF) is present during photon decoupling and afterward, a finite neutrino mass can affect all  modes of the CMB.  In this work, we expand on  earlier studies of the scalar mode effects by constructing  the vector and tensor mode equations in the presence of  massive neutrinos and a PMF. We compute the power spectrum of the various modes in an illustrative example and  find that the neutrino mass can significantly affect the vector and tensor modes  when a PMF exists, while the effects are negligible for no PMF. The most prominent result of the present analysis is the behavior of the EE component of the tensor mode at low multipoles.  For massive neutrinos the EE mode  can become comparable to the observed primary anisotropy.  Therefore, if and when the EE mode power spectrum is measured at low multipoles  the possibility exists  to place a strong  constraint on the sum of the neutrino masses. ",
        "introduction": "The possibility that a primordial magnetic field (PMF) could affect the CMB power spectrum is a subject of considerable recent interest, e.g.~ \\cite{2006ApJ...646..719Y,2004PhRvD..70d3518L,2004PhRvD..70d3011L,2002PhRvD..65l3004M,2007PhRvD..75b3002K}. Such effects arise in two ways.  Baryons are affected by the Lorentz force from a PMF.  They then  influence  the CMB spectrum indirectly through Thomson scattering.  Moreover, the Lorentz force from a PMF also has  a significant  effect on the later development of large-scale structure.   Among other things a PMF could explain  the possible power excess observed by ACBAR \\cite{2004ApJ...600...32K,2008arXiv0801.1491R} for large multipoles ($l>2500$) in the CMB power spectrum. It could also explain the BB mode CMB anisotropies observed by CBI \\cite{2003ApJ...591..540M} as well as  the origin of the  magnetic field inferred from the observed \\cite{2006ApJ...637...19X,2001ApJ...547L.111C,1992ApJ...388...17W,1992ApJ...387..528K} polarization of light from  galactic  clusters. Such studies  suggest that a PMF may indeed exist, and that it is worthwhile to examine other possible consequences of its existence. In previous work \\cite{kojima}, we deduced  neutrino mass effects from the  scalar mode of the CMB power spectrum using the WMAP-3yr data \\cite{2007ApJS..170..377S}, when the PMF effect is also taken into consideration. Indeed, the neutrino mass constraint from the scalar mode of the CMB is the recent focus in the literature \\cite{2007ApJS..170..377S,2005PhRvD..71d3001I,2006PhRvD..74b7302F}. In this work, we expand on this discussion by introducing new studies into the neutrino-mass effects  on the vector and tensor modes of the CMB as well as the scalar mode.   The CMB power spectrum induced by the PMF separates into 3 parts, i.e.~the  scalar, vector, and tensor modes. The scalar mode affects the power spectrum at  large angular scales $l<10$, while the vector mode primarily affects $l>1000$. In previous work \\cite{kojima}, we deduced  neutrino mass constraints from the  scalar mode of the CMB power spectrum using the WMAP-3yr data.  We showed that a  PMF can decrease the upper limit on the neutrino mass. Our ultimate  goal, however,  is to constrain the neutrino mass using all PMF modes and using all available observations including all multipoles. However, before calculating the neutrino mass constraint, we need to derive a framework in which to  study the effects of the neutrino mass on the CMB power spectrum in the presence of a  PMF.  Hence, the purpose of  this work, is to construct the relevant equations for the effects of massive neutrinos on the  vector, and tensor modes in addition to the scalar modes of the CMB power spectrum.   With this framework in place, we then study the possible neutrino mass  effects in the context of an illustrative model.  We show that  the  neutrino mass can cause the EE mode at low multipole  to become comparable to the observed primary anisotropy. We also find that the BB mode is affected by the neutrino mass and may be detectable  if the gravitational wave contribution is not too large. Thus, if and when the EE and BB mode power spectra are measured  the possibility exists  to better constrain  the sum of the neutrino masses. ",
        "conclusions": "In this work, we have expanded on earlier studies of the scalar CMB anisotropies in the presence of a PMF.  In particular we have derived new  vector and tensor mode equations in the presence of  massive neutrinos and a PMF.  We find a large  effect from a finite neutrino mass on the vector and tensor modes  when a PMF exists.  In particular, the effect of  massive neutrinos on the EE mode become comparable to the observed primary anisotropy. Therefore, if and when the polarization power spectrum is ever measured at low multipoles, the possibility may exist  to place a much stronger constraint on the sum of the neutrino masses than presently exists, though the effect of cosmic variance should be carefully taken into consideration."
    },
    "0806/0806.2532_arXiv.txt": {
        "abstract": "{ We present the results of the ``Quasars near Quasars'' (QNQ) survey, a CCD-based slitless spectroscopic survey for faint $V\\la 22$ quasars at $1.7\\la z\\la 3.6$ on 18 $26\\farcm 2\\times 33\\farcm 5$ fields centred on bright quasars at $2.76<z<4.69$. In total 169 quasar candidates with emission lines were selected from the extracted flux-calibrated spectra on the basis of well-defined automatic selection criteria followed by visual inspection and verification. With follow-up spectroscopy of 81 candidates that were likely to reside at $z>1.7$ we were able to confirm 80 new quasars at $0.580\\le z\\le 3.586$ on 16 of our fields. 64 of the newly discovered quasars are located at $z>1.7$. The overall high success rate implies that most of the remaining 88 candidates are quasars as well, although the majority of them likely resides at $z<1.7$ on the basis of the observed line shapes and strengths. Due to the insufficient depth of the input source catalogues needed for extraction of the slitless spectra our survey is not well defined in terms of limiting magnitude for faint $2.5\\la z\\la 3.6$ quasars whose Ly$\\alpha$ emission is detectable well beyond $V=22$, albeit at a continuum $S/N\\la 1$. While not useful for characterising the evolving space density of quasars, our sample provides many new closely spaced quasar sightlines around intensely studied quasars for further investigations on the three-dimensional distribution of the intergalactic medium. ",
        "introduction": "Recent large optical surveys, such as the Sloan Digital Sky Survey (SDSS) and the 2dF QSO Redshift Survey (2QZ), have revealed thousands of previously unknown quasars selected on the basis of their broadband optical colours \\citep{schneider07,croom04}.  Colour selection is efficient if quasar candidates are well separated from normal stars in multidimensional colour space, most notably at $z\\la 2.2$ (UV excess) and at $z\\ga 3.5$. However, even optical multicolour surveys are systematically incomplete at $2.5\\la z\\la 3.5$, where the colours of quasars and stars are similar \\citep[e.g.][]{warren91a,richards02b}. Incompleteness in this redshift range can be significantly reduced by a better tracing of the spectral energy distributions with additional filters, e.g.\\ by incorporating mediumband filters as in the COMBO-17 survey \\citep{wolf03}.  Alternatively, slitless spectroscopy is a particularly efficient way to find quasars at redshifts $z\\ga 2$ because of the prominent Ly$\\alpha$ emission line redshifted into the optical wavelength regime. Early surveys generated hundreds of quasars by visual scanning of objective-prism photographic plates for emission-line objects \\citep[e.g.][]{osmer80b,crampton85}. Subjecting such plates to digitisation with fast measuring machines made it possible to employ the automated selection of quasar candidates \\citep{clowes84,hewett85} and to build substantial quasar samples at redshifts $0 \\la z \\la 3.2$ \\citep{hewett95,wisotzki00}. The systematic CCD-based slitless survey for $2.7\\la z\\la 4.8$ quasars by \\citet{schneider94} was among the first to quantify the declining space density of high-redshift quasars \\citep{schmidt95}.  Apparent pairs or close groups of high-redshift quasars are very attractive targets to study the three-dimensional distribution of the intergalactic medium (IGM). But high-resolution studies have so far been limited to a small number of suitable groups of bright quasars \\citep[e.g.][]{dodorico02,dodorico06}. Going fainter than $V\\sim 19$ immediately limits the achievable spectral resolution and S/N. A possible compromise lies in combining high-resolution spectra of bright quasars with lower resolution, lower S/N data of fainter quasars in the surroundings. \\citet{pichon01} argued that one can significantly improve the recovery of the 3-dimensional topology of the IGM this way.  Although most known quasars were colour-selected from either SDSS or 2QZ, these surveys produced relatively few useful quasar groups because of the reduced selection efficiency at $z\\ga 2.5$ combined with a bright magnitude limit. In fact, most well-studied close groups of quasars at $z>2$ were found by slitless spectroscopy. Follow-up spectroscopy of candidates by \\citet{bohuski79} revealed 13 $z>1.5$ quasars on $2.1$~deg$^2$ \\citep{jakobsen92} with two showing correlated complex intergalactic \\ion{C}{iv} absorption at $1.48<z<2.15$, indicative of an elongated superstructure extending over $17\\farcm 9$ on the sky \\citep[e.g.][]{jakobsen86,dinshaw96}. Another group discovered by \\citet{sramek78} contains now 6 QSOs at $2.49\\le z\\le 3.45$ within a radius of 20\\arcmin\\ around Q~1623$+$2653, all of which are bright enough for correlation studies of the IGM \\citep[e.g.][]{crotts98}. \\citet{williger96} reported 25 quasars at $1.5\\la z\\la 3.4$ within a $\\sim 1$~deg$^2$ region and used these to reveal large-scale structure in the IGM \\citep{williger00,liske00b}. At lower redshifts, slitless surveys revealed large associations of quasars at similar redshift \\citep{crampton90,clowes99}.  Groups of quasars, or more generally speaking, active galactic nuclei, have also been discovered by recent deep multi-wavelength surveys in selected fields, such as in the Chandra Deep Fields North \\citep{barger03,cowie04} and South \\citep{szokoly04,wolf04}, the Marano field \\citep{zitelli92,krumpe07}, or the COSMOS field \\citep{prescott06,trump07}. However, the majority of the AGN thus found are too faint for follow-up studies at a spectral resolution allowing for meaningful IGM studies.  Complementary to slitless spectroscopy, quasar pairs can be found from similarities in multi-colour space. Using this approach, \\citet{hennawi06b} confirmed 40 new associated and 73 projected quasar pairs with separations $\\vartheta<1\\arcmin$ from a sample of faint $i\\la 21$ quasar pair candidates selected via SDSS photometry. Follow-up spectra revealed transverse clustering of optically thick absorption systems near foreground quasars \\citep{hennawi06a,hennawi07}.  Targeted deep surveys of sky regions around well-studied high-redshift quasars are rare. An exception are many of the fields selected for the Lyman-break galaxy survey by \\citet{steidel03} which were centred on bright quasars to correlate the galaxies with the intergalactic absorption along the sightline \\citep{adelberger05}. Here we describe a systematic search for apparent quasar groups at $1.7\\la z\\la 3.6$ in the southern hemisphere, targeting fields centred on known bright $z>2.7$ quasars that had been observed at high resolution with the UV-Visual Echelle Spectrograph (UVES) at the VLT.  We already reported results for two special fields \\citep[][hereafter Papers~I and II, respectively]{worseck06,worseck07}. The present paper is devoted to present the entire survey. In Sect.~\\ref{qnq_surveyobs} we describe the slitless spectroscopic survey observations. Section~\\ref{qnq_surveyred} outlines the automatic reduction pipeline developed for these data. Section~\\ref{qnq_candselection} describes the semi-automatic selection of quasar candidates. We report on the follow-up slit spectroscopy of candidates in Sect.~\\ref{qnq_followup}, followed by a brief discussion of the properties of confirmed quasars and the remaining candidates (Sect.~\\ref{qnq_results}). In Sect.~\\ref{qnq_discussion} we present the resulting quasar groups and discuss the efficiency and completeness of our survey.  We conclude in Sect.~\\ref{qnq_conclusions}.   \\begin{figure} \\centering \\includegraphics[scale=0.80]{0157f1.eps} \\caption{\\label{qnq_wfidither} Schematic view of the WFI dither pattern. The full lines denote the edges of the $25\\farcm 1\\times 32\\farcm 7$ field of view in the dithered 600~s exposures. Dashed lines mark the 6 2k$\\times$4k chips with the inter-chip gaps. After rotating the instrument by $10\\degr$ about its origin in the focal plane (cross) three further 600~s exposures were taken.} \\end{figure}  \\begin{table*} \\caption{Observing log of the slitless survey observations. The columns list the quasar at the field centre with redshift and celestial coordinates, the night of observation, the employed filters, instrument rotations, total exposure time per band and the seeing.}  \\label{qnq_wfiobservinglog} \\centering\\scriptsize \\begin{tabular}[t]{lcccccccc}\\hline\\hline\\noalign{\\smallskip} Field\t\t&$z$\t&$\\alpha$~(J2000) \t\t&$\\delta$~(J2000) \t\t&Night\t\t&Filters&Rotations~[\\degr]&Exposure~[s]&Seeing~[\\arcsec]\\\\\\noalign{\\smallskip}\\hline\\noalign{\\smallskip} \\object{Q~0000$-$263}\t&$4.125$&00$^{\\mathrm{h}}$03$^{\\mathrm{m}}$22\\fs91\t&$-$26\\degr03\\arcmin16\\farcs8\t&03~Oct~2002\t&$B$, $V$&0\t&1800\t&0.7--1.0\\\\ &\t&\t\t\t\t\t\t&\t\t\t\t\t&04~Oct~2002\t&$B$, $V$&10\t&1800\t&0.9--1.4\\\\ \\object{Q~0002$-$422}\t&$2.767$&00$^{\\mathrm{h}}$04$^{\\mathrm{m}}$48\\fs11\t&$-$41\\degr57\\arcmin28\\farcs8\t&01~Oct~2002\t&$B$, $V$&0, 10\t\t&3600\t&1.3--1.6\\\\ \\object{Q~0055$-$269}\t&$3.665$&00$^{\\mathrm{h}}$57$^{\\mathrm{m}}$57\\fs92\t&$-$26\\degr43\\arcmin14\\farcs2\t&02~Oct~2002\t&$B$, $V$&0, 10\t\t&3600\t&0.9--1.5\\\\ \\object{Q~0302$-$003}\t&$3.285$&03$^{\\mathrm{h}}$04$^{\\mathrm{m}}$49\\fs86\t&$-$00\\degr08\\arcmin13\\farcs4\t&03~Oct~2002\t&$B$, $V$&0\t&1800\t&0.9\\\\ &\t&\t\t\t\t\t\t&\t\t\t\t\t&03~Oct~2002\t&$B$\t&10\t&1800\t&0.9\\\\ &\t&\t\t\t\t\t\t&\t\t\t\t\t&04~Oct~2002\t&$V$\t&10\t&1800\t&0.9--1.1\\\\ \\object{Q~0347$-$383}\t&$3.220$&03$^{\\mathrm{h}}$49$^{\\mathrm{m}}$43\\fs68\t&$-$38\\degr10\\arcmin31\\farcs3\t&27~Feb~2003\t&$B$, $V$&10\t&1800\t&1.0\\\\ &\t&\t\t\t\t\t\t&\t\t\t\t\t&27~Feb~2003\t&$V$\t&0\t&1800\t&1.0\\\\ &\t&\t\t\t\t\t\t&\t\t\t\t\t&28~Feb~2003\t&$B$\t&0\t&1800\t&1.9\\\\ \\object{CTQ~0247}\t&$3.025$&04$^{\\mathrm{h}}$07$^{\\mathrm{m}}$17\\fs99\t&$-$44\\degr10\\arcmin13\\farcs4\t&30~Sep~2002\t&$B$, $V$&0\t&1800\t&1.3--1.8\\\\ &\t&\t\t\t\t\t\t&\t\t\t\t\t&01~Oct~2002\t&$B$, $V$&10\t&1800\t&1.4--1.8\\\\ \\object{Q~0420$-$388}\t&$3.120$&04$^{\\mathrm{h}}$22$^{\\mathrm{m}}$14\\fs81\t&$-$38\\degr44\\arcmin52\\farcs9\t&26~Feb~2003\t&$B$, $V$&0, 10\t&3600\t&0.8\\\\ \\object{PKS~0528$-$250} &$2.813$&05$^{\\mathrm{h}}$30$^{\\mathrm{m}}$07\\fs96\t&$-$25\\degr03\\arcmin29\\farcs9\t&03~Nov~2002\t&$V$\t&10\t&1800\t&1.1\\\\ &\t&\t\t\t\t\t\t&\t\t\t\t\t&04~Nov~2002\t&$V$\t&0\t&1800\t&1.2\\\\ &\t&\t\t\t\t\t\t&\t\t\t\t\t&28~Feb~2003\t&$B$\t&0, 10\t&3600\t&1.2\\\\ \\object{HE~0940$-$1050} &$3.088$&09$^{\\mathrm{h}}$42$^{\\mathrm{m}}$53\\fs40\t&$-$11\\degr04\\arcmin25\\farcs0\t&26~Feb~2003\t&$B$, $V$&0, 10\t&3600\t&0.8\\\\ \\object{CTQ~0460}\t&$3.139$&10$^{\\mathrm{h}}$39$^{\\mathrm{m}}$09\\fs51\t&$-$23\\degr13\\arcmin25\\farcs7\t&27~Feb~2003\t&$B$, $V$&0, 10\t&3600\t&1.0--1.5\\\\ \\object{BR~1117$-$1329}\t&$3.958$&11$^{\\mathrm{h}}$20$^{\\mathrm{m}}$10\\fs30\t&$-$13\\degr46\\arcmin25\\farcs0\t&28~Feb~2003\t&$B$, $V$&0, 10\t&3600\t&$>2$\\\\ \\object{BR~1202$-$0725}\t&$4.690$&12$^{\\mathrm{h}}$05$^{\\mathrm{m}}$23\\fs12\t&$-$07\\degr42\\arcmin32\\farcs5\t&26~Feb~2003\t&$B$, $V$&0\t&1800\t&0.8\\\\ &\t&\t\t\t\t\t\t&\t\t\t\t\t&27~Feb~2003\t&$V$\t&10\t&1800\t&1.6\\\\ \\object{Q~1209$+$093}\t&$3.291$&12$^{\\mathrm{h}}$11$^{\\mathrm{m}}$34\\fs95\t&$+$09\\degr02\\arcmin20\\farcs9\t&27~Feb~2003\t&$B$, $V$&0, 10\t&3600\t&1.6\\\\ \\object{Q~1451$+$123}\t&$3.246$&14$^{\\mathrm{h}}$54$^{\\mathrm{m}}$18\\fs61\t&$+$12\\degr10\\arcmin54\\farcs8\t&28~Feb~2003\t&$B$, $V$&0, 10\t&3600\t&0.8--2.0\\\\ \\object{PKS~2126$-$158}\t&$3.285$&21$^{\\mathrm{h}}$29$^{\\mathrm{m}}$12\\fs18\t&$-$15\\degr38\\arcmin41\\farcs0\t&30~Sep~2002\t&$B$, $V$&0, 10\t&3600\t&0.8--1.0\\\\ \\object{Q~2139$-$4434}\t&$3.214$&21$^{\\mathrm{h}}$42$^{\\mathrm{m}}$25\\fs81\t&$-$44\\degr20\\arcmin17\\farcs2\t&03~Oct~2002\t&$B$, $V$&0, 10\t&3600\t&1.0--1.5\\\\ \\object{HE~2243$-$6031}\t&$3.010$&22$^{\\mathrm{h}}$47$^{\\mathrm{m}}$09\\fs10\t&$-$60\\degr15\\arcmin45\\farcs0\t&02~Oct~2002\t&$B$, $V$&0, 10\t&3600\t&1.0--1.7\\\\ \\object{HE~2347$-$4342}\t&$2.885$&23$^{\\mathrm{h}}$50$^{\\mathrm{m}}$34\\fs21\t&$-$43\\degr25\\arcmin59\\farcs6\t&04~Oct~2002\t&$B$, $V$&0, 10\t&3600\t&0.8--1.2\\\\ \\noalign{\\smallskip}\\hline \\end{tabular} \\end{table*} ",
        "conclusions": "\\label{qnq_conclusions}  We performed the ``Quasars near Quasars'' (QNQ) survey, a CCD-based slitless spectroscopic survey for faint $V\\la 22$ quasars at $1.7\\la z\\la 3.6$ around 18 well-studied bright quasars at $2.76<z<4.69$, covering a total area of $\\simeq 4.39$~deg$^2$. In order to analyse the slitless data we developed a data reduction pipeline that performs an optimal extraction of the spectra and that is able to cope with contaminating spectral orders of slitless grism data. From the $\\sim 29000$ extracted flux-calibrated spectra we selected 169 previously unknown quasar candidates and $>100$ likely low-redshift emission line galaxy candidates on the basis of emission features that fall in the covered wavelength range 4200~\\AA$\\le\\lambda\\le$5800~\\AA. A semi-automatic selection routine limited potential biases of purely visual selection and allowed to quantify selection effects.  Follow-up spectroscopy confirmed 80 out of 81 selected quasar candidates on 16 fields. 64 of these newly established quasars reside at $z>1.7$. The highest redshift quasar is QNQ~J22484$-$6002 at $z=3.586$. The brightest newly discovered high-redshift quasar is QNQ~J11197$-$1340 ($z=2.220$, $B=18.3$). Given the high success rate of the follow-up, the vast majority of the remaining 88 candidates will be quasars as well, although most of them likely reside at lower redshifts.  The primary aim of this survey was to provide new groups of quasars for medium-resolution spectroscopy in the vicinity of well-studied known quasars. Originally, we had not expected to reach as faint as $V\\sim 22$ in the survey observations. At these faint magnitudes our survey is not well defined because of photometric incompleteness in the source catalogues needed for automatic extraction of the slitless spectra. Together with the fact that not all selected quasar candidates could be included in the spectroscopic follow-up, our survey probably is of limited use for constraining quasar evolution. However, focusing on survey efficiency rather than on completeness at the faint end was justified in order to accomplish the main goal of this survey. In fact, the faintest quasars discovered will for now remain beyond the limits for obtaining high-quality spectra. But at least the $V\\la 21$ quasars are well suited for follow-up studies with current (e.g.\\ ESI at Keck) or upcoming (e.g.\\ X-shooter at VLT) high-throughput spectrographs at 8--10~m-class telescopes.  Together with the central quasars in the fields already observed at high resolution, these quasar groups can be used to perform a tomography of the intergalactic medium. Large-scale clustering of the Ly$\\alpha$ forest or correlations of metal line systems can be investigated as well \\citep{williger00,dodorico02,dodorico06}. Some of the discovered quasars reside at similar redshifts or approximately at the same redshift of the central quasar in the field, giving potentially insights to quasar clustering along overdense filaments in the plane of the sky. We have identified two quasars coinciding with a damped Ly$\\alpha$ absorber on one central line of sight, as well as a large-scale group of quasars at $z=2.70$--$2.78$. Moreover, our study provides new foreground quasars to investigate the transverse proximity effect of quasars (Papers~I and II). In combination with the already available high-resolution spectra of the central quasars, medium-resolution spectra of these quasar groups will offer great opportunities to study the large-scale cosmic web in three dimensions."
    },
    "0806/0806.0471_arXiv.txt": {
        "abstract": "We have modelled 38 barred galaxies by using near-IR and optical data from the Ohio State University Bright Spiral Galaxy Survey. We constructed the gravitational potentials of the galaxies from $H$-band photometry, assuming constant mass-to-light ratio. The halo component we chose corresponds to the so called universal rotation curve. In each case, we used the response of gaseous and stellar particle disc to rigidly rotating potential to determine the pattern speed.  We find that the pattern speed of the bar depends roughly on the morphological type. The average value of corotation resonance radius to bar radius, $\\mathcal{R}$, increases from $1.15 \\pm 0.25$ in types SB0/a -- SBab to $1.44 \\pm 0.29$ in SBb and $1.82\\pm 0.63$ in SBbc -- SBc. Within the error estimates for the pattern speed and bar radius, all galaxies of type SBab or earlier have a fast bar ($\\mathcal{R} \\le 1.4$), whereas the bars in later type galaxies include both fast and slow rotators. Of 16 later type galaxies with a nominal value of $\\mathcal{R} > 1.4$, there are five cases, where the fast rotating bar is ruled out by the adopted error estimates.  We also study the correlation between the parameter $\\mathcal{R}$ and other galactic properties. The clearest correlation is with the bar size: the slowest bars are also the shortest bars when compared to the galaxy size. A weaker correlation is seen with bar strength in a sense that slow bars tend to be weaker. These correlations leave room for a possibility that the determined pattern speed in many galaxies corresponds actually that of the spiral, which rotates more slowly than the bar. No clear correlation is seen with either the galaxy luminosity or colour. ",
        "introduction": "\\label{introduction}  Studies in near-IR, where the extinction is lower than in visual wavelengths and majority of light comes from the old stellar population, have shown that about 60-70\\% of all spiral galaxies have a large scale stellar bar \\citep{eskridge2000,whyte2002,laurikainen2004,menendez2007,marinova2007}. According to recent analysis of over 2000 spiral galaxies \\citep{sheth2008}, the bar fraction decreases from about 65\\% in the local universe to about 20\\% at redshift $z=0.84$ \\citep[see also][]{abraham99,elmegreen2004,jogee2004,menendez2007}. Anyhow, bars are so common that they are either very robust or they represent a recurrent phenomenon in the life of a spiral galaxy \\citep{bournaud2005}. In contrast with some earlier studies \\citep{thompson81,elmegreen90c}, \\citet{hernandeztoledo2007} claim the bar frequency is roughly the same in different environments -- modest interactions do not seem to play a major role in bar formation and destruction.  Bars can be roughly divided into two classes based on their light distribution: flat and exponential bars \\citep{elmegreen85,elmegreen96b}. In the first type, which is more typical to early type barred galaxies, the radial surface brightness profile along the bar major axis is flatter than in the surrounding disc, whereas in the exponential bars the profile is quite similar to the surrounding disc. Furthermore, the flat bars can display twisting of isophotes \\citep{elmegreen96a}. Bar morphology can also be either classical or of ansae-type, characterized by blops at the both ends of the bar \\citep{laurikainen2007,martinez2007}.  Another approach to characterize a bar is its strength. Some attempts have based on the ellipticity of the deprojected bar \\citep[e.g. ][]{martin95b,whyte2002,laurikainen2002b}. Recently, there have been attempts to estimate the actual gravitational perturbation of the bar by using near-IR photometry \\citep{buta2001b,laurikainen2002b,laurikainen2004,laurikainen2005,buta2005}. There is also another bar strength estimate, namely the $A_2$-Fourier amplitudes of density \\citep{laurikainen2004,laurikainen2005}, which is an approximation of the relative mass of the bar.  All these bar strength estimates are discussed with respect to the Hubble sequence by \\citet{laurikainen2007}.  Perhaps the most important parameter defining a bar is its pattern speed, $\\Omega_{bar}$, or how fast the bar rotates. In principle, this determines how far the orbits of stars and gas clouds are affected by the bar. The pattern speed has a physical upper limit -- a bar cannot reach beyond its corotation resonance (CR) radius $R_{CR}$, i.e.\\ the region in the disc where the angular speed of circular rotation equals the bar pattern speed. This limitation is based on the studies of stellar orbits in barred potentials -- the orientation of the major axes of closed orbits becomes perpendicular with the bar beyond $R_{CR}$, thus the orbits in the outer disc are not able to support the bar \\citep{contopoulos80a}. On the other hand, there is no evident lower limit based on stellar orbits for the bar pattern speed.  In the literature an often used nomenclature is based on the value of a dimensionless parameter $\\mathcal{R} = R_{CR} / R_{bar}$, where $R_{bar}$ is the semi major axis of the bar. The cases where $\\mathcal{R} \\le 1.4$ are usually called ``fast bars'', whereas those with a larger ratio are ``slow bars'' \\citep{debattista2000}. This is shown as a schematic drawing in Fig.~\\ref{barschema}. Considerable effort has been devoted to determine $\\mathcal{R}$ for individual galaxies and to study if it depends on other properties of the bar itself or the other galactic components.  \\begin{figure} \\resizebox{\\hsize}{!}{\\includegraphics{fig1.eps}} \\caption{A schematic drawing of fast and slow bars.} \\label{barschema} \\end{figure}  One way to study the evolution of galaxies is to make self-consistent N-body simulations, where the original particle distribution forms different kinds of structures depending on its initial stage. If a bar forms by a global bar instability in these simulations, it tends to be fast rotating \\citep{sellwood81}. On the other hand, it has been suggested that if a bar forms in a galaxy due to an interaction with other galaxies, it can be a slow rotator, perhaps extending only near its inner Lindblad resonance \\citep{miwa98}. This may also be the case with gradual bar growth \\citep{lynden79}, or when the initial bulge-to-disc mass ratio is low \\citep{combes93}. The pattern speed of the bar does not depend only on its initial value -- the bar interacts with the other components of the galaxy. Secular evolution due to interaction between the bar and the outer disc decelerates the bar rotation, but this seems to be compensated by the growth of the bar length by capturing more particles \\citep{sellwood81}. The situation is different with the interaction between the bar and the dark halo: if the halo density is high in the inner parts of the galaxy, the resonant interaction can decelerate an initially fast rotating bar so much that it becomes a slow rotator \\citep{debattista98,athanassoula2003}. Thus, the observed bar pattern speed could be considered as a rough estimator for the halo contribution.  Much weight has been given to a model-independent kinematic method, the so-called Tremaine--Weinberg method, hereafter TW \\citep{tremaine84}. It is based on photometric and spectroscopic observations along slits parallel to the galaxy's major axis and assumes a rigidly rotating component, which satisfies the continuity equation. TW-method has been used to determine pattern speeds of about 20 large scale stellar bars. The results range from probably unphysical $\\mathcal{R}=0.6$ to slow bar regime with $\\mathcal{R}=1.8$. This has been considered to be in accordance with fast bars, at least when the quite large error estimates \\citep[e.g.\\ due to sensitivity to errors in the position angle of the disc, ][]{debattista2003} are taken into account \\citep*{kent87,merrifield95,gerssen99,gerssen2003,aguerri2003,debattista2004,corsini2007,treuthardt2007}. This apparent lack of slowly rotating bars has been interpreted to indicate that the density of a dark halo is small in the inner parts of barred galaxies \\citep{debattista2000,aguerri2003}, which is in disagreement with some cosmological simulations of structure formation, typically producing centrally peaked dark matter component \\citep{navarro96}. However, due to limitations set by the initial assumptions, measurements using stellar absorption line spectrometry have been made almost exclusively for SB0-galaxies, which are almost free of gas and dust.  Although some attempts \\citep{zimmer2004,rand2004} have been made to use TW-method with CO observations, the pattern speeds of late type barred galaxies have been usually determined by various indirect methods. Many of these rely on morphological features in barred galaxies such as spiral arms, rings or leading offset dust lanes. In a classic scenario, a two-armed spiral starts from the ends of the bar, but exceptions to this are not rare: there can be an ``empty'' region between the bar and the spiral (e.g.\\ NGC 210) or an offset between the bar position angle and the starting points of the spiral arms (e.g.\\ NGC 799). There can also be more than two arms either in the whole scale of spiral structure (e.g. ESO 566-24) or as a multiply armed outer disc (e.g.\\ NGC 4303). The clearest correlation with the bar properties is observed in the sizes and shapes of rings \\citep[for a detailed review, see ][]{buta96b}: outer rings whose semi major axis is about twice the bar radius and inner rings whose major axis is usually the same as the major axis of the bar. There is also a third ring type, the nuclear rings, whose radius is roughly 1/10 of the bar radius, but with a large scatter. Leading offset dust lanes are straight or curved dust features inside the bar and are located in its leading side.  Models of individual galaxies have been constructed either by fitting analytical potential components to observations or by determining the potential from photometry, making some assumption about mass-to-luminosity ratio, orientation and internal geometry of galaxies. In principle, modelling is more economical than using TW-method - the pattern speed can be estimated by comparing the simulated and observed morphology. Naturally, kinematical observations can help to determine the acceptable parameter range. Modelling has produced both high and low pattern speed estimates for the bars \\citep{hunter88,sempere95a,lindblad96b,salo99,rautiainen2004}. The given error estimates of models are typically smaller than with TW-method.  Pattern speed estimates have also been done by identifying various morphological or photometric features with resonances. Inner and outer rings are usually considered to form by gravitational torque of the rotating bar, which causes gas to flow in the radial direction. The net torque vanishes in major resonances, where gas then accumulates \\citep{buta96b}. Comparison with results of gas dynamical simulations and analysis of orbits in barred galaxies have led to the following identifications: an outer ring should be located near the Outer Lindblad resonance (OLR) and the inner ring near the inner 4/1-resonance, which is located inside $R_{CR}$ \\citep{schwarz81,schwarz84b,byrd94,salo99,rautiainen2000,patsis2003}. When the rotation curve is known, these resonance identifications can be used to determine the pattern speed \\citep{buta98a}. Another approach is based on two-colour photometry \\citep{puerari97}: the location of shock induced star forming regions should change side related to spiral arms when crossing $R_{CR}$. This crossing should be detectable from two-colour photometry. Especially, the results of \\citet{aguerri98} gave a hint (the total sample consisted of 10 galaxies) that late type spirals could be slower rotators than early types.  In addition there are at least two other methods worth mentioning. A kinematical method to determine the corotation radius, which is based on the residual patterns in the velocity field after removal of circular velocities, was suggested by \\citet{canzian93}. However, it has been rarely applied. Recently, \\citet{zhang2007} suggested that the calculated phase shifts between potential and density could be used to determine $R_{CR}$. The validity of this approach is still questionable: for several galaxies it found corotation radii well inside the bar, which is in disagreement with the analysis of orbits in barred potentials.  Using the previously published pattern speed estimates in studying a possible dependency on the morphological type is problematic. Very different methods have been used, the data is of uneven quality or the definition of the bar differs. There are several papers where quite successful simulation models are presented, but only one pattern speed has been tried (e.g.\\ by assuming $\\mathcal{R}=1$) -- this cannot be taken as a {\\it determination} of the pattern speed.  Large galaxy surveys and the increase of computing power makes it possible to improve the situation by mass production of galaxy models: here we present a simulation series where we estimate the pattern speeds of 38 moderately inclined barred galaxies, using data from the Ohio State University Bright Spiral Galaxy Survey \\citep[][hereafter OSUBSGS]{eskridge2002}. To our knowledge, this is the largest sample of barred spiral galaxies whose pattern speeds are determined with a consistent method. Furthermore, the morphological types of these galaxies range from SB0/a to SBc, based on Third Reference Catalogue of Bright Galaxies \\citep[][hereafter RC3]{devaucouleurs91}, so that each morphological type is represented by several galaxies. The initial modelling results for this sample were published in \\citet[][ hereafter RSL2005]{rautiainen2005}, and part of the analysis in this paper was also presented in \\citet{salo2007}.  Whereas we find all the bars of early type galaxies (SBab or earlier) of the sample are fast rotators, our models also suggest that galaxies of later morphological types include both fast and slow rotating bars. The slow rotation seems to be related to the small size of the bar. An alternative interpretation to this is that the spiral arms in these galaxies have a lower pattern speed than the bar, a situation often seen in N-body models \\citep{sellwood88,rautiainen99}. ",
        "conclusions": "We have modelled 38 barred galaxies with simulations using collisionless test particles and inelastatically colliding test particles. The gravitational potentials used were derived from $H$-band images from Ohio State University Bright Spiral Galaxy Survey.  Our findings are as follows:  1) the average value of the dimensionless pattern speed, $\\mathcal{R}=R_{CR}/R_{bar}$ depends on the morphological type, being about 1.2 for the early type barred galaxies (SB0/a--SBab), 1.4 to intermediate type (SBb) and 1.8 to late type (SBbc--SBc).  2) When error estimates are considered, all early type spiral galaxies in the sample ($T \\le 2$) are in accordance with having fast bars ($\\mathcal{R} < 1.4$).  3) If the derived pattern speed corresponds to that of the bar, then intermediate and late type galaxies include both fast and slow bars. In five cases the bar remains in slow bar domain even if a ``worst case scenario'' allowed by the error estimates in pattern speed and bar size determinations is considered. However, the existence of multiple pattern speeds, i.e.\\ spiral rotating more slowly than the bar, is a viable alternative in some cases, especially for NGC 4303.  4) Slow bars are short when compared to isophotal radius $R_{25}$. Especially, with only one exception, galaxies with $R_{bar}/R_{25} \\le 0.25$ have a slow bar.  5) With one exception, the slow bars have bar strength $Q_B \\le 0.3$, although several bars with similar strength are also found to be fast.  6) There is no clear correlation with the absolute $B$- or $K$-magnitudes of the galaxies or the $B-K$ color index.  7) Galaxies with fast and slow bars seem to occupy different regions when plotted in a two-dimensional coordinate system defined by bulge-to-total flux ratio and absolute $K$-magnitude.  8) Omitting possible systematic errors, the error estimates of the model-based pattern speeds are typically smaller than with Tremaine-Weinberg method. If the results by these two methods are found to be consistent, then modelling would enable pattern speed estimation for large galaxy samples."
    },
    "0806/0806.1995.txt": {
        "abstract": "{}{We study the variability of the cataclysmic variable DO Dra, on time-scales of between minutes and decades.} %the outburst, orbital, spin and transient periodic {The observations were obtained at the Korean 1m telescope at the Mt. Lemmon in 2006-2007, 14 observational runs cover 45 hours. The table of individual observations is available electronically. Additionally, we have used 1509 patrol observations from the international AFOEV and VSOLJ databases.} {The characteristic decay time dt/dm=0.902(3) days/mag was estimated from our 3 nights of CCD R observations, which cover the descending branch of the outburst in 2006. The range of the outburst cycle  is from $311^{\\rm d}$ to $422^{\\rm d},$ contrary to a previous estimate of $870^{\\rm d}.$ The ``quiescent\" data show a photometric wave with a cycle $\\sim 303(15)^{\\rm d}.$ We analyzed the profile of the \"composite\" (or \"mean\") outburst. We discovered however, that a variety of different outburst heights and durations had occurred, contrary to theoretical predictions. The analysis of the historical data has shown a correlation between the decay time dt/dm and the outburst maximum brightness with a slope d(dt/dm)/dm=0.37(9). With increasing maximum brightness, we find that the decay time also increases; this is in contrast to the model predictions, which indicate that outbursts should have a constant shape. This is interpreted as representing the presence of outburst-to-outburst variability of the magnetospheric radius. A presence of a number of missed weak narrow outbursts is predicted from this statistical relationship. We tabulate characteristics of the \"quasi-orbital\" variations, which indicate that an amplitude maximum occurs between quiescence and the outburst peak. The semi-amplitude of the spin variability does not exceeded 0.02 mag. A new type of variability is detected, during 3 subsequent nights in 2007: periodic (during one nightly run) oscillations with rapidly-decreasing frequency from 86 to 47 cycles/day and a semi-amplitude increasing from $0\\fm06$ to $0\\fm10,$ during a monotonic brightness increase from $14\\fm27$ to $14\\fm13.$ This phenomenon was observed only during an unusually prolonged event of $\\sim1$ mag brightening in 2007 (lasting till autumn), during which no (expected) outburst was detected. We refer to this behaviour as to the ``transient periodic oscillations\" (TPO).  We attribute the frequency decrease to \"beat\"-type of the variability, probably caused by irradiation of a cloud that is spiralling down to the white dwarf. Its frequency would then increase and coverge towards the spin frequency. To study this new and interesting phenomenon, new regular photometric and spectral (in a \"target of opportunity\" mode) observations are required. }{} ",
        "introduction": "DO Dra belongs to a class of cataclysmic variables that have accretion quite unaffected by the magnetic field. Such objects are called ``intermediate polars\" or ``DQ Her - type stars\" (see Patterson (1994), Warner (1995), Norton et al. (2004) and Hellier (2001) for more detailed description).  This object was detected as an X-ray source 2A 1150+720 and later classified as a cataclysmic variable star by Patterson et al. (1982). It was also detected as a cataclysmic variable in the Palomar-Green Survey, and listed as PG 1140+719 (Green et al. 1982). They suggested an identification of this object with a previously-registered variable YY Dra.  The designation \"YY Dra\" was however assigned to an eclipsing variable that had a brightness range of 12\\fm9 -- $<14\\fm5$, and period 4\\fd21123 as measured by Tsesevich (1934), and was almost coincident in co-ordinates with the X-ray source. Possibly due to a misprint of coordinates, the ``true\" (eclipsing) YY Dra was not found until now. Wenzel (1983) failed to find variability of the $12^{\\rm m}$ star close to the published position of YY Dra, but found an eruptive object at the position of  PG 1140+719. He detected this object on only two plates from 700, and thus classified the object as a dwarf nova with an extremely long cycle length. This is not an eclipsing variable with a well-defined period, and thus a separate official GCVS name ``DO Dra\" was assigned to PG 1140+719 (Kholopov et al. 1985, Samus' et al., 2007). The designation of the star was discussed by Patterson and Eisenman (1987) and Kholopov and Samus (1987). In the literature, DO Dra can still be referred to as \"YY Dra\" or \"DO/YY Dra\".  The physical nature of this object may be inferred from the photometric behaviour of the system as of the dwarf nova. These systems are close binaries with a red dwarf filling its Roche lobe and a white dwarf. The plasma stream from the secondary forms an accretion disk, which becomes cyclically unstable after reaching some critical viscosity (cf. Warner 1995). Smak (1984), based on theoretical models, distinguished between two types of dwarf nova outbursts - with an onset in the {\\it outer} (type A) and {\\it inner} (type B) parts of the accretion disk.  An unusual feature of DO Dra is a short outburst as compared to other systems. This is interpreted by a relatively large inner radius of the accretion disk, which, contrary to ``non-magnetic\" dwarf novae, is equal to the Alfven radius $R_A,$ rather than to the radius of the white dwarf (cf. Angelini and Verbunt 1989). The detailed study of the long-term and outburst behaviour of DO Dra was presented by \\v{S}imon (2000) on the base of visual amateur observations. He estimated an outburst cycle length of $\\sim868^{\\rm d},$ but further monitoring presented in the international databases of AFOEV (2007) and VSOLJ (2006) imply much shorter time intervals (as discussed in Section 2).  So large outburst cycle is in addition implied by the relatively large inner radius $R_{in}$ of the accretion disk. For non-magnetic cataclysmic variables, $R_{in}$ is compared with the radius of the white dwarf $R_{wd}.$ For intermediate polars $R_{in}\\underline{\\sim} R_{A},$ where $R_{A}$ is the radius of magnetosphere (Angelini \\& Verbunt 1989). Another bright magnetic dwarf nova with short outbursts and long outburst cycle is GK Per (see \\v{S}imon 2007 for a recent review). Cannizzo \\& Mattei (1992, 1998) have thoroughly studied characteristics of 705 outbursts of the prototype dwarf nova SS Cyg observed during 95 years. They found a bimodal distribution. The slope dm/dt for the descending branch (the decay of the outburst) is nearly constant (within a dozen per cent) for different outbursts, in agreement with theoretical models.  Patterson et al. (1992) found a ``fundamental\" period of $550^s\\pm3^s$ with a double-peak structure of the phase curve. The highest peak in the amplitude spectrum therefore occurs at a half of the period and has an amplitude equal to $275^s\\pm1^s.$ For some nights, the shorter period was found to be $266^s.$ This value was interpreted as a half of the siderial period of the magnetic white dwarf. This interpretation has been made in other papers, even though it is accepted that the \"fundamental\" spin period was determined as $P_{spin}=529.31(2)$ s (Haswell et al. 1997). The number in parentheses corresponds to an accuracy estimate (i.e. standard error). Haswell et al. (1997) precisely determined the orbital period $P_{orb}=0\\fd16537398(17),$ the initial epoch $T_0=244683.4376(5)$ for the inferior conjunction of the secondary (we use these values here to compute the orbital phases), the masses of both the secondary $(M_2=0.375(14)M_\\odot,$ and the white dwarf $M_{wd}=0.83(10)M_\\odot,$ and the inclination angle $i=45^\\circ(4^\\circ).$  With such a moderate inclination, there are no eclipses of the accretion disk/columns by the red dwarf, so the variability with orbital phase may be attributed mainly to the secondary star - either the ellipticity effect (cf. detection for EF Eri by Allen and Cherepashchuk (1982)) or irradiation (cf. Basko and Sunyaev 1973, King and Lasota 1984). Results of X-ray/optical studies of two outbursts (1999 and 2000) were presented by Szkody et al. (2002), who noted that the behaviour of the spin pulse amplitude with luminosity was unusual.  Norton et al. (1999) have split the group of intermediate polars into two subclasses with relatively large and relatively weak magnetic field and classified DO Dra as belonging to the second group (with DQ Her, V709 Cas et al.). In the catalogue of Ritter and Kolb (2003), there are 48 objects classified as intermediate polars (IP) or DQ Her-type stars (DQ), from which 36 (75\\%) are nova-like (NL) variables, in 7 (15\\%), the Nova outbursts were detected, and 5 (10\\%) show dwarf nova-type outbursts. Consequently, the latter group may be called either \"magnetic dwarf novae\", or \"outbursting intermediate polars\".  In this paper, we study the time and luminosity variability in observations of DO Dra, acquired over 13 nights, at quiescence, outburst peak, and at the descending branch of the light curve. Additionally, we reanalyze patrol observations published in the AFOEV (2007) and VSOLJ (2006) databases.  \\begin{figure} \\psfig{file=9056fig1.eps} \\caption{Finding chart for DO Dra. The size of the field is $7.3'\\times4'.$} \\end{figure}  %\\unitlength=1in %\\begin{figure} %\\label{f1} %\\begin{center} %\\begin{picture}(3.47,1.89) %\\put(0,1.89){\\special{em:graph 9056fig1.bmp}} %\\end{picture} %\\caption{Finding chart for DO Dra. The size of the field is $7.3'\\times4'.$} %\\end{center} %\\end{figure}  \\begin{table} \\label{t1} \\caption{Mean brightness of DO Dra and the comparison stars from the Mt.Lemmon observations (R). The values of $\\sigma$ correspond to the r.m.s. scatter of a single observation in respect to the \"artificial comparison star\". The statistical accuracy of the mean value is by a factor of $n^{1/2}=1350^{1/2}\\approx37$ times smaller, i.e. does not exceed 0\\fm001. } \\centerline{\\begin{tabular}{cccccccccc} \\hline Star&R&$\\sigma_R$&Star&R&$\\sigma_R$\\\\ \\hline DO &13.747 & 1.320 & C4 &13.754 & 0.024 \\\\ C1 &13.840 & 0.014 & C5 &14.631 & 0.033 \\\\ C2 &15.016 & 0.020 & C6 &14.831 & 0.026 \\\\ C3 &15.378 & 0.024 & C7 &13.543 & 0.009 \\\\ \\hline \\end{tabular}} \\end{table} ",
        "conclusions": ""
    },
    "0806/0806.3158_arXiv.txt": {
        "abstract": "The 2PI effective action formalism for quantum fields out of equilibrium is set up in an expanding (Friedmann-Robertson-Walker) background. We write down and solve the evolution equations for a $\\varphi^4$ model at $\\mathcal{O}(\\lambda^2)$ in a coupling expansion. We comment on issues of renormalization, lattice discretization and the range of applicability of the approach. A number of example calculations are presented, including thermalization and (p)reheating. Generalizations to more complicated systems and applications are discussed. ",
        "introduction": "}  In recent years, significant attention has been drawn to the process of thermalization of quantum fields. Quantitative description of the physics of the very early Universe and of heavy-ion collision experiments requires an understanding of the real-time dynamics of quantum fields at finite energy density, but out of equilibrium.  One very promising development is the application of the 2PI-formalism \\cite{Cornwall:1974vz,Berges:2000ur,Berges:2001fi}, which allows the derivation and explicit numerical solution of a set of equations of motion for the mean field and propagator in the full quantum theory. This is realized through the truncation of a controlled diagram expansion in terms of 2PI diagrams. Already at next-to-leading order (NLO) in either a coupling or $1/N$ expansion, interacting systems exhibit equilibration, effective dissipation and thermalisation to the {\\it quantum} equilibrium state \\cite{Berges:2000ur,Aarts:2002dj,Berges:2002wr,Juchem:2003bi,Berges:2004ce,Arrizabalaga:2005tf,Lindner:2005kv}.  The physics of the early Universe is described not in a static background, but in expanding space. This is often approximated by a homogeneous, flat Friedmann-Robertson-Walker (FRW) space-time, parametrized by the scale factor $a(t)$ of the metric $ds^2=dt^2-a^2(t)d{\\bf x}^2$. In some cases, expansion can be neglected when the time scale of a phenomenon is very short on the time scale of expansion. But in general, and of course in principle, expansion should be included in the description of early Universe physics.  Many processes at high temperature or energy density are well described by the classical approximation, where a Monte-Carlo sample of initial field configurations are evolved using Hamiltonian equations of motion (see \\cite{Smit:2005vp} for a brief review). The observables of interest are then averages over this classical ensemble. The approximation must however break down eventually, as classical fields equilibrate to a classical equilibrium, which suffers from the Rayleigh-Jeans problem: in the continuum temperature will go to zero, and on the lattice it will be cut-off dependent.  In this paper, we set out the 2PI formalism in a FRW space-time, and solve the resulting equations numerically for some example applications. A number of studies have been carried out in this context in the Hartree approximation \\cite{Boyanovsky:1993xf,Khlebnikov:1996mc,Khlebnikov:1996wr,Boyanovsky:1996fz,Boyanovsky:1997cr} (which is also leading order (LO) in a 2PI coupling expansion), and even before that, the formalism was set out in \\cite{Calzetta:1986ey,Calzetta:1986cq}. Recently, attempts have been made to partly include the effect of an expanding background for specific applications also at NLO \\cite{Rajantie:2006gy,Aarts:2007qu,Aarts:2007ye}.  Two main issues present themselves. Firstly, as we are discretizing the system on a finite co-moving lattice, there is only a finite number of momentum modes available, and as the lattice expands in time these will be redshifted towards the IR in physical units. This means that the physical cut-off changes in time. Therefore, there is a limit on how many e-folds one can run the simulation before running out of ``dynamical range''. In practice, this means that at some point discretization errors become important, and results can no longer be relied upon. As a result, reliably simulating cosmological inflation proper is a daunting task, as the Universe expands many e-folds. Still, most of the inflationary stage is often well described by semi-analytical tools and the slow-roll approximation, and only the couple of e-folds around the end of inflation, reheating and the transition to radiation domination requires numerical treatment. Post-inflationary phenomena typically only span a few e-folds.  Secondly, since we are doing quantum physics, the theory has to be renormalized. In particular, the energy density which enters in the semi-classical Friedmann equation (see below) needs appropriate counterterms. Fortunately, features of the 2PI formalism include that it is renormalizable at any level of diagram truncation  \\cite{vanHees:2001ik,Blaizot:2003br,Berges:2004hn,Berges:2005hc,Arrizabalaga:2006hj}, and that there is a similarly truncated energy density which is conserved. Hence by introducing (scale-factor dependent) counterterms for the energy density, mass and couplings, we can in principle cancel all divergences, and construct a well-behaved Friedmann equation.  We study a self-interacting real scalar, and go to NLO ($\\mathcal{O}(\\lambda^2)$) in a 2PI coupling expansion. By showing how to apply the procedure in practice, we expect it will be clear how to generalize to more complicated systems. In the conclusions we point out some issues, applications and ways of refining the approach.    \\subsection{Setup\\label{sec:setup}}  We are concerned with a single scalar field with $\\varphi^4$ interaction. The action is \\be S= \\int dt\\,d^{3}\\vcx\\, a^{3}(t)\\bigg[ \\frac{1}{2}(\\partial_t\\varphi)^{2}-\\frac{1}{2a^2(t)}(\\partial_{\\bf x}\\varphi)^{2}-\\frac{1}{2}m^2\\varphi^2-\\frac{\\lambda}{24}\\varphi^4\\bigg], \\label{eq:action1} \\ee written in terms of co-moving spatial coordinates {\\bf x}. $a(t)$ is the scale factor, and we assume $a(0)=1$. Correspondingly, we will consider co-moving and physical momenta, denoted $\\vck$ and $\\tilde\\vck=\\vck/a(t)$, respectively.  The evolution of the scale factor is in turn given by the Friedmann equation in terms of the Hubble rate $H$\\footnote{We use $\\dot{a}$ to denote $\\partial_t a(t)$. $a'$ will denote $\\partial_\\eta a(\\eta)$, with $\\eta$ conformal time.}, \\be H^2(t)=\\frac{1}{3M_{\\rm pl}^2}\\langle T^{00}(t)\\rangle_{\\rm ren}, \\qquad H(t)=\\frac{\\dot{a}(t)}{a(t)}. \\label{eq:friedmann} \\ee Here we equate a classical quantity on the left-hand side to a quantum expectation value on the right-hand side. This only makes sense when the energy density is appropriately renormalized, an issue we will return to below.  The system can be recast in comoving (conformal) time $\\eta$, with $dt= a(\\eta)d\\eta$ and we can rescale the field\\footnote{The equations of motion will be solved in terms of the ``conformal'' field $\\phi(\\eta)$, but results converted back to the ``physical'' field $\\varphi(t)$.} $\\varphi(x)=\\phi(x)/a(\\eta)$, in which case the action becomes \\be S= \\int d\\eta\\,d^{3}\\vcx\\,\\bigg[ \\frac{1}{2}(\\partial_\\eta \\phi-\\mathcal{H}\\phi)^{2}-\\frac{1}{2}(\\partial_{\\bf x} \\phi)^2-\\frac{1}{2}a^2(\\eta)m^2\\phi^2-\\frac{\\lambda}{24}\\phi^4\\bigg]. \\label{eq:action2} \\ee We have introduced a new ``comoving Hubble rate'', $\\mathcal{H}=a'/a=aH$. In terms of the canonical momentum $\\pi=\\partial_\\eta \\phi-\\mathcal{H}\\phi$, the corresponding Friedmann equation is \\footnote{The right-hand side is not the Hamiltonian corresponding to the action (\\ref{eq:action2}), but the Hamiltonian of (\\ref{eq:action1}), written in terms of the rescaled fields.} \\be \\frac{(a')^2}{a^4}=\\frac{1}{3a^4M_{\\rm pl}^2}\\langle\\bigg[ \\frac{1}{2}\\pi^{2}+\\frac{1}{2}(\\partial_{\\bf x} \\phi)^2+\\frac{1}{2}a^2(\\eta)m^2\\phi^2+\\frac{\\lambda}{24}\\phi^4\\bigg]\\rangle. \\label{eq:friedmann2} \\ee  In passing, it is useful to recall the classical equation of motion \\be \\left[\\partial_\\eta^2-\\partial_\\vcx^2-\\frac{a''}{a}+a^2m^2+\\frac{\\lambda}{6}\\phi^2(x)\\right]\\phi(x)=0. \\label{eq:classical} \\ee The classical approximation amounts to generating a set of random initial conditions, solve for the evolution using (\\ref{eq:classical}) together with the Friedmann equation and then to average observables over initial conditions. In addition to using approximate dynamics, also the classical averaging procedure is different from (\\ref{eq:friedmann2}) in that the Hubble rate, and hence $a(t)$, is derived from each individual initial condition rather than the average energy density. This could be resolved by simulating all initial conditions simultaneously using a common $a(t)$ determined through the ensemble averaged energy density. Still, it would be a classical average rather than a quantum one. ",
        "conclusions": "}  We have seen that extending the 2PI formalism of out-of-equilibrium quantum fields to expanding backgrounds amounts to introducing a time dependent mass $m^2\\rightarrow a^2m^2-a''/a$ in the conformal time, rescaled field, equations of motion (10-19). Having solved these equations on the lattice, observables are translated back to physical fields and physical time. The scale factor is derived from the semi-classical Friedmann equation involving the renormalized energy $\\langle T^{00}\\rangle_{\\rm ren}$. We have here opted for an approximate renormalization strategy, where counterterms for the mass and the energy density are calculated in the LO/Hartree approximation in terms of a particular vacuum, the adiabatic free-field solution to second order, both in WKB and in $\\dot{a}$. We argued that going beyond this order is possible, although only really necessary when aiming at taking the continuum limit or using very large expansion rates. A fully 2PI renormalization beyond LO is much harder.  On a finite comoving lattice, modes are redshifted towards the IR. Therefore the number of e-folds of expansion available is limited; a simulation can only be trusted as long as there is a range of UV modes that stay in the vacuum. Otherwise cut-off effects will influence the physics and presumably the renormalization. One way of quantifying this is for the time-dependent mass $am$  to stay less than unity. In the ``conformal'' case $m=0$ another mass scale (temperature, initial mean field) will play a similar role.  At the end of the day this is a practical question of computer capacity. 2PI simulation are memory intensive in that the memory of past time-steps must be saved to generate the self-energy kernels (right hand sides of (10-12)). On the other hand, no statistical averaging is necessary as the solutions to the equation are the full correlators. The total simulations performed here amount to about 5000 CPU hours.  Possible applications range over all the topics already studied in Minkowski space: Thermalisation \\cite{Berges:2000ur,Berges:2002wr,Juchem:2003bi,Berges:2004ce,Arrizabalaga:2005tf,Lindner:2005kv}, which enters in early Universe physics as well as heavy ion collisions\\footnote{In heavy ion collisions the expansion is somewhat different from FRW.}; reheating and preheating, both the resonant variety \\cite{Berges:2002cz} and tachyonic preheating \\cite{Felder:2000hj,Arrizabalaga:2004iw}. In the latter cases, it may be prudent to use a diagram expansion in $1/N$ rather than $\\lambda$, to be sure corrections are under control.  In the present paper we have made test-runs of many of these cases and pointed out the main effects of the cosmological expansion. All of these applications deserve further scrutiny, also of combined effects of smaller/larger couplings, expansion rate, temperature and masses. A study of the inflationary regime was not attempted, mainly because it would require very large lattices (lots of expansion) and/or $m=0$ which is a very special case.  One interesting result is the difference between the kinetic and chemical equilibration timescales, and the possibility of having one but not the other happen. Clearly, one must be careful when assuming instantaneous thermalisation, as is sometimes done when considering reheating, preheating and phase transitions in the early Universe. Using a criterion like $\\Gamma/H\\gg 1$ to ensure instant thermalisation presumes careful consideration of which $\\Gamma$ is the relevant one.  We were also able to confirm that the amount of resonant preheating depends sensitively on the rate of expansion. This is because field modes are redhifted in and out of resonance bands. This shortens the resonance time, making even exponential growth much less effective.  The quantum 2PI equations have a classical counterpart, i.e. 2PI-truncated equations for classical correlators, reproducing classical dynamics \\cite{Aarts:2001yn}. The relation between the two amounts to neglecting terms like $\\rho^2$ compared to terms like $F^2$ in the self-energies (14-16), and ignoring renormalization. Although the full classical approximation is in principle exact, it relies on  statistical averaging over initial conditions. Classical 2PI has no statistical errors, but diagram expansion truncation introduces a different type of approximation. As such, it contributes an alternative way of doing classical simulations.  Extension to more complicated models than a single self-interacting scalar is straightforward. Possible applications include reheating and preheating with multiple fields (including fermions \\cite{Berges:2002wr}), departure from and return to equilibrium for systems with heavy particles decaying into light ones. In the context of multi-field preheating this may allow to calculate non-gaussian signatures in the CMB \\cite{Chambers:2007se,Chambers:2008gu}.  In conclusions, we believe that given sufficient numerical capacity, and observing certain simple rules, the 2PI formalism for out-of-equilibrium fields provides a convenient tool for quantitative calculations in cosmology.    \\vspace*{0.5cm} \\noindent {\\bf Acknowledgments.} It is a pleasure to thank Gert Aarts, Kari Rummukainen and Szabolcs Borsanyi for many comments, and Arttu Rajantie and Jan Smit for enlightening discussions and work on related topics. I am indebted to Julien Serreau for a host of useful suggestions and for pointing out a couple of errors. The numerical work was conducted on the Murska cluster at the Finnish center for computational sciences, CSC. This work was supported by Academy of Finland Grant 114371.     \\appendix"
    },
    "0806/0806.4587_arXiv.txt": {
        "abstract": "{This paper reports on the detection of a planetary system with three Super-Earths orbiting HD\\,40307. HD\\,40307 is a K2\\,V metal-deficient star at a distance of only 13\\,parsec, part of the HARPS GTO high-precision planet-search programme. The three planets on circular orbits have very low minimum masses of respectively 4.2, 6.9 and 9.2 Earth masses and periods of 4.3, 9.6 and 20.5 days.  The planet with the shortest period is the lightest planet detected to-date orbiting a main sequence star. The detection of the correspondingly low amplitudes of the induced radial-velocity variations is completely secured by the 135 very high-quality HARPS observations illustrated by the radial-velocity residuals around the 3-Keplerian solution of only 0.85\\,ms$^{-1}$.  Activity and bisector indicators exclude any significant perturbations of stellar intrinsic origin, which supports the planetary interpretation.  Contrary to most planet-host stars, HD\\,40307 has a marked sub-solar metallicity ([Fe/H]\\,=\\,$-$0.31), further supporting the already raised possibility that the occurrence of very light planets might show a different dependence on host star's metallicity compared to the population of gas giant planets.  In addition to the 3 planets close to the central star, a small drift of the radial-velocity residuals reveals the presence of another companion in the system the nature of which is still unknown. ",
        "introduction": "The planet-search programme conducted at high precision with the HARPS spectrograph on the ESO 3.6-m telescope at La Silla aims at the detection of very low-mass planets in a sample of solar-type stars already screened for giant planets at a lower precision with CORALIE on the 1.2-m Swiss telescope on the same site. About 50\\,\\% of the HARPS GTO time is dedicated to this survey. After 4.5 year of the programme, we are starting to see a large population of Neptune-mass and super-Earth planets emerging from the data, including the system presented here.  Several reasons motivate our interest to search for very low-mass planets, with masses in the range of the Neptunes or the so-called Super-Earths ($\\sim$\\,2\\,M$_{\\oplus} \\leq m_2\\sin{i} \\leq 10$\\,M$_{\\oplus}$).  i) Over the past decade, several statistical distributions of the orbital elements of gaseous giant planets have emerged from the nearly 300 detected planetary systems \\citep[see e.g.][]{Udry-2007:b,Marcy-2005}. These statistical properties provide constraints to complex physical scenarios of planetary system formation. One of the most obvious example of that dialogue between planetary formation theory and observations is illustrated by the comparison of the planetary mass vs semi-major axis ($m_2-a$) diagram \\citep{Ida-2004:a,Mordasini-2008}. Comparison can be made for specific categories of host stars by selecting different primary masses ($m_1$) or metallicities ([Fe/H]). In all cases, global features of planet formation directly affect the overall topology of the ($m_2-a$) diagram. In particular, the location of the large population of very low-mass planets predicted by the models  \\citep{Mordasini-2008,Ida-2008} depends upon the extend of migration undergone by the planets during their formation. The detection of a large sample of planets with masses less than roughly 25\\,M$_{\\oplus}$ at relatively close distances is therefore an important indicator of the efficiency of type I migration (assuming that the planets are not too close so that evaporation can be neglected). Despite the still very limited number of planets detected in the range of Neptune masses, already a few interesting characteristics are emerging \\citep{Mayor-2008}:  - The distribution of planetary masses appears as bimodal. A new population of light planets, although more difficult to detect, is differentiating itself from the distribution of giant-planet masses. The Neptune- and super-Earth mass distribution seems not to be the extrapolation towards lower masses of the distribution for gaseous giant planets.  - The very strong correlation observed between the host star metallicity and occurrence frequency of giant planets \\citep{Santos-2001:a,Santos-2004:b,Fischer-2005} seems to be vanishing or at least to be reduced \\citep{Udry-2006}.  - Neptune-mass planets and Super-Earths are found most of the time in multiplanetary systems ($>$\\,80\\,\\% of the known candidates).  ii) Simulations of planetary formation do not only provide the statistical distributions of masses and semi-major axes. For every planet, we have, in addition, a prediction of its internal structure. The internal composition of the planet at the end of the formation/migration process carries a fossil signature of the system history. The end-state diversity is broad: rocky planets, icy planets, ocean planets, evaporated gaseous giant planets, or possibly objects with variable percentages of these ingredients.  The predicted distributions of the planetary internal composition, as a function of the different significant parameters ($m_1$, $m_2$, $a$, [Fe/H]) can be observed in the corresponding radius-mass ($m_2 - R$) diagrams. These predicted ($m_2 - R$) distributions can then be compared to the observed distributions derived from combined radial-velocity and transit searches.  The coming soon results from space missions searching for planetary transits combined with ground-based, high-precision, radial-velocity follow-up, will provide rich observational constraints to the planetary formation theory via the $m_2-R$ distributions of low-mass planets.  iii) In a more distant future, space missions will be developed to search for life signatures in the atmosphere of terrestrial-type planets. Before the detailed design of such ambitious missions, it would be wise to have first insights in the occurrence frequency of terrestrial planets and on the statistical properties of their orbits. Still more valuable is the detection of planets in the habitable zone of our closest neighbours, initiating the preparation of an \"input catalogue\" for these future missions.  Already a few planets have been detected with masses between 3 and 10\\,M$_{\\oplus}$.  In 1992, from precise timing, \\citep{Wolszczan-92} have discovered two planets with masses of 2.8 and 3.4\\,M$_{\\oplus}$ orbiting the pulsar PSR1257+12 on almost perfect circular orbits. Microlensing technique has demonstrated as well its potential to detect low-mass planets. Two planets with masses possibly in the range of super-Earths have been announced: a planet of about 5.5\\,M$_{\\oplus}$ orbiting a low-mass star \\citep{Beaulieu-2006} and a still less massive object, of only 3.3\\,M$_{\\oplus}$, probably gravitationally bound to a brown dwarf \\citep{Bennett-2008}.  Doppler spectroscopy also revealed quite a few planets with $m_2\\sin{i}$ less than 10\\,M$_{\\oplus}$: GJ\\,876\\,d, $m_2\\sin{i}=5.9$\\,M$_{\\oplus}$ \\citep{Rivera-2005}; GJ\\,581\\,c and d, with $m_2\\sin{i}$ of 5.1 and 8.2\\,M$_{\\oplus}$, respectively \\citep{Udry-2007:a}; HD\\,181433\\,b with $m_2\\sin{i}=7.5$\\,M$_{\\oplus}$ \\citep{Bouchy-2008}. GJ\\,876 as well as GJ\\,581 are both stars at the bottom of the main sequence with spectral types M4\\,V and M3\\,V, respectively. HD\\,181433 and HD\\,40307 are on the other hand both K dwarfs.  The detection of planetary systems around bright stars by Doppler spectroscopy is very important because it concerns mostly nearby stars for which we can obtain interesting further information (e.g. on their orbital, planetary and stellar properties), opening the possibility of rich complementary studies. This is even more true when the planet is transiting in front of its parent star (see for illustration the wealth of studies concerning GJ\\,436).  In this paper, we characterize the orbits of three new super-Earth planets on short-period trajectories orbiting the main sequence star HD\\,40307. Section\\,2 will briefly describe the improvement of the HARPS radial velocities in term of observational strategy and software developments. The stellar characteristics of HD\\,40307 are presented in Sect.\\,3. Section\\,4 deals with the derived orbital solution while summary and conclusion are given in Sect.\\,5. ",
        "conclusions": "We report the detection of 3 super-Earth planets orbiting the low metallicity K dwarf HD\\,40307, a star located at only 13\\,pc from the Sun. The high precision radial velocities acquired with the HARPS spectrograph on the ESO 3.6-m telescope enabled this discovery. The 3-Keplerian fit reveals the presence of 3 low-mass planets. The closest one HD\\,40307\\,b with $m_2\\sin{i}$\\,=\\,4.2\\,M$_{\\oplus}$ is presently the lightest exoplanet detected around a main sequence star. The two other planets, with masses of 6.9 and 9.2\\,M$_{\\oplus}$ also belong to the category of super-Earths. All the 3 planets are on circular orbits.  It is amazing to notice that the global rms of the 135 measurements, before fitting any planets, was only 2.94\\,ms$^{-1}$.  The sigma(O-C) after the 3-planet Keplerian fit and drift is 0.85\\,ms$^{-1}$.  We will continue the velocity monitoring to better characterize the longer period 4th object bound to the system, revealed by the additional observed linear drift of the radial velocities.  Available Spitzer IRS data of HD\\,40307 do not show any IR excess in the 10--40\\,$\\mu$m region of the spectra (Augereau, private communication). No warm dust disk is thus detected in the inner regions of the system, unlike the case of HD\\,69830, the star harbouring a trio of Neptune planets \\citep{Lovis-2006} and for which an observed IR excess indicates the presence of a debris disk, possibly under the form of an asteroid belt \\citep{Beichman-2005}.  The characterization of multi-planetary systems with very low-mass planets require a rather large number of measurements.  After 4.5 years of the HARPS programme and with the improved reduction software, several dozens of planets with masses less than 30 Earth-masses and period less than 50 days have been detected. Coming observations will confirm these detections and allow us to fully characterize the systems.  The domain of Neptune-type and rocky planets will be drastically boosted in a near future with these detections.  In particular we expect to have enough systems to revisit the emerging properties for these low mass planets as tentatively discussed by \\citet{Mayor-2008}:\\\\ \\indent - is the mass-distribution of exoplanets bimodal?  \\\\ \\indent - Is the correlation between host star metallicity and occurrence frequency of Neptune-type planets (or smaller) still existing?\\\\ \\indent - What is the frequency of multiplanetary systems with these low mass planets?\\\\ \\indent - What is the frequency of Neptune or rocky planets orbiting G and K dwarfs?  A first estimate based on the HARPS high-precision survey suggests a frequency of 30\\,$\\pm$\\,10\\,\\% in the narrow range of periods shorter than 50 days.   One of the most exciting possibility offered by this large emerging population of low-mass planets with short orbital periods is the related high probability to have transiting super-Earths among the candidates. If detected and targeted for complementary observations, these transiting super-Earths would bring a tremendous contribution to the study of the expected diversity of the structure of low-mass planets."
    },
    "0806/0806.3969_arXiv.txt": {
        "abstract": "We apply a new, second-order Godunov code, Athena, to studies of the magnetorotational instability (MRI) using unstratified shearing box simulations with a uniform net vertical field and a sinusoidally varying zero net vertical field.  The Athena results agree well with similar studies that used different numerical algorithms, including the observation that the turbulent energy decreases with increasing resolution in the zero net field model.  We conduct analyses to study the flow of energy from differential rotation to turbulent fluctuations to thermalization.  A study of the time-correlation between the rates of change of different volume-averaged energy components shows that energy injected into turbulent fluctuations dissipates on a timescale of $\\Omega^{-1}$, where $\\Omega$ is the orbital frequency of the local domain.  Magnetic dissipation dominates over kinetic dissipation, although not by as great a factor as the ratio of magnetic to kinetic energy. We Fourier-transform the magnetic and kinetic energy evolution equations and, using the assumption that the time-averaged energies are constant, determine the level of numerical dissipation as a function of length scale and resolution.  By modeling numerical dissipation as if it were physical in origin, we characterize numerical resistivity and viscosity in terms of effective Reynolds and Prandtl numbers.  The resulting effective magnetic Prandtl number is $\\sim 2$, independent of resolution or initial field geometry.  MRI simulations with effective Reynolds and Prandtl numbers determined by numerical dissipation are not equivalent to those where these numbers are set by physical resistivity and viscosity. These results serve, then, as a baseline for future shearing box studies where dissipation is controlled by the inclusion of explicit viscosity and resistivity. ",
        "introduction": "\\label{introduction}  The process of accretion powers a wide range of astrophysical systems, from protostars to quasars.  In accretion disks, gravitational energy is converted into other forms including bulk outflows, heat, and radiation. In the traditional time-stationary thin disk model of \\cite{shak73}, the $r,\\phi$ component of the stress, $\\tau_{r\\phi}$, is proportional to the local pressure, $\\tau_{r\\phi} = \\alpha P$.  The $\\alpha$ model assumes that the accretion energy is deposited as heat locally and radiated rapidly, providing a relation between disk emissivity and accretion rate. While the $\\alpha$ model has proven valuable in interpreting many aspects of accretion systems, advancing beyond it will require a more detailed understanding of the stress that produces angular momentum transport as well as the physical processes involved in the subsequent thermalization and radiation of the orbital energy released by those stresses.  It is now understood that magnetohydrodynamic (MHD) turbulence generated by the magnetorotational instability (MRI) \\cite[]{balb91,balb98} produces significant Maxwell stresses, $-B_rB_\\phi/4\\pi$, and Reynolds stresses, $\\rho \\delta v_r \\delta v_\\phi$, that account for transport within accretion disks.  The absence of an analytic theory for MHD turbulence, however, means that direct numerical simulations play an essential role in investigating accretion physics.  In this regard, local simulations, which reduce the problem to the simplest form that can sustain MRI-driven turbulence, have proven very useful.  The ``shearing box'' model is a representation of a small patch of the disk constructed by boosting to a local co-rotating Cartesian frame that ignores geometric curvature but retains all rotational forces.  MRI shearing box simulations were introduced by \\cite{haw95} and have been extensively used since then both without \\cite[e.g.,][]{haw96,balb98} and with vertical stratification \\cite[e.g.,][]{bran95,stone96,hir06}.    Shearing box simulations can investigate several key questions including the functional dependence of the stress on disk properties and the turbulent energy flow that leads to dissipation as heat. These simulations have made it increasingly clear, for example, that the basic $\\alpha$ stress parameterization is not only too simplistic, it is actually misleading.  Shearing boxes have provided ample evidence that stress is {\\it not} determined by pressure, at least in the usual manner of the $\\alpha$ disk \\cite[][]{haw95, sano04}. Early studies showed instead that stress is (in some cases) proportional to the {\\it magnetic} pressure, but the magnetic energy is not itself directly determined by the gas and radiation pressure.  \\cite{black08} recently reviewed a large number of shearing box results and found that this result holds across the full ensemble of simulations with only small differences in the constant of proportionality from one run to another.  The implications of these results are significant.  For example, recent local simulations using stratified shearing boxes and radiation transport \\cite[][]{bla07, kro07} have found no evidence of the thermal instability long believed to be present in radiation-pressure supported $\\alpha$ disks.    If stress is proportional to magnetic rather than total pressure, what determines the magnetic pressure in a disk?  Apart from the expectation that the field will remain subthermal, this remains uncertain. The simplest shearing box simulations using ideal MHD have a limited range of significant parameters; this is both a strength and a weakness of that model.  The magnetic energy in the saturated state could depend upon such factors as box size, the amplitude and geometry of the imposed initial magnetic field, and the ratio of the gas pressure to magnetic pressure (the plasma $\\beta$ value).  \\cite{haw95} and \\cite{haw96} studied the effect of initial magnetic field topology on the resulting stress  and found that although the MRI leads to turbulence regardless of the initial field, simulations that had an imposed net vertical field produce higher turbulence levels than an imposed toroidal field or a simulation that began with zero net magnetic flux within the domain.  \\cite{haw95} found that the total magnetic energy and the resulting stress in the saturated turbulent state was a function of the initial plasma $\\beta$ with a uniform vertical field, namely that larger $\\beta$ (i.e., weaker fields) leads to smaller saturation levels.  Other initial field configurations do not yield so direct a correlation between background field strength and saturation.  Many simulations have failed to find any noticeable correlation between mean turbulent magnetic energy and the gas pressure. A comprehensive parameter study by \\cite{sano04} observed at best only a very weak gas pressure dependence.     Since the mean magnetic energy at saturation is presumably a balance between continued driving by the MRI and loss due to magnetic dissipation and reconnection, there has been interest in going beyond ideal MHD to include explicit physical dissipation in the form of kinematic viscosity, $\\nu$, and Ohmic resistivity, $\\eta$.  Both of these properties have been shown to be important in determining the mean energies and stresses in MRI turbulence. Simulations by \\cite{haw96}, \\cite{sano98}, \\cite{flem00}, \\cite{sano01}, \\cite{zieg01}, and \\cite{sano02b} have investigated the impact of a nonzero $\\eta$.  The main result of these studies is that increasing the resistivity leads to a decrease in turbulence, independent of the initial field configuration.  In zero net field models, the effect of resistivity on the turbulence is larger than one might expect from the linear MRI relation \\cite[]{flem00}.  On the other hand, \\cite{haw96} found that increasing the viscosity increased the magnetic energy in the saturated state.  Recent work has clarified the situation by demonstrating a dependence of the saturation level on both $\\eta$ and $\\nu$ in terms of the magnetic Prandtl number, $\\pr = \\nu/\\eta$.  In particular, the level of angular momentum transport increases with increasing $\\pr$ for simulations initiated with a uniform as well as vanishing mean magnetic field in the vertical direction \\cite[]{from07b,lesur07}.    Determining the stress levels in MRI turbulence is only one aspect of the problem; another is exploring how that turbulence is dissipated into heat.  This question has direct relevance to phenomenological disk models as well as observations.  The $\\alpha$ model assumes that the accretion energy is deposited as heat locally and rapidly, and \\cite{balb99} showed that this property should hold for the energetics of MHD turbulence as well. In the simulations, we can determine the rate at which turbulent energy is thermalized and the path that energy takes as it moves from the free energy of the shear flow to turbulence and then to heat.  Such issues were briefly touched on by \\cite{bran95} who found that the turbulent magnetic energy was $\\sim 6$ times greater than the perturbed kinetic energy, but dissipational heating resulted from roughly equal contributions of magnetic and kinetic energy dissipation.  This result led them to suggest that there was a net transfer of magnetic energy to turbulent kinetic energy. \\cite{sano01} studied energy flow in the context of MRI channel modes, which are strong radial streaming motions that result from the linear growth of the vertical field MRI \\cite[]{haw92,balb98}.  Their work included Ohmic resistivity (but not viscosity) and showed that resistive heating dominated the thermalization of energy stored in these channel modes. Dissipational heating also plays an important role in radiative effects and determining disk structure, both of which may be observable properties of disks \\cite[e.g.,][]{beck08}.    In any study that depends on simulations, there remain factors which cannot be overlooked:  the effects due to numerics and finite resolution.  The majority of the results to-date were obtained with numerical codes based on the finite-difference ZEUS algorithm \\cite[]{stone92a,stone92b}, carried out at relatively low resolution. ZEUS is effectively first-order in asymptotic convergence, and in its most widely used form, evolves the internal rather than the total energy equation.  There have been improvements in both the available computational power, which makes higher resolutions and longer evolution times possible, and in the algorithms for compressible MHD. In this work, we will reexamine the properties of MHD turbulence in the shearing box using a higher-order, Godunov scheme.    The new code, $\\at$, \\cite[see][]{stone08} represents an improvement over ZEUS in several ways including true second-order convergence, increased effective resolution \\cite[see][]{stone05}, accurate shock capturing, and conservation of total energy.  The energy-conserving properties of $\\at$ allow us to study energy flow and dissipation within the shearing box in greater detail than allowed for by the ZEUS algorithm.  The version of $\\at$ we use in this paper does not include explicit resistivity or viscosity and instead relies on numerical dissipation to thermalize the turbulent energy.  Nevertheless, this work will serve as a starting point for planned studies of nonideal effects, including the influence of $\\pr$ on the turbulence \\cite[][]{from07b, lesur07}.  As an important part of establishing a baseline of simulations, we will characterize the numerical resistivity and viscosity of $\\at$ for the shearing box problem.  To do so, we will follow the recent work of \\cite{from07a} who studied the numerical effects of ZEUS on the saturated state of MRI shearing box simulations that begin with zero net field.  They found that the amplitude of the turbulence decreases with increasing resolution and developed several useful diagnostics with which to quantify the effective numerical resistivity and viscosity in the problem.    The structure of the paper is as follows.  In \\S~\\ref{method}, we describe the algorithm employed and our simulations.  In \\S~\\ref{general_properties}, we reexamine some of the results from previous MRI studies and provide a comparison with these studies.  In \\S~\\ref{energy_fluctuations}, we present the first of two diagnostics used to study turbulent energy flow and dissipation.  The second of these diagnostics is applied in \\S~\\ref{trans_funcs}.  Finally, we discuss our results and summarize our conclusions in \\S~\\ref{conclusions}. ",
        "conclusions": "\\label{conclusions}  We have carried out a series of local, unstratified shearing box simulations with the recently developed $\\at$ code to study the characteristics of MRI driven turbulence.  $\\at$ uses a second-order, conservative, compressive MHD algorithm, which is significantly different from the algorithms employed in many of the previous MRI studies. In our work, we have run several standard models for comparison with previous work, and characterized the numerical dissipation of the $\\at$ code for the shearing box problem.  Furthermore, we have exploited the energy conservation property of $\\at$ to carry out a study of energy flow within MRI-driven turbulence.  To compare with previous numerical results, we have investigated the effects of different initial field geometries (uniform or sinusoidal $B_z$), varying domain aspect ratio, and numerical resolution.  In all of our simulations, the MRI is initiated and sustained over many orbits. The time- and volume-averaged properties of the resulting turbulent flow, such as stress levels and magnetic and kinetic energies, are consistent with previous results.  As in previous work, we find that boxes containing net vertical field saturate at higher amplitudes compared to those without net fields.  The total stress is proportional to the magnetic pressure with a constant of proportionality $\\sim 0.5$, but is independent of the gas pressure.  In the net field simulation, the gas pressure increases by a factor of 100, due to thermalization of the turbulence, without affecting the stress.  The consistency of these results with past work indicate that these properties do not result from details of the employed algorithm.  Fourier analysis of the turbulence shows that the largest scales in the box dominate the energetics.  In the presence of a net field, the amplitude of the spatial power spectra is largely independent of resolution on the largest scales.  This is not true for the zero net flux simulations however.  For those simulations, the amplitude decreases as resolution increases, which is consistent with the overall resolution behavior. For net field simulations, the averaged turbulent magnetic and kinetic energies increase slightly with resolution, whereas for the zero net field simulations, the energies decrease with increasing resolution roughly in proportion to the grid zone size.  This apparent lack of convergence for the zero net field shearing box simulations was previously demonstrated by \\cite{from07a} using the ZEUS code.  The net field simulation shows intermittent channel flows which cause temporary increases in stress through amplification of large-scale MRI modes.   The parasitic modes described by \\cite{good94} destroy the channel flow within about one orbit of time, but the rapid increase in stress produces a subsequent increase in thermal energy.  The presence of these discrete channel flow events is a consequence of the box size---larger boxes do not experience them---but we use their presence to study the subsequent energy flow following a rapid increase in stress.  Because $\\at$ evolves the total energy equation, magnetic and kinetic energy losses due to numerical grid-scale effects are added to the internal energy.  This makes $\\at$ well suited to examining the turbulent energy flow and subsequent dissipation.  The recurring channel flows in the net flux model provide a sudden injection of energy into the box by increasing the stress operating on the shearing boundaries of the box.  The injected energy appears as heat after $\\sim 0.2$ orbits. This corresponds to a timescale $\\Omega^{-1}$, which equals $L_z/c_{\\rm s}$ where $c_{\\rm s}$ is the initial soundspeed.  This timescale determines the amplitude of the $\\alf$ speed, $v_{\\rm A}$, and its fundamental MRI wavelength, $\\lambda_{\\rm MRI}$; $L_z/c_{\\rm s} \\sim \\lambda_{\\rm MRI}/v_{\\rm A}$. The timescale is thus on the order of the eddy turnover time, indicating that dissipational heating is a local process and that energy is not carried over large distances before it is thermalized.  In the fiducial zero net magnetic flux simulation, $\\zone$, there are no recurring channel modes, making it more difficult to trace the flow of injected energy.  The analysis is further complicated by the presence of compressive waves that dominate the time derivative of the thermal energy, $\\td$.  These waves are also present in the net field simulations, but their amplitude is smaller relative to the larger turbulent kinetic energy found with a net field.  A detailed examination of the components of the internal energy equation indicate that the compressive waves do not appear to contribute significantly to irreversible heating.  By averaging $\\td$ for the zero net flux simulation, we find a correlation of $\\td$ with $\\ein$ on the same timescale of $\\sim$ 0.2~orbits.  In the net field simulation, the dissipation of magnetic energy is larger than that for the kinetic energy, not unexpected as the ratio of the average magnetic to perturbed kinetic energy is $\\sim 3.4$.  But the ratio of the magnetic to kinetic dissipation rate is roughly constant at $\\sim 1.7$. The fact that the ratio of dissipation rates does not equal the ratio of energies may result from a couple of possibilities.  First, there could be a net transfer of magnetic to {\\it perturbed} kinetic energy as was suggested in \\cite{bran95}.\\footnote{\\S~\\ref{sustained} shows that there is in fact a net transfer of kinetic to magnetic energy.  However, this kinetic energy includes the shear flow, and thus, this result tells us nothing of the energy transfer between magnetic and {\\it perturbed} kinetic energy.} Second, the difference in the ratios could arise from the effective Prandtl number being larger than one.  In particular, if $\\qk \\propto \\nueff \\delta v^2/2$ and $\\qm \\propto \\etaeff B^2/2$, then $(B^2/\\delta v^2)(\\qk/\\qm) \\sim \\pmeff$. With the above values for the energy and dissipation ratios, we find $(B^2/\\delta v^2)(\\qk/\\qm) \\sim 2$, which is consistent with the determination of $\\pmeff$ from the Fourier analysis (see discussion below).  The agreement between the two separate calculations of $\\pmeff$ may be coincidental, but it is suggestive of $\\qk \\propto \\nueff \\delta v^2/2$ and $\\qm \\propto \\etaeff B^2/2$.  The turbulence is sustained by the continued action of the MRI in extracting energy from the differential rotation.  This can be removed from the simulations allowing us to study the decay of the turbulence in detail (simulations $\\ndone$ and $\\zdone$). Figure~\\ref{turb_decay} shows that magnetic losses dominate over kinetic losses during this decay.  In both simulations nearly 50\\% of the magnetic energy and 20\\% of the kinetic energy has been dissipated after 0.2~orbits.  By one orbit into the decay, most of the magnetic and kinetic energy has been lost.  Although these decay timescales arise in a turbulent flow that lacks power input from the MRI, the results are consistent with the conclusion that turbulent energy dissipation occurs on a rapid timescale of order $\\Omega^{-1}$.  \\cite{from07a} used a detailed Fourier analysis (\\S~\\ref{trans_funcs}) to study magnetic energy flow and thermalization as a function of length scale in the shearing box.  In this analysis, the individual terms in the evolution equation for the magnetic energy are examined in Fourier space.  Averaging over time and assuming that the magnetic energy is in a statistical steady state, one sets the sum of these terms equal to a remainder, which is credited to numerical effects. These numerical losses can then be modeled as an effective resistivity (and viscosity for the kinetic energy), allowing one to characterize the numerical dissipation in the simulation.  We repeated their analysis with $\\at$ and extended it to the kinetic energy.  The dominant effect at large scales is the generation of magnetic field by the background shear.  This energy is transferred to other scales by the turbulence.  Net positive field creation by the turbulent flow and energy gains by the transfer between scales only happens at small wavelengths.  This point of transition from loss to gain happens at smaller scales for the zero net field simulation compared to the net field model.  Magnetic dissipation dominates over kinetic dissipation at small scales (i.e., $k L/(2\\pi) \\gtrsim 20$).  Modeling these as an effective resistivity $\\eta$ and viscosity $\\nu$ shows that $\\eta$ and $\\nu$ drop with increasing resolution with a power that lies between first- and second-order in grid resolution.  The effective Prandtl number, on the other hand, is nearly constant as a function of resolution with a value between $\\sim 1.5$ and 2.  \\cite{from07a} observed what they described as ``negative\" resistivity in an analysis restricted to the poloidal field alone.  In repeating their exact analysis with $\\at$, we also observed such an ``anti-dissipation\" at large scales.  This indicates that this effect is not associated with a numerical algorithm limitation associated with ZEUS.  More likely, it arises from the statistical uncertainty at large scales and from the failure of the assumptions that go into the definition of the dissipation term.  We note that the inclusion of the toroidal field $B_y$ in the analysis shows net dissipation at all scales, although again the statistical variation is large at large scales.  In conclusion, what do these results imply for shearing box simulations and the MRI?  First, as observed by  \\cite{from07a}, the scales over which turbulent energy generation occurs are not well-separated from those where there is significant dissipation; the MRI operates over a wide range of scales.  The MRI grows at a rate $\\sim k v_{\\rm A}$ for all $k$ less than $\\Omega/v_{\\rm A}$.  At large scales, a weak field will grow more slowly than the timescale over which energy is transferred between scales, between magnetic and kinetic forms, and ultimately thermalized.  If a field is chopped up by reconnection, it may be reduced to small scales where the MRI no longer operates.  In the presence of a net field, there will always be a significant driving term at the scales set by that imposed field. In the absence of such a field, however, the outcome will be determined by the complex interplay of loss due to dissipation and amplification by the MRI.  In the numerical simulations with zero net field, increasing the resolution causes an overall decrease in the saturation energies. \\cite{from07a} attribute this to higher resolution enabling the MRI to operate at intermediate scales which facilitates the transfer of energy to small scales and promotes reconnection and dissipation. What is perhaps surprising is that resolving the MRI at these scales leads to greater field dissipation than would otherwise be accomplished by the numerical losses that would occur if those scales were underresolved. Because the same effect is observed with both $\\at$ and ZEUS, it seems likely that this ability of the MRI to transfer energy away from the largest scales in the shearing box and to increase the total dissipation is a physical rather than numerical effect.  In related work, \\cite{from07b} and \\cite{lesur07} studied the effect of varying the physical (not numerical) magnetic Prandtl number, $P_m$, on the turbulence.  They found that the saturation amplitudes were increased with increased $P_m$.  \\cite{from07b} found evidence that there exists a critical $P_m > 1$ below which zero net field simulations would die out rather than achieve a steady turbulent state.  Our results in this investigation show that this Prandtl number dependence is a distinct effect from the observed dependence of the turbulence on resolution. We find the numerical $P_m$ to be largely independent of resolution in $\\at$.  Taken together, however, the dependence on physical $P_m$ and the dependence on resolution point to the importance of small and intermediate scale magnetic dissipation and reconnection to establishing saturation amplitudes in MRI-driven turbulence.  As discussed by \\cite{from07a}, numerical dissipation can deviate significantly from physical dissipation.  In \\S~\\ref{fid_z}, we showed that $\\etaeff$ and $\\nueff$ are relatively flat at small scales, suggesting a resemblance to physical dissipation.   However, consider the numerical Reynolds number as calculated from equation~(\\ref{re}) for our zero net flux simulations.  For $N_x$ = 128, we found $\\reeff \\sim$ 12000, and $\\pmeff \\sim$ 1.6 for all of zero net flux simulations.  From the parameter space studies of \\cite{from07b}, these values for the Reynolds and Prandtl numbers correspond to marginal MRI turublence; that is, they lie very close to the critical line between sustained and decaying turbulence.  For $N_x$ = 64, $\\reeff \\sim$ 4100, and the Reynolds number is even smaller for the lower resolutions.  These values are well within the decaying turbulence regime, but we find active MRI turbulence in all of our simulations. These results show that the effective Reynolds and Prandtl numbers of $\\at$ as measured at large wavenumbers does not apply at smaller $k$ values where there are many grid zones per wavelength. Thus, the Reynolds numbers and Prandtl numbers that we calculate should be taken as a measure of the effective numerical dissipation of the code and not equated to a flow with the same Reynolds and Prandtl number as determined by a simple physical resistivity and viscosity.  This result highlights an uncertainty associated with any MRI simulation that depends only on numerical rather than physical dissipation. It is apparent that the numerical Prandtl number can play an important role in determining the ratio of magnetic to kinetic dissipation. More speculatively, the Prandtl number may also play a role in the timescale over which thermalization occurs.  In the present study, we found that both the thermalization timescale and the effective numerical Prandtl number were largely independent of resolution.   However, the turbulent energy thermalization timescales and properties we measure may be subject to change when explicit dissipation is included. It will be a very important next step in this work to include physical dissipation and verify these results.  This work is only the first step in applying $\\at$ to the problem of the energetics of MRI turbulence.  The present study provides a calibration of the numerical dissipation, which will be important in future studies that include explicit resistivity and viscosity.  Furthermore, the unstratified shearing box has the virtue of simplicity and allows a detailed study of MRI turbulence without too many confounding factors, but it also may prove too limited for predictive application to accretion flows.  The inclusion of vertical stratification and radiative cooling are both straightforward extensions to the present study.  The detailed diagnostics developed and applied in this study should prove valuable in this planned work."
    },
    "0806/0806.4314_arXiv.txt": {
        "abstract": "So far radial velocity (RV) measurements have discovered $\\sim~25$ stars to host multiple planets.  The statistics imply that many of the known hosts of transiting planets should have additional planets, yet none have been solidly detected.  They will be soon, via complementary search methods of RV, transit-time variations (TTV) of the known planet, and transits of the additional planet.  When they are found, what can transit measurements add to studies of multiplanet dynamical evolution?  First, mutual inclinations become measurable, for comparison to the solar system's disk-like configuration.  Such measurements will give important constraints to planet-planet scattering models, just as the RV measurements of eccentricity have done.  Second, the Rossiter-McLaughlin effect measures stellar obliquity, which can be modified by two-planet dynamics with a tidally evolving inner planet.  Third, TTV is exquisitely sensitive to planets in mean motion resonance.  Two planets differentially migrating in the disk can establish such resonances, and tidal evolution of the planets can break them, so the configuration and frequency of these resonances as a function of planetary parameters will constrain these processes. ",
        "introduction": "This contribution is about systems with at least two planets, at least one of which transits, and it is regretably a ``what to expect'' talk because no such systems are yet known.  That fact, however, ought to be surprising.  In Figure 1 is plotted systems discovered by radial velocity (RV) which display no transits, rank-ordered by the period of the inner planet (just the first 33 shown), and the same plot for the transiting planets made public before the conference.  For systems in which they are known, planetary companions with full RV orbits are plotted on the same line.  The periods covered by the inner planets are comparable between the samples.  But there is a remarkable difference in the number of companions between the samples: 11 companions are known for RV-only planets; none are known for systems with a transiting planet.  Therefore the first expectation we can draw is that the \\emph{known} transiting planets have quite a few planetary companions that have until now gone undetected.  \\begin{figure}[b] \\begin{center} \\includegraphics[width=3.0in]{pdist.ps} \\caption{ The orbital periods of planetary systems (one system per line), rank-ordered by the period of the inner planet, and separated into two groups.  For RV-only systems, only the 33 with the shortest periods are shown, which compare well to the periods of transiting planets, and thus give us baseline expectations for the (yet undetected) companions to transiting planets. } \\label{fig:pdist} \\end{center} \\end{figure}  One potential bias for this comparison is that in three of the five systems, the inner planet has an $m \\sin i$ more similar to Neptune than to Jupiter, and such planets are not amenable to discovery by transit surveys.  Moreover, those planets may not have even been discovered by RV unless the outer massive planets were intensely observed to constrain their orbits.  Also, lower mass planets may be intrinsically multiple more often (Lovis et al., these proceedings).  On the other hand, dynamical interaction with a transiting planet, as measured by TTV, could in principle reveal companions to much smaller masses than RV.  Instead of trying to unravel such biases, let us proceed with the knowledge that statistics derived from this comparison will only be good at the factor of $\\sim 2$ level. ",
        "conclusions": "The main expectations we have for multiplanet systems with at least one transiting planet are (1) that they exist and should be plentiful among the current crop of transiting exoplanets, (2) there are already methods in use---RV, TTV, transits---that will find them, probably soon, and (3) such systems will be enormously powerful probes of planetary formation and evolution: the theorists are salivating."
    },
    "0806/0806.1531_arXiv.txt": {
        "abstract": "Neutron stars and black holes are the astrophysical systems with the strongest gravitational fields in the universe. In this article, I review the prospect of probing with observations of such compact objects some of the most intriguing General Relativistic predictions in the strong-field regime: the absence of stable circular orbits near a compact object and the presence of event horizons around black-hole singularities. I discuss the need for a theoretical framework within which future experiments will provide detailed, quantitative tests of gravity theories. Finally, I summarize the constraints imposed by current observations of neutron stars on potential deviations from General Relativity. ",
        "introduction": "\\label{sec:int}  Over the past 90 years, the basic ingredients of General Relativity have been tested in many different ways and in many different settings. From the solar eclipse expedition of 1917 to the modern observations of double neutron stars, General Relativity has passed all tests with flying colors~\\cite{Will06}. Yet, our inability to devise a renormalizable quantum gravity theory as well as the mathematical singularities found in many solutions of Einstein's equations suggest that we should look harder for gravitational phenomena not described by General Relativity.  The search for such deviations has been very fruitful in the regime of very weak fields. Observations of high-redshift supernovae~\\cite{Perl97, Riess98} and of the cosmic microwave background with WMAP~\\cite{Spergel03} have measured a non-zero cosmological constant (or a slowly rolling field that behaves as such at late times). This discovery can be incorporated within the framework of General Relativity, if interpreted simply as a constant in the Einstein--Hilbert action. It nevertheless brought to the surface a major problem in trying to connect gravity to basic ideas of quantum vacuum fluctuations~\\cite{Weinberg89,Carroll01}.  In the strong-field regime, on the other hand, which is relevant for the evolution of the very early universe and for determining the properties of black holes and neutron stars, little progress has been made in testing the predictions of general relativity~\\cite{Stairs03}. There are two reasons that have been responsible for this lag. First, phenomena that occur in strong gravitational fields are complex and often explosive, making it very difficult to find observable properties that depend cleanly on the gravitational field and that allow for quantitative tests of gravity theories. Second, there exists no general theoretical framework within which to quantify deviations from general relativistic predictions in the strong-field regime.  During the current decade, technological advances and increased theoretical activity have led to developments that promise to make strong-field gravity tests a routine in the near future. The first generation of earth-based gravitational wave observatories (such as LIGO~\\cite{LIGO}, GEO600~\\cite{GEO}, TAMA300~\\cite{TAMA}, and VIRGO~\\cite{VIRGO}) as well as the Beyond Einstein Missions (such as Constellation-X, LISA, and the Black Hole Imager~\\cite{beyondEins}) will offer an unprecedented look into the near fields of black holes and neutron stars. Moreover, recent ideas on quantum gravity~\\cite{Burgess04}, brane-world gravity~\\cite{Maartens04}, or other Lagrangian extensions of general relativity~\\cite{Woodard06,Sotiriou08} will provide the means with which the experimental results will be interpreted.  In this article, I review the theoretical and experimental prospects of testing strong-field General Relativity with observations in the electromagnetic spectrum. In the first few sections, I discuss the motivation for performing such tests and then describe the astrophysical settings in which strong-field effects can be measured. In Section~\\ref{sec:need}, I elaborate on the need for a theoretical framework within which strong-field gravity tests can be performed and in Section~\\ref{sec:tests} I review the current quantitative tests of General Relativity in the strong-field regime that use neutron stars. Finally, in Section~\\ref{section:beyond} I discuss the prospect of probing and testing strong gravitational fields with upcoming experiments and observatories.   \\newpage ",
        "conclusions": ""
    },
    "0806/0806.4122_arXiv.txt": {
        "abstract": "Formation of massive stars by accretion requires a high accretion rate of $\\dot{M}_\\ast > 10^{-4}~M_{\\odot}/{\\rm yr}$ to overcome the radiation pressure barrier of the forming stars. Here, we study evolution of protostars accreting at such high rates, by solving the structure of the central star and the inner accreting envelope simultaneously. The protostellar evolution is followed starting from small initial cores until their arrival at the stage of the Zero-Age Main Sequence (ZAMS) stars. An emphasis is put on evolutionary features different from those with a low accretion rate of $\\dot{M}_\\ast \\sim 10^{-5}~M_{\\odot}/{\\rm yr}$, which is presumed in the standard scenario for low-mass star formation. With the high accretion rate of $\\dot{M}_\\ast \\sim 10^{-3}~M_{\\odot}/{\\rm yr}$, the protostellar radius becomes very large and exceeds $100~R_{\\odot}$. Unlike the cases of low accretion rates, deuterium burning hardly affects the evolution, and the protostar remains radiative even after its ignition. It is not until the stellar mass reaches $\\simeq 40~M_{\\odot}$ that hydrogen burning begins and the protostar reaches the ZAMS phase, and this ZAMS arrival mass increases with the accretion rate. These features are similar to those of the first star formation in the universe, where high accretion rates are also expected, rather than to the present-day low-mass star formation. At a very high accretion rate of $> 3 \\times 10^{-3}~M_{\\odot}/{\\rm yr}$, the total luminosity of the protostar becomes so high that the resultant radiation pressure inhibits the growth of the protostars under steady accretion before reaching the ZAMS stage. Therefore, the evolution under the critical accretion rate $3 \\times 10^{-3}~M_{\\odot}/{\\rm yr}$ gives the upper mass limit of possible pre-main-sequence stars at $\\simeq 60~M_{\\odot}$. The upper mass limit of MS stars is also set by the radiation pressure onto the dusty envelope under the same accretion rate at $\\simeq 250~M_{\\odot}$. We also propose that the central source enshrouded in the Orion KL/BN nebula has effective temperature and luminosity consistent with our model, and is a possible candidate for such protostars growing under the high accretion rate. ",
        "introduction": "\\label{sec:intro}  Massive ($ > 8~M_{\\odot}$) stars make significant impacts on the interstellar medium via various feedback processes:, e.g., UV radiation, stellar winds, and supernova explosions. Such feedback processes sometimes trigger or regulate nearby star-formation activity. Their thermal and kinetic effects are important factors in the phase cycle of the interstellar medium. Strong coherent feedback by many massive stars causes galactic-scale dynamical phenomena, such as galactic winds. Furthermore, massive stars dominate the light from distant galaxies. The cosmic star formation rate is mostly estimated with the light observed from massive stars \\citep[e.g.,][]{Md96, HB06}. Despite such importance, the question of the formation process of massive stars still remains open. As for lower-mass stars, there is a widely-accepted formation scenario, where gravitational collapse of molecular cloud cores leads to subsequent mass accretion onto tiny protostars \\citep[e.g.,][]{SAL87}. If one directly applies this scenario to massive star formation, however, some difficulties arise in the main accretion phase \\citep[e.g., see][for recent reviews]{ZY07, MO07}. The main difficulty in the formation of massive stars is the very strong radiation pressure acting on a dusty envelope. The radiative repulsive effect becomes quite strong at the dust destruction front, where the accretion flow gets much of the outward momentum of radiation. Therefore, a necessary condition for massive star formation is to overcome this barrier at the dust destruction front. However, some theoretical work has shown that the accretion rate of $\\dot{M}_\\ast \\sim 10^{-5}~M_{\\odot}/{\\rm yr}$ expected of lower-mass protostars is too low to overcome the barrier \\citep[e.g.,][hereafter WC87]{WC87}. Because of this difficulty, various scenarios different from the standard accretion paradigm have been proposed such as stellar mergers \\citep[e.g.,][]{SPH00} and competitive accretion \\citep[e.g.,][]{BBZ98}.  An origin of this disputed situation is uncertainty concerning the initial condition for massive star formation. This uncertainty partly originates from the scarcity of massive stars; most of the massive star forming regions are distant. In addition, the initial condition is easily disrupted by feedback from newly formed massive stars. Since the accretion rate reflects the thermal state of the original molecular core, it is still uncertain which rates should be favored for growing high-mass protostars. In fact, some observations of young high-mass sources suggest high accretion rates. Signatures of infall motion are detected with various line observations toward high-mass protostellar objects (HMPOs) and hyper-/ultra-compact H~II regions, and the derived accretion rates are $10^{-4} - 10^{-3}~M_{\\odot}/{\\rm yr}$ \\citep[e.g.,][]{Fl05, Bl06, KW06}. Similar high accretion rates are also inferred by SED fitting of hot cores \\citep{OLD99} and HMPOs \\citep{Fz07, KG08}. Molecular outflows are also ubiquitous in high-mass star-forming regions, and estimated high mass outflow rates also suggest high accretion rates \\citep[e.g.,][]{B02, Zh05}. Such high accretion rates have the advantage of overcoming the radiation pressure barrier \\citep{LS71, KH74}. Theoretically, \\citet{Nk00} have suggested a protostar growing at the very high accretion rate of $\\sim 10^{-2}~M_{\\odot}/{\\rm yr}$ to explain the low radiation temperature of scattered light from Orion KL region \\citep{Mn98}. \\citet{MT02, MT03} have predicted that massive molecular cloud cores, from which a few massive stars form, should be dominated by supersonic turbulence, envisioning that such cores, if they exist, are embedded in a high-pressure environment. \\citet{KKM07} have simulated collapse of the turbulent cores performing radiation-hydrodynamical calculations, and demonstrated that the accretion rates attain more than $10^{-4}~M_{\\odot}/{\\rm yr}$ in their simulations. Some recent observations are getting a close look at the initial condition of massive star formation, and have found strong candidates for high-mass pre-stellar cores \\citep[e.g.,][]{RSJ07, Mt07}. Further detailed observations will verify the scenario.  The main targets of this paper are protostars growing at the high accretion rates of $\\dot{M}_\\ast \\geq 10^{-4}~M_{\\odot}/{\\rm yr}$. Detailed modeling of such protostars should be useful for future high-resolution observations and simulations of massive star formation. However, previous studies on such protostars are fairly limited. The protostellar evolution has been well studied for low-mass ($< 1~M_{\\odot}$) and intermediate-mass ($< 8~M_{\\odot}$) protostars by detailed numerical calculations solving the stellar structure, but these studies have focused on evolution at the low accretion rate of $\\sim 10^{-5}~M_{\\odot}/{\\rm yr}$ (e.g., Stahler, Shu \\& Taam 1980a,b, hereafter SST80a,b; Palla \\& Stahler 1990, 1991, hereafter PS91, 1992; Beech \\& Mitalas 1994) Maeder and coworkers have calculated the protostellar evolution under accretion rates growing with the stellar mass, which finally exceeds $10^{-4}~M_{\\odot}/{\\rm yr}$ \\citep{BM96, BM01, NM00}, but it is still uncertain how the evolution changes with accretion rates, and which physical mechanisms cause differences. Some authors have used polytropic one-zone models originally invented for low accretion rates even at such a high rate \\citep[e.g.,][]{Nk00, MT03, KT07}, but its validity remains obscure owing to the lack of detailed calculations. By coincidence, the similar high accretion rates are expected for protostars forming in the early universe \\citep[e.g.,][]{ON98, Ys06}. This simply reflects high temperatures in primordial pre-stellar clouds. Evolution of such primordial protostars has been studied by \\citet[][hereafter SPS86]{SPS86} and \\citet{OP01, OP03}.  Therefore, our other motivation is to investigate how the protostellar evolution changes with metallicities.  To summarize, our goal in this paper is to answer the following questions; \\begin{itemize} \\item What are the properties of massive protostars (e.g., radius, luminosity, and effective temperature) growing at high accretion rates? Can we observe any signatures of their characteristic properties? \\item How different are high-mass protostars from low- and intermediate-mass ones, for which lower accretion rates are presumed? How about the difference from the protostellar evolution of primordial stars? Furthermore, if the protostellar evolution significantly varies with accretion rates and metallicities, what causes such differences? \\item What are the consequences of protostellar evolution with high accretion rates for the formation and feedback processes of massive stars? How massive can a star get before its arrival at the Zero-Age Main Sequence (ZAMS)? What is the maximum mass of stars that can be formed? \\end{itemize} In order to answer these questions, we solve the structure of accreting protostars with various accretion rates and metallicities with detailed numerical calculations \\citep[also see][for a similar effort]{YB08}.  The organization of this paper is as follows: In \\S~\\ref{sec:num}, we briefly review the basic procedure to construct numerical models of accreting protostars and their surrounding envelopes. The subsequent \\S~\\ref{sec:result} is the main part of this paper, where our numerical results are presented. First, we investigate the protostellar evolution of two fiducial cases with the accretion rates of $10^{-3}$ and $10^{-5}~M_{\\odot}/{\\rm yr}$ in \\S~\\ref{ssec:md_1em3} and \\S~\\ref{ssec:md_1em5}. After that, we show more general variations of protostellar evolution, e.g., mass accretion rates in \\S~\\ref{ssec:md_dep}, and metallicities in \\S~\\ref{ssec:metal}. In \\S~\\ref{ssec:edd}, we present the protostellar evolution at very high accretion rates exceeding $10^{-3}~M_{\\odot}/{\\rm yr}$, and show that the steady accretion is limited by a radiation pressure barrier acting on a gas envelope. We further discuss some implications based on our numerical results. In \\S~\\ref{sec:stopacc}, we examine which feedback effects can affect the growth of protostars for a given accretion rate. We also explore some observational possibilities of detecting signatures of the supposed high accretion rates in \\S~\\ref{sec:obs}. Finally, \\S~\\ref{sec:sum} is assigned to summary and conclusions. ",
        "conclusions": "\\label{sec:sum}  We have studied evolution of accreting protostars by solving the structure of the central growing stars and the surrounding accreting envelopes simultaneously. Particular attention is paid to cases with high accretion rates of $\\dot{M}_\\ast \\geq 10^{-4}~M_{\\odot}/{\\rm yr}$, which are envisaged in some current scenarios of massive star formation. The protostellar evolution at a high mass accretion rate of $\\dot{M}_\\ast = 10^{-3}~M_{\\odot}/{\\rm yr}$ (run MD3 in Table 1) can be summarized as follows:  The entire evolution is divided into four characteristic phases; (I) adiabatic accretion ($M_{\\ast} \\lesssim 6~M_{\\odot}$), (II) swelling ($6~M_{\\odot} \\lesssim M_{\\ast} \\lesssim 10~M_{\\odot}$), (III) KH contraction ($10~M_{\\odot} \\lesssim M_{\\ast} \\lesssim 30~M_{\\odot}$), and (IV) main-sequence accretion ($M_{\\ast} \\gtrsim 30~M_{\\odot}$) phases.  A main driver of transition between the evolutionary phases above is a decrease in opacity in the stellar interior with increasing temperature, and thus protostellar mass. Due to the fast accretion, the matter accreted onto the stellar surface is embedded into the interior before radiatively losing the entropy produced at the accretion shock. Early in the evolution, the opacity in the stellar interior, which is mainly by free-free absorption, is very high owing to low temperature. As a result of the high opacity, the star keeps a large amount of entropy imported by the accreted matter. This is the adiabatic accretion phase (I). With the decrease in opacity, the entropy is transported outward gradually. When the matter near the surface temporarily has a large amount of entropy, the star expands as large as $\\simeq 200~R_{\\odot}$. This is the swelling phase (II). After that, all parts of the protostar lose energy and the protostar contracts. This is the KH contraction phase (III). With contraction, the temperature increases and the hydrogen burning eventually begins. When the nuclear burning provides sufficient energy to balance with that lost from the star radiatively, the star stops the KH contraction and reaches the zero-age main-sequence (ZAMS) phase at $\\simeq 30~M_{\\sun}$ (IV).  Deuterium burning plays little role in the evolution. Even if the deuterium burning is turned off by hand, the result hardly changes. The swelling (II) is caused by outward transfer of entropy within the star, which is observed as a propagating luminosity wave.  In general, at the higher accretion rate, the protostellar radius is larger and then the maximum temperature is lower at the same stellar mass, owing to the higher entropy within the protostar. As a result of the lower temperature, the onset of the nuclear burning is postponed to the higher protostellar mass.  For very high accretion rate $\\dot{M}_\\ast > 3 \\times 10^{-3}~M_{\\odot}/{\\rm yr}$, in the course of the KH contraction, the radiation pressure onto the inner accreting envelope becomes so strong that the flow is retarded before hitting the stellar surface. The reduced ram pressure onto the stellar surface causes abrupt expansion of the star. The steady-state accretion is not possible thereafter. At this moment, the accretion might be halted, or might continue at a reduced rate or in a sporadic way. In either case, the evolution at the critical accretion rate of $M_{\\rm cr} = 3 \\times 10^{-3}~M_{\\odot}/{\\rm yr}$ gives the upper mass limit of PMS stars at $\\simeq 60~M_\\odot$. Further growth of the star should be finally limited by the radiation pressure onto the outer dusty envelope. We have found this limit with $M_{\\rm cr}$ gives the maximum mass of MS stars at $\\simeq 250~M_\\odot$.  Such a high accretion rate has also been expected and studied for the first star formation in the universe. The evolutions of primordial and solar-metallicity protostars are very similar at the same accretion rate. However, the lower opacity of the primordial gas results in the earlier transitions between the evolutionary phases above. For example, with $\\dot{M}_\\ast = 10^{-3}~M_{\\odot}/{\\rm yr}$, the accreting star enters the KH contraction phase at $10~M_{\\sun}$ in the solar metallicity case, while it enters at $7M_{\\sun}$ in the primordial case.  Distinguishing features of those massive protostars accreting at high rates include the large stellar radius sometimes exceeding $100~R_{\\odot}$, and then low color temperature $\\simeq 6000$K. A massive protostar in Orion BN/KL nebula indeed has such features and is a possible candidate of such objects.  {"
    },
    "0806/0806.0121_arXiv.txt": {
        "abstract": "Massive protostars dramatically influence their surroundings via accretion-induced outflows and intense radiation fields. They evolve rapidly, the disk and infalling envelope being evaporated and dissipated in $\\sim$~ 10$^5$ years. Consequently, they are very rare and investigating this important phase of early stellar evolution is extremely difficult. Here we present the discovery of a key transient phase in the emergence of a massive young star, in which ultraviolet radiation from the new-born giant has just punctured through its natal core. The massive young stellar object AFGL 961 II is readily resolved in the near infrared. Its morphology closely resembles a cat's eye and is here dubbed as the Rosette Eye. Emerging ionized flows blow out an hourglass shaped nebula, which, along with the existence of strong near-infrared excess, suggests the existence of an accretion disk in the perpendicular direction. The lobes of the hourglass, however, are capped with arcs of static H$_{2}$ emission produced by fluorescence. This study has strong implications for our understanding of how massive stars embark on their formation. ",
        "introduction": "Protostellar objects are, in most cases, deeply embedded in molecular clouds and are enshrouded by heavy foreground extinction. This makes the very early stages of stellar gestation notoriously illusive~(Zinnecker \\& Yorke 2007). The birth of stars with masses above 10~M$_\\odot$ is particularly intriguing as their radiation is apparently sufficient to resist the accretion of gas and, hence, further mass growth unless somehow mitigated~(Mckee \\& Tan 2002; Behrend \\& Maeder 2001). Evidence of potential disks and/or envelopes associated with massive star formation has been accumulated through various probes especially in the radio domain~(Chini et al. 2004; Patel et al. 2005). However, the formation process of massive stars remains far from being resolved. The critical growth period of massive stars lasts only tens of thousand of years but is usually accompanied by spectacular ejections of gas in opposite directions. Since such jets are circumstantial evidence of an accretion disk, it is possible and crucial to gain robust evidence of ongoing accretion associated with massive young stellar objects (YSO) in their early stages of evolution.  The Rosette Molecular Complex (RMC) is a famous isolated massive star forming region with an extent of about 100 pc. It is located at a distance of $\\sim$~1.39 kpc (Hensberge et al. 2000) and harbors a gas reservoir of $\\sim$10$^5~M_\\odot$(Blitz \\& Thaddeus 1980). New generation OB star formation is in evidence in the densest ridge of the Complex (Li \\& Smith, 2005). The well known high-mass protostellar system AFGL 961 is situated well within this region (Cohen 1973; Grasdalen et al. 1983; Castelaz et al. 1985; Lenzen et al. 1984; Hodapp 1994). The third component of the system, designated as AFGL 961 II (Li \\& Smith 2005), was first noticed to be associated with a small nebulosity by Eiroa (1981) and later briefly discussed by Hodapp (1994). The nebulosity was considered as a cavity with a partial shell surrounding the central young star~(Aspin 1998; Alvarez et al. 2004). However, the origin of the intriguing YSO is far from clear. ",
        "conclusions": "This study substantiates that the exciting source of the Rosette Eye, AFGL 961 II, has just finished its original collapse, that UV radiation from shocked massive accretion has recently turned on and that ionized stellar winds have begun to emerge in the polar directions. The shell left by the original collapse is now prepared to face the UV ionization from the newly born star. We suggest that the Eye may closely correspond to a transient phase immediately preceding that of the famous Orion OMC-1 outflow~(Stone et al. 1995). In this scenario, the arcs will be subjected to fluid instabilities as an ensuing fast wind of low density emerges, which leads to the so-called fireworks as the shell gas is driven out in the form of dense bullets~(McCaughrean \\& Mac Low, 1997). The results presented in this paper corroborate the recent onset of the formation of a massive star in the ridge of the RMC, which is in a key transient phase of its emergence from the natal cloud and the development of the HII region."
    },
    "0806/0806.0067_arXiv.txt": {
        "abstract": "We report the discovery of three very late T~dwarfs in the UKIRT Infrared Deep Sky Survey (UKIDSS) Third Data Release: ULAS~J101721.40+011817.9 (ULAS1017), ULAS~J123828.51+095351.3 (ULAS1238) and ULAS~J133553.45+113005.2 (ULAS1335). We detail optical and near-infrared photometry for all three sources, and mid-infrared photometry for ULAS1335. We use near-infrared spectra of each source to assign spectral types T8p (ULAS1017), T8.5 (ULAS1238) and T9 (ULAS1335) to these objects. ULAS1017 is classed as a peculiar T8 (T8p) due to appearing as a T8~dwarf in the $J$-band, whilst exhibiting $H$~and $K$-band flux ratios consistent with a T6 classification. Through comparison to BT-Settl model spectra we estimate that ULAS1017 has $750 {\\rm K} \\ltsimeq T_{\\rm eff} \\ltsimeq 850 {\\rm K}$, and $5.0 \\ltsimeq \\log {\\it g }(cms^{-2}) \\ltsimeq  5.5$, assuming solar metallicity. This estimate for gravity is degenerate with varying metallicity. We estimate that ULAS1017 has an age of 1.6--15~Gyr, a mass of 33--70~M$_J$ and lies at a distance of 31--54~pc. We do not estimate atmospheric parameters for ULAS1238 due to a lack of $K$-band photometry. We extend the unified scheme of \\citet{burgasser06} to the type T9 and suggest the inclusion of the $W_J$ index to replace the now saturated $J$-band indices. We propose ULAS1335 as the T9 spectral type standard. ULAS1335 is the same spectral type as ULAS~J003402.77-005206.7 and CFBDS~J005910.90-011401.3. We argue that given the similarity of the currently known $>$T8~dwarfs to the rest of the T~dwarf sequence, the suggestion of the Y0 spectral class for these objects is premature. Comparison of model spectra with that of ULAS1335 suggest a temperature below 600K, possibly combined with low-gravity and/or high-metallicity. We find ULAS1335 to be extremely red  in near to mid-infrared colours, with $H -[4.49]~=~4.34 \\pm 0.04$ . This is the reddest near to mid-infrared colour yet observed for a T~dwarf. The near to mid-infrared spectral energy distribution of ULAS1335 further supports $T_{\\rm eff} < 600{\\rm K}$, and we estimate $T_{\\rm eff} \\sim 550-600{\\rm K}$ for ULAS1335. We estimate that ULAS1335 has an age of 0.6--5.3~Gyr, a mass of 15--31~M$_J$ and lies at a distance of 8--12~pc. ",
        "introduction": "\\label{sec:intro} Extending the known sample of field dwarfs to ever lower effective temperatures ($T_{\\rm eff}$) is important not only for the determination of the field mass function, but also for probing temperature and pressure regimes that have hitherto been unexplored observationally. The study of extremely cool brown dwarfs opens a window on atmospheric physics which will be fundamental for understanding the processes within a broad range of substellar atmospheres, including those of giant exoplanets.   At the time of writing there are 142 T~dwarfs published in the literature \\citep[e.g., DwarfArchives.org; ][]{pinfield08}. Of these, only five are classified as T8 or later \\citep[using the scheme of][]{burgasser06}, two of which have spectral types later than T8: ULAS~J003402.77-005206.7 \\citep[hereafter ULAS0034; ][]{warren07} and CFBDS~J005910.90-011401.3 \\citep[hereafter CFBDS0059; ][]{delorme08}. ULAS0034 has an inferred $T_{ \\rm eff}$ of 600-700~K \\citep{warren07}, and \\citet{delorme08} infer a $T_{ \\rm eff}$ for CFBDS0059 that is $\\sim 50$~K cooler than that of ULAS0034.   Expanding the sample of very cool brown dwarfs, and exploring possible factors that could motivate the implementation of new spectral classes \\citep[e.g Y dwarfs; ][]{kirkpatrick99} are primary science drivers for the UKIRT Infrared Deep Sky Survey (UKIDSS) Large Area Survey \\citep[LAS; see][]{ukidss}. Our collaboration, the UKIDSS Cool Dwarf Science Working Group (CDSWG) is engaged in a substantial effort to achieve this aim \\citep[e.g.,][]{pinfield08,kendall07,lod07,warren07}. Here we present the recent discovery by CDSWG of three very cool brown dwarfs. These are ULAS~J101721.40+011817.9 (ULAS1017), ULAS~J123828.51+095351.3 (ULAS1238) and ULAS~J133553.45+113005.2 (ULAS1335). The last of these, ULAS1335, may be the coolest brown dwarf yet discovered and we here discuss the properties of this object, estimating a $T_{\\rm eff}$ of approximately 550--600~K. ",
        "conclusions": "\\label{sec:summ}  We have identified three very late-type T~dwarfs in the UKIDSS LAS DR3: ULAS1017, ULAS1238 and ULAS1335, for which we have adopted the spectral types T8p, T8.5 and T9 respectively. ULAS1017 has a peculiar spectrum, with a $J$-band typical of a T8 dwarf, but $H$~and $K$-band peaks of a T6. In assigning spectral types to the our targets, we have defined the extension of the T-dwarf sequence to the type T9, using the H$_2$O--$H$ and $W_J$ indices, and ULAS1335 as the spectral standard.   To estimate atmospheric parameters for ULAS1017 and ULAS1335, we have performed a detailed comparison with BT-Settl model spectra. For ULAS1017, we have estimated that its temperature lies in the range $750 \\leq T_{\\rm eff} \\leq 850$, with $\\log {\\it g} = 5.0-5.5$ (assuming solar metallicity).   Our model comparison for ULAS1335 has suggested a $T_{\\rm eff}$ that is cooler than the coolest available BT-Settl models ($T_{\\rm eff} < 600{\\rm K}$), and it may have low surface gravity possibly combined with high metallicity. As a result we have based our ``best-guess'' of the atmospheric parameters on an examination of the near- to mid-infrared SED. We have found ULAS1335 to be the reddest T-dwarf yet identified, with $H-{\\rm [4.49]}~=~4.34 \\pm 0.04$. By extrapolating the empirical $T_{\\rm eff}$ -- $H$--[4.49] relation to lower $T_{\\rm eff}$, and guided by model trends, we have estimated that ULAS1335 has $T_{\\rm eff} \\sim 550-600{\\rm K}$. This estimate suggests that ULAS1335 is the coolest brown dwarf yet discovered, although we await the determination of its distance by parallax with interest."
    },
    "0806/0806.0298_arXiv.txt": {
        "abstract": "We have carried out a search for the 2.14 \\micron\\ spectroscopic signature of the close orbiting extrasolar giant planet, HD 179949b. High cadence time series spectra were obtained with the CRIRES spectrograph at VLT1 on two closely separated nights. Deconvolution yielded spectroscopic profiles with mean S/N ratios of several thousand, enabling the near infrared contrast ratios predicted for the \\hd\\ system to be achieved.  Recent models have predicted that the hottest planets may exhibit spectral signatures in emission due to the presence of TiO and VO which may be responsible for a temperature inversion high in the atmosphere. We have used our phase dependent orbital model and tomographic techniques to search for the planetary signature under the assumption of an absorption line dominated atmospheric spectrum, where T and V are depleted from the atmospheric model, and an emission line dominated spectrum, where TiO and VO are present.  We do not detect a planet in either case, but the \\hbox{$2.120$\\ \\micron\\,-\\,$ 2.174$\\ $\\micron$} wavelength region covered by our observations enables the deepest near infrared limits yet to be placed on the planet/star contrast ratio of any close orbiting extrasolar giant planet system. We are able to rule out the presence of an atmosphere dominated by absorption opacities in the case of HD 179949b at a contrast ratio of $F_p/F_* \\sim 1/3350$, with 99 per cent confidence. ",
        "introduction": "\\protect\\label{section:intro}  Recent models \\citep{burrows08cegp} have focused on the role of a parametrized stratospheric absorbing molecular species which leads to abundant molecules such as H$_2$O appearing in emission in certain wavelength bands. Similarly, \\cite{fortney08unified} have more specifically investigated the effects of TiO and VO which absorb much of the incoming stellar radiation high in the atmospheres of the hottest CEGPs. \\cite{fortney08unified} expect that \\hdb\\ belongs to this class of hot CEGPs where the absorbing species lead to a temperature inversion and the formation of a stratosphere. It is claimed that such models give better agreement with multi-wavelength mid-infrared Spitzer observations for those CEGPs which are most highly irradiated. Systems such as HD 209458b \\citep{burrows07hd209458b,knutson08hd209458b}, HD 149026b and $\\upsilon$ And are among those systems which possess atmospheres most consistent with the presence of temperature inversions.  In the near infrared, the strong \\hbox{2.2 \\micron}\\ bump due to the presence of strong H$_2$O and CO molecular bands for \\hbox{$2.2\\ \\micron\\ < \\lambda > 2.2\\ \\micron$} may not be present in systems which exhibit a stratosphere \\citep{burrows08cegp}. This possibility was indicated by observations of \\hbox{HD 209458b} \\citep{richardson03} which failed to detect the \\hbox{2.2 \\micron}\\ bump in the atmosphere of \\hbox{HD 209458b}. The weak H$_2$O and CO absorption transitions close to \\hbox{2.2 \\micron}\\ are rather seen in emission while the 2.2 \\micron\\ $F_p/F_*$ flux ratio is lower and flatter owing to a spectral energy distribution which more closely resembles that of a blackbody.   \\begin{table*} \\caption{CRIRES/VLT1 observations of \\hd\\ for UT 2007 July 26 and August 3. } \\protect\\label{tab:journal} \\vspace{5mm} \\begin{center} \\begin{tabular}{lcccccl} \\hline Object\t\t\t& UT start of\t& UT start of\t& Time per\t  & Number of\t\t& Number of \t& Comments  \\\\ & first frame\t& last frame\t& exposure [secs] & co-adds per frame\t& observations\t& \t    \\\\ \\hline \\multicolumn{7}{c}{UT 2007 July 26/27} \\\\ \\hline HD 179949\t\t& 23:51:46\t& 00:06:52\t&\t25\t  & \t4\t\t& 8\t\t&  Main F8V Target\t \\\\ HD 182645\t\t& 00:18:04\t& 00:24:35\t&\t50\t  & \t2\t\t& 4\t\t&  B7IV standard\t\\\\ HD 179949\t\t& 00:39:49\t& 03:26:40\t&\t25\t  & \t4\t\t& 76\t\t&  Main F8V Target\t \\\\ HD 158643\t\t& 00:40:07\t& 03:51:41\t&\t10\t  & \t6\t\t& 4\t\t&  A0V Standard\t \t\\\\ HD 179949\t\t& 04:04:28\t& 07:54:52\t&\t25\t  & \t4\t\t& 104\t\t&  Main F8V Target\t \\\\ HD 212581\t\t& 08:09:05\t& 08:35:48\t&\t50\t  & \t4\t\t& 8\t\t&  B9.5V Standard\t \\\\ HD 212301\t\t& 08:48:28\t& 09:51:06\t&\t50\t  & \t4\t\t& 16\t\t&  F8V target\t \t\\\\ HD 212581\t\t& 10:05:19\t& 10:10:31\t&\t12\t  & \t6\t\t& 4\t\t&  B9.5 Standard \t\\\\ \\hline \\multicolumn{7}{c}{UT 2007 Aug 03} \\\\ \\hline HD 158643\t\t& 00:18:55\t& 00:23:17\t&\t15\t  & \t4\t\t& 4\t\t&  A0V Standard\t\t\\\\ HD 179949\t\t& 01:08:19\t& 02:44:18\t&\t25\t  & \t4\t\t& 44\t\t&  Main F8V Target\t\\\\ HD 158643\t\t& 03:05:54\t& 03:17:00\t&\t10\t  & \t6\t\t& 8\t\t&  A0V Standard\t\t\\\\ HD 179949\t\t& 03:31:11\t& 07:22:23\t&\t25\t  & \t4\t\t& 104\t\t&  Main F8V Target\t\\\\ HD 212581\t\t& 07:39:34\t& 07:53:31\t&\t15\t  & \t6\t\t& 8\t\t&  B9.5V Standard\t\\\\ HD 212581\t\t& 07:56:14\t& 08:36:11\t&\t20\t  & \t6\t\t& 16\t\t&  B9.5V Standard\t\\\\ HD 212301\t\t& 08:49:44\t& 09:35:04\t&\t50\t  & \t4\t\t& 12\t\t&  F8V target\t\t\\\\ HD 212581\t\t& 09:44:38\t& 10:24:05\t&\t20\t  & \t6\t\t& 16\t\t&  B9.5V Standard\t \\\\  \\hline \\end{tabular} \\end{center} \\end{table*}  \\subsection{HD 179949 and its planet} \\protect\\label{section:HD179949}  The presence of a close orbiting planetary system around the \\hbox{$M_v = 4.09 \\pm 0.04$}\\ F8 dwarf star, \\hd\\ was first reported by \\cite{tinney01hd179949}. \\cite{eggenberger07} included \\hd\\ in a search for stellar duplicity but found no companion star, indicating it to be a single star system. Further, a survey of stars known to harbour planets \\citep{santos04metallicity} revealed that the [Fe/H] $= 0.22 \\pm 0.05$ dex for \\hdb\\ places it in the most common metallicity band, in a distribution significantly more metal rich than than for stars which are not known to harbour planets. The reported velocity amplitude induced by the planet on its parent star, at \\hbox{$K = 101.3 \\pm 3.0$ \\kms}, with orbital period of $3.093 \\pm 0.001$ d\\ and orbital radius of $0.045 \\pm 0.004$   placed it among the closest orbiting planets. More recently, \\cite{butler06catalogue} have published revised system parameters as a result of further monitoring of the system, finding $K = 112.6 \\pm 1.8$ \\kms, $P = 3.092514 \\pm 0.000032$ d and $a = 0.0443 \\pm 0.0026$ AU in a low eccentricity ($e = 0.022 \\pm 0.015$) orbit. By studying emission in the Ca {\\sc ii} H \\& K lines, \\cite{shkolnik03hd179949} found evidence for perturbation of the stellar magnetic field with a periodicity which coincides with the orbital period of \\hdb\\ while \\cite{wolf04hd179949} found the stellar rotation period of \\hd\\ to be independent of the planetary orbital period, with $P_{rot} = 7.06549 \\pm 0.00061$ d. With no transit of the system reported, a minimum mass of $M sin i = 0.916 \\pm 0.076$ \\msun\\ is found. The time of inferior conjunction is given by an ephemeris of \\hbox{HJD = $2451001.510 \\pm 0.020$} d \\citep{butler06catalogue}. \\cite{wittenmyer07} found no evidence for long period objects in the \\hd\\ system and reported updated system parameters which are consistent with those of \\cite{butler06catalogue} at the 1-$\\sigma$ level.   \\cite{cowan07hotnights} presented mid-infrared Spitzer observations of the \\hd\\ system which revealed a phase dependent light curve in phase with the planet's orbit and with a relative peak-to-trough amplitude of 0.00141 at \\hbox{8 \\micron}. This implies that less than 21 per cent of the incident stellar radiation is recirculated to the night side of the planet and contrasts with the other systems in their study, HD 209458 and 51 Peg, which did not reveal photometric variation, suggesting a higher level of redistribution of incident energy. While these observations  were used by \\cite{burrows08cegp} to compare model fits, the data carry sufficiently high uncertainties that they enable a number of models with varying degrees of heat redistribution, absorber opacity and inclination to give reasonable fits. A hint that this degeneracy may be partially broken comes from the similarity between \\hdb\\ and $\\upsilon$ And b \\citep{harrington06}, both showing only a small shift between the superior conjunction ephemeris and lightcurve maximum. This may indicate that heat is re-radiated without being carried downstream, as suggested by the results of \\cite{harrington06,harrington07}, due to the presence of a hot stratosphere which inhibits advection in the lower atmosphere.   In this paper, we present 2.14 \\micron\\ observations taken with the Cryogenic high-resolution InfraRed \\'{E}chelle Spectrograph (CRIRES) at the Very Large Telescope and search for the signature of a planet which exhibits either absorption or emission features. We first describe observations and data extraction procedures (\\S \\ref{section:obs}) before describing the method (\\S \\ref{section:analysis}) which involves searching for the faint planetary spectrum in a mean spectrum subtracted timeseries of spectra. A Gaussian matched filter mimics the radial velocity motion and varying strength of the planetary spectrum which is used to search for the best fitting model. The results are presented in \\S \\ref{section:results} and discussed in \\S \\ref{section:discussion}. ",
        "conclusions": "\\protect\\label{section:conclusion}  We have applied a planetary atmospheric model with two physically different conditions in an attempt to detect the planetary signature of \\hdb. We are able to reject the scenario in which Ti and V are depleted from the atmosphere of \\hdb\\ in agreement with the predictions of recent models. We have also assessed our ability to test the models in which TiO and VO are present in the upper atmosphere, resulting in a temperature inversion and lines appearing in emission. We are however unable to detect a planet using this model, or to reject the model since our data do not achieve the necessary S/N ratios required. The 99 per cent confidence limit is 2 times higher than the expected flux ratio at the most probable velocity amplitude of the planet and would thus necessitate a further six nights of observation before this level of sensitivity could be reached.  In order to reliably test the most recent models which include a stratosphere and weak emission, cross dispersed spectrographs which offer greater wavelength coverage will afford higher S/N gain and thus improve our chances of detection. Equally, study of transiting systems reduce the degeneracy considerably in any analysis since the inclination is known and will be close, or equal to 90\\degs. The greater maximum brightness of a transiting CEGP system will increase our chances of detecting and characterising the planet."
    },
    "0806/0806.1061_arXiv.txt": {
        "abstract": "I compute the effect on the power spectrum of tracers of the large-scale mass-density field (e.g., galaxies) of primordial non-Gaussianity of the form $\\Phi=\\phi+\\fnl\\left(\\phi^2-\\left<\\phi^2\\right>\\right)+\\gnl\\phi^3+...$, where $\\Phi$ is proportional to the initial potential fluctuations and $\\phi$ is a Gaussian field, using beyond-linear-order perturbation theory. I find that the need to eliminate large higher-order corrections necessitates the addition of a new term to the bias model, proportional to $\\phi$, i.e., $\\delta_g=b_\\delta \\delta+b_\\phi\\fnl\\phi+...$, with all the consequences this implies for clustering statistics, e.g., $P_{gg}\\left(k\\right)=b_\\delta^2 P_{\\delta\\delta} \\left(k\\right)+2 b_\\delta b_\\phi\\fnl P_{\\phi\\delta}\\left(k\\right)+b_\\phi^2 \\fnl^2 P_{\\phi\\phi}\\left(k\\right)+...~$.  This result is consistent with calculations based on a model for dark matter halo clustering, showing that the form is quite general, not requiring assumptions about peaks, or the formation or existence of halos. The halo model plays the same role it does in the usual bias picture, giving a prediction for $b_\\phi$ for galaxies known to sit in a certain type of halo. Previous projections for future constraints based on this effect have been very conservative -- there is enough volume at $z\\lesssim 2$ to measure $\\fnl$ to $\\sim \\pm 1$, with much more volume at higher $z$. As a prelude to the bias calculation, I point out that the beyond-linear (in $\\phi$) corrections to the power spectrum of mass-density perturbations are naively infinite, so it is dangerous to assume they are negligible; however, the infinite part can be removed by a renormalization of the fluctuation amplitude, with the residual $k$-dependent corrections negligible for models allowed by current constraints. ",
        "introduction": "Inflation \\cite{1979JETPL..30..682S,1981PhRvD..23..347G,1982PhLB..108..389L, 1982PhRvL..48.1220A} has been tested successfully mainly through measurements of the power spectrum of primordial density perturbations \\cite{2006JCAP...10..014S,2008arXiv0803.0547K}. These perturbations are expected to be nearly Gaussian and nearly scale invariant \\cite{1981JETPL..33..532M,1982PhLB..115..295H,1982PhRvL..49.1110G, 1982PhLB..117..175S,1983PhRvD..28..679B}.  Testing the Gaussianity of the perturbations with increasing accuracy will be a major goal of future work aimed at distinguishing different models (see \\cite{2004PhR...402..103B} for a review of non-Gaussianity from inflation). The simplest models of inflation predict that non-Gaussianity will be undetectably small \\cite{1991PhRvD..43.1005S,1993PhRvL..71.1124L, 1994ApJ...430..447G,1995ApJ...454..552M, 2003NuPhB.667..119A,2003JHEP...05..013M,2004PhR...402..103B}, but multifield models \\cite{1997PhRvD..56..535L,2002PhRvD..66j3506B, 2003PhRvD..67b3503L,2007PhRvD..76d3526B,2006PhRvD..73h3522R, 2008PhRvD..78b3513I}, models where non-Gaussianity is generated during reheating \\cite{2004PhRvD..69b3505D,2008PhRvD..77b3505S} or preheating \\cite{2005PhRvL..94p1301E,2006PhRvD..73j6012B,2007PhRvD..75h6004B, 2008PhRvL.100d1302C}, bouncing/ekpyrotic/cyclic models \\cite{2007JCAP...11..010C,2008PhRvD..77f3533L, 2008PhRvL.100q1302B,2008PhRvD..78b3506L}, or inflation models based on nonlocal field theory \\cite{2007JCAP...07..017B,2008JCAP...06..030B} can predict levels of non-Gaussianity near the present detection limits ($\\fnl \\sim 100$, as defined below). While there have been some hints of non-Gaussianity in the cosmic microwave background (CMB) \\cite{2007arXiv0710.2371J,2008PhRvL.100r1301Y}, the general consensus seems to be that nothing convincingly primordial has been detected \\cite{2003ApJS..148..119K,2007JCAP...03..005C, 2008arXiv0803.0547K}.  Recently, \\cite{2008PhRvD..77l3514D,2008ApJ...677L..77M,2008arXiv0806.1046A} showed that there should be a distinctive signature of non-Gaussianity in the large-scale power spectrum of dark matter halos, observable as galaxies, in the local model for non-Gaussianity with curvature perturbations proportional to \\begin{equation} \\Phi = \\phi +  \\fnl \\left(\\phi^2 -\\left<\\phi^2\\right>\\right) +\\gnl \\phi^3+... \\end{equation} \\cite{2001PhRvD..63f3002K,2004PhR...402..103B,2006PhRvD..74j3003S, 2006PhRvD..74l3519B}, where I include the 3rd order term which is necessary to compute the power spectrum to 4th order in the small, Gaussian, perturbations $\\phi$. This form is only a special case of non-Gaussianity, but it is simple, and calculations using it should point the way to generalizations. The idea of \\cite{2008PhRvD..77l3514D}, to use the power spectrum, was a departure from the previously standard approach of studying non-Gaussianity in large-scale structure (LSS) using the bispectrum \\cite{2000MNRAS.313..141V,2001PhRvL..86.1434F,2002MNRAS.335..432V, 2004PhRvD..69j3513S,2005MNRAS.361..824G, 2005MNRAS.364..620G,2005JCAP...10..010B,2006PhRvD..74b3522S, 2007PhRvD..76h3004S} or other higher-order statistics \\cite{2005PASJ...57..709H, 2006ApJ...653...11H,2007MNRAS.376..343K, 2007ApJ...658..669S,2008JCAP...04..014L,2008MNRAS.385.1613H}.  It was somewhat surprising that this signal would be detectable, because the nonlinear corrections to the power spectrum in this model were generally assumed to be negligibly small (because fluctuations in $\\phi$ are of order $10^{-5}$). Reference \\cite{2008JCAP...08..031S} verified and further explored this idea and applied it to real galaxy and quasar data sets.  They found observational constraints $-1~(-23)<\\fnl<+70~(+86)$, at 95\\% (99.7\\%) confidence, by combining LSS and CMB data. The projected rms error from the Planck satellite, measuring the CMB anisotropy, is $\\sigma_{\\fnl} \\sim 5$ \\cite{2001PhRvD..63f3002K, 2008PhRvD..77l3006C}.  The calculations of \\cite{2008PhRvD..77l3514D,2008ApJ...677L..77M,2008JCAP...08..031S, 2008arXiv0806.1046A} were all based on models for the clustering of dark matter halos. The purpose of this paper is to investigate the result using a different approach -- the renormalized perturbative bias model of \\cite{2006PhRvD..74j3512M}. Instead of focusing on dark matter halos, this approach starts by assuming that the tracer density field is a completely general, unknown, function of the local mass-density field and Taylor expanding this function in the mass-density fluctuations, leading to the form \\cite{1993ApJ...413..447F} \\begin{equation} \\delta_g\\left(\\delta\\right)=c_1 \\delta + \\frac{1}{2}c_2 \\left(\\delta^2-\\sigma^2\\right)+ \\frac{1}{6}c_3 \\delta^3+\\epsilon+\\orderfour~, \\end{equation} where $\\delta$ is the mass-density perturbation, $\\sigma^2=\\left<\\delta^2\\right>$, and $\\delta_g$ is the tracer density perturbation (I will usually refer to the tracer as galaxies, although there are many other possibilities like quasars \\cite{2008JCAP...08..031S}, the \\lyaf\\ \\citep{2005ApJ...635..761M,2006ApJS..163...80M,2006MNRAS.365..231V}, galaxy cluster/Sunyaev-Zel'dovich effect measurements \\citep{2008RPPh...71f6902A}, and possibly future 21cm surveys \\citep{2005MNRAS.364..743N,2008PhRvL.100i1303C}). The unknown coefficients of the Taylor series have become the bias parameters $c_i$ and a random noise variable $\\epsilon$ has been added to allow for stochasticity in the relation (e.g., from shot noise, at least). This type of bias model has been explored in many papers \\cite{1995ApJS..101....1M, 1998MNRAS.301..797H, 1998ApJ...504..607S,1999ApJ...510..541T,1999ApJ...522...46T, 1999ApJ...520...24D,2000ApJ...537...37T,2001ApJ...556..641H}. The new development in \\cite{2006PhRvD..74j3512M} was to show that large higher-order corrections in this perturbative approach can be eliminated by redefining the bias parameters to absorb them, motivated by the way that masses and coupling constants are redefined to absorb divergent loop corrections in quantum field theory \\cite{Peskin:1995ev}. I argued that this renormalization approach is an improvement over the previous approach of defining the galaxy density to be a function of a smoothed mass-density field (so the higher-order terms in the Taylor series can be kept explicitly small \\cite{1993ApJ...413..447F,1998MNRAS.301..797H}), because the smoothing needed is so extreme that it directly affects the scales of interest \\cite{2007PhRvD..75f3512S,2008PhRvD..78b3523S}, and because, in the unrenormalized approach, the higher-order corrections modify the power on arbitrarily large, ideally truly linear, scales \\cite{1998MNRAS.301..797H}. Reference \\cite{2008arXiv0805.2632J} found extremely good agreement between the renormalized model and the clustering of galaxies in numerical simulations. Generally, the value of perturbation theory (PT) for describing LSS will only increase as observational measurements become more precise, because this increases the range of scales where corrections to linear theory are important but are still small enough to be treated perturbatively.  The value of PT has been further enhanced recently by the introduction of several renormalization methods applicable to the mass-power spectrum calculation \\cite{2006PhRvD..73f3519C,2006PhRvD..73f3520C, 2007PhRvD..75d3514M,2007JCAP...06..026M, 2007PhRvD..76h3517I,2008ApJ...674..617T,2008PhRvD..77b3533C, 2008MPLA...23...25M,2008PhRvD..77f3530M, 2008PhRvD..78h3503B}.  To complete the relation between primordial non-Gaussianity and final density, note that the initial density field is related to $\\Phi$ through the transfer function from primordial to late-time linear fluctuations, $T(k)$, and the Poisson equation \\begin{equation} \\delta_1\\left(\\vk,a\\right) = \\frac{2}{3}\\frac{c^2 D\\left(a\\right)}{\\Omega_{m,0}H_0^2 } k^2 T\\left(k\\right)\\Phi_\\vk \\equiv M\\left(k,a\\right) \\Phi_\\vk ~, \\label{eqPoisson} \\end{equation} using the definitions of \\cite{2007PhRvD..76h3004S}, where I will use the subscript on $\\delta$ to indicate the order in the initial density perturbations, not the Gaussian field $\\phi$ (I will use $\\delta_L$ to indicate the fully linear, including in $\\phi$, density perturbation).  The growth factor is normalized so $D\\left(a\\right)=a$ in the matter dominated era. The transfer function is time independent and normalized by $T\\left(k\\rightarrow 0\\right)\\rightarrow 1$.  These definitions make $\\Phi_\\vk$ time independent and close to scale invariant. The perturbative density field is written as a Taylor series in $\\delta_1$, i.e., $\\delta=\\delta_1+\\delta_2+\\delta_3+...$, with $\\delta_i \\sim \\mathcal{O}\\left(\\delta_1^i\\right)$, e.g., \\begin{equation} \\delta_2\\left(\\vk\\right)= \\int \\frac{d^3\\vq}{\\left(2 \\pi\\right)^3}\\delta_1\\left(\\vq\\right) \\delta_1\\left(\\vk-\\vq\\right) J_S^{\\left(2\\right)}\\left(\\vq,\\vk-\\vq\\right)~. \\end{equation} where \\begin{equation} J_S^{(2)}(\\vk_1,\\vk_2)= \\frac{5}{7}+\\frac{1}{2}\\frac{\\vk_1\\cdot\\vk_2}{k_1 k_2} \\left(\\frac{k_1}{k_2}+\\frac{k_2}{k_1}\\right)+\\frac{2}{7} \\left(\\frac{\\vk_1\\cdot\\vk_2}{k_1 k_2}\\right)^2 \\end{equation} is given by standard LSS perturbation theory (see \\cite{2002PhR...367....1B} for a review).  The rest of the paper goes as follows: To put the perturbative bias calculation on a firm foundation, in Sec. \\ref{secmassdensity}, I compute the mass power spectrum in the non-Gaussian model. In Sec. \\ref{secbias}, I calculate the power spectrum of a biased tracer. Finally, in Sec. \\ref{secconclusions}, I discuss the results. ",
        "conclusions": "}  In Sec. \\ref{secmassdensity}, we found that the corrections to the mass power spectrum arising from the non-Gaussian initial conditions are naively infinite, but the divergence can be removed by a renormalization of the amplitude of perturbations, after which there is no significant effect of non-Gaussianity for realistic models. It is interesting to note that a sufficiently large numerical simulation would have to deal with this issue explicitly, because the most natural method of implementing initial conditions effectively implements the unrenormalized version \\cite{2008PhRvD..77l3514D}; however, the dynamical range of typical simulations is small enough that this effect is unlikely to be noticeable (the box size and finite resolution provide large and small-scale cutoffs). The issues discussed in Sec. \\ref{secmassdensity} should be revisited when considering other forms of non-Gaussianity.  In Sec. \\ref{secbias} I computed the power spectrum of biased tracers of mass density and found a large higher-order perturbative correction that can only be removed by renormalization if we add a term to the standard linear bias model, proportional to $\\phi$, producing the model \\begin{equation} \\delta_g= c_\\delta~\\delta +c_\\phi \\fnl \\phi+ \\frac{1}{2}~ c_{\\delta^2}~\\left( \\delta^2-\\sigma^2\\right) + \\frac{1}{3!}~ c_{\\delta^3}~ \\delta^3 +\\epsilon+\\orderfour~. \\end{equation} This new term will propagate into the computation of all statistics, not just the power spectrum.  Note that $\\phi$ should not be confused with the Newtonian potential at the time of observation -- it is defined in terms of $\\delta$ by Eq. (\\ref{eqPoisson}).  One might ask at this point: ``once there is a linear term, why did we not add higher-order terms in $\\phi$?'' The short answer is simply that we did not need to.  We added the minimal term necessary to produce a well-behaved calculation.  A calculation to even higher order would no doubt require the addition of higher order counter terms, and guide their form.  It seems likely that the higher-order terms would be completely negligible, like the corrections to the mass power spectrum, but, as we have learned, one should really compute them instead of assuming.  Measuring $\\fnl$ accurately depends on predicting $b_\\phi$ for an observable in just the same way as measuring the amplitude of the matter power spectrum depends on a prediction of the usual bias $b_\\delta$. By comparison to the results of \\cite{2008PhRvD..77l3514D, 2008ApJ...677L..77M,2008JCAP...08..031S}, we can determine the value of $b_\\phi$ predicted by the halo model.  As emphasized by \\cite{2008JCAP...08..031S}, however, it should be kept in mind that this kind of prediction makes some assumptions that are not guaranteed to be true, e.g., the galaxy population in a halo might depend on something other than halo mass alone, like merging history, which affects the clustering of the population. Halo model predictions will give a good estimate for the size of effect we expect for a given $\\fnl$ and type of galaxy, but if we are fortunate enough to make a detection it will be very difficult to measure $\\fnl$ to high precision using only this power spectrum effect.  With that caveat, Eq. (18) of \\cite{2008JCAP...08..031S} leads to \\begin{equation} b_\\phi = 2 \\left(b_\\delta-1\\right) \\delta_c=3.372 \\left(b_\\delta-1\\right)~. \\label{halobphi} \\end{equation}  Reference \\cite{2008PhRvD..77l3514D} computed that a very modest future survey, extending only to $z=0.7$ (while the bulk of the volume of the universe is at higher redshift) can constrain $\\fnl$ to $\\sim \\pm 10$. It is interesting to know at least roughly how well a larger survey can do. For a well-sampled survey (i.e., no shot noise), the signal to noise of a single mode is \\begin{equation} \\frac{S}{N}_{\\rm mode}=\\frac{\\left[1+ \\left(\\frac{k_\\star}{k}\\right)^2\\right]^2-1} {\\left[1+\\left(\\frac{k_\\star}{k}\\right)^2\\right]^2} ~, \\end{equation} (note that a null test would not include the signal in the denominator, but this makes very little difference, as I will show) where \\begin{equation} k_\\star=\\left(\\frac{3~b_\\phi \\fnl \\Omega_{m,0}} {2~ b_\\delta D\\left(a\\right)}\\right)^{1/2} \\frac{h~ {\\rm Mpc}^{-1}}{3000} \\end{equation} I am assuming $T\\left(k\\right)\\simeq 1$ on the relevant scales, as discussed below. I will ignore redshift space, geometric, and evolutionary distortions, and any confounding systematic errors or parameter degeneracies. The total signal to noise of a survey is approximately \\begin{equation} \\left(\\frac{S}{N}\\right)^2\\simeq \\frac{V}{4 \\pi^2} \\int_{k_{\\rm min}}^{k_{\\rm max}} dk k^2 \\left[\\frac{S}{N} \\left(k\\right)\\right]^2 \\end{equation} where $k_{\\rm min}\\sim 2 \\pi / V^{1/3}$ (the detection limit is not very sensitive to $k_{\\rm min}$, as long as $k_{\\rm min}$ is basically set by the survey volume). $k_{\\rm max}$ should be something less than $\\infty$, but, for interesting $k_\\star$ the integral converges within the linear regime (the integral would not converge if the $k$ dependence of the transfer function was included, however, the information gained this way would come from a tiny change in signal on scales where the signal would not be uniquely distinguishable from changes in other parameters, and linearity cannot be safely assumed, so a projection including it would be unreliable). The bulk of the signal comes from the range $2 k_\\star \\lesssim k \\lesssim 10 k_\\star$.  Consequently, the bulk of the signal comes from the $\\fnl \\phi~\\delta$ cross term, rather than the $\\fnl^2 \\phi~\\phi$ term (cross correlating with an unbiased tracer is worse than autocorrelation, but essentially only because of the factor of 2 that is lost in the cross term). To a very good approximation, $S/N=1$ when $k_\\star \\sim \\pi V^{-1/3} \\sim k_{\\rm min}/2$ (e.g., this would be $\\sim 0.8~\\pi V^{-1/3}$ if I used the null test $S/N$ per mode and set $k_{\\rm min}$ based on the diameter of a sphere instead of the side length of a cube, which makes it slightly smaller), which gives $\\fnl b_\\phi/ b_\\delta D(a)\\simeq 200 \\left({\\rm Gpc}/h\\right)^2/V^{2/3}$ (for $\\Omega_m=0.3$).  Assuming Eq. (\\ref{halobphi}) and $b_\\delta D(a)=2$, we have finally that $S/N=1$ when $\\fnl \\sim 125 \\left[\\left({\\rm Gpc}/h\\right)^2/V^{2/3}\\right]  / \\left[2/D(a)-1\\right]$. I plot the value of this function for a survey covering all volume up to $z$ in Fig. \\ref{figfnlcon}. \\begin{figure} \\resizebox{\\textwidth}{!}{\\includegraphics{fnlcon.eps}} \\caption{ Rms error on $\\fnl$ for a well-sampled all-sky survey out to $z$, ignoring redshift-space, geometric, and evolutionary distortions.  Rms detection significance is essentially linear in $\\fnl$ (out to a $\\sim 5 \\sigma$ level of detection), and scales roughly as ${\\rm volume}^{2/3}$ (as long as the volume remains compact). } \\label{figfnlcon} \\end{figure} Note that the assumption $b_\\delta D(a)=2$ appears to make $b_\\delta$ unreasonably large at high redshift, however, the value of $b_\\delta$ cancels in $b_\\phi/b_\\delta$ if $b_\\delta$ is very large, so the result is not actually very sensitive to this assumption. We see that there is plenty of volume in the Universe to measure $\\fnl$ even if it is less than 1, although this will of course take a huge amount of work. Reference \\cite{2008arXiv0806.1046A} found similar results, using a variety of slightly different assumptions, which testifies to the robustness of the projection (which ultimately just amounts to the statement that the survey should be large enough to resolve the $k$ where the non-Gaussian-induced power becomes roughly equal to the Gaussian power). Twenty-one cm intensity mapping \\cite{2008MNRAS.383.1195W,2008PhRvL.100i1303C} may be a route to surveying all of the post-reionization volume.  Even if one's tracer is weakly biased, so the halo model predicts little signal, it may be possible to make a nonlinear transformation of the density field to enhance the bias; however, it is not clear that any such transformation would be better than simply measuring the bispectrum.  It is interesting to note that the power spectrum measurement proposed by  \\cite{2008PhRvD..77l3514D} and discussed in this paper can be seen as nature's implementation of the poor-person's bispectrum measurement suggested by \\cite{2006PhRvL..97z1301C}, i.e., the cross-correlation of the field with the square of the field, where here the squaring is done for us, as part of the formation of biased structure. The significance of the difference between the natural and artificial versions should not be underestimated -- exploiting the natural version we can observe the effect on large scales even if limited resolution or noise prevents us from probing the small scales that would be required to perform a similar transformation artificially. \\cite{2006PhRvL..97z1301C,2008PhRvD..77j3506C} found that $\\fnl$ may be measurable to $\\sim 0.01$ using the bispectrum of high redshift 21cm observations and probing down to very small scales.  In this paper I have excluded the possibility of measuring the signal using small scales (by assuming the transfer function is 1 in the signal-to-noise ratio calculation), but it may be useful, especially at high redshift, to explore the possibility of using smaller scales and combining the bispectrum and power spectrum.  I thank Niayesh Afshordi and Latham Boyle for helpful discussions."
    },
    "0806/0806.1311_arXiv.txt": {
        "abstract": "Recent Chandra and Newton observations indicate that there are two-temperature components ($T \\sim$ 8 keV, 0.8 keV) of the diffuse x-rays emitted from deep inside the center of Milky Way. We show that this can be explained by the existence of sterile neutrinos, which decay to emit photons that can be bound-free absorbed by the isothermal hot gas particles in the center of Milky Way. This model can account for the two-temperature components naturally as well as the energy needed to maintain the $\\sim$ 8 keV temperature in the hot gas. The predicted sterile neutrino mass is between 16-18 keV. ",
        "introduction": "Recently, a large amount of diffuse x-ray data have been obtained by Chandra, BeppoSAX, Suzaku and XMM-Newton, giving a complex picture near the Milky Way center \\citep{Muno,Rockefeller,Senda,Hamaguchi,Sidoli,Sakano}. The data indicates that there exists a high temperature ($\\sim$ keV) hot gas near the Milky Way center. \\citet{Kaneda} suggested that the existence of two-temperature components of the hot gas can explain the observed x-ray spectrum. \\citet{Sakano} analysed the XMM-Newton data to get 1 keV and 4 keV hot gas components in Sgr A East, a supernova remnant located close to the Milky Way center. Later \\citet{Muno} used the data from Chandra to model the temperature of the two components within 20 pc as 0.8 keV and 8 keV. The temperature of the soft component ($T_1 \\sim$ 0.8 keV) can be explained by 1 percent of kinetic energy by one supernova explosion in every 3000 years near the Milky Way center, corresponding to power $\\sim 3 \\times 10^{36}$ erg s$^{-1}$ \\citep{Muno}. However, the high temperature of the hard component ($T_2 \\sim$ 8 keV) cannot be explained satisfactorily. The power needed to sustain the high temperature is about $10^{40}$ erg s$^{-1}$. Moreover, the emission of the hard component is distributed much more uniformly than the soft component and the intensity of the two components are correlated, which suggest that they are produced by related physical processes \\citep{Muno}. In this article, we present a model to explain the two-temperature composition, as well as the high temperature and uniform emission of the hard component obtained by Chandra within 20 pc. We assume that there are sterile neutrinos with rest mass $m_s \\ge 16$ keV in the deep galactic center and their decay photons continuously supply the energy of the hard x-ray emission. The sterile neutrino halo is located at the center of a 20-pc radius gas cloud which is heated by the decayed photons. Most photons are absorbed near the center ($\\sim 0-1$ pc) of Milky Way and the energy is subsequently transferred to the surrounding gas. The major heating mechanism is the bound-free collisions between the decay photons and the ions in the hot gas. In this optically thick region, the hot gas particles are in photoionization equilibrium. The high metallicity of the gas, $\\tilde{Z}_{\\rm{metal}} \\geq 2-3$ solar metallicity, enhances the heating rate of the gas and provides enough energy to sustain the high temperature of the hard component. The energy absorbed in the central region is transferred to the outside optically thin region ($1-20$ pc) by collisions among the electrons to share their energy (collisional equilibrium). Also, we assume that the soft and hard components are in equilibrium with different uniform temperatures and they are bounded hydrostatically.  Although the recent MiniBooNE data challenges the LSND result that suggests the existence of eV scale sterile neutrinos \\citep{Aguilar}, more massive sterile neutrinos (eg. keV) may still exist. The fact that active neutrinos have rest mass implies that right-handed neutrinos should exist which may indeed be massive sterile neutrinos. The existence of the sterile neutrinos has been invoked to explain many phenomena such as reionization \\citep{Hansen}, missing mass \\citep{Dodelson} and the high temperature of the hot gas in clusters \\citep{Chan}. Therefore, it is worthwhile to discuss the consequences of the existence of massive sterile neutrinos, which may decay into light neutrinos and photons.  The existence of the small size keV sterile neutrino halo is first suggested by \\citet{Viollier}. The size of a self-gravitating degenerate sterile neutrino halo depends on $m_s$ and total mass $M_s$: $R_s=0.0006(M_s/10^6M_{\\odot})^{-1/3}(m_s/16~{\\rm keV})^{-8/3}$ pc. Including the contribution of the baryons, the size will be even smaller. The size of the sterile neutrino halo is upper bounded by $R_s \\le 0.0005$ pc \\citep{Schodel}. This size is very small compared to that of a galaxy and therefore the sterile neutrino halo will hide deeply inside the galactic center \\citep{Munyaneza}. Sterile neutrinos may decay into active neutrinos and become a strong energy source to galaxies and clusters. The decays of keV order sterile neutrinos may also help to solve the cooling flow problem in clusters \\citep{Chan}. ",
        "conclusions": "To explain the origin of the hard component, \\citet{Muno} make use of magnetic reconnection driven by the turbulence that supernovae generate in the interstellar medium. Magnetic reconnection can heat the hot gas to $kT \\sim B_{\\rm center}^2/8 \\pi n_g$. For $n_g \\sim 0.1$ cm$^{-3}$, $B_{\\rm center} \\sim 0.2$ mG, $kT \\sim 8$ keV. However there is not enough evidence to support whether this mechanism can maintain the high temperature of the hard component \\citep{Muno}.  In our model, we have assumed that there exists a sterile neutrino halo with $m_s=16-18$ keV in the Milky Way center, which decay to emit $\\gamma$ with life-time of cosmological order. It provides a large amount of energy to the hot gas and maintains the extremely high temperature. The bound-free collisions provide enough energy to the two different temperature components and maintain their temperatures. At the same time, a stable two-temperature structure in the hot gas can be explained by this heating mechanism naturally. The uniform emission of the soft and hard components suggests that they may come from similar physical processes \\citep{Muno}. In our model, both components indeed share the same source of energy - the 8-9 keV photons emitted by the decays of sterile neutrinos. In our model, the sterile neutrinos may not be the major component of dark matter. Therefore, any bounds on $m_s$ assuming they are the major dark matter candidate does not constrain our model severely. The heating rate in the Milky Way center is time dependent as there is a decreasing number of sterile neutrinos. Therefore, if two galaxies have similar chemical compositions, the heating rate is greater for the large redshift one, which has more particles in the higher temperature component. We therefore predict that the hard component of the x-rays would be stronger for large redshift and metal-rich galaxies. Moreover, if a galaxy has lower metallicity, then only a single temperature component may be observed instead of two (see Fig.~(2))."
    },
    "0806/0806.2725_arXiv.txt": {
        "abstract": "A new self-similar solution describing the dynamical condensation of a radiative gas is investigated under a plane-parallel geometry. The dynamical condensation is caused by thermal instability. The solution is applicable to generic flow with a net cooling rate per unit volume and time $\\propto \\rho^2 T^\\alpha$, where $\\rho$, $T$ and $\\alpha$ are density, temperature and a free parameter, respectively. Given $\\alpha$, a family of self-similar solutions with one parameter $\\eta$ is found in which the central density and pressure evolve as follows: $\\rho(x=0,t)\\propto (t_\\mathrm{c}-t)^{-\\eta/(2-\\alpha)}$ and $P(x=0,t)\\propto (t_\\mathrm{c}-t)^{(1-\\eta)/(1-\\alpha)}$, where $t_\\mathrm{c}$ is an epoch when the central density becomes infinite. For $\\eta\\sim 0$, the solution describes the isochoric mode, whereas for $\\eta\\sim1$, the solution describes the isobaric mode. The self-similar solutions exist in the range between the two limits; that is, for $0<\\eta<1$. No self-similar solution is found for $\\alpha>1$. We compare the obtained self-similar solutions with the results of one-dimensional hydrodynamical simulations. In a converging flow, the results of the numerical simulations agree well with the self-similar solutions in the high-density limit. Our self-similar solutions are applicable to the formation of interstellar clouds (HI cloud and molecular cloud) by thermal instability. ",
        "introduction": "Thermal instability (TI) is an important physical process in astrophysical environments that are subject to radiative cooling. \\citet{FR85} investigated the effect of TI caused by Bremsstrahlung on structure formation at galactic and subgalactic scales. TI also plays an important role in the interstellar medium (ISM). The ISM consists of a diffuse high-temperature phase (warm neutral medium, or WNM) and a clumpy low-temperature phase (cold neutral medium, or CNM), in supersonic turbulence. It is known that the ISM is thermally unstable in the temperature range between these the two stable phases; that is, in the range $T\\sim 300-6000$ K \\citep{DM72}. Koyama \\& Inutsuka (2000, 2002) suggested that the structure of the ISM is provided by the phase transition that is caused by TI that occurs in a shock-compressed region.   The basic properties of TI were investigated by \\citet{F65}, who investigated the TI of a spatially uniform gas using linear analysis for the unperturbed state in thermal equilibrium. The net cooling function per unit volume and time was assumed to be $\\Lambda_0\\rho^2T^\\alpha - \\Gamma_0 \\rho$. \\citet{F65} showed that gas is isobarically unstable for $\\alpha<1$ and is isochorically unstable for $\\alpha<0$. In thermal non-equilibrium, however, where the heating and the cooling rates are not equal in the unperturbed state, the stability condition of TI is different from that in thermal equilibrium. This situation is, for example, generated in a shock-compressed region. \\citet{B86} derived a general criterion for the TI of gas in a non-equilibrium state. A flow dominated by cooling is isobarically unstable for $\\alpha\\la2$ and is also isochorically unstable for $\\alpha\\la 1$. Therefore, the phase with $T\\sim 300-6000$K is both isobarically and isochorically unstable.  The non-linear evolution of TI has been investigated by many authors \\citep{KI02,AH05,HA07,H06,V07} using multi-dimensional simulations. However, the nonlinear behaviour has only been investigated analytically by a few authors, using some simplifications in an attempt to gain a deeper insight into the nature of TI. \\citet{M89} investigated analytic solutions under the isobaric approximation. However, this approximation is valid only in the limit of short or intermediate-scale.  In this paper, a new semi-analytic model for the non-linear hydrodynamical evolution of TI for a thermally non-equilibrium gas is considered without the isobaric approximation. Once TI occurs, the density increases drastically until a thermally stable phase it reached. During this dynamical condensation, the gas is expected to lose its memory of the initial condition, and is expected to behave in a self-similar manner \\citep{ZR67,LL59}. In particularly, in star formation, self-similar (S-S) solutions for a runaway collapse of a self-gravitating gas have been investigate by many authors \\citep{P69,L69,S77,H77,WS85,BL95}. In this paper, S-S solutions describing dynamical condensation by TI are investigated by assuming a simple cooling rate without self-gravity. A family of S-S solutions which have two asymptotic limits (the isobaric and the isochoric modes) is presented.  In Section \\ref{formulation}, we derive the S-S equations, and describe the mathematical characteristic of our S-S equations as well as the numerical methods. In Section \\ref{result}, our S-S solutions are presented and their properties are discussed. In Section \\ref{discuss}, the S-S solutions are compared with the results of one-dimensional numerical simulations. The astrophysical implications of the S-S solutions and the effect of dissipation are also discussed. Our study is summarized in Section \\ref{summary}. ",
        "conclusions": "\\label{discuss} \\subsection{Comparison with the results of time-dependent numerical hydrodynamics} We compare the S-S solutions with results of time-dependent numerical hydrodynamics using one-dimensional Lagrangian 2nd-order Godunov method \\citep{vL97}. A converging flow of an initially uniform and thermal-equilibrium gas ($\\rho(x)=P(x)=1$), is considered. An initial velocity field is given by $v(x)= -2\\tanh(x/L)$, where $L$ is a parameter. The net cooling rate is assumed to be $\\rho{\\cal L}(\\rho,T)=(\\rho^{3/2}P^{1/2} - \\rho)/(\\gamma-1)$, where $\\alpha$=0.5. The inflow creates two shock waves, which propagate outward. Runaway condensation of the perturbation owing to TI occurs in the post-shock region because the cooling rate dominates the heating rate. The central density continues to grow and ultimately becomes infinite. During the runaway condensation, the flow in the central region is expected to lose the memory of the initial and the boundary conditions, and to converge to one of the S-S solutions. The scale of perturbation in the post-shock region can be controlled by the value of $L$. The larger $L$ is, the larger the scale of perturbation becomes. We perform numerical simulations for the case with $L$=0.005, 0.01, 0.02, 0.2 and 0.5.  Fig. \\ref{ve2 00} shows the time evolution of $\\rho_\\mathrm{cen}(t)$ and $P_\\mathrm{cen}(t)$ as a function of $1-t/t_\\mathrm{con}$, where $\\rho_\\mathrm{cen}$ and $P_\\mathrm{cen}$ are the central density and pressure, and $t_\\mathrm{con}$ is an epoch when $\\rho_\\mathrm{cen}$ becomes infinity estimated from the numerical results. The panels represent (a)$L$=0.01, (b)0.02, (c)0.2 and (d)0.5, respectively. Using equation (\\ref{den P 00}), the time evolution of $\\rho_\\mathrm{cen}(t)$ and $P_\\mathrm{cen}(t)$ of each simulation provides the corresponding $\\eta$. As a result, the corresponding values of $\\eta$ are found to be (a)0.99, (b)0.98, (c)0.76 and (d)0.50. The thin solid lines in Fig. \\ref{ve2 00} indicate the corresponding S-S solutions. From the value of $\\eta$ and Fig. \\ref{ve2 00}, cases (a) and (b) correspond to the isobaric mode, case (c) corresponds to the intermediate mode ($\\eta\\sim\\eta_\\mathrm{eq}$) and case (d) corresponds to the isochoric mode. There is a difference in the convergence speed from the corresponding S-S solution. Fig. \\ref{ve2 00} shows that the convergence speed of (a) and (b) is faster than that of (c) and (d). This is because the scale-length of the condensed region of (a) and (b) is smaller and the sound-crossing time-scale is shorter.  Fig. \\ref{ve2} shows the time evolution of the rescaled density $\\rho(1-t/t_\\mathrm{con})^{\\eta/(2-\\alpha)}$ and pressure $P(1-t/t_\\mathrm{con})^{-(1-\\eta)/(1-\\alpha)}$ as a function of the rescaled coordinate $x(1-t/t_\\mathrm{con})^{-n}$. The parameters $L$ and $\\eta$ in Fig. \\ref{ve2} are the same as those in Fig. \\ref{ve2 00}. The thin solid lines in Fig. \\ref{ve2} indicate the corresponding S-S solutions. In all the cases (a-d), the central regions are well approximated by the corresponding S-S solutions. From Fig. \\ref{ve2}, it can be seen that the region that is well approximated by the S-S solution spreads with time in each panel.   From the above discussion, the non-linear evolution of perturbation is well approximated by one of the S-S solution in high density limit. Which S-S solution ($0<\\eta<1$) is more likely to be realised in actual environments where various scale perturbations exist? To answer this question, the dependence of $t_\\mathrm{con}$ on $L$ is investigated. Fig. \\ref{ve2 collapse} shows the dependence of $1/t_\\mathrm{con}$ on $1/L$. Here, $1/L$ and $1/t_\\mathrm{con}$ correspond to the wavenumber and the growth rate of the perturbation, respectively. Therefore, Fig. \\ref{ve2 collapse} can be interpreted as an approximate dispersion relation for the non-linear evolution of TI. For larger $1/L$, $1/t_\\mathrm{con}$ is larger and $1/t_\\mathrm{con}$ appears to approach a certain asymptotic value. A similar behaviour is also found in the dispersion relation of \\citet{F65} without heat conduction. In Fig. \\ref{ve2 collapse}, the condensation for $1/L=5.0$ corresponds to the intermediate mode. Furthermore, although the isobaric condensation grows faster, there is little difference between $1/t_\\mathrm{con}$ for $1/L$=200 (the isobaric mode) and 2 (the isochoric mode). Therefore, both the isobaric and the isochoric condensations are expected to coexist in actual astrophysical environments. \\subsection{Astrophysical implication} In astrophysical environment, there are roughly two ranges of temperature with $\\alpha<1$ where our S-S solutions are applicable.  In the high-temperature case, $T\\ga 2\\times 10^5$K \\citep{SD93}, the dominant coolants are metal lines and Bremsstrahlung. In a shock with a velocity greater than \\begin{equation} V\\ga 120\\;\\mathrm{km\\;s^{-1}}\\left(\\frac{T}{2\\times10^5\\;\\mathrm{K}}\\right)^{1/2}, \\end{equation} the post-shock region is thermally unstable ($\\alpha<1$). For example, a typical supernova satisfies this condition, and the S-S condensation is expected to occur. In reality, the density does not become infinite because the cooling time-scale increases when the gas reach the thermally stable phase. In this case, the cooled layer will rebound and overstability may take place \\citep{CI82,YN01,S03}.  In the low-temperature case, 300$\\la T\\la$6000K, the dominant coolant is neutral carbon atom. \\citet{KI00} suggested that the structure of the ISM is caused by TI during the phase transition (WNM $\\rightarrow$ CNM). The phase transition is induced by a shock wave. \\citet{HP99} considered the converging flow of a WNM under a plane-parallel geometry using one-dimensional hydrodynamical calculations, and investigated the condensation condition of the Mach number and pressure in the pre-shock region. In situations which satisfy their condition, S-S condensation is expected to occur. \\begin{figure} \\centering \\includegraphics[width=80mm]{figure7.eps} \\caption{The dependence of the growth rate, $1/t_\\mathrm{con}$, on the wavenumber, $1/L$. } \\label{ve2 collapse} \\end{figure} \\subsection{Effects of dissipation} In the actual gas, the effect of viscosity and heat conduction becomes important on small scales. The importance of dissipation can be evaluated by the ratio between the advection and the dissipation terms, or Raynolds number, which is given by \\begin{equation} {\\cal R}= \\frac{C_\\mathrm{V} \\rho v\\partial_x T}{\\partial_x(K(T)\\partial_x T)} \\sim \\frac{C_\\mathrm{V}\\rho_{00} U x_\\mathrm{cen}}{K(T_{00})}, \\label{raynolds} \\end{equation} where $C_\\mathrm{V}$, $K(T)$ and $U$ are the specific heat at constant volume, the heat conduction coefficient and the typical velocity, respectively. The Plandtl number is implicitly assumed to be of order unity. A typical length-scale is assumed to be $x_\\mathrm{cen}(t)$. If the flow converges to one of the S-S solutions, the ratio, $\\tau$, between the sound-crossing and the cooling timescales is constant. Therefore $\\tau$ is defined by \\begin{equation} \\tau = \\frac{t_\\mathrm{sound}^\\mathrm{cen}}{t_\\mathrm{cool}^\\mathrm{cen}} = x_\\mathrm{cen}\\gamma (\\gamma-1)\\Lambda_0 \\rho_{00} c_{00}^{2\\alpha-3}. \\end{equation} Using $\\tau$, $x_\\mathrm{cen}$ is written as \\begin{equation} x_\\mathrm{cen}= \\frac{\\tau}{\\gamma (\\gamma-1)\\Lambda_0} \\rho_{00}^{-1} c_{00}^{3-2\\alpha}. \\end{equation} Substituting this equation into equation (\\ref{raynolds}), one obtains \\begin{equation} {\\cal R}\\sim\\frac{C_\\mathrm{V}f_v \\tau}{\\gamma (\\gamma-1)\\Lambda_0K_0}c_{00}^{4-2\\alpha-\\kappa}, \\label{ray} \\end{equation} where $U=f_v c_{00}$$(f_v<1)$ and $K(T_{00})=K_0c_{00}^\\kappa$.  For the low-temperature case, we adopt $K = 2.5\\times10^3 \\sqrt{T}\\;\\;$ ergs cm$^{-1}$ K$^{-1}$ s$^{-1}$ \\citep{P53} and the following cooling function \\citep{KI02}: \\begin{equation} \\Lambda(T) \\sim 1.0\\times10^{20} \\rho^2\\sqrt{T}\\;\\;\\;\\mathrm{erg\\;cm^{-3}\\;s^{-1}}. \\end{equation} From equation (\\ref{ray}), the Raynolds number is given by \\begin{equation} {\\cal R}\\sim 84\\left(\\frac{f_v}{0.1}\\right)\\left(\\frac{\\tau}{0.1}\\right) \\left(\\frac{T}{10^3\\mathrm{K}}\\right). \\end{equation} In the isobaric mode, $\\tau \\sim 0.1$ from Fig. \\ref{00}. Because the Raynolds number is much larger than unity, the dynamical condensation of the post-shock region is expected to be well described by the S-S solution with $\\tau\\ga 0.1$.  For the high-temperature case, the cooling rate of metal lines ($10^5$K $<T<$ $10^7$K) is given by \\begin{equation} \\Lambda(T) \\sim 2.2\\times10^{29}\\rho^2 T^{-0.6} \\;\\;\\mathrm{ergs\\;cm^{-3}\\;s^{-1}}. \\end{equation} We adopt $K=1.24\\times 10^{-6} T^{5/2}$ ergs cm$^{-1}$ K$^{-1}$ s$^{-1}$ \\citep{P53}. Using this formula and eqution (\\ref{ray}), the Raynolds number is given by \\begin{equation} {\\cal R}\\sim85\\left(\\frac{f_v}{0.1}\\right)\\left(\\frac{\\tau}{0.1}\\right) \\left(\\frac{T}{10^6\\mathrm{K}}\\right)^{0.1}. \\end{equation} Because of ${\\cal R}\\gg1$, in the high-temperature case, the S-S solution with $\\tau \\ga 0.1$ is expected to well describe the dynamical condensation.  As the temperature decreases, ${\\cal R}$ decreases and becomes less than unity at a certain epoch. Thereafter, our S-S solutions are invalid and the effect of dissipation becomes important."
    },
    "0806/0806.0720_arXiv.txt": {
        "abstract": "{The external destruction of protoplanetary discs in a clustered environment acts mainly due to two mechanisms: gravitational drag by stellar encounters and evaporation by strong stellar winds and radiation. It is a fundamental question whether either of these mechanisms is important for the stellar evolution and planet formation process in young star clusters.} {We focus on the effect of stellar encounters in young dense clusters and investigate whether there are any observables that could trace this mechanism and its impact on disc evolution. If encounters play a role in disc destruction, one would expect that stars devoid of disc material would show unexpectedly high velocities as an outcome of close interactions. We want to quantify this effect by numerical simulations and compare it to observations.} {As a model cluster we chose the Orion Nebula Cluster (ONC). We reanalyzed observational data of the ONC to find encounter signatures in the velocity distribution and a possible correlation with signatures of circumstellar discs. We use the {\\textsc{\\mbox{nbody6\\raise.2ex\\hbox{\\tiny{++}}}}} code to model the dynamics of an ONC-like cluster and analyze the velocity distribution and the disc-mass loss due to encounters.} {We found from the observational data that 8 to 18 stars leave the ONC with velocities several times the velocity dispersion. The majority of these high-velocity stars are young low-mass stars ($t$\\,$\\lesssim$\\,$10^5$\\,yr, $m$\\,$\\approx$\\,0.2-0.3\\,{\\mbox{M$_{\\odot}$}}), among them several lacking infrared excess emission. Interestingly, the high-velocity stars are found only in two separate regions of the ONC - i) close to the cluster centre and ii) in the outer cluster region. Our simulations give an explanation for the location of the high-velocity stars and provide evidence for a strong correlation between location and disc destruction.} {The high-velocity stars can be explained as the outcome of close three-body encounters; the partial lack of disc signatures is attributed to gravitational interaction. The spatial distribution of the high-velocity stars reflects the initial structure and dynamics of the ONC. Our approach can be generalized to study the evolution of other young dense star clusters, like the Arches cluster, back in time. } ",
        "introduction": "According to current knowledge, planetary systems form from the accretion discs around young stars. These young stars are in most cases not isolated but are part of a cluster \\citep[e.g.][]{2003ARA&A..41...57L}. Densities in these cluster environments vary considerably, spanning a range of 10\\,pc$^{-3}$ (e.g. $\\eta$ Chameleontis) to $10^6$\\,pc$^{-3}$ (e.g. Arches Cluster) . It is still an open question as to how far interactions with the surrounding stars influence planet formation in dense young clusters ($n$\\,$\\gtrsim$\\,$10^4$\\,pc$^{-3}$). These discs disperse over time \\citep{2001ApJ...553L.153H,2002astro.ph.10520H,2006ApJ...638..897S,2008ApJ...672..558C} and in dense clusters the disc frequency seems to be lower in the core \\citep[e.g.][]{2007ApJ...660.1532B}. This is attributed to external violent processes like photo-evaporation or encounter-induced disc mass loss, mainly caused by the massive stars that are concentrated in the cluster core.  Earlier numerical investigations seemed to indicate that photoevaporation should by far dominate the external disc destruction \\citep{2001MNRAS.325..449S,2004ApJ...611..360A}. The question of whether encounters play a vital role in the formation process of stars and planets has been studied far less and is still open \\citep[e.g.][]{2006ApJ...641..504A}. Only recently has it been shown that stellar encounters do have an effect on the discs surrounding stars in a young dense cluster \\citetext{\\citealp[][]{2006ApJ...642.1140O,2006A&A...454..811P,2006ApJ...652L.129P,2006ApJ...653..437M, 2007ApJ...661L.183M,2007ApJ...656..275M,2007A&A...462..193P,2007A&A...475..875P}; \\citealp[see also the review by][]{2007ARA&A..45..481Z}}. The importance of encounters has been underestimated because in previous studies the focus has been on encounters between solar mass stars \\citep[e.g.][]{1991MNRAS.249..584C,1995ApJ...455..252H}. However, discs are most affected when the masses of the stars involved in the encounter are unequal \\citep{2006ApJ...642.1140O,2007ApJ...656..275M}. Moreover, the massive stars in the center of a stellar cluster act as gravitational foci for the lower mass stars and are thus subject to repeated encounters. These accumulated perturbations can lead to a total destruction of a massive star's disc \\citep{2006A&A...454..811P}.  The main focus of this paper is the question of whether there is direct \\emph{observational} evidence for encounters among star-disc systems in young clusters and for disc-mass loss due to encounters. It is difficult to distinguish observationally whether photoevaporation or gravitational interaction are responsible for the loss of (outer) discs. The reason is that in both cases interaction with a massive star is the most destructive process. Thus the observation of a decreased disc frequency in the cluster core, as mentioned above, does not allow one to favor either of the two mechanisms. The difficulty of tracing stellar encounters directly by observations is their short duration. Direct imaging of tidal tails of a disc would be the best proof but the probability of such an event is very low: the perturbed disc circularizes and tidal tails dissipate on a time scale of $\\lesssim$\\,1000\\,yr. Nevertheless, there exist observations of tidal tails or spiral arms in star-disc systems \\citep[e.g.][]{2005prpl.conf.8092B,2006A&A...452..897C,2006ApJ...645.1297L}, but often it is unclear whether these can be attributed to the passage of an unbound perturber or a binary companion or gravitational instabilities caused by a giant planet. Due to the short dissipation time scale of encounter-induced tidal tails, there would be a high probability for the encounter partner still to be located close ($\\lesssim$\\,$10^4$\\,AU) to the perturbed disc. However, source confusion in dense young clusters like the ONC, where stellar encounters are frequent enough to study the effect on stellar discs, would prevent a clear identification. Set in relation to the rate of encounters in which prominent tidal tails may be formed, we estimate an upper limit of just four candidate stars for a direct observation of tidal tails in a protoplanetary disc due to encounters in the ONC (see Appendix \\ref{app:enc_rate} for a detailed calculation). However, an unambiguous imprint of an encounter among stars is the high velocity of a star which has been expelled in a close gravitational interaction, mostly as a result of a three-body encounter \\citep[see][]{1975MNRAS.173..729H}. Thus an analysis of the velocity distribution of a cluster is the key to finding candidates of close encounters between young stars.  Here the ONC is used as a prototype young cluster because it is one of the best-studied regions in our galaxy, and the only young dense cluster for which velocities of more than 1000 of its members have been determined \\citep{1988AJ.....95.1755J}. In addition, its high density suggests that stellar encounters might be relevant for the evolution of circumstellar discs. Throughout this work we assume that initially all stars are surrounded by protoplanetary discs. This is justified by observations that reveal disc fractions of nearly 100\\,\\% in very young star clusters \\citep[e.g.][]{2000AJ....120.1396H,2000AJ....120.3162L,2001ApJ...553L.153H,2005astro.ph.11083H}.  In \\S\\ref{sec:obs} we present results from a search for candidate stars of close encounters in the publicly available observational data of the ONC. For this purpose we have reanalyzed the investigations of \\citet{1988AJ.....95.1755J}, \\citet{1997AJ....113.1733H}, and \\citet{1998AJ....116.1816H} for stellar velocities and infrared excess. The basic properties of the ONC used for our numerical model are described in \\S\\ref{sec:onc_dyn}. There we also discuss theoretical estimates and observational data related to binary populations and encounters. Afterwards we present results from a numerical approach to this problem (\\S\\ref{sec:sim}). Observational and numerical results are summerized and discussed in \\S\\ref{sec:discussion}. ",
        "conclusions": "\\label{sec:discussion}  Combining observational data and numerical simulations, we have shown that even after 1\\,Myr of dynamical evolution a cluster of the size and density of the ONC is still dynamically active. Young stars, preferentially of low-mass, are expelled in close $N$-body encounters, losing a large fraction of their circumstellar matter.  The encounters are in most cases interactions with massive stars in the cluster centre. This finding highlights the significant effect that encounters in (massive) young stellar clusters can have on the evolution of protoplanetary discs. This is even more evident if one addresses not only the effect on the mass, but the even stronger effect on the angular momentum of the disc as shown in previous investigations \\citep{2006ApJ...652L.129P,2007A&A...462..193P}.  Using the observational data from \\citet{1988AJ.....95.1755J} we show (Fig.~\\ref{fig:obs_gauss}c) that in the ONC there is a small population of stars with proper motions $\\mu_{\\rm{x,y}}$ larger than three times the one-dimensional velocity dispersion of the cluster, $\\mu_{\\rm{x,y}}$\\,$\\gtrsim$\\,7.5\\,{\\mbox{km\\,s$^{-1}$}}.  From numerical simulations we find that the so-called ``high-velocity stars'' form two dynamically disjoint groups. One group is composed of stars expelled in a close encounter, moving on radial tracks directly outward from the cluster centre, termed ``radial escapers''. The other contains unperturbed stars running on wide, non-closed orbits around the cluster centre (on average with lower velocities), termed ``orbital escapers''. The dynamical state of the stars has its origin in their initial phase space location: radial escapers are initially located close to the cluster centre ($r$\\,$\\lesssim$\\,0.3\\,pc) and characterized by a low angular momentum (relative to the cluster centre). In contrast, orbital escapers were formed in the outer cluster with large angular momentum relative to the cluster centre.  The different dynamics of the high-velocity stars implies a signature in their circumstellar characteristics: stars being subject to close encounters are expected to lose their disc material faster and to a higher degree \\citep{2006A&A...454..811P}. Using our prescription of disc-mass loss in stellar encounters, indeed we find a clear correlation from numerical simulations: radial escapers, initially located close to the cluster centre and later expelled in a close encounter, lose more than 90\\,\\% of their disc material in 1\\,Myr of dynamical evolution, while orbital escapers, initially distant stars, moving on wide orbits around the cluster centre, do not experience strong perturbations and keep most of their disc material.  We compare our numerical results with infrared observations from \\citet{1997AJ....113.1733H}, \\citet{2002ApJ...573..366M}, and \\citet{2004AJ....128.1254L}, tracing the (inner) circumstellar material. Though strict conclusions are not possible due to the small observational sample and some stars that cannot be classified according to our scheme, we observe the same trend: stars that have been classified as orbital escapers do show near-infrared excess emission, indicative of circumstellar matter, while those classified as radial escapers show pure photospheric colors, lacking evidence of (inner) circumstellar discs.  This view is strongly supported by the work of \\citet{2004ApJ...607L..47T} on the dynamics of {\\mbox{$\\theta^1{\\rm{C}}\\:{\\rm{Ori}}$}}. The most massive star in the ONC, {\\mbox{$\\theta^1{\\rm{C}}\\:{\\rm{Ori}}$}}, has a proper motion several times greater than the dispersion of bright ONC stars and much larger than the velocity expected if it were in equipartition with the other cluster stars \\citep{1988AJ.....95.1744V}. \\citet{2004ApJ...607L..47T} showed that the direction of {\\mbox{$\\theta^1{\\rm{C}}\\:{\\rm{Ori}}$}}'s motion is consistent with being exactly opposite to that of a B-type star embedded in the background molecular cloud, the so-called BN object, and concludes that it is most probably a runaway star originating from the Trapezium ejected about 4000 years ago after a close encounter with {\\mbox{$\\theta^1{\\rm{C}}\\:{\\rm{Ori}}$}}. This supports the idea that stellar encounters may be likely events in such dense regions as the ONC, or at least in their cores. The evidence that the most massive stellar object was involved in a close encounter with another massive star is in best accordance with the work of \\citet{2007ApJ...661L.183M} and \\citet{2007A&A...475..875P}.   Considering the dynamical age of the ONC of several crossing times, encounters must have had much stronger impacts on a stellar disc at an early age of the cluster, when densities were much higher but massive stars already had formed. Indeed, our simulations confirm this expectation, giving rise to an era of strong and frequent interactions among star-disc systems at the onset of massive star formation.  We expect that in even denser clusters such as the Arches cluster, high-velocity stars should be even more frequent than in the ONC. In such systems the current spatial and velocity distribution of the high-velocity stars should give strong indications of how the cluster developed in former times."
    },
    "0806/0806.2663_arXiv.txt": {
        "abstract": "New observations from the Hubble ultra deep field suggest that the star formation rate at $z>7$ drops off faster than previously thought. Using a newly determined star formation rate for the normal mode of Population II/I stars  (PopII/I), including this new constraint,  we compute the  Thomson scattering optical depth and find a result that is marginally consistent with WMAP5 results. We also reconsider the role of Population III stars (PopIII) in light of cosmological and stellar evolution constraints. While this input may be needed for reionization, we show that it is essential in order to account for cosmic chemical evolution in the early Universe. We  investigate the consequences of PopIII stars on the local metallicity distribution function of the Galactic halo (from the recent Hamburg/ESO survey of metal-poor stars) and on the evolution of abundances with metallicity (based on the ESO large program on very metal-poor stars), with special emphasis on carbon-enhanced metal-poor stars. The metallicity distribution function shape is well reproduced at low iron abundance ([Fe/H]$\\gtrsim-4$), in agreement with other studies. However, the Hamburg/ESO survey hints at a sharp decrease of the number of low-mass stars at very low iron abundance, which is not reproduced in models with only PopII/I stars. The presence of PopIII stars, of typical masses 30-40 \\msun, helps us to reproduce this feature, leading to a prompt initial enrichment before the onset of PopII/I stars. The metallicity at which this cut-off occurs is  sensitive to the lowest mass of the massive PopIII stars, which makes the metallicity distri\\-bution function a promising tool to constrain this population. Our most important results show that the nucleosynthetic yields of PopIII stars lead to abundance patterns in agreement with those observed in extremely metal-poor stars. This can be demonstrated by the transition discri\\-mi\\-nant (a criterion for low-mass star formation taking into account the cooling due to \\ion{C}{ii} and \\ion{O}{i}). In this chemical approach to cosmic evolution, PopIII stars prove to be a compulsory ingredient, and extremely metal-poor stars are inevitably born at high redshift. ",
        "introduction": "\\label{s:introduction}  Recent data from the Hubble Ultra Deep Field (HUDF) have allowed the determinations of the cosmic star formation rate (SFR) to be extended from redshifts $z\\sim4$ up to $z\\sim7-8$ \\citep{2007ApJ...670..928B}.  In the context of the  $\\Lambda$CDM  model, which has by and large  been confirmed by the observation of anisotropies in the diffuse cosmological microwave background (CMB) by WMAP \\citep{2007ApJS..170..377S,2009ApJS..180..306D}, a remaining challenge lies in understanding the first structures, including galaxies and stars. This step requires either direct observations (of black holes, gamma-ray bursts\\ldots) or indirect constraints through luminosity functions, metal pollution and relic metal-poor stars (e.g. \\citealt{2006ApJ...641....1T,2007MNRAS.381..647S}).  Our work develops the latter approach, and focuses on the nucleosynthesis pollution from the first stars and their consequences for relic metal-poor stars.   The role of massive PopIII stars as ionization sources at high redshift is poorly understood. Indeed, their role has been clouded by the fact that the measured value of the Thomson scattering optical depth (from $z=0$ to the redshift of emission of the CMB) has decreased significantly from WMAP1 to WMAP5. Current measurements imply a redshift for an instantaneous reionization of $z=11.0\\pm1.4$ ($\\tau=0.087\\pm0.017$) \\citep{2009ApJS..180..306D}. In this context, the new observations of the cosmic SFR at high redshift bring a fresh perspective on the role played by PopIII stars.  Since PopIII stars have a very specific impact on nucleosynthesis, it  is useful to incorporate constraints from stellar observations. From the nucleosynthetic point of view, halo stars have long been used to constrain galactic chemical evolution models, but were quite disconnected from cosmological models. Yet  PopIII stars might be the primary source for the early metal enrichment of the interstellar medium (ISM) as well as for the intergalactic medium (IGM). The different abundance patterns observed in extreme metal-poor stars (EMPS) may well be explained in terms of these stars \\citep{2007ApJ...663..687Y,2008ASPC..393...63F}. It has been shown \\citep{2003Natur.425..812B} that  the abundances of ionized carbon and neutral atomic oxygen are important for the transition from PopIII to PopII/I. \\cite{2007MNRAS.380L..40F} have  defined a  transition discriminant, \\Dtrans,  and we show below that this quantity clearly reveals the nucleosynthetic imprint of PopIII stars.   The paper is organized as follows. In \\Sec{models}, we summarize our model for the global chemical and cosmic evolution. New constraints related to the SFR and reionization are described in \\Sec{SFR}. We first reconsider the reionization constraint from WMAP5 data in terms of a single normal mode of PopII/I star formation, by fitting the SFR to the most recent data \\citep{2007ApJ...670..928B}. We compare this result to a model which includes a contribution from PopIII stars. Stellar observations are described in \\Sec{stellar}. To address these issues, we use some of the abundances (Fe, O, C, Si) of individual stars of the European Southern Observatory Large Program (ESO-LP) 'First Stars'  \\citep{2004A&A...416.1117C} and of peculiar extremely-low metallicity stars. We also compute the metallicity distribution function (MDF, number of observed low-mass stars at a specific metallicity) and compare this to recent observations  \\citep{2008arXiv0809.1172S}. A grid study of the model parameters allows us to draw confidence contours for the SFR of PopIII stars and on the transition discriminant in \\Sec{grid}. We also consider  abundances derived from 1D and 3D model atmospheres and compare results in each case. Our best model is used as an illustration in Sections~\\ref{s:SFR} and \\ref{s:stellar}. We summarize our results and conclude in \\Sec{conclusion}. ",
        "conclusions": "\\label{s:conclusion}  Following the approach developed in \\citet{2004ApJ...617..693D}, we have modeled  the evolution of individual element abundances in the ISM assuming homogeneous star formation and stellar yields. Recent observations at z$\\sim$ 7-8 \\citep{2007ApJ...670..928B} were used to better constrain one ingredient of the model, namely the SFR for PopII/I at high redshift. We have shown that a homogeneous scenario of hierarchical structure formation reproduces many different  observations, from reionisation and first star abundances to local abundance observations. We found that using the most recent results on the optical depth from WMAP, a massive mode is not absolutely required. Nevertheless, the data can accommodate a PopIII contribution, responsible for the gradual reionization starting from $z\\simeq 20$.  Although the cosmological importance  of PopIII stars cannot be fully constrained by the integrated Thomson optical depth, this question may be better tackled in the future with the help of accurate measurements of the CMB polarization data \\citep{2009ApJS..180..306D}.  We have also considered stellar constraints, in particular the MDF and the evolution of \\Dtrans\\ with [Fe/H]. In the literature, the slope of the MDF can be reproduced by different models of chemical enrichment, from galactic models to hierarchical analytic models including merger trees, as shown in Fig. 12 of \\citet{2008arXiv0809.1172S} (see also \\citealt{2006ApJ...641....1T}). This indicates that the slope of the  MDF does not discriminate between the methods used and the associated   level of heterogeneity. In contrast, the pattern at [Fe/H]$<-3.5$ is more difficult to reproduce. \\citet{2008arXiv0809.1172S} claim that no model considered in their paper can reproduce the MDF at very low iron abundances. We find that the modification of the tail at low values of [Fe/H] may be  related to the presence of  PopIII stars.  In our study, we have tested the impact  of the massive mode by varying the typical PopIII minimal mass from 8 to 40 \\msun. This appears to be an important parameter as found in the study of \\citet{2007MNRAS.381..647S}. They used a mass range from 140 to 200~\\msun\\ which we consider disfavoured by the very specific yields of the stars within this mass range \\citep{2006ApJ...647..773D}. In this paper, we demonstrate that the minimum mass must be larger than about 30 \\msun\\ in order to reproduce the observed part of the MDF.  We have shown that the existence of  PopIII stars at high redshift is required to explain the abundance pattern observed in the CEMP stars. In addition, we have shown that a massive mode with a typical mass of  40 \\msun\\ reproduces the evolution of observed \\Dtrans. In contrast, it is known that PISN with masses 140-200~\\msun\\ do not provide the correct chemical pattern and cannot reach high values of \\Dtrans\\ for low values of [Fe/H] (the ratio C/Fe in their yields, \\citealt{2003ApJ...591..288H}, is not high enough). However, if  it could be established that the CEMP stars were particular cases, such as belonging to binaries \\citep{2005ApJ...635..349R,2007ApJ...665.1361T} or due to the preferential depletion of iron in grains \\citep{2008ApJ...677..572V}, the SFR related to PopIII stars would be diminished, at least as far as chemical evolution is concerned.  In conclusion, our analysis  hints at a massive mode at $z\\simeq 20-30$, which becomes sub-dominant at lower $z$ ($z\\sim 15$).\\\\  Going further requires improvements from the observational and theoretical point of view. On the theoretical part, the existence of PopIII stars with masses of a few tens of solar masses is also suggested by recent hydrodynamical simulations \\citep{2003Natur.425..812B,2006MNRAS.373..128G,2006MNRAS.366..247J,2007MNRAS.374.1557J,2007ApJ...665...85J,2007ApJ...663..687Y,2008MNRAS.387.1021G,2009ApJ...691..441S}. The explosion of these high mass stars disrupts the surrounding environment and delays the formation of lower mass stars  {\\em within the same structure}.  Merger trees \\citep{2007MNRAS.381..647S} account for such a mechanical feedback which delays  PopII/I star formation along a single branch.  In a homogeneous picture, this delay translates into a reduction of the SFR between the two modes. We note that  this epoch is very short in our model (100-200 Myr). However, the exact evolution of the structures are still uncertain. One also needs  improved yields provided by stellar evolution models at all metallicities  and, more specifically, SN calculations that are even more uncertain. Note however that uncertainties related to the determination of abundances (derived from 1D or 3D model atmospheres) do not change our conclusion on the cosmological importance of PopIII stars.  Many improvements are possible in the future with regard to  observations. $(i)$ Additional stellar constraints and abundance measurements. Better statistics are necessary to construct the MDF at very low metallicity, with next generation optical telescopes such as GMT. It will then be possible to constrain in a better way the epoch of  the massive population (whose duration is related  the critical metallicity as used in merger trees;  \\citealt{2007MNRAS.381..647S,2006ApJ...641....1T})  and  the typical mass range of PopIII stars. $(ii)$  More complete spectroscopic observations in very metal-poor stars are needed in order to obtain a complete set of observations in the \\Dtrans\\ diagram. Of particular importance  is the continued search for HMP stars and UMP stars to test their statistical significance relative to the bulk of the VMP/EMP stars. $(iii)$ In future studies, one could attempt to better constrain the model using  other chemical elements, in particular r and s-process elements which can bring new constraints on the mass range of massive stars, PopIII \\citep[e.g.][]{2008arXiv0812.1227F}. $(iv)$ Additional cosmological constraints at high redshift. Massive objects may be directly observable, as more and more quasars and galaxies are detected at high redshift \\citep{2003AJ....125.1649F,2004ApJ...607..697K,2006Natur.443..186I} and with JWST in the future (e.g. \\citealt{2009AAS...21342603S}).\\ However, the extreme brevity of the PopIII epoch makes the connection unclear \\citep{2007AAS...211.9114A}. A more promising probe would be the observations of high redshift gamma-ray bursts, that correspond, for the longer bursts, to the deaths of massive single stars \\citep{2003ApJ...591..288H,2004ApJ...604L...1B,2006ApJ...642..382B,2006MNRAS.372.1034D}. Recent observations of GRBs at $z>6$ (GRB 050904 at $z=6.3$, \\citealt{2006Natur.440..184K}; GRB 080913 at $z=6.7$, \\citealt{2009ApJ...693.1610G}) and even $z>8$ (GRB 090423 at $z\\sim 8.2$, \\citealt{2009arXiv0906.1577T}) are  extremely encouraging. $(v)$  In addition, massive stars would have polluted their environment with an initial enrichment of heavy elements which could be compared to the one observed  in the Ly$\\alpha$ forest along quasar absorption spectra at $z\\lsim 6$ \\citep{2001ApJ...561L.153S,2002ApJ...576....1A,2004A&A...419..811A} or in the Damped Lyman $\\alpha$ systems \\citep{2003MNRAS.346..209L,2003ApJ...595L...9P}. This requires a better understanding of the outflows of metals into  the IGM which is beyond the scope of this work. In addition, it will be important to better understand the role of inhomogeneities  on the different populations of the first stars. The uniformity and extent of the metal pollution is also under debate, and could in the future be used to distinguish between local and recent pollution and global pollution by an earlier population of stars \\citep{2001ApJ...555...92M,2002ApJ...571...40M,2003ApJ...591...38W}."
    },
    "0806/0806.2380_arXiv.txt": {
        "abstract": "The gravitational collapse of a star is an important issue both for general relativity and astrophysics, which is related to the well known ``frozen star\" paradox. Following the seminal work of Oppenheimer and Schneider (1939), we present the exact solution for two dust shells collapsing towards a pre-existing black hole. We find that the inner region of the shell is influenced by the property of the shell, which is contrary to the result in Newtonian theory and and the clock inside the shell becomes slower as the shell collapses towards the pre-existing black hole. This result in principle may be tested experimentally if a beam of light travels across the shell. We conclude that the concept of the ``frozen star\" should be abandoned, since matter can indeed cross a black hole's horizon according to the clock of an external observer. Since matter will not accumulate around the event horizon of a black hole, we predict that only gravitational wave radiation can be produced in the final stage of the merging process of two coalescing black holes. Our results also indicate that for the clock of an external observer, matter, after crossing the event horizon, will never arrive at the ``singularity\" (i.e. the exact center of the black hole). ",
        "introduction": "The ``frozen star\" paradox is a well known novel phenomenon predicted by general relativity, i.e. a distant observer ($O$) sees a test particle falling towards a black hole moving slower and slower, becoming darker and darker, and is eventually frozen near the event horizon of the black hole. However, as we discussed in \\cite{20}, several possible explanations to this phenomenon have been proposed, but none of these is completely satisfactory. In the previous work, we found that an external observer should be able to observe matter falling in a black hole, based on our {\\it stationary} solution of a dust shell around a pre-existing black hole. In this paper we further obtain the exact {\\it dynamical} solution for two shells collapsing towards a pre-existing black hole, in particular to find the difference between this solution and that for the gravitational collapse of a uniform dust ball \\cite{14}, and the implication for the ``frozen star\" paradox. ",
        "conclusions": "In the Newtonian gravitation theory, if the matter is spherically symmetric, the outer matter will not influence the inner region. However, as shown in Sec. II, in general relativity case, the clock in the inner region of the shell is slower compared with the case without the shell. This effect may be testable experimentally in principle, e.g. if a beam of light travels across the shell and a parallel beam of light passes outside the shell (far enough from the shell), then the two beams of light will travel at different velocities with respect to the observers outside the shell.  The inner shell can  cross the Schwarzschild radius in the two shell case. In this sense we could observe the matter falls into a black hole and the ``frozen star\" paradox discussed in \\cite{10} is naturally solved. As pointed out previously by us, the origin of the ``frozen star\" is the ``test particle\" assumption for the infalling matter, which neglects the influence of the infalling matter itself on the metric  \\cite{20}. In fact, for some practical astrophysical settings, the time taken for matter falling into a black hole is quite short, even for the external observer \\cite{20}. We find that the infalling matter will not accumulate outside the event horizon, and thus the quantum radiation and Gamma Ray bursts predicted in \\cite{11} and \\cite{12} are not likely to be generated. We predict that only gravitational wave radiation can be produced in the final stage of the merging process of two coalescing black holes. Future simultaneous observations by X-ray telescopes and gravitational wave telescopes shall be able to verify our prediction.  It is also interesting to note that, as can be seen from Fig. \\ref{fig6}b, in ordinary coordinates, the matter will not collapse to the singularity  $(r = 0)$ even with infinite coordinate time (if $r_0' = 0$, the inner boundary of the inner shell will take infinite time to arrive at  $r = 0$). It means that in real astrophysics sense, matter can never arrive at the singularity (i.e., the exact center of the black hole) with respect to the clock of the external observer. Therefore, no gravitational singularity exists physically, even within the framework of the classical general relativity.     \\begin{theacknowledgments} We thank many discussions with Sumin Tang, Richard Lieu, Kinwah Wu, Kazuo Makishima, Neil Gehrels, Masaruare Shibata, Ramesh Narayan, Zheng Zhao, Zonghong Zhu and Chongming Xu. SNZ is grateful to Prof. Sandip Chakrabarti for his great effort in organizing this conference, and to the great hospitality of S.N. Bose National Centre for Basic Sciences, Kolkata, India. Many participants of this conference, especially Prof. Roy Kerr, are greatly appreciated for interesting discussions. SNZ acknowledges partial funding support by the Yangtze Endowment from the Ministry of Education at Tsinghua University, Directional Research Project of the Chinese Academy of Sciences under project No. KJCX2-YW-T03 and by the National Natural Science Foundation of China under project no. 10521001, 10733010 and 10725313. \\end{theacknowledgments}"
    },
    "0806/0806.3981_arXiv.txt": {
        "abstract": "We use semi-analytic models of structure formation to interpret gravitational lensing measurements of substructure in galaxy cluster cores ($R{\\le}250{\\hkpc}$) at $z=0.2$.  The dynamic range of the lensing-based substructure fraction measurements is well matched to the theoretical predictions, both spanning $\\fsub\\sim0.05-0.65$.  The structure formation model predicts that $\\fsub$ is correlated with cluster assembly history.  We use simple fitting formulae to parameterize the predicted correlations: $\\Delta_{90}=\\tau_{90}+\\alpha_{90}\\log(\\fsub)$ and $\\Delta_{50}=\\tau_{50}+\\alpha_{50}\\log(\\fsub)$, where $\\Delta_{90}$ and $\\Delta_{50}$ are the predicted lookback times from $z=0.2$ to when each theoretical cluster had acquired 90\\% and 50\\% respectively of the mass it had at $z=0.2$.  The best-fit parameter values are: $\\alpha_{90}=(-1.34\\pm0.79)\\,{\\rm Gyr}$, $\\tau_{90}=(0.31\\pm0.56)\\,{\\rm Gyr}$ and $\\alpha_{50}=(-2.77\\pm1.66)\\,{\\rm Gyr}$, $\\tau_{50}=(0.99\\pm1.18)\\,{\\rm Gyr}$.  Therefore (i) observed clusters with $\\fsub\\ls0.1$ (e.g.\\ A\\,383, A\\,1835) are interpreted, on average, to have formed at $z\\gs0.8$ and to have suffered $\\le10\\%$ mass growth since $z\\simeq0.4$, (ii) observed clusters with $\\fsub\\gs0.4$ (e.g.\\ A\\,68, A\\,773) are interpreted as, on average, forming since $z\\simeq0.4$ and suffering $>10\\%$ mass growth in the $\\sim500\\,{\\rm Myr}$ preceding $z=0.2$, i.e.\\ since $z=0.25$.  In summary, observational measurements of $\\fsub$ can be combined with structure formation models to estimate the age and assembly history of observed clusters.  The ability to ``age-date'' approximately clusters in this way has numerous applications to the large clusters samples that are becoming available. ",
        "introduction": "\\label{sec:intro}  The mass growth of clusters is sensitive to the dark energy equation of state parameter $w$, the matter density of the universe $\\Omega_M$ and the normalization of the matter power spectrum $\\sigma_8$ (e.g.\\ \\citealt{Evrard93,Smith03,Mantz07}).  Clusters are inferred to grow hierarchically via the ingestion of smaller dark matter halos (that host galaxies) into the more massive parent halo (cluster).  The structure of galaxy clusters, specifically the internal substructure of clusters, therefore contains a wealth of cosmological information, including possible clues about the physics of the dark matter particle itself (e.g.\\ \\citealt{Natarajan02b}).  From an astrophysical point of view, the mass growth of clusters brings new (generally gas rich) galaxy populations into clusters (e.g.\\ \\citeauthor{Moran07} \\citeyear{Moran07}) and may lead to shock-heating of the intracluster medium and/or disruption of cooling in cluster cores (e.g.\\ \\citealt{Poole08}).  Reliable measurement and intepretation of cluster substructure is therefore of broad interest.  The most direct way to detect substructure within clusters is via gravitational lensing.  Group-scale substructures within individual clusters were detected in early ground-based strong-lensing studies of individual clusters \\citep{Pello91,Kneib93,Kneib95} and subsequently measured to high precision using \\emph{Hubble Space Telescope (HST)} data \\citep{Kneib96}.  \\citeauthor{Smith05a} (\\citeyear{Smith05a} -- hereafter Sm05 -- see \\S\\ref{sec:obs}) then measured the structure of a sample of 10 clusters at $z{\\simeq}0.2$.  In this article we use \\citeauthor{Taylor04}'s (\\citeyear{Taylor04} -- hereafter TB04) semi-analytical models of structure formation to interpret Sm05's cluster substructure measurements, as a means of exploring lensing-based substructure measurements as a quantitative probe of cluster age and assembly history.  We summarize Sm05 and TB04 in \\S\\ref{sec:obs} and \\S\\ref{sec:theory} respectively, and then synthesize observations and theory in \\S\\ref{sec:results}.  We discuss caveats in \\S\\ref{sec:discuss} and summarize our conclusions and discuss future prospects in \\S\\ref{sec:conc}.  We assume $H_0{=}70{\\rm km\\,s^{-1}Mpc^{-1}}$, ${\\Omega}_{\\rm M}{=}0.3$, ${\\Omega}_{\\rm \\Lambda}{=}0.7$ and $\\sigma_8=0.9$ throughout.  The lookback time from $z=0$ to $z=0.2$ is $t_{z=0.2}=2.44\\,{\\rm Gyr}$ in this cosmology. ",
        "conclusions": "\\label{sec:conc}  We have combined theoretical models of structure formation (TB04) with gravitational lens models of galaxy clusters (Sm05) to explore how measurements of cluster substructure from lensing observations can be interpreted in the context of the age and assembly history of clusters.  The main result is that $\\fsub$, the fraction of cluster mass within a projected cluster-centric radius of $R=250\\hkpc$ associated with substructure (galaxies and group-scale halos), as can be measured from lensing data, is predicted to be strongly correlated with the age and recent mass growth of galaxy clusters.  We fitted the the following simple formulae to the theoretical data to quantify the predicted behavior in a convenient form: $\\Delta_{90}=\\tninety-t_{z=0.2}=\\tau_{90}+\\alpha_{90}\\log(\\fsub)$ and $\\Delta_{50}=\\tninety-t_{z=0.2}=\\tau_{50}+\\alpha_{50}\\log(\\fsub)$, where $\\tninety$ and $\\tfifty$ are the lookback times at which a cluster had acquired 90\\% and 50\\% of its mass at $z=0.2$.  The best fit parameter values are: $\\alpha_{90}=(-1.34\\pm0.79)\\,{\\rm Gyr}$, $\\tau_{90}=(0.31\\pm0.56)\\,{\\rm Gyr}$ and $\\alpha_{50}=(-2.77\\pm1.66)\\,{\\rm Gyr}$, $\\tau_{50}=(0.99\\pm1.18)\\,{\\rm Gyr}$.  Low-$\\fsub$ clusters ($\\fsub\\ls0.1$; e.g.\\ A\\,383, A\\,1835) are therefore interpreted as, on average, having formed at $z{\\gs}0.8$ and having suffered ${\\le}10\\%$ mass growth in the 2\\,Gyr preceding $z=0.2$, i.e.\\ since $z\\simeq0.4$.  In contrast, high-$\\fsub$ clusters ($\\fsub\\gs0.4$; e.g.\\ A\\,68, A\\,773) are interpreted, on average, to have formed just $\\sim2\\,{\\rm Gyr}$ before $z=0.2$, i.e.\\ since $z\\simeq0.4$, and suffered $10\\%$ mass growth in the $\\sim0.5\\,{\\rm Gyr}$ preceding $z=0.2$, i.e.\\ since $z{\\simeq}0.25$.  Our synthesis therefore demonstrate that lensing-based measurements of $\\fsub$ can be combined with semi-analytic structure formation models to estimate the average age and assembly history of observed clusters. This suggests numerous avenues for further exploration, including: (i) expansion of the observed samples by at least an order of magnitude, (ii) calibration of the completeness of the lensing-based mass function of sub-halos in clusters, (iii) investigation of how lensing-based cluster age and assembly history estimates might allow new cosmological constraints to be derived, for example, on the dark energy equation of state parameter $w$, (iv) analysis of cluster galaxy populations and cluster scaling relations as a function of cluster age."
    },
    "0806/0806.3386_arXiv.txt": {
        "abstract": "{ Scientific exploitation of large variability databases can only be fully optimized if these archives contain, besides the actual observations, annotations about the variability class of the objects they contain. Supervised classification of observations produces these tags, and makes it possible to generate refined candidate lists and catalogues suitable for further investigation.} {We aim to extend and test the classifiers presented in a previous work against an independent dataset. We complement the assessment of the validity of the classifiers by applying them to the set of OGLE light curves treated as variable objects of unknown class. The results are compared to published classification results based on the so-called extractor methods.} {Two complementary analyses are carried out in parallel. In both cases, the original time series of OGLE observations of the Galactic bulge and Magellanic Clouds are processed in order to identify and characterize the frequency components. In the first approach, the classifiers are applied to the data and the results analyzed in terms of systematic errors and differences between the definition samples in the training set and in the extractor rules. In the second approach, the original classifiers are extended with colour information and, again, applied to OGLE light curves.} {We have constructed a classification system that can process huge amounts of time series in negligible time and provide reliable samples of the main variability classes. We have evaluated its strengths and weaknesses and provide potential users of the classifier with a detailed description of its characteristics to aid in the interpretation of classification results. Finally, we apply the classifiers to obtain object samples of classes not previously studied in the OGLE database and analyse the results. We pay specific attention to the B-stars in the samples, as their pulsations are strongly dependent on metallicity.} {} ",
        "introduction": "In the last decade, astronomy witnessed several major advances. The advent of large detection arrays, the operation of robotic telescopes and the consolidation of high duty cycle space missions have provided astronomers with a wealth of observations with unprecedented sensitivity in virtually the whole electromagnetic spectrum during long uninterrupted periods of time. At the same time, the ever-growing storage capacity of digital devices has made it possible to archive and make these enormous datasets available. The consolidation of the Virtual Observatory (VO) technology and the interoperability provided by its services make it possible for the astronomer to work consistently on large portions of the electromagnetic spectrum, combining different data models (magnitudes, colours, spectra, radial velocities, etc).  The traditional procedures for data reduction and analysis do not scale with the sizes of the available data warehouses. Some of its components have been automated and can now be carried out in a systematic way, but it is becoming evident that optimal scientific exploitation of these databases requires the addition of information inferred from the observed data to enable the extraction of homogeneous (in some sense) samples of observations for further specific studies that could not be applied to the entire database. The process by which this added value is extracted is widely known as Knowledge Discovery and relies mostly on recent advances in the artificial intelligence fields of pattern recognition, statistical learning or multi-agent systems.  The use of these new techniques has the particular advantage that, once accepted that every search for a given type of object is biased {\\sl ab initio} by the adopted definition of that class, automatic classifiers produce consistent object lists according to the same objective and stable criteria openly declared in the so-called training set. We thus eliminate subjective and unquantifiable considerations inherent to, for example, visual inspection and produce object samples comparable across different surveys.  Altogether, the integration of Computer Science techniques (Grid computing, Artificial Intelligence and VO technology) and domain knowledge (physics in this case), and the new possibilities that this synergy offers are known as e-Science. Science proceeds in much the same way as before; the e- prefix only provides the basis to approach more ambitious scientific challenges, feasible on the grounds of more and better quality data.  In \\citet[ hereafter paper I]{PaperI} we introduced the problem of the scientific analysis of variable objects and proposed several methods to classify new objects on the basis of their photometric time series. The OGLE database (see section \\ref{OGLE} for a summary of its objectives and characteristics) exemplifies some of the difficulties described in previous paragraphs. Although not its principal target, the OGLE survey has produced as a by-product hundreds of thousands of light curves of objects in the Galactic bulge and in the Large and Small Magellanic Clouds. These light curves have been analysed using the so-called extractor methods. Extractor methods can be assimilated to the classical rule-based systems where the target objects are identified by defining characteristic attribute ranges (where attribute is to be interpreted as any of the parameters used to describe the object light curves such as the significant frequencies, harmonic amplitudes or phase differences) where these objects must lie. In a subsequent stage, individual light curves are visually inspected and the object samples refined on a per object basis.   In this work we also present an extension of the classifiers defined in Paper I, to handle photometric colours. In section \\ref{OGLE} we summarize the objectives and characteristics of the OGLE survey; section \\ref{colours} describes the sources and criteria used for the assignment of colours to the training set and section \\ref{results} compares the results of the application of the classifiers (both with and without colours) to the OGLE database (bulge and Magellanic Clouds) with object lists available in the literature (obtained by means of extractor methods and human intervention) for a reduced set of classes. Finally, we analyse the object lists obtained with our classifiers for special classes in the realm of multiperiodic variables, not previously studied in an extensive way (to the best of our knowledge) in the context of the OGLE database. ",
        "conclusions": "In the past few years, the world of astronomy has seen a revolution taking place with the advent of massive sky surveys and large scale detectors. This revolution cannot be fully exploited unless automatic methods are devised in order to preprocess the otherwise unmanageably large databases. Otherwise, the efforts of the astronomical community will have to focus on repetitive uninteresting data processing rather than in the solution of the scientific questions that motivate the efforts. In this work we have presented a scenario with many interesting open questions for research (distance estimator calibration, stellar interiors, galactic evolution...), i.e. that of stellar variability, where automatic procedures for data processing can help astronomers concentrate on the solution to these problems. We have developed automatic classifiers that, in a matter of seconds or minutes, can automatically assign class probabilities to hundreds of thousands of variable objects, and we have proved that these probabilities are highly reliable for the set of classical variables best studied in the literature. These experiments are repeatable and thus free from human subjectivity. The classifiers show minor discrepancies with the classifications used as a reference in this work (as explained in previous sections) and these discrepancies, when due to the classifiers themselves, need to be corrected for. Until then, users of the publicly available classifiers have to be aware of these minor pitfalls when interpreting their results.  The results presented here suggest that further steps can be taken in the analysis of the resulting samples. Two obvious steps are the search for correlations between subsets of attributes not necessarily of dimension 2, and the study of density plots and clustering results in order to explore the substructure within each variability class. This is the subject of ongoing research in the framework of the CoRoT, Kepler and Gaia missions.  The training set and the classifiers are only the first operational versions developed for the optimization of on-going and future databases such as CoRoT, Kepler or Gaia. Obviously, both the training set and the classifiers will greatly benefit from the analysis of these future databases, especially for those classes underrepresented in terms of the real prevalences. This is where the improvement and correction of the discrepancies mentioned in the previous paragraph will take place. They must be oriented towards obtaining a class definition (training) set that better reproduces the real probability densities in parameter space (the probability of a variable object of class $C_k$ having a certain set of attributes such as frequencies, amplitudes, phase differences, colours, etc). Furthermore, it must be made more robust against overfitting by combining data from various surveys/instruments in such a way that the sampling properties (including measurement errors) have as little an impact on the inference process as possible. We believe that this paper is a crucial starting point in the sense that we have proved the validity of the classifier predictions, and, at the same time, we have identified and pointed out the source of its limitations, thus showing the path to more complete and accurate classifiers. Obviously, it is in the non-periodic and rarer classes that there is more room for improvement.  Finally, there is ongoing development of new versions of the classifiers adapted to handle spectral information making use of VSOP data \\citep{2007A&A...470.1201D} and including one of the features of Bayesian Networks that make them especially suitable for their integration in the framework of Virtual Observatories, i.e. their capacity to draw inferences based on incomplete (missing) data. We strongly believe that the probabilistic foundations of these models (at the basis of these capabilities) provide astronomers with explanations of the inference process very much in line with the reasoning usually used in astronomy.\\\\  In this work We have concentrated on the validation of the developed classifiers, using the OGLE database. This database contains a large number of light curves of different variability types. Existing extractor-type results for the classical pulsators and eclipsing binaries allowed us to judge the quality of our classification results. Our classifiers also identified candidate new members for some of those classes. Little had been done up to now on the multiperiodic pulsators, the most interesting targets from an asteroseismological point of view. The OGLE data are not optimally suited to study those variables, but some types could be studied and discovered. Our classifiers have identified $107$ candidate B-type pulsators (SPB, BCEP and PVSG) in the Magellanic clouds. Those candidates were placed on the HR diagram, to see how they are situated with respect to the instability strips of B-type pulsators. This allowed us to conclude that the present instability computations are incomplete and that their improvement probably needs new input physics. In practice, we provide here a list of new candidate variables of multiperiodic classes (DSCUT, BCEP, SPB and GDOR), including several in the Bulge. A more in-depth analysis of these candidates is needed, but this is outside the scope of this classification work."
    },
    "0806/0806.4273_arXiv.txt": {
        "abstract": "We consider sterile neutrinos with rest masses $\\sim 0.2\\,{\\rm GeV}$ and with vacuum flavor mixing angles $\\sin^2\\theta >10^{-8}$ for mixing with $\\tau$-neutrinos, or $10^{-8}<\\theta^2 <10^{-7}$ for mixing with muon neutrinos.  Such sterile neutrinos could augment core collapse supernova shock energies by enhancing energy transport from the core to the vicinity of the shock front. The decay of these neutrinos could produce a flux of very energetic active neutrinos, detectable by future neutrino observations from galactic supernova. The relevant range of sterile neutrino masses and mixing angles can be probed in future laboratory experiments. ",
        "introduction": "Neutrino masses are usually incorporated into the Standard Model (SM) by the addition of SU(3)$\\times$SU(2)$\\times$U(1) singlet fermions, often called right-handed neutrinos~\\cite{seesaw}. These gauge singlets can have Majorana masses  as low as a few eV~\\cite{deGouvea:2005er}, or as large as the Grand Unified scale~\\cite{seesaw}. If the Majorana mass terms are large, the particles associated with the singlet fields are very heavy. However, if the Majorana masses are below the electroweak scale, the corresponding degrees of freedom appear in the low-energy effective theory as so called \\textit{sterile neutrinos}. Light sterile neutrinos with masses of a few keV could be the cosmological dark matter~\\cite{dw,nuMSM}, their production in a supernova could result in large pulsar kicks~\\cite{pulsars}, and they could affect supernovae in a variety of  ways~\\cite{supernova_misc}. The same particles can play an important role in the formation of the first stars~\\cite{reion}, baryogenesis~\\cite{baryogenesis}, and other astrophysical phenomena~\\cite{Biermann:2007ap}.   In this paper we investigate the effect of heavier sterile neutrinos in astrophysical systems. We show that sterile neutrinos with masses $\\sim 0.2 \\GeV$ and small mixing $\\sin^2 \\theta \\sim 10^{-8}$, with either the muon or tau neutrinos, could be produced in supernova cores and subsequently augment, via their decay, the energy transport from the core to the region around the stalled shock, thereby increasing the prospects for a core collapse supernova explosion. This scenario could be testable, both by future laboratory searches for heavy sterile neutrinos and by observations of the neutrino signal from supernovae.   Our analysis differs from earlier work. Heavy neutral leptons, with masses 10~keV -- 10~MeV, produced in supernovae, have been considered for setting limits and for powering supernova explosions~\\cite{Lepton Decays in SN}. However, the mass range we consider here, 145 -- 250 MeV, is qualitatively different in that the sterile neutrinos decay predominantly into a pion and a light fermion. If the mixing angle with the electron neutrino is negligible, and the sterile neutrino mass $M_s$ is in the range $m_{\\pi^0} < M_s < (m_{\\pi^0}+m_\\mu)$, the daughter pion is the neutral pion, which decays into  two photons: $\\nu^{\\rm (s)} \\rightarrow \\nu^{\\rm (a)} \\pi^0 \\rightarrow \\nu^{\\rm (a)} \\gamma \\gamma$~\\cite{Dolgov:2000jw}, where $a=(\\mu,\\tau) $.  This decay mode, with lifetime $\\sim 0.1\\,{\\rm s}$, changes the impact of sterile neutrinos on the supernova  explosion. To distinguish sterile neutrinos that decay mainly into photons from the other types, we call them {\\em eosphoric}.\\footnote{From the ancient Greek god $E \\omega\\sigma\\varphi \\acute{o} \\rho o \\varsigma$, the bearer of light.} As a consequence of small mixing angles, $\\sin^2 \\theta \\sim 10^{-8} - 10^{-7}$, the eosphoric sterile neutrino production and propagation history inside the supernova environment could be significantly different from those of the neutral leptons with electroweak scale interactions which have been considered previously \\cite{Lepton Decays in SN}.   We assume that the particle physics lagrangian at low energies comprises the Standard Model, albeit modified as indicated. The Standard Model was originally formulated with massless neutrinos $\\nu_i$ transforming as components of the electroweak SU(2) doublets $L_\\alpha$ ($\\alpha =1,2,3$), but here we will extend it to include seesaw mass terms for neutrinos~\\cite{seesaw}, in which we allow the Majorana masses to be below the electroweak scale: \\beq  {\\cal L} = {\\cal L_{\\rm SM}} + i \\bar N_a \\slashed{\\partial} N_a - y_{\\alpha a} H^{\\dag} \\,  \\bar L_\\alpha N_a - \\frac{M_a}{2} \\; \\bar N_a^c N_a + h.c. \\label{L} \\eeq The neutrino mass eigenstates $\\{ \\nu^{\\rm (a)}_1 \\nu^{\\rm (a)}_2 \\nu^{\\rm (a)}_3, \\nu^{\\rm (s)}_1, \\nu^{\\rm (s)}_2, ..., \\nu^{\\rm (s)}_n \\}$  are linear combinations of the weak eigenstates $\\{\\nu_\\alpha, N_a \\}$.  The states $ \\nu^{\\rm (a)}_{1,2,3}$ are {\\em active} and have masses below 0.2~eV, while $ \\nu^{\\rm (s)}_{1,...,n}$ are {\\em sterile}. In particular, several recent studies focus on the $\\nu$MSM~\\cite{nuMSM}, a model with $n=3$ sterile neutrinos: one with  a few ${\\rm keV}$ rest mass (dark matter), and two nearly degenerate, heavier states. This model facilitates leptogenesis via neutrino oscillations~\\cite{baryogenesis}.    We will consider sterile neutrinos with masses 145--250 MeV, and with vacuum flavor mixing angle $\\sin^2\\theta \\sim 10^{-8} - 10^{-7}$ for mixing with either muon or tau neutrinos in vacuum. The current bounds~\\cite{sterile_constraints} allow this range of mixing angles with both muon and tau neutrinos.  For sterile neutrinos mixing only with tau neutrinos an even broader range is allowed~\\cite{sterile_constraints,sterile_exp_constraints}. Sterile neutrinos with somewhat smaller vacuum mixing could produce a pre-nucleosynthesis matter-dominated epoch in the early universe, and this is being investigated separately~\\cite{us-BBN}. In section \\ref{production}, we investigate the production of these neutrinos in the core of a hot proto-neutron star. In section \\ref{shock}, we discuss how the decay of eosphoric neutrinos leads to energy deposition in the mantle above the post-collapse neutron star which could help in shock revival. Section \\ref{nu signal} is devoted to the neutrino signal produced by these decays, and the associated observational signature of this scenario. We show that this is consistent with the neutrino detection from SN1987A. ",
        "conclusions": "Heavy sterile neutrinos could prove to be very important in compact astrophysical systems, such as supernovae, essentially because they could move significant amounts of energy around in ways that ordinary active neutrinos and hydrodynamic motions cannot. In this paper, we have showed that sterile neutrinos with mass $\\sim 200 \\MeV$ and small mixing $\\sin^2 \\theta \\sim 10^{-8}$ with either $\\nu_\\mu$ or $\\nu_\\tau$ could facilitate energy transport from the supernova core to the shock front, possibly ultimately leading to a successful explosion, but in any case altering supernova energetics in ways which change the standard core collapse paradigm, yet produce signatures and behaviors that could remain within existing observational bounds.  The eosphoric sterile neutrino scenario can be tested by future neutrino observations from galactic supernova explosions. A supernova should  produce a short burst of very energetic neutrinos followed by a longer (and much more powerful) signal of lower-energy neutrinos. The non-observation of very energetic neutrinos from SN1987A is consistent with our model because of the low flux. However, present and future neutrino detectors should be well positioned to detect the (less numerous) 80~MeV neutrinos. Observations of a Galactic supernova neutrino signal in these detectors could be used to constrain the eosphoric sterile neutrino parameter space in ways which could extend or be complimentary to future laboratory-based neutrino mixing probes.    \\medskip  The work of A.K. and K.P. was supported in part by the DOE grant DE-FG03-91ER40662 and by the NASA ATFP grant  NNX08AL48G; while G.M.F. was supported in part by NSF grant PHY-04-00359 at UCSD."
    },
    "0806/0806.4103_arXiv.txt": {
        "abstract": "BD+53$^{\\circ}$2790, an O9.5\\,Vp star, is the optical counterpart to the HMXRB 4U~2206+54. This system was classified initially as a BeX, but observational evidence soon stressed the need to revise this classification. The permanent asymmetry in the H$\\alpha$ line profiles (in contrast with the cyclic variations shown by Be stars), the variations in the profile of this line in time scales of hours (while time scales from weeks to months are expected in Be stars), and the lack of correlation between IR observables and H$\\alpha$ line parameters, strongly suggest that, while BD+53$^{\\circ}$2790 contains a circunstellar disc, it is not like the one present in Be stars \\citep{blay05}. Furthermore, there is evidence of overabundance of He in BD+53$^{\\circ}$2790. Together with the presence of an anomalous wind, found through UV spectroscopy, the possibility to link this star with the group of He rich stars is open. We will discuss the work done with {\\it IUE} data from BD+53$^{\\circ}$2790 and the unexpected finding of a slow and dense wind, very rare for an O\\,9.5V star. ",
        "introduction": "BD+53$^{\\circ}$2790 is an early  type star whose association to the High Mass X-Ray Binary system (HMXRB) 4U~2206+54 was proposed for the first time in the work of \\citet{steiner84}. Through {\\it UBV} photometry they estimated an spectral type B1 and a distance between 3.5 and 1.5 kpc. The red spectrum of BD+53$^{\\circ}$2790 showed the H$\\alpha$ line in emission with a shell-like profile (an absorption core over imposed on an emission profile), therefore they concluded that BD+53$^{\\circ}$2790 was a Be star. They failed to realize that a preliminary classification of this star existed in the work of \\citet{hiltner56}. \\citet{hiltner56} classified BD+53$^{\\circ}$2790 as a O9III., but he may not be sure about this classification as he reported the luminosity classification (III) with a question tag. This classification implies a distance of $\\sim$6 kpc to BD+53$^{\\circ}$2790, larger than the one derived by \\citet{steiner84}.  \\cite{negueruela01} proposed a O\\,9.5Vp spectral classification for BD+53$^{\\circ}$2790 and concluded that this star is not a typical Be, but a very peculiar late O star. \\cite{blay05} confirmed this hypothesis and tentatively linked BD+53$^{\\circ}$2790 to the group of He-rich stars. BD+53$^{\\circ}$2790 could be the second representative of this group among O-type stars, being $\\theta^1$Ori C the first one \\citep{donati02,smith05}.  The International Ultraviolet Explorer ({\\it IUE}) observed BD+53$^{\\circ}$2790 in high and low resolution modes. We present in this work a detailed UV analysis of the high resolution {\\it IUE} spectrum SWP39112M, described in \\cite{negueruela01}. We will investigate the stellar wind in BD+53$^{\\circ}$2790 from the resonance ultra-violet lines. ",
        "conclusions": "\\begin{figure}[t] \\includegraphics[width=0.35\\textwidth]{xrays.eps} \\caption{{\\bf Top:} {\\it RXTE}/ASM 2--10 keV binned light curve to the 9.6 d period known from 4U~2206+54. {\\bf Bottom:} Simulated light curve using the \\cite{bondi44} approximation and a dense ($\\dot M \\sim 5\\times10^{-8}$~M$_{\\sun}$~yr$^{-1}$) and slow ($\\sim$350 km s$^{-1}$) wind, figures extracted from \\cite{ribo06}} \\label{fig:wind_and_xray} \\end{figure}  We have seen that BD+53$^{\\circ}$2790 presents a very peculiar wind structure. Contrary to what we would expect for a typical late type main sequence O star \\citep{howarth96} we find in BD+53$^{\\circ}$2790 a very slow and dense wind. This result, as shown in \\cite{ribo06}, is compatible with the observed X-Ray variability from the HMXRB system 4U~2206+54, to which BD+53$^{\\circ}$2790 belongs. In this system the X-ray emission is produced when the wind from BD+53$^{\\circ}$2790 is accreted by its neutron star companion. The potential energy of the wind matter is released as high energy photons when it falls and collides with the neutron star surface. Fig.~\\ref{fig:wind_and_xray} shows how the wind parameters derived from the {\\it IUE} spectrum can reproduce the observed X-Ray variability when we model it according to the \\cite{bondi44} approximation. Only when mass loss rates on the order of 10$^{-8}$~M$_{\\sun}$~yr$^{-1}$ and terminal wind velocities well below 1000 km s$^{-1}$ are used, the observed light curve can be reproduced by the model. {\\nopagebreak We have analyzed here only one {\\it IUE} spectrum and we have found astonishing properties in the wind of BD+53$^{\\circ}$2790. We can expect that the behavior seen in this spectrum is representative of the overall behavior of the source, but for a detailed analysis more observations in the UV range would be needed. The capabilities of the World Space Observatory (WSO) will offer a unique opportunity to study this kind of systems and will add crucial information to unveil the secrets of the mass transfer form the massive companion onto the compact object in HMXRBs."
    },
    "0806/0806.4335_arXiv.txt": {
        "abstract": "A quantization procedure without Hamiltonian is reported which starts from a statistical ensemble of particles of mass $m$ and an associated continuity equation. The basic variables of this theory are a probability density $\\rho$, and a scalar field $S$ which defines a probability current $\\vec{j}=\\rho\\nabla S/m$. A first equation for $\\rho$ and $S$ is given by the continuity equation. We further assume that this system may be described by a linear differential equation for a complex-valued state variable $\\chi$. Using these assumptions and the simplest possible Ansatz $\\chi(\\rho,\\,S)$, for the relation between $\\chi$ and $\\rho,\\,S$, Schr\\\"odinger's equation for a particle of mass $m$ in a mechanical potential $V(q,t)$ is deduced. For simplicity the calculations are performed for a single spatial dimension (variable $q$). Using a second Ansatz $\\chi(\\rho,\\,S,\\,q,\\,t)$, which allows for an explicit $q,t$-dependence of $\\chi$, one obtains a generalized Schr\\\"odinger equation with an unusual external influence described by a time-dependent Planck constant. All other modifications of Schr\\\"odinger' equation obtained within this Ansatz may be eliminated by means of a gauge transformation. Thus, this second Ansatz may be considered as a generalized gauging procedure. Finally, making a third Ansatz, which allows for a \\emph{non-unique} external $q,t$-dependence of $\\chi$, one obtains Schr\\\"odinger's equation with electrodynamic potentials $\\vec{A},\\,\\phi$ in the familiar gauge coupling form. This derivation shows a deep connection between non-uniqueness, quantum mechanics and the form of the gauge coupling. A possible source of the non-uniqueness is pointed out. ",
        "introduction": "\\label{intro} Usually a physical system, which one wants to describe quantum mechanically, is first identified in the context of classical physics and then somehow transferred to the quantum mechanical domain; this process is referred to as \"quantization\". Often the first step in the quantization process is the tacit assumption that a Hilbert space is associated with the examined system. Then, the remaining task is to find the proper algebra of operators. A more direct method which avoids this assumption, is to ``derive'' Schr\\\"odinger's equation, i.e. to find premises which imply Schr\\\"odinger's equation. This may be done in several ways. The method which was historically at the beginning of quantum mechanics~\\cite{schrodinger:quantisierung_I} (\"wave mechanics\") starts from the Hamilton-Jacobi equation of a classical system and tries to deduce from it - with the help of suitable modifications~\\cite{cook:probability,castro.dutra:quantum,elizalde:quantum-hamilton-jacobi} - the corresponding quantum-mechanical equation for the time development of the system. Other premises leading to Schr\\\"odinger's equation include special assumptions about the structure of momentum fluctuations~\\cite{hall.reginatto:schroedinger} and the principle of minimum Fisher information~\\cite{reginatto:derivation}.  In this paper, a new quantization procedure is reported which shares with the last two examples the property that it does not start from a single-particle picture but from a statistical ensemble. The simplest nontrivial system, a spinless particle of mass $m$ in nonrelativistic approximation is investigated. In order to define the subject of this work more precisely we start from the classical Hamilton-Jacobi equation for the action function $S(\\vec{q},t)$, which depends on the particle coordinates $q_k$ and the time $t$. It is given by \\begin{equation} \\label{eq:HJCL} \\frac{\\partial S(\\vec{q},t)}{\\partial t}+\\frac{1}{2m} \\left(\\frac{\\partial S(\\vec{q},t)}{\\partial\\vec{q}}  \\right)^{2} +V(\\vec{q},t)=0 \\mbox{,} \\end{equation} if the movement takes place under the influence of a potential $V(\\vec{q},t)$. The momentum field, that appears in Eq.~(\\ref{eq:HJCL}) is given by \\begin{equation} \\label{eq:momenta} p_k(\\vec{q},t)=\\frac{\\partial S(\\vec{q},t)}{\\partial q_k} \\mbox{.} \\end{equation}  The fact that the Hamilton-Jacobi equation is the ideal starting point for the transition from classical physics to quantum mechanics is formally based on the following well-known reformulation of the time-dependent Schr\\\"odinger equation \\begin{equation} \\label{eq:TDSCHR} \\frac{\\hbar}{\\imath}\\frac{\\partial }{\\partial t}\\psi(\\vec{q},t) + V(\\vec{q},t) \\psi(\\vec{q},t) = \\frac{\\hbar^2}{2m}\\nabla\\psi(\\vec{q},t) \\mbox{.} \\end{equation} If the complex-valued variable $\\psi(\\vec{q},t)$ is, without any restrictions of generality, written in the form \\begin{equation} \\label{eq:PRODANS} \\psi(\\vec{q},t)=\\sqrt{\\rho(\\vec{q},t)}\\mathrm{e}^{\\frac{\\imath}{\\hbar}S(\\vec{q},t)} \\mbox{,} \\end{equation} then one obtains from Eq.~(\\ref{eq:TDSCHR}), by calculating the real parts of both sides, the relation \\begin{equation} \\label{eq:CONT} \\frac{\\partial \\rho(\\vec{q},t)}{\\partial t}+\\frac{\\partial}{\\partial\\vec{q}} \\frac{\\rho}{m} \\frac{\\partial S(\\vec{q},t)}{\\partial\\vec{q}}=0 \\mbox{,} \\end{equation} which is a classical (no $\\hbar$ occurs) continuity equation for the probability density $\\rho$ and the probability current $\\rho\\,\\vec{p}/m$. Equating the imaginary parts of both sides of Eq.~(\\ref{eq:TDSCHR}) one obtains the relation \\begin{equation} \\label{eq:QHJ} \\frac{\\partial S(\\vec{q},t)}{\\partial t}+\\frac{1}{2m} \\left( \\frac{\\partial S(\\vec{q},t)}{\\partial\\vec{q}} \\right)^{2}+V(\\vec{q},t)= \\frac{\\hbar^2}{2m}\\frac{\\triangle\\sqrt{\\rho(\\vec{q},t)}}{\\sqrt{\\rho(\\vec{q},t)}} \\mbox{,} \\end{equation} which differs from the Hamilton-Jacobi equation~(\\ref{eq:HJCL}) only by the single term on the right hand side. Eq.~(\\ref{eq:QHJ}) is sometimes referred to as quantum Hamilton-Jacobi equation.  Due to this similarity there have been attempts~\\cite{castro.dutra:quantum,elizalde:quantum-hamilton-jacobi} to use Eq.~(\\ref{eq:HJCL}) as a starting point and to derive Eq.~(\\ref{eq:TDSCHR}), which presents the basis of quantum mechanics, by justifying introduction of the crucial quantum term appearing in Eq.~(\\ref{eq:QHJ}). In the present work we go a different route, starting from assumptions which are simpler in certain respects. We postulate the existence of a statistical ensemble but do \\emph{not} start from the Hamilton-Jacobi equation itself. Instead, our basic postulate is the validity of a continuity equation (of the above type), interpreted as a local conservation law of probability. Since we have two unknown functions ($\\rho$ und $S$) and only one single equation, we clearly need further assumptions - and a second differential equation - in order to arrive at a mathematically well-defined problem. Our second assumption is very simple and of a purely formal nature. We require that both equations - the one already known and the second one still to be found - may be expressed mathematically as a single equation for a single complex state variable $\\chi$. This second assumption expresses something like the postulate of maximal mathematical simplicity. As we know, this postulate may be quite successful in physics, in particular if combined with other ideas.   Thus, the continuity equation has to be ``extended'' to the complex domain. This task may be described briefly as follows: Consider a complex-valued variable $\\chi$ which depends in an unspecified way on the real variables $S$ and $\\rho$. Which differential equations for $\\chi$ exist, whose real (or imaginary) part agrees with the continuity equation and which functional dependencies $\\chi(\\rho,S)$ are compatible with this requirement ? Note that both the functional dependence of $\\chi$ and the shape of the differential equation are unknown ``variables'' of this problem. A detailed formulation of this problem is reported in the next section~\\ref{sec:2}). The calculation, reported in section~\\ref{sec:3}) and appendix~\\ref{sec:8}, has been performed for simplicity for a single spatial dimension and for a reduced class of differential equations obeying several additional constraints. The result is Schr\\\"odinger's equation for a particle in an external mechanical potential. Further, we formulate and justify in section~\\ref{sec:3}) the conjecture that all additional constraints except linearity may be omitted.  In section~\\ref{sec:4}) the original Ansatz is extended by allowing for an additional, explicit space-time dependence of the state variable $\\chi$. This is our first attempt to derive the minimal coupling rule, which is obeyed by essentially all fundamental interactions, in the present context. It turns out (details of the calculation are reported in appendix~\\ref{sec:9}) that a second, very unusual external influence, besides the potential $V$, appears in the Schr\\\"odinger equation. It takes the form of a time-dependent Planck constant. All other modifications due to the extended Ansatz are spurious, because they may be eliminated from the Schr\\\"odinger equation by a gauge transformation. The gauge field itself cannot be derived by means of this Ansatz. In section~\\ref{sec:5}) our second attempt is undertaken to derive a gauge field. The Ansatz of the last section is once more extended by allowing for a \\emph{non-unique} space-time dependence of $\\chi$. The result is Schr\\\"odinger's equation for a charged particle in an external electromagnetic field.  Section ~\\ref{sec:6}) contains a detailed discussion of all assumptions and results and may be consulted in a first reading to obtain an overview of this work. It also contains remarks of a speculative nature concerning the relation between the classical theory of charged particles and fields on the one hand and the form of quantum mechanics and gauge coupling on the other hand. In the last section~\\ref{sec:7}) one finds concluding remarks. ",
        "conclusions": "\\label{sec:7} In this work Schr\\\"odinger's equation with gauge coupling has been derived. The construction was based on the validity of a continuity equation and a linear differential equation for a complex-valued state variable. As an heuristic recipe this scheme may  probably be extended to more general situations. Obvious possibilities to do that include a relativistic formulation or a multi-component version, which should be the analogon of non-abelian gauge theories. Other types of variables (quaternions) or even an application beyond the realm of special relativity seem conceivable. From a physical point of view, the most important generalization or modification of the present theory concerns the assumption of a complex state variable. This is a purely mathematical assumption which  - despite of the outstanding structural properties of the field of complex numbers - should be replaced by an different, equivalent assumption which can be directly interpreted  in physical terms. This is the most challenging question to be asked in the context of the present theory. \\begin{appendix}"
    },
    "0806/0806.2333_arXiv.txt": {
        "abstract": "We explore the ability of measurements of the 21-cm power spectrum during reionization to enable the simultaneous reconstruction of the reionization history and the properties of the ionizing sources. For various sets of simulated 21-cm observations, we perform maximum likelihood fits in order to constrain the reionization and galaxy formation histories. We employ a flexible six-parameter model that parametrizes the uncertainties in the properties of high-redshift galaxies. The computational speed needed is attained through the use of an analytical model that is in reasonable agreement with numerical simulations of reionization. We find that one-year observations with the MWA should measure the cosmic ionized fraction to $\\sim 1\\%$ accuracy at the very end of reionization, and a few percent accuracy around the mid-point of reionization. The mean halo mass of the ionizing sources should be measurable to $10\\%$ accuracy when reionization is 2/3 of the way through, and to $20\\%$ accuracy throughout the central stage of reionization, if this mass is anywhere in the range 1/3 to 100 billion solar masses. ",
        "introduction": "The earliest generations of galaxies are thought to have heated and reionized the universe. Ly$\\alpha$ absorption shows that the IGM has been a hot plasma at least since $z \\sim 6.5$ \\citep{fan06}, while the five-year WMAP measurements of the large-angle polarization of the cosmic microwave background imply that the universe was significantly ionized within the redshift range $z \\sim 8$--14 \\citep{wmapRei}. The same WMAP measurements are also consistent with the $\\Lambda$CDM model and, together with distance measurements from supernovae and baryon acoustic oscillations from galaxy surveys, have helped to determine rather accurately the standard cosmological parameters of this model \\citep{wmap}; we thus assume the $\\Lambda$CDM model with density parameters $\\Omega_m=0.28$ (dark matter plus baryons), $\\Omega_\\Lambda=0.72$ (cosmological constant), and $\\Omega_b=0.046$ (baryons), together with $h=0.7$ (dimensionless Hubble constant), $n=0.96$ (power spectrum index), and $\\sigma_8=0.82$ (power spectrum normalization).  Observational study of reionization promises to teach us a great deal about the early generations of galaxies that formed when the universe was between $\\sim 300$ and 800 Myr in age. An important feature of reionization is that it indirectly probes whatever are the dominant sources of ionizing radiation, even if these are otherwise unobservable; e.g., the universe may have been ionized by large numbers of very small galaxies that are too faint to be detected individually (or even by unexpected sources such as miniquasars or decaying dark matter). The overall timing of reionization versus redshift mainly constrains the overall cosmic efficiency of ionizing photon production. In comparison, a detailed picture of reionization as it happens can teach us a great deal about the sources that produced this cosmic phase transition, particularly in the case of the most natural source, stars in galaxies.  A key point is that the spatial distribution of ionized bubbles is determined by clustered groups of galaxies and not by individual galaxies. At such early times galaxies were strongly clustered even on rather large scales (tens of comoving Mpc), and these scales therefore dominate the structure of reionization \\citep{BLflucts}. Overdense regions fully reionize first because the number of ionizing sources in these regions is increased very strongly \\citep{BLflucts}. The large-scale topology of reionization is therefore inside out, with underdense voids reionizing only at the very end of reionization, with the help of extra ionizing photons coming in from their surroundings. This picture has been confirmed and quantified more precisely in detailed analytical models that account for large-scale variations in the abundance of galaxies \\citep{fzh04}, and in a number of large-scale numerical simulations of reionization \\citep[e.g.,][]{zahn, iliev, santos}.  Cosmic reionization represents an extreme challenge for numerical simulations, because of the enormous range of spatial scales involved \\citep{BLflucts}. On the one hand, individual galaxies as small as a comoving kpc may contribute to the early stages of reionization, while ionizing photons may travel as much as 100 Mpc before being absorbed, near the end of reionization (or earlier if X-rays make a significant contribution). Even smaller scales must be resolved in order to probe the dense gas clumps that likely determine the ionizing mean free path and the recombination rate, while star formation and stellar feedback are still further out of reach. While simulations are improving and may soon include hydrodynamics (in addition to the current N-body plus radiative transfer codes), some have tried to bridge the gap of scales by developing fast semi-numerical methods that make substantial use of approximate analytical models \\citep{zahn, mesinger}.  The most promising probe of the cosmic reionization history is to use new low-frequency radio telescope arrays to detect emission in the redshifted 21-cm line corresponding to the hyperfine transition of atomic hydrogen. 21-cm cosmology is potentially also a great source of fundamental cosmological information, especially if observations reach small scales and high redshifts. The 21-cm fluctuations can in principle be measured down to the smallest scales where the baryon pressure suppresses gas fluctuations, while the cosmic microwave background (CMB) anisotropies are damped on much larger scales (through Silk damping and the finite width of the surface of last scattering). Since the 21-cm technique is also three-dimensional (while the CMB yields a single sky map), there is a much larger potential number of independent modes probed by the 21-cm signal, which could help to detect primordial non-Gaussianity and test inflation \\citep{zald}. However, ionization fluctuations dominate the 21-cm fluctuations during the epoch of reionization, and thus the first generation of 21-cm experiments are expected to bring new discoveries related to the reionization history rather than the fundamental cosmological parameters \\citep[e.g.,][]{mcquinn06}. Furthermore, since the observational noise will be too large to produce 21-cm maps, the fluctuations must be measured statistically, and the power spectrum is both the most natural and the highest signal-to-noise statistic. Upcoming experiments include the Murchison Widefield Array (MWA)\\footnote{http://www.haystack.mit.edu/ast/arrays/mwa/} and the Low Frequency Array (LOFAR)\\footnote{http://www.lofar.org/}.  While theorists and numerical simulators have begun to elucidate the relation between the properties of the ionizing sources and the 21-cm power spectrum, a key question has not been addressed thus far: assuming the upcoming observations measure the 21-cm power spectrum as expected, how is the power spectrum to be inverted into a determination of the properties of the sources? Such an inversion problem is usually solved by a maximum likelihood (or $\\chi^2$) procedure whereby a model is fit to the observed power spectrum in order to determine the best-fit parameters and their uncertainties. In order to explore maximum likelihood fitting of simulated observations, a flexible model is needed that can quickly yield the 21-cm power spectrum predicted for given parameters of the ionizing galaxies. It is important to have a flexible model that does not presume that we can theoretically predict the properties of the ionizing galaxies, which depend on many complex feedback processes. This ``reionize-fast'' code would essentially play the same role that CMBFAST \\citep{cmbfast} did for analyses of measurements of the CMB angular power spectrum.  Ultimately, this type of code will most likely be developed from an analytical model that includes as much of the detailed physics as possible and is also partly tuned to fit more accurately the results of numerical simulations, analogous to the way that the formula of \\citet{Sheth} for the halo mass function was developed from the original model of \\citet{PS}.  In this paper we employ the model from \\citet{b07} in which we solved for the correlated two-point distribution of density and ionization based on the one-point model of \\citet{fzh04}. This is currently the most realistic fully analytical model of the 21-cm power spectrum. In the next section we briefly review the model before comparing its predictions to 21-cm power spectra from several numerical simulations of reionization. In the following section we use the model to summarize which galaxy parameters affect the 21-cm power spectrum, and then calculate the expected uncertainties from maximum likelihood fits to simulated sets of observed power spectra. ",
        "conclusions": "We have used the most realistic fully analytical model available for the 21-cm power spectrum \\citep{b07} to fit models of the galaxy population during reionization to simulated 21-cm power spectrum observations.  The model assumes at each redshift a fixed ionizing efficiency $\\zeta$ and a minimum halo circular velocity $V_{\\rm c}$ for galactic halos.  Allowing each of these quantities to vary linearly or quadratically with redshift yields reionization models with up to 6 parameters, which we allow to vary freely without any restrictions based on specific models of feedback.  Before proceeding, we compared the analytical model of \\citet{b07} to the results of three different groups that have run N-body simulations, processed the outputs with radiative transfer, and calculated the 21-cm power spectrum at various redshifts during the reionization epoch. The simulations are all in reasonable agreement with the analytical model (Figures~\\ref{f:test1}--\\ref{f:test3}), with typical differences in $\\Delta_{21}$ in the range $\\sim 10-30\\%$. While the analytical model makes various simplifying assumptions, the simulations are also limited by box size and radiative transfer resolution. The analytical model generally captures the evolution of the 21-cm power spectrum during reionization as seen in the simulations, but a more precise comparison must await a demonstration that the simulations have numerically converged and are consistent with each other.  As discussed in earlier sections, the model used here should not be considered the ultimate model to use in fitting the 21-cm power spectrum during reionization, but as an important step towards a final model to be constructed using guidance from numerical simulations. The main qualitative limitation of the model used here is that while it allows a redshift dependence in the properties of galactic halos, the parameters are limited to being spatially uniform at each redshift. A more realistic model would add the possibility of spatially inhomogeneous (in particular, density-dependent) feedback. However, we note that any such extension, which will likely add substantial complexity to the model, must still allow a relatively quick calculation of the 21-cm power spectrum in order to permit a maximum likelihood analysis. It will also be important to keep the model flexible without relying too heavily on results of particular models or simulations, where feedback processes can only be included with limited and approximate methods. We note that \\citet{mcquinn07} considered the effect of minihalos and Lyman-limit systems in limiting the mean free paths of ionizing photons, and showed that their effect on the 21-cm power spectrum at a given $\\bar{x}^i$ is rather small and is much less significant than the effect of the halo mass of the ionizing sources.  The maximum likelihood fitting yields good grounds for optimism. While our ignorance regarding the properties of the ionizing sources has a substantial effect on the expected errors, we still conclude that the expected measurements of the 21-cm power spectrum will enable us to reconstruct both the reionization history and the properties of the sources. In particular, even with a 6-parameter model that allows for a fairly large parameter space of galaxy properties, the one-year MWA observations allow a remarkably precise reconstruction of reionization.  As a specific example of the expected errors, if reionization ends at $z=6.5$ and is dominated by intermediate-mass halos (minimum halo circular velocity $V_{\\rm c} = 35$ km/s, corresponding to a mean halo mass of ionizing sources $\\langle M \\rangle \\sim 3 \\times 10^9 M_\\odot$ at the midpoint of reionization), then the cosmic mean ionized fraction can be measured to $0.3\\%$ accuracy at the very end of reionization, to a relative accuracy of a few percent around the mid-point of reionization, and better than $10 \\%$ as early as a cosmic mean ionized fraction of $\\bar{x}^i = 10\\%$. Also, the mean halo mass of the ionizing sources can be measured in this case to $5\\%$ accuracy when reionization is 2/3 of the way through, and to $20\\%$ accuracy throughout the last 2/3 of reionization (i.e., when $\\bar{x}^i$ between 1/3 and 1). The errors in general increase with the redshift at which reionization ends, and decrease with the halo mass of the dominant ionizing sources.  The best-measured point of reionization is around the 2/3 mark in terms of precision in $\\langle M \\rangle$, and near the very end of reionization in terms of precision in $\\bar{x}^i$. The errors, though, are fairly small in the central and late stages of reionization for all the models that we have examined (see especially Figure~\\ref{f:hiz}), which include halos that range from the atomic-cooling minimum ($V_{\\rm c} = 16.5$ km/s, $\\langle M \\rangle \\sim 4 \\times 10^8 M_\\odot$) to 200 times more massive halos ($V_{\\rm c} = 125$ km/s, $\\langle M \\rangle \\sim 8 \\times 10^{10} M_\\odot$), examined for reionization that ends at $z \\sim 6.5$ or $z \\sim 8$.  We thus conclude that if the upcoming 21-cm experiments, after foregrounds are removed and instrumental systematics are dealt with, reach anywhere near their expected sensitivity, then they will allow us to study high-redshift astrophysics in unprecedented detail."
    },
    "0806/0806.2796_arXiv.txt": {
        "abstract": " ",
        "introduction": "Two meetings of interest to close binaries took place during the reporting period: A full day session on short-period binary stars -- mostly CV's -- (\\cite[Milone et al. 2008]{Milone2008}) during the 2006 AAS Spring meeting in Calgary and the very broadly designed IAU Symposium 240 in Prague in 2006, with many papers on close binaries \\cite[(Hartkopf et al.\\ 2007)]{Hartkopf:07}. In addition, the book by \\cite{Eggleton2006}, which is a comprehensive summary of evolutionary processes in binary and multiple stars, was published.  The report that follows consists of individual contributions of the Commission 42 Organizing Committee members. Its goal has been to give very personal views of a few individuals who are active in the field, so the report does not aim at covering the whole field of close binaries. ",
        "conclusions": ""
    },
    "0806/0806.0793_arXiv.txt": {
        "abstract": "{} {We study the timing and spectral properties of the intermediate polar MU Camelopardalis (1RXS J062518.2+733433) to determine the accretion modes and the accretion geometry from multi-wavelength, multi-epoch observational data.} {Light curves in different observed energy ranges (optical, UV, X-ray) are extracted. The timescales of variability in these light curves are determined using Analysis of Variance. Phase-resolved X-ray spectra are created with respect to the most prominent detected periodicities and each fitted with an identical model, to quantify the differences in the fitted components.} {The published tentative value for the spin period is unambiguously identified with the rotation period of the white dwarf. We detect a distinct soft X-ray component that can be reproduced well by a black body. The analysis of data obtained at different epochs demonstrates that the system is changing its accretion geometry from disk-dominated to a combination of disk- plus stream-dominated, accompanied with a significant change in brightness at optical wavelengths.} {} ",
        "introduction": "\\begin{figure*}[htb] \\begin{center} \\includegraphics[height=0.9\\linewidth, angle=-90,clip=]{lc_comp_rx0625_2_60s.eps} \\end{center} \\caption{\\label{fig:lc2005}The light curves of MU Cam, obtained simultaneously on March 31, 2005, by XMM (PN and OM) and with the AIP 70cm-telescope (R). The data shown here are binned in time, 60s for the first, second, and forth panel (counted from top to bottom), 120s for the lower two. The R-band data are unbinned.} \\end{figure*}  \\begin{figure*}[htb] \\begin{center} \\includegraphics[height=0.9\\linewidth, angle=-90,clip=]{lc_OM_PN_MOS_comb.eps} \\end{center} \\caption{\\label{fig:lc2006}The light curves of MU Cam, obtained on April 6, 2006, by XMM (PN, MOS1, MOS2, and OM). The data shown here are binned to 120s.} \\end{figure*}  Intermediate Polars (IPs) are close binaries that consist of a magnetic white dwarf and a Roche lobe-filling late-type main-sequence star. Mass loss from the secondary star occurs via the inner Lagrangian point, and the mass is accreted by the white dwarf. In contrast to the true Polars, the white-dwarf rotation is unsynchronised with the orbital motion but is more rapid due to the angular momentum gained from the accreted matter.  In the case of a high accretion rate or a weak magnetic field, the magnetosphere will be smaller than the circularisation radius and an accretion disk will form. The inner disk is truncated by the magnetic field of the white dwarf, which channels the matter via accretion curtains to the magnetic poles. Accretion arcs are formed along the footpoints of accreting field lines. Along these arcs, the accretion rate is assumed to vary.  The heated plasma in the accretion zone is a prominent site of hard X-ray radiation, which due to the lighthouse effect is modulated on the spin period, $\\omega$. Because of irradiation, the inner-disk rim may emit radiation that is modulated on this period.  For stronger magnetic fields, some or all of the material from the infalling stream might be directly captured by the field without going through an intermediary accretion disk. In the case of this stream-fed accretion, material transferred to the white dwarf retains the orbital motion of the secondary star, producing periodic variability of the beat frequency $\\omega$-$\\Omega$, or other sidebands with the orbital frequency $\\Omega$.  Most IPs accrete predominantly by means of an accretion disk, V2400 Oph being a rare example of a system with pure stream-fed accretion \\citep{hellierbeardmore2002}. A number of IPs show evidence for accretion by means of both mechanisms simultaneously: via accretion curtains beginning at the inner disk rim, and via so-called disk-overflow when a part of the ballistic stream bypasses the accretion disk outside the orbital plane. In some systems repeated changes between the different accretion states were observed (TX\\,Col, \\citealt{norton1997}; FO\\,Aqr, \\citealt{evans2004a}). Principal evidence for stream-overflow in IPs has been based on indirect means, i.e.~the detection of certain sideband frequencies in their power spectra, which is not always unambiguous. For instance, sideband frequencies may also be produced by spin-pulsed light reprocessed in orbitally-fixed structures (a problem mainly in the optical and UV); alternatively, amplitude modulation of the orbital period (e.g.~due to X-ray absorption) may introduce spurious side-bands \\citep{warner1986}.  An article by \\citet{patterson1994} introduces IPs in general, whereas \\citet{hellier2002} discusses the mechanisms of stream-fed accretion.  The X-ray source MU Camelopardalis (the new name according to \\citet{2006IBVS.5721....1K} for the source formerly known as 1RXSJ062518.2+733433) was identified as an intermediate polar by \\citet{sab2003} and \\citet{staude2003a}. Optical photometry showed strong variability on mainly two timescales, which were tentatively identified with the orbital period of the system and the spin period of the white dwarf, \\mbox{$P_{\\rm orb} = 16987(23)$\\,s$~= 0.19661(27)$\\,d} and \\mbox{$P_{\\rm spin} = 1187.246(4)$\\,s$~= 0.01374127(5)$\\,d}, respectively \\citep{staude2003a}. It remained an open question whether the measured spin period provided the white-dwarf rotation, if it was twice the true rotation value, or if it was a side-band period with the orbital variation.  MU Cam is one of a handful of IPs that show a prominent soft X-ray component additional to the hard component \\citep{staude2003a}. An XMM-Newton observation was proposed to identify undoubtedly the spin period of the white, and to determine the origin of the soft X-ray radiation.  We report on X-ray and UV observations of MU Cam with XMM-Newton on both March 31, 2005, and April 6, 2006, and coordinated ground-based optical photometry. ",
        "conclusions": "We have presented the analysis of comprehensive data sets obtained simultaneously at optical, ultraviolet, soft and hard X-rays in 2005 and 2006 with the XMM-Newton observatory and the AIP 70cm telescope. Our main results and conclusions may be summarized as following:  \\begin{itemize}  \\item The period tentatively derived as the white dwarf spin by \\citet{staude2003a} is the dominant signal in the hard X-ray light curves. We claim that this provides a measure of the white-dwarf rotation. By linear interpolation of the timings of the mean hard X-ray pulse maxima in the two XMM-Newton observations we measure a spin ephemeris of $BJED(\\rm max,PN_{hard}) = 2\\,453\\,461.4475(8) + E\\times0.01374120(6)$  \\item A distinct soft component in the X-ray emission, which was already visible in the ROSAT spectrum, was unambiguously detected and could be fitted well with a black body of temperature 59\\,eV and 54\\,eV in 2005 and 2006, respectively. MU Cam could be reliably identified as a soft IP.  \\item MU Cam shows changes in its accretion state between disk-dominated and disk with an additional stream component during the observing interval from 2003 to 2006. These are displayed by a change in the mean optical brightness of about 1\\,mag, and by fundamental changes in the power spectra throughout the complete observed energy range. Such changes were also reported for other IPs (e.g.~TX Col: \\citealt{norton1997}, FO Aqr: \\citealt{beardmore1998}, \\citealt{evans2004a}), and therefore appear to be a prevalent feature of intermediate polars.  \\item The occurrence of the $2\\omega$ frequency in the X-ray data from 2006 and the strength of $2\\omega$-$\\Omega$ in 2005 suggest two accretion regions at opposite positions with similar properties and visibility, at least in certain energy ranges. This may be achieved if the dipole axis of the magnetic field of the white dwarf was strongly inclined away from the rotation axis to the orbital plane.  \\item The strong modulation in all wavelength ranges implies a rather high inclination, whereas the absence of an X-ray (white dwarf) eclipse limits it to $\\lesssim75\\degs$. It is probably even lower, because no eclipse is detected in the optical, which would be a sign of a partial obscuration of the accretion disk.  \\item To improve our understanding of the properties of the system, future observations should include more optical photometry to refine the spin and orbital ephemerides. For studying the accretion geometry, time-resolved optical high-resolution spectroscopy will be crucial, and it may help to determine an accurate orbital ephemeris. To determine the distance, near-infrared spectroscopy should be performed to be able to determine the spectral type of the secondary star and its contribution to the spectral energy distribution.  \\end{itemize}"
    },
    "0806/0806.2872_arXiv.txt": {
        "abstract": "We analyse a sample of 52,000 Milky Way (MW)  type galaxies drawn from the publicly available galaxy catalogue of the Millennium Simulation with the aim of studying statistically the differences and similarities of their properties in comparison to our Galaxy. Model galaxies are chosen to lie in haloes with maximum circular velocities in the range 200-250 km s$^{-1}$ and to have bulge-to-disk ratios similar to that of the Milky Way. We find that model MW galaxies  formed `quietly' through the accretion of cold gas and small satellite systems. Only $\\approx 12$  per cent of our model galaxies experienced a major merger during their lifetime. Most of the stars formed `in situ', with only about 15 per cent of the final mass gathered through accretion. Supernovae and AGN feedback play an important role in the evolution of these systems. At high redshifts, when the potential wells of the MW progenitors are  shallower, winds driven by supernovae explosions blow out a large fraction of the gas and metals. As the systems grow in mass, SN feedback effects decrease and AGN feedback takes over, playing  a more important role in   the regulation of the star formation activity at lower redshifts. Although model Milky Way galaxies have been selected to lie in a narrow range of maximum circular velocities,  they nevertheless exhibit a significant dispersion in the  final stellar masses and  metallicities. Our analysis suggests that this dispersion results from the different accretion histories of the parent dark matter haloes. Statically, we also find evidences to support the Milky Way as a typical Sb/Sc galaxy in the same mass range, providing a suitable benchmark to constrain numerical models of galaxy formation. ",
        "introduction": "\\label{intro}   Our own Galaxy (the Milky Way) has always represented a challenge for galaxy formation theories. Being the only system for which we can access full phase-space information for a significant number of individual stars, the Milky Way represents an important benchmark for theoretical models (Perryman et al. 2001; Beers et al. 2004; Everdasson et al. 1993; Steinmetz et al. 2006; Ivezi\\'c et al. 2008). However, it is only one galaxy among many, and may not be representative of `typical' spiral galaxies in the same mass range (Hammer et al. 2007). Hence, a statistical analysis of MW type galaxies could help us to understand to which extent we can rely on the Milky Way to set constrains for galaxy formation models.  The first model for the formation of our Galaxy was proposed by Eggen, Lynden-Bell \\& Sandage (1962), who argued that the observed relation between the metallicity and the orbital eccentricity of a sample of about 200 stars could be interpreted as a signature of a rapid radial collapse which led to the formation of the stellar halo. Searle \\& Zinn (1978) later proved this scenario to be inconsistent with the observation of a negligible metallicity gradient for the globular cluster population at large galactocentric distances. These authors proposed an alternative scenario in which the stellar halo of the Galaxy formed through accretion of smaller galactic systems.  This picture is in qualitative agreement with expectations from the Cold Dark Matter (CDM) model and with the observed signatures of substructure in the stellar halo of the Milky Way, which appears to be a complex dynamical system still being shaped by merging of smaller neighbouring galaxies (e.g. Vivas \\& Zinn 2006).  Numerical simulations of structure formation in a CDM Universe indicate the important role of the merging histories of dark matter haloes in determining the structure and motions of stars within galaxies. These simulations imply that the last major merger event in our Galaxy should have occurred at $z >1$, otherwise the very thin cold disc observed in the Galaxy would have been destroyed (Navarro et al. 2004; Kazantzidis et al. 2008). Bekki \\& Chiba (2001) have also shown that dissipative mergers with gas rich systems could have generated halo stars before the formation of the Galactic disk. More recent studies using an hybrid approach that combines $N$-body simulations and semi-analytic techniques (Font et al. 2006; De Lucia \\& Helmi 2008) have suggested that the stellar halo of the Galaxy formed from the accretion of a few relatively massive satellites ($10^{8} - 10^{9} M_{\\odot}$) at early times ($>$ 9 Gyr).  Although the basic cosmological paradigm appears to be well established, and supported by a large number of observational results, our understanding of the physics of galaxy formation is still far from complete. Within the currently accepted paradigm, galaxies form when gas condenses at the centre of dark matter haloes, which assemble in a `bottom-up' fashion with smaller systems forming first and  merging later into larger structures. The evolution of the baryonic components is dominated by complex physical processes (e.g. star formation, supernovae and AGN feedback, chemical enrichment, etc.) which are poorly understood from both the observational and the theoretical viewpoint. The morphology, dynamics and chemistry of a galaxy is the result of many intertwined processes.  In this complex framework, a number of questions still remain to be answered: how did the Galaxy assemble? How `typical' is the Galaxy in the Local Universe?  Which physical processes play a role in determining its physical and chemical properties? Which kind of merging histories lead to the formation of galactic systems similar to our Galaxy?  In this work we will address some of these questions by taking advantage of one of the largest cosmological simulations of structure formation carried out so far, the Millennium Simulation, which is combined with a semi-analytic model of galaxy formation (for a recent review on these techniques, see Baugh 2006). The aim of this paper is to explore the formation histories of Milky Way-type galaxies and to analyse the origin of the dispersion in their physical properties. In order to achieve this goal we study simultaneously the assembly and chemical evolution of model galaxies, and their location on the well-known correlation between stellar mass and metallicity (e.g. Lequeux et al. 1979; Tremonti et al. 2004; Lee et al. 2006).  This strong correlation has been proved to evolve with redshift in such a way that, at a given stellar mass, the gas-phase metallicities of galaxies were lower in the past (e.g. Savaglio et al.  2005; Erb et al. 2006).  Studying the evolution of the mass-metallicity relation as a function of cosmic time can provide important information on the physical processes responsible for the joint evolution of the chemical and dynamical properties of galaxies, e.g supernovae and AGN feedback, star formation and mergers (e.g. Tissera et al. 2005; Brooks et al. 2007; De Rossi et al. 2007; Finlator et al. 2008).  This paper is organised as follows.  In Section 2 we give a brief description of the simulation and of the semi-analytic model used in our study. In Section 3 we study the main physical properties of model Milky-Way type galaxies at $z=0$, while in Section 4 we study their assembly and formation histories. In Section 5 we analyse the influence of different assembly histories on the chemical properties of model galaxies, as a function of redshift. Finally, we summarise our findings in Section 6. ",
        "conclusions": "In this paper we have studied the assembly and chemical evolution of MW type galaxies. We used the galaxy catalogue built by DLB07 for the Millennium Simulation, and selected MW type galaxies by imposing circular velocity and bulge-to-disk ratio constrains.  Our model MW type galaxies have mean properties such stellar mass, gas fraction,  age and star formation rate,  in good agreement with the estimations obtained  for our own Galaxy. Interestingly, there is also substantial dispersion in these quantities which, accordingly to our work, can be related to their history of assembly.  Taking advantage of the publicly available database of merger trees, we studied the assembly and chemical enrichment histories of our model MW type galaxies. Our main results can be summarised as follows:  1. Most of the final stellar mass of model MW type galaxies is formed {\\it in situ} from gas infalling from the surrounding halo. A small fraction of the final stellar mass (about 15 per cent) is formed in smaller galaxies that are accreted over the lifetime of the Milky Way. Only 12 per cent of our model MW type galaxies experienced a major merger during their lifetime, and for only 3.23 per cent of our sample was this major merger the last accretion event.  2. The distribution of baryons in different components is regulated by feedback processes. At $z > 4$, supernovae feedback is effective in ejecting large fractions of the gas outside the haloes (due to the shallower potential wells). By $z\\sim 1$, a large part of this gas has been re-incorporated and the suppression of cooling flows by AGN feedback starts playing a more important role, keeping an important  fraction of the baryons in the hot phase.  3. The MZR of model MW  type galaxies has a dispersion of $\\sim 0.10$ dex, in agreement with the observed results by Tremonti et al. (2004). We found that, at a given stellar mass, the main parameter determining this dispersion is the gas richness of the systems. Gas-rich systems tend to be more metal-poor, while gas-poor galaxies have converted most of their cold gas component into stars in the past and, therefore have  reached a higher level of chemical enrichment.  4. The accretion histories of the haloes hosting our MW type galaxies exhibit a large dispersion.  Haloes hosting gas-rich MW progenitors at high redshifts tend to experience a higher accretion rate  at later times. These differences in the accretion histories of the parent haloes introduce differences in the  star formation rates of the progenitors which, on their turn,   modulate the impact of supernovae and AGN feedbacks. This leads to an important dispersion in the stellar masses and metallicities of the z=0 Milky Way systems.  5.  If we restrict the model MW galaxies to satisfy also observational constrains on stellar mass and gas fractions, we get a smaller sample with similar trends to those of the complete MW type sample reflecting the fact that the dark matter halo is the dominant factor determining the history of assembly of galaxies.  However, the dispersion in the metallicity and gas fraction at a given mass is produced by slightly differences in the accretion rates of substructures which regulate the star formation activity.   Our findings suggest that  part of the dispersion observed in the MZR could be revealing differences in the  histories of formation. Also, our results suggest that the Galaxy may be consider a typical Sb/Sc galaxy in the same mass range, providing  a suitable benchmark for numerical models of galaxy formation."
    },
    "0806/0806.3045_arXiv.txt": {
        "abstract": "Observations from the Hinode/XRT telescope and STEREO/SECCHI/EUVI are utilized to study polar coronal jets and plumes. The study focuses on the temporal evolution of both structures and their relationship. The data sample, spanning April 7-8 2007, shows that over $90\\%$ of the 28 observed jet events are associated with polar plumes. EUV images (STEREO/SECCHI) show plume haze rising from the location of approximately $70\\%$ of the polar X-ray (Hinode/XRT) and EUV jets, with the plume haze appearing minutes to hours after the jet was observed. The remaining jets occurred in areas where plume material previously existed causing a brightness enhancement of the latter after the jet event. Short-lived, jet-like events and small transient bright points are seen (one at a time) at different locations within the base of pre-existing long-lived plumes. X-ray images also show instances (at least two events) of collimated-thin jets rapidly evolving into significantly wider plume-like structures that are followed by the delayed appearance of plume haze in the EUV. These observations provide evidence that X-ray jets are precursors of polar plumes, and in some cases cause brightenings of plumes. Possible mechanisms to explain the observed jet and plume relationship are discussed. ",
        "introduction": "Recent space missions, such as Hinode \\citep{Kosugi07} and STEREO \\citep{Kaiser08}, and ground-based facilities such as SOLIS \\citep{Keller03} provide a set of data unprecedented in quality and cadence. The complementary observations from the different instruments provide the necessary spatial, temporal and temperature coverage to observe the dynamics of jets and polar plumes, helping to form a more complete picture of these structures.  X-ray jets occur almost everywhere in the solar corona \\citep[see][]{Shibata92}, in particular in the polar holes. They are characterized by their transient nature and often appear as collimated high-temperature emissive beam guided by open magnetic flux \\cite[length of $10^5-10^6$ km and collimated widths of $\\sim10^4$ km; see][]{Cirtain07}. \\cite{Cirtain07} reported that the plasma outflow speeds within X-ray jets range from $\\sim100$ to $\\sim1000$ km~s$^{-1}$ and that Alfv\\'en waves are responsible for the high outflow velocities.  In contrast, polar coronal plumes are observed to be hazy in nature without sharp edges, as seen in extreme ultra-violet (EUV) images from SOHO/EIT \\citep{Boudin95} and STEREO/SECCHI/EUVI \\citep{Howard08}. Plumes are also observed to be significantly wider than X-ray jets \\cite[$\\sim20-40$~Mm; see][]{Wilhelm06} and reach several solar radii in height \\cite[see][]{DeForest97}. Plumes are brighter, cooler and the plasma outflows are smaller than in inter-plumes \\citep[see][]{DeForest97,Wilhelm98,Raouafi07}.   Recent studies of jets and polar plumes \\citep[X-ray and EUV; see][]{Wang98,MInsertis08} treat these coronal structures independently and the relationship between them is not investigated. The present research is motivated by the fact that polar X-ray and EUV jets and plumes usually share common properties. Both are episodic in nature and occur at magnetic field concentrations that coincide with the chromospheric network where both structures form through flux emergence \\citep[see][]{Canfield96,Wang98}. Studying the relationship between jets and plumes is important to understand their formation processes, evolution and the eventual contributions to the solar wind and heating of the plasma in the polar coronal holes. The present work is motivated by the observations of polar jets evolving into plumes such as the one shown in Fig.~\\ref{EUV_fig_jet_plume_070407_2200}. The aim of the paper is to investigate the relationship between these prominent coronal structures. ",
        "conclusions": "X-ray and EUV observations indicate that more than 90\\% of the jets observed in the southern polar hole on April 7-8, 2007 are associated with plume haze. 70\\% of these jets are followed by polar plumes with a time delay ranging from minutes to tens of minutes. Emission of pre-existing plumes is enhanced after every jet eruption within their base. A number of prominent plumes (e.g., P$_{07}$ and P$_{08}$) show evidence for short lived, jet-like events in the EUV that occur within the plume base (see the sharp structures Fig.~\\ref{fig_letter_080407_0332UT}(f) \\& (h) and the several bright points in panel (i)). Jet-like events ensure the continuous rise of haze and may contribute to the change in plume brightness \\cite[see][]{DeForest97}.   The event xj$_7$ in Fig.~\\ref{fig_letter_080408_2240UT} is an interesting case. It was observed in X-rays from 21:58 - 22:16 UT on April 8, 2007. Fig.~\\ref{fig_letter_080408_2240UT}(d-e) shows an EUV collimated structure similar to the one observed in X-rays more than three hours earlier. This may be caused by the plasma being heated to several MK and then becoming visible in X-rays, then gradually cooling down until it appears in the EUV range. More data needs to be analyzed to confirm the plausibility of this hypothesis.  The event xj$_9$, illustrated by Fig.~\\ref{fig_letter_080408_2240UT}(c), is also peculiar and lasted less than 30 minutes. A narrow, collimated beam of plasma rose from the left edge of the large bright point with a shape typical of X-ray jets. It evolved rapidly and after 4-5 minutes the base width of the emission began to widen to cover the whole bright point. The width of the emitting structure exceeded 20~Mm, which is the typical width of polar plumes \\cite[see][]{Wilhelm06}. EUV images showed a faint haze several hours after the X-ray event (see Fig.~\\ref{fig_letter_080408_2240UT}d-i). GONG magnetograms show that the flux at the base of xj$_9$ weakened during the event's lifetime. We believe that the initial jet event evolved into a plume due to significant emerging magnetic flux causing a catastrophic magnetic reconnection on a relatively short time scale but over a large spatial area. This may allow dissipation of the magnetic energy budget of the structure over a short period of time with an associated ejection of a significant amount of material over a relative large spatial scale, unlike other jet-plume events that develop over intervals of several hours. This type of event is recorded twice in the data set utilized here.  It is likely that jets play a key role in the formation process of polar plumes. Both coronal structures share numerous common characteristics, i.e., a magnetic field of mixed polarities at the base, leading to magnetic reconnection. We believe that the magnetic flux emergence causes the jet, opening of previously closed flux results in plume. Jet eruption seems to be the result of gradually emerging magnetic flux from the solar interior that suddenly reconnects on a small scale with the ambient photospheric field, leading to a collimated beam of plasma rising in the corona \\citep[e.g.,][]{Yokoyama95}. EUV images show that coronal plume haze is observed following the jet events. They also provide evidence for several small bright points and short-lived, jet-like events within the base of the plume. These may be the results of magnetic reconnection at smaller spatio-temporal scales that modulate and sporadically brighten pre-existing polar plumes. This is most often seen in long-lived polar plumes, since several phases of reconnection can develop in a single long-lived structure. However, fast opening of magnetic flux can allow a plume to develop almost immediately such as in the case of the xj$_9$ event.  The transition from fast, impulsive, magnetically-driven dynamics of reconnection to the thermal expansion of newly liberated gas along opened magnetic field could explain the time delay observed between the jet and plume events. On the one hand, the jet eruption is the result of fast and explosive dissipation of magnetic energy on a short time scale. On the other hand, the plume might be a result of a pressure gradient within the open flux, which would lift the plume material in the corona. This hypothesis is supported by the fact that plasma outflow velocities in plumes are measured to be rather low up to $\\sim1~R_\\sun$ above the solar surface. The continuous emergence of magnetic flux at a slow rate and relatively large scale might ultimately create a sizable bundle of newly opened flux, allowing in turn a significant plume of escaping plasma to develop.  It is beyond us to simulate the development of a jet into a plume in an MHD model.  However, some basic physics of such a development can be anticipated.  If a bipolar field emerges into a unipolar, open field region, then the two fields are not, in general, exactly parallel across the boundary between them.  Then, according to Parker's (1994) theory, a magnetic tangential discontinuity forms and current dissipation and field reconnection become inevitable at this boundary. Any two non-parallel fields can be resolved into parallel and anti-parallel components. The anti-parallel components will mutually annihilate at the discontinuity. The dissipated magnetic energy is partially converted to kinetic and thermal energy, which would cause a jet of energized plasma to escape along the open field next to the dissipating current sheet. Whenever some quantity of open flux is locally annihilated along the current sheet an equal quantity of closed flux must become open for magnetic flux continuity (${\\bf\\nabla}\\cdot{\\bf B}=0$).  This open flux can allow a plume of thermally expanding plasma, formerly trapped by its closed field, to escape.  A jet model with a single magnetic neutral point such as Yokoyama \\& Shibata's (1996) anenome jet model (see their Figure 1) could also result in a plume. Energy gained from emerging flux is converted to kinetic and thermal energy at the X-type neutral point during reconnection producing a jet of energized plasma. When the field has reconnected, there is a bundle of newly-opened magnetic flux through which hitherto trapped coronal plasma can escape as a plume.  The present results would benefit from future, more extensive analysis of larger data samples recorded by different instruments in a simultaneous fashion over large time intervals."
    },
    "0806/0806.4453_arXiv.txt": {
        "abstract": "We have carried out an intensive survey of the northern region of the Fornax dwarf spheroidal galaxy with the aim of detecting the galaxy's short--period pulsating stars ($P<0.25$~days). Observations collected over three consecutive nights with the Wide Field Imager of the 2.2m MPI telescope at ESO allowed us to detect 85 high--amplitude (0.20--1.00 mag in $B$-light) variable stars with periods in the range from 0.046 to 0.126~days, similar to SX Phoenicis stars in Galactic metal-poor stellar populations. The plots of the observed periods vs. the $B$ and $V$ magnitudes  show a dispersion largely exceeding the observational errors.  To disentangle the matter, we separated the first-overtone from the fundamental-mode pulsators and tentatively identified a group of subluminous variables, about 0.35~mag fainter than the others. Their nature as either metal-poor intermediate-age stars or stars formed by the merging of close binary systems is discussed. The rich sample of the Fornax variables also led us to reconstruct the Period--Luminosity relation for short--period pulsating stars. An excellent linear fit, $M_V=-1.83(\\pm0.08)-3.65(\\pm0.07)\\times\\log P_F$, was obtained using 153 $\\delta$ Scuti and SX Phoenicis stars in a number of different stellar systems. ",
        "introduction": "In globular clusters, the systematic detection of the short--period ($P < 0.25$~days) pulsating stars located below the horizontal branch (namely, SX Phe and $\\delta$ Sct stars) started about twenty years ago (e.g., \\citealt{baas,ngc5053}). At that time, Nemec, Nemec \\& Lutz (1994) listed 14 SX Phe stars in NGC~5053, $\\omega$~Cen and NGC~5466. The use of these variable stars as distance indicators has made some progress in the last years (e.g., \\citealt{pych,mazur,olech}). However, due to their intrinsic faintness, not many of these pulsating stars have been discovered so far in stellar systems outside the Milky Way: \\cite{matcarina} and \\cite{carina} describe the first results obtained for SX Phe stars in the Carina dwarf spheroidal (dSph) galaxy; McNamara, Clementini \\& Marconi (2007) have used a sample of $\\delta$ Sct stars to estimate the distance to the Large Magellanic Cloud. To exploit in a more complete way the potential of these variable stars as distance indicators and stellar population tracers we have carried out a search for pulsating stars below the horizontal branch as starting point for an extensive project on the Fornax dSph. This galaxy is an ideal target since: (i) it is known to host a mix of old and intermediate--age stars with different metal abundances \\citep[see, e.g.,][and refs. therein]{battaglia06, coleman}; (ii) its size and distance make the observation of large fractions of the galaxy, down to a limiting  magnitude of about 3.0 mag below the horizontal branch, practical using a medium--class telescope equipped with a wide--field imager, using exposure times fully adequate to reveal the galaxy's short--period pulsators.  Our project also allowed a more comprehensive study of several classes of variable stars in the Fornax dSph galaxy, and it was extended over the course of the years to cover the galaxy's field and its system of five globular clusters, through the acquisitions of new observations with other telescopes (\\citealt{clementini06,for4}).  The nomenclature of the short--period pulsating stars below the horizontal branch is confusing. In the Milky Way there is a physical distinction between $\\delta$ Sct and SX Phe stars: the former are Population~I stars (Pop.~I), the latter are Population~II (Pop.~II) objects. Low--amplitude, nonradial modes are typical of the $\\delta$ Sct stars, but intensive and accurate surveys in globular clusters have provided observational evidences that they  are excited also in the SX Phe stars (\\citealt{pych,mazur,olech}).  Therefore, the amplitude can vary from a few 0.001~mag to several 0.1~mag both in the $\\delta$ Sct and in the SX Phe variables. The $\\delta$ Sct stars showing amplitude larger than 0.20~mag are generally referred to as High--Amplitude $\\delta$ Sct (HADS) stars. As in the case of high--amplitude SX Phe stars, the main pulsation period is a radial mode. The pulsation period can provide a rough separation between HADS and SX Phe stars, since periods of HADS vary from 0.07 to 0.25~days \\citep{ogle}, while most of the SX Phe stars in globular clusters have  $P<0.10$~days. However, since there is some overlap in the period range spanned by the two types of variables, disentangling HADS from SX Phe stars may not be an easy task if no details about the metallicity are known \\footnote{In the past, it was quite common to use the term ``Dwarf Cepheids\" to identify both HADS and SX Phe stars, since their light curves are reminding those of the Classical Cepheids. This term has an unclear meaning in an astrophysical context, since it groups  stars belonging to different populations, and has not been used here.}. In this respect, it is appropriate to use the SX Phe term to identify the short--period Pop.~II variables found in globular clusters. We will adopt this name also in the case of Fornax, since the range of metal abundance observed in Fornax is $-2.2<$[Fe/H]$<-0.7$ \\citep{saviane,pont04,battaglia06,gullieus07,coleman}, so that all the short--period pulsating stars in this galaxy are more similar to the SX Phe stars according to the nomenclature scheme described above. However, we note that in a different environment such as the Fornax dSph galaxy the criteria used in the Milky Way may not be adequate to describe the mixture of stars having such different ages and metallicities. ",
        "conclusions": "The intensive observation of a large part of the Fornax dSph galaxy down to 3.0~mag below the horizontal branch yielded us  a suitable tool to investigate the properties of the Fornax stellar populations. The peculiar characteristics of the sample of SX Phe variables, as the group of subluminous variables, provide useful clues on the stellar formation processes. The numerous sample of Fornax variables significantly increased the number of SX Phe variables known in different stellar systems, allowing us the possibility to obtain a comprehensive $P-L$ relation.  We plan to determine the physical parameters of a subset of  stars fitting their well defined light curves  by means of  theoretical pulsational models. Moreover, the  Fornax project  will be completed by future works on the other variables (namely, RR Lyr stars, anomalous and Pop.~II Cepheids, eclipsing binaries)."
    },
    "0806/0806.3273_arXiv.txt": {
        "abstract": "\\noindent The semi-classical nature of braneworld cosmological models does not account for any quantum gravitational effects. In this letter we use the gauge/gravity correspondence to argue that quantum string corrections {\\it cannot} be ignored in any study of braneworld stability. As an example, we find, by analysing the quantum gravitational backreaction, that a closed universe is unstable to radiation into the bulk. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0806/0806.3932_arXiv.txt": {
        "abstract": "We studied the complete randomness of the angular distribution of gamma-ray bursts (GRBs) detected by BATSE. Since GRBs seem to be a mixture of objects of different physical nature we divided the BATSE sample into 5 subsamples (short1, short2, intermediate, long1, long2) based on their durations and peak fluxes and studied the angular distributions separately. We used three methods, Voronoi tesselation, minimal spanning tree and multifractal spectra to search for non-randomness in the subsamples. To investigate the eventual non-randomness in the subsamples we defined 13 test-variables (9 from the Voronoi tesselation, 3 from the minimal spanning tree and one from the multifractal spectrum). Assuming that the point patterns obtained from the BATSE subsamples   are fully random we made Monte Carlo simulations taking into account the BATSE's sky-exposure function. The MC simulations enabled us to test the null hypothesis i.e. that the angular distributions are fully random. We tested the randomness by binomial test and introducing  squared Euclidean distances in the parameter space of the test-variables. We concluded that the short1, short2 groups deviate significantly (99.90\\%, 99.98\\%) from the fully randomness in the distribution of the squared Euclidean distances but it is not the case at the long samples. At the intermediate group the squared Euclidean distances  also give significant deviation (98.51\\%). ",
        "introduction": "Recently, there is no doubt about the cosmological origin of the gamma-ray bursts (hereafter GRBs) \\citep{zame04,fox05,me06}. Then, assuming a large scale isotropy for the Universe, one expects the same property for the GRBs as well. Another property, which is also expected to occur that GRBs should appear fully randomly, i.e.  if a burst is observed it does not give any information about the place of the next one. If both properties are fulfilled, then the distribution is called completely random (for the astronomical context of spatial point processes see \\citet{pato95}). There are several tests for checking the complete randomness of point patterns, however, these procedures do not always give information for both properties simultaneously.  There are increasing evidence that all the GRBs do not represent a physically homogeneous group \\citep{kou93,ho98,muk98,hak00,ho02,bal03,hak03,ho06}. Hence, it is worth investigating that the physically different subgroups are also different in their angular distributions. In the last years the authors provided \\citep{ba98,ba99,me00a,me00b}  several different tests probing the intrinsic isotropy in the angular sky-distribution of GRBs collected in BATSE Catalog \\citep{mee00}. Shortly summarizing the results of these studies one may conclude: A. The long subgroup ($T_{90} > 10\\; s$) seems to be distributed isotropically; B. The intermediate subgroup ($2\\;s \\lid T_{90} \\lid 10\\;s $) is distributed anisotropically on the $\\simeq (96-97)$\\% significance level; C. For the short subgroup ($2\\; s > T_{90}$) the assumption of isotropy is rejected only on the $92$\\% significance level; D. The long and the short subclasses, respectively, are distributed differently on the $99.3$\\% significance level.  (About the definition of subclasses see \\citet{ho98}; $T_{90}$ is the duration of a GRB, during which time the $90$\\% of the radiated energy is received \\citep{mee00}.)  Independently and by different tests, \\citet{li01}  confirmed the results A., B. and C. with one essential difference: for the intermediate subclass a much higher - namely $99.89$\\% - significance level of anisotropy is claimed. Again, the short subgroup is found to be \"suspicious\", but only the $\\simeq (85-95)$\\% significance level is reached. The long subclass seems to be distributed isotropically (but see \\citet{mest03}). \\citet{magl03} found significant angular correlation on the $2^{\\circ} - 5^{\\circ}$ scale for GRBs with $T_{90}< 2s$ durations. \\citet{tan05}  reported a correlation between the locations of previously observed short bursts and the positions of galaxies in the local Universe, indicating that between 10 and 25 per cent of short GRBs originate at low redshifts ($z < 0.025$).  It is a reasonable requirement to continue these tests using more sophisticated procedures in order to see whether the angular distribution of GRBs is completely random or has some sort of regularity. This is the subject of this article. New tests will be presented here. Mainly the clarification of the short subgroup's behaviour is expected from these tests. In this paper, similarly to the previous studies, the {\\it intrinsic\\/} randomness is tested; this means that the non-uniform sky-exposure function of BATSE instrument is eliminated.  The paper is organized as follows. In Section \\ref{mat} the three new tests are described. This Section does not contain new results, but - because the methods are not widely familiar - this minimal survey may be useful. Section \\ref{tests} contains the statistical tests on the data. Section  \\ref{disc} summarizes the results of the statistical tests,  and Section \\ref{conc} presents the main conclusions of the paper. ",
        "conclusions": "\\label{conc} We made additional studies on the degree of the randomness in the angular distribution of samples selected from the BATSE Catalog. According to the $T_{90}$ durations and $P_{256}$ peak fluxes of the GRBs in the Catalog we defined five groups: $short1$ ($T_{90}<2$ s \\& $0.65 < P_{256} < 2$), $short2$ ($T_{90}<2$ s \\& $0.65 < P_{256}$ ), $intermediate$ ($2\\;s \\lid T_{90} \\lid 10\\;s $ \\& $0.65 < P_{256}$), $long1$ ($T_{90}>2$ s \\& $0.65 < P_{256} < 2$) and $long2$ ($T_{90}>10$ s \\& $0.65 < P_{256}$).  To characterize the statistical properties of the point patterns, given by the samples, we defined 13 test-variables based on the Voronoi tesselation (VT), Minimal spanning tree (MST) and Multifractal spectra. For all five GRB samples defined we  made 200 numerical simulations assuming fully random angular distribution and taking into account the BATSE exposure function. The numerical simulations enabled us to define empirical probabilities  for testing the null hypothesis, i.e. the assumption that the angular distributions of the BATSE samples are fully random.  Since we performed 13 single tests simultaneously on each subsamples the significance obtained by calculating it separately for each test can not be treated as a true indication for deviating from the fully random case. At first we supposed that the test-variables were independent and making use the binomial distribution computed the probability of obtaining significant deviation  in at least one of the variables only by chance. In fact, some of the test-variables are strongly correlated. To concentrate the information on the non-randomness experienced by the test-variables, we assumed that they can be represented as a linear combination of non-correlated hidden factors of less in number. Actually, we estimated $k=8$ as the number of hidden factors. Making use the hidden factors we computed the distribution of the squared Euclidean distances from the mean of the simulated variables. {\\it Comparing  the distribution of the squared Euclidean distances of the simulated with the BATSE samples we concluded that the short1, short2 groups deviate significantly (99.90\\%, 99.98\\%)  from the fully randomness, but it is not the case at the long samples. At the intermediate group squared Euclidean distances also give significant deviation (98.51\\%).}"
    },
    "0806/0806.3103_arXiv.txt": {
        "abstract": "We discuss  the connection between the chemistry of dense interstellar clouds  and those characteristics of cometary matter that could be remnants of it. The chemical evolution observed to occur in molecular clouds is summarized and a model for dense core collapse that can plausibly account for  the isotopic fractionation of hydrogen, nitrogen, oxygen and carbon measured in primitive solar system materials is presented. ",
        "introduction": "\\label{sec:intro}  Comets have probably retained some material that originated in the molecular cloud from which the Sun formed\\cite{Pascale04}. Determining how much pristine interstellar material is in comets will help answer many important questions relating to the origin of our Solar System and, as comets are strong candidates for seeding planets with complex organic molecules, understanding the details of the interstellar-comet connection will have important implications for astrobiology\\cite{ARAA00}. Recent data from the {\\it Stardust} mission\\cite{Brownlee06}, and ground-based observations\\cite{Wooden02,Wooden07}, indicate that some materials present in cometary dust experienced very high temperatures ($\\sim 800$ K), relative to those typically found in molecular clouds ($\\sim 10$ K).  Nevertheless, the organic inventory and isotopic signatures measured for cometary molecules do provide a tantalizing connection with interstellar chemistry.  In this paper we discuss the chemical structure and evolution of interstellar matter prior to its incorporation into protoplanetary disks. We outline a model whereby many of the key cosmogonic markers in cometary matter can be explained as being of interstellar origin. ",
        "conclusions": "\\label{sec:conc}    We have reviewed the putative contribution of  interstellar chemistry to the volatile composition of comets.  In doing so, we have purposely neglected discussion of the possible contribution of nebular chemistry as a source of  cometary volatiles  - either from the warm innner nebula or from an essential continuation of interstellar chemistry in the cold outer disk; these have been reviewed elsewhere\\cite{AJM04,Ciesla06}.  Comets could contain material from  several different stages of molecular  cloud evolution and we have illustrated the related chemistry with some recent observations. Many of the molecules detected in comets, including several  organics, are either directly observed in interstellar ices (e.g. methanol, formic acid) or are believed to form on interstellar grains (e.g. formamide)\\cite{ARAA00}.  A direct comparison of the organic inventories is however complicated by the fact that, in comets, additional sources of simple molecules appear  to contribute to their coma abundances (e.g. CO, CS, CN, formaldehyde). These so-called {\\it extended sources} are believed to be due to the thermal or photolytic break-up of large organic macromolecules\\cite{Cottin07} and have no readily identifiable parallel in interstellar chemistry. Interstellar isotopic fractionation in dense gas with temperatures in the range $\\sim 10$--35~K  could account for the currently known, albeit meagre,  fractionation  ratios measured in comets.  We have presented the outline of a model that may account for all these characteristics as arising in the prestellar core from which the Sun formed.  As observed in many galactic sources, such cores can develop very strong chemical gradients. Gravitational collapse of such a core may deliver chemically distinct  regions (i.e.\\ mass shells)  onto the central disk during the Class 0/I epoch, and later, to the extent that all the `interstellar' characteristics of comets are delivered during this infall.  Finally,  future observations of more bright comets, especially short-period ones,  as well as  of protoplanetary disks, are necessary (e.g.\\ \\refcite{Lahuis 06}).   The advent of the {\\it Atacama Large Millimetre Array} promises a great advances in further investigating the ISM-comet connection\\cite{vanD07,WilsonALMA}."
    },
    "0806/0806.1923_arXiv.txt": {
        "abstract": "{In most global fits of the constrained minimal supersymmetric model (CMSSM) to indirect data, the {\\it a priori}\\/ likelihoods of   any two points in $\\tan \\beta$ are treated as equal, and   the more fundamental $\\mu$ and $B$ Higgs potential parameters are fixed by potential minimization conditions. We find that, if instead a flat (``natural'') prior measure on $\\mu$ and $B$ is placed, a strong preference exists for the focus point region from fits to particle physics and cosmological data. In particular, we find that the lightest neutralino is strongly favored to be a mixed bino-higgsino ($\\sim 10\\%$ higgsino). Such mixed neutralinos have large elastic scattering cross sections with nuclei, leading to extremely promising prospects for both underground direct detection experiments and neutrino telescopes. In particular, the majority of the posterior probability distribution falls within parameter space within an order of magnitude of current direct detection constraints. Furthermore, neutralino annihilations in the sun are predicted to generate thousands of neutrino induced muon events per years at IceCube. 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\\put(0,5.8){(a)} \\put(2,5.8){(b)} \\put(4,5.8){(c)} \\put(0,3.8){(d)} \\put(2,3.8){(e)} \\put(4,3.8){(f)} \\put(1,1.8){(g)} \\put(3,1.8){(h)} \\end{picture} }  \\preprint{DAMTP-2008-51\\\\ FERMILAB-PUB-08-153-A}  \\begin{document} ",
        "introduction": "For a variety of reasons, supersymmetry is considered to be among the most attractive extensions of the Standard Model. In particular, weak-scale supersymmetry provides an elegant solution to the hierarchy problem~\\cite{susyreview}, and enables grand unification by causing the gauge couplings of the Standard Model to evolve to a common scale~\\cite{gut}. From the standpoint of providing a dark matter candidate, the lightest neutralino is naturally stable by virtue of R-parity conservation \\cite{neutralinodm}, and in many models is thermally produced in the early universe in a quantity similar to the measured density of cold dark matter~\\cite{wmap5}.  In addition to collider searches for superpartners, a wide range of astrophysical experiments are currently operating and being developed in the hopes of detecting neutralino dark matter~\\cite{Bertone:2004pz}. These techniques can be classified as direct and indirect detection. While the former efforts are designed to observe the elastic scattering of neutralinos with target nuclei, the latter techniques attempt to detect the annihilation products of neutralinos, including gamma-rays~\\cite{gamma}, neutrinos~\\cite{neutrinos}, positrons~\\cite{positron}, antiprotons~\\cite{antiproton}, antideuterons~\\cite{antideu}, and synchrotron radiation~\\cite{syn}. In addition to astrophysical inputs, the prospects for direct and indirect dark matter detection depend on the mass and couplings of the lightest neutralino, and in turn on the many parameters which define the masses and couplings of the superpartners.  Weak-scale supersymmetry could take a great variety of forms, depending on the details of how supersymmetry is broken. Empirically, our insights into this question are limited to the measurements of observables indirectly related to the supersymmetric spectrum, such as the anomalous magnetic moment of the muon, the $b \\rightarrow s \\gamma$ branching fraction, the $B_s \\rightarrow \\mu^+  \\mu^-$ branching fraction, the mass of the $W$ boson, the effective leptonic mixing angle, Higgs boson and sparticle search constraints, and the cosmological dark matter abundance. Such observables have been used in the past to constrain the properties of the CMSSM spectrum (see, for example, Refs.~\\cite{ellis,Profumo:2004at,Allanach:2005kz,Roszkowski:2006mi,weather}). Ultimately, this information can be used to determine the posterior probability distribution over the parameter space of supersymmetry. In Refs.~\\cite{deAustri:2006pe,Roszkowski:2007fd,Roszkowski:2007va}, it was used to examine the prospects for dark matter detection.  In this paper, we consider another input that can play a significant role in determining the posterior distribution over supersymmetric parameters. In particular, we consider the measure which is associated with each point in parameter space and define a prior measure which is flat in terms of fundamental CMSSM parameters. In our analysis, we closely follow Ref.~\\cite{weather}, but focus on the phenomenology of neutralino dark matter in the regions of supersymmetric parameter space favored by indirect constraints {\\em and}\\/ naturalness considerations. When a natural prior measure (flat in more fundamental CMSSM parameters, rather than in $\\tan \\beta$) is included in the analysis of the parameter space of the constrained minimal supersymmetric standard model (CMSSM), we find that the focus point region is highly preferred. In this region, the lightest neutralino $\\chi_1^0$ is a mixed bino-higgsino ($\\sim 10\\%$ higgsino fraction) and, therefore, has relatively significant couplings to the Standard Model.  The prospects for the direct and indirect detection of neutralino dark matter in the favored regions are highly promising. In particular, about \\fpPercentage\\% of the posterior probability distribution predicts a neutralino-nucleon elastic scattering cross section of $\\sigma_{\\chi^0 N} \\approx 10^{-8}-10^{-7}$ pb, which is within one order of magnitude of the current direct detection constraints. The remaining \\hpolePercentage\\% of the posterior probability distribution corresponds to parameter space in which the lightest neutralino has somewhat smaller couplings (and direct detection rates) but still annihilates efficiently in the early universe via the light Higgs resonance ($2 m_{\\chi^0} \\approx m_h$). The projected rates at neutrino telescopes are also extremely promising, with most of the posterior probability distribution being made up of models which predict thousands of events per year at a kilometer-scale neutrino telescope such as IceCube. Current constraints from Super-Kamiokande and Amanda/IceCube already exclude a sizable fraction of the otherwise favored probability distribution. We also discuss the prospects for indirect detection using gamma-rays and charged cosmic ray particles. ",
        "conclusions": "}  By considering measurements of quantities such as the anomalous magnetic moment of the muon, the $b \\rightarrow s \\gamma$ branching fraction, the $B_s \\rightarrow \\mu^+  \\mu^-$ branching fraction, the mass of the $W$ boson, the effective leptonic mixing angle, Higgs boson and sparticle search constraints, and the cosmological dark matter abundance, it is possible to constrain the parameter space of supersymmetry. The results of global fits to such indirect data currently depend, however, on the choice of priors which are adopted. In most of the global fits of supersymmetric parameter space which have been performed to date, priors have been used which are flat in the derived quantity, $\\tan \\beta$. A far more natural choice would be to use priors which are flat (or perhaps, logarithmic) in the fundamental parameters $\\mu$ and $B$. In this article, we have considered the impact of adopting such natural priors upon fits of the CMSSM to indirect data, focusing on the phenomenology of neutralino dark matter that is found in the parameter space favored by such fits.  Using natural priors and updated indirect data, we find a that two regions of the CMSSM parameter space are strongly favored. Firstly, about \\fpPercentage\\% of the posterior probability distribution corresponds to the focus point region. Of the remainder, \\hpolePercentage\\% of the posterior probability distribution corresponds to the light Higgs-pole region in which the lightest neutralino annihilates on resonance with the light Higgs boson and \\stauCoanPercentage\\% corresponds to the stau co-annihilation region, where staus and other sleptons efficiently annihilate with the lightest neutralinos. In contrast to the results found using priors flat in $\\tan \\beta$, we find that the stau-co-annihilation and $A$-funnel regions of the CMSSM parameter space contribute negligibly to the posterior probability distribution.  In the favored focus point region, the lightest neutralino is a mixed gaugino-higgsino ($\\sim$10\\% higgsino fraction) with a mass less than approximately 300 GeV. Such a neutralino has a very distinctive dark matter phenomenology and is nearly optimally suited for the purposes of direct and indirect detection. In particular, a mixed gaugino-higgsino neutralino possesses large couplings to Standard Model fermions, and thus has large elastic scattering cross sections with nuclei. In the light Higgs-pole region, the lightest neutralino can have considerably smaller couplings.  We find that neutralinos in the favored focus point region have a spin-independent elastic scattering cross section with nucleons of $\\sim 3 \\times 10^{-8}$ pb, which is within a factor of 2 (5) of the current limit from CDMS for a 100 GeV (300 GeV) neutralino. We, therefore, expect direct detection experiments to probe the majority of the posterior probability distribution of the CMSSM parameter space in the very near future.  The prospects for neutrino telescopes found using natural priors are also very promising. In particular, most of the favored parameter space predicts thousands of events to be observed per year in a kilometer-scale neutrino telescope such as IceCube. Current constraints from Super-Kamiokande already exclude 38\\% of the posterior probability distribution, assuming a local dark matter density of 0.3 GeV/cm$^3$.  Although searches for dark matter using gamma-rays and charged particles depend strongly on unknown astrophysical inputs, our analysis finds that the majority of the favored parameter space predicts a neutralino annihilation cross section near the maximum possible for a thermal relic ($\\sim 3 \\times 10^{-26}$ cm$^3$/s). This along with the relatively light mass range favored for the lightest neutralino makes the prospects for GLAST and PAMELA to detect neutralino dark matter near optimal.  We believe that a prior that is flat in $\\mu, B$ is a much more natural choice than one flat in $\\tan \\beta$.If the naturalness prior were complemented with an additional hyper-parameter prior that enforces that all soft terms are ``of the same order''~\\cite{weather}, the focus point is {\\em dis}favoured. However one can consider differences in the derived posterior probability distributions from the different priors as evidence that more data is needed to constrain the model. Thus, fit predictions that are robust (i.e. approximately invariant) with respect to changes in assumed prior distributions are not attained for mSUGRA, since it has many parameters and the data constraining it are rather indirect. While this is undeniably true, it is still interesting to examine the effect of the more natural prior as it gives us our ``best bet'' for quantities such as the dark matter-nucleon cross-sections relevant for direct detection, or galactic annihilation cross sections relevant for indirect detection. Our neutralino-nucleon cross-sections coming from the fit for the flat $\\tan \\beta$ prior are similar to other previous determinations in the literature, providing validation of our calculations. Our best guess for this quantity leads to a good chance that a further increase of a factor of 10 in sensitivity by the experiments will lead to a direct detection discovery. Had, instead of priors flat in $\\mu, B$, we had chosen priors that are flat in $\\log \\mu$ and $\\log B$, we expect that our fits would show even stronger preference for the focus point: an additional factor of $1/(\\mu B)$ in the integrand of Eq.~\\ref{initial2} would have lead to even more preference for the focus point region, with an associated extra boost in detection cross-sections.   Taken together, the results presented in this article are very encouraging for the prospects for direct and indirect efforts to detect neutralino dark matter. If natural choices are made in constructing priors, fits to the currently available data predict that, within the context of the CMSSM, the lightest neutralino is likely to have large elastic scattering and annihilation cross sections. In particular, the majority of the posterior probability distribution of the CMSSM parameter space (about \\fpPercentage\\%) should be within the reach of very near future direct detection experiments, and should be detectable in the near future by IceCube.  Voltaire's satirical philosopher Pangloss long held the position that we live in the ``best of all possible worlds''. We find that if naturalness considerations are taken into account, then (modulo the usual astrophysical uncertainties) the prospects for the direct and indirect detection of neutralino dark matter in the CMSSM are, if not Panglossian, are at least extremely encouraging."
    },
    "0806/0806.1112_arXiv.txt": {
        "abstract": "Radiative and Auger decay data have been calculated for modelling the K lines in ions of the nickel isonuclear sequence, from Ni$^+$ up to Ni$^{27+}$. Level energies, transition wavelengths, radiative transition probabilities, and radiative and Auger widths have been determined using Cowan's Hartree--Fock with Relativistic corrections (HFR) method. Auger widths for the third-row ions (Ni$^+$--Ni$^{10+}$) have been computed using single-configuration average (SCA) compact formulae. Results are compared with data sets computed with the AUTOSTRUCTURE and MCDF atomic structure codes and with available experimental and theoretical values, mainly in highly ionized ions and in the solid state. ",
        "introduction": "The nickel K lines have arguable diagnostic potential in X-ray astronomy. They reside in a relatively unconfused part of the X-ray spectrum where they can be used to estimate quantities of interest such as the redshift, temperature, abundance, and the velocity of the emitting gas. Sensitive observations from the currently active {\\em Suzaku} satellite, and those to be expected from future space missions, e.g. the {\\em Constellation-X} observatory, are bringing about further spectroscopic attention to the Ni K features.  On the other hand, clear detections of nickel have been obtained only in a few astrophysical X-ray spectra. This is undoubtedly due to a combination of limited counting statistics at energies above 7 keV, where the nickel lines occur, and to poor spectral resolution in this band which makes nickel $K\\alpha$ and iron $K\\beta$ lines indistinguishable. The first detection of Ni K$\\alpha$ in an X-ray binary was reported by \\cite{Sak99} in the {\\em ASCA} spectrum of Vela X-1. Fluorescence line energies and intensities were measured mostly for neutral elements, including nickel in eight ionization stages in spite of the uncertainties in the line energy ($7.52^{+0.06}_{-0.05}$ keV). This detection was later confirmed by \\cite{Gol04} using three data sets obtained with {\\em Chandra/HETGS} during a binary orbit, representing the orbital eclipse phases of $\\phi=0$, $\\phi=0.25$, and $\\phi=0.5$. They found relative strong Ni K$\\alpha$ emission with an increase in the line flux of a factor of 15 between the $\\phi=0$ and $\\phi=0.5$ phases.  X-ray absorption features from Ni have been found in X-ray binary systems. \\cite{Boi04} presented evidence of K$\\alpha$ absorption in Ni~{\\sc xxvii} at 7.82 keV in the {\\em XMM-Newton} spectrum of the low-mass X-ray binary (LMXB) XB 1916-053. In the case of the LMXB GX 13+1, \\cite{Sid02} came across two absorption features at 7.85 and 8.26 keV which they respectively identified as either K$\\beta$ lines from Fe~{\\sc xxv} and Fe~{\\sc xxvi} or K$\\alpha$ from Ni~{\\sc xxvii} and Ni~{\\sc xxviii}.  Ni K$\\alpha$ emission in an active galactic nucleus (AGN) was first detected in the {\\em XMM-Newton} X-ray spectra of the Circinus galaxy \\citep{Mol03}. Fe K$\\alpha$, Fe K$\\beta$, and Ni K$\\alpha$ were identified at 6.4, 7.058, and 7.472 keV, respectively. From the Ni and Fe K$\\alpha$ line fluxes, a nickel-to-iron abundance ratio was estimated at 0.055--0.075, a factor of 1.5--2 larger than the cosmic values \\citep{And89}. Similar studies of the Seyfert~2 galaxy NGC 1068 were reported by \\cite{Mat04} using {\\em Chandra} and {\\em XMM-Newton} observations. Fe and Ni fluorescence emission in both neutral and highly ionized material was identified and, from the nickel-to-iron flux ratio, Fe was found to be overabundant by a factor of 2 with respect to solar and Ni by the same factor with respect to iron. \\cite{Pou06} confirmed the nickel emission in NGC 1068 as well as in Mkn 3, a Seyfert~2 galaxy previously studied by \\cite{Pou05}. \\cite{Pou06} also showed evidence of high-velocity shifts in the lines coming from ionized nickel. \\cite{Mar07} detected with {\\em Suzaku} narrow K$\\alpha$ fluorescence emission lines from Fe, Si, S, Ar, Ca, and Ni in the radio-loud AGN Centaurus A (NGC 5128).   \\cite{Mol98} analyzed data collected with {\\em BeppoSAX} from the Perseus cluster of galaxies, to find that the ratio of the flux of the 7--8 keV line complex to the 6.7 keV line was significantly larger than that predicted by optically thin plasma codes, and that it diminishes with increasing cluster radius. It is argued that this effect is due to resonance scattering in an optically thick plasma at the energies of the Fe K$\\alpha$ line. However, \\cite{Gas04} used {\\em XMM-Newton} observations of the same cluster to measure an overabundance for nickel. These authors claimed that the excess in the flux of the 7--8 keV line complex was due to Ni K$\\alpha$ emission rather than resonance scattering. \\cite{deP04} also detected an overabundant nickel in the {\\em XMM-Newton} spectrum of the cluster Abel 478, but they noticed that this could be overestimated due to errors in the nickel line energies. The determination of the Ni abundance profile in the ICM is also important because Ni is almost exclusively produced in SNIa; therefore, high resolution spectral analysis, together with accurate atomic data, are needed to understand these observations.  The first object detected by {\\em INTEGRAL} in the Galactic plane, IGR J16318-4848, is a possible X-ray binary where the obscuring matter has a column density as large as the inverse Thompson cross section. This source also shows strong Fe K$\\alpha$ and Fe K$\\beta$ emission accompanied by a weaker, but nevertheless distinct, Ni K$\\alpha$ feature at 7.5 keV \\citep{Wal03}. Observations of the same source with {\\em XMM-Newton} reveal similar properties \\citep{Mat03}, and a subsequent monitoring campaign with both {\\em INTEGRAL} and {\\em XMM-Newton} confirmed column densities of $1.2 \\times 10^{24}$ cm$^{-2}$ and large equivalent widths for the K$\\alpha$ emission lines of Fe and Ni \\citep{Iba07}.  Perhaps the first unambiguous detection of Ni K lines in an astrophysical spectrum is due to \\cite{koy07}. They have observed the diffuse X-ray emission from the Galactic center using the X-ray Imaging Spectrometer (XIS) on {\\em Suzaku}. They have detected, for the first time, lowly ionized nickel and He-like nickel K lines, and have measured the K$\\alpha$ line flux ratio of Ni~{\\sc xxvii} to that of Fe~{\\sc xxv} to determine a plasma temperature of $\\sim$5.4 keV, assuming solar relative Ni and Fe abundances.  Following work by \\cite{pal02,pal03a,pal03b}, \\cite{bau03,bau04}, \\cite{men04}, and \\cite{kal04} on the Fe K lines, by \\cite{gar05} on the K-shell photoabsorption of O ions, and the recent study of the K lines in the Ne, Mg, Si, Ar, and Ca isonuclear sequences \\citep{pal08}, we report new atomic data for K-vacancy levels in the nickel isonuclear sequence. Prime objectives are to improve the atomic database of the XSTAR modelling code \\citep{bau01} and to prepare ionic targets (configuration expansions and orbitals) for the lengthy computations of the K-shell photoabsorption and photoionization cross sections, where both radiative and Auger dampings are key effects. With respect to Ni, available atomic structure data sets---namely K-vacancy level energies, wavelengths, $A$-values, and radiative and Auger widths---for first-row ions with electron number $2 \\le N \\le 9$ are far from complete while for the second and third-row ions ($10 \\le N \\le 27$) they are certainly lacking. The only exception is Ni$^+$ where measurements in the solid state have been reported \\citep{sal70, sal72, sli72, ber78, rao86, per87, hol97, raj01, egr08} and the K$\\alpha $ unresolved transition array (UTA) centroid wavelength have been calculated by \\cite{hou69}.  Previous study on the K-shell structure of nickel includes that by \\cite{hsu87} on the satellite spectra of He-like nickel. They have recorded spectra emitted from the plasma of the Tokamak Fusion Test Reactor (TFTR) with a high-resolution crystal spectrometer, providing a K-line list for Ni~{\\sc xxvii} ($N=2$) and  Ni~{\\sc xxvi} ($N=3$) and interpreting the observed spectra with the aid of a Hartree-Fock-Slater (HFS) calculation. \\cite{bom88} have measured the wavelengths of the K lines in Ni~{\\sc xxiv}--Ni~{\\sc xxvii} ($N=$ 2--5) emitted by a hot plasma at the Joint European Torus (JET). In order to analyze these observations, they have computed wavelengths, $A$-values, and radiative widths with the SUPERSTRUCTURE atomic structure code \\citep{eis74} and Auger rates with the AUTOLSJ \\citep{tfr81}. \\cite{vai78} have calculated wavelengths, radiative transition probabilities, and autoionization (Auger) rates for ions with atomic numbers $Z=$ 4--34 using the $1/Z$ expansion technique. They have considered the ${\\rm 1s-2p}$ transitions in the H-like sequence, ${\\rm 1s2l-2p2l}$ and ${\\rm 1s^2-1s2l}$ in the He-like sequence, and ${\\rm 1s^22l-1s2p2l}$ in the Li-like sequence. Energies for K-vacancy levels of the type ${\\rm 1s}nl$ in the He-like isoelectronic sequence have been calculated by \\cite{vai85} using the same technique. \\cite{gor03} have audited the fluorescence database by \\cite{kaa93}, which is widely used in modelling codes, in particular their scaling procedures along isoelectronic sequences. They have found serious flaws which appear to compromise the application of this database in plasma modelling.  The outline of the present report is as follows. The numerical methods are briefly decribed in Section~2 while an analysis of the results based on comparisons with previous experimental and theoretical values is carried out in Section~3. The two supplementary electronic tables are explained in Section~4 while some conclusions are finally discussed in Section~5. ",
        "conclusions": "Extensive data sets containing energy levels, wavelengths, radiative transition probabilities, absorption oscillator strengths, radiative and Auger widths, and fluorescence yields for K-vacancy levels in the nickel isonuclear sequence have been computed with the atomic structure codes HFR, AUTOSTRUCTURE, and GRASP. For the ionic species Ni$^{2+}$-Ni$^{22+}$, with electron numbers $6 \\le N \\le 26$, this is the first time that such data become available. For ions with $2 \\le N \\le 5$ and $N=27$, detailed comparisons have been carried out with available measurements and theoretical values which bring forth the consistency and accuracy of the present data.  Comparisons of HFR K-level energies with those reported in the NIST database \\citep{nist} for He- and Li-like nickel and the calculated values in Ni$^{26+}$ by \\cite{vai85} support an accuracy rating for ions with electron number $N \\le 3$ of better than 3~eV. With regards to wavelengths, comparisons between HFR and the tokamak data \\citep{hsu87, bom88} for ions with $2\\leq N\\leq 5$, with the solid-state measurements \\citep{hol97} for $N=27$, and with other published theoretical values \\citep{vai78, vai85, hsu87, bom88} result in a general agreement within 1~m\\AA.  Based on the level of agreement of the present HFR $A$-values with those by \\citet{vai78} and \\citet{bom88} for the highly ionized species ($2\\leq N\\leq 5$), we are confident that they show an accuracy within 10\\% for transitions with $A\\gtrsim 10^{13}$ s$^{-1}$. Regarding radiative widths, similar considerations have led us to the conclusion that our HFR widths are also accurate to 10\\%. For the $K\\alpha_2/K\\alpha_1$ and $K\\beta/K\\alpha$ line ratios in the second and third-row ions ($12\\leq N\\leq 27$), the trends along the isonuclear sequence are consistent (within 10\\%) for the three independent approaches we have used. For $N=27$, our HFR $K\\alpha_2/K\\alpha_1$ ratio agrees within 5\\% with the solid-state measurements \\citep{sal70, hol97}, but the HFR $K\\beta/K\\alpha$ ratio is 15\\% lower than the bulk of the available experimental values located at 0.139$\\pm$0.015.  Comparisons of the HFR Auger widths with those by \\citet{bom88} for species with $3\\leq N\\leq 5$ and by \\cite{vai78} for $N=3$ show a fair consistency for widths greater than 10$^{13}$ s$^{-1}$, although the latter is higher by $\\sim$15\\% and the former by 10\\% for $N=4$ and 15\\% for $N=5$. The systematic deviations with the data by \\cite{bom88} may be due to the orbital optimization procedure they used. Moreover, recent measurements of $KLM$ Auger channel relative intensities \\citep{egr08} support our HFR Auger rates. Finally, HFR fluorescence yield and natural K-level width for $N=27$ are consistent within 5\\% with the recommended values found in the literature \\citep{hub94,cam01}.  The present radiative and Auger widths will be used in the computations of the K-shell photoionization cross sections of these ions which are required in XSTAR \\citep{kal01} for the modelling of interesting Ni spectral features."
    },
    "0806/0806.4392_arXiv.txt": {
        "abstract": "We present time--resolved V--photometry of the old nova RR\\,Pic. Apart from the hump--like variability, the light curves show the strong flickering and random variation typical for RR\\,Pic. We do not find any convincing evidence for the previously reported eclipse. The extrapolated eclipse phase coincides with a broad minimum, but comparing the overall shape of the light curve suggests that the eclipse should actually be located around phase 0.2. The orbital period which we derive from these data agrees well with the old one, any uncertainty is too small to account for the possible phase shift. Apart from the 3.48\\,h period, which is usually interpreted as the orbital one, we find an additional period at $P=3.78$\\,h, which we interpret as the superhump period of the system; the corresponding precession period at 1.79\\,d is also present in the data. We also find indications for the presence of a 13\\,min quasi--periodic oscillation. ",
        "introduction": "Classical novae are a subclass of cataclysmic variables (CVs), close interacting binary stars with a white dwarf primary receiving matter from a Roche--lobe--filling late--type star. They are distinguished by the observation of a thermonuclear runaway outburst, the nova explosion. As such, RR Pic was discovered by \\citet{spen31} at maximum light in 1925 and, although it was a slow nova, it is supposed to be in its quiescence state by now. The orbital period of 0.145025 days \\citep{vogt75} places it just above the period gap and into the regime of the SW\\,Sex type stars. Indeed, RR\\,Pic has been found to show several observational features typical for SW\\,Sex stars \\citep{schm+03} and can thus be regarded as a nova--like CV with very high mass transfer rate.  Vogt found the lightcurve dominated by a very broad hump, often interrupted by superimposed minima. He explained this behaviour by an extended hot spot region with an inhomogeneous structure. \\citet{haef+82} however, explained their own time--resolved photometric and polarimetric observations together with radial velocity variations of the He\\,II (4686\\,\\AA) emission line \\citep{wyck+77} by suggesting the presence of an additional source of radiation in the disc opposite of the hot spot. The presence of such an emission source was confirmed via Doppler tomography by \\citet{schm+03} and \\citet{ribe+06}. From high speed photometry, \\citet{warn86} concluded that during the 1970s (about 50\\,yr after outburst) structural changes have taken place in the system, resulting in a more isotropic  distribution of the emitted radiation. In addition, he has found evidence for a shallow, irregular eclipse, showing RR\\,Pic to be a high inclination system. Note that no signature of an eclipse had been found in the previous lightcurves. In addition to the orbital period, \\citet{kubi84} found a periodic modulation in the optical with a 15\\,min period. He interpreted this as the rotation of the white dwarf and concluded that RR\\,Pic is an intermediate polar. \\citet{haef+85} however, repeated the high time--resolved photometry on a longer time--scale and could not find any sign of this short period. Since no 15\\,min period variation is found in X--ray measurements \\citep{beck+81} either, they concluded that Kubiak's variation was more likely a transient event in the disc rather than a rotating white dwarf. Also Warner's high-speed photometry does not reveal any period other than the orbital one.  \\citet{schm+05} compared radial velocity curves of different epochs and noticed a shift of about 0.1 phases of the radial velocity curves of data taken two years apart. They argue that this might be due to unstable emission sources in the accretion disc or might indicate a change in the orbital period. To test these alternatives, we performed new time--resolved photometry of RR\\,Pic with the aim to determine the orbital period and look for a possible change that could account for the observed phase shift. These data and the results are presented in this paper. ",
        "conclusions": "We have presented optical lightcurves of RR\\,Pic and shown that they are dominated by a strong orbital variation. The orbital period derived from this data is consistent with the previous reported ones. I.e. a change of this period can not be responsible for the previously observed shift in the phases of radial velocity curves. Instead, this shift is rather due to structural changes that are known to occur in the accretion disc of RR\\,Pic.  In addition to the orbital variation, a superhump is found that was used to derive the mass ratio $q = 0.31$. This value does not agree with a previously reported lower one, further observations are needed for clarification.  QPOs of 13\\,min are present in all our data taken between February and April 2005. Older data from 2004, do not show this oscillation. While our analysis thus confirms the variations reported earlier, it also shows that these QPOs are a transient phenomenon. Their presence might be connected with the accretion disc's structure if they occur due to an illumination of blobs in the inner accretion disc from a spinning white dwarf.  From our data we can not confirm the presence of an eclipse. Instead, we note that at least in one historical case the eclipse occurs when the minima of orbital light curve and superhump lightcurve fall together. This might indicate that the observed eclipse is a resonance phenomenon between the two lightcurves and its existence or not-existence is modulated with the precession period. We would like to clarify that we do not rule out the presence of a shallow eclipse but insist that either its visibility is enhanced by the resonance or a favourable structure of the accretion disc is needed for an eclipse to be observed.  In general we conclude that RR\\,Pic is a highly variable system. The previously reported changes that happen in the accretion disc are probably responsible for the various features in the lightcurve that are not present at all times. To really understand what is going on in this system, parallel time--resolved spectroscopy and photometry would be needed over several cycles."
    },
    "0806/0806.3267_arXiv.txt": {
        "abstract": "Recent work suggests that Type Ia supernovae (SNe) are composed of two distinct populations: prompt and delayed.  By explicitly incorporating properties of host galaxies, it may be possible to target and eliminate systematic differences between these two putative populations. However, any resulting {\\em post}-calibration shift in luminosity between the components will cause a redshift-dependent systematic shift in the Hubble diagram.  Utilizing an existing sample of 192 SNe Ia, we find that the average luminosity difference between prompt and delayed SNe is constrained to be $(4.5 \\pm 8.9)\\%$.   If the absolute difference between the two populations is 0.025 mag, and this is ignored when fitting for cosmological parameters, then the dark energy equation of state (EOS) determined from a sample of 2300 SNe Ia is biased at $\\sim1\\sigma$. By incorporating the possibility of a two-population systematic, this bias can be eliminated. However, assuming no prior on the strength of the two-population effect, the uncertainty in the best-fit EOS is increased by a factor of 2.5, when compared to the equivalent sample with no underlying two-population systematic.  To avoid introducing a bias in the EOS parameters, or significantly degrading the measurement accuracy, it is necessary to control the post-calibration luminosity difference between prompt and delayed SN populations to better than 0.025 mag. ",
        "introduction": "\\label{sec:introduction}  The discovery of the accelerating expansion of the universe~\\citep{Riess:98, Perl:99} has led to an explosion of interest in the underlying physics responsible for this acceleration.  A favored model characterizes the acceleration by an unknown energy density component, dubbed the dark energy. While there exist a variety of probes to explore the nature of this dark energy, one of the most compelling entails the use of type Ia supernovae (SNe henceforth) to map the expansion history of the universe. Several present and future SN surveys are aimed at constraining the dark energy equation-of-state (EOS) to better than 10\\%. With increasing sample sizes, SN distances can potentially provide multiple independent estimates of the EOS of dark energy when binned in redshift \\citep{Huterer:05,sullivan:07,sarkar:08a}.  Given the importance of dark energy measurements, it is then useful to quantify various systematics that impact SN cosmology \\citep{huigreene:06,cooraycaldwell:06,cooray:06,sarkar:08b}.  Recently, suggestions have been made that the SN population consists of two components, with a ``prompt'' component proportional to the instantaneous host galaxy star formation rate, and a ``delayed'' (or ``extended'') component that is delayed by several Gyrs \\citep{hamuy:95, livio:00, scannapieco:05,mannucci:06, sullivan:06, strovink:07}.  The former is expected to be more luminous, and thus, prompt SN lightcurves are broader than those of the delayed population. By classifying SNe by host galaxy type, \\citet{howell:07} found a {\\em pre}-calibration intrinsic luminosity difference of $\\sim (12 \\pm 4)$\\% between the two components, based on a difference of $(8.1 \\pm 2.7)\\%$ in the width of lightcurves. Since the SN lightcurves are used to calibrate the intrinsic luminosity \\citep{phillips:93, riess:96, perlmutter:97, tonry:03, prieto:06, guy:07, jha:07}, a systematic difference in intrinsic luminosity could conceivably be calibrated out, if the SN lightcurve calibration relation is the same for both populations.  However, it is unclear whether the full intrinsic difference in luminosity between the two populations is captured by a calibratable difference in the lightcurves.  A residual in the calibrated luminosity could potentially remain, leading to a redshift-dependent shift in the Hubble diagram, and systematic errors in the best-fit cosmological parameters. For example, it is likely that the intrinsic colors of Type Ia SNe are not uniquely determined by light-curve shape~\\citep{conley08}. Even if this is not the case, differences in intrinsic color between the populations might introduce systematic differences in post-calibration luminosity (e.g., through differences in the extinction corrections). We model the two-population systematic, constraining the magnitude of the effect with current data. With large SN samples it may be possible to estimate the magnitude of the systematic directly from the data (e.g., by correlating observed SN brightnesses with properties of the host galaxies \\citep{hamuy:96b,Riess:99,sullivan:06,jha:07, gallagher:08}). We quantify the level of calibration required to avoid significantly degrading the determination of the dark energy EOS.  The paper is organized as follows: in $\\S$2 we discuss a model for incorporating a two-population systematic residual luminosity into the Hubble diagram. In $\\S$3 we investigate the possibility of detecting this systematic from both current and future SN data, and the impact on dark energy parameter estimation.  \\begin{figure*}[!t] \\centerline{ \\includegraphics[scale=0.5]{f1a.eps} \\includegraphics[scale=0.5]{f1b.eps} } \\caption{{\\it Left panel:} Correlation between $\\delta_D$ and $w$, corresponding to a $w$CDM model, for the 192 SN dataset discussed in the text. There is a strong degeneracy between the two-population systematic, $\\delta_D$, and the dark energy EOS, $w$. This leads to a potential bias in the measurement of the EOS, and increased errors when the fits are marginalized over $\\delta_D$. {\\it Right panel:} Histograms showing the distributions of the best-fit $w$, from 200 mock data sets with an inherent $\\delta_D=0.025$, after marginalizing over the WMAP $w$CDM 5-Year priors. The shaded histogram represents the case where $\\delta_D$ is allowed to vary freely, while for the hatched histogram the two-population systematic is ignored ($\\delta_D=0$). The dot-dashed line corresponds to the case where the data is fit assuming a $\\delta_D=0.025\\pm0.025$ prior. For clarity we omit the underlying histogram.} \\label{fig:1} \\end{figure*} ",
        "conclusions": ""
    },
    "0806/0806.3098_arXiv.txt": {
        "abstract": "\\noindent With applications in astroparticle physics in mind, we generalize a method for the solution of the nonlinear, space homogeneous Boltzmann equation with isotropic distribution function to arbitrary matrix elements. The method is based on the expansion of the matrix element in terms of two cosines of the ``scattering angles''. The scattering functions used by previous authors in particle physics for matrix elements in Fermi-approximation are retrieved as lowest order results in this expansion. The method is designed for the unified treatment of reactive mixtures of particles obeying different scattering laws, including the quantum statistical terms for blocking or stimulated emission, in possibly large networks of Boltzmann equations. Although our notation is the relativistic one, as it is used in astroparticle physics, the results can also be applied in the classical case. ",
        "introduction": "Non-equilibrium processes in astroparticle physics, such as Big Bang nucleosynthesis (BBN), neutrino decoupling and more speculative ones, like baryogenesis through leptogenesis or the freeze-out of hypothetical relic particles, \\cite{Hannestad:1995rs,Gnedin:1998,Dodelson:1992,Kolb:1990,Fukugita:1986,Buchmuller:2004,Serpico:2004gx} are usually computed through the solution of the corresponding coupled system of Boltzmann equations \\cite{Bernstein:1988,Groot:1980,Mischler:2003,Cercignani:2002,Liboff:2003} describing the time evolution of the one-particle distribution functions $f^i(t,k)$. Typically, in cosmology, it is anticipated that the relevant particle distributions are in exact kinetic equilibrium and of Maxwell-Boltzmann type. These assumptions, together with others, allow the Boltzmann equation to be linearized and integrated, which leads to coupled sets of chemical rate equations (mostly themselves dubbed Boltzmann equations in this context). This procedure drastically simplifies the numerical computation of the particle abundances, such that even the approximate solution of very large networks of Boltzmann equations, as in the case of BBN, becomes possible. However, in doing so, one looses the spectral information contained in the definition of the distribution functions and other fundamental properties of the Boltzmann equation are neglected as well. It is well-known that the solution of the full Boltzmann equations can lead to relevant corrections to the equilibrium results in several cases \\cite{Dolgov:1997mb,Hannestad:1999fj,Basboll:2006yx}. In the era of precision cosmology the inclusion of such non-equilibrium effects gains in importance. Regarding the use of classical kinetic theory for the description of phenomena in the (very) early universe there are concerns, originating in the belief that these calculations should be performed in the framework of non-equilibrium quantum field theory. These concerns are supported by recent results, revealing differences between the two approaches for simple toy models, at least in extreme non-equilibrium situations, see e.g. \\cite{Lindner:2005,Berges:2002wr}. However, it seems natural to attempt to include quantum effects in modified effective kinetic equations. Boltzmann equations will continue to play an important role in cosmology at least at the relatively low energies of neutrino decoupling or nucleosynthesis, where the standard calculations give already quite good results.  In general, a network of Boltzmann equations can be written as: \\bigformula{L[f^i]=\\sum_{l} C^{il}[f^1,\\ldots ,f^i ,\\ldots ,f^N]\\kend}{system_of_boltzmann_equation} where there is one equation for each of the $N$ participating particle species ($i=1\\ldots N$) and one collision term $C^{il}$ for each interaction with particles of the same and of other species. $L$ denotes the Liouville operator, most commonly in Minkowski space-time, \\bigformula{L[f^i](k)=\\partd{f^i(t,k)}{t}\\kend}{rel_liouville_uncurved} or in Robertson-Walker space-time, \\bigformula{L[f^i](k)=\\partd{f^i(t,k)}{t}-\\hubblerate k\\partd{f^i(t,k)}{k}\\kend}{rel_liouville} with Hubble rate $\\hubblerate ={\\dot{a}}/{a}$. By writing the collision integrals as $C^{il}[f^1\\ldots f^i\\ldots f^N]$, we have formally taken the possibility of multi-particle scattering processes into account. Usually only decays, inverse decays and $2-2$ scattering processes,  $a+b\\leftrightarrow A+B$, are considered. For the latter ones the collision integral reads \\bigformularray{&&\\hspace{-6mm}C^{al}[f^af^bf^Af^B](k)=\\nonumber\\\\ &&\\qquad\\frac{1}{2E^a_k}\\int\\int\\int\\frac{\\dif^3{p}}{(2\\pi)^32E^b_p} \\frac{\\dif^3{q}}{(2\\pi)^32E^A_q} \\frac{\\dif^3{r}}{(2\\pi)^32E^B_r}  (2\\pi)^4\\delta^4(k+p-q-r)\\meavsqu\\nonumber\\\\ &&\\qquad\\qquad\\qquad\\times \\left[(1-\\xi^a f_k^a)(1-\\xi^b f_p^b)f_q^Af_r^B-f_k^af_p^b(1-\\xi^A f_q^A)(1-\\xi^B f_r^B)\\right]\\kend\\nonumber\\\\}{collision_integral} where we have used the short-hand notation $f^i_k=f^i(t,k)$ and $\\xi^i$ to specify the quantum statistics of particle species $i$, i.e. $\\xi^i=+1$ for Fermi-Dirac, $\\xi^i=-1$ for Bose-Einstein and $\\xi^i=0$ for Maxwell-Boltzmann statistics. $\\meavsqu$ denotes the invariant matrix element squared and averaged over initial and final spin states. Note that we take $\\meavsqu$ to include possible symmetrization factors of $1/2$ for identical particles in the initial or final state. There is a vast number of different methods for the solution of the Boltzmann equation (mostly applied in different fields of physics), out of which only a few exploit the homogeneity and isotropy as imposed by the cosmological principle. So-called direct integration methods, where the collision term \\eqnref{collision_integral} is integrated numerically, seem to be most advantageous because they are characterised by high precision This is desirable since one wants to keep track of only small deviations from equilibrium. The direct numerical solution using deterministic methods is numerically expensive, mainly because of the multiple integrals in the collision terms. Therefore, it relies on the successful reduction of the collision integral for isotropic distribution functions.  In the present paper a technique for this reduction of the collision integral is presented which generalizes previous results in high energy and astro physics \\cite{Dolgov:1997mb,Yueh:1976} for matrix elements in Fermi-approximation to (in principle) arbitrary matrix elements, relying on a series expansion of the matrix element. The resulting reduced Boltzmann equation contains only a two-fold integral over the magnitudes of the post-collisional momenta. The method is applicable to Boltzmann equations with and without quantum statistical terms and can be used independent of the dispersion relation, i.e. it can be used for massive and massless relativistic particles as well as for non-relativistic ones. The loss and gain terms can be treated collectively or independently. Thus the method represents an approach to treat reactive mixtures of all kinds of particles with different interactions, in a unified manner.  The outline is as follows: In \\sect\\secref{reduction_of_the_collision_integral} we show how the nine-dimensional collision integral for $2-2$ scattering processes can be reduced to a two-dimensional one, integrating out the energy and momentum conserving $\\delta$-functions. (Collision integrals for decays and inverse decays can be integrated in the same way.) In doing so, a certain angular integral over the matrix element arises. In \\sect\\secref{a_simple_numerical_model} we establish a simple numerical model for the reduced Boltzmann equation. The integral of \\sect\\secref{reduction_of_the_collision_integral} is solved by expanding the matrix element in terms of the cosines of two ``scattering angles'', see \\sect\\secref{expansion_of_the_scattering_kernel}. In \\sect\\secref{numerical_integration} we derive a formula, suitable for numerical integration of the full matrix element, which we employ in the last section to demonstrate the convergence of the series torwards the exact result for a simple example. We conclude in \\sect\\secref{conclusion}. ",
        "conclusions": "}  In state-of-the-art computations in astroparticle physics usually all species apart from the neutrinos, which experience only weak interactions, are assumed to be in exact kinetic equilibrium and heterogeneous (at best) networks of Boltzmann- and rate equations are solved instead of the full system of kinetic equations. The number of species involved in realistic systems is large and requires a unified treatment of the different particles and interactions.  A method for the solution of the space homogeneous Boltzmann equation (with isotropic distribution function), for general scattering laws was presented here. The method relies on the expansion of the matrix element in terms of cosines of two ``scattering angles''. For the separate terms in this expansion the full angular integration was carried out. The functions $\\Dnmind{0}{0}, \\Dnmind{0}{1}$ and $\\Dnmind{0}{2}$, corresponding to matrix elements in Fermi-Approximation, were used in previous literature to compute non-equilibrium corrections to the neutrino-distribution functions and have been obtained in the lowest order of the expansion.  Though our starting point was the relativistic form of the Boltzmann equation, as encountered in astroparticle physics, the method can be used for the non-relativistic equation as well. In any case, it allows for the full angular integration of the scattering kernel, reducing the collision integral from effective dimension $5$ to dimension $2$. The only prerequisite is that the matrix element can be expanded into a series of the scattering angles and that this series converges rapidly enough. The quantum statistical terms for blocking and stimulated emission can be carried along.  In the introduction we mentioned that, at high densities and temperatures, modifications of the Boltzmann equation might become necessary (if these are sufficient at all). One such modification, which has been suggested on various occasions, is the inclusion of higher order scattering processes. Since the representation of the angular integral in terms of spherical Bessel functions (\\ref{eqn:Dnm_function_1}-\\ref{eqn:Dnm_function_2}) and the integral \\eqnref{intfour_result}, by means of \\eqn\\eqnref{int_four_bessel_to_int_three_bessel}, can be generalized for higher order processes (in which case more than four Bessel functions appear in the integrals) the method, described above, can in principle be used to reduce the corresponding collision integrals from dimension $3\\,(n-1)$ to $n-2$. ($n$ is the number of particles involved)  The functions $\\Dnm$, though of very simple structure, can become lengthy for higher orders. Moreover, due to the presence of very different relaxation time-scales the system of ODE's, corresponding to the numerical method  presented in \\sect\\secref{a_simple_numerical_model}, tends to behave stiff. Therefore, in possible implementations, careful optimisation for efficiency and stability is necessary.  Let us add, that the expansion of the scattering kernel in terms of the cosines of the angles is not a new idea. For example in \\cite{Kuegerl:1989} an expansion of the scattering kernel has been combined with a moment method for the non-relativistic, inhomogeneous Boltzmann equation. The expansion of generic kernels with full integration of the angular part, in the space-homogeneous and isotropic case, seems to be new, however.  \\subsection*{Acknowledgements} \\noindent Andreas Hohenegger was supported by the ``Sonderforschungsbereich'' TR27.  \\newpage  \\begin{appendix}  \\renewcommand{\\theequation}{A.\\arabic{equation}}  \\setcounter{equation}{0}"
    },
    "0806/0806.1865_arXiv.txt": {
        "abstract": "The MAGIC collaboration has recently reported the discovery of $\\gamma$-ray emission from the binary system \\lsi\\ in the TeV energy region. Here we present new observational results on this source in the energy range between $300\\textrm{ GeV}$ and $3\\textrm{ TeV}$. In total 112 hours of data were taken between September and December 2006 covering 4 orbital cycles of this object. This large amount of data allowed us to produce an integral flux light curve covering for the first time all orbital phases of \\lsi. In addition, we also obtained a differential energy spectrum for two orbital phase bins covering the phase range $0.5 < \\phi<0.6$ and $0.6 < \\phi<0.7$. The photon index in the two phase bins is consistent within the errors with an average index $\\Gamma=2.6\\pm0.2_{stat}\\pm0.2_{sys}$. \\lsi\\ was found to be variable at TeV energies on timescales of days. These new MAGIC measurements allowed us to search  for intra-night variability of the VHE emission; however, no evidence for flux variability on timescales down to 30 minutes was found. To test for possible periodic structures in the light curve, we apply the formalism developed by Lomb and Scargle to the \\lsi\\ data taken in 2005 and 2006. We found the \\lsi\\ data set to be periodic with a period of (26.8$\\pm$0.2)~days (with a post-trial chance probability of 10$^{-7}$), close to the orbital period. ",
        "introduction": "The $\\gamma$-ray binary system \\lsi\\ is located at a distance of $\\sim$2 kpc and is composed of a compact object of unknown nature (neutron star or black hole) orbiting a Be star in a highly eccentric orbit ($e=0.72\\pm0.15$ or $e=0.55\\pm0.05$ following~\\cite{Casares:2005wn} and \\cite{Grundstrom:2006} respectively).  \\lsi\\ was found to display periodic variability in the radio, infrared, optical, and X-ray bands (\\citealt{1982ApJ...255..210T}, \\citealt{Paredes:1995}, \\citealt{Mendelson:1989} and \\citealt{Paredes:1997}, respectively).  The orbital period of the system is 26.4960 days long~\\citep{Gregory:2002}. The periastron passage, derived from the optical spectra, is found to be at phase $\\phi=0.23\\pm0.02$ in~\\cite{Casares:2005wn} and $\\phi=0.301\\pm0.011$ in~\\cite{Grundstrom:2006}, adopting a zero-phase at $T_0=\\textrm{JD } 2443366.775$.  Radio outbursts are observed every orbital cycle at phases varying between 0.45 and 0.95 with a 4.6 years modulation~\\citep{Gregory:2002}. Radio imaging techniques have shown extended, radio-emitting structures with angular extensions of $\\sim$0.01 to $\\sim$0.1 arc-sec, where the radio emission originates in a two-sided, possibly precessing, relativistic jet ($\\beta/c=0.6$)~\\citep{Massi:2004}. These extended radio structures have led some authors to adopt the microquasar scenario to explain the non-thermal emission in \\lsi\\ ~\\citep[e.g.,][]{Bosch-Ramon:2006, Bednarek:2006}. Recent high resolution VLBA measurements show a complex and changing morphology different from what is expected for a typical microquasar jet (see radio images in~\\citealt[e.g.,][]{2006smqw.confE..52D,magiclsi.2006MW}). Furthermore no solid evidence for the presence of an accretion disk (i.e. a thermal X-ray component) has been observed~\\citep{2006MNRAS.372.1585C}. This seems to favor a scenario in which the non-thermal emission in \\lsi\\ is powered by the interaction between a pulsar and the primary star winds~\\citep{Maraschi:1981}.  At higher energies, \\lsi\\ was found to be spatially coincident with the EGRET $\\gamma$-ray source \\eg~\\citep{kniffen}. Variable emission at TeV energies was observed with the MAGIC telescope~\\citep{Albert:2006vk} and was recently confirmed by VERITAS~\\citep{Veritas_lsi}. The system showed the peak TeV $\\gamma$-ray flux at phase $\\phi\\sim0.65$, while no very high-energy emission was detected around the periastron passage.  Here we present new MAGIC telescope observations of \\lsi. We briefly discuss the observational technique and the data analysis procedure, investigate the very high energy (VHE) $\\gamma$-ray spectrum during the high emission phase of the source, and put the results into perspective with previous VHE $\\gamma$-ray observations of this system. Finally we analyzed the temporal characteristics of the TeV emission and find a periodic modulation of the signal with the orbital period. ",
        "conclusions": "We find that \\lsi\\ is a periodic $\\gamma$-ray binary with an orbital period of 26.8$\\pm$0.2~days (and chance probability $\\sim 10^{-7}$), compatible with the optical, radio and X-ray period. This result implies that the flux modulation is tied to the orbital period. The high state in VHE $\\gamma$-rays occurs in the same phases as the X-ray high state. This is especially interesting since we found a additional hint for X-ray/$\\gamma$-ray variability correlation in the orbital phase 0.85. A strictly simultaneous multi-wavelength campaign is needed to investigate this correlation in more detail.  We looked for possible intranight variability and found the flux consistent with being constant within errors in 30-75 minutes time-scales.  We produce energy spectra for two phase bins $0.5 < \\phi < 0.6$ and $0.6 < \\phi < 0.7$ and averaged flux values for several phase bins. There is clear evidence for a significant change in the VHE $\\gamma$-ray flux level between different phase bins of \\lsi. The spectral photon index does not show this dependence on the phase. All derived spectral photon indices are compatible with $2.6\\pm0.2$ , obtained from the most significant phase bin of \\lsi.  We can put constraints to the emission at the periastron passage and conclude that the system is detected in $\\gamma$-rays only in the phases $0.5-0.9$. Since significant emission is only detected in an orbital sector off the phases at which the maximum gamma ray flux should occur under photon-photon absorption (see fig. 5 in~\\citealt{2006A&A...451....9D}), the latter can hardly be the only source of variability in the emission.  Thorough multiwavelength observations will allow us to probe the intrinsic variability of the non-thermal emission from \\lsi\\ along the orbit and can proof possible correlations between the X-ray and TeV energy bands. This is a necessary step for understanding the source nature, and the physics underlying the VHE radiation."
    },
    "0806/0806.2866_arXiv.txt": {
        "abstract": "The $\\OVI$ ion observed in quasar absorption line spectra is the most accessible tracer of the cosmic metal distribution in the low redshift ($z<0.5$) intergalactic medium (IGM).  We explore the nature and origin of $\\OVI$ absorbers using cosmological hydrodynamic simulations including galactic outflows with a range of strengths.  We consider the effects of ionization background variations, non-equilibrium ionization and cooling, uniform metallicity, and small-scale (sub-resolution) turbulence.  Our main results are (1) IGM $\\OVI$ is predominantly photo-ionized with $T\\approx 10^{4.2\\pm0.2}$ K.  A key reason for this is that $\\OVI$ absorbers preferentially trace over-enriched (by $\\sim\\times 5$) regions of the IGM at a given density, which enhances metal-line cooling such that absorbers can cool to photo-ionized temperatures within a Hubble time.  As such, $\\OVI$ is not a good tracer of the Warm-Hot Intergalactic Medium.  (2) The predicted $\\OVI$ properties fit observables if and only if sub-resolution turbulence is added, regardless of any other model variations.  The required turbulence increases with $\\OVI$ absorber strength.  Stronger absorbers arise from more recent outflows, so qualitatively this can be understood if IGM turbulence dissipates on the order of a Hubble time.  The amount of turbulence is consistent with other examples of turbulence observed in the IGM and galactic halos.  (3) Metals traced by $\\OVI$ and $\\HI$ do not trace exactly the same baryons, but reside in the same large-scale structure.  Our simulations reproduce observed alignment statistics between $\\OVI$ and $\\HI$, yet aligned absorbers typically have $\\OVI$ arising from cooler gas, and for stronger absorbers lower densities, than $\\HI$.  Owing to peculiar velocities dominating the line structure, coincident absorption often arises from spatially distinct gas.  (4) Photo-ionized $\\OVI$ traces gas in a variety of environments, and is not directly associated with the nearest galaxy, though is typically nearest to $\\sim 0.1L_*$ galaxies.  Weaker $\\OVI$ components trace some of the oldest cosmic metals.  (5) Very strong absorbers ($EW\\ga 100$m\\AA) are more likely to be collisionally ionized, tracing more recent enrichment ($\\la 2$~Gyr) within or near galactic halos. ",
        "introduction": "The exploration of metals in the low redshift ($z<0.5$) intergalactic medium (IGM) has taken a large leap forward with the advent of high-resolution space-based ultraviolet (UV) spectroscopy with the {\\it Hubble Space Telescope Imaging Spectrograph} (STIS) and {\\it Far Ultraviolet Spectroscopic Explore} (FUSE).  The strongest and most common metal transition seen in quasar absorption line spectra is the $\\OVI$ doublet at 1032,1038\\AA.  Many recent papers \\citep[e.g.][]{tri00, sav02, pro04, ric04, sem04a, leh06, coo08} have examined $\\OVI$ absorbers along single sight lines, and have fit ionization models to individual systems with associated $\\lya$ and lower ionization metal transitions.  Three recent studies (Tripp et al. 2008, hereafter T08; Danforth \\& Shull 2008, hereafter DS08; and Thom \\& Chen 2008a, 2008b) have compiled samples of $\\OVI$, providing the largest statistical surveys to date.  Understanding $\\OVI$ in the low-$z$ Universe is extremely important, because it may hold the key to locating a significant reservoir of cosmic baryons and metals \\citep[e.g.][]{tri00b}.  The inventory of observed baryons \\citep[e.g.][]{fuk98} falls well short of the predicted cosmological values from the cosmic microwave background \\citep{hin08}, leading to the well-known ``missing baryons problem.'' Cosmological simulations \\citep{cen99,dav99,dav01a} suggest that a large fraction of baryons today ($>30\\%$) reside in a hard-to-observe warm-hot intergalactic medium (WHIM), with temperatures of $T=10^{5}-10^{7}$~K.  The WHIM results primarily from shocks created during large-scale structure formation.  More recent simulations including feedback \\citep[hereafter OD08]{cen06a,opp08} find that galactic winds may increase the fraction of cosmic baryons in the WHIM to 50\\% or more.  A possibly related problem is that only $\\frac{1}{3}$ of metals have been accounted for observationally~\\citep{fuk04}, judging from the mismatch between the amount of metals nucleosynthesized and ejected by observed stars and those observed in various cosmic baryonic components.  $\\OVI$ potentially represents the Holy Grail of all things missing in the low-$z$ Universe, because its collisionally ionized equilibrium (CIE) maximum temperature, $10^{5.45}$ K, provides a unique and relatively easily accessible tracer of WHIM metals and baryons.  The incidence of $\\OVI$ with broad $\\lya$ absorbers (BLAs) thought to trace gas at $T>10^{5}$ K \\citep[e.g.][]{ric04} supports this notion, while $\\OVI$ seen with very broad $\\HI$ having line widths $>100\\;\\kms$ may indicate even hotter $\\OVI$ at $T\\sim10^6$ K \\citep[e.g.][]{dan06}.  Early investigations into $\\OVI$ using cosmological simulations \\citep{cen01,fan01,che03} predicted that stronger $\\OVI$ absorbers tend to be collisionally ionized while weaker ones tend to be photo-ionized, with the cross-over equivalent width of $\\sim 30-50$~m\\AA.  \\citet{cen06b} added non-equilibrium ionization and used higher resolution simulation, finding similar behavior compared to their earlier work.  The above simulations generally focused on fitting only the observed cumulative equivalent width ($EW$) distribution.  Recent observations now provide new challenges for simulations to fit a wider range of low-$z$ $\\OVI$ observables.  Recent surveys have renewed confusion about the nature of $\\OVI$ absorbers.  Ionization models for absorbers showing aligned $\\OVI$, $\\HI$, and $\\CIII$ are usually forced to invoke multi-phase gas, with both $\\OVI$ in CIE and lower ionization species at photo-ionized temperatures \\citep{pro04, dan06, coo08}.  Yet T08 find that the majority of aligned $\\HI-\\OVI$ absorbers can also be explained by photo-ionization alone.  The ``blind'' searches of $\\OVI$ without $\\HI$ by \\citet{tho08a} indicate $\\ga 95\\%$ of $\\OVI$ is associated with $\\HI$.  These studies suggest that $\\OVI$ traces baryons at least partially accounted for by $\\HI$ absorbers, and is not necessarily a good tracer of the WHIM.  Studies of galaxy-absorber correlations find a wide variety of environments for $\\OVI$ absorbers including voids, filaments, galaxies, and groups \\citep[e.g.][]{tri01,tum05b,pro06,tri06}.  Furthermore, it is not clear how the $\\OVI$ in intermediate- and high-velocity clouds (IVCs and HVCs) associated with the Milky Way (MW) halo \\citep[e.g.][]{sem03,sav03,fox05} relate to $\\OVI$ observed in quasar absorption line spectra.  A unified model was proposed by \\citet{hec02} where $\\OVI$ absorbers are all radiative cooling flows passing through the coronal temperature regime, with a modification by \\citet{fur05} for long cooling times in the IGM.  However, observations are inconclusive as to whether this scenario also applies to IGM $\\OVI$ \\citep[e.g][]{dan06,leh06}.  Hence there remains much uncertainty in regards to which cosmic gas and metal phases $\\OVI$ absorbers actually trace.  In this paper we explore $\\OVI$ and $\\HI$ absorbers in a range of physical models using our modified version of the cosmological hydrodynamics code \\gad.  We pay close attention to three key observables confirmed in multiple studies: the cumulative $EW$ distribution, $\\OVI$ linewidths ($b$-parameters) as a function of $\\OVI$ column density, and the alignment of $\\OVI$ with $\\HI$.  Our two main goals are (1) to see how self-consistent metal enrichment via galactic superwinds reproduces the $\\OVI$ observations, including an honest evaluation of the short-comings of our simulations, and (2) to understand the physical conditions and environments of $\\OVI$ absorbers.  This study combines state-of-the-art modeling with the most recent $\\OVI$ data to further understand the nature of IGM $\\OVI$, in light of the much anticipated installation of the {\\it Cosmic Origins Spectrograph} (COS) and the re-activation of STIS on {\\it Hubble}.  The structure of the paper flows as follows.  In \\S2, we introduce our simulations all run to $z=0$ with three different galactic outflow models, plus other post-run input physics variations.  Then, in \\S3 we see how the various models fit the $\\lya$ forest and our three $\\OVI$ observables, including a test of resolution convergence. The crux of this paper lies in the following sections where we dissect our simulated absorber population, acknowledging the imperfections in our modeling while assessing what we believe is physically significant.  \\S4 advocates that IGM $\\OVI$ absorbers are primarily photo-ionized, with an in-depth analysis of the association with $\\HI$, and forwards an explanation for the non-thermal component of $\\OVI$ line profiles as arising from small-scale turbulence.  \\S5 examines the origin of $\\OVI$ absorbers with an emphasis on environment.  We discuss the minority population of collisionally ionized $\\OVI$ here, and examine tell-tale signs of such WHIM absorbers with an eye towards COS.  We summarize in \\S6.  Throughout this paper, we adopt \\citet{asp05} solar abundances, as most low-$z$ $\\OVI$ observations use this or a similar value.  This contrasts to our use of \\citet{and89} values in previous publications (Oppenheimer \\& Dav\\'e 2006, hereafter OD06; Dav\\'e \\& Oppenheimer 2007; OD08), and results in a significant decrease in solar oxygen mass fraction (0.00541 versus 0.00962).  The Type II supernovae (SNe) yields we use as simulation inputs are from \\citet{lim05}, and are independent of solar values.  For ease of discussion, we classify $\\OVI$ absorbers into three categories according to their column density: weak ($N(\\OVI)<10^{13.5} \\cms$), intermediate ($N(\\OVI)=10^{13.5-14.5} \\cms$), and strong ($N(\\OVI)=10^{14.5-15.0} \\cms$).  We often split the intermediate absorbers into two bins, since this comprises the majority of observed $\\OVI$. ",
        "conclusions": "We examine the nature and origin of $\\OVI$ absorbers in the low redshift ($z<0.5$) IGM using \\gad~cosmological hydrodynamic simulations including outflows, exploring a variety of wind models and input physics variations. We determine the best-fitting model by comparing to a suite of observed $\\OVI$ and $\\HI$ absorber statistics, including the cumulative equivalent width distribution, $b(\\OVI)$ as a function of $N(\\OVI)$, and the $N(\\OVI)$ as a function of $N(\\HI)$.  Our first main result is that {\\it only a model where we explicitly add sub-resolution turbulence is able to match all observations.}  The observations motivating this conclusion are the progressively larger line widths for stronger $\\OVI$ systems, along with the high incidence of large equivalent width absorbers.  We discuss this further later on.  The second main result is that {\\it the vast majority of $\\OVI$ absorbers are photo-ionized, with only a few strong systems being collisionally ionized.}  This result does not depend on turbulence or any other input physics variations.  The governing variable in photo-ionized $\\OVI$ strength is density, which steadily increases from overdensities of $\\sim10$ for the weakest observed absorbers tracing metals in extended filaments, to $\\sim200$ for the strongest observed absorbers tracing metals in or around galactic halos.  Increasing metallicity with density also helps $\\OVI$ column densities become stronger.  Our third main result is that {\\it metals in the IGM are distributed inhomogeneously, and that $\\OVI$ absorbers preferentially arise from over-enriched regions.}  The average absorber traces metals 5$\\times$ above the mean metallicity-density relationship in our simulation ($0.15-1.0 \\Zsolar$).  Only 1.3\\% of the IGM volume enriched to greater than $0.1 \\Zsolar$, with a filling factor of only 0.3\\% resulting from weaker winds working nearly as well.  We also test a uniform distribution of IGM metals, finding that this creates too many weak absorbers.  Our results imply that applying the mean IGM metallicity-density gradient to gas everywhere in order to study metal absorption is not appropriate; the {\\it scatter} in the global metallicity-density relationship critically influences $\\OVI$ absorbers.  The clumpy metal distribution is an important reason why most $\\OVI$ is at photo-ionization temperatures, because it strongly enhances metal-line cooling.  Oxygen is a powerful coolant and is (not coincidentally) particularly efficient in the narrow temperature regime where $\\OVI$ has a collisional ionization maximum.  This creates a ``zone of avoidance\" for enriched IGM gas between $10^{4.5}-10^{5.5}$ K.  Previous theoretical studies that did not account for metal line cooling incorrectly predict more collisionally ionized systems.  Non-equilibrium ionization allows collisionally ionized $\\OVI$ to extend to lower temperatures \\citep{gna07}, but the differences are only important for relatively rare high density regions, and do little to alter the overall $\\OVI$ statistics.  The high fraction of aligned absorbers in the \\citet{tho08a} and \\citet{tri08} datasets supports most $\\OVI$ absorbers being photo-ionized.  While observed complex, multi-phase systems with mis-aligned components indicating CIE $\\OVI$ may be under-represented in our simulations, it is unlikely that would dominate typical $\\OVI$ systems under any circumstance.  A key implication of our findings is that {\\it $\\OVI$ is generally a poor tracer of the Warm-Hot Intergalactic Medium (WHIM), and cannot be used to infer the WHIM baryonic content.}  In depth study of the properties of $\\OVI$ absorbers and their associated $\\HI$ reveals that multi-phase temperatures and densities are often needed to explain their properties.  Metal-line cooling allows inhomogeneously distributed $\\OVI$ to cool slightly below $\\HI$-weighted temperatures, while stronger $\\OVI$ absorbers often trace lower overdensities than aligned strong $\\HI$, which typically traces gas inside halos.  Such a multi-density photo-ionized model appears necessary to explain strong $\\OVI$ absorbers aligned with $N(\\HI)>10^{15} \\cms$ when considering absorbers on the $N(\\HI)-N(\\OVI)$ plane.  The \\citet{tri08} dataset finds more weak $\\HI$ absorbers aligned with $\\OVI$ than our simulations, possibly indicating local fluctuations in the meta-galactic photo-ionizing background around 1~Rydberg.  We find good agreement with observations for the distribution of velocity separations between $\\OVI$ and $\\HI$ absorbers.  The typical length of a photo-ionized $\\OVI$ absorber is $50-100$~kpc, although the peculiar velocity field often dominates resulting in mis-aligned $\\HI$ components, especially those tracing higher densities.  Only 42\\% of our strong absorbers $N(\\OVI)\\ge 10^{14} \\cms$ are aligned within $4\\kms$ while 91\\% are aligned within $80\\kms$.  This indicates a clumpy distribution of $\\OVI$ that does not exactly trace the smoothly varying gas in the $\\lya$ forest, but still arises in the same underlying large-scale structure.  A clear anti-correlation exists between the density traced by $\\OVI$ and its age as determined by the last time the metals were launched in a wind.  This trend arises in our simulations because of the outside-in pattern of metal enrichment; outer, lower overdensities are enriched first when the Universe is young and physical distances are small, and inner, higher overdensities are progressively more enriched by later winds extending a smaller comoving distance. This is a consequence of our finding in \\citet{opp08} that winds travel similar physical distances relatively independent of redshift or galaxy mass. While a majority of metals falls back into galaxies, our work here shows that low-$z$ $\\OVI$ absorbers below $N(\\OVI)\\sim 10^{14} \\cms$ provide one of the best ways to observe the oldest metals that remain in the IGM.  We suggest the fascinating possibility of deriving nucleosynthetic yields tracing the earliest stars if another high ionization species, such as $\\CIV$, can be observed at low densities with the increased throughput of COS or future facilities.  Finally, we study the galaxy environment around $\\OVI$ absorbers. Photo-ionized $\\OVI$ usually has little to do with either the mass of the neighboring galaxy, averaging $\\sim 0.03-0.1 M^*$, or the distance, averaging $\\ga 2r_{vir}$, because of the significant galactic and large-scale structural evolution occurring while these metals reside in the IGM.  This helps explain why observed photo-ionized absorbers appear to trace a variety of environments, with a typical distance of 100-300 kpc to the nearest galaxy.  The opposite is true of the small minority of collisionally ionized absorbers.  These strong absorbers are $<2r_{vir}$ from $M^*$ galaxies and have ages indicating activity within the last 2~Gyr, and are perhaps analogous to $\\OVI$ observed in HVCs and IVCs in the Milky Way halo.  The relatively young age for strong $\\OVI$ signifies that these metals recycle back into galaxies multiple times in what can be described as halo fountains~\\citep{opp08}.  The most uncertain and controversial part of this paper is the claim that significant sub-resolution turbulence is present in $\\OVI$ absorbers.  We determine the amount of turbulence needed by directly fitting the observed $N(\\OVI)-b(\\OVI)$ relation.  This increases $b$-parameters to observed levels by construction, while simultaneously increasing large-$EW$ absorbers to the observed frequency.  We attempt to justify this ``magic bullet\", as well as compare the implied turbulence to other instances of observed IGM turbulence.  We point out that various authors have argued for clumpiness in IGM metals extending well below the mass and spatial resolution of our simulations~\\citep{sim06,sch07}.  If such clumpiness exists, then surely those clumps must have some relative motion, so assuming completely static velocities below the SPH smoothing scale of $20-100$~kpc is unrealistic.  Our model states that strong absorbers can be dominated by turbulent motions, just as the line profiles associated with molecular clouds, $\\HII$ regions, IVCs, and HVCs associated with our Galaxy are.  The energy dissipation required using a Kolmogorov spectrum is a fraction of the turbulence observed by \\citet{rau01} in high-$z$ $\\CIV$ absorbers, which is encouraging considering the IGM at late times should be calmer, although it is unclear whether Kolmogorov theory accurately describes IGM turbulence. We show that this trend is quantitatively consistent with the idea that turbulence dissipates as metals reside longer in the IGM.  Our turbulent scenario makes the case that metal-line absorbers are made up of numerous cloudlets with lower ionization species tracing high-density clouds making multiple thin profiles, and $\\OVI$ tracing more of the low-density regions in between creating a single broader profile; such is often observed to be the structure of low-$z$ metal line systems.  In short, we believe that small-scale turbulence is a viable possibility, although more detailed modeling is required to fully understand its implications.  Applying a volume renormalization zoom technique to cosmological simulations, allowing the resampling of individual haloes in galactic-scale simulations, provides a realistic, near-term possibility for following the evolution of turbulence over a Hubble time.  While it may be disappointing that $\\OVI$ is a poor WHIM tracer and provides only a weak handle on the IGM metal distribution, our results open up some new and interesting possibilities for using $\\OVI$ absorbers to understand cosmic metals and feedback.  For instance, low-$z$ $\\OVI$ may provide a fascinating opportunity to study some of the oldest IGM metals in the Universe through weak absorbers.  The strongest absorbers appear related to the recycling of gas between the IGM and galaxies, providing a unique window into how galaxies get their gas \\citep[e.g.][]{ker05, ker08}.  Detecting multiple ions in weaker absorbers will provide a good handle on physical conditions owing to the photo-ionized nature of $\\OVI$.  Our work here is only a first step towards understanding the nature and origin of $\\OVI$ in the low-$z$ Universe using numerical simulations, which we hope to extend further by considering other metal species in a greater evolutionary context.  We stress the need to model in detail the complex metal-line systems that will undoubtedly be observed by COS. The computational models presented, while state of the art, still fall well short of what is necessary, particularly for stronger collisionally ionized systems.  We anticipate that future observational and modeling improvements will shed new light on the metal distribution at the present cosmic epoch with all its important implications.  This work provides a first step in that direction."
    },
    "0806/0806.0863_arXiv.txt": {
        "abstract": "{ Stellar collisions are % an important formation channel for blue straggler stars in globular and old open clusters. Hydrodynamical simulations have shown that the remnants of such collisions are out of thermal equilibrium, are not strongly mixed and can rotate very rapidly. Detailed evolution models of collision products are needed to interpret observed blue straggler populations and to use them to probe the dynamical history of a star cluster. We expand on previous studies by presenting an efficient procedure to import the results of detailed collision simulations into a fully implicit stellar evolution code. Our code is able to evolve stellar collision products in a fairly robust manner and allows for a systematic study of their evolution.  Using our code we have constructed detailed models of the collisional blue stragglers produced in the $N$-body simulation of M67 performed by Hurley \\emph{et al.} in 2005. We assume the collisions are head-on and thus ignore the effects of rotation in this paper. Our detailed models are more luminous than normal stars of the same mass and in the same stage of evolution, but cooler than homogeneously mixed versions of the collision products. The increased luminosity and inefficient mixing decrease the remaining main-sequence lifetimes of the collision products, which are much shorter than predicted by the simple prescription commonly used in $N$-body simulations. } ",
        "introduction": "\\label{sec:intro}  Blue stragglers are stars that appear above and blueward of the main sequence turnoff in the colour-magnitude diagrams (CMDs) of star clusters, with masses larger than that of a turnoff star. They were first identified by \\citet{sandage_m3} in the globular cluster M3 and soon afterwards in the old open cluster M67 \\citep{1955ApJ...121..616J}. Blue stragglers have since been found in other globular clusters and in open clusters of all ages \\citep{1995A&AS..109..375A}.  Various mechanisms have been proposed for their formation \\citep{1993PASP..105.1081S}. Currently, there are % three accepted formation channels: mass transfer due to Roche-lobe overflow in binary systems, and stellar mergers, either due to dynamical collisions or through coalescence of close binaries.   \\citet{1976ApL....17...87H} first showed that physical collisions between main-sequence stars are likely to occur in the dense cores of some globular clusters. In most environments collisions between single stars are very rare, but binary systems can significantly enhance the rate of stellar collisions because of their much larger collisional cross sections. Stellar collisions are thus an important formation channel for blue stragglers even in old open clusters such as M67, as is demonstrated by direct $N$-body calculations \\citep{article:hurley_first_m67,article:hurley_m67}. These simulations indicate that, while all the above-mentioned formation mechanisms operate in different regions of the cluster and all are needed to reproduce the observed blue straggler population, all formation paths -- including binary mass transfer and binary coalescence -- are strongly affected by the dynamical evolution of the cluster.  The blue straggler population therefore contains important information on the dynamical history of a cluster. Extracting this information requires understanding not only how blue stragglers are formed but also how they subsequently evolve.   In $N$-body models collisions between two main-sequence stars are usually approximated by assuming that the stars merge without mass loss and replacing the end product by a normal (evolved) main-sequence star with its age reset in accordance with the assumption of complete mixing \\citep{article:tout_evolution_models, article:hurley_bse}. Smoothed particle hydrodynamics (SPH) simulations of stellar mergers \\citep{1995ApJ...445L.117L,1996ApJ...468..797L} show that some mass is lost during the collision and that the collision products are generally far from being fully mixed. To understand how this affects their further evolution and the predicted blue straggler population, detailed stellar evolution models of the collision products are needed.   In early attempts to model the evolution of stellar merger remnants simplifying assumptions were made: the two stars were assumed to mix completely \\citep{article:bailyn_pinsonneault}, a specified chemical profile was imposed on the remnant \\citep{1995ApJ...455L.163P} or a simplified prescription for heating during the collision was applied to mimic the presumed pre-main sequence contraction phase of evolution \\citep{article:sandquist_bolte_hernquist}.  The first realistic evolution calculations of collision products were performed by \\citet{article:sills_evcolprod1} who took SPH simulations of head-on collisions between detailed stellar models and imported the SPH results into a stellar evolution code. Their models demonstrated that none of the previously made simplifications are valid: although the collision products are inflated due to shock heating, they do not develop substantial convection zones during thermal relaxation and do not undergo significant mixing during their evolution.  This situation changes when the angular momentum of the collisions is considered. For collisions that are even slightly off-axis, the remnants retain too much angular momentum to relax into thermal equilibrium without reaching break-up velocity \\citep{1996ApJ...468..797L}.  The evolution of such collisions was studied by \\citet{sills_evcolprod2}. In the absence of a clear mechanism by which the stars can lose their excess angular momentum, they artificially removed a large fraction of the angular momentum to allow thermal relaxation. Nevertheless, the remnants continue to rotate rapidly throughout their main-sequence evolution and rotational mixing makes the remnants bluer and brighter and significantly extends their lifetimes. These conclusions were confirmed by higher-resolution calculations \\citep{2002MNRAS.332...49S}.  The above studies have focussed on blue stragglers in globular clusters, and investigated only a few interesting cases. This limitation was imposed by the computation time required and numerical difficulties in the evolution calculations.  However, the importance of stellar collisions for the evolution of star clusters calls for a more systematic approach, covering a larger parameter space and extending to higher masses and younger clusters. This is the aim of our current work. We have developed a flexible evolution code that is able to evolve stellar collision products under a wide range of circumstances in a fairly robust manner. As our code cannot yet treat rotation and rotational mixing properly, for the moment we consider only head-on collisions and ignore the effects of rotation.  As a first step in a systematic study of stellar merger remnants we have constructed detailed models of the blue stragglers formed by stellar collisions in the $N$-body model of M67 of \\cite{article:hurley_m67}. They evolved a cluster of 36\\,000 stars from zero age to the age of M67 (4~Gyr) taking into account both cluster dynamics and stellar and binary evolution. In their simulation the cluster evolution resulted in 20 blue stragglers at 4~Gyr, eight of which had a collisional origin. They formed either as a result of dynamical perturbation of a primordial binary, or as a result of three-body (binary-single star) or four-body (binary-binary) interactions. In two of the latter cases, three stars merged in subsequent collisions with the fourth star ending up as a binary companion to the blue straggler. Hence in total ten collisions were involved in the formation of these eight blue stragglers.   We evolve these collision products with our detailed stellar evolution code and compare these models with the evolution tracks of normal detailed stellar models and fully mixed detailed models, as well as with the parametric models used by \\cite{article:hurley_m67}.  In particular we investigate the effect on the main-sequence lifetime of the merger product (\\emph{i.e.}, the time during which it will be visible as a blue straggler), its position in the Hertzsprung-Russell diagram and the effect on the chemical abundances of the remnant. In a companion paper \\citep{GlebbeekPols2008} (\\refereebf{paper II}) we study the influence of varying the collision parameters, in particular the masses of the two stars and their evolutionary stage. ",
        "conclusions": "Hydrodynamical simulations of stellar collisions produce remnants that are out of thermal equilibrium and are not fully mixed.  We have developed an efficient procedure for importing the structural and chemical profiles of such collision products into a fully implicit stellar evolution code. Our evolution code is fairly robust and can evolve the collision remnants until the tip of the giant branch with a minimum of human intervention. We have applied our code to construct detailed models of collisional blue stragglers formed in the $N$-body simulation of M67 by \\citet{article:hurley_m67}.  The evolution of collision products depends on the amount of mixing during the collision and the thermal relaxation phase. Assuming the collision product has been homogeneously mixed produces evolution tracks that are too blue while replacing the collision product with a normal star of the same mass (as done in the simulations of \\citealp{article:hurley_m67}) produces an evolution track that is not bright enough. Both approximations overestimate the lifetime of the collision product. These considerations will affect the predicted colour-magnitude diagram distribution of collisional blue stragglers from a cluster simulation.  Our code is suitable for a systematic exploration of the wide parameter space of collisions in clusters of different ages. This will be the topic of future papers. Eventually we hope to integrate our code into a full $N$-body code to allow for more realistic and self-consistent star cluster simulations."
    },
    "0806/0806.0925_arXiv.txt": {
        "abstract": "During our Swift/XRT program to obtain X-ray positions at arcsecond level for a sample of Galactic X-ray binaries, we discovered that \\saxsource{} is not a binary, but rather BeppoSAX/WFC+GRBM X-ray Rich \\grbsource. Here we report on this discovery and on the properties of this long, X-ray rich gamma-ray burst, from prompt to (very) late followup. ",
        "introduction": "Several catalogued X--ray binaries still have sky positions measured with uncertainties at arcminute level. Such large error boxes prevent a fruitful multiwavelength study; in particular, they make it impossible to establish a firm association with optical/IR/radio counterparts. Therefore, we have a standing Swift fill-in target proposal to observe a sample of such objects drawn from the catalogues of \\cite{Liu2001:lmxbcat,Liu2006:hmxbcat}. Among them is the Low Mass X--ray Binary \\saxsource, which was at first classified as a so-called ``burst-only''  source \\cite{Cocchi2001,Cornelisse2002}. It was  observed only once \\cite{Heise:1998IAUC6892} through a bright ($\\sim 1$ Crab peak intensity in the 2--25 keV energy range), $\\sim 100$\\,s long burst with the Wide Field Cameras \\cite[WFC,][]{Jager1997:SAXWFC} on-board BeppoSAX \\cite{Boella1997:SAX}, at the position RA(J2000$)=08^{\\rm h}$ $40^{\\rm m}$ $40^{\\rm s}$, Dec(J2000$)=+22^{\\circ}$ $48\\arcmin$ $18\\arcsec$ (error radius 3\\arcmin{}). Here we describe the discovery, based on Swift data and reported in \\cite{Sidoli:2007ATel1089}, that \\saxsource{} was actually a gamma ray burst, \\grbsource.  Throughout this paper the quoted uncertainties are given at 90\\% confidence level for one interesting parameter unless otherwise stated. We use $\\Gamma$  as the power-law photon index, $N(E) \\propto E^{-\\Gamma}$ (ph keV$^{-1}$ cm$^{-2}$ s$^{-1}$). ",
        "conclusions": ""
    },
    "0806/0806.2920_arXiv.txt": {
        "abstract": "{Measurements of intracluster gas temperatures out to large radii, where much of the galaxy cluster mass resides, are important for using clusters for precision cosmology and for studies of cluster physics. Previous attempts to measure robust temperatures at cluster virial radii have failed.} {The goal of this work is to measure the temperature profile of the very relaxed symmetric galaxy cluster Abell 2204 out to large radii, possibly reaching the virial radius.} {Taking advantage of its low particle background due to its low-Earth orbit, \\suz\\ data are used to measure the outer temperature profile of Abell 2204. These data are combined with \\cha\\ and \\xmm\\ data of the same cluster to make the connection to the inner regions, unresolved by \\suz, and to determine the smearing due to \\suz's point spread function.} {The temperature profile of Abell 2204 is determined from $\\sim$10 kpc to $\\sim$1800 kpc, close to an estimate of $r_{200}$ (the approximation to the virial radius). The temperature rises steeply from below 4 keV in the very center up to more than 8 keV in the intermediate range and then decreases again to about 4 keV at the largest radii. Varying the measured particle background normalization artificially by $\\pm$10\\% does not change the results significantly. Several additional systematic effects are quantified, e.g., those due to the point spread function and astrophysical fore- and backgrounds. Predictions for outer temperature profiles based on hydrodynamic simulations show good agreement. In particular, we find the observed temperature profile to be slightly steeper but consistent with a drop of a factor of 0.6 from 0.3 $r_{200}$ to $r_{200}$, as predicted by simulations.} {Intracluster gas temperature measurements up to $r_{200}$ seem feasible with \\suz, after a careful analysis of the different background components and the effects of the point spread function. Such measurements now need to be performed for a statistical sample of clusters. The result obtained here indicates that numerical simulations capture the intracluster gas physics well in cluster outskirts.} ",
        "introduction": "Cosmologically, the most important parameter of galaxy clusters is their total gravitational mass. X-rays offer an attractive way to determine this mass through measurements of the intracluster gas temperature and density structures. X-rays are also a unique tool to study the physics of the hot cluster gas; e.g., gas temperature profiles allow constraints on (the suppression of) heat conduction. Consequently, constraining cluster temperature profiles has been the subject of many (partially contradictory) works in the recent past \\citep[e.g.,][]{mmi96,f97,mfs98,ibe99,w00,ib00,asf01,dm02,zfb04,vmm05,app05,kv05,pjk05,hr07,pbc07,smk08,lm08}.  Unfortunately, making these measurements is quite challenging in outer cluster regions. Even with \\xmm\\ and \\cha, it is very difficult to determine temperature profiles reliably out to more than about 1/2 the cluster virial radius, $\\rv$ \\citep[see the references above but also][for a different view]{sas07}. As a result, only about 1/8 of the cluster volume is actually probed. The primary reason is not an insufficient collecting area or spectral resolution of current instruments: the limiting factor is the high particle background. Here, the X-ray CCDs onboard \\suz\\ come into play. Owing to its low-Earth orbit and short focal length, the background is much lower and more stable than for \\cha\\ and \\xmm\\ \\citep{mbi07}, making \\suz\\ a very promising instrument for finally settling the cluster temperature profile debate. And indeed, some of us have recently succeeded in measuring the temperature and metal abundance in the outskirts of the merging clusters A399/A401 \\citep{fth07}. Here, we go one step further and also confirm this great prospect for outer cluster regions whose emissivity is not enhanced by merging activity and determine the temperature profile of the regular cluster Abell 2204 out to $\\sim$1800 kpc with \\suz. This radius is close to an estimate of $r_{200}$; i.e., the radius within which the mean total density equals 200 times the critical density, often used as approximation to the virial radius. Note that shortly after this paper was submitted, another study of a cluster temperature profile towards very large radii with \\suz\\ was submitted by \\citet{gfs08}.  Throughout, we assume $\\om=0.3$, $\\ol=0.7$, and $H_0=71$ km/s/Mpc; i.e., at the redshift of Abell 2204 ($z = 0.1523$), 1' = 157 kpc. ",
        "conclusions": "\\begin{figure*} \\centering \\includegraphics[width=10.cm]{0404fg3a.eps} \\includegraphics[width=6.5cm,angle=270,]{0404fg3b.eps} \\caption{\\emph{Left panel:} Observed outer temperature profile compared to profile and scatter predicted by hydrodynamical simulations of \\citet[solid lines]{red06}. Symbols have the same meaning as in Fig.~\\ref{fig}. For clarity, only the two \\cha\\ and \\xmm\\ data points are shown that were used to determine $k\\tx(0.3 r_{200})$. The other \\cha\\ and \\xmm\\ data points further out are consistent with the simulation results. The dashed error bars of the outer \\suz\\ bin indicate the combined statistical plus systematic error range, calculated as described in the text and using the uncertainties given in Tab.~\\ref{sys}. \\emph{Right panel:} Contamination fractions of the XIS1 bin 7.5$'$--11.5$'$ determined from ray tracing simulations. 75\\% of the photons detected in this bin also originated in this region. The energy dependence is negligible (the lines for all five energies overlap) and the results are very similar for the other detectors. } \\label{fig_simu}% \\end{figure*}  We showed clear evidence that the temperature of this relaxed cluster declines significantly when going from the inner to the outer regions. This is expected theoretically \\citep[e.g.,][]{fwb99}. To compare, in detail, the result obtained here to predictions, we overplotted in Fig.~\\ref{fig_simu} (left) the average temperature profile in cluster outskirts as determined with hydrodynamical simulations of massive clusters by \\citet[][Sample A in their Tab.~2]{red06}. These authors specifically studied the regions around the virial radius. Note that magnetic fields and cosmic rays were not included in their simulations. For the comparison, we used $r_{200}=11.75'$ and $k\\tx(0.3 r_{200})=7.42$ keV. The latter was calculated by taking the average of the best fit \\cha\\ and \\xmm\\ temperatures of the bins that include $0.3 r_{200}$ (shown in Fig.~\\ref{fig_simu}, left). We excluded the \\suz\\ measurement because its broad PSF makes an accurate determination at this radius difficult.  The temperature of the inner \\suz\\ bin shown is slightly higher but consistent with the prediction (keep in mind that, in general, regions left of bin centers carry more weight in the temperature determination than regions right of bin centers, because the surface brightness decreases rapidly with radius). As mentioned above and discussed in more detail below, this bin is strongly affected by PSF smearing, so the actual projected temperature may be lower. Due to the steep surface brightness profile, deprojection would likely result in only a minor increase of the temperature estimate. The important outer bin is much less affected by PSF and projection effects and is slightly lower but consistent with the prediction. In conclusion, the observed outer temperature profile is slightly steeper but consistent with a drop of a factor of 0.6 from 0.3 $r_{200}$ to $r_{200}$, as predicted by simulations.  We tested illustratively what the improvement in the uncertainty of a total mass estimate is due to the improvement of the uncertainty of the temperature profile provided by the \\suz\\ data. We fitted powerlaws to the outermost \\xmm\\ data point and the \\cha\\ and \\suz\\ outermost data points, respectively, taking into account the upper and lower statistical (90\\%) errors of the latter two. Then we determined a fiducial total mass using the best fit single $\\beta$ model parameters for the density profile \\citep[from][]{rb01} and the powerlaw model parameters for the temperature profiles, under the assumption of hydrostatic equilibrium. Using the \\cha\\ temperature we found $M(<1840\\,{\\rm kpc})=5.34^{+2.97}_{-2.10}\\times 10^{14}\\msun$. Using \\suz\\ we found $M(<1840\\,{\\rm kpc})=6.18^{+0.64}_{-0.56}\\times 10^{14}\\msun$; i.e., the uncertainty due to the temperature profile decreased from 40--60\\% (\\cha)\\ to 10\\% (\\suz). So, this represents a significant improvement of a factor of $\\sim$5 (factor of $\\sim$3 when taking into account the different exposure times). Note that a full analysis of the total mass profile will be performed in a more in-depth work.  In the following, we discuss and quantify several systematic effects. We discuss these issues here at some length because this is one of the first papers determining cluster properties out to $r_{200}$, where the surface brightness is very low and systematic effects potentially quite important. In summary, we found that the combination of all quantified systematic effects is of the same order as the statistical uncertainty in the most relevant bin (7.5$'$--11.5$'$). The individual results of the tests are summarized in Tab.~\\ref{sys}.  \\begin{table} \\caption[]{Systematic uncertainties of the cluster temperature measurement in the bin 7.5$'$--11.5$'$.} $$ \\begin{array}{p{0.7\\linewidth}c} \\hline \\noalign{\\smallskip} Test & \\Delta k\\tx/[\\mathrm{keV}] \\\\ \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} 1 Particle background ($-$5\\%)\t& +0.18\t\\\\ \\noalign{\\smallskip} 2 Particle background ($+$5\\%)\t& -0.17\t\\\\ \\noalign{\\smallskip} 3 PSF (25\\% contamination)\t& -0.68\t\\\\ \\noalign{\\smallskip} 4 Foreground modeling (double thermal)\t& +0.61\t\\\\ \\noalign{\\smallskip} 5 ROSAT foreground (thermal)\t& +0.32\t\\\\ \\noalign{\\smallskip} 6 ROSAT background (powerlaw)\t& +0.46\t\\\\ \\noalign{\\smallskip} 7 ARF from ROSAT surface brightness\t& +0.01\t\\\\ \\noalign{\\smallskip} 8 Extrapolated cluster emission beyond $r_{200}$\t& +0.24\t\\\\ \\noalign{\\smallskip} \\hline \\end{array} $$ \\label{sys} \\end{table}  We start with the particle background. The reproducibility of this background component for \\suz\\ is better than about 3\\% \\citep[e.g.,][]{thn08}. Here, we conservatively assumed 5\\% and found best fit temperatures and errors for the third bin as $4.66_{-0.65}^{+0.81}$ keV and $4.32_{-0.54}^{+0.77}$ keV. So, for \\suz, this effect is very small.  Next, we discuss the influence of \\suz's PSF. We performed ray tracing simulations using the xissim tool \\citep{imf07} using $10^7$ photons per energy and detector. We assumed Abell 2204's surface brightness profile to follow the best double $\\beta$ model fit to the \\cha\\ data. We followed a procedure very similar to the one described by \\citet{syi07} and found that the second bin (3.5$'$--7.5$'$) is significantly (49\\%) contaminated by emission originating in the cluster center ($<$3.5$'$). The most important region, the third bin (7.5$'$--11.5$'$) is contaminated by 25\\%; i.e., a relatively small fraction (see Fig.~\\ref{fig_simu}, right, for the contamination fractions of this bin determined for XIS1, as a representative example). The energy and detector dependence of the PSF were found to be small enough to be negligible here. The same is true for the fine details of the model for the surface brightness profile, since we found very similar fractions when repeating this analysis using the best fit double $\\beta$ model as obtained from \\ro\\ \\ps\\ pointed data.  Naively, one could assume that a simple PSF contamination correction could be performed, starting with the assumption that 49\\% of the emission in the 3.5$'$--7.5$'$ bin comes from a plasma at $k\\tx=6.65$ keV (Tab.~\\ref{tab}). However, this would not be accurate, as the \\cha\\ and \\xmm\\ data reveal a wide range of temperatures at different distances from the cluster center (Fig.~\\ref{fig}). A detailed correction for the PSF effects, therefore, requires a large number of ray tracing simulations to be performed, using fine \\cha\\ temperature bins, to determine precisely how much emission from plasma at what temperature contaminates each of the \\suz\\ bins. A detailed correction for the PSF effects is beyond the scope of this paper. Such a correction will be performed in a more in-depth analysis, using the longer \\cha\\ observation of this cluster \\citep{sft09} and the recently updated effective area calibration. Here, we follow a rather conservative route and quantify the changes in temperature and uncertainty by assuming the temperature for the emission, contaminating the important third bin, to lie in the range 6.0--8.0 keV. Furthermore, we freeze the metallicities of the third bin to 0.2 and that of the contaminating emission to 0.3. This results in the following best fit temperatures and uncertainties: $k\\tx(7.5'$--$11.5') = 4.13_{-0.77}^{+0.96}$ keV (for $kT_{\\rm contami}=6$ keV) and $k\\tx(7.5'$--$11.5') = 3.81_{-0.67}^{+0.90}$ keV (for $kT_{\\rm contami}=8$ keV). Changing the metallicities of the third bin to 0.15 and 0.25 has a quite negligible influence on the best fit temperatures. The same is true when changing the metallicity from the contaminating emission from the second bin to 0.25 and 0.35. In Tab.~\\ref{sys} and the calculation of the total systematic uncertainty we use the more conservative ($kT_{\\rm contami}=8$ keV) result.  We specifically designed the \\suz\\ observation in such a way as to have a cluster free region available to help constrain Galactic fore- and cosmic X-ray background directly and locally, using the same observation. This ensures that the results are not affected by systematic calibration differences between different satellites and varying point source subtraction fractions. Still, we checked whether the resulting parameter values are in a reasonable range. We found that the normalization of the powerlaw component expected from deep CXB studies using several different satellites is lower by factors 0.77--0.62 compared to our best fit normalization. We then froze the powerlaw normalization to the highest \\citep{vmg99} and lowest \\citep{rgs03} normalization we found in the literature, resulting in temperatures $6.08^{+0.98}_{-0.82}$ and $6.84^{+1.17}_{-0.72}$ keV, respectively, in the third bin. This procedure results in worse fits but note that, in the latter case, the temperature is significantly hotter than when leaving the powerlaw normalization free. Since intensity variations of the CXB may be correlated with large scale structure, we decided to use the independent \\ro\\ \\ps\\ observation of Abell 2204 and fitted the same fore-/background model in the energy band 0.3--2.4 keV to the outskirts. We found that the normalization is fully consistent (factor 0.81--1.04) with the higher powerlaw normalization from the \\suz\\ data of Abell 2204. Moreover, we also found the temperature and normalization of the thermal component to be fully consistent with the \\suz\\ results. In particular, for the Galactic emission, we found temperatures of $0.25_{-0.01}^{+0.02}$ keV (\\suz)\\ and $0.26_{-0.01}^{+0.01}$ keV (\\ro), quite typical of this foreground component; i.e., there was no indication of significant hot cluster or accretion shock emission ``contaminating'' annulus 4. We, therefore, conclude that our Galactic fore- and cosmic X-ray background modeling is adequate.  Recall that the purely statistical uncertainties on these fore- and background components are already included in our quoted statistical errors on the cluster temperatures because in our standard analysis we performed a simultaneous fit of all the astrophysical fore- and background components together with the cluster component; i.e, the relevant parameters for these models were all free to vary. For the calculation of the total systematic uncertainty, we additionally included effects, which relate to the modeling of the Galactic thermal emission as well as the spatial variation of the fore- and background components, as described in the following.  Several authors used more than one spectral component to model the Galactic foreground emission \\citep[e.g.,][]{syi07,smk08,hs08}. Therefore, we tried a two component foreground model (phabs*apec+apec) instead of a single component model (apec) for the Suzaku observation. As best fit temperatures for the Galactic components we found 0.28 keV for the absorbed component and 0.13 keV for the unabsorbed one, in good agreement with the typical values quoted by \\citet{smk08}. Furthermore, the new powerlaw model normalization (representing the unresolved cosmic X-ray background) is only 2.5\\% lower than in the original fit. This change is much smaller than the statistical error on this normalization.  We found that the resulting new best fit cluster temperature in the most important and most affected third bin (7.5$'$--11.5$'$) gets a bit higher but not significantly so (5.10 keV). On the one hand, this lends confidence to our approach; on the other hand, this result is also not too surprising since in our spectral analysis we ignored all photons with energies E$<$0.7 keV, so the cluster temperature measurements are by construction less sensitive to uncertainties in the soft fore- and background.  Also, $\\chi_{\\rm red}^2$ does not change by adding this additional model component, showing that no significant improvement can be achieved by adding this second thermal component for the data under consideration here.  Moreover, the relative statistical uncertainties stay very similar, e.g., for the best fit cluster temperature in the third bin the errors at the 90\\% confidence level are $-$13\\% +18\\% (single thermal Galactic foreground model) and $-$15\\% +21\\% (double thermal Galactic foreground model); therefore, the simple single thermal model does not result in a significant underestimate of statistical errors. In our systematic error analysis we do include the model dependence in the final error.  Last not least, we tested the double thermal model in the \\ro\\ analysis of the foreground emission. We found very similar best fit temperatures for the Galactic components (absorbed 0.25 keV, unabsorbed 0.13 keV), although \\ro's poor energy resolution results in significantly enlarged errors for the temperatures of the two thermal components (the two temperature ranges overlap; degeneracies cannot be broken, resulting in an unstable error analysis). The change in the powerlaw normalization is again well within the statistical uncertainty.  Next, we tested the influence of a possible spatial variation of the fore- and background estimates. The regions selected for the fore- and background analysis in the \\ro\\ observation differ from those in the \\suz\\ observation. Therefore, we refit the \\suz\\ data but this time not letting the powerlaw and thermal component vary freely but, separately, freezing them to the values determined from the \\ro\\ observation (after correction for the different covered areas). The best fit cluster temperatures for the third bin do not change significantly. They are $4.81_{-0.65}^{+0.81}$ keV (foreground apec model frozen) and $4.95_{-0.66}^{+0.85}$ keV (background powerlaw model frozen). Both effects are included in the total systematic error analysis.  The Monte Carlo ARF calculation using xissimarfgen requires a priory knowledge of the cluster surface brightness distribution, which we provided using a double $\\beta$ model fit to the surface brightness measured with \\cha. We tested the influence of deviations from this assumed distribution on the cluster temperature measurements by instead generating ARFs using the best fit double $\\beta$ model from the \\ro\\ observation. The resulting best fit temperature in the third bin was almost unchanged (4.50 keV).  We assumed that the ``cluster free'' region (the forth bin) contains negligible cluster emission and argued that the best fit temperature for the Galactic foreground emission we found supports this assumption. Nevertheless, we performed an additional test to estimate the contribution of this assumption to the total systematic error. To this end, we determined a rough surface brightness profile using bins 1 through 3 and extrapolated it by conservatively assuming the surface brightness to drop by a factor of 5 from bin 3 to bin 4. Furthermore, we assumed a temperature of 2 keV and a metallicity of 0.2 solar in bin 4. We then included this cluster component in the model for bin 4, in addition to the fore- and background components, and redid the full simultaneous fit. The resulting new best fit cluster temperature in bin 3 is 4.73 keV; i.e., slightly higher than the original fit result but not significantly so. This rather small change is expected because the extrapolated surface brightness of this fourth bin is much lower than the surface brightness of the powerlaw background component. Therefore, the new normalization of the powerlaw background component is also only slightly lower (7\\%) compared to the original one. A detailed determination and discussion of the surface brightness profile and especially its error will be presented in a subsequent paper. We do add the effect of the possible cluster emission in bin 4 to the systematic error budget.  Some of the systematic effects described above raise the temperature (e.g., using a double thermal model for the Galactic foreground emission) others lower it (e.g., roughly accounting for PSF effects). To estimate the total systematic error as well as the combined statistical plus systematic error, we used the following scheme. We separately added all positive and negative errors from Tab.~\\ref{sys} in quadrature, resulting in a total systematic error $_{-0.70}^{+0.88}$ keV. The total systematic error is, therefore, slightly larger but of the same order as the statistical uncertainty ($_{-0.59}^{+0.79}$ keV). Then, we added the systematic and statistical errors in quadrature, resulting in a combined statistical plus systematic error $_{-0.91}^{+1.18}$ keV (shown as dashed error bars in Fig.~\\ref{fig_simu}).  For the comparison to simulated temperature profiles, the uncertainty of $r_{200}$ as estimated from the extrapolated \\xmm\\ density and temperature profiles may be important. For instance, using a cluster mean temperature and the relation of \\citet{emn96} would result in a larger value for $r_{200}$. Also, \\citet{app05} determined $r_{200} = 2075$ kpc using the same \\xmm\\ data, probably due to the flatter temperature profile they found. This would make the observed temperature drop even steeper in terms of $r_{200}$, potentially resulting in tension between observation and simulations. We will redetermine $r_{200}$ from the cluster mass profile, taking advantage of the information on the outer temperature from \\suz\\ in the more in-depth analysis of this cluster. It is reassuring to note that further corroboration for the $r_{200}$ estimate used here comes from the independent weak lensing analysis of this cluster \\citep{cs02}, yielding $r_{200}=11.8'$, in perfect agreement with our estimate.  The confirmation of the predictions from hydrodynamical simulations for the gas physics in cluster outskirts indicated here rests on the analysis of a single cluster. What is required for a general confirmation is the analysis of a statistical sample of clusters. In the future, this can be performed with dedicated cluster observations accumulating rapidly in the \\suz\\ archive."
    },
    "0806/0806.2488_arXiv.txt": {
        "abstract": "We present {\\it Spitzer} Space Telescope IRAC and MIPS observations toward a sample of nine high-mass star forming regions at a distance of around 2 kpc. Based on IRAC and MIPS 24 $\\mu$m photometric results and 2MASS \\emph{JHKs} data, we carry out a census of young stellar objects (YSOs) in a $5'\\times5'$ field toward each region. Toward seven out of the nine regions, we detect parsec sized clusters with around 20 YSOs surrounded by a more extended and sparse distribution of young stars and protostars. For the other two regions, IRAS 20126+4104 and IRAS 22172+5549, the former has the lowest number of YSOs in the sample and shows no obvious cluster, and the latter appears to be part of a larger, potentially more evolved cluster. The deep IRAC imaging reveals at least twelve outflows in eight out of the nine regions, with nine outflows prominent in the 4.5 $\\mu$m band most probably attributed to shocked H$_2$ emission, two outflows dominated by scattered light in the 3.6 and 4.5 $\\mu$m bands, and one outflow standing out from its hydrocarbon emission in the 8.0 $\\mu$m band. In comparison with previous ground-based observations, our IRAC observations reveal new outflow structures in five regions. The dramatically different morphologies of detected outflows can be tentatively interpreted in terms of possible evolution of massive outflows. The driving sources of these outflows are deeply embedded in dense dusty cores revealed by previous millimeter interferometric observations. We detect infrared counterparts of these dusty cores in the IRAC or MIPS 24 $\\mu$m bands. Reflection nebulae dominated by the emission from UV heated hydrocarbons in the $8\\,\\mu$m band can be found in most regions and they may imply the presence of young B stars. ",
        "introduction": "It is well know that high-mass stars form in dense clusters \\citep{Lada03}. The research on embedded young clusters around young intermediate- to high-mass stars has advanced significantly via near-infrared surveys \\citep[e.g.,][]{Testi99, Gutermuth05, Kumar06}. The {\\it Spitzer} space telescope with its unprecedented sensitivity at mid-infrared wavelengths provides a new means for identifying young stellar objects (YSOs) in star forming regions through detection of dusty disks and envelopes. High sensitivity in the $3-24$ $\\mu$m bands is invaluable for distinguishing young stars from reddened background stars through infrared excess arising from circumstellar materials \\citep{Allen04, Megeath04, Whitney04, Gutermuth04, Muzerolle04}. This capability is particularly valuable for studies of high-mass star forming regions for two reasons: 1) such regions are typically located near the Galactic plane where the density of background stars is high and 2) the member YSOs are often faint due to the large distances to these regions ($>1$kpc). The faint magnitude of the YSOs complicates not only the detection of members but also their identification since the density of background stars rises rapidly with increasing magnitudes. With the Infrared Array Camera (IRAC) and Multiband Imaging Photometer for \\emph{Spitzer} (MIPS), a significant census of YSOs can be carried out, which provides a fossil record of the distribution of star formation sites and information on the environments in which stars and planets form. This method is particularly effective for finding extended, lower density distributions of stars surrounding the dense regions of embedded clusters.  Extended structures, such as reflection nebulae and outflows, can also be mapped with the IRAC. Outflows emanating from the central young stars or protostars are usually detectable in rotational transitions of CO and other molecules at (sub)millimeter wavelengths. In IRAC observations, outflows can be revealed from shocked H$_2$ emission. Hydrodynamic simulations of shock models predict that H$_2$ emission is particularly strong in the 4.5 $\\mu$m band \\citep{Smith05}. Since the 4.5 $\\mu$m band is relatively free of the emission from hydrocarbons (probably in the form of Polycyclic Aromatic Hydrocarbons, or PAHs) compared with the other three bands, and the 5.8 and 8.0 $\\mu$m bands are significantly less sensitive than the shorter wavelength bands, composite images with emission from different bands coded in different colors can be a diagnostic tool for H$_2$ emission in outflows. A remarkable example of such a detection can be found in the IRAC observations of the HH 46/47 outflow, a collimated bipolar outflow emanating from a low-mass protostar, where the southwestern lobe was clearly seen in the 4.5 $\\mu$m band as a limb-brightened bow-shock cavity \\citep{Raga04, Noriega04b}. \\citet{Smith06} reported the IRAC detection of a spectacular outflow in the high-mass star forming region DR 21 by virtue of its brightness in the 4.5 $\\mu$m band. Further investigations of the \\emph{ISO} SWS spectra of the DR 21 outflow verified that the H$_2$ lines contribute significantly to all four IRAC bands and the emission in the 4.5 $\\mu$m band is mostly attributed to the H$_2$ v=0-0 $S(9)$ (4.695 $\\mu$m) line. Outflow activities can also be revealed from scattered light in the IRAC bands. By comparison of the deep IRAC imaging of L1448 with the radiative transfer modelling, \\citet{Tobin07} demonstrated that the observed infrared emission in the IRAC bands is consistent with the scattered light from the central low-mass protostars escaping the cavities carved by molecular outflows. \\citet{Valusamy07} reprocessed the IRAC observations of the HH 46/47 outflow using a deconvolution algorithm and found a wide-angle biconical component in scattered light in addition to the previously reported bow-shock cavity. Finally, in some cases hydrocarbon emission is visible from outflow cavities, presumably due to the illumination of the cavity walls by UV radiation from the central source \\citep{vandenAncker00}.  We have observed a sample of nine high-mass star forming regions with the IRAC and MIPS instruments (PID: 3528, PI: Qizhou Zhang). These regions were chosen according to the following criteria: 1) already imaged with millimeter interferometers; 2) known to have molecular outflows; 3) outside of the GLIMPSE survey except IRAS 19410+2336; 4) at a distance of $<2.5$kpc ($<5000$AU at the resolution of IRAC). These regions exhibit far-infrared luminosities ranging from $10^3$ to $10^5\\,L_{\\odot}$ (see Table \\ref{source} for source parameters). We carry out a census of YSOs toward each region, and simultaneously, zoom in on the central massive star formation sites with the aim of searching for infrared counterparts of millimeter continuum sources and infrared emission from outflows. In this paper we present initial results of the observations. In {\\S} 2, we describe the observations. We present observational results in {\\S} 3 and subsequent discussion in {\\S} 4. A brief summary is given in {\\S} 5. ",
        "conclusions": "\\subsection{Clustering} \\label{dis_cluster} One of the distinctive features of massive star formation is that it occurs in clusters. Consequently, understanding the role of clustering in the formation of high-mass stars is an essential step toward a theory of massive star formation. A controversial issue is whether clusters are necessary for the formation of high-mass stars, or whether clusters are merely the byproduct of the formation of high-mass stars in dense massive molecular cores. In a near-infrared survey of intermediate-mass Herbig Ae/Be stars, \\citet{Testi99} found a smooth transition between the low density aggregates of young stars associated with stars of spectral type A or later, and the dense clusters associated with O and early-B type stars. They argued that the presence of dense clusters is required for the formation of high-mass stars. The smooth transition suggested that there may be a fundamental relationship between the mass of the most massive star and the number and/or density of stars in a cluster. However, \\citet{deWit05} found evidence that 4\\% of O stars form in isolation. This is evidence that clusters are not necessary for the formation of O stars, although \\citet{Parker07} argued that the isolated O stars may actually form in small clusters.  To test whether massive stars can form without low-mass stars requires identification of young massive objects without associated clusters of low-mass stars; only in these cases, can we rule out the possibility that an attendant cluster dispersed due to dynamical evolution. In a $K$-band survey of eight high-mass star forming cores, \\citet{Megeath04} found one core, NGC 6334 I(North), showing evidence of high-mass star formation without evidence for a cluster of embedded low-mass stars \\citep[also see ][]{Megeath99, Hunter06}, although a cluster may still exist deeply embedded in the cloud \\citep{Persi05}. It is important to search for other potential examples of massive star formation without the presence of a dense cluster. With its ability to identify YSOs through infrared excesses and detect deeply embedded protostars, \\emph{Spitzer} is well suited to this task. Although such regions may yet form clusters, the identification of high-mass young or protostars without a dense cluster would demonstrate that low-mass stars are not required for the formation of massive stars.  Our survey of young intermediate- to high-mass (proto)stars ($\\sim10^3$ to $10^5L_{\\odot}$) can directly address the role of clustering in the early stages of intermediate- and high-mass forming stars. The {\\it Spitzer} data provides the ability to directly identify likely YSOs by the detection of mid-infrared excesses. The advantage of this approach over previous near-infrared source counting methods is that the sample does not significantly suffer from contamination by background stars which can be significant for high-mass star forming regions in the Galactic plane \\citep{Pratap99}. Also, the mid-infrared colors are much more sensitive to protostellar objects, and candidate protostellar objects can be identified through their colors. However, the bright mid-infrared nebulosity in high-mass star forming regions, the $\\sim2-3''$ angular resolution of the IRAC data, and the modest sensitivity of the 2MASS photometry limit the detection of sources, in particular in regions of bright nebulosity. We expect the census of YSOs to be incomplete in all of the regions. Furthermore, we expect that a certain fraction of the stars do not have disks and will not be identified in our analysis; the fraction of stars without disks is approximately 30\\% for 1 Myr old low-mass stars \\citep{Hernandez07}.  We group the objects by their total far-infrared luminosity, as measured from the IRAS point source catalog; this luminosity will be dominated by the most massive objects in the regions. In all the regions, associated low-mass YSOs are detected.  The $\\sim10^3L_{\\odot}$ regions, AFGL 5142, G192, IRAS 05358 and IRAS 20293, show parsec sized clusters with around 20 YSOs surrounded by a more extended and sparse distribution of young stars and protostars. The $\\sim10^4L_{\\odot}$ regions IRAS 19410 and HH 80-81 also show clusters and extended components. Except AFGL 5142, the clusters associated with these regions do not show concentrations of stars toward the central massive objects. This may be due to the decrease in completeness toward the bright nebulosities in the central regions. Finally, the $10^{5.1}L_{\\odot}$ region W75\\,N shows a cluster that is both richer and more spatially extended. This region also shows a paucity of sources in the bright nebulous center of the region, further indicating the observations are significantly limited by completeness.  In Figure \\ref{fig17} we plot the number of sources as a function of the source luminosity. To mitigate the effects of incompleteness on our samples, we plot the number of sources with 3.6 $\\mu$m magnitudes lower than 13 and 14. Given a typical $Ks-[3.6]$ color of 0.25, an extinction of $A_K\\sim0.25$ at 1.8 kpc, and an age of 1 Myr, these limits correspond to masses of 1.2 and 0.5 $M_{\\odot}$ \\citep{Baraffe98}. The two plots show a trend of increasing number of associated YSOs with increasing total luminosity. In particular, the number of YSOs in the $10^{5.1}L_{\\odot}$ region (W75\\,N) is significant higher than those in $\\sim10^4L_{\\odot}$ regions, which again are somewhat higher in number than the $\\sim10^3L_{\\odot}$ regions. It is possible that a trend of increasing incompleteness with higher far-infrared luminosities is decreasing the slope of this trend.  There are two main exceptions to this trend: IRAS 22172 and IRAS 20126. IRAS 22172 appears to be part of a larger, potentially more evolved region which has undergone significant gas clearing. The clearing of the gas would also lower the fraction of light absorbed by dust and re-emitted in the infrared; hence this region shows a large number of stars for its measured IRAS far-infrared luminosity. The other exception is IRAS 20126, which has the lowest number of associated YSOs in the sample. Unlike the other regions, IRAS 20126 shows no obvious cluster in the field. This is unlikely to be solely the effects of incompleteness, with the average signal of about 2.5 MJy/sr at 4.5 $\\mu$m and 52 MJy/sr at 8.0 $\\mu$m in the central parsec area, compared to 6.5 MJy/sr at 4.5 $\\mu$m and 67 MJy/sr at 8.0 $\\mu$m for HH 80-81 and 3.8 MJy/sr at 4.5 $\\mu$m and 72 MJy/sr at 8.0 $\\mu$m for IRAS 19410.  More rigorous studies of this sample await deeper near-infrared data to complement the {\\it Spitzer} 3.6 and 4.5 $\\mu$m imaging. However, the analysis here already illustrates several features. First, there does seem to be a trend of the number of associated YSOs with luminosity. Second, these regions show both clusters in the central 1 pc as well as more extended halos of sources around the clusters. Finally, there does seem to be a significant dispersion in the number of associated YSOs for a given total far-infrared luminosity, with one luminous source, IRAS 20126, showing no obvious cluster. The lack of a cluster warrants future work; if confirmed this would demonstrate that the formation of $\\sim10^4L_{\\odot}$ sources is not dependent on the presence of deep clusters.  \\subsection{Outflows and Outflow Cavities} For a sample of nine high-mass star forming regions associated with CO and/or SiO outflows, the deep IRAC imaging reveals outflows and outflow cavities toward eight regions, illustrating that the IRAC imaging can be an effective tool for outflow detection. For the jets in IRAS 05358, AFGL 5142, IRAS 20293, the bow-shock shells in W75\\,N, and the nebulosities in the eastern lobe of the G192 outflow, the prominence in the 4.5 $\\mu$m band and coincidence with the ground-based 2.12 $\\mu$m H$_2$ detections (if detected in the 2.12 $\\mu$m H$_2$ line) suggest the emission is mostly attributed to shocked H$_2$ emission, although the contribution from H{\\scriptsize I} Br$\\alpha$ (4.052 $\\mu$m), CO v=1-0 (4.45-4.95 $\\mu$m) lines cannot be ruled out without spectroscopic observations. A detailed investigation of the emission in the biconical cavity in IRAS 20126 shows strong evidence for the scattered light being dominant in the 3.6 and 4.5 $\\mu$m bands (Qiu et al., in prep.). The elliptical structure in IRAS 19410 appears relatively diffuse and most prominent in the 3.6 $\\mu$m band. The emission in this outflow can be dominated by the scattered light as well. The clear prominence of the biconical structure in HH 80-81 suggested is dominated by UV heated hydrocarbon emission.  \\subsubsection{An evolutionary picture of massive outflows?} Both theoretical models and high spatial resolution observations suggest an evolutionary scenario for low-mass outflows: as low-mass stars evolve from class 0 through class I to class II stages, the mass loading in the winds and density structure in the cores conspire the widening of outflows \\citep{Fuller02, Arce06, Shang07}. While for massive outflows, whether an evolutionary picture exists is unknown \\citep{Beuther05}.  The detected outflows in the sample show dramatically different morphologies. We detect highly collimated jets, bow-shock shells, and biconical cavities. Given the extreme complexity of the CO observations \\citep{Beuther03} and relatively confused emission in the IRAC imaging compared with the other regions, we leave out the IRAS 19410 outflow in the following discussion. Considering Jet1 in IRAS 05358, the short jet in AFGL 5142, and the jet in IRAS 20293, the IRAC observations reveal the internal driving agents of well collimated CO outflows in these regions. The powering sources of Jet1 in IRAS 05358 and the short jet in AFGL 5142 are suggested to be deeply embedded proto-B stars, toward which only very weak radio continuum was detected \\citep{Beuther07,Zhang07}. The jet in IRAS 20293 is most likely driven by an intermediate-mass protostar showing no detectable centimeter emission \\citep{Beuther04b}. In contrast, the large scale CO outflows in W75\\,N and G192 appear to be poorly collimated within $\\lesssim0.5$ pc from the central driving sources. In W75\\,N, proper-motion observations of water masers by \\citet{Torrelles03} delineated a non-collimated, shell outflow at a 160 AU scale expanding in multiple directions with respective to the UC H{\\scriptsize II} region VLA 2, which may drive the large scale CO outflow as suggested by \\citet{Shepherd03}. The bipolar bow-shaped structure beyond $\\sim1$ pc detected in the 2.12 $\\mu$m H$_2$ and the IRAC 4.5 $\\mu$m imaging could be the remanent of a collimated component. For the G192 outflow, no collimated component can be found within $\\sim2.4$ pc from the driving source. The proposed central driving sources of these two poorly collimated outflows are found to be surrounded by UC H{\\scriptsize II} regions. If we adopt the centimeter free-free emission as an indicator of the relative evolution between these sources, there seems to be a trend of evolution for the related outflows: both internal driving agents and entrained molecular outflows appear to be highly collimated for the youngest sources (e.g., IRAS 05358, AFGL 5142, IRAS 20293); as the central source evolves to form a significant UC H{\\scriptsize II} region, only poorly collimated outflow structures can be found around the central driving source (e.g., W75\\,N, G192).  For the IRAS 20126 outflow, it remains ambiguous whether it is merely composed of a precessing jet and consequently larger structures are all entrained/swept-up by this jet or it has a jet-like component as well as a biconical component with a moderate opening angle. Toward the central driving source of this outflow, very weak centimeter continuum was detected, presumably suggesting the source being at a very young evolutionary stage. For the HH 80-81 outflow, our \\emph{Spitzer} data and previous centimeter continuum reveal a remarkable two-component outflow with a biconical cavity surrounding an axial radio jet. Strong centimeter continuum was detected toward the central driving source of the HH 80-81 outflow, presumably suggesting the source being relatively more evolved compared with that of IRAS 20126. The outflow cavity of HH 80-81 also shows emission in the mid-infrared hydrocarbon features, suggesting that UV radiation from the central sources are heating the cavity walls. Therefore, if the IRAC biconical cavity in IRAS 20126 is a complementary less collimated component of the previously detected SiO jet, there will be a large evolutionary time scale for the existence of both a highly collimated jet-like component and a less collimated biconical component for outflows from proto-B stars. With the existing data it is very difficult to determine whether the central source of IRAS 20126 is more evolved than those of AFGL 5142 and IRAS 05358, and whether the central source of HH 80-81 is younger than those of W75\\,N and G192. Presumably it is likely that the IRAS 20126 and HH 80-81 outflows, with the former being younger, represent evolutionary stages between that of the AFGL 5142 \\& IRAS 05358 outflows, for which both internal driving agents and entrained gas appear well collimated, and that of the W75\\,N \\& G192 outflows, for which only poorly collimated structures can be found around the central driving sources.  However, such a trend is only a tentative interpretation based on current observations of a small sample. The question whether there is an evolutionary picture for massive outflows is far from conclusive, and remains an observational challenge. In high-mass star forming regions, multiple outflows with complicated structures are often detected. Due to relatively larger distances and crowded clustering mode, it is very difficult with the current facilities to reliably identify the driving source of a massive outflow. Also resolving the launching zone of a massive outflow, which is of great importance for investigate driving mechanisms, requires extremely high spatial resolution and high sensitivity. In addition, observations of outflows in more luminous objects ($L\\,{\\gtrsim}\\,10^5L_{\\odot}$), in particular high spatial resolution observations, are still rare (the combined bolometric luminosity of UC H{\\scriptsize II} regions in MM-1 in W75\\,N is $10^{4.6}L_{\\odot}$). Extensive and comprehensive studies of outflows in this luminosity regime would be crucial for testing possible evolutionary scenarios of massive outflows and understanding formation processes of O type stars."
    },
    "0806/0806.0819_arXiv.txt": {
        "abstract": "We numerically explore planet formation around $\\alpha$ Centauri A by focusing on the crucial planetesimals-to-embryos phase. Our approach is significantly improved with respect to the earlier work of \\cite{mascho00}, since our deterministic N-body code computing the relative velocities between test planetesimals handles bodies with different size. Due to this step up, we can derive the accretion vs. fragmentation trend of a planetesimal population having any given size distribution. This is a critical aspect of planet formation in binaries since the pericenter alignment of planetesimal orbits due to the gravitational perturbations of the companion star and to gas friction strongly depends on size. Contrary to \\cite{mascho00}, we find that, for the nominal case of a Minimum Mass Solar Nebula gas disc, the region beyond $\\sim 0.5\\,$AU from the primary is strongly hostile to planetesimal accretion. In this area, impact velocities between different-size bodies are increased, by the differential orbital phasing, to values too high to allow mutual accretion. For any realistic size distribution for the planetesimal population, this accretion-inhibiting effect is the dominant collision outcome and the accretion process is halted. Results are relatively robust with respect to the profile and density of the gas disc. Except for an unrealistic almost gas-free case, the inner \"accretion-safe\" area never extends beyond 0.75AU. We conclude that planet formation is very difficult in the terrestrial region around $\\alpha$ Centauri A, unless it started from fast-formed very large ($>$30km) planetesimals. Notwithstanding these unlikely initial conditions, the only possible explanation for the presence of planets around 1 AU from the star would be the hypothetical outward migration of planets formed closer to the star or a different orbital configuration in the binary's early history. Our conclusions differ from those of several studies focusing on the later embryos-to-planets stage, confirming that the planetesimals-to-embryos phase is more affected by binary perturbations. ",
        "introduction": "\\label{sec:intro}   A majority of stars are members of binary or multiple systems \\citep[e.g.][]{duq91}, making the presence and frequency of planets in such systems, in particular terrestrial planets, a very relevant issue in our ongoing search for other worlds. This issue has been made all the more timely by the discovery of more than 40 extrasolar planets in multiple systems, amounting to $\\sim 20\\%$ of all detected exoplanets \\citep{ragha06,desi07}. Most of these planets are located in very wide binaries ($>100$AU) where the companion star's influence at the planet's location is probably weak \\citep[except for possible high inclination effects between the planet's and the binary's orbits, see e.g.][]{takeda08}. However, a handfull of planets, like Gliese 86, HD41004 or $\\gamma$-Cephei, inhabit binaries with separations $\\sim$20 AU, for which the binarity of the system probably cannot be ignored.  The long term stability of planetary orbits in binaries has been explored in several studies \\citep[e.g.][]{holw99,dav03,mud06}, aimed at estimating $a_\\mathrm{crit}$, the distance from the primary beyond which orbits become unstable. This critical distance depends on several parameters, mainly the binary's semi-major axis $a_b$ and eccentricity $e_b$ as well as the mass ratio $\\mu=m_2/(m_1+m_2)$. However, one can use as a rule of thumb that, except for extremely eccentric systems, $a_\\mathrm{crit}$ is in the 0.1-0.4$a_b$ range \\citep[see for example the empirical formulae derived by][]{holw99}. As a consequence, the terrestrial planet region around 1\\,AU is almost always \"safe\" for binaries of separations $>10$AU.   The way the planet \\emph{formation} process is affected by stellar binarity is a more difficult issue. This process is indeed  a complex succession of different stages \\footnote{we consider here the  widely accepted \"core-accretion\" scenario} \\citep[e.g][]{lis93}, each of which could react in a different way to the companion's star perturbations.  The last stage of planet formation, leading from planetary embryos to planets, has been investigated in several papers \\citep[e.g][]{barb02,quin02,quin07,hag07}. These studies have revealed that the ($a_b,e_b,\\mu$) parameter space for which this stage can proceed unimpeded is slightly, but not too much more limited than that for orbital stability. As an illustration, the numerical exploration of \\citet{quin07} showed that planetary embryos can mutually accrete at 1AU from the primary if the companion's periastron $q_2$ is $\\geq\\,5\\,$AU.  In contrast, the stage preceding this final step, i.e. the mutual accretion of kilometre-sized planetesimals leading to the embryos themselves, is much more sensitive to binarity effects. A series of numerical studies \\citep{mascho00,theb04,theb06,paard08} have indeed shown that perturbations by the companion star can stir up encounter velocities $\\langle \\Delta v \\rangle$ to values too high for small planetesimals to have accreting impacts. Such accretion-inhibiting configurations are obtained for a wider range of binary parameters than for the later embryo-to-planet phase. As an example, \\citet{theb06} (hereafter TMS06) have shown that planetesimal accretion in the 1\\,AU region can be perturbed in binaries with separations as large as 40 AU. These results are not surprising in themselves, and this for two reasons. 1) The growth mode in the planetesimal-accumulation phase is the so-called runaway growth, whose efficiency is very sensitive to increases of the encounter velocities (which reduce the gravitational focusing factor, e.g. TMS06). 2) It takes a much smaller perturbation to stop accretion of a 1km body than of a lunar-sized embryo: $\\langle \\Delta v \\rangle>$10m.s$^{-1}$ is typically enough to achieve the former, while large embryos will hardly feel any difference in their accretion rate for such small velocity disturbances. A crucial mechanism driving the $\\langle \\Delta v \\rangle$ evolution of the planetesimals is friction due to the primordial gas remaining in the protoplanetary disc. In a circumprimary disc, its main effect is to induce a strong phasing of neighboroughing planetesimal orbits, which cancels out the large orbital parameters oscillations around the forced eccentricity $e_f$ induced by the companion's perturbations \\citep{mascho00}. A crucial point is that this effect is \\emph{size dependent}, so that while $\\langle \\Delta v \\rangle$ between equal-sized objects are reduced, impact velocities between objects of \\emph{different sizes} are increased. TMS06 have shown that the differential phasing effect appears already for small size differences, suggesting that this effect may dominate in \"real\" planetesimal systems. However, the exact balance between this accretion-inhibiting effect and the accretion-friendly effect on equal-sized objects depends on the size distribution in the initial planetesimal population, a parameter which could not be thoroughly investigated in the global investigation of TMS06.   \\subsection{planet formation in \\alp\\ A}  We follow here the approach of TMS06 to investigate in detail one specific system: \\alp\\ A. Focusing on one single system allows to take the TMS06 model a step further and explore realistic size distributions of the planetesimal population as well as crucial parameters such as gas disc density and radial profile.  At a distance of 1.33\\,pc, the $\\alpha$ Cen AB binary system, with a possible third distant companion M--star Proxima \\citep{wert06}, is our closest neighbour in space. \\alp\\ A and B are G2V and K1V stars, of masses $M_A = 1.1 M_{\\odot}$ and $M_B = 0.93 M_{\\odot}$ respectively. The binary's semi-major axis is $a_b=23.4\\,$AU and its eccentricity $e_b=0.52$ \\citep{pourb02}. \\citet{holw97} have shown that the region within $\\sim 3\\,$AU from the primary can harbour stable planetary orbits. However, the search for planetary companion has so far been unsuccessful. The radial velocity analysis of \\citet{endl01} has shown that the upper limit for a planet (on a circular orbit) is $2.5\\,M_{Jupiter}$ around $\\alpha$ Cen A, and this at any radius from the star. This leaves an open possibility for the presence of terrestrial planets in the crucial $\\sim 1\\,$AU region.  Planet formation around $\\alpha$ Cen A has so far been specifically investigated in 4 studies. \\citet{barb02} and \\citet{quin02,quin07} have focused on the late embryos-to-planets stage and have found that it can proceed unimpeded in a wide inner region ($\\leq 2.5\\,$AU). For the planetesimal-accretion phase, \\citet{mascho00} have found that orbital phasing induced by gas drag allowed km-sized bodies to accrete in the $r\\leq2\\,$AU region. However, this study only considered impacts between equal-sized objects, thus implicitly overlooking the accretion-inhibiting differential phasing effect. We reinvestigate here this crucial phase, exploring in detail gas drag effects on realistic size distributions for the planetesimal population. ",
        "conclusions": "\\label{sec:concl}  Our numerical exploration has shown that for our nominal case set-up the $r\\geq 0.5\\,$AU region around \\alp\\ is hostile to kilometre-sized planetesimal accretion. The coupling of the companion star's secular perturbations with the size-dependent phasing imposed by gas drag leads to high $\\langle \\Delta v \\rangle_{\\rm(s_1,s_2)}$ destructive collisions for all $s_1 \\neq s_2$ pairs. For any realistic size distribution for the planetesimal population (except very unlikely Dirac-like size distributions), these erosive impacts outnumber by far the low-velocity $s_1 \\sim s_2$ encounters and are the dominant collision outcome. Only the innermost region within $\\sim 0.5\\,$AU has $\\langle \\Delta v \\rangle$ low enough to allow planetesimal accretion.  Interestingly, these results are relatively robust when varying the gas disc density. The main difference is that the outer limit of the inner \"accretion-safe\" region varies between 0.35 and 0.75\\,AU. Only for very gas poor systems do we get a qualitatively different behaviour, with the safe region extending up to $\\sim 1.3\\,$AU, where a high-$\\langle \\Delta v \\rangle$ zone due to pure secular perturbations begins.  These results are further strengthened by the fact that our simplified assumption for the gas disc probably results in a slight underestimate of $\\langle \\Delta v \\rangle_{\\rm(s_1,s_2)}$ values (see discussion in Sec.\\ref{sec:gas}). As a consequence, our approach can be regarded as conservative regarding the extent of the impact velocity increase which is reached, the \"real\" system being probably even more accretion-hostile than in the presented results.  This hostility to kilometre-size planetesimal accretion is of course a major threat to planet formation as a whole, of which it is a crucial stage. However, before reaching any definitive conclusions, one has to see if there might be some ways for planetary formation to get round this hurdle. In the next 3 subsections we critically examine three such possibilities.     \\subsection{Accretion from large initial planetesimals?}  We have seen that, apart from an unrealistic gas-free case, the only way to get an accretion friendly environment at 1\\,AU is to assume large, $\\geq 30\\,$km planetesimals which can sustain the high velocity regime without being preferentially eroded (Fig.\\ref{accbig}). This solution to the high velocity problem raises however important issues. One of them is how realistic it is to assume a population of large initial planetesimals. As already pointed out, the planetesimal formation process is still too poorly understood to reach any definitive conclusion regarding what an initial planetesimal population looks like, so that values in the 30-50\\,km range cannot be ruled out \\emph{a priori}. Nevertheless, one has to ask the question of what was \\emph{before} these large planetesimals. Should these objects arise from the progressive accumulation of grains, rocks and pebbles in a turbulent nebula \\citep[as in the standard coagulation model, e.g.][]{dul05}, then there is no clear cut separation between the planetesimal-formation phase and the subsequent planetesimal-to-embryo phase \\footnote{In this case, the very concept of an \"initial\" planetesimal population might be questioned \\citep[see for instance][]{wetina00}}. In this case, the stage where the biggest planetesimals are $\\sim$50\\,km in size should be preceded by a stage where the biggest objects are in the 1-10\\,km range, which leads us back to the accretion-inhibiting nominal case. The only solution would be to bypass these intermediate accretion-inhibiting stages by directly forming large planetesimals. This could in principle be achieved with the alternative gravitational instability scenario \\citep[e.g][]{gold73,you04}, but this scenario still has several problems to overcome \\citep[see for instance][ for a brief review]{arm07}. Moreover, it is not clear how gravitational instabilities should proceed in the highly perturbed environment of a binary.  It should also be noted that even \\emph{if} at some point this accretion-friendly large-planetesimals stage is reached, the way accretion will proceed from there would be very different from the standard runaway-accretion scheme. Indeed, as can be clearly seen in Fig.\\ref{accbig}, the ($s_1$,$s_2$) parameter space for which accretion can proceed unimpeded is restricted to a narrow $s_1 \\sim s_2$ band. For most impacts on $s\\geq 30\\,$km targets, even if relative velocities are below $v*_{(s_1,s_2){\\rm low}}$ (the \"blue\" zone), they are nevertheless increased relative to the unperturbed case. These increased $\\langle \\Delta v \\rangle$ would strongly reduce the gravitational focusing factor which is the key process driving runaway accretion. In this case, the growth mode could be the so-called \"type II\" runaway identified by \\citet{kort01}, corresponding to a slowed-down runaway growth in a high $\\langle \\Delta v \\rangle$ environment.   \\subsection{Accretion once gas has dissipated?}  An interesting issue is what happens \\emph{after} gas dispersal. Indeed, it is believed that primordial gas discs dissipate, due to photo-evaporation and/or viscous accretion onto the central star, in a few $10^{6}$ years \\citep[e.g][]{alex06}, a timescale which could even be shortened in binaries because of the outer truncation of the protoplanetary disc. One might thus wonder if, despite of a long accretion-hostile period while gas is present, planetesimal accretion could start anew once the environment becomes gas free and differential phasing effects disappear. However, for this to occur, there is a big difficulty to overcome. Indeed, when the environment becomes gas free, all planetesimals are still on the differentially-phased orbits they were forced onto by gas friction. This means that the \\emph{initial} dynamical conditions are such as $\\langle \\Delta v \\rangle_{\\rm(s_1,s_2)}$ are high. Under the sole influence of the companion star's secular perturbations (the only external force left in the gas free system), these initial orbital differences cannot be erased or damped. The size-independent orbital phasing which will be forced by secular effects will come \\emph{in addition} to the initial orbital differences, which can in this respect be regarded as an initial free component. This behaviour can be easily checked by gas free test simulations.  Moreover, even if a hypothetical additional mechanism acts to damp the initial differential phasing, it has to act quickly in order to lead to low impact velocities. Indeed, if given enough time, secular perturbations \\emph{alone} will eventually induce high $\\langle \\Delta v \\rangle$ due to orbital crossing of neighbouring orbits. The inner limit of $\\sim$1.3\\,AU for such orbital crossing given in Sec.\\,\\ref{sec:accero} is indeed the one reached at $t=10^{4}yrs$, but this limit is moving inwards with time. Using the empirical analytical expressions derived by TMS06 (see Eq.15 of that paper), we find that the orbital crossing \"front\" reaches the 0.5\\,AU region after $\\sim 1.5\\times 10^{5}$yrs.  As a consequence, it seems unlikely that $\\langle \\Delta v \\rangle_{\\rm(s_1,s_2)}$ can be significantly reduced in the later gas free system. Collisional erosion, started in the gas-rich initial disc, should thus probably continue once the gas has dissipated \\footnote{As this paper was being reviewed, a yet unpublished study by Xie and Zhou \\citep[ApJ, submitted, with a prelimimary summary in][]{xie08} was brought to our attention. It explores the possibility, for the specific case of $\\gamma$ Cephei and under possibly extreme assumptions for gas dispersal timescales, that some \"re-phasing\" might occur {\\it during} the gas dissipation phase}.    \\subsection{Outward planet migration and different initial binary configuration}\\label{sec:diss}  It could be argued that planets could have formed elsewhere and later migrated to the $r\\geq 0.75$AU regions (the maximum extent of the accretion-friendly region obtained in our gas-disc profile exploration). Early migration of the protoplanets by interaction with the primordial gas disc might be ruled out since such a migration \\citep[be it of type I, II or III, e.g][]{papa07} would be inwards. In the present case, the planets would thus have to come from further out in the disc, i.e, from regions which are even more hostile to accretion. Outward drift of planets could however occur for another type of migration, the one triggered, after gas-disc dispersal, by interactions between a multi-planets system and a massive disc of remaining planetesimals \\citep[as in the so-called Nice-model, e.g.][]{tsig05}. Netherless, it remains to see how this scenario, which was studied for mutually interacting giant planets at radial distances between 5 and 30\\,AU, would work in the present case, $\\sim 10$ times closer to the star and with probably much less massive planets. This important issue will be addressed in a forthcoming paper (Scholl et al., in preparation).  An alternative possibility would be that the binary's orbital configuration was not the same in the past as it appears today. If the presently observed binary system was part of an unstable multiple system (triple or more), planets might have formed when the binary had a larger semi-major axis and it was less eccentric \\citep{marbar07}. Both these conditions could lead to a more favourable environment for planetesimal accretion than present. In this scenario planets should have formed prior to the onset of the stellar chaotic phase which inevitably ends with the ejection of one star, leaving the remaining binary system with smaller separation and higher eccentricity. The discussion of the possible consequences of a transient triple system stage for \\alp\\ A and B is complicated by the presence of Proxima Cen, a red dwarf located roughly a fifth of a light-year from the AB binary. A recent paper by \\citet{wert06} argues that the binding energy of Proxima Centauri relative to the centre of mass of the $\\alpha$ Centauri binary is indeed negative, even if the star is close to the outer border of the Hill's sphere of the AB system in the galactic potential. This might be an indication of a violent past of the system when the AB binary was possibly part of a higher multiplicity unstable stellar system. This issue exceeds the scope of the present paper and should be explored in future studies.   \\subsection{Conclusion and Perspectives}  Because of the hypothetical additional effects listed in Sec.\\ref{sec:diss}, the presence of planets beyond 0.75\\,AU cannot be fully ruled out. However, we think that our results on planet \\emph{formation}, with the present binary configuration, are relatively robust. Our main result is that it is very difficult for $s\\leq30$\\,km planetesimals to have accreting encounters beyond 0.5-0.75\\,AU from the primary, which makes planet formation very unlikely in these regions. These results are in sharp contrast to those of \\citet{mascho00}, who found the terrestrial region aournd \\alp\\ A to be favourable to planetesimal accumulation. The main reason for this discrepancy is that \\citet{mascho00} restricted their pioneering study to collisions between \\emph{equal-size} bodies, for which gas friction indeed decreases $\\langle \\Delta v \\rangle$ and favours accretion, thus implicitly neglecting the dominant effect of eroding $s_1 \\neq s_2$ impacts. Our results are also in contrast to the conclusions of \\citet{barb02} and \\citet{quin02,quin07}, who regarded planet accretion as been possible in the inner $\\leq 2.5\\,$AU region. This is because these studies focused on the later embryo-to-planet phase, thus implicitly assuming that the preceding planetesimal-to-embryo phase was successful. The present study shows that this is probably not the case. This confirms that the planetesimals-to-embryos phase is more affected by the binary environment than the last stages of planet formation.  Due to the complexity of the coupling between secular perturbations and gas drag, and their strong dependence on binary orbital elements, the present results cannot be directly extrapolated to other systems. Clearly, such detailed studies of planetesimal accumulation in other close binary systems, either those known to harbour planets (like $\\gamma$ Cephei \\footnote{$\\gamma$ Cephei has already been investigated by \\citet{theb04}, but with a simplified version of the present code, in particular with respect to the role of planetesimal size distributions} or HD41004) or those being potential candidates, should be undertaken in the close future. In this respect, one object of particular interest could be $\\alpha$ Cen A's companion, Alpha Centauri B, for which \\citet{guedes08} very recently explored the possibilities for observational detection of terrestrial planets."
    },
    "0806/0806.3278_arXiv.txt": {
        "abstract": "We present and analyze the correlations between mid-infrared (MIR), far-infrared (FIR), total-infrared (TIR), H$\\alpha$, and FUV luminosities for star-forming galaxies, composite galaxies and AGNs, based on a large sample of galaxies selected from the $Spitzer$ SWIRE fields. The MIR luminosities of star-forming galaxies are well correlated with their H$\\alpha$, TIR and FUV luminosities, and we re-scaled the MIR-derived SFR formulae according to the above correlations with differences less than 15\\%. We confirm the recent result by \\citet{calzetti07} that the combined observed H$\\alpha$ $+$ 24$\\mu$m luminosities L(H$\\alpha$$_{\\rm obs}$$+$ 24$\\mu$m) possess very tight correlation with the extinction-corrected H$\\alpha$ luminosities L(H$\\alpha$$\\_$corr) for star-forming and even for dwarf galaxies, and show that the combined L(H$\\alpha$$_{\\rm obs}$$+$ 8$\\mu$m[dust]) are also tightly correlated with L(H$\\alpha$$\\_$corr) for the above sample galaxies. Among all the L(MIR)-L(FIR) correlations for star-forming galaxies, the L(24$\\mu$m) vs. L(70$\\mu$m) and L(8$\\mu$m[dust]) vs. L(160$\\mu$m) are the tightest and also nearly linear. The former could be related to young massive star formation, while the latter might be relevant to diffuse dust emissions heated by old stellar populations. Composite galaxies and AGNs have higher MIR-to-H$\\alpha$/MIR-to-FUV luminosity ratios than star-forming galaxies, nevertheless their correlations among MIR, FIR and TIR luminosities are completely following those of star-forming galaxies. ",
        "introduction": "Since the successful launch of $Infrared~Astronomical~Satellite$ ($IRAS$) in 1983, more advanced space-based infrared telescopes had been launched, such as $Infrared~Space~Observatory$ \\citep[$ISO$;][]{kessler96} and $Spitzer~Space~Telescope$ \\citep{werner04}, to help as to study the MIR to FIR emission properties of galaxies. MIR to FIR emissions are crucial for quantifying extragalactic star formation activities, since infrared emissions are arising from the dust re-radiation of UV photons emitted by young massive stars. FIR luminosity has been proved to be a good star formation rate (SFR) indicator of galaxies by $IRAS$ observations \\citep[e.g.,][]{hunter86, lehnert96}. However, the low sensitivity and spatial resolution of FIR observations obstructed us from applying it to most galaxies at intermediate and high redshifts, even though there are some degree of improvements in the $ISO$ and $Spitzer$ era. Optical/UV SFR tracers, such as H$\\alpha$-line emission and FUV continuum luminosities, are important because they are still the only way to measure SFR in regions of low extinction such as faint, dwarf galaxies and especially extended UV disk (XUV-disk) galaxies \\citep{thilker07}. But these tracers usually suffer from serious extinction effects for normal and especially infrared-bright galaxies which are difficult to be corrected \\citep{jonsson04}. MIR emissions of normal galaxies are dominated by dust continuum from very small grains (VSGs), together with some broad emission features \\citep{gillett73, willner77} which were realized to be from polycyclic aromatic hydrocarbons \\citep[PAHs;][]{leger84, puget89}. Early analysis of correlations between MIR luminosity (derived from $ISO$ observations) and other SFR tracers \\citep[e.g.,][]{elbaz02, roussel01, schreiber04, flores04} have shown that the MIR luminosity could be used to estimate SFRs of galaxies.  Mid-Infrared observations by $Spitzer$ with higher sensitivity and better angular resolution than $ISO$, provide a new opportunity to study young stellar populations and star formation processes in galaxies. The 8$\\mu$m band of $Spitzer$ Infrared Array Camera \\citep[IRAC;][]{fazio04} is designed to cover the 7.7$\\mu$m PAH feature; whilst the 24$\\mu$m band of Multiband Imaging Photometer for $Spitzer$ \\citep[MIPS;][]{rieke04} just covers the dust continuum from VSGs for local galaxies, avoiding the contamination of most of the MIR emission or absorption lines. The pioneering work by \\citet{wu05} (hereafter Wu05) has shown that, for star-forming galaxies, both $Spitzer$ IRAC 8$\\mu$m(dust) and MIPS 24$\\mu$m luminosities are strongly correlated with the 1.4~GHz radio and H$\\alpha$ luminosities. Then the new formulae for estimating SFRs of galaxies were derived from MIR luminosities based on above correlations. Detailed studies of nearby galaxies M~51 \\citep{calzetti05} and M~81 \\citep{perez06} revealed that the 24$\\mu$m luminosities of star-forming regions in the above two galaxies were well correlated with the extinction-corrected Pa$\\alpha$ luminosities. Based on the 8$\\mu$m and 24$\\mu$m data for HII regions in 33 nearby SINGS \\citep[$Spitzer$ Infrared Nearby Galaxies Survey;][]{kennicutt03} galaxies, \\citet{calzetti07} found that both 8$\\mu$m(dust) and 24$\\mu$m luminosities are correlated with L(Pa$\\alpha$$\\_$corr), but non-linearly. They and \\citet{kennicutt07} also demonstrated that the combination of observed H$\\alpha$ (uncorrected for dust extinctions) and 24$\\mu$m luminosities could be one of the best SFR indicator for HII regions in these galaxies.  Besides being heated by young massive stars, the energy source of MIR and FIR emissions could also be the active galactic nuclei (AGN) in the galactic center, since AGNs are also powerful UV emitters. Weak AGN-hosting galaxies are very popular in the local universe \\citep{ho97}, thus the possible AGN contribution to MIR and FIR emissions could have significant impact on the estimation of cosmic SFR density and the understanding of its evolution \\citep[see, e.g,][]{bell05,perez06,elbaz07,daddi07}. Previous works showed that powerful AGNs can strengthen MIR VSG emissions, but weaken or even avoid PAH emissions \\citep{weedman05,siebenmorgen04,verma05}. Weak AGNs are found to have MIR properties different from that of star-forming galaxies \\citep{wen07,li07}, and the degree of such differences could be related to the AGN activity \\citep{wu07}.  The $Spitzer$ Wide-area Infrared Extragalactic Survey \\citep[SWIRE;][]{lonsdale03} is the largest extragalactic survey program among the six $Spitzer$ cycle-1 Legacy Programs. With a total field of $\\sim$49 deg$^{2}$, it is much larger than the $Spitzer$ First Look Survey (FLS) field ($\\sim$3.7 deg$^{2}$) studied by \\citet{wu05}. In this paper, we perform statistical analysis on the correlations between MIR, FIR, H$\\alpha$, and FUV luminosities for star-forming galaxies, composite galaxies (which were affected by both central nuclear and star formation activities), and AGNs selected from three northern SWIRE fields.  The structure of the paper is as follows. We describe the construction of our sample and the estimation of multi-wavelength luminosities in $\\S$2 and $\\S$3. The major results on correlation analysis are presented in $\\S$4. Some discussions and a summary of this work are given in $\\S$5 and $\\S$6. Throughout this paper, we adopted a $\\Lambda$CDM cosmology with $\\Omega_{\\rm m}=0.3$, $\\Omega_{\\rm \\Lambda}=0.7$ and $H_{\\rm 0}=70\\,{\\rm km \\, s^{-1} Mpc^{-1}}$. ",
        "conclusions": "\\subsection{Discrepancies in L(MIR)-L(H$\\alpha$) Correlations} The aperture correction to H$\\alpha$ luminosities (see $\\S$3.3) were performed based on the assumption that the distributions of H$\\alpha$ emission in galaxies are similar to those of the continuum (SDSS-$r$) emission. All of our sample galaxies have both detectable H$\\alpha$ and H$\\beta$ emission lines above 5$\\sigma$ significance level, thus it could bias towards galaxies with stronger (circum-)nuclear star formation activities. In the outer parts of these types of galaxies, star formation activity would be much weaker and the optical continuum emissions could be dominated by old stellar populations. In such cases, we would overestimate the $total$ H$\\alpha$ luminosity if performing corrections from SDSS $r$-band magnitude. The aperture correction effect discussed above could explain the shifts between best fits of $total$ and $aperture$ MIR-to-H$\\alpha$ luminosity correlations shown in Figures~\\ref{fig4} and described in $\\S$4.1.  Besides the aperture correction effect discussed above, differential dust obscuration could be another reason for such discrepancies. The central region of galaxies is more dusty than their outer parts \\citep[e.g.,][]{popescu05,prescott07}, and a majority of UV continuum and Balmer line emissions might be completely enshrouded by dust thus of course can't be corrected. There is no doubt that the above effect will be more severe for the central regions than for the entire galaxies, thus it could cause a higher $aperture$ MIR-to-H$\\alpha$ luminosity ratio than the $total$ value, in accordance with the discrepancies shown above.  Dwarf galaxies were shown to be obviously deviated from the star-forming galaxies in the L(MIR)-L(H$\\alpha$) correlations (see Figure~\\ref{fig4}(a) and (b)). This large deviation could be mainly caused by the aperture correction effect discussed above, since many dwarf galaxies have intensive star formation in their nuclear regions \\citep[e.g.,][]{gu06,lisker06}. This interpretation can also be proved by the $aperture$ L(8$\\mu$m[dust])-L(H$\\alpha$) correlation shown in Figure~\\ref{fig4}(d), in which dwarf galaxies are only slightly deviated from the star-forming galaxies. Alternatively, this slight deviation could be explained by the fact that PAH and warm dust emissions are relatively weak in low metallicity environments \\citep[e.g.,][]{madden00, engelbracht05, wyl06, galliano05, wu07, engelbracht08}.  \\subsection{Estimation of SFRs from MIR Luminosities} Based on a sample of $Spitzer$ FLS-field galaxies, \\citet{wu05} showed that both $Spitzer$ 8$\\mu$m(dust) and 24$\\mu$m luminosities could be used to calculate SFRs of normal star-forming galaxies, and they gave corresponding MIR-derived SFR formulae calibrated by either H$\\alpha$ or radio luminosities. However, the size of Wu05's sample is a little small (only have $<$ 80 galaxies) and the MIR luminosities of their sample galaxies only cover a range of two orders of magnitude. In this work, we construct a large sample of $379$ star-forming galaxies with the MIR luminosities range of three orders of magnitude. Besides correlating well with H$\\alpha$ luminosities, the MIR luminosities of our sample galaxies are also found to have tight correlations with TIR and FUV luminosities. These correlations allow us to estimate and compare SFRs from MIR luminosities calibrated by different SFR indicators.  Although the correlations between L(MIR) with L(H$\\alpha$) \\& L(TIR) are somewhat nonlinear, the nonlinearities of the above two correlations are generally less than 10\\%. Therefore, we can adopt linear correlations for calibrating SFRs (SFRs derived from non-linear correlations between L(MIR) with other luminosities also were illustrated in Figure~\\ref{fig13}) . Based on the three linear correlations of L(MIR)-L(H$\\alpha$), L(MIR)-L(TIR) and L(MIR)-L(FUV), we obtained that the SFR conversion factors of L(24$\\mu$m) are 7.15$\\times$10$^8$, 7.79$\\times$10$^8$, and 7.04$\\times$10$^8$ L$_{\\odot}$, and those of L(8$\\mu$m[dust]) are 1.58$\\times$10$^8$, 1.70$\\times$10$^8$, and 1.46$\\times$10$^8$ L$_{\\odot}$, respectively. These factors are in good agreement with those of L(MIR)-L(H$\\alpha$) and L(MIR)-L(radio) correlations by Wu05 with differences less than 10\\% (Figure~\\ref{fig13}). In fact, the differences in L(8$\\mu$m[dust])-L(H$\\alpha$) correlations between Wu05's sample and ours are mainly because of the galaxies with very low MIR luminosities. However, it must be noted that the MIR SFR formulae (especially for 8$\\mu$m(dust)) derived above can not be applied to dwarf galaxies (and other galaxies with low metallicity; \\citealt{wu07}), since they deviate much from the linear correlations for normal star-forming galaxies. However, the deviation of dwarfs to normal star-forming galaxies in Figure~\\ref{fig5} is apparently smaller than in Figure~\\ref{fig4}. Therefore, it may be a simple and feasible method of using the combined observed H$\\alpha$ + MIR luminosities to compute SFRs. The black solid line in Figure~\\ref{fig13} show the equation to compute SFR based on 24$\\mu$m luminosity given by \\citet{alonso06a} (their equation 3), which was derived from the correlation between 24$\\mu$m and H$\\alpha$ luminosities for 30 LIRGs in the local universe, and this equation is similar to the one by \\citet{calzetti07} except for different IMFs. We should keep in mind such large differences for infrared luminous galaxies.  For TIR emission, the contribution from dust heating by the old stellar population could not be ignored ($\\sim$30\\% for normal galaxies, \\citet{bell03}); while for UV emission, except for the pollution from old star, the effect of intrinsic extinction is even stronger than that for H$\\alpha$ emission. Although the standard deviations of the fitting residuals from TIR (0.07 for 24$\\mu$m, 0.06 for 8$\\mu$m(dust)) are smaller than from H$\\alpha$ (0.16, 0.18 respectively), the H$\\alpha$ emission trace the HII region surrounded massive star more directly. Therefore, we suggest that the MIR-derived SFR based on the MIR vs. H$\\alpha$ correlation is ``better'' than those derived from other correlations.  \\subsection{Correlations among L(MIR), L(FIR) and L(TIR)} The MIR luminosities L(MIR) are mainly from PAH and warm dust (VSGs) emissions, while FIR luminosities L(FIR) are dominated by cold dust emissions from large grains. From Figure~\\ref{fig6}, and comparing with the correlations between L(MIR) and L(H$\\alpha$) shown in Figure~\\ref{fig4} and Table~\\ref{tab3}, we find that the MIR luminosities correlate better with L(FIR) than L(H$\\alpha$), and this result is consistent with that from \\citet{boselli04} for late-type galaxies observed by $ISO$.  However, correlations shown in Figures~\\ref{fig4} and~\\ref{fig6} are on the basis of two different (MIR \\& FIR) samples. Therefore, we firstly checked the distributions of absolute B-band magnitude, redshift, $u-r$ color, and EW(H$\\alpha$) for both MIR and FIR sample galaxies (see Figure~\\ref{fig2}), and found the two samples have similar distributions. To avoid possible selection effect, the L(MIR)-L(H$\\alpha$) correlations for FIR sample galaxies are also shown in Figure~\\ref{fig14}. The best fits for galaxies in the FIR sample are consistent with those for Wu05 and MIR sample galaxies. The standard deviations for the FIR sample are about 0.15, a little smaller than those for the MIR sample, since the FIR sample do not include galaxies with L(MIR) $<$ 10$^8$$L_{\\odot}$; while their Spearman Rank-order correlation analysis coefficients are similar to those of the MIR sample. From the comparison described above, we confirm that the MIR luminosities are better correlated with FIR than H$\\alpha$ luminosities for star-forming galaxies.  As described in $\\S$4.3 and shown in Figure~\\ref{fig6}, among the four L(MIR)-L(FIR) correlations, L(24$\\mu$m)-L(70$\\mu$m) and L(8$\\mu$m[dust])-L(160$\\mu$m) have much smaller scatters (the standard deviation of the fitting residuals are both 0.07) than the rest two correlations of L(24$\\mu$m)-L(160$\\mu$m) and L(8$\\mu$m[dust])-L(70$\\mu$m) (the standard deviation of the fitting residuals are 0.12 and 0.10 respectively). Furthermore, the slopes of both L(24$\\mu$m)-L(70$\\mu$m) and L(8$\\mu$m[dust])-L(160$\\mu$m) correlations are close to unity, which indicates that such correlations are almost linear. The above results hint that the 24$\\mu$m vs. 70$\\mu$m, and 8$\\mu$m vs. 160$\\mu$m emissions may have similar physical origins:\\\\ $~~~~$(a) The 24$\\mu$m emission is known to be dominated by warm dust emissions from VSGs, which are mainly heated by young massive stars. Recent results by $Spitzer$ \\citep{wu05,perez06,calzetti07} have shown that the 24$\\mu$m luminosity is one of the best SFR indicators. Previous works based on $IRAS$ and $ISO$ observations have proved that the FIR luminosity is also a good SFR tracer because of the emission peak around 60$\\mu$m of dust heated by star formation, thus the 70$\\mu$m emission must be closely related to star formation activities and should have tight and linear correlation with 24$\\mu$m warm dust emission.\\\\ $~~~~$(b) The 8$\\mu$m(dust) luminosity was also thought to be a SFR tracer by some authors \\citep[e.g.,][]{wu05,alonso06b}. 8$\\mu$m dust emissions include emissions from both PAHs and VSGs. In star-forming galaxies, the 7.7$\\mu$m PAH emission dominates the 8$\\mu$m band \\citep{smith07} and its strength is likely to be affected by many factors, such as the radiation field and metallicity \\citep[e.g.,][]{madden00,galliano05,cao07,wu07}. Moreover, PAH emissions trace B star better than young O stars \\citep{peeters04}. Therefore, the 8$\\mu$m(dust) luminosity may not be an optimal SFR tracer, and this could explain why the 70$\\mu$m luminosities have better correlation with 24$\\mu$m rather than 8$\\mu$m(dust) luminosities.\\\\ $~~~~$(c) The 160$\\mu$m emission is mainly from cold dust heated by old stars and is thought to locate in the diffuse regions of galaxies \\citep[e.g.,][]{gordon04}; while the PAH molecules are generally located in the photo-dissociation regions (PDRs), at the interface of HII regions \\citep{hollenbach97,draine03}. The tight L(8$\\mu$m[dust])-L(160$\\mu$m) correlation indicates that the 8$\\mu$m(dust) and 160$\\mu$m emissions might be from the same regions of galaxies, i.e., cold dust emission could exists in PDRs and it is also possible that there exists PAH emission in the diffuse regions because PAHs can be pumped by cool stars with less UV radiation \\citep{li02}. Due to the low spatial resolution and shallowness of $Spitzer$ 160$\\mu$m observations, firm conclusions on the physical connection between 8$\\mu$m(dust) and 160$\\mu$m emissions and their distributions in galaxies must await high quality FIR data from future space telescopes (e.g., $Herschel$).  Figure~\\ref{fig6} shows that both MIR and FIR luminosities are well correlated with TIR luminosities for star-forming galaxies. Except for L(24$\\mu$m), all of the 8$\\mu$m(dust), 70$\\mu$m and 160$\\mu$m luminosities are almost linearly correlated with the L(TIR). This result indicates that the monochromatic MIR and FIR luminosities can be used for estimating the TIR luminosity, in agreement with previous studies \\citep{elbaz02, takeuchi05, marcillac06}.  \\subsection{Properties of Composite galaxies and AGNs} Investigating the properties of composite galaxies and AGNs will provide us some information on the origin of MIR emission in AGN-hosting galaxies. Figures~\\ref{fig8} and~\\ref{fig10} show that composite galaxies and AGNs have distinct L(MIR)-L(H$\\alpha$) and L(MIR)-L(FUV) correlations, i.e., they (especially for the ones with low H$\\alpha$ luminosity) have relatively higher MIR-to-H$\\alpha$/MIR-to-FUV luminosity ratios than star-forming galaxies do. The high MIR-to-H$\\alpha$/MIR-to-FUV luminosity ratios can not be caused by heavier dust obscuration in the host galaxies of composites and AGNs, since from Figure~\\ref{fig3}(c) it is clear that the host galaxies of composites and AGNs tend to be earlier Hubble types with redder $u-r$ colors, and it is well-known that early-type galaxies should have less gas and dust \\citep[e.g.,][]{wen07, li07}. Alternatively, for early-type galaxies in which the diffuse IR emissions can not be neglected, their MIR-to-H$\\alpha$/MIR-to-FUV luminosity ratios would tend to be higher due to additional dust heating by abundant old stellar populations. Furthermore, the higher luminosity (Figure~\\ref{fig3}(a)) could refer to higher metallicity\\citep[e.g.,][]{lamareille04}, while higher metallicity in early-types \\citep{wu07, li07} could be another explanation to the higher MIR-to-H$\\alpha$ luminosity ratios of composite galaxies and AGNs except the re-emission of dust heated by AGN.  From Figure~\\ref{fig11} it is clear that all the four L(MIR)-L(FIR) correlations for composite galaxies and AGNs are completely following those for star-forming galaxies. The L(MIR)-L(TIR) and L(FIR)-L(TIR) correlations for composite galaxies and AGNs are also shown to follow those for star-forming galaxies (see Figure~\\ref{fig12}). Hence, the contribution from the active nuclear to both MIR and FIR emissions is more smaller than the contribution from the HII region, and it could be neglected for those weak AGNs. Nevertheless, the similarity of dust heating by AGN or HII region could be another possible explanation. These results indicate that we can estimate the total infrared luminosity accurately from monochromatic MIR and FIR luminosities, not only for star-forming galaxies (see also $\\S$5.3), but also for AGN-hosting galaxies. However, firm conclusion must await a quantitative analysis of a large, well-defined, and unbiased sample of FIR-selected galaxies from $AKARI$ and $Herschel$ observations."
    },
    "0806/0806.3752_arXiv.txt": {
        "abstract": "{Whether or not there exists a physical upper mass limit for star clusters is as yet unclear. For small cluster samples the mass function may not be sampled all the way to the truncation, if there is one. Data for the rich cluster population in the interacting galaxy M51 enables us to investigate this in more detail.} {Using \\textit{HST/ACS} data, we investigate whether the cluster luminosity function (LF) in M51 shows evidence for an upper limit to the mass function. The variations of the cluster luminosity function parameters with position on the disk are addressed.} {We determine the cluster LF for all clusters in M51 falling within our selection criteria, as well as for several subsets of the sample. In that way we can determine the properties of the cluster population as a function of galactocentric distance and background intensity. By comparing observed and simulated LFs we can constrain the underlying cluster initial mass function and/or cluster disruption parameters. A physical upper mass limit for star clusters will appear as a bend dividing two power law parts in the LF, if the cluster sample is large enough to sample the full range of cluster masses. The location of the bend in the LF is indicative of the value of the upper mass limit. The slopes of the power laws are an interplay between upper mass limits, disruption times and evolutionary fading.} {The LF of the cluster population of M51 is better described by a double power law than by a single power law. We show that the cluster initial mass function is likely to be truncated at the high mass end. We conclude from the variation of the LF parameters with galactocentric distance that both the upper mass limit and the cluster disruption parameters are likely to be a function of position in the galactic disk. At higher galactocentric distances the maximum mass is lower, cluster disruption slower, or both.} {} ",
        "introduction": "The distinction between star clusters and galaxies used to be fairly clear, but recently the gap between both has been slowly filled up with objects in the intermediate mass range. The most massive star clusters were the Milky Way's old globular clusters, with masses going up to a few $10^6 \\,\\msun$. The least massive galaxies known were galaxies of the type of Fornax, having masses of about $10^7 \\,\\msun$. Many objects are now determined to be in between these distinct mass ranges, e.g. `Ultra Compact Dwarfs' \\citep[e.g.][]{mieske04, hilker99}, `stellar superclusters' \\citep[e.g.][]{hilker99} and `Dwarf-Globular Transition Objects' \\citep[e.g.][]{hasegan05}. Whether these are all different names for the same objects or not remains to be seen.  Are these new classes of objects really formed in distinct formation processes, or is there rather a continuous transition between all, and are star clusters just a more common mode of formation? \\citet{kissler-patig06} put young massive star clusters ($M = 10^{6.5 - 7}\\,\\msun$) on the scaling relations (relations between mass, velocity dispersion and radius) of other hot stellar systems and find that these clusters do follow scaling relations derived for the least massive galaxies, i.e. dwarf elliptical galaxies. Less massive clusters do not show a mass-radius relation. Simulations by e.g. \\citet{fellhauerkroupa05} and \\citet{bekki04} support the idea that the most massive cluster-like objects may have formed from merging smaller clusters, thereby establishing a mass-radius relation \\citep{kissler-patig06}.  As this would imply that the formation scenarios of systems considerably more massive than $\\sim 10^6 \\,\\msun$ are really distinct from `normal' star clusters, it is interesting to see whether we can also find evidence for a maximum mass for the star cluster-like objects. As pointed out by e.g. \\citet{jordan07} and \\citet[and references therein]{waters06}, the universality of the globular cluster luminosity function (GCLF) of globular clusters in other galaxies is much easier to achieve when including a fundamental upper mass limit, in the form of an exponential cut-off, for star clusters in dynamical cluster evolution models. This also holds for the globular cluster system of the Milky Way \\citep[e.g.][]{burkertsmith00}.  The universal turn-over in the globular cluster mass function is achieved more easily with cluster evolution models when including a exponential cut off in the mass function around $10^6$ \\msun.  To obtain detailed information (for example the mass, age and extinction) of a single cluster of which the stars are not resolved requires photometry in a wide range of broadband filters (UV to IR, e.g. \\cite{degrijs05}, and references therein), or spectroscopy. In the study of star cluster populations, where spectroscopy is a very time consuming process, the cluster LF is a useful tool. This is especially true when photometry is only available in a small number of passbands and obtaining the mass, metallicity and extinction of separate clusters is hampered by the age-metallicity-extinction degeneracy. In such cases, the LF can be used as a statistical tool, giving only information about the population as a whole.  In various environments the LF can be approximated well by a power law distribution ($N \\, \\d L \\propto L^{-\\alpha}\\, \\d L$), with a slope ($-\\alpha$) between -1.8 and -2.4 (e.g. \\citet{larsen02,degrijs03a}). The translation from the LF to the cluster initial mass function (cluster IMF) is non-trivial, however, especially when clusters spanning a range of ages are present. During its lifetime a cluster fades due to stellar evolution \\citep[e.g.][]{schulz02} and loses stars due to a variety of disruption effects (e.g. \\citet{baumgardtmakino03,lamers05,gieles06c,gieles06b}). Nevertheless, using assumptions for the cluster formation rate, the cluster IMF and cluster disruption parameters, one can construct expected LFs and subsequently compare these with observed LFs \\citep{gieles05a}.  In three galaxies the LF of the rich star cluster systems appears to be better described by a double power law, i.e. two distinct parts, both described by a power law, which are separated by a bend at some absolute magnitude. This bend magnitude differs from galaxy to galaxy. \\citet{whitmore99} found for the ``Antennae'' (NGC 4038/4039) a bend at $M_V \\simeq -10$ (standard UBVI), with on the faint side a shallower slope ($\\sim -1.7$) than on the bright side ($\\sim -2.7$). For M51, \\citet{gieles06a} found hints for a double power law behavior of the LF for a sub-sample of clusters in the inner part of the disk, which were later confirmed by \\citet{gieles06} for a 5 times larger sample, based on the same data as used in this study. The slopes on both sides are similar to the slopes found by \\citet{whitmore99}, but the bend occurs about 1.6 magnitudes fainter. With the distance modulus of the Antennae newly derived by \\citet{saviane08}, which differs by 1.7 magnitudes from the older determination by \\citet{whitmoreschweizer95}, the difference in bend magnitudes between M51 and the Antennae reduces to 0.1 magnitude. In \\citet{gieles06a} it is shown that the LF of the star cluster population of NGC 6946 is also better fit with a double power law, with parameters comparable to M51. Note that the slopes at the faint end of the LF for all these galaxies are similar to the slopes found for populations with a single power law distribution. If the observed or formed number of clusters is not high enough to sample the LF up to the point where the bend occurs, the bend will not be detected. The galaxies for which a double power law LF for their clusters has been reported are all typified by a high star formation rate.  Whereas the bend in the LF of the ``Antennae'' was originally interpreted as being caused by a bend in the mass function, \\citet{zhangfall99} already mentioned a truncated mass function as a possible explanation for the double power law behavior. \\citet{gieles06a} have shown with analytic cluster population models that such a bend can occur if the highest occurring cluster mass is no longer determined by the size of the sample (i.e. the highest mass present in a cluster population is determined statistically).  For example, in the case of the LMC and SMC \\citep{hunter03}, the cluster IMF does not show signs of an upper mass limit, and therefore sampling this mass function results in a maximum mass in the sample that is determined by the number of clusters and the slope of the power law cluster IMF. This is argued to be generally true by \\citet{weidner04}, who argue that the maximum cluster mass is a function of the star formation rate only. On the other hand, if there exists a \\textit{physical upper mass limit} for star clusters and the cluster IMF is sampled up to this limit the LF will appear in a shape more closely resembling a double power law \\citep{larsen06}.  If the cluster IMF is truncated, the physical reason for the maximum mass may have its imprints on the LF of subsets of the population. Variations of the maximum mass with galactocentric distance or surrounding ambient densities can hold important clues about the formation and evolution of star clusters.  Using deep observations by \\textit{HST} using the \\textit{Advanced Camera for Surveys (ACS)} of the interacting, face-on, spiral galaxy M51, which cover a large fraction of the disk of the galaxy and have a resolution of 0$\\farcs$05, we are able to select a sample of thousands of clusters \\citep{scheepmaker06}. Those clusters range in luminosity from the magnitude of single bright stars (distinguished on the basis of their spatial extent) to the brightest clusters found in spiral galaxies. The size of this sample gives us the opportunity to divide the clusters in smaller subsets and still have enough statistics for the obtained LFs. In this way we are able to study trends with e.g. galactocentric distance and the spiral arm structure of the galaxy. The same observations are used by \\citet{hwanglee08} to determine the LF of luminous clusters. They find a single power law, but only include clusters brighter than the bend magnitude, as found by \\citet{gieles06a} and \\citet{gieles06}, with a power law slope which agrees with the bright end slope of \\citet{gieles06a, gieles06}.  In \\S~\\ref{sec:obs} we describe the data, source selection and photometry.  The luminosity function of clusters in M51 will be presented in \\S~\\ref{sec:lumfunc}, for the whole population as well as for certain subsets. In \\S~\\ref{sec:modcomp} we perform Monte Carlo simulations of star cluster populations in order to investigate the dependence of obtained LF parameters on maximum cluster mass, disruption time and cluster formation history. There, we also draw qualitative conclusions about the cluster IMF and cluster disruption parameters of M51. The variation of the bend luminosity with position on the disk is further investigated in \\S~\\ref{sec:discussion} and our conclusions are summarized in \\S~\\ref{sec:conclusions}. ",
        "conclusions": " \\begin{enumerate} \\item The bend in the LF occurs at \\textit{brighter} magnitudes, \\textit{closer} to the center of the galaxy. This may imply a higher possible maximum mass closer to the center of the galaxy. This conclusion is guided by the Monte Carlo simulations, which show that for given cluster formation history and disruption time, the bend magnitude is brighter for higher maximum masses. Varying the cluster formation history, and/or cluster disruption across the disk makes this inference less straightforward. However, assuming a constant maximum mass would require large variations in the disruption time as a function of galactocentric distance. \\item The magnitude of the bend in the LF varies slowly with background intensity. This variation may be solely due to the difference in average galactocentric distance between both background intensity regions. \\item The other deduction from the LF parameters concerns disruption parameters, which also depend on galactocentric distance. A higher surrounding density will be more destructive for a cluster, as encounters with massive objects, clouds and clusters, will be more frequent. Also, the tidal field of the galaxy is decreasing outwards. Due to the degeneracy of a decreasing upper mass limit and stronger disruption to flatten the faint end of the LF, the dependence of the disruption time on galactocentric distance can not be constrained. This degeneracy can be lifted if the ages and masses of the clusters can be measured. \\item The slope of the faint end side of the LF is \\textit{shallower, closer} to the center of the galaxy and supports the suggestion that star clusters in the center are more easily disrupted. The higher density, stronger tidal field and larger relative velocities between several massive components (clusters and clouds) make the disruption of star clusters easier in those parts of the galaxy \\citep{gieles06c, gieles06b}. \\item The slope of the faint end side of the LF is \\textit{shallower} in \\textit{high} background regions. This is counter-intuitive, as one expects the younger population in the spiral arms to have a faint-end LF slope more closely resembling the cluster IMF slope compared to the population in the interarm regions. Selection effects may play an important role here. Clusters have a high number density, causing confusion, blending and the rejection of clusters from the sample as a result of a very nearby neighbor. \\end{enumerate}  Whether galactocentric distance or background intensity is the driving parameter for the variations of the faint-end slopes and bend magnitudes is unclear, due to the correlation of background intensity with galactocentric distance."
    },
    "0806/0806.1611_arXiv.txt": {
        "abstract": "In this contribution we outline plans for identifying and characterising numerous young, low-mass stars within 150 pc of the Sun using the new SkyMapper telescope and Southern Sky Survey. We aim to learn more about the star formation history of the solar neighbourhood over the past 5--50 Myr, the dispersal processes involved, as well as testing pre-main sequence evolutionary models and the universality of the stellar Inital Mass Function. Searching for the dispersed halo of low-mass objects predicted to surround the $\\eta$ Chamaeleontis cluster will be one of the first goals of the project. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0806/0806.1427_arXiv.txt": {
        "abstract": "{} {In order to understand the first stages of planet formation, when tiny grains aggregate to form planetesimals, one needs to simultaneously model grain growth, vertical settling and radial migration of dust in protoplanetary disks. In this study, we implement an analytical prescription for grain growth into a 3D two-phase hydrodynamics code to understand its effects on the dust distribution in disks.} {Following the analytic derivation of Stepinski \\& Valageas (1997), which assumes that grains stick perfectly upon collision, we implement a convenient and fast method of following grain growth in our 3D, two-phase (gas+dust) SPH code. We then follow the evolution of the size and spatial distribution of a dust population in a classical T~Tauri star disk.} {We find that the grains go through various stages of growth due to the complex interplay between gas drag, dust dynamics, and growth. Grains initially grow rapidly as they settle to the mid-plane, then experience a fast radial migration with little growth through the bulk of the disk, and finally pile-up in the inner disk where they grow more efficiently. This results in a bimodal distribution of grain sizes. Using this simple prescription of grain growth, we find that grains reach decimetric sizes in $10^5$~years in the inner disk and survive the fast migration phase.} {} ",
        "introduction": "\\label{Intro}  The first steps of planet formation are governed by the build-up of planetesimals due to the dust coagulation in protoplanetary disks \\citep{Dominik2007}. Observational evidence for grain growth in disks is now common \\citep{Apai2004,Rodmann2006,Muzerolle2006,Lommen2007,Graham2007}. Grains must grow from sub-$\\mu$m sizes to planetesimal scale (kilometer size) objects in a fraction of the lifetime of the disk, which is estimated to be a few $10^7$ years \\citep{Haisch2001,Carpenter2005}. The timescales of grain growth, however, are unclear: some young disks show signatures of grain growth while old disks can show signatures of unprocessed grains and coeval disks can show a range of grain sizes and dust processing \\citep{Kessler-Silacci2006}.  Grains can grow via collisions and depending on their relative velocity and on their chemical and physical properties \\citep{Chokshi1993,Blum2006}, colliding grains can rebound, shatter or stick. Grains will settle vertically and migrate radially at different rates according to their size \\citep{Weidenschilling1977,Garaud2004,BF05}, leading to local density enhancements in the disk. Since grain growth is dependent on density, changes in the dust distribution will affect growth rates, which in turn will affect the dynamics of the dust \\citep{Weidenschilling1980,Haghi2005}. Therefore, growth, settling and migration need to be simulated together.  Various models have been developed to describe the grain growth process. One approach is to use the time-dependent Smoluchowski coagulation equation \\citep{Weidenschilling1980,Weidenschilling1997,SuttnerYorke2001, DullemondDominik2005,Tanaka2005,Nomura2006,Ciesla2007} which describes the number density evolution of particles of a given mass range. The numerical solution of the Smoluchowski equation is challenging. Another approach is to use an analytic expression for the grain growth rate as a function of local disk conditions \\citep{Stepinski1997,Haghi2005}. In this study, we use this second approach and implement the analytical prescription of \\citet[hereafter SV97]{Stepinski1997}\\defcitealias{Stepinski1997}{SV97} in our 3D, two-phase (gas+dust) hydrodynamics code and follow the evolution of the grain size distribution in a protoplanetary disk. We validate our method on an axisymmetric disk here, before applying it to non-axisymmetric complex problems in future work. ",
        "conclusions": "\\label{SectConclusion}  We have implemented a mechanism able to treat grain growth in protoplanetary disks via the analytical expression of \\citetalias{Stepinski1997} into our two-phase SPH code. We simulated for the first time the full 3D evolution of a typical T Tauri disk, following the simultaneous radial migration, vertical settling, and growth of solid particles. Their interplay is complex: dynamics affects grain growth by modifying local physical quantities such as density or relative velocity. Conversely, grain growth also acts on dust dynamics: where non-growing grains would either stay well mixed with the gas or settle and migrate according to their sizes \\citepalias{BF05}, growing grains will go through various stages and produce totally different spatial distributions. They initially grow rapidly as they settle to the mid-plane, then experience a fast radial migration with little growth through the bulk of the disk, and finally pile-up in the inner disk where they grow more efficiently. This results in a bimodal distribution of grain sizes, with the largest grains found in the denser inner disk, where growth is most efficient. The survival times of the solid particles are longer than previously found, which has an implication on planet formation.  We find that grains grow very quickly: they reach decimetric size in $10^5$~yr. This is in general agreement with the results of \\citet{DullemondDominik2005} where shattering is neglected. They used the different approach of solving the Smoluchowski equation to study the growth of settling, but non-migrating, dust and in $10^5$~yr formed grains of maximum sizes ranging from  1~cm to over 1~m depending on their model parameters. Similar to them, we also find that the small grains only survive in the very outer disk and are depleted too rapidly elsewhere to be consistent with infrared observations of disks, highlighting the importance of shattering.  In order to compute synthetic images from our simulations and compare them to the observations, one would have to assume that the collisional cascade resulting from the inclusion of shattering would produce a whole particle distribution from the maximum size at a given radius shown in Fig.~\\ref{FigSizeDist}d down to sub-$\\mu$m size, described by a quasi-steady power law as argued by \\citet{Garaud2007}.  The method we used to treat grain growth can easily be applied to other analytical prescriptions.  The development of a more detailed model is necessary for a realistic description of grain growth in protoplanetary disks. In addition to shattering, one needs to take into account other processes such as microscopic interactions between the grains, kinetic energy dissipation and grain porosity. This is the subject of a forthcoming paper."
    },
    "0806/0806.3108_arXiv.txt": {
        "abstract": "We discuss the possibility to observe  the products of dark matter annihilation that was going on in the early Universe. Of all the particles that could be generated by this process we consider only photons, as they are both uncharged and easily detectable. The earlier the Universe was, the higher the dark matter concentration $n$ and the annihilation rate (proportional to $n^2$) were. However, the emission from the very early Universe cannot reach us because of the opacity. The main part of the signal was generated at the moment the Universe had just become transparent for the photons produced by the annihilation. Thus, the dark matter annihilation in the early Universe should have created a sort of relic emission. We obtain its flux and the spectrum.  If weakly interacting massive particles (WIMPs) constitute dark matter, it is shown that we may expect  an extragalactic gamma-ray signal in the energy range 0.5 - 20~{MeV} with a maximum near 8~{MeV}. We show that an experimentally observed excess in the gamma-ray background at 0.5 - 20~{MeV} could be created by the relic WIMPs annihilation only if the dark matter structures in the universe had appeared before the universe became transparent for the annihilation products ($z \\simeq 300$). We discuss in more detail physical conditions whereby this interpretation could be possible. ",
        "introduction": "Though the cosmological measurements \\citep{cosmol} show that there must be approximately five times as much dark matter as all baryon one, its physical nature remains unknown. Now the most commonly used hypothesis is that it consists of some elementary particles generated in the early Universe (hereafter we will call them Dark Matter Particles, DMPs). These particles are uncharged and do not interact strongly; there are telling arguments to believe that they were cold ($\\upsilon\\ll c$) in the epoch when the relic radiation was generated. It is worth mentioning that a particle with suitable properties hasn't been discovered yet, in spite of no lack of theoretical candidates predicted by various quantum field theory models.  If the premise is true, the dark matter is a mixture of equal quantities of particles and antiparticles, and they must collide and annihilate wherever the dark matter is present. Experimental observation of such a process would give us some valuable information about the DMP nature. In this article, we consider the dark matter annihilation in the epoch near the relic radiation formation ($z\\sim 1000$). At that time, the average dark matter particle concentration $n\\propto (z+1)^3$ was nine orders higher than now and four orders higher than in our Galaxy near the Sun system. So we may expect that the annihilation, the rate of which is proportional to $n^2$, was very intensive in that epoch.  Of all the particles that can be generated by the dark matter annihilation we will consider only photons, as they are both uncharged and easily detectable. Uncharged particles do not interact with the magnetic field of the Galaxy, which allows one to measure the extragalactic background reliably enough.  The earlier the Universe was, the higher the DM density and the annihilation rate were. However, the emission from the very early Universe cannot reach us because of the opacity. The main part of the signal was generated at the moment the Universe had just become transparent for the photons produced by the annihilation. Later the dark matter density rapidly dropped, decreasing the signal. The moment (and its redshift) depends, of course, on the characteristic energy of the photons, in other words, on the nature of the dark matter.  Thus, the DM annihilation in the early Universe should have produced a sort of relic emission. We obtain its flux and the spectrum. On the one hand, such an emission can be detected in the spectrum of the extragalactic background. A distinguishing feature of this radiation should be its high isotropy. On the other hand, the absence of such a signal can impose a severe limitation on a dark matter model.  In calculations we assume that the dark matter is homogeneous. This is the simplest, but by no means the most natural supposition. Possible amplitude of the dark matter density perturbations and their influence on the annihilation signal will be discussed at the end of the article.  In the second part of the article, the case of the most popular dark matter candidate - Weakly Interacting Massive Particles (WIMPs) - is considered in more detail. We demonstrate that the relic signal from the WIMPs annihilation might have been already observed. ",
        "conclusions": "In order to fulfil the calculations, the cosmological parameters should be concretized. We use the following set (see \\citet{gorbrub} and references therein): $\\Omega_\\Lambda = 0.75$, $\\Omega_m = \\Omega_{DM} + \\Omega_b=0.25$, $\\Omega_b = 0.042$ (of course, $\\Omega_{\\Lambda}+\\Omega_{DM}+\\Omega_b=1$), the Hubble constant $H_0=2.4\\cdot 10^{-18} \\mbox{s}^{-1}$, the relic radiation temperature $2.725$~K, the baryon-photon ratio $\\eta\\equiv n_b/n_{ph}=6.1 \\cdot 10^{-10}$. We obtain the present baryon concentration $n^b_0 = 2.5 \\cdot 10^{-7}\\, \\mbox{cm}^{-3}$. The DMP concentration, with DMP mass taken as $M_{DMP}=50$~{GeV}, is $n_0=2.5 \\cdot 10^{-8}\\, \\mbox{cm}^{-3}$. Equation (\\ref{a22}) gives $\\langle\\sigma\\upsilon\\rangle\\simeq 2 \\cdot 10^{-26}\\: (\\mbox{cm}^3 /\\mbox{s})$.  The influence of the relic annihilation on the ionization history of the universe is negligible. Indeed, for the recombination epoch ($z\\simeq 1200$) we have:  $n^b \\simeq n^b_0 z^3= 4.3\\cdot 10^2 \\, \\mbox{cm}^{-3}$, $ n \\simeq n_0 z^3 =43\\, \\mbox{cm}^{-3}$. The number of annihilations in a unit volume per unit time is $\\dfrac12 \\langle\\sigma\\upsilon\\rangle\\: n^2 = 1.9\\cdot 10^{-23} {cm}^{-3} {s}^{-1}$. In the characteristic time (of the order of the hydrogen ion recombination time at that epoch {435}~{years}~$\\simeq 1.4\\cdot 10^{10}$~s \\citep{gorbrub}) we have $2.6 \\cdot 10^{-13}$ annihilations that produce $2.6 \\cdot 10^{-2}$~{eV} of energy. A hydrogen atom ionization requires $\\sim 14$~{eV}. So, even all the energy produced by the annihilation is enough to ionize only $4\\cdot 10^{-4}$~\\% of atoms.  We take $\\aleph=\\aleph(2 \\mbox{GeV}) =1.63 \\cdot 10^{-26}\\,\\mbox{cm}^2$ (\\ref{a23}). Then the constant $\\wp$ is equal to $1.0 \\cdot 10^{-4}$. \\begin{figure} \\resizebox{\\hsize}{!}{\\includegraphics[angle=270]{fig2.eps}} \\caption{The spectrum of the relic gamma-ray background, calculated according to (\\ref{a30}), (\\ref{a31}), (\\ref{a20}), and (\\ref{a29}). Separate contributions of the primary and the scattered photons are represented by the dashed and the dotted lines, respectively.} \\label{fig2} \\end{figure} The resulting spectrum of the relic gamma-ray background, obtained with the aid of (\\ref{a30}), (\\ref{a31}), (\\ref{a20}), and (\\ref{a29}), is represented in Figure~\\ref{fig2}. Instead of the photon flux $\\tilde Q$ we have plotted the quantity $\\varepsilon^2 * \\mbox{flux}$ (that is $\\varepsilon^2 \\cdot \\tilde Q$), which is traditionally used in experimental data picturing. The primary (\\ref{a20}) and the secondary (\\ref{a29}) photon separate contributions are represented by the dashed and the dotted lines, respectively. The contribution of secondary photons is small in these coordinates, even though their total number is very large. Our main conclusions are that the spectrum grows up to $\\sim 8$~{MeV}, and the bulk of the signal lies in the range from 0.5 to 20~{MeV}. Characteristic redshift of the relic gamma-rays can be easily calculated. According to (\\ref{a20}), the quantity $\\varepsilon^2 \\cdot \\tilde Q$ has its maximum at $\\frac{\\varepsilon}{\\beta} = a = (2 \\wp)^{2/3}$. According to (\\ref{b1}), it corresponds to $z \\simeq 300$.  The cosmic gamma-ray background reportedly (see, for instance, \\citet{exp1, ahn2005a, ahn2005b, rasera2006, strigari2005}) has a peculiarity in the energy range 0.5 - 20~{MeV}. A ledge-like feature is visible in the extragalactic gamma-ray spectrum (Fig.~3 in \\citet{strong2004a}). The photon index here is markedly distinct from those of the softer or harder parts of the spectrum \\citep{sreekumar1998, strong2004a, weidenspointner2000}, indicating its different origin.  Moreover, this spectral band can be formed neither by too soft emission of normal active galactic nuclei, nor by too hard blazar-type AGNs contribution (see \\citet{exp1} and references therein). Attempts to consider the nuclear-decay gamma rays from Type Ia supernovae as the source have not been successful: the flux expected from the supernovae is several times weaker than the observed \\citep{ahn05, rasera2006, strigari2005}. It might be well to point out that the precise determination of the excess boundaries and intensity is model-dependent, and the literature values vary considerably \\citep{exp1, ahn2005a, ahn2005b, rasera2006}. In any case, however, the excess becomes apparent near 0.5~{MeV} and disappears at the energies $\\gtrsim$~20~{MeV}. One can see that the energy range of the feature corresponds closely to the interval characteristic for the relic gamma emission from the WIMPs annihilation. This coincidence looks promising when it is considered that the WIMP is now one of the most probable dark matter candidates.  At the same time, the relic gamma emission predicted by equations (\\ref{a30}), (\\ref{a31}), (\\ref{a20}), and (\\ref{a29}), is approximately five orders fainter than the observed feature (as we can see in Figure~\\ref{fig2}, the relic emission near the maximum at 10~{MeV} has $\\varepsilon^2 * \\mbox{flux}\\simeq 10^{-8}\\, \\mbox{MeV} /(\\mbox{sr}\\, \\mbox{cm}^2\\, \\mbox{s})$ while the total extragalactic gamma-ray background at 10~{MeV} has $\\varepsilon^2 * \\mbox{flux}\\simeq 2\\cdot 10^{-3}\\, \\mbox{MeV} /(\\mbox{sr}\\, \\mbox{cm}^2\\, \\mbox{s})$ \\citep{sreekumar1998, rasera2006}). This discrepancy might result from inapplicability of the assumption of homogeneous dark matter distribution. In fact it cannot be so. The modern structure of the Universe appeared from some initial perturbations that had already existed, beyond any doubt, in the epoch $z \\sim 300$. According to WMAP measurements \\citep{wmap}, in the recombination epoch $z \\simeq 1100 \\div 1400$ relative variations of the baryonic matter density were of the order of $10^{-5}$. Dark matter perturbations could be much more intensive, they were not suppressed by the radiation pressure in the pre-recombination epoch. Moreover, they must have been significantly stronger (not less than $10^{-3}$) to explain the modern Universe structure \\citep{gurevichzybin}. Since the recombination happened in the matter-dominated epoch, the perturbations rapidly grew and at the moment $z \\sim 300$ could attain very big amplitude. The presence of density inhomogeneities does not affect the spectrum of the annihilation signal but increases its intensity. This effect is usually described by the quantity $$ C\\equiv \\dfrac{<\\rho^2>}{<\\rho>^2} $$ that appears as a multiplier in the expression for the intensity (see, for instance, \\citet{ahn2005a}). Of course, $C$ is a function of $z$.  This brings up two points: first, is it possible that the structure growth in the early universe proceeded so fast that $C$ was as large as $10^5 - 10^6$ by the moment $z=300$? Second, if an intensive structure formation took place at some moment $\\tilde z < 300$, the coefficient $C(z)$ could grow so rapidly that it far outweighed the signal diminution owing to average density decrease. As a result a hard tail or even a secondary hard maximum on the energy (2-4)/$\\tilde z$~{GeV} can appear in the spectrum of the relic emission, which is not observed. Unfortunately, the theory of evolution of dark matter perturbations is still far from accurate. They evolved from some primordial fluctuations existed in the very early universe. While the universe was radiation-dominated their growth was slow. The smallest perturbations were destroyed by free-streaming (for the instance of neutralinos of the mass $\\sim 100$~{GeV} this limit is estimated as $10^{-12}-10^{-6} M_\\odot$ \\citep{10-12, 10-6}). When the universe transits into the matter-dominated stage, the perturbations start to grow rapidly, and eventually they become nonlinear and collapse. As this takes place, the smallest clumps collapse the first (at a time we denote as $\\tilde z$) \\citep{gurevichzybin}. The overwhelming majority of these small clumps originated at that moment were subsequently destroyed by tidal interaction with larger clumps originated later. But up to now it is the small clumps formed at the epoch $\\tilde z$ that makes the main contribution to the annihilation rate of the dark matter since they are the densest \\citep{berezinsky2003, gurzybsir}. At the moment $\\tilde z$ function $C(z)$ underwent a rapid increase from a value close to $1$ to a very big value. In what follows creation of the larger structures was accompanied by smaller clump destruction, and $C(z)$ changes more smoothly. In order to be allowed to suggest that the extragalactic gamma-ray 0.5 - 20~{MeV} excess is related to the neutralino annihilation we must assume that $\\tilde z > 300$, i.e. the first structures started to form before the universe became transparent for the photons produced by the annihilation.  How realistic is such an assumption? Unfortunately, present estimations are very vague. Even the minimal possible clump mass for the neutralino dark matter is determined extremely uncertainly (from $10^{-12} M_\\odot$ \\citep{10-12} to $10^{-6} M_\\odot$ \\citep{10-6}, to say nothing of the clumps density profile and the moment when the fluctuations become nonlinear. Experimental data as well as numerical simulations essentially cover the range of very large structures $10^{15}-10^6 M_\\odot$ (for instance, WMAP can observe only the biggest perturbations with masses corresponding to a cluster of galaxies $\\sim 10^{15} M_\\odot$ \\citep{wmap}). Properties of smaller clumps are usually obtained by approximation \\citep{ahn2005a}. However, in order to obtain any parameters for the tiny clumps of mass $10^{-6}-10^{-12} M_\\odot$ one has to extend the results by $12$-$18$ orders. Another source of uncertainties is the spectrum of primordial fluctuations. Usually it is deemed that it has flat Harrison-Zeldovich shape. In this case the moment of the first intensive clump creation is estimated as $\\tilde z \\simeq 80$, though some individual clumps collapsed much earlier \\citep{green2005}. If such a scenario was indeed realized, the interpretation of 0.5 - 20~{MeV} excess as a result of neutralino annihilation is out of the question: otherwise a strong maximum at the energy 4/$\\tilde z$~{GeV}~$\\simeq$~50~{MeV} would appear (which contradicts to the observations).  On the other hand, let us suppose that the spectrum of primordial fluctuations is not exactly flat, and the intensity of the fluctuations slightly builds up as their scale decreases. In the matter-dominated stage the perturbations grow as $\\delta\\rho/\\rho\\propto t^{2/3}$ and $a(t)\\propto t^{2/3}$, therefore $\\delta\\rho/\\rho\\propto a$ \\citep{gurzybsir}. If we assume that the small-scale perturbations left the radiation-dominated stage with amplitudes more than expected from the Harrison-Zeldovich spectrum, they collapsed earlier (in so doing we imply that the amplitudes of large-scale perturbations are fixed in such a way as they reproduce the observed large-scale structure of the Universe). Considering that the largest and the smallest clump mass scales differ by more than $20$ orders, even a small tilt of the spectrum of primordial fluctuations with respect to the Harrison-Zeldovich shape can be sufficient. A kindred scenario was considered, for instance, by \\citet{gurzybsir}. In the context of the model examined by the authors small clumps are extremely dense and collapse just after the universe transition to the matter-dominated epoch ($\\tilde z>1500$). It is worthy of note that the situation when $\\tilde z > 300$ can appear in much softer scenarios than those similar to \\citet{gurzybsir}.  Let us give a more specific form to the above reasoning. The spectrum of primordial fluctuations is usually considered to have a power-law shape: \\begin{equation} \\label{c1} |\\delta^{2}({\\bf k})|\\sim k^n \\end{equation} The case when $n=1$ corresponds to the Harrison-Zeldovich spectrum. The moment when a perturbation mode becomes nonlinear and collapses is determined by the spectrum function \\citep{gurzybsir} \\begin{equation} \\label{c2} \\Gamma(k)\\propto\\dfrac{k^3 |\\delta^{2}({\\bf k})|}{1+(k/k_{eq})^4} \\end{equation} (in \\citet{gurzybsir} $\\Gamma(k)$ is symbolized by $F(k)$). For all the perturbations we consider $k\\gg k_{eq}$, and \\begin{equation} \\label{c3} \\Gamma(k)\\propto\\dfrac{|\\delta^{2}({\\bf k})|}{k}\\propto k^{n-1} \\end{equation} The moment $a^*=\\frac{1}{1+z^*}$ when clumps collapse is \\begin{equation} \\label{c4} a^*\\propto\\dfrac{1}{\\sqrt{\\Gamma(k)}}\\propto k^{-\\frac{n-1}{2}} \\end{equation} which correlates with the results of \\citet{bullock}. In the case of Harrison-Zeldovich spectrum $\\Gamma(k)$ is flat; the structures of various scales appear almost simultaneously and have similar densities. If $n>1$, smaller clumps appear much earlier and have much higher density (since the universe density is much higher at the moment they collapse).  We can now construct a toy model of structure formation with the help of an approach used in \\citep{ahn2005a, bullock}. We take $|\\delta^{2}({\\bf k})|\\sim k^2$, i.e. $n=2$, and the minimal possible clump mass $M_{min}=10^{-7} M_\\odot$. Since the clamp mass relating to a perturbation mode is proportional to $k^{-3}$, we obtain from (\\ref{c4}): \\begin{equation} \\label{c41} a^*= \\left(\\dfrac{M}{M_{max}}\\right)^{\\frac{n-1}{6}} = \\left(\\dfrac{M}{M_{max}}\\right)^{\\frac{1}{6}} \\end{equation} Here we have introduced the maximum perturbation mass $M_{max}$ that collapses at $a=1$. This equation is consistent with \\citet{bullock}, where it was accepted $M_{max}\\simeq 1.5 \\cdot 10^{13} h^{-1} M_\\odot$ ($h=0.7$). Then the first clumps in our model appear at $\\tilde a=\\frac{1}{2450}$, i.e $\\tilde z = 2450$. We can rewrite (\\ref{c4}) as \\begin{equation} \\label{c8} a^*= \\tilde a\\cdot\\left(\\dfrac{M}{M_{min}}\\right)^{\\frac{1}{6}} \\end{equation} In accordance with \\citet{ahn2005a} we can represent the boost factor $C(z)$ as a product of three multipliers: \\begin{equation} \\label{c5} C(z) =\\Delta(z) \\cdot F_{coll}(z) \\cdot [C^{halo}] \\end{equation} (see all the details of the model and the notation in \\citep{ahn2005a, bullock}). For a flat $\\Lambda$CDM universe $\\Delta(z)=(18\\pi^2+82x-39 x^2)/(x+1)$, where $x=(\\rho_m/\\rho_{cr})-1$ and $\\rho_m$, $\\rho_{cr}$ are the matter and the critical universe densities at given $z$.  The matter fraction collapsed into cosmological halos is: \\begin{equation} \\label{c6} F_{coll}(z) = \\int_{M_{min}}^{M^*}\\frac{dn}{dM} M dM/\\rho_0 \\end{equation} Hereafter $M^*$ is the maximum having collapsed at given $z$. Theoretical model \\citep{berezinsky2003} and numerical simulations \\citep{diemand} give for the differential number density of clumps in the comoving frame of reference $\\frac{dn}{dM}\\propto M^{-2}$. Substituting it in (\\ref{c6}) and using (\\ref{c8}), we obtain: \\begin{equation} \\label{c7} F_{coll}(z) = \\frac{F^0_{coll}}{46.8} \\ln\\left(\\dfrac{M^*}{M_{min}}\\right)=\\frac{F^0_{coll}}{7.8} \\ln\\left(\\dfrac{a}{\\tilde a}\\right) \\end{equation} where $F^0_{coll}\\equiv F_{coll}(a=1)$. We will adopt the value obtained by \\citet{ahn2005a} $F^0_{coll}\\simeq 0.8$.  The factor $[C^{halo}]=\\int C^{halo} M \\frac{dn}{dM} dM/ \\int dM \\frac{dn}{dM}M$ represents the \"halo clumping\". If we adopt the Navarro-Frenk-White clump profile than \\begin{equation} \\label{c9} C^{halo}=\\frac{c_{vir}^{3}(1-1/(1+c_{vir})^3)}{9(\\ln(1+c_{vir})-c_{vir}/(1+c_{vir}))^2} \\end{equation} In order to describe the halo concentration parameter $c_{vir}$ evolution we will use an equation from \\citet{ahn2005a, bullock} with $K=8$, $F=0.01$ \\begin{equation} \\label{c10} c_{vir}(M,a)=K a \\left(F \\dfrac{M}{M^*}\\right)^{-1/6}= K \\dfrac{a}{\\tilde a} \\left( \\dfrac{M_{min}}{F M}\\right)^{1/6} \\end{equation} This formula was obtained as a fit of N-body simulations and is valid only for a limited range of $M$ and $z$ covered by them. We have to consider a much wider mass and red shift range, and (\\ref{c10}) gives too large $c_{vir}$ for the smallest clumps (for instance, for the minimal mass clumps at the present epoch $c_{vir}\\simeq 20000$). Such a huge value seems unlikely and indicates that equation (\\ref{c10}) should be corrected. Following \\citet{ahn2005a}, we assume that for any chosen clump mass the concentration parameter rises only up to $c_{vir}=100$, and after remains constant. \\begin{figure} \\resizebox{\\hsize}{!}{\\includegraphics[angle=270]{fig3.eps}} \\caption{The boost factor evolution with $z$.} \\label{fig3} \\end{figure}  Now the toy model is defined. The boost factor $C(z)$ evolution curve predicted by it is represented in Figure~\\ref{fig3}. We can make two principal conclusions. First of all, $\\tilde z = 2450$ and $C(z=300)\\simeq 1.5\\cdot 10^{5}$ in the model considered, i.e. the first structures appear very early ($\\tilde z \\gg 300$), and at the moment $z=300$ the boost factor is large enough to explain the discrepancy between the observed and the predicted signal intensities. Second, the boost factor $C(z)$ grows respectively slowly after $z=300$. From the moment when the universe became transparent for the annihilation photons up to the present moment it increases only on an order, while the signal without regard for the clumpiness rapidly falls as $z^{2.5}$ owing to the dark matter density decrease (\\ref{a12}). It means that almost all the signal appears at the moment $z\\simeq 300$ and the resulting spectrum has neither a hard tail nor a secondary harder maximum. In closing we remark that the foregoing structure formation scheme is no more than a toy model presented here only to illustrate that the structures in the universe could appear early, and the interpretation of the excess as a result of WIMP annihilation is principally possible.   The dark matter annihilation has already been invoked to explain the 0.5 - 20~{MeV} excess \\citep{ahn2005a, ahn2005b, rasera2006}. Since the photon spectrum is relatively soft, the authors introduced a low-mass dark matter candidate ($M_{DMP} < 100$~{MeV}). Its annihilation cross-section is sizable to provide the observed signal $\\langle\\sigma\\upsilon\\rangle\\simeq 2.5\\cdot 10^{-26}\\: (\\mbox{cm}^3 /\\mbox{s})$ \\citep{ahn2005a} that is at least no less than the typical weak interaction cross-section at this energy scale. Such a low-mass dark matter candidate with such a significant cross-section is now ruled out (by the accelerator experiments) in ordinary schemes like MSSM, but there is an interesting possibility to introduce it in more sophisticated scenarios \\citep{boehm1, boehm2}.  In our case, we can manage with usual heavy WIMP candidates. As we could see, the relic gamma-ray signal redshift is $z \\sim 300$. Originally, the photons had the energy $1 \\div 5$~{GeV} and they were produced by the annihilation of ordinary WIMP particles like the lightest neutralinos. On the other hand, if our interpretation is true, it counts in favour of a relatively light WIMP ($M_{DMP} \\sim 100$~{GeV}). If $M_{DMP} \\gg 100$~{GeV}, the typical energy of the producing photons is higher, and the feature in the spectrum must be harder than the observed one.  Thus, the critical point for the 0.5 - 20~{MeV} excess interpretation as a result of relic neutralino annihilation is determination of the moment $\\tilde z$ when the first structures in the universe appeared. If the structures had appeared before the universe became transparent for the annihilation products (i.e. $\\tilde z > 300$), then there are strong arguments to believe that the excess 0.5 - 20~{MeV} is created by the relic WIMPs annihilation. Above all, the characteristic energies of the spectra agree. Besides, the WIMP is now one of the most probable dark matter candidates, and the coincidence does not look random. The discrepancy of the predicted and observed signal intensities can be naturally explained by the nonuniform structure of the dark matter. If the first clumps appeared later ($\\tilde z < 300$), the excess undoubtedly could not be produced by neutralino annihilation. Further progress in the Universe structure formation understanding will be able to shed light on this problem."
    },
    "0806/0806.4591.txt": {
        "abstract": "{}{}{}{}{} % 5 {} token are mandatory  \\abstract % context heading {Asymptotic Giant Branch (AGB) phases mark the end of the evolution for Low- and Intermediate-Mass Stars. Our understanding of the mechanisms through which they eject the envelope and our assessment of their contribution to the mass return to the Interstellar Medium and to the chemical evolution of Galaxies are hampered by poor knowledge of their Luminosities and mass loss rates, both for C-rich and for O-rich sources.} % aims heading {We plan to establish criteria permitting a more quantitative determination of luminosities (and subsequently of mass loss rates) for the various types of AGB stars on the basis of infrared fluxes. In this paper, in particular, we concentrate on O-rich and s-element-rich MS, S stars and include a small sample of SC stars.} % methods heading {We reanalyze the absolute bolometric magnitudes and colors of MS, S, SC stars on the basis of a sample of intrinsic (single) and extrinsic (binary) long period variables. We derive bolometric corrections as a function of near- and mid-infrared colors, adopting as references a group of stars for which the Spectral Energy Distribution could be reconstructed in detail over a large wavelength range. We determine the absolute HR diagrams, and compare luminosities and colors of S-type giants with those, previously derived, of C-rich AGB stars. Luminosity estimates are also verified on the basis of existing Period-Luminosity relations valid for O-rich Miras.} %results heading {S star bolometric luminosities are almost indistinguishable from those of C-rich AGB stars. On the contrary, their circumstellar envelopes are thinner and less opaque. Despite this last property the IR wavelengths remain dominant, with the bluest stars having their maximum emission in the H or K(short) bands. Near-to-Mid infrared color differences are in any case smaller than for C stars. Based on Period-Luminosity relations for O-rich Miras and on Magnitude-color relations for the same variables we show how approximate distances (hence intrinsic parameters) for sources of so far unknown parallax can be inferred. We argue that most of the sources have a rather small mass ($<$ 2 $M_{\\odot}$); dredge-up might then be not effective enough to let the C/O ratio exceed unity.} {} ",
        "introduction": "}  The Asymptotic Giant Branch phases (hereafter AGB) represent the second ascent along the Red Giant Branch, occurring after the exhaustion of core He burning for all stars between $\\sim$~0.8 and 8.0\\,M$_{\\sun}$. In these evolutionary stages, stars are powered by two nuclear shells, burning H and He alternatively. In particular, in the final 1-2 Myr of the AGB, the He-burning shell remains mainly quiescent, if not for recurrent explosive ignitions during which a lot of C (from 20 to 25\\% by mass) is produced and spread over the whole He-rich layer, in short phases of convective mixing (the so-called {\\it thermal pulses}). Convective penetration of the envelope follows, in repeated episodes collectively called ''the third dredge-up'', and carries the new carbon to the surface, together with other nucleosynthesis products, in particular s-elements generated by efficient neutron captures \\citep{busso99}.  %\\begin{landscape} \\begin{table*}[t!] \\caption{Sample A $-$ First part. Spectral Type is from the GCVS catalogue whenever possible, otherwise it is obtained from the SIMBAD Astronomical Database. }             % title of Table \\label{table:1}      % is used to refer this table in the text \\centering                          % used for centering table \\begin{tabular}{c c c c c c}        % centered columns (4 columns) \\hline \\hline IRAS    &   Other   &   Stephenson  &   Coordinates &   Spectral Type   &   Var. Type   \\\\ name    &   name    &   name    &   ICRS    &       &   (GCVS)  \\\\ \\hline \\hline 01159+7220  &   \\object{S Cas}  &   CSS 28  &   01 19 41.97 +72 36 39.3 &   S3,4e$-$S5,8e   &   M   \\\\ 19126$-$0708    &   \\object{W Aql}  &   CSS 1115    &   19 15 23.44 $-$07 02 49.9   &   S3,9e$-$S6,9e   &   M   \\\\ 19354+5005  &   \\object{R Cyg}  &   CSS 1150    &   19 36 49.381 +50 11 59.46   &   S2.5,9e$-$S6,9e(Tc) &   M   \\\\ 19486+3247  &   \\object{chi Cyg}    &   CSS 1165    &   19 50 33.9220 +32 54 50.610     &   S6,2e$-$S10,4e/MSe  &   M   \\\\ 23595$-$1457    &   \\object{W Cet}  &   CSS 1346    &   00 02 07.3891 $-$14 40 33.065   &   S6,3e$-$S9,2e   &   M   \\\\ 22196$-$4612    &   \\object{pi1 Gru}    &   CSS 1294    &   22 22 44.2053 $-$45 56 52.598   &   S5,7e   &   SRB \\\\ 20026+3640  &   \\object{AA Cyg} &   CSS 1188    &   20 04 27.6055 +36 49 00.465 &   S7,5$-$S7.5,6(MpTc) &   SRB \\\\ 20120$-$4433    &   \\object{RZ Sgr} &   CSS 1196    &   20 15 28.4049 $-$44 24 37.480   &   S4,4ep  &   SRB \\\\ 03452+5301  &   \\object{WX Cam} &   CSS 82  &   03 49 03.77 +53 10 59.2 &   S5,8    &   LB  \\\\ \\hline 23070+0824  &   \\object{GZ Peg} &   CSS 1322    &   23 09 31.4570 +08 40 37.778 &   M4SIII  &   SRA \\\\ 15492+4837  &   \\object{ST Her} &   CSS 903 &   15 50 46.6248 +48 28 58.856 &   M6$-$7IIIaS &   SRB \\\\ 00192$-$2020    &   \\object{T Cet}  &   CSS 8   &   00 21 46.2737 $-$20 03 28.885   &   M5$-$6SIIe  &   SRC \\\\ 05374+3153  &   \\object{NO Aur} &   CSS 149 &   05 40 42.0504 +31 55 14.187 &   M2SIab  &   LC  \\\\ 22476+4047  &   \\object{RX Lac} &   CSS 1308    &   22 49 56.8992 +41 03 04.312     &   M7.5Se  &   SRB \\\\ \\hline 00213+3817  &   \\object{R And}  &   CSS 9   &   00 24 01.9469 +38 34 37.328 &   S3,5e$-$S8,8e/M7e   &   M   \\\\ 22521+1640  &   \\object{HR Peg} &   CSS 1315    &   22 54 35.6272 +16 56 30.601     &   S5,1/M4 &   SRB \\\\ 17553+4521  &   \\object{OP Her} &   $-$ &   17 56 48.5274 +45 21 03.063 &   M5IIb$-$IIIa/S  &   SRB \\\\ \\hline 13372$-$7136    &   \\object{LY Mus} &   CSS 826 &   13 41 13.5883 $-$71 52 05.767   &   M4III   &   LB  \\\\ 18058$-$3658    &   $-$ &   \\object{CSS 1023}   &   18 09 17.1853 $-$36 57 57.614   &   M2II$-$III  &   $-$ \\\\ \\hline 19111+2555  &   \\object{S Lyr}  &   CSS 1112    &   19 13 11.79 +26 00 28.3 &   SCe &   M   \\\\ \\hline 15194$-$5115    &   \\object{II Lup} &   CSS 886 &   15 23 04.91 $-$51 25 59.0   &   C   &   M   \\\\ \\hline \\hline \\end{tabular} \\end{table*} %\\end{landscape}  %\\begin{landscape} \\begin{table*}[t!] \\caption{Sample A $-$ Second part.}             % title of Table \\label{table:2}      % is used to refer this table in the text \\centering                          % used for centering table \\begin{tabular}{c c c c c c c c c c c c}        % centered columns (4 columns) \\hline      \\hline Source  &   J   &   H   &   K   &   [8.8]   &   [9.8]   &   [11.7]  &   [12.5]  &   D   &   E   &   Mid$-$IR Data Origin    &   ISO \\\\ name    &   [Jy]    &   [Jy]    &   [Jy]    &   [Jy]    &   [Jy]    &   [Jy]    &   [Jy]    &   [Jy]    &   [Jy]    &       &   TDT Number  \\\\ \\hline      \\hline S Cas   &   56.6    &   168 &   244 &   232 &   316 &   304 &   281 &   235 &   185 &   ISO$-$SWS1  &   41602133    \\\\ W Aql   &   388 &   822 &   1113    &   969 &   1099    &   1043    &   789 &   625 &   488 &   ISO$-$SWS1  &   16402335    \\\\ R Cyg   &   200 &   288 &   302 &   77.3    &   87.8    &   79.3    &   69.2    &   52.8    &   29.7    &   ISO$-$SWS1  &   42201625    \\\\ chi Cyg &   1365    &   2823    &   3176    &   1408    &   1655    &   1579    &   1095    &   736 &   487 &   ISO$-$SWS1  &   15900437    \\\\ W Cet   &   74.1    &   113 &   101 &   12.7    &   12.4    &   10.5    &   10.9    &   8.3 &   4.5 &   ISO$-$SWS1  &   37802225    \\\\ pi1 Gru &   3080    &   $-$ &   5812    &   544 &   632 &   679 &   683 &   505 &   365 &   ISO$-$SWS1  &   34402039    \\\\ AA Cyg  &   238 &   389 &   375 &   42.6    &   42.5    &   42.9    &   36.8    &   27.5    &   16.9    &   ISO$-$SWS1  &   36401817    \\\\ RZ Sgr  &   139 &   206 &   190 &   25.0    &   26.2    &   25.7    &   24.5    &   21.8    &   21.0    &   ISO$-$SWS1  &   14100818    \\\\ WX Cam  &   43.5    &   79.3    &   87.6    &   8.1 &   7.9 &   8.4 &   16.6    &   9.3 &   4.0 &   ISO$-$SWS1  &   81002721    \\\\ \\hline GZ Peg  &   735 &   1100    &   965 &   88.6    &   78.9    &   62.4    &   56.8    &   40.9    &   21.0    &   ISO$-$SWS1  &   37600306    \\\\ ST Her  &   804 &   1162    &   1098    &   166 &   186 &   200 &   187 &   149 &   104 &   ISO$-$SWS1  &   41901305    \\\\ T Cet   &   1009    &   1596    &   1403    &   172 &   163 &   172 &   171 &   134 &   83.0    &   ISO$-$SWS1  &   55502308 $-$ 37801819   \\\\ NO Aur  &   226 &   362 &   273 &   33.2    &   41.6    &   46.3    &   36.1    &   23.3    &   17.2    &   ISO$-$SWS1  &   86603434    \\\\ RX Lac  &   465 &   737 &   681 &   91.0    &   85.4    &   81.9    &   68.6    &   50.3    &   28.8    &   ISO$-$SWS1  &   78200427    \\\\ \\hline R And   &   247 &   506 &   596 &   193 &   264 &   248 &   210 &   176 &   135 &   ISO$-$SWS1  &   40201723    \\\\ HR Peg  &   191 &   327 &   256 &   27.0    &   24.4    &   20.9    &   20.0    &   12.6    &   7.8 &   ISO$-$SWS1  &   37401910    \\\\ OP Her  &   416 &   731 &   556 &   63.5    &   58.1    &   46.4    &   37.8    &   30.9    &   16.0    &   ISO$-$SWS1  &   77800625    \\\\ \\hline LY Mus  &   221 &   334 &   300 &   28.4    &   24.7    &   19.7    &   16.4    &   12.9    &   7.6 &   ISO$-$SWS1  &   13201304    \\\\ CSS 1023    &   33.2    &   51.7    &   39.9    &   2.3 &   2.2 &   1.5 &   1.4 &   0.83    &   0.46    &   ISO$-$SWS1  &   14100603    \\\\ \\hline S Lyr   &   8.2 &   13.2    &   17.9    &   20.0    &   23.2    &   25.5    &   24.4    &   22.1    &   15.7    &   ISO$-$SWS1  &   52000546    \\\\ \\hline II Lup  &   3.7 &   25.5    &   99.0    &   860 &   852 &   860 &   681 &   506 &   391 &   ISO$-$SWS6  &   29700401    \\\\ \\hline      \\hline \\end{tabular} \\end{table*} %\\end{landscape}  %\\begin{landscape} \\begin{table*}[t!] \\caption{Sample A $-$ Third part. The indication I. $-$ E. (Intrinsic $-$ Extrinsic) is given according to the suggestions from \\citet{vaneck00} and \\citet{yang}. In discordant cases we prefer the choice by \\citet{vaneck00} and show the one of \\citet{yang} in parenthesis. For few sources, for which neither study offers a suggestion, we infer that they are extrinsic from their low intrinsic Luminosity: in the tables they are underlined.}             % title of Table \\label{table:3}      % is used to refer this table in the text \\centering                          % used for centering table \\begin{tabular}{c c c c c c c c}        % centered columns (4 columns) \\hline              \\hline Source  &   Var. Type   &   Period  &   Distance    &   Min. $-$ Max.   &   Ref.    &   Bol. Magnitudes &   I. $-$ E.   \\\\ name    &   (GCVS)  &   (GCVS)  &   (kpc)   &   (kpc)   &   Distance    &   ISO Integration &       \\\\ \\hline              \\hline S Cas   &   M   &   612.43  &   0.85    &   $-$ &   P$-$L / this paper  &   $-$5.71 &   I   \\\\ W Aql   &   M   &   490.43  &   0.34    &   $-$ &   P$-$L / this paper  &   $-$5.44 &   I   \\\\ R Cyg   &   M   &   426.45  &   0.55    &   $-$ &   P$-$L / this paper  &   $-$5.42 &   $-$ \\\\ chi Cyg &   M   &   408.05  &   0.18    &   0.15 $-$ 0.22   &   Hip. / \\cite{vleu2007}  &   $-$5.39 &   I   \\\\ W Cet   &   M   &   351.31  &   0.83    &   $-$ &   P$-$L / this paper  &   $-$5.18 &   I   \\\\ pi1 Gru &   SRB &   150 &   0.16    &   0.15 $-$ 0.19   &   Hip. / \\cite{vleu2007}  &   $-$5.75 &   I   \\\\ AA Cyg  &   SRB &   212.7   &   $-$ &   $-$ &   $-$ &   $-$ &   I   \\\\ RZ Sgr  &   SRB &   223.2   &   $-$ &   $-$ &   $-$ &   $-$ &   I   \\\\ WX Cam  &   LB  &   $-$ &   $-$ &   $-$ &   $-$ &   $-$ &   I   \\\\ \\hline GZ Peg  &   SRA &   92.66   &   0.24    &   0.22 $-$ 0.26   &   Hip. / \\cite{vleu2007}  &   $-$5.02 &   E   \\\\ ST Her  &   SRB &   148 &   0.30    &   0.25 $-$ 0.36   &   Hip. / \\cite{vleu2007}  &   $-$5.64 &   I   \\\\ T Cet   &   SRC &   158.9   &   0.27    &   0.24 $-$ 0.31   &   Hip. / \\cite{vleu2007}  &   $-$5.63 &   I   \\\\ NO Aur  &   LC  &   $-$ &   0.60    &   0.47 $-$ 0.83   &   Hip. / \\cite{vleu2007}  &   $-$5.73 &   I   \\\\ RX Lac  &   SRB &   650 &   $-$ &   $-$ &   $-$ &   $-$ &   I   \\\\ \\hline R And   &   M   &   409.33  &   0.41    &   $-$ &   P$-$L / this paper  &   $-$5.19 &   I   \\\\ HR Peg  &   SRB &   50  &   0.41    &   0.36 $-$ 0.50   &   Hip. / \\cite{vleu2007}  &   $-$4.75 &   I   \\\\ OP Her  &   SRB &   120.5   &   0.30    &   0.27 $-$ 0.32   &   Hip. / \\cite{vleu2007}  &   $-$4.92 &   $-$ \\\\ \\hline LY Mus  &   LB  &   $-$ &   0.29    &   0.26 $-$ 0.33   &   Hip. / \\cite{vleu2007}  &   $-$4.12 &   E   \\\\ CSS 1023    &   $-$ &   $-$ &   $-$ &   $-$ &   $-$ &   $-$ &   E   \\\\ \\hline S Lyr   &   M   &   438.4   &   2.27    &   $-$ &   P$-$L / this paper  &   $-$5.50 &   I   \\\\ \\hline II Lup  &   M   &   $-$ &   0.59    &   $-$ &   \\citet{groenewegen02b}  &   $-$4.82 &   $-$ \\\\ \\hline              \\hline \\end{tabular} \\end{table*} %\\end{landscape}  %\\begin{landscape} \\begin{table*} \\caption{Sample B $-$ First part. Suggestions from \\citet{sl86} for Spectral Type: \\emph{1} RS Cnc: M6eIIIaS. \\emph{2} V1743 Cyg: M5IIIaS. \\emph{3} V1981 Cyg: M4IIIaS.}             % title of Table \\label{table:4}      % is used to refer this table in the text \\centering                          % used for centering table \\begin{tabular}{c c c c c c}        % centered columns (4 columns) \\hline \\hline IRAS    &   Other   &   Stephenson  &   Coordinates &   Spectral Type   &   Var. Type   \\\\ name    &   name    &   name    &   ICRS    &       &   (GCVS)  \\\\ \\hline \\hline 05199$-$0842    &   \\object{V1261 Ori}  &   CSS 133 &   05 22 18.6453 $-$08 39 58.034   &   S\u0085  &   Algol Type  \\\\ \\hline 04497+1410  &   \\object{omi Ori}    &   CSS 114 &   04 52 31.9621 +14 15 02.311 &   M3.2IIIaS   &   SRB \\\\ 10226+0902  &   \\object{DE Leo} &   $-$ &   10 25 15.1951 +08 47 05.441 &   M2IIIabS    &   SRB \\\\ 07245+4605  &   \\object{Y Lyn}  &   CSS 347 &   07 28 11.6109 +45 59 26.207 &   M6SIb$-$II  &   SRC \\\\ 07392+1419  &   \\object{NZ Gem} &   CSS 382 &   07 42 03.2185 +14 12 30.612 &   M3II$-$IIIS &   SR  \\\\ \\hline 06457+0535  &   \\object{V613 Mon}   &   CSS 260 &   06 48 22.2963 +05 32 30.050 &   M2/S5,1 &   SRB \\\\ 09076+3110  &   \\object{RS Cnc} &   CSS 589 &   09 10 38.7990 +30 57 47.300 &   M6eIb$-$II/S \\footnote{}    &   SRC \\\\ 03377+6303  &   \\object{BD Cam} &   CSS 79  &   03 42 09.3250 +63 13 00.501 &   S5,3/M4III  &   LB  \\\\ 07095+6853  &   \\object{AA Cam} &   CSS 312 &   07 14 52.0703 +68 48 15.380 &   M5/S    &   LB  \\\\ 13079$-$8931    &   \\object{BQ Oct} &   CSS 804 &   14 35 29.5001 $-$89 46 18.182   &   M4III/S5,1  &   LB  \\\\ \\hline 12272$-$4127    &   \\object{V928 Cen}   &   CSS 796 &   12 29 57.8871 $-$41 44 09.242   &   M2II$-$III  &   SRB \\\\ 19323+4909  &   \\object{V1743 Cyg}  &   $-$ &   19 33 41.6068 +49 15 44.347 &   M4.5III \\footnote{} &   SRB \\\\ $-$ &   \\object{V1981 Cyg}  &   $-$ &   21 02 24.1993 +44 47 27.528     &   M4s... \\footnote{}  &   SRB \\\\ 08214$-$3807    &   \\object{V436 Pup}   &   CSS 500 &   08 23 16.9344 $-$38 17 09.884   &   M1III   &   LB  \\\\ $-$ &   \\object{V2141 Cyg}  &   CSS 1254    &   20 57 53.1771 +44 47 17.336 &   M1  &   LB  \\\\ 12106$-$3350    &   \\object{V335 Hya}   &   $-$ &   12 13 12.9423 $-$34 07 30.981   &   M4III   &   LB  \\\\ 14510$-$6052    &   \\object{CR Cir} &   CSS 867 &   14 54 56.9389 $-$61 04 33.027   &   M2/M3II &   LC  \\\\ 16418$-$1359    &   $-$ &   \\object{CSS 937}    &   16 44 42.1936 $-$14 04 48.553   &   M1III   &   $-$ \\\\ \\hline 13136$-$4426    &   \\object{UY Cen} &   CSS 816 &   13 16 31.8300 $-$44 42 15.741   &   SC  &   SR  \\\\ \\hline 16425$-$1902    &   $-$ &   \\object{CSS 938}    &   16 45 30.1769 $-$19 08 12.939   &   K5II    &   $-$ \\\\ 20076+3331  &   $-$ &   \\object{CSS 1194}   &   20 09 32.9873 +33 40 53.851 &   K5III   &   $-$ \\\\ \\hline \\hline \\end{tabular} \\end{table*} %\\end{landscape}  %\\begin{landscape} \\begin{table*} \\caption{Sample B $-$ Second part.}             % title of Table \\label{table:5}      % is used to refer this table in the text \\centering                          % used for centering table \\begin{tabular}{c c c c c c c c c c c}        % centered columns (4 columns) \\hline      \\hline Source  &   J   &   H   &   K   &   [8.8]   &   [9.8]   &   [11.7]  &   [12.5]  &   D   &   E   &   Mid$-$IR Data Origin    \\\\ name    &   [Jy]    &   [Jy]    &   [Jy]    &   [Jy]    &   [Jy]    &   [Jy]    &   [Jy]    &   [Jy]    &   [Jy]    &       \\\\ \\hline      \\hline V1261 Ori   &   73.8    &   111 &   93.1    &   16.4    &   16.5    &   19.6    &   21.3    &   $-$ &   $-$ &   IRAS$-$LRS  \\\\ \\hline omi Ori &   1022    &   1587    &   1227    &   103 &   82.4    &   64.8    &   60.9    &   $-$ &   $-$ &   IRAS$-$LRS  \\\\ DE Leo  &   178 &   234 &   184 &   34.0    &   30.7    &   34.4    &   36.4    &   $-$ &   $-$ &   IRAS$-$LRS  \\\\ Y Lyn   &   876 &   1448    &   1256    &   132 &   150 &   121 &   107 &   $-$ &   $-$ &   IRAS$-$LRS  \\\\ NZ Gem  &   369 &   566 &   399 &   38.9    &   34.5    &   32.6    &   32.5    &   $-$ &   $-$ &   IRAS$-$LRS  \\\\ \\hline V613 Mon    &   49.6    &   74.3    &   68.9    &   5.2 &   $-$ &   $-$ &   3.3 &   2.0 &   $-$ &   MSX \\\\ RS Cnc  &   3065    &   4324    &   3742    &   512 &   693 &   493 &   436 &   $-$ &   $-$ &   TIRCAM2 \\\\ BD Cam  &   439 &   630 &   521 &   62.3    &   52.5    &   43.6    &   41.2    &   $-$ &   $-$ &   IRAS$-$LRS  \\\\ AA Cam  &   146 &   220 &   185 &   $-$ &   $-$ &   $-$ &   $-$ &   $-$ &   $-$ &   $-$ \\\\ BQ Oct  &   139 &   196 &   170 &   24.8    &   22.5    &   23.5    &   23.6    &   $-$ &   $-$ &   IRAS$-$LRS  \\\\ \\hline V928 Cen    &   196 &   267 &   225 &   25.4    &   21.9    &   20.3    &   20.5    &   $-$ &   $-$ &   IRAS$-$LRS  \\\\ V1743 Cyg   &   259 &   421 &   335 &   40.4    &   36.4    &   34.5    &   34.2    &   $-$ &   $-$ &   IRAS$-$LRS  \\\\ V1981 Cyg   &   197 &   279 &   245 &   28.9    &   $-$ &   $-$ &   16.3    &   10.8    &   4.7 &   MSX \\\\ V436 Pup    &   139 &   198 &   165 &   17.1    &   $-$ &   $-$ &   10.3    &   7.0 &   2.7 &   MSX \\\\ V2141 Cyg   &   147 &   216 &   195 &   20.8    &   $-$ &   $-$ &   12.2    &   8.2 &   3.5 &   MSX \\\\ V335 Hya    &   611 &   946 &   833 &   88.6    &   78.7    &   67.4    &   63.3    &   $-$ &   $-$ &   IRAS$-$LRS  \\\\ CR Cir  &   75.1    &   117 &   99.6    &   9.9 &   $-$ &   $-$ &   5.8 &   3.8 &   $-$ &   MSX \\\\ CSS 937 &   104 &   169 &   151 &   15.7    &   15.9    &   20.1    &   21.6    &   $-$ &   $-$ &   IRAS$-$LRS  \\\\ \\hline UY Cen  &   184 &   360 &   350 &   66.7    &   62.9    &   61.4    &   55.6    &   $-$ &   $-$ &   IRAS$-$LRS  \\\\ \\hline CSS 938 &   142 &   200 &   175 &   20.4    &   19.7    &   21.4    &   23.2    &   $-$ &   $-$ &   IRAS$-$LRS  \\\\ CSS 1194    &   31.1    &   41.5    &   37.0    &   2.9 &   $-$ &   $-$ &   2.4 &   0.96    &   $-$ &   MSX \\\\ \\hline      \\hline \\end{tabular} \\end{table*} %\\end{landscape}  %\\begin{landscape} \\begin{table*} \\caption{Sample B $-$ Third part.}             % title of Table \\label{table:6}      % is used to refer this table in the text \\centering                          % used for centering table \\begin{tabular}{c c c c c c c c}        % centered columns (4 columns) \\hline              \\hline Source  &   Var. Type   &   Period  &   Distance    &   Min. $-$ Max.   &   Ref.    &   Bol. Magnitudes &   I. $-$ E.   \\\\ name    &   (GCVS)  &   (GCVS)  &   (kpc)   &   (kpc)   &   Distance    &   Bol. Corrections    &       \\\\ \\hline              \\hline V1261 Ori   &   Algol Type  &   $-$ &   0.29    &   0.23 $-$ 0.38   &   Hip. / \\cite{vleu2007}  &   $-$2.71 &   E   \\\\ \\hline omi Ori &   SRB &   30  &   0.20    &   0.17 $-$ 0.23   &   Hip. / \\cite{vleu2007}  &   $-$4.89 &   I   \\\\ DE Leo  &   SRB &   $-$ &   0.31    &   0.27 $-$ 0.37   &   Hip. / \\cite{vleu2007}  &   $-$3.60 &   \\textbf{$\\underline{E}$}    \\\\ Y Lyn   &   SRC &   110 &   0.25    &   0.20 $-$ 0.33   &   Hip. / \\cite{vleu2007}  &   $-$5.33 &   I   \\\\ NZ Gem  &   SR  &   $-$ &   0.39    &   0.32 $-$ 0.51   &   Hip. / \\cite{vleu2007}  &   $-$5.06 &   E   \\\\ \\hline V613 Mon    &   SRB &   $-$ &   0.50    &   0.35 $-$ 0.84   &   Hip. / \\cite{vleu2007}  &   $-$3.76 &   E   \\\\ RS Cnc  &   SRC &   120 &   0.14    &   0.13 $-$ 0.15   &   Hip. / \\cite{vleu2007}  &   $-$5.21 &   I   \\\\ BD Cam  &   LB  &   $-$ &   0.16    &   0.15 $-$ 0.17   &   Hip. / \\cite{vleu2007}  &   $-$3.40 &   E   \\\\ AA Cam  &   LB  &   $-$ &   0.78    &   0.50 $-$ 1.82   &   Hip. / \\cite{vleu2007}  &   $-$ &   I   \\\\ BQ Oct  &   LB  &   $-$ &   0.49    &   0.41 $-$ 0.60   &   Hip. / \\cite{vleu2007}  &   $-$4.55 &   I   \\\\ \\hline V928 Cen    &   SRB &   $-$ &   0.23    &   0.21 $-$ 0.25   &   Hip. / \\cite{vleu2007}  &   $-$3.27 &   E   \\\\ V1743 Cyg   &   SRB &   40  &   0.41    &   0.38 $-$ 0.44   &   Hip. / \\cite{vleu2007}  &   $-$4.94 &   $-$ \\\\ V1981 Cyg   &   SRB &   $-$ &   0.30    &   0.27 $-$ 0.33   &   Hip. / \\cite{vleu2007}  &   $-$3.96 &   \\textbf{$\\underline{E}$}    \\\\ V436 Pup    &   LB  &   $-$ &   0.33    &   0.30 $-$ 0.38   &   Hip. / \\cite{vleu2007}  &   $-$3.75 &   E   \\\\ V2141 Cyg   &   LB  &   $-$ &   0.38    &   0.31 $-$ 0.51   &   Hip. / \\cite{vleu2007}  &   $-$4.24 &   \\textbf{$\\underline{E}$}    \\\\ V335 Hya    &   LB  &   $-$ &   0.36    &   0.31 $-$ 0.43   &   Hip. / \\cite{vleu2007}  &   $-$5.68 &   $-$ \\\\ CR Cir  &   LC  &   $-$ &   0.31    &   0.25 $-$ 0.42   &   Hip. / \\cite{vleu2007}  &   $-$3.08 &   E   \\\\ CSS 937 &   $-$ &   $-$ &   0.42    &   0.31 $-$ 0.66   &   Hip. / \\cite{vleu2007}  &   $-$4.11 &   E   \\\\ \\hline UY Cen  &   SR  &   114.6   &   0.69    &   0.47 $-$ 1.33   &   Hip. / \\cite{vleu2007}  &   $-$6.05 &   I   \\\\ \\hline CSS 938 &   $-$ &   $-$ &   0.25    &   0.20 $-$ 0.33   &   Hip. / \\cite{vleu2007}  &   $-$3.15 &   E   \\\\ CSS 1194    &   $-$ &   $-$ &   0.36    &   0.30 $-$ 0.45   &   Hip. / \\cite{vleu2007}  &   $-$2.34 &   \\textbf{$\\underline{E}$}    \\\\ \\hline              \\hline \\end{tabular} \\end{table*} %\\end{landscape}  Due to the above phenomena, the atmospheres of AGB stars are characterized by an increasing enrichment of $^{12}$C (up to C/O $>$ 1 by number, in which case we speak of C stars) and of s-process nuclei, in particular revealing the recent nucleosynthesis through the short-lived $^{99}$Tc \\citep{merrill}. Sometimes Tc itself offers actually the only real evidence of ongoing s-processing in stars that do not show other remarkable chemical anomalies \\citep{utt07}. When instead the third dredge-up process is efficient enough, changes in the photospheric abundances of other s-elements begin to occur, first of all for Zr, which has various isotopes on the main s-process path. In such cases the appearance of ZrO bands in the spectra (at wavelengths 464.1, 462.0, 530.4, 537.9, and 555.1 nm) tells us that the star, although still richer in O than in C, is mixing to the surface the products of shell-He burning. The cool giants presenting these signatures are called MS and S stars (the second group showing more prominent features). They also have a C/O abundance ratio by number higher than in the Sun, but lower than unity. There is still some confusion about the exact values of the C/O ratios in MS and S stars, which is mainly induced by the remaining uncertainty in the calibrating solar oxygen abundance. If one excludes for this calibration the recent, still debated suggestion \\citep{apla} and adopts instead the previous more traditional reference \\citep{ag89}, then MS stars are found typically at C/O = 0.5 to 0.7, and S stars are found above this range and up to more that 0.95. Around the border between O-rich and C-rich giants, the so-called SC stars represent a rare, but important, transition group.  In some cases, red giants showing s-element enhancement are Tc-poor. This is a clear indication that a sufficiently long time interval has passed since the production of neutron capture nuclei, so that Tc has decayed. In general, this is the case when the nucleosynthesis phenomena occurred not in the same star we see today, but in a more massive companion, which is now evolved to the white dwarf stage, and whose mass loss enriched in the past the photosphere of the observed object \\citep{busso01}. In these cases we speak more properly of an {\\it extrinsic} AGB star \\citep{smilamb}.  AGB stars lose mass very effectively, and their winds replenish the Interstellar Medium (ISM) guaranteeing up to 70\\% of the mass return from stars \\citep{sedlmayr94}. Before being dispersed over the Galaxy, the material thus lost forms cool envelopes \\citep{winters03} where dust grains condense \\citep{carciofi04}. These solid particles carry the elemental and isotopic composition generated by AGB nucleosynthesis; they have been found in ancient meteorites offering the possibility of high precision isotopic abundance measurements on matter coming from circumstellar environments \\citep[e.g.][]{zinner00}.  The cool AGB photospheres radiate most of their flux at red-infrared wavelengths. The infrared component of the Spectral Energy Distribution (SED) grows in importance (and in average wavelength) as far as the evolution proceeds, because of the increased extinction of the photospheric flux operated by dust, which then re-radiates at long wavelengths \\citep[see e.g.][]{habing96}. This correlation between extinction and evolutionary stage is however confused by the stars switching from Semiregular to Mira-type surface variability, which fact modulates the mass loss efficiency and hence the extinction properties. Due to these complicacies, large surveys of infrared (IR) observations play a fundamental role in studying luminosities and mass loss rates of AGB stars and in disentangling the variability and evolution effects \\citep[see e.g.][]{woodcohen,lebertre01,lebertre03,groenewegen02b,cioni03,omont03,olofsson03}.  Longstanding efforts have been devoted to describe the mass loss mechanisms, either with phenomenological models or with sophisticated hydrodynamical approaches \\citep{salpeter,knappmorris,winters03,wachter02,sandin3a,sandin3b}. Despite this, our quantitative knowledge of AGB winds is still poor and forces us to adopt parametric treatments, where observations play a crucial role in fixing the (otherwise free) parameters \\citep{wood03,olivier,wood04,andersen03}.  Similar problems affect the estimates of the stellar luminosity. Observations are hampered by the difficulties of measuring the distances for single, often obscured objects like AGB stars. On the other hand, the luminosities derived from full stellar evolutionary models are affected by the uncertainties in the choice of the mixing parameters (in particular of the extension of convective overshoot) and of surface atomic and molecular opacities \\citep{marigo}. Models adopting large overshoot parameters \\citep[see e.g.][]{izz07} derive large values for the mass dredged-up after each thermal pulse, thus obtaining the surface enrichment in C and s-elements earlier and at a lower luminosity than models based on the Schwarzschild's criterion can do.  %\\begin{landscape} \\begin{table*}[t!] \\caption{Sample C $-$ First part.}             % title of Table \\label{table:7}      % is used to refer this table in the text \\centering                          % used for centering table \\begin{tabular}{c c c c c c}        % centered columns (4 columns) \\hline \\hline IRAS    &   Other   &   Stephenson  &   Coordinates &   Spectral Type   &   Var. Type   \\\\ name    &   name    &   name    &   ICRS    &       &   (GCVS)  \\\\ \\hline \\hline 04352+6602  &   \\object{T Cam}  &   CSS 103 &   04 40 08.8768 +66 08 48.654     &   S4,7e$-$S8.5,8e &   M   \\\\ 06571+5524  &   \\object{R Lyn}  &   CSS 283 &   07 01 18.0093 +55 19 49.766     &   S2.5,5e$-$S6,8e:    &   M   \\\\ 07043+2246  &   \\object{R Gem}  &   CSS 307 &   07 07 21.2744 +22 42 12.736 &   S2,9e$-$S8,9e(Tc)   &   M   \\\\ 07092+0735  &   \\object{WX CMi} &   CSS 316 &   07 11 57.45 +07 29 59.3 &   Se  &   M   \\\\ 12417+6121  &   \\object{S UMa}  &   CSS 803 &   12 43 56.676 +61 05 35.51   &   S0,9e$-$S5,9e   &   M   \\\\ 15030$-$4116    &   \\object{GI Lup} &   CSS 872 &   15 06 16.31 $-$41 28 14.1   &   S7,8e   &   M   \\\\ 23554+5612  &   \\object{WY Cas} &   CSS 1345    &   23 58 01.30 +56 29 13.5 &   S6,5pe  &   M   \\\\ 00135+4644  &   \\object{X And}  &   CSS 6   &   00 16 09.57 +47 00 44.8 &   S2,9e$-$S5,5e   &   M   \\\\ 00435+4758  &   \\object{U Cas}  &   CSS 12  &   00 46 21.371 +48 14 38.72   &   S3,5e$-$S8,6e   &   M   \\\\ 06062+2830  &   \\object{GH Aur} &   CSS 191 &   06 09 27.71 +28 29 43.4 &   S   &   M   \\\\ 07197$-$1451    &   \\object{TT CMa} &   CSS 341 &   07 22 02.00 $-$14 56 56.7   &   S   &   M   \\\\ 07584$-$2051    &   \\object{EX Pup} &   CSS 443 &   08 00 38.31 $-$20 59 35.3   &   S2,4e   &   M   \\\\ 11179$-$6135    &   \\object{RY Car} &   CSS 742 &   11 20 11.39 $-$61 52 16.8   &   S7,8e   &   M   \\\\ 13226$-$6302    &   \\object{NZ Cen} &   CSS 820 &   13 26 02.52 $-$63 18 28.5   &   Se  &   M   \\\\ 14212+8403  &   \\object{R Cam}  &   CSS 856 &   14 17 51.0439 +83 49 53.861     &   S2,8e$-$S8,7e   &   M   \\\\ 17478$-$2957    &   \\object{V762 Sgr}   &   CSS 1001    &   17 51 04.04 $-$29 58 30.9   &   S6,4    &   M   \\\\ 17490$-$3502    &   \\object{V407 Sco}   &   CSS 1004    &   17 52 25.53 $-$35 03 17.5   &   Se  &   M   \\\\ 19166+0318  &   \\object{ER Aql} &   CSS 1121    &   19 19 06.99 +03 24 05.2 &   S   &   M   \\\\ 23376+6304  &   \\object{V441 Cas}   &   CSS 1338    &   23 39 58.92 +63 20 55.1 &   S   &   M   \\\\ 23489+6235  &   \\object{EO Cas} &   CSS 1342    &   23 51 27.30 +62 51 47.0 &   Se  &   M   \\\\ 00001+4826  &   \\object{IW Cas} &   CSS 1347    &   00 02 44.22 +48 42 50.9 &   S4.5,9e &   M   \\\\ $-$ &   $-$ &   \\object{CSS2 10}    &   02 51 33.00 +57 50 34.5 &   S   &   M   \\\\ \\hline 10237$-$6135    &   \\object{AU Car} &   CSS 679 &   10 25 29.74 $-$61 50 59.1   &   MS  &   M   \\\\ 10349$-$6203    &   \\object{RX Car} &   CSS 690 &   10 36 45.82 $-$62 19 16.8   &   MS  &   M   \\\\ \\hline 02143+4404  &   \\object{W And}  &   CSS 49  &   02 17 32.9606 +44 18 17.766 &   S6,1e$-$S9,2e/M4$-$M1   &   M   \\\\ 07149+0111  &   \\object{RR Mon} &   CSS 326 &   07 17 31.54 +01 05 41.5 &   S7,2e$-$S8,2e/M6$-$10   &   M   \\\\ 07545$-$4400    &   \\object{SU Pup} &   CSS 436 &   07 56 12.0813 $-$44 08 33.254   &   M/S4,2e &   M   \\\\ 09338$-$5349    &   \\object{UU Vel} &   CSS 614 &   09 35 33.21 $-$54 03 25.9   &   M2e/S7,8e   &   M   \\\\ 17001$-$3651    &   \\object{RT Sco} &   CSS 954 &   17 03 32.56 $-$36 55 13.7   &   S7,2/M6e$-$M7e  &   M   \\\\ 20213+0047  &   \\object{V865 Aql}   &   CSS 1211    &   20 23 54.6422 +00 56 44.794 &   M6$-$M7/S7,5e:  &   M   \\\\ 00445+3224  &   \\object{RW And} &   CSS 14  &   00 47 18.92 +32 41 08.6 &   M5e$-$M10e/S6,2e    &   M   \\\\ 07103$-$0258    &   \\object{AK Mon} &   CSS 319 &   07 12 49.91 $-$03 03 29.0   &   M5/S5,1 &   M   \\\\ 17521$-$2907    &   \\object{V745 Sgr}   &   CSS 1007    &   17 55 19.00 $-$29 07 54.4   &   Se/M    &   M   \\\\ 20044+5750  &   \\object{S Cyg}  &   CSS 1191    &   20 05 29.85 +57 59 09.1 &   S2.5,1e/M3.5$-$M7e  &   M   \\\\ 20369+3742  &   \\object{FF Cyg} &   CSS 1232    &   20 38 51.71 +37 53 23.2 &   S6,8e/M4e   &   M   \\\\ \\hline 03499+4730  &   \\object{FG Per} &   CSS 85  &   03 53 30.2 +47 39 04    &   M9  &   M   \\\\ \\hline 13163$-$6031    &   \\object{TT Cen} &   CSS 817 &   13 19 35.016 $-$60 46 46.26 &   CSe &   M   \\\\ 18586$-$1249    &   \\object{ST Sgr} &   CSS 1096    &   19 01 29.20 $-$12 45 34.0   &   C4,3e$-$S9,5e   &   M   \\\\ 21540+4806  &   \\object{LX Cyg} &   CSS 1286    &   21 55 57.03 +48 20 52.6 &   SC3e$-$S5,5e:   &   M   \\\\ \\hline 18575$-$0139    &   \\object{VX Aql} &   CSS 1093    &   19 00 09.61 $-$01 34 56.8   &   C9,1p/M0ep  &   M   \\\\ \\hline 01097+6154  &   \\object{V418 Cas}   &   CSS 23  &   01 12 59.89 +62 10 47.6 &   $-$ &   M   \\\\ \\hline \\hline \\end{tabular} \\end{table*} %\\end{landscape}  %\\begin{landscape} \\begin{table*}[t!] \\caption{Sample C $-$ Second part.}             % title of Table \\label{table:8}      % is used to refer this table in the text \\centering                          % used for centering table \\begin{tabular}{c c c c c c c c c c c}        % centered columns (4 columns) \\hline      \\hline Source  &   J   &   H   &   K   &   [8.8]   &   [9.8]   &   [11.7]  &   [12.5]  &   D   &   E   &   Mid$-$IR Data Origin    \\\\ name    &   [Jy]    &   [Jy]    &   [Jy]    &   [Jy]    &   [Jy]    &   [Jy]    &   [Jy]    &   [Jy]    &   [Jy]    &       \\\\ \\hline      \\hline T Cam   &   152 &   306 &   315 &   62.1    &   53.4    &   48.8    &   47.1    &   $-$ &   $-$ &   IRAS$-$LRS  \\\\ R Lyn   &   51.7    &   77.7    &   91.1    &   29.9    &   29.9    &   29.2    &   28.5    &   $-$ &   $-$ &   IRAS$-$LRS  \\\\ R Gem   &   155 &   226 &   174 &   24.1    &   $-$ &   $-$ &   21.4    &   16.3    &   7.5 &   MSX \\\\ WX CMi  &   8.9 &   13.8    &   17.3    &   8.3 &   $-$ &   $-$ &   9.2 &   6.3 &   4.1 &   MSX \\\\ S UMa   &   26.3    &   43.4    &   41.4    &   4.5 &   $-$ &   2.8 &   3.1 &   $-$ &   $-$ &   TIRCAM2 \\\\ GI Lup  &   61.7    &   115 &   128 &   $-$ &   $-$ &   $-$ &   $-$ &   $-$ &   $-$ &   $-$ \\\\ WY Cas  &   64.5    &   94.4    &   120 &   48.8    &   55.6    &   58.2    &   55.0    &   $-$ &   $-$ &   IRAS$-$LRS  \\\\ X And   &   18.1    &   37.3    &   43.3    &   23.6    &   23.6    &   24.9    &   26.6    &   $-$ &   $-$ &   IRAS$-$LRS  \\\\ U Cas   &   29.6    &   45.0    &   47.1    &   15.2    &   15.6    &   19.4    &   20.3    &   $-$ &   $-$ &   IRAS$-$LRS  \\\\ GH Aur  &   4.4 &   7.8 &   10.2    &   1.5 &   $-$ &   $-$ &   1.8 &   $-$ &   $-$ &   MSX \\\\ TT CMa  &   29.6    &   50.1    &   54.3    &   12.3    &   $-$ &   $-$ &   14.1    &   10.8    &   7.7 &   MSX \\\\ EX Pup  &   3.9 &   5.9 &   5.8 &   0.54    &   $-$ &   $-$ &   $-$ &   $-$ &   $-$ &   MSX \\\\ RY Car  &   8.3 &   18.2    &   25.1    &   11.0    &   $-$ &   $-$ &   8.9 &   5.6 &   3.2 &   MSX \\\\ NZ Cen  &   8.3 &   15.7    &   17.2    &   9.1 &   $-$ &   $-$ &   9.3 &   6.3 &   5.2 &   MSX \\\\ R Cam   &   47.2    &   70.8    &   69.5    &   7.1 &   $-$ &   $-$ &   $-$ &   2.8 &   $-$ &   MSX \\\\ V762 Sgr    &   15.0    &   35.9    &   44.5    &   13.5    &   $-$ &   $-$ &   14.4    &   9.3 &   5.1 &   MSX \\\\ V407 Sco    &   9.8 &   18.1    &   19.4    &   6.1 &   $-$ &   $-$ &   5.6 &   3.8 &   $-$ &   MSX \\\\ ER Aql  &   34.4    &   61.5    &   62.3    &   10.1    &   $-$ &   $-$ &   6.9 &   4.7 &   $-$ &   MSX \\\\ V441 Cas    &   5.9 &   11.6    &   14.3    &   2.0 &   $-$ &   $-$ &   1.5 &   $-$ &   $-$ &   MSX \\\\ EO Cas  &   11.8    &   21.4    &   30.9    &   9.3 &   $-$ &   $-$ &   7.5 &   4.8 &   3.0 &   MSX \\\\ IW Cas  &   23.3    &   47.4    &   50.9    &   52.9    &   58.5    &   61.8    &   60.5    &   $-$ &   $-$ &   IRAS$-$LRS  \\\\ CSS2 10 &   1.3 &   2.8 &   3.1 &   0.36    &   $-$ &   $-$ &   $-$ &   $-$ &   $-$ &   MSX \\\\ \\hline AU Car  &   5.9 &   9.0 &   10.3    &   2.1 &   $-$ &   $-$ &   $-$ &   $-$ &   $-$ &   MSX \\\\ RX Car  &   5.0 &   7.7 &   8.9 &   2.1 &   $-$ &   $-$ &   1.7 &   $-$ &   $-$ &   MSX \\\\ \\hline W And   &   368 &   643 &   591 &   185 &   198 &   163 &   143 &   $-$ &   $-$ &   IRAS$-$LRS  \\\\ RR Mon  &   24.1    &   42.5    &   52.8    &   21.3    &   $-$ &   $-$ &   18.0    &   12.7    &   7.2 &   MSX \\\\ SU Pup  &   28.5    &   43.8    &   45.5    &   22.3    &   26.6    &   29.1    &   29.3    &   $-$ &   $-$ &   IRAS$-$LRS  \\\\ UU Vel  &   23.5    &   46.9    &   51.6    &   13.6    &   $-$ &   $-$ &   11.5    &   7.5 &   4.1 &   MSX \\\\ RT Sco  &   285 &   460 &   512 &   162 &   $-$ &   $-$ &   163 &   108 &   67.4    &   MSX \\\\ V865 Aql    &   126 &   191 &   199 &   36.5    &   33.8    &   35.2    &   35.8    &   $-$ &   $-$ &   IRAS$-$LRS  \\\\ RW And  &   95.9    &   132 &   131 &   48.8    &   52.8    &   53.0    &   50.1    &   $-$ &   $-$ &   IRAS$-$LRS  \\\\ AK Mon  &   9.3 &   13.1    &   14.7    &   3.1 &   $-$ &   $-$ &   2.8 &   2.2 &   $-$ &   MSX \\\\ V745 Sgr    &   38.2    &   75.8    &   82.5    &   18.5    &   $-$ &   $-$ &   16.1    &   11.3    &   7.4 &   MSX \\\\ S Cyg   &   11.8    &   14.0    &   16.5    &   2.0 &   $-$ &   $-$ &   1.5 &   0.74    &   $-$ &   MSX \\\\ FF Cyg  &   37.8    &   65.8    &   82.0    &   8.5 &   $-$ &   $-$ &   6.4 &   4.9 &   $-$ &   MSX \\\\ \\hline FG Per  &   3.2 &   6.0 &   6.8 &   1.1 &   $-$ &   $-$ &   $-$ &   $-$ &   $-$ &   MSX \\\\ \\hline TT Cen  &   19.5    &   48.9    &   54.7    &   15.5    &   $-$ &   $-$ &   16.3    &   10.1    &   8.0 &   MSX \\\\ ST Sgr  &   97.7    &   148 &   149 &   61.6    &   63.2    &   61.1    &   56.4    &   $-$ &   $-$ &   IRAS$-$LRS  \\\\ LX Cyg  &   14.2    &   28.0    &   44.9    &   10.2    &   $-$ &   $-$ &   9.4 &   5.1 &   $-$ &   MSX \\\\ \\hline VX Aql  &   13.5    &   33.6    &   37.7    &   9.3 &   $-$ &   $-$ &   10.6    &   6.8 &   3.5 &   MSX \\\\ \\hline V418 Cas    &   11.0    &   16.0    &   22.8    &   12.5    &   $-$ &   $-$ &   12.2    &   7.9 &   5.5 &   MSX \\\\ \\hline      \\hline \\end{tabular} \\end{table*} %\\end{landscape}  %\\begin{landscape} \\begin{table*}[t!] \\caption{Sample C $-$ Third part.}             % title of Table \\label{table:9}      % is used to refer this table in the text \\centering                          % used for centering table \\begin{tabular}{c c c c c c c}        % centered columns (4 columns) \\hline              \\hline Source  &   Var. Type   &   Period  &   Distance    &   Ref.    &   Bol. Magnitudes &   I. $-$ E.   \\\\ name    &   (GCVS)  &   (GCVS)  &   (kpc)   &   Distance    &   P$-$L Method    &       \\\\ \\hline              \\hline T Cam   &   M   &   373.2   &   0.50    &   P$-$L / this paper  &   $-$5.22 &   I   \\\\ R Lyn   &   M   &   378.75  &   0.95    &   P$-$L / this paper  &   $-$5.19 &   I   \\\\ R Gem   &   M   &   369.91  &   0.66    &   P$-$L / this paper  &   $-$5.23 &   I   \\\\ WX CMi  &   M   &   420.1   &   2.35    &   P$-$L / this paper  &   $-$5.31 &   $-$ \\\\ S UMa   &   M   &   225.87  &   0.96    &   P$-$L / this paper  &   $-$4.56 &   I   \\\\ GI Lup  &   M   &   326.2   &   0.80    &   P$-$L / this paper  &   $-$ &   I   \\\\ WY Cas  &   M   &   476.56  &   0.97    &   P$-$L / this paper  &   $-$5.50 &   I   \\\\ X And   &   M   &   346.18  &   1.31    &   P$-$L / this paper  &   $-$5.03 &   I   \\\\ U Cas   &   M   &   277.19  &   1.08    &   P$-$L / this paper  &   $-$4.74 &   I   \\\\ GH Aur  &   M   &   349 &   2.64    &   P$-$L / this paper  &   $-$5.13 &   I   \\\\ TT CMa  &   M   &   314 &   1.08    &   P$-$L / this paper  &   $-$4.95 &   I   \\\\ EX Pup  &   M   &   289 &   3.02    &   P$-$L / this paper  &   $-$4.92 &   $-$ \\\\ RY Car  &   M   &   424.3   &   1.95    &   P$-$L / this paper  &   $-$5.34 &   I   \\\\ NZ Cen  &   M   &   382 &   2.22    &   P$-$L / this paper  &   $-$5.17 &   I   \\\\ R Cam   &   M   &   270.22  &   0.84    &   P$-$L / this paper  &   $-$4.81 &   E   \\\\ V762 Sgr    &   M   &   444 &   1.50    &   P$-$L / this paper  &   $-$5.42 &   I   \\\\ V407 Sco    &   M   &   396 &   2.11    &   P$-$L / this paper  &   $-$5.26 &   I   \\\\ ER Aql  &   M   &   337.6   &   1.04    &   P$-$L / this paper  &   $-$5.10 &   I   \\\\ V441 Cas    &   M   &   175.6   &   1.40    &   P$-$L / this paper  &   $-$4.17 &   \\textbf{$\\underline{E}$}    \\\\ EO Cas  &   M   &   455 &   1.83    &   P$-$L / this paper  &   $-$5.47 &   I   \\\\ IW Cas  &   M   &   396.38  &   1.34    &   P$-$L / this paper  &   $-$5.19 &   $-$ \\\\ CSS2 10 &   M   &   250 &   3.78    &   P$-$L / this paper  &   $-$4.69 &   $-$ \\\\ \\hline AU Car  &   M   &   332 &   2.54    &   P$-$L / this paper  &   $-$5.05 &   $-$ \\\\ RX Car  &   M   &   332.8   &   2.75    &   P$-$L / this paper  &   $-$5.04 &   $-$ \\\\ \\hline W And   &   M   &   395.93  &   0.38    &   P$-$L / this paper  &   $-$5.27 &   I   \\\\ RR Mon  &   M   &   394.7   &   1.28    &   P$-$L / this paper  &   $-$5.24 &   I   \\\\ SU Pup  &   M   &   339.8   &   1.26    &   P$-$L / this paper  &   $-$5.01 &   I   \\\\ UU Vel  &   M   &   408.9   &   1.31    &   P$-$L / this paper  &   $-$5.32 &   I   \\\\ RT Sco  &   M   &   449.04  &   0.45    &   P$-$L / this paper  &   $-$5.44 &   I   \\\\ V865 Aql    &   M   &   364.8   &   0.62    &   P$-$L / this paper  &   $-$5.18 &   $-$ \\\\ RW And  &   M   &   430.3   &   0.86    &   P$-$L / this paper  &   $-$5.36 &   $-$ \\\\ AK Mon  &   M   &   328.6   &   2.12    &   P$-$L / this paper  &   $-$5.02 &   $-$ \\\\ V745 Sgr    &   M   &   380.2   &   0.99    &   P$-$L / this paper  &   $-$5.23 &   $-$ \\\\ S Cyg   &   M   &   322.93  &   1.94    &   P$-$L / this paper  &   $-$5.06 &   I   \\\\ FF Cyg  &   M   &   323.82  &   0.87    &   P$-$L / this paper  &   $-$5.07 &   I   \\\\ \\hline FG Per  &   M   &   340.3   &   3.16    &   P$-$L / this paper  &   $-$5.10 &   $-$ \\\\ \\hline TT Cen  &   M   &   462 &   1.39    &   P$-$L / this paper  &   $-$5.48 &   $-$ \\\\ ST Sgr  &   M   &   395.12  &   0.76    &   P$-$L / this paper  &   $-$5.24 &   I   \\\\ LX Cyg  &   M   &   465.3   &   1.53    &   P$-$L / this paper  &   $-$5.51 &   I   \\\\ \\hline VX Aql  &   M   &   604 &   1.99    &   P$-$L / this paper  &   $-$5.87 &   I   \\\\ &       &   \\citet{zijlstra}    &       &       &       &       \\\\ \\hline V418 Cas    &   M   &   480 &   2.24    &   P$-$L / this paper  &   $-$5.50 &   $-$ \\\\ \\hline              \\hline \\end{tabular} \\end{table*} %\\end{landscape}  In a previous paper of this series, hereafter referred to as \"Paper I\" \\citep{guandalini}, we analyzed a sample of C stars reconstructing their SEDs up to far infrared, on the basis of space-borne infrared observations from the ISO and MSX missions. We found evidence for a relatively large average C-star luminosity, thus suggesting that the so-called \"C-star luminosity problem\" \\citep{cohen} might not be real, being simply an effect of poor estimates of the luminosity, due to insufficient knowledge of the mid-infrared emission. We also reviewed the available mass loss rates and showed their correlation with infrared colors.  We want now to extend that analysis, considering those thermally-pulsing AGB stars where the enhancement of C (and s-elements) is more moderate than in C stars: it is the case of MS and S giants \\citep{busso92,busso95}. In Sect. \\ref{sect2} we present the sample stars, and we discuss the choices made in selecting and organizing them in sub-samples, according to the quality of the available data. In Sect. \\ref{sect3} we present the IR colors and derive the bolometric corrections, based on a set of sources whose magnitude can be estimated safely through the integral of detailed SEDs. We also use these corrections for inferring the critical parameters (absolute Magnitudes or distances) of sources for which either i) we have incomplete IR coverage but reliable distance estimates; or ii) Period-Luminosity relations yield the Luminosity, and the distance needs to be inferred from the distance modulus. (For the sake of clarity, the adopted Period-Luminosity relations are discussed in Appendix \\ref{app1}). Once the absolute Magnitudes are known, in Sect. \\ref{sect4} we can analyze HR diagrams and luminosity functions and on this base we also attempt a rough estimate of photometric parallaxes for Mira variables with no other available data on Luminosities. Then, in Sect. \\ref{sect5} some preliminary conclusions are derived, while we postpone to a forthcoming dedicated work the analysis of stellar winds. ",
        "conclusions": "}  \\begin{figure*}[t!!] \\centering \\resizebox{\\textwidth}{!} {\\includegraphics[height=6cm,angle=-90]{gbf8.eps}} \\caption{The magnitude variations of two model stars of Population I, during the thermally pulsing AGB phase, as computed by the FRANEC code. The abscissa shows the time passed after the ignition of He-burning in the core, in Million years (once an offset is deduced, whose value is indicated).  Most sources with features typical of S or C stars and magnitudes in the range $-5.2 \\pm 0.4$ should belong to the mass and metallicity range illustrated here, but only the star in the right panel does achieve the C/O $>$ 1 condition.} \\label{fig8} \\end{figure*}  In this paper we have presented a reanalysis of the properties of MS-S-SC stars, based on a sample of about 600 sources, whose infrared fluxes from 1.25 to 21 $\\mu$m were measured by the 2MASS, IRAS, ISO and MSX experiments. A 'best' group of 21 stars (for which detailed ISO-SWS spectra are available up to long wavelengths) allowed us to obtain the bolometric Magnitudes from an effective integral of the spectral energy distribution up to 45 $\\mu$m. Correlations with near-to-mid IR colors then allowed us to infer bolometric corrections suitable to be applied to other groups of sources, with a less detailed coverage of the energy distribution. We can thus estimate with sufficient accuracy the apparent bolometric magnitudes of more than 500 sources. The whole analysis was performed in the photometric system suggested in \\citet{busso96} and subsequently used in Paper I and Paper II.  Criteria for obtaining the distance have then been discussed, from the simple use of revised astrometric measurements to a reformulation of the known Period-Luminosity relations for O-rich Long Period variables. The results of our analysis suggest that Mira variables of the S type have on average magnitudes in the range $-5.15\\pm0.4$, showing a well-defined linear correlation with infrared colors, especially the K-[8.8] one. Low-mass AGB stars do not appear to proceed beyond the upper limit of the Mira luminosity range. Inside this range $P-L$ relations have become rather tight and accurate; they now form a very useful tool for determining intrinsic stellar parameters. From the $P-L$ relations and the $M_{bol} - (K-[8.8])$ relation of Equation \\ref{eq5} we now have tools to estimate the absolute Magnitudes, hence the distances, of Mira S variables for which either the period or the infrared colors have been determined.  From statistical considerations it is also argued that Mira variables should occupy mainly the final part of the AGB track, as the simple intermittency between low-luminosity post flash dips and high-luminosity H-shell-powered stages is not sufficient to explain the available numbers of Semiregulars and Miras.  Luminosity functions confirm that intrinsic S stars (especially if in the Mira class) are distributed over a narrow range around the above average magnitude, and that this last datum is very close to the one previously found for C-stars. This is so to the point that it appears unlikely that a single AGB star can follow the whole M-MS-S-SC-C sequence, by simply increasing gradually its content of carbon and heavy elements as time passes, new dredge-up episodes enrich the envelope and the luminosity increases. Most probably, small differences in the initial mass and metallicity, almost indistinguishable once on the AGB, determine the final chemical fate of a star, which, for increasing initial mass, can end its life either as a MS-S giant, or a SC giant, or reaching effectively the C(N) stage. As an example of the effects of small mass (and metallicity) differences, Fig. \\ref{fig8} shows two AGB luminosity sequences obtained from the FRANEC code by \\citet{busso03}. The model star of the left panel does not reach the C-star phase, ending as an MS/S star. The evolutionary phases during which S-type chemical peculiarities are exhibited by the envelope are shaded in the plot. In the right panel we instead show the results of a model producing a real C-star: it has a slightly higher mass. Given the magnitudes we found in this paper for the sources of types MS and S, and those found in Paper I for C stars, most of the objects in our sample should belong to the mass and metallicity range covered by Fig. \\ref{fig8}. If we broadly consider the typical magnitudes of S and C stars as being in the range $-5.2 \\pm$ 0.4, then it is clear that this range includes both the S-star phase in the first panel of Fig. \\ref{fig8} and the S- and C-star phases in the right panel. It would be very difficult or impossible to distinguish between the two cases on the basis of their magnitudes, given the remaining uncertainties on distances and on possibly variable bolometric magnitudes.  The above results, and the typical initial mass implied for S and C stars (1.5 - 2 $M_{\\odot}$) might appear rather peculiar, in view of the fact that about half of Planetary Nebulae are carbon rich. This cannot be explained by progenitors of around 2 $M_{\\odot}$, at least according to the most common choices for the Initial Mass Function. However we must remember that the efficiency of dredge up strongly increases with decreasing metallicity. In halo stars one solar mass might be sufficient for forming a C star (also because the abundance of oxygen is very small). Hence the results presented here should be considered as valid only for galactic disc stars, at relatively high metallicities."
    },
    "0806/0806.0617_arXiv.txt": {
        "abstract": "Dark Stars are the very first phase of stellar evolution in the history of the universe: the first stars to form (typically at redshifts $z \\sim 10-50$) are powered by heating from dark matter (DM) annihilation instead of fusion (if the DM is made of particles which are their own antiparticles).  We find equilibrium polytropic configurations for these stars; we start from the time DM heating becomes important ($M \\sim 1-10~M_\\odot$) and build up the star via accretion up to 1000~ M$_\\odot$.  The dark stars, with an assumed particle mass of 100 GeV, are found to have luminosities of a few times $10^6$ L$_\\odot$, surface temperatures of 4000--10,000 K, radii $\\sim 10^{14}$ cm, lifetimes of at least $ 0.5$ Myr, and are predicted to show lines of atomic and molecular hydrogen.  Dark stars look quite different from standard metal-free stars without DM heating: they are far more massive (e.g. $\\sim 800 M_\\odot$ for 100 GeV WIMPs), cooler, and larger, and can be distinguished in future observations, possibly even by JWST or TMT. ",
        "introduction": "The first stars in the Universe mark the end of the cosmic dark ages, reionize the Universe, and provide the enriched gas required for later stellar generations.  They may also be important as precursors to black holes that coalesce and power bright early quasars.  The first stars are thought to form inside dark matter (DM) halos of mass $ 10^5 M_\\odot$--$ 10^6 M_\\odot$  at redshifts $z \\sim 10-50$ (Yoshida et al. 2003).  These halos consist of 85\\% DM and 15\\% baryons in the form of metal-free gas made of H and He.  Theoretical calculations indicate that the baryonic matter cools and collapses via H$_2$ cooling (Peebles \\& Dicke 1968, Matsuda et al. 1971, Hollenbach \\& McKee 1979) into a single small protostar (Omukai \\& Nishi 1998) at the center of the halo (for reviews see Ripamonti \\& Abel 2005; Barkana \\& Loeb 2001; Bromm \\& Larson 2004).  Previously, Spolyar et al. (2008; hereafter, Paper I) first considered the effect of DM  particles on the first stars during their formation.  Any DM particle which is capable of annihilating with itself in such a way as to give the correct relic abundance today will also annihilate wherever the DM density is high. The first protostars and stars are particularly good sites for annihilation because they form at high redshifts (density scales as $(1+z)^3$) and in the high density centers of DM haloes.  Paper I  found that DM annihilation provides a powerful heat source in the first stars, a source so intense that its heating overwhelms all cooling mechanisms. Paper I suggested that the very first stellar objects might be {\\it Dark Stars} (DS), a new phase of stellar evolution in which the DM -- while only a negligible fraction of the star's mass -- provides the key power source for the star through DM heating. Note that the term 'Dark' refers to the power source, not the luminosity.  In this paper, we continue the work originally suggested in Paper I by studying the DS structure.  The canonical example of particle DM is Weakly Interacting Massive Particles (WIMPs), which automatically provide the right amount of DM, i.e.\\ $\\sim$ 24\\% of the current energy density of the Universe.  In many theories WIMPs are their own antiparticles and annihilate with themselves in the early universe, leaving behind this relic density. In particular, the neutralino, the supersymmetric partner of the W, Z, and Higgs bosons, is a strong candidate (reviewed by Jungman et al. 1996).  As our canonical values, we use the standard $\\langle \\sigma v \\rangle = 3 \\times 10^{-26}\\,{\\rm cm^3/s}$ for the annihilation cross section and $m_\\chi = 100\\,{\\rm GeV}$ for the particle mass.  A companion paper will generalize to other masses and cross sections. The analysis in this paper could apply equally well to other DM candidates.  WIMP annihilation produces energy at a rate per unit volume \\begin{equation} \\hat Q_{DM} = \\langle \\sigma v \\rangle \\rho_\\chi^2/m_\\chi , \\label{eq:Q} \\end{equation} where $\\rho_\\chi$ is the energy density of the WIMPs. In the early stages of Pop III star formation, when the gas density is low ($n \\lesssim 10^4 {\\rm cm}^{-3}$), most of the annihilation products simply escape from the protostar without heating it (Ripamonti et al. 2007). However, a crucial transition takes place (Paper I)  when the gas density of the collapsing protostar exceeds a critical value at which point most of the annihilation energy is trapped in the star.  For a 100 GeV particle,  at hydrogen density  $\\sim 10^{13}\\, {\\rm cm}^{-3}$, typically 1/3 of the energy is lost to neutrinos that escape the star, while the other 2/3 of the energy is trapped inside the star. Hence the luminosity from the DM heating is \\begin{equation} \\label{DMheating} L_{DM} \\sim {2 \\over 3} \\int \\hat Q_{DM} dV \\end{equation} where $dV$ is the volume element.  The properties of the collapsing protostellar clouds have been given by 3D simulations (Abel et al. 2002; Gao et al. 2007).  At the time when the density reaches $n = 10^{13}$cm$^{-3}$, the critical value for 100 GeV particles, Paper I found a proto-DS in equilibrium with a radius of 17 AU and a mass of $0.6 M_\\odot$, giving a DM luminosity of $\\sim 140$ L$_\\odot$. As more mass accretes onto the DS, the protostellar luminosity begins to exceed the DM heating, so that the protostar is no longer in thermal equilibrium.  Thus it must contract which increases the DM density until the DM heating as given in equation (\\ref{DMheating}) matches its radiated luminosity.  In this calculation, we assume that such a situation can be reached, and we then build up the Dark Star from a few solar masses up to 1000 $M_\\odot$, finding its structure as a polytrope in hydrostatic and thermal equilibrium at each step in mass.  As we build up the star more DM is pulled into the star via adiabatic contraction and subsequently annihilates; we find that the annihilation fuel contained in the star can thereby last $\\sim 10^6$yr.  While the results of this paper were being written, a paper appeared by Iocco et al. (2008) which included DM heating in Pop III pre-main-sequence evolution of a set of stars of {\\it fixed mass}, finding that the quasi-hydrostatic contraction is halted for times of $2 \\times 10^3$ ($2 \\times 10^4$) yr for stars of mass 600 (9) M$_\\odot$, at radii $\\approx$ a few AU.  During the evolution of a DS, additional WIMPs could be captured via scattering off of nuclei.  The cross section for scattering ($\\sigma_s$) is very uncertain. For $\\sigma_s <10^{-39}$ cm$^2$ we find that a DM particle undergoes less than one scattering event in 1 Myr in the evolutionary stage considered in this paper. The experimental bounds for 100 GeV particles from DM searches are $\\sigma_s \\lesssim 2 \\times 10^{-43}$ cm$^2$ for the spin-independent case (Gaitskill et al. 2008) and $\\sigma_s \\lesssim 3.5 \\times 10^{-39}$ cm$^2$ for the spin-dependent case (Savage et al. 2004). Hence we assume negligible scattering here. However at later stages of the evolution, once the DM density becomes too low to support the star via heating, the DS contracts until nuclear burning sets in. At these higher densities scattering at the experimentally allowed limit would become important. DM passing through the star could be captured and again drive DM heating.  These effects have been considered for main-sequence and pre-main-sequence DS (Freese et al. 2008; Iocco 2008; Iocco et al. 2008), who find that the DM heating could dominate nuclear fusion as long as the background DM density (from which the capture takes place) remains high enough. Future work will further consider scattering in the DS.  We also cite previous work on DM annihilation in today's stars (less powerful than in the first stars): Krauss et al (1985); Bouquet \\& Salati (1989); Salati \\& Silk (1989); Moskalenko \\& Wai (2007); Scott et al. (2007); Bertone \\& Fairbairn (2007). ",
        "conclusions": "We have followed the growth of equilibrium Dark Stars, powered by DM annihilation, up to 1000 M$_\\odot$. The objects have sizes of a few AU and central $T_c\\approx 10^5-10^6$ K.  Sufficient DM is brought into the star by contraction from the DM halo to result in a DS which lives at least 0.5 Myr  (the lifetime could be significantly longer if DM capture becomes important at the later stages, as long as the background DM density is high enough for capture to take place). Because of the relatively low $T_{\\rm eff}$ (4000--10,000 K), feedback mechanisms for shutting off accretion of baryons, such as the formation of HII regions or the dissociation of infalling H$_2$ by Lyman-Werner photons, are not effective.  The implication is that main-sequence stars of Pop. III are very massive. This conclusion depends on uncertain parameters such as the DM particle mass, the accretion rate, and scattering, effects that will be studied in future work.  Although DS shine with a few $10^6 L_\\odot$ they would be very difficult to observe at $z \\sim 10-50$. One can speculate that pristine regions containing only H and He might still exist to lower redshifts; then DS forming in these regions might be easier to detect. One may hope that the ones that form most recently are detectable by JWST or TMT and differentiable from the standard metal-free Pop. III objects.  DS are also predicted to have atomic hydrogen lines originating in the warmer photospheres, and H$_2$ lines arising from the infalling material, which is still relatively cool.  It has been argued that Pop III.1 stars (the very first metal-free stars) may constitute at most $\\sim 10\\%$ of metal poor stars on observational grounds.  Heger \\& Woosley (2002; HW) showed that for $140 M_\\odot < M < 260 M_\\odot$, pair instability (SN) lead to odd-even effects in the nuclei produced that are strongly constrained by observations.  Thus if Pop III.1 stars are really in this mass range one would have to constrain their abundance. For $M > 260 M_\\odot$, HW find that no SN occurs, and the end result of stellar evolution is collapse of the entire star into a black hole. We expect, based on extension of the $n=3$ calculation, that our DS runs out of DM at about 700--900 M$_\\odot$ (for $m_\\chi = 100$ GeV).  Then it must contract to the main sequence, where nuclear burning sets in, and further evolution would proceed as in HW. Alternatively, the evolution could proceed as described by Ohkubo et al. (2006) who found that metal-free stars of 500 and 1000 $M_\\odot$, taking into account two-dimensional effects, did blow up as SN, leaving about half their mass behind in a black hole.  In this case the SN might be observable signatures of DS, distinguishable since they arise from such high mass stars.  The end product in either case would be a plausible precursor of the otherwise unexplained $10^9 M_\\odot$ black holes at $z=6$ (Yoshida et al., in prep).  We acknowledge support from: the DOE and MCTP via the Univ.\\ of Michigan (K.F.); NSF grant AST-0507117 and GAANN (D.S.); NSF grant PHY-0456825 (P.G.). K.F. acknowledges the hospitality of the Physics Dept. at the Univ. of Utah.  K.F. and D.S. are extremely grateful to Chris McKee and Pierre Salati for their encouragement of this line of research, and to A. Aguirre, L. Bildsten, R. Bouwens, J. Gardner, N. Murray, J. Primack, M. Rieke, C. Savage, J. Sellwood, J. Tan, and N. Yoshida for helpful discussions."
    },
    "0806/0806.0567_arXiv.txt": {
        "abstract": "Thanks to its sharp view, HST has significantly improved our knowledge of tens of gravitationally lensed quasars in four different respects: (1) confirming their lensed nature; (2) detecting the lensing galaxy responsible for the image splitting; (3) improving the astrometric accuracy on the positions of the unresolved QSO images and of the lens; (4) resolving  {\\em extended} lensed structures from the QSO hosts into faint NIR or  optical rings or arcs. These observations have helped to break some degeneracies on the lens potential, to probe the galaxy evolution and to reconstruct the true shape of the QSO host with an increased angular resolution. ",
        "introduction": "\\label{sec:intro}  Strong lensing, i.e. the splitting of the image of a background source into several, magnified but distorted images, occurs each time an intervening massive object lies nearly on the same line-of-sight, provided its surface mass density is large enough (typically larger than 0.5 gr/cm$^2$ at cosmological distances). The angular resolution of the telescope must be sufficient to resolve the images ($\\Delta \\theta \\simeq 1'' \\sqrt{M_E/10^{11}M_\\odot}$). This is a consequence of the curvature of space-time around massive objects, as predicted by General Relativity.  However, a massive object such as an isolated galaxy can be considered as an imperfect optical lens (like the foot of a glass of wine). Instead of a focus, a generic diamond-shaped closed curve is produced -- the caustic. All observed configurations of multiply imaged QSOs as well as giant arcs can be explained by the exact relative position of the (possibly extended) source with respect to that caustic (see Fig. \\ref{fig:lens_config}). Of course, besides being an exotic curiosity, this phenomenon betrays the nature of the deflector and can be used as an astrophysical tool. Reviews on gravitational optics and its astrophysical applications may be found in \\cite{rev06} and references therein. Here, we concentrate on selected results derived from HST observations of lensed QSOs.  \\begin{figure}[htb] \\hbox to \\columnwidth{  \\includegraphics[width=2.8cm]{h1413_obs.eps} \\includegraphics[width=2.8cm]{pg1115_obs.eps} \\includegraphics[width=2.8cm]{b1422_obs.eps} \\includegraphics[width=2.8cm]{he1104_obs.eps} }  \\hbox to \\columnwidth{ \\includegraphics[width=2.8cm]{h1413.eps} \\includegraphics[width=2.8cm]{pg1115.eps} \\includegraphics[width=2.8cm]{b1422.eps} \\includegraphics[width=2.8cm]{he1104.eps} } \\caption{{\\em Top:} Deconvolved HST NICMOS observations of 4 gravitationally lensed QSOs (CASTLES project, {\\sl http://cfa-www.harvard.edu/glensdata}); {\\em bottom:} fitted source position w.r.t. the caustics of the Singular Isothermal Ellispoid model. From left to right: 4 image cross, fold and  cusp configurations, 2 image configuration.} \\label{fig:lens_config} \\end{figure}   In 1979, the first gravitationally lensed source was discovered serendipitously. The source was the radio-loud QSO Q0957+561A\\&B. However, as demonstrated later\\cite{rat99} with HST, only relying on {\\em chance} to discover lensed sources leads to preferentially select {\\em faint} ones, which are then more difficult to follow up and to characterize and which are further embedded in complicated observational biases.  However, as first mentioned in the 80s\\cite{tog84}, {\\em luminous} quasars are excellent candidates to look for multiple images produced by isolated galaxies: the optical depth for lensing is high towards those distant sources {\\em and} their luminosity function is steep, introducing an amplification bias, which boosts the probability of lensing by a factor of 10 or more. Although about 100 such systems are presently known, HST has only contributed to one discovery through a QSO survey (i.e. Q1208+1011\\cite{bah92})! Indeed, ground based optical and radio QSO surveys (and recently the SDSS survey) were much more efficient and faster. But starting from here, HST has played a crucial role. ",
        "conclusions": ""
    },
    "0806/0806.2879_arXiv.txt": {
        "abstract": "{The compact association Cygnus OB2 is known to contain a large population of massive stars, but its total mass is currently a matter of debate. While recent surveys have uncovered large numbers of OB stars in the area around Cyg~OB2, detailed study of the optically brightest among them suggests that most are not part of the association.} {We observed an additional sample of optically faint OB star candidates,  with the aim of checking if more obscured candidates are correspondingly more likely to be members of Cyg~OB2.} {Low resolution spectra of 9 objects allow the rejection of one foreground star and the selection of four O-type stars, which were later observed at higher resolution. In a subsequent run, we observed three more stars in the classification region and three other stars in the far red.} {We identify five (perhaps six) new evolved very massive stars and three main sequence O-type stars, all of which are likely to be members of Cyg~OB2. The new findings allow a much better definition of the upper HR diagram, suggesting an age $\\sim2.5$~Myr for the association and hinting that the O3--5 supergiants in the association are blue stragglers, either younger or following a different evolutionary path from other cluster members. Though the bulk of the early stars seems to belong to an (approximately) single-age population, there is ample evidence for the presence of somewhat older stars at the same distance.} {Our results suggest that, even though Cyg~OB2 is unlikely to contain as many as 100 O-type stars, it is indeed substantially more massive than  was thought prior to  recent infrared surveys.} ",
        "introduction": "Among Galactic OB associations, Cyg OB2 is special in many respects. For a start,  it is known to host a large population of massive stars, including a significant fraction of the earliest spectral types in the Galaxy \\citep{wal02}. The optical extinction to Cyg OB2 is high, but not sufficiently so that it prevents spectra of its stars in the classification region being taken (something impossible for  other very massive open clusters with a large population of massive stars, such as Westerlund~1 \\citep{clark05} or the Arches Cluster \\citep{fig02}). Because of this, Cyg~OB2 is a very useful laboratory, since, on  one hand, it provides a large homogeneous population of OB stars that can be analysed \\citep{her99,her02} and, on the other, can be used as a template to compare optical and infrared investigations \\citep[e.g.,][]{hanson}. Finally, because of its compactness and high stellar content, Cyg~OB2 seems to occupy a more or less unique position somewhat intermediate between an open cluster and a normal OB association \\citep[cf.][]{kno00}.  These properties have led to a great deal of interest in Cyg~OB2, from the ``classical'' study of \\citet{jm54} to the comprehensive investigation by \\citet{mt91}, who identified $\\sim60$ stars more massive than $15M_{\\sun}$. More recently,  based on star counts in the 2MASS observations of the region, \\citet{kno00}  proposed that the number of O-type stars in Cyg OB2 was much larger. Building on this result, \\citet{com02} preselected a large number of possible OB members of Cyg~OB2 from their 2MASS colours and obtained low-resolution $H$- and $K$-band spectroscopy of the candidates. Candidates that lacked molecular bands were selected as very likely early-type stars. Of 77 candidates so selected, 31 stars for which optical spectra existed were OB stars, suggesting that most, if not all, of the other 46 objects were also OB stars in Cyg OB2.  From this list of candidates, \\citet{hanson} selected those brightest in the optical (14 objects with $B= 12$ to 14), for which she obtained classification spectra, finding that all of them were indeed  OB stars. However, \\citet{hanson} argues that most of these objects are not members of Cyg~OB2. For a start, they all lie at some distance from the previously defined boundaries of Cyg~OB2, as most of the sources located by \\citet{com02} do. Moreover, about half of the objects observed are late O and early B supergiants, indicating ages rather larger than the 2~Myr that \\citet{hanson} derives for Cyg~OB2 from  isochrone fitting to the location of the main sequence. Finally, one star (A39, B2\\,V) appears far too bright for its spectral type and is almost certainly a  foreground object.  It is therefore an open question as to whether the list of candidates from \\citet{com02} really contains a high fraction of actual Cyg~OB2 members. Here we investigate this issue with new spectra of several other fainter optical candidates. We also make use of the recent publication of a large catalogue of accurate spectral types for Cyg~OB2 members \\citep{kiminki}, which combined with our results and those of \\citet{hanson}, allows an enormous improvement in the characterisation of the HR diagram for the association.  In what follows, we will use the notation of \\citet{com02} for stars within their list (A\\#\\# for OB candidates and B\\#\\# for emission-line stars). For other members, we will use the numbering system of \\citet{mt91}, with prefix MT, except for the twelve stars with the classical numbering of \\citet{jm54}, which are given with the symbol \\# followed by their number. ",
        "conclusions": "Though the candidate sample of \\citet{com02} contains a high fraction of likely non-members, as discussed by \\citet{hanson}, it has also allowed the detection of a number of obscured O stars and very luminous B0\\,Ia supergiants very likely to be members of Cyg~OB2.  When these objects are included in the HR diagram, it becomes clear that there is a sequence of moderately evolved stars detaching from the main sequence exactly at the position where we stop seeing luminosity class V objects, i.e., around O6\\,V. These two facts combined support an age of $\\sim2.5$~Myr for the bulk of the association.  The HR diagram presented in Figure~\\ref{fig:hr} contains the largest number of Cyg~OB2 members ever displayed in such a diagram. It contains $\\sim 50$ stars that may have started their lives as main-sequence O-type stars and only  a few of these are unlikely to be members. Unless a population of extremely obscured O-type stars is lying at fainter magnitudes than probed by 2MASS, the total number of O-type stars in the association is almost certain to be in the 60\\,--\\,70 range.  The number of stars that have already left the main sequence and lie above the O6\\,V members that define the turn-off is more securely determined.  If the main association is basically co-eval, these represent the subset of stars  that were originally more massive than $35\\,M_{\\sun}$. Counting \\#12, which is not shown in Fig.~\\ref{fig:hr} because of its claimed spectral variability \\citep{kiminki}, there are 21 such stars. The evolved interacting binaries \\#5 and B17 (not in Fig.~\\ref{fig:hr}) should be counted too (perhaps doubly). The resulting number is certainly only a lower limit. Apart from possible unrecognised close doubles and binaries, at least two of the  Wolf-Rayet stars in the area are likely to be descendants of very massive stars \\citep{pas02}. Also, \\cite{com07} identify BD~$+53\\degr$3654 as a likely runaway O4\\,If member of the association. Therefore, we have identified a population of at least 25 stars that were originally more massive than $35\\,M_{\\sun}$. Given the uncertainties - in particular the very high binary fraction \\citep{kiminki} - we refrain from trying to derive a total mass for the association by assuming an IMF.  The brightest members, with spectral types in the O3--O5 range, may technically be considered blue stragglers. Though a real age difference cannot be ruled out, it does not seem to be borne out by the spatial distribution of stars, perhaps suggesting that we are seeing stars of similar mass evolve in very different ways.  Three luminous supergiants (\\#10 O9.5\\,Ia, A12 B0\\,Ia and A27 B0\\,Ia) seem to follow the 2.5~Myr isochrone and so appear to be the descendants of stars more massive than $\\sim40\\,M_{\\sun}$. This is in agreement with an initial mass estimate of $48\\,M_{\\sun}$ for \\#10 \\citep{her02}, which may have to be slightly reduced if the lower $DM=10.8$ is adopted. These objects will probably soon reach the LBV instability, which \\#12 is perhaps already encountering. A large population of O9--B1\\,Ia supergiants descended from stars with $M_{*}\\approx35\\,M_{\\sun}$ is found in the older ($\\sim4.5$~Myr) cluster Westerlund~1 together with a number of LBVs and Yellow Hypergiants \\citep{clark05}.  In summary, even if Cyg~OB2 falls short of the proposed 100 O-type stars by a factor of $\\sim 2$, its nuclear region still represents one of the most massive collections of early-type stars known in the Galaxy and its relatively low reddening cements its value as a laboratory for the study of their properties."
    },
    "0806/0806.0798_arXiv.txt": {
        "abstract": "Following an earlier analysis which examined the effect of the self-gravity of pressure on big-bang nucleosynthesis (BBN), we explore the effect of pressure's self-gravity on the structure of neutron stars.  We construct an {\\em ad hoc} modification of the Tolman-Oppenheimer-Volkoff equation wherein pressure's self-gravity is parameterized by a constant, $\\chi$, with $0 \\le \\chi \\le 1$.  The full general relativistic contribution to the gravity of pressure is recovered with $\\chi = 1$, and is eliminated when $\\chi = 0$.  This formulation is not proposed as an alternative theory of gravity, but is merely used to quantify the extent to which the self-gravity of pressure contributes to the structure of dense objects.  As can be surmised qualitatively, neutron star masses can be quite sensitive to $\\chi$, with higher values of neutron-star mass (by $\\sim$20--25\\%) allowed for smaller values of $\\chi$.  However, for a given equation of state, neither the range of neutron star radii nor the radii at fixed central density depend sensitively on $\\chi$.  Over the neutron star mass range measured so far, the presence or absence of pressure's self-gravity yields a nearly immeasurable change in radius --- much smaller than the variations in radius due to the uncertainty in the equation of state. In contrast to the result for BBN, we thus find that neutron stars are not likely to be useful testbeds for examining the self-gravity of pressure. ",
        "introduction": "One of the more fascinating predictions of general relativity is that pressure is self gravitating.  This is a strictly non-Newtonian result, whereby the pressure of a field contributes to gravity, increasing the effective density by an amount $3 P/c^2$ for a homogeneous perfect fluid.  There are four prominent physical situations in which the self-gravity of pressure could potentially lead to a measurable effect (three of which involve the expansion history of the Universe): (i) the acceleration of the Universe during inflation; (ii) the expansion of the Universe during the radiation-dominated epoch; (iii) the current acceleration of the Universe due to dark energy; and (iv) the mass-radius relation for neutron stars.  In an earlier paper, Rappaport et al.\\ \\cite{Rappaport08} addressed (ii) by setting a constraint on the self-gravity of pressure during the radiation-dominated epoch of big bang nucleosynthesis (BBN). They introduced an {\\it ad hoc} multiplicative parameter $\\chi$ to the $3 P/c^2$ contribution in the Einstein field equations.  For $\\chi = 1$, the full general relativistic contribution to the self-gravity of pressure is retained; for $\\chi = 0$, there is no such contribution. This yielded a modified set of Friedmann-like equations.  The contribution to $(\\dot a/a)^2$ for a species of matter with equation of state (EOS) $w = P/\\rho c^2$ appears as \\begin{equation} \\left(\\frac{\\dot a}{a}\\right)^2 = \\left(\\frac{1+3w\\chi}{1+3w}\\right)\\frac{\\Omega}{a^{1+3w}}\\;. \\label{eq:scaling} \\end{equation} The gravitational effect of radiation ($w = 1/3$) is therefore scaled by $(1+\\chi)/2$.  The absence of presure's self-gravity means that the Hubble constant at the time of BBN would be smaller by a factor of $\\sqrt{2}$, potentially changing the abundance of light elements. Rappaport et al.\\ used a standard BBN code and performed light element abundance calculations for a range of $\\chi$ to quantify this effect. When combined with current light element observations, the data were shown to be consistent with $\\chi = 1$, and strongly exclude $\\chi = 0$  Our goal here is to see whether similar limits can be set by observations of neutron stars.  At neutron star densities, $P \\sim (0.1 - 0.6)\\rho c^2$, suggesting that pressure is high enough that a measureable self-gravity effect might exist.  We introduce a parameterization of self-gravity very similar to that used to modify the Friedmann equations.  The result is a parameterized equation of stellar structure, with $\\chi = 1$ reproducing the Tolman-Oppenheimer-Volkoff equation, and $\\chi = 0$ ``turning off'' the pressure self-gravity of that equation.  Although at fixed central density $\\chi$ has a significant impact on a neutron star's mass, it has very little impact on its radius.  If measurements were to determine both mass and radius, it would still be extremely difficult to tell the difference between models with $\\chi = 0$ and $\\chi = 1$ due to the uncertainty in the neutron star EOS. Neutron stars thus appear to {\\it not} be very useful for testing the self-gravity of pressure.  This is not to say that one cannot make interesting statements about gravity with neutron stars {\\cite{psaltis07}}; but, in contrast to the situation with BBN, the pressure self-gravity aspect cannot be usefully tested.  The difference between the two cases is simple: the EOS of the universe during BBN is well understood, but the EOS of neutron stars is not.  Testing pressure's self gravity is thus degenerate with testing the EOS for neutron stars, but is not degenerate during BBN. ",
        "conclusions": ""
    },
    "0806/0806.0751_arXiv.txt": {
        "abstract": "{} {We determine the nature of the intermediate polar candidates \\object{XSS~J00564+4548}, \\object{IGR~J17195--4100}, and \\object{XSS~J12270--4859}.} {Pointed {\\em RXTE\\em} observations searched for intermediate polar characteristics in these candidate systems.} {\\object{XSS~J00564+4548} exhibits a period of 465.68$\\pm$0.07~s, which we interpret as the spin period, an energy dependent modulation depth, and a spectrum that is fit by a 22~keV photoelectrically absorbed bremsstrahlung with an iron line profile. \\object{IGR~J17195--4100} shows several candidate periodicities and a spectrum that is fit by a power law with an iron line. \\object{XSS~J12270--4859} exhibits a candidate spin period of 859.57$\\pm$0.64~s and a spectrum that is fit by a power law with no evidence of an iron line.} {\\object{XSS~J00564+4548} is confirmed to be an intermediate polar. \\object{IGR~J17195--4100} and \\object{XSS~J12270--4859} both show some properties of intermediate polars, but cannot be confirmed as definite members of the class here.} ",
        "introduction": "Intermediate polars (IPs) belong to the class of systems known as cataclysmic variables (CVs). They occupy the phase space, in terms of magnetic field strength, between the polars and the non-magnetic CVs. This intermediate strength magnetic field alters the accretion flow from the main sequence donor star to the white dwarf (WD). Eventually, most of the accreting material is channelled to accretion curtains above the WD magnetic poles. The temperature and density of this region causes the emission of bremsstrahlung radiation, which varies at the spin period of the WD. It is this variation that most consider to be the defining characteristics of IPs. For a review of IPs see e.g.~\\citet{warner95}. There are at least 30 confirmed IPs\\footnote{asd.gsfc.nasa.gov/Koji.Mukai/iphome/iphome.html as at 30/04/08}. \\citet{ramsay08}, however, have recently pointed out that the commonly used criteria to certify CVs as IPs may be too restrictive. It is possible that many of the 84 candidates\\footnotemark[1] are indeed IPs, and if classes such as SW~Sex systems are in fact IPs then the true number may be several hundred.  In recent years there have been 16 IPs found to emit in the hard X-ray/soft gamma-ray part of the spectrum, with the {\\em INTEGRAL}/IBIS survey \\citep{barlow06,bird07}. With this in mind we have embarked on a campaign to observe some hard X-ray sources and determine their credentials as potential IPs. In the first paper in this campaign, \\object{SWIFT~J0732.5--1331} was confirmed as an IP \\citep{butters07}. Here the results of pointed {\\em RXTE} observations of \\object{XSS~J00564+4548} (hereafter J0056), \\object{IGR~J17195--4100} (hereafter J1719) and \\object{XSS~J12270--4859} (hereafter J1227) are presented. ",
        "conclusions": "The unambiguous detection of an X-ray spin period of 465.68$\\pm$0.07~s in J0056 and its decreasing modulation depth with increasing energy, along with its spectral properties, confirm its inclusion into the IP class. Both J1719 and J1227 clearly exhibit some properties seen in IPs, but not to an extent for us to definitively classify them as such. We do note that it is likely these latter two are IPs, and that their true nature is being masked by the presence of contamination from other sources. X-ray imaging of these sources will definitively decide their fate, allowing their true spectral characteristics to be revealed. All three sources would benefit from long base line optical campaigns to determine their orbital periods and ratify the validity of the periods in J1719 and J1227. If J1227 does turn out to be an IP, then the presence of an X-ray iron feature will have to be reconsidered as a defining characteristic of IPs, since it is not present here."
    },
    "0806/0806.3491_arXiv.txt": {
        "abstract": "The Galactic black hole X-ray binary \\xte~entered a quiescent regime following the decline from the 2001--2002 outburst that led to its discovery. Here we report on the first detection of its quiescent counterpart in a 36 ks observation taken in 2007 July with the \\cxo~{\\it X-ray Observatory}. The inferred 0.5--10 keV unabsorbed flux is in the range $2.5$--$5.0\\times 10^{-15} $ \\ergscm. Notwithstanding large distance uncertainties, the measured luminosity is comparable to that of the faintest detected black hole X-ray binaries, all having orbital periods close to the expected bifurcation period between $j$- and $n$-driven low-mass X-ray binaries. This suggests that a few $10^{30}$ \\ergs~might be a limiting luminosity value for quiescent black holes. ",
        "introduction": "\\label{sec:intro}  Black hole (BH) X-ray transients -- close binary systems in which a low-mass donor transfers mass via Roche-lobe overflow onto a black hole accretor -- spend most of their lifetimes in a low-luminosity state, where the boundary between `quiescence' and more active regime can be set around $10^{33.5}$ erg \\se, corresponding to a few $10^{-6}$ times the Eddington luminosity (\\ledd) for a 10~\\msun~object (e.g. McClintock \\& Remillard 2006). Broadly speaking, the low Eddington ratios can be due to low radiative efficiency or low net accretion rate in the inner regions, or a combination of the two (see e.g. Narayan 2005, and references therein). First explored in their basic ideas by Ichimaru (1977), and Rees \\etal (1982), stable, radiatively inefficient accretion flow models were later formalized and gained vast popularity owing to the works of Narayan \\& Yi (1994, 1995), and Abramowicz \\etal (1995). Since the mid 1990s, they have been widely employed to account for the broadband (radio/optical/UV/X-ray) spectra of low-luminosity BH candidates, such as A0620--00, GS 1124--68 and V404 Cyg (Narayan \\etal 1996, 1997a), as well as the Galactic center source Sgr A$^{\\star}$ (Narayan \\etal 1995). The increased sensitivity of \\cxo~and \\xmm~with respect to previous X-ray telescopes has eventually permitted detailed X-ray studies of quiescent Galactic BHs down to Eddington ratios as low as a few $10^{-9}$ (Garcia \\etal 2001; Kong \\etal 2002; Hameury \\etal 2003; Tomsick \\etal 2003; McClintock \\etal 2003; Gallo \\etal 2006; Homan \\etal 2006; Corbel \\etal 2006; Bradley \\etal 2007). In the framework of advection-dominated accretion flow models (ADAF; Narayan \\& Yi 1994, 1995), the observed luminosity difference between quiescent BHs and quiescent neutron stars -- the former being fainter by about one order of magnitude at comparable orbital periods -- has been interpreted as evidence for the existence for an event horizon in BHs (Narayan \\etal 1997b; Menou \\etal 1999; Garcia \\etal 2001; Narayan \\& McClintock 2008). At the same time, recent studies at lower frequencies, in the radio and mid-infrared bands, suggest that BHs and neutron stars may channel different fractions of the total accretion power into relativistic outflows, with a substantially reduced jet contribution in quiescent neutron stars with respect to BHs (Fender \\etal 2003; Gallo \\etal 2006, 2007; Migliari \\etal 2006; K{\\\"o}erding \\etal 2007).  Of the 40 candidate BH X-ray binaries\\footnote{Additional likely BH candidates have been since discovered, including: XTE J1817--330, Swift J1753.5--0127, IGR J17091--3624, IGR J17098--3628, IGR J17497--2821 and IGR J18539+0727.} listed by Remillard \\& McClintock (2006), 15 have sensitive measurements/upper limits on their quiescent X-ray luminosities. In this Letter, we report on a 36 ks observation of the quiescent BH XTE J1650--500 performed in 2007 July with the \\cxo~{\\it X-ray Observatory}, and briefly discuss it in the context of quiescent BH X-ray binaries and how the advent of high sensitivity/resolution X-ray telescopes has improved our understanding of such systems.  The Galactic X-ray binary system \\xte~entered a quiescent regime following the 2001--2002 outburst that led to its discovery with the {\\it Rossi X-ray Timing Explorer} (Remillard 2001).  Observations conducted with \\xmm~and {\\it BeppoSAX} in late 2001, right after the outburst peak, revealed a broad, asymmetric Fe K$\\alpha$ emission line, interpreted as due to irradiation of the inner accretion disk around a rapidly spinning Kerr BH (Miller \\etal 2002; Miniutti \\etal 2004; see Done \\& Gierlinski 2006 for a different interpretation). The prolonged quiescent regime has allowed for the derivation of the system orbital period and optical mass function: P$_{\\rm orb}$=7.7 hr, and $f(M)$=2.73$\\pm$0.56 \\msun, respectively (Orosz \\etal 2004). The mass of the BH in \\xte~is highly uncertain. The amplitude of the phased $R$-band light curve results in a lower bound to the orbital inclination $i>50$\\degree$\\pm$3\\degree, which, in the limiting case of no disk contribution, implies in an upper limit of 7.3 \\msun~to the accretor mass (Orosz \\etal 2004). However --caveat the poor signal-to-noise of the adopted stellar templates-- these authors find that the accretion disk contribution in the $R$ band can be as high as $\\simgt$80$\\%$. This results in an orbital inclination $i=70$\\degree~or higher, that is a mass of only 4 \\msun~for the BH.  The non-detection of \\xte~with \\xmm, in 2005 March, yielded an upper limit on its quiescent unabsorbed flux of 1.1--1.3$\\times 10^{-14}$ \\ergscm~(0.5--10 keV), depending on the assumed spectral model (Homan \\etal 2006).  At a (by no means certain\\footnote{Similarly to its mass, the distance to \\xte~is also uncertain: the quoted value was estimated by Homan \\etal (2006) based on the system Eddington-scaled X-ray luminosity at the soft-to-hard X-ray state transition, following Maccarone \\& Coppi (2003). As such, it suffers from uncertainties of at least a factor 2.}) distance of 2.6 kpc, this corresponds to a X-ray luminosity \\lx$\\simlt 10^{31}$ \\ergs, comparable to the lowest luminosities inferred for quiescent BH X-ray binaries (i.e. GRO J1655--40, GRO J0422+32, XTE J1118+480, A0620--00 and GS 2000+25).    \\def\\errtwo#1#2#3{$#1^{+#2}_{-#3}$}  \\begin{deluxetable*}{crcccc} \\setlength{\\tabcolsep}{0.07in} \\tabletypesize{\\scriptsize} \\tablewidth{0pt} \\tablecaption{\\cxo~detected sources in the \\xte~field.\\label{tab:list}} \\tablehead{  \\colhead{Source} & \\colhead{Name} & \\colhead{R.A.} & \\colhead{Decl.} & \\colhead{Count rate} & \\colhead{Notes} \\\\ \\colhead{(1)} & \\colhead{(2)} & \\colhead{(3)} & \\colhead{(4)} & \\colhead{(5)} & \\colhead{(6)} } \\startdata 1  & CXOU J165007.6-495623 & 16 50 07.69 & -49 56 23.8 & 1.9 (0.7) \\\\ 2  & CXOU J165000.2-495723 & 16 50 00.21 & -49 57 23.4 & 1.9 (0.7) \\\\ 3  & {\\bf XTE J1650-500}   & 16 50 00.92 & -49 57 44.1 & 2.0 (0.7) \\\\ 4  & CXOU J165006.3-495814 & 16 50 06.37 & -49 58 14.3 & 2.0 (0.7) \\\\ 5  & CXOU J165013.1-495709 & 16 50 13.14 & -49 57 09.1 & 2.0 (0.7) \\\\ 6  & CXOU J165005.4-495427 & 16 50 05.48 & -49 54 27.6 & 2.1 (0.8) \\\\ 7  & CXOU J165002.5-495333 & 16 50 02.56 & -49 53 34.0 & 2.2 (0.8) \\\\ 8  & CXOU J164943.9-495901 & 16 49 43.97 & -49 59 01.3 & 2.2 (0.8) \\\\ 9  & CXOU J165012.1-495715 & 16 50 12.12 & -49 57 15.2 & 2.4 (0.8) \\\\ 10 & CXOU J165000.3-495655 & 16 50 00.38 & -49 56 55.9 & 2.5 (0.8) \\\\ 11 & CXOU J165006.0-495455 & 16 50 06.05 & -49 54 55.1 & 2.5 (0.8) & USNO B1  \\\\ 12 & CXOU J165021.3-495632 & 16 50 21.38 & -49 56 32.0 & 3.0 (0.9) \\\\ 13 & CXOU J164950.3-495931 & 16 49 50.30 & -49 59 31.0 & 3.1 (0.9) \\\\ 14 & CXOU J165006.3-495743 & 16 50 06.35 & -49 57 44.0 & 3.6 (1.0) \\\\ 15 & CXOU J165009.1-495442 & 16 50 09.10 & -49 54 42.9 & 4.2 (1.1) & USNO B1  \\\\ 16 & CXOU J164959.7-495518 & 16 49 59.77 & -49 55 18.7 & 4.2 (1.1) \\\\ 17 & CXOU J164958.5-495614 & 16 49 58.54 & -49 56 14.2 & 4.8 (1.2) \\\\ 18 & CXOU J164951.8-495653 & 16 49 51.87 & -49 56 53.8 & 5.2 (1.2) \\\\ 19 & CXOU J165013.9-495726 & 16 50 13.94 & -49 57 26.7 & 5.3 (1.2) \\\\ 20 & CXOU J164947.8-500119 & 16 49 47.82 & -50 01 19.9 & 5.4 (1.2) \\\\ 21 & CXOU J164955.5-495705 & 16 49 55.50 & -49 57 05.6 & 5.6 (1.2) \\\\ 22 & CXOU J165005.2-495622 & 16 50 05.27 & -49 56 22.6 & 5.9 (1.3) \\\\ 23 & CXOU J164943.2-495450 & 16 49 43.28 & -49 54 50.6 & 7.2 (1.4) \\\\ 24 & CXOU J165005.1-495624 & 16 50 05.13 & -49 56 24.3 & 7.2 (1.4) & USNO B1 \\\\ 25 & CXOU J164956.0-495711 & 16 49 56.04 & -49 57 11.2 & 8.0 (1.5) \\\\ 26 & CXOU J164953.5-495747 & 16 49 53.59 & -49 57 47.6 & 8.6 (1.5) \\\\ 27 & CXOU J164948.8-495509 & 16 49 48.82 & -49 55 09.8 & 9.4 (1.6) \\\\ 28 & CXOU J164948.7-495936 & 16 49 48.72 & -49 59 36.2 &11.5 (1.8) \\\\ \\enddata \\tablecomments{(1) Target numeration (2) Source name following the \\cxo~convention (3) Units of right ascension (equinox J2000.0) are hours, minutes, seconds; (4) Units of declination (equinox J2000.0) are degrees, arcminutes and arcseconds; (5) Net count rate, in units of $10^{-4}$ cps, as measured by {\\tt wavdetect}, with errors given in parenthesis, over the energy interval 0.3--7.0 keV. Notice that {\\tt wavdetect} is designed to be used as a detection algorithm, and only secondarily as a source flux measurement tool. Count rates are generally reliable, but can be slightly under-estimated for very low number of counts. } \\end{deluxetable*} ",
        "conclusions": "As illustrated in Figure~\\ref{fig:qbh} (updated from Tomsick \\etal 2003 after Corbel \\etal 2006, Gallo \\etal 2006 and Homan \\etal 2006), the quiescent X-ray luminosity of \\xte~as measured by \\cxo~sits in the range of values inferred for systems with similar orbital periods. Out of 15 candidate BH X-ray binary systems with sensitive observations while in the quiescent regime, 12 have now been detected in X-rays.  For those 12, the quiescent luminosities range between a few $10^{30}$ and $10^{33}$ \\ergs. The nearest BH, A0620--00, has been steadily emitting at $\\simeq 2-3\\times 10^{30}$ erg \\se at least for the past 5 years (Kong \\etal 2002; Gallo \\etal 2006); this is approximately the same luminosity level as XTE J1650--500 (this work), XTE J1118+480 (McClintock \\etal 2003) and GS~2000+25 (Garcia \\etal 2001), suggesting that this might be some kind of limiting value.  Indeed, for low-mass X-ray binaries one can make use of binary evolution theory, combined with a given accretion flow solution, to predict a relation between the minimum quiescent luminosity and the system orbital period \\porb~(see e.g. Menou \\etal 1999; Lasota 2000). As an example, Figure 4 of Menou \\etal (1999) illustrates how the predicted luminosity of quiescent BHs powered by ADAFs depends on the ratio between the outer mass transfer rate and the ADAF accretion rate.  The lower band, which roughly reproduces the observed luminosities of 3 representative systems spanning the whole range of detected luminosities (A0620--00, GRO J1655--40 and V404 Cyg), corresponds to $\\sim$1/3 of the mass transferred being accreted via the ADAF, whereas the remaining 2/3 would be either accumulated in an outer thin disk or lost to an outflow.  Most importantly, independently of the actual solution for the accretion flow in quiescence, the existence of a minimum luminosity in low-mass X-ray binaries stems directly from the existence of a bifurcation period, P$_{\\rm bif}$, below which the mass transfer rate is driven by gravitational wave radiation ($j$-driven systems), and above which it is dominated by the nuclear evolution of the secondary star ($n$-driven systems). Specifically, the outer mass transfer rate $\\dot{M}_{\\rm T}$ increases with decreasing orbital period below P$_{\\rm bif}$, while increases with increasing orbital period above P$_{\\rm bif}$ (the same applies to quiescent neutron star X-ray binaries, although with higher normalization). For a wide range of donor masses, Menou \\etal find P$_{\\rm bif}\\simeq 10$ hr. {\\it As long as the luminosity expected from a given accretion flow model scales with a positive power of $\\dot{M}_{\\rm T}$, systems with orbital periods close to the bifurcation period should display the lowest quiescent luminosities}. Caveat the large distance uncertainties which affect most BH X-ray binaries (e.g. Jonker \\& Nelemans 2004), this is indeed observed, as illustrated in the right panel of Figure~\\ref{fig:qbh}. \\begin{figure*}[t!] \\begin{center} \\includegraphics[angle=0,scale=.4]{f2a.eps} \\hspace{0.75cm}\\includegraphics[angle=0,scale=.4]{f2b.eps} \\caption{ {\\it Left}:  Quiescent X-ray luminosities/upper limits for 15 BH X-ray binaries with sensitive observations (down to a minimum threshold of $\\sim 10^{33.5}$ \\ergs). {\\it Right}: The measured luminosities are plotted against the systems' orbital periods. The faintest systems (GS~2000+25, A0620--00 and XTE J1650--500) have orbital periods close to 10 hr, the expected bifurcation period between $j$- and $n$-driven low-mass X-ray binaries according to Menou \\etal (1999). \\label{fig:qbh}} \\end{center} \\end{figure*}  Deep observations of short ($\\simlt$5 hr), as well as long ($\\simgt$200 hr) orbital period systems will hopefully add to the slim statistics in support of this argument. We wish to point out that, given the distance/absorption column to the 3 BH systems with upper limits (H1705--250, GRS1009--45 and 4U~1543--47; see Figure~\\ref{fig:qbh}), several tens of ks of integration with \\cxo~would be needed in order to detect them, should their luminosity be close to that of the faintest objects. Quiescent ultra-compact X-ray binaries (with orbital periods \\porb$\\simlt$1 hr) could be primary targets to test this conclusion at the short \\porb~end. However, the argument has been made that the mass-transfer rate evolution has not been studied in any detail for systems with such very low mass ratios. Extreme mass ratios could result in a drastically reduced $\\dot{M}_{\\rm T}$, making the direct comparison between quiescent ultra-compact and longer period~binaries somewhat dubious (see Jonker \\etal 2006, 2007, and related discussion in Lasota 2007). For long orbital period systems, the ideal testbed would be obviously provided by the power off of the as yet super-luminous GRS~1915+105, with an orbital period of over 800 hr. Similarly, the newly discovered BH candidate Swift J1753.5--0127, with an inferred orbital period of 3.2 hr (Zurita \\etal 2008), represents the ideal short period target."
    },
    "0806/0806.1800_arXiv.txt": {
        "abstract": "\\leftskip1.0cm \\rightskip1.0cm The conversion of neutron star to strange star is argued to be a two step process. The first process involves the deconfinement of nuclear to two-flavour quark matter. The GR results shows that the propagating front breaks up into fragments which propagate with different velocities along different directions. The time taken for this conversion to happen is of the order of few $ms$. This calculation indicates the inadequacy of non-relativistic (NR) or even Special Relativistic (SR) treatments for these cases. ",
        "introduction": "Strange Quark Matter (SQM), consisting of approximately equal numbers of up ({\\it u}), down ({\\it d}) and strange ({\\it s}) quarks, is conjectured to be the {\\it true} ground state of strong interaction \\cite{key-1,key-2}. SQM could naturally occur in the cores of compact stars, where central densities of about an order of magnitude higher than the nuclear matter saturation density. The transition from nuclear (hadronic) to quark matter should proceed through a conversion to an initial stage of (metastable) two flavour quark matter, which should decay to the stable SQM. Thus, neutron stars with sufficiently high central densities ought to get converted to strange, or at least hybrid, stars. These transitions could have observable signatures in the form of a jump in the breaking index and gamma ray bursts  \\cite{key-3,key-4}. \\par There are several plausible scenarios where neutron stars could convert to quark stars, through a \"seed\" of external SQM \\cite{key-5}, or triggered by the rise in the central density due to a sudden spin-down in older neutron stars \\cite{key-6}. Several authors have studied the conversion of nuclear matter to strange matter under different assumptions \\cite{key-7,key-11,key-17}. \\par In our recent work \\cite{key-18}, we have argued that the conversion process is a two step process. The first process involves the deconfinement of nuclear to two-flavour quark matter and the conversion process takes some milliseconds to occur. GR effects in such processes is studied by us \\cite{key-18a}. In this article, we make a detailed study of this process. ",
        "conclusions": "Having constructed the density profile of the star for a fixed central density, the respective flow velocities $v_{1}$ and $v_{2}$ of the nuclear and quark matter in the rest frame of the front, at a radius infinitesimally close to the center of the star. This would give us the initial velocity of the front ($-v_{1}$), at that radius, in the nuclear matter rest frame. With this, we integrate eqn. (2) outwards along the radius of the star. The solution gives the variation of the velocity with position as a function of the time of arrival of the front. Using this velocity profile, we can calculate the time required to convert the whole star using the relation $\\frac{dr}{d\\tau}=vG$.  \\begin{figure} \\vskip 0.4 in \\centering \\includegraphics[width=2.75in]{v_chi.eps} \\caption{Variation of velocity of the front along the radial direction for different  $\\chi$.} \\end{figure}  For a rotating star, due to the asymmetry, we introduce a new parameter $\\chi=cos\\theta$, along the vertical axis of the star. We start our calculation by choosing the central density of the star to be $7$ times the nuclear matter density, for which the Keplerian velocity of the star is $0.67 \\times 10^{-4} sec^{-1}$ (the rotational velocities given in fig. (2) are all in units of $10^{-4} sec^{-1}$). For this central density, the initial velocity of the front comes out to be $0.45  $. In fig. (1) the propagation velocity of the front along the radial direction of the star for three cases. The unbroken curve is for the SR case, the broken curve for non-rotating GR case and the dotted curve for the rotating GR case with $\\chi=0$, {\\it i.e} at the equator. Due to the asymptotic behaviour the velocity shoots up at the centre and saturates at larger radii. It can also be clearly seen that the GR effect increases the velocity of the front considerably (maximum by $30 \\%$) and the effect is most pronounced for the static case. The rotational effect of the star seems to suppress the GR effect and therefore the velocity of the front decreases. The result becomes clearer if we look at fig. (2) where we have plotted the front velocity with equatorial radius for different rotational velocities; as the rotational velocity increases, the velocity of the front decreases. \\begin{figure} \\vskip 0.4 in \\includegraphics[width=2.75in]{time.eps} \\caption{Variation of time of arrival of the front at certain radial distance for different cases.} \\end{figure}  From fig. (3), we find that the front velocity is maximum along the polar direction and minimum along the equator. Therefore, at any particular instant of time, we may have a situation where the polar part of the star has been converted while along the equatorial direction, the front is still propagating. \\par From fig. (4) we can see that the time taken by the conversion front to convert the neutron star to two-flavour quark star is of the order of few $ms$. The static star takes the minimum time ($3.3 ms$) whereas the rotating star takes the maximum time ($5.1 ms$) due to the enlarged equatorial radius. The polar part of the star needs much lesser time for conversion ($3.1 ms$), even less than static star, as its radius gets compressed.  \\par To summarize, we have shown in this article that the conversion of nuclear matter to quark matter in compact stars, especially rotating stars which are more realistic than static stars, is strongly affected by GR effects. The emergence of different conversion fronts, propagating with different velocities along different radial directions was not anticipated by Newtonian or SR calculations. It remains to be explored whether the incorporation of dissipative effects materially changes the results. Though the calculation is much involved, such an investigation is on our immediate agenda. \\par  \\noindent"
    },
    "0806/0806.4778.txt": {
        "abstract": "{ We present the optical identifications and a multi-band catalogue of a sample of 478 X-ray sources detected in the \\xmm~  and \\chandra~ surveys of the central 0.6 deg$^2$ of the ELAIS-S1 field. The most likely optical/infrared counterpart of each X-ray source was identified using the chance coincidence probability in the R and IRAC 3.6 \\micron~ bands. This method was complemented by the precise positions obtained through \\chandra~ observations. We were able to associate a counterpart to each X-ray source in the catalogue. Approximately 94\\% of them are detected in the R band, while the remaining are blank fields in the optical down to R$\\sim$24.5, but have a near-infrared counterpart detected by IRAC within 6\\arcsec~ from the \\xmm~ centroid. The multi-band catalogue, produced using the positions of the identified optical counterparts, contains  photometry in ten photometric bands, from B to the MIPS 24 \\micron~ band. The spectroscopic follow-up allowed us to determine the redshift and classification for 237 sources ($\\sim 50 \\%$ of the sample) brighter than $R=24$. The spectroscopic redshifts were complemented by reliable photometric redshifts for 68 sources. We classified 47\\% of the sources with spectroscopic redshift as broad-line active galactic nuclei (BL AGNs) with z$=0.1-3.5$, while sources without broad-lines (NOT BL AGNs) are about 46\\perc~ of the spectroscopic sample and are found up to $z=2.6$. The remaining fraction is represented by extended X-ray sources and stars. We spectroscopically identified 11 type 2 QSOs among the sources with \\xo$>$8, with redshift between 0.9 and 2.6, high 2-10 keV luminosity  (log\\Lx$\\ge$43.8 \\ergs) and hard X-ray colors suggesting large absorbing columns at the rest frame (log$N_H$ up to 23.6 cm$^{-2}$). BL AGNs show on average blue optical-to-near-infrared colors, softer X-ray colors and X-ray-to-optical colors typical of optically selected AGNs. Conversely, narrow-line sources show redder optical colors, harder X-ray flux ratio and span a wider range of X-ray-to-optical colors. On average the Spectral Energy Distributions (SEDs) of high-luminosity BL AGNs resemble the power-law typical of unobscured AGNs. The SEDs of NOT BL AGNs are dominated by the galaxy emission in the optical/near-infrared, and show a rise in the mid-infrared which suggests the presence of an obscured active nucleus. We study the infrared-to-optical colors and near-infrared SEDs to infer the properties of the AGN host galaxies. %On average the hosts of high-luminosity (log\\Lx$>$43.0 \\ergs) NOT BL AGNs show redder R-[3.6 \\micron] colors, compared to NOT BL AGNs with $41.8<$log\\Lx$<$43.0 \\ergs. ",
        "introduction": "One of the primary goals of observational cosmology is the determination of the census of the \\agn~ population in the Universe, along with their cosmic evolution and the assembly of super-massive black-holes (SMBHs) in galaxy nuclei.  Several theoretical and observational results indicate that the assembly of SMBHs is tightly linked to the evolution of the galaxy bulge component.  The discovery of SMBHs in the center of most nearby bulge-dominated galaxies and the correlation existing between the black-hole mass and the bulge properties (the $M_{BH}-\\sigma$ relation, \\citet{Gebhardt2000}, \\citet{Ferrarese2000}) suggest that the assembly of bulge masses is tied to the evolution of the accretion processes in AGNs.  In such scenario, the details of the co-evolution of black-holes and their host galaxies depend on feedback mechanisms between the AGNs and their host galaxies (e.g. \\citet{Granato2001,Granato2004}, \\citet{DiMatteo2005}, \\citet{Menci2006}, \\citet{Croton2006}). Luminous AGNs are more efficient in inhibiting the star-formation in their host galaxies, heating the interstellar matter through winds, shocks, and high energy radiation, thus making their colors redder. In this picture, the AGN phase should precede the phase when a galaxy is caught in a passive phase. Indeed, \\citet{Pozzi2007} and \\citet{Mignoli2004} using Spitzer photometry found that a sample of highly obscured (log$N_H=22.5-23.5$ cm$^{-2}$) QSOs at z=1-2 are hosted by red passive galaxies, suggesting a later stage in their evolution. Spectroscopy of type 2 QSOs with the Infrared Spectrograph (IRS) onboard \\emph{Spitzer} finds similar results (\\citet{Weedman2006}, \\citet{Houck2005}). This would suggest that the red colors of the IR selected heavily obscured AGNs (N$_H \\ge 10^{23}$ cm$^{-2}$) may be associated to a passive host galaxy. The study of the optical and infrared colors of AGN host galaxies can therefore put constraints on AGN feedback mechanisms and on their relative time-scales. This study is of course easier in the cases where the nuclear light does not over-shine the stellar light, i.e. in optically obscured AGNs.  Hard X-ray surveys performed by \\chandra~ and \\xmm~ in 2-10 keV band are the primary and most efficient tool to detect unobscured and moderately obscured (N$_H$ up to a few 10$^{23}$ cm$^{-2}$) AGNs up to high redshift, as they are less affected by dust and gas obscuration, compared to optical surveys and soft X-ray surveys.  In particular, deep pencil-beam surveys and shallower, large area surveys have been recently used to select and study the hard X-ray population, which dominates the cosmic X-ray background (see \\citet{Brandt2005} and references therein). These surveys are particularly efficient in selecting optically obscured AGNs, which show X-ray-to-optical colors (\\xo\\footnote{$log \\xo = log f_X + R/2.5 +5.4176$}, hereafter X/O) larger than 10, and are usually missed by optical selection (see e.g. \\citet{Fiore2003}, \\citet{Cocchia2007}, \\citet{Caccianiga2007}). The majority of these sources are thought to be obscured QSO at z$\\gtrsim$1, but many of them remain spectroscopically unidentified due to the faintness of their optical counterparts, which makes them difficult to access even with 10 m class telescopes.  Several studies suggest that the amount of obscuration decreases with increasing intrinsic luminosity (e.g., \\citet{LaFranca2005},\\citet{Ueda2003}). However, even the 2-10 keV selection becomes highly incomplete in selecting highly obscured and Compton-thick AGNs, with log\\nh$\\ge 24 cm^{-2}$.  The \\xmm~ survey of ELAIS-S1 is a large-area, medium depth survey, and therefore is particularly suited to select a large number of luminous AGNs. The ELAIS-S1 field covers an area of about 4 deg$^2$ in the southern hemisphere, and includes the minimum in the Galactic 100 $\\mu$m emission in that hemisphere (0.37 MJy/sr, \\citet{Schlegel1998}). A central 0.6 deg$^2$ contiguous area (centered at 00 34 40.4, -43 28 44.6) has been surveyed with \\xmm~ for a total of 400 ks. The regions with highest XMM sensitivity ($\\sim$65\\% of the full XMM area) were target of  6 \\chandra~ pointings (two of them centered on the same coordinates) with the aim of obtaining precise positions for the X-ray sources.  The field has a multi-band photometric coverage from the optical B band to the mid-infrared \\spitzer~ bands. A spectroscopic follow-up has been performed with VIMOS/VLT, FORS2/VLT and EFOSC/ESO3.6m for redshift determination and classification of the sources.  This paper presents the optical identifications and a multi-band catalogue of the X-ray sample described in \\citet{Puccetti2006}, along with the optical spectroscopy (redshift and source classification) for a sub-sample of sources, and an analysis of their multiwavelength properties. In this paper we concentrate on two main issues: 1) the identifications of a sizable sample of high luminosity, optically obscured QSOs; 2) the study of the optical and infrared properties of the host galaxies of the highly obscured AGNs.  The paper is organized as follows: Section 2 presents the sample, the multi-band photometry and optical spectroscopy.  Section 3 discusses source identification and presents the multi-band catalogue. Section 4 presents our results, redshifts, source classification and an analysis of the spectral energy distributions.  In Section 5 we present our conclusions. Magnitudes are given in Vega system unless otherwise stated. A $H_0=70$ km s$^{-1}$ Mpc$^{-1}$, $\\Omega_M$=0.3, $\\Omega_{\\Lambda}=0.7$ cosmology is adopted throughout. ",
        "conclusions": "We have presented the optical and infrared identifications of 478 X-ray sources detected by \\xmm~ in the central 0.6 deg$^2$ of the ELAIS-S1 field. The identification process was validated by precise (arcsec) source positions obtained with \\chandra~ in a fraction of the area covered by the \\xmm~ survey. We found  that \\chandra~ observations are crucial to identify the correct counterpart for optically faint sources. Comparing the \\chandra~ identifications with the IRAC ones, we find that the latter miss only 4\\perc~ of the real counterparts. Therefore, we feel confident to use the IRAC identifications for the area not covered by \\chandra.  We compiled a multiwavelength catalogue, with photometric data ranging from the mid-infrared to the optical bands. We obtained optical spectra to measure redshifts and to obtain a first classification of the counterparts. The spectroscopy was performed using VIMOS/VLT, FORS2/VLT and EFOSC/ESO3.6m. We obtained reliable redshifts and classification for 237 sources with optical counterparts brighter than  R$=$24. 47\\perc~ of the sample turned out to be broad-line AGNs, while the other sources are narrow-line AGNs (12\\perc), ELGs (27\\perc) and absorption-line galaxies (8\\perc). We find 47 sources showing \\xo$>$8 (23\\perc~ of the 2-10 keV sample) . Out of the 16 spectroscopically identified sources with \\xo$>$8, 11 ($\\sim$70\\perc) turned out to be type 2 QSOs at z=[0.9-2.6], with log\\Lx$\\ge$43.8 \\ergs. All these type 2 QSO have hard X-ray colors, suggesting large absorbing columns at the rest frame. The ELAIS-S1 sample therefore confirms that type 2 QSOs are efficiently selected among the sources with high \\xo.  We classified  empirically the broad band SEDs of the X-ray sources, from pure power-laws to galaxy-dominated SEDs. We find a generally good agreement between the SED classification and the optical spectroscopy. 80\\perc~ of the BL AGNs have power-law SEDs, and  $\\sim$71\\perc~ of the sources with  a power-law SED and a spectroscopic redshift are classified as BL AGNs.  Broad band SEDs have also been used to compute photometric redshifts. Reliable photometric redshifts  were obtained for a sample of 68 sources without a spectroscopic redshift and which show a SED dominated by the host galaxy stellar light.  We computed average rest-frame SEDs, finding that BL and NOT BL AGNs show similar L$_{10 \\mu m}$/L(2-10 keV) ratios ($\\sim$0.4 in logarithm), and consistent with \\citet{Pozzi2007}.  By comparing near-infrared colors, we find that low luminosity NOT BL AGNs (log\\Lx$<$43 \\ergs) are hosted in star forming galaxies, which show bluer rest-frame U-V and R-[3.6 \\micron] colors. High-luminosity NOT BL AGNs hosts (log\\Lx$>$43 \\ergs) have, on average, redder R-[3.6 \\micron] colors. This colors can be due either to a dusty star-forming host galaxy, or to a passive early-type host (\\citet{Pozzi2007}, \\citet{Mignoli2004}), but the quality of our photometry does not allow us to distinguish between the two. Deeper multiwavelength surveys, such as COSMOS and GOODS, are needed to asses the nature of these objects."
    },
    "0806/0806.3017.txt": {
        "abstract": "I work out  the Newtonian and general relativistic effects due to an isotropic mass loss $\\dot M/M$ of a body on the orbital motion of a test particle around it; the present analysis is also valid  for a variation $\\dot G/G$ of the Newtonian constant of gravitation. Concerning the Newtonian case, I use the Gauss equations for the variation of the elements and obtain negative secular rates for the osculating semimajor axis $a$, the eccentricity $e$ and the mean anomaly $\\mathcal{M}$, while the argument of pericenter $\\omega$  does not experience secular precession; the longitude of the ascending node $\\Omega$ and the inclination $i$ remain unchanged as well. The anomalistic period is different from  the Keplerian one: it turns out to be larger than it. The true orbit, instead, expands, as shown by a numerical integration of the equations of motion with MATHEMATICA; in fact, this is in agreement with the seemingly counter-intuitive decreasing of $a$ and $e$ because they refer to the osculating Keplerian ellipses which approximate the trajectory at each instant.  A comparison with the results obtained with different approaches by other researchers is made. % General relativity induces positive secular rates of the semimajor axis  and the eccentricity completely negligible in the present and future evolution of the solar system. % %By assuming for the Sun $\\dot M/M = -9\\times 10^{-14}$  yr$^{-1}$ it turns out that the Earth's perihelion position is displaced outward by 1.3 cm %along the fixed line of apsides  after each revolution. By applying our results  to the phase in which the radius of the Sun, already moved to the Red %Giant Branch of the Hertzsprung-Russell Diagram, will become as large as 1.20 AU in about $1$ Myr, I find that the Earth's perihelion position on the %fixed line of the apsides  will increase by  $\\approx 0.22-0.25$ AU (for $\\dot M/M = -2\\times 10^{-7}$ yr$^{-1}$); other researchers point towards an %increase of $0.37-0.63$ AU. Mercury will be destroyed already at the end of the Main Sequence, while Venus should be engulfed in the initial phase of %the Red Giant Branch phase; the orbits of the outer planets will increase by $1.2-7.5$ AU.  Simultaneous long-term numerical integrations of the %equations of motion of all the major bodies of the solar system, with the inclusion of a mass-loss term in the dynamical force models as well, are %required to check if the mutual N-body interactions may substantially change the picture analytically outlined here, especially in the Red Giant %Branch phase in which Mercury and Venus may be removed from the integration. %%Thus, even without invoking tidal effects and drag, the Earth should not avoid the engulfment in the expanded solar photosphere. ",
        "introduction": "In this paper I investigate the classical orbital effects induced by an isotropic variation $\\dot M/M$ of the mass of a central body  on the motion of a test particle; my analysis is valid also for a change $\\dot G/G$ of the Newtonian constant of gravitation. This problem, although interesting in itself, is not only an academic one because of the relevance that it may have on the ultimate destiny of planetary companions in many stellar systems in which the host star experiences a mass loss, like our Sun \\cite{Sch08}. With respect to this aspect, my analysis may be helpful  in driving future researches towards the implementation of long-term N-body simulations including the temporal change of $GM$ as well, especially over timescales covering paleoclimate changes, up to the Red Giant Branch (RGB) phase in which some of the inner planets should  be engulfed by the expanding Sun. Another problem, linked to the one investigated here, which has recently received attention is the observationally determined secular variation of the Astronomical Unit \\cite{Kra04,Sta05,Nor08,Kli08}. Moreover, increasing accuracy in astrometry pointing towards microarcsecond level \\cite{IAU07}, and long-term stability in clocks \\cite{Osk06} require to consider the possibility that smaller and subtler perturbations will be soon detectable in the solar system.  Also future planetary ephemerides should take into account $\\dot M/M$. Other phenomena which may, in principle, show connections with the problem treated here are the secular decrease of the semimajor axes of the LAGEOS satellites, amounting to 1.1 mm d$^{-1}$, \\cite{Ruby} and the increase of the lunar orbit's eccentricity of $0.9\\times 10^{-11}$ yr$^{-1}$ \\cite{luna}. %However, a detailed analysis of all such issues is beyond the scope of this paper.  Many treatments of the mass loss-driven  orbital dynamics in the framework of the Newtonian mechanics, based on different approaches and laws of variation of the central body's mass, can be found in literature; see, e.g., \\cite{Stro,Jea24,Jea29,Arm53,Haj63,Haj66,Khol,Dep83,Kev96,Kra04,Nor08} and references therein. However, they are sometimes rather confused and involved, giving unclear results concerning the behavior of the Keplerian orbital elements and the true orbit.  The plan of the paper is as follows. Section \\ref{minchia} is devoted to a theoretical description of the phenomenon in a two-body scenario. By working in the Newtonian framework, I will  analytically work out the changes after one orbital revolution experienced by all the Keplerian orbital elements of a test particle moving in the gravitational field of a central mass experiencing a variation of its $GM$ linear in time. Then, I will clarify the meaning of the results obtained by  performing a numerical integration of the equations of motion in order to visualize the true trajectory followed by the planet.  Concerning the method adopted, I will use the Gauss perturbation equations \\cite{Bert,roy}, which are valid for generic disturbing accelerations depending on position, velocity and time,  the \\virg{standard} Keplerian orbital elements (the Type I according to, e.g., \\rf{Khol}) with the eccentric anomaly $E$ as \\virg{fast} angular variable. Other approaches and angular variables like, e.g. the Lagrange perturbation equations \\cite{Bert,roy}, the Type II orbital elements \\cite{Khol} and the mean anomaly $\\mathcal{M}$ could be used, but, in my opinion, at a price of major conceptual and computational difficulties\\footnote{Think, e.g.,  about the cumbersome expansions in terms of the mean anomaly and the Hansen coefficients, the subtleties concerning the choice of the independent variable in the Lagrange equations for the semimajor axis and the eccentricity \\cite{Bert}.}. With respect to possible connections with realistic situations, it should be noted that, after all, the Type I orbital elements are usually determined or improved in standard data reduction analyses of the motion of planets and (natural and artificial) satellites. Instead, my approach should, hopefully, appear more transparent and easy to interpret, although, at first sight, some counter-intuitive results concerning the semimajor axis and the eccentricity will be obtained; moreover, for the chosen time variation of the mass of the primary, no approximations are used in the calculations which are quite straightforward.  However, it is important to stress that such allegedly puzzling features are only seemingly paradoxical because they will turn out to be in agreement with numerical integrations of the equations of motion, as explicitly shown by the numerous pictures depicted.  Anyway, the interested reader is advised to look also at \\rf{Khol} for a different approach. % In Section \\ref{due} I will work within the general relativistic gravitoelectromagnetic framework by calculating the gravitoelectric effects on all the Keplerian orbital elements of a freely falling test particle in a non-stationary gravitational field. %In Section \\ref{evol} I will apply our results to the future Sun-Earth scenario and to the other planets of the solar system. Section \\ref{tre} is devoted to a discussion of the findings of other researchers and contains some numerical calculations concerning the previously mentioned orbital phenomena of LAGEOS and the Moon. Section \\ref{quattro} summarizes my results. % %I wish to make a final remark concerning the field of applicability of our results to  realistic astrophysical contexts. Indeed, throughout the paper %I will consider only a two-body configuration in which the primary undergoes a time-variation of its $GM$. If I want to apply this scenario to the %evolution of the Sun-Earth system over timescales of the order of 0.1-1 Myr and more it should be taken into account that, in principle, also the %other planets  induce relatively large changes in the eccentricity (and the other orbital parameters) of the terrestrial orbit (see \\cite{KK} and %references therein; \\cite{Lask08}). Simulations looking back in time have shown that this happens on timescales of the order of just 0.1 Myr, and it %even appears to be an important forcing factor for climate changes \\cite{Lask04}. Thus, in extending our results to deep-future scenarios,  I might %be wrong, in principle, about how representative the present-day Earth's eccentricity is for any very long timescale (as I will show, the magnitude %of the changes depends on the eccentricity).    % % ",
        "conclusions": "\\lb{quattro} I started in the framework of the two-body Newtonian dynamics by using a radial perturbing acceleration linear in time and straightforwardly treated it with the standard Gaussian scheme. I found that  the osculating semimajor axis $a$, the eccentricity $e$ and the mean anomaly $\\mathcal{M}$ secularly decrease while the argument of pericentre $\\omega$ remains unchanged; the longitude of the ascending node $\\Omega$ and the inclination $i$ are not affected by the phenomenon considered. %Moreover, the Keplerian period $P^{\\rm Kep}$ decreases; it is no longer coincident with the apsidal period $P$ which turns out to be larger than %$P^{\\rm Kep}$ and increasing. The radial distance from the central body, taken on the fixed line of the apsides, experiences a secular increase $\\Delta r$. For the Earth, such an effect amounts to about 1.3 cm yr$^{-1}$. By numerically integrating the equations of motion in Cartesian coordinates I found that the real orbital path expands after every revolution, the line of the apsides does not change and the apsidal period is larger than the unperturbed Keplerian one. I have also  clarified that such results are not in contrast with those analytically obtained for the Keplerian orbital elements which, indeed, refer to the osculating ellipses approximating the true trajectory  at each instant. % I also computed the orbital effects of a secular variation of the Sun's mass in the framework of the general relativistic linearized gravitoelectromagnetism  which predicts a perturbing gravitoelectric tangential force proportional to  $\\bds{v}/r$. I found that both the semimajor axis and the eccentricity secularly increase; the other Keplerian elements remain constant. Such effects are completely negligible in the present and future evolution of, e.g., the solar system. % %I applied our results to the evolution of the Sun-Earth system in the distant future with particular care to the phase in which the Sun, moved to the %RGB of the HR, will expand up to 1.20 AU in order to see if the Earth will avoid to be engulfed by the expanded solar photosphere. Our answer is %negative because, even considering a small acceleration in the process of the solar mass-loss, it turns out that at the end of such a dramatic phase %lasting about $1$ Myr the perihelion distance will have increased by only $\\Delta r\\approx 0.22-0.25$ AU, contrary to the estimates by \\rf{Sch08} %who argue an increment of about $0.37-0.63$ AU. In the case of a circular orbit, the osculating semimajor axis remains unchanged, as confirmed by a %numerical integration of the equations of motion which also shows that the true orbital period increases and is larger than the unperturbed Keplerian %one which remains fixed. Concerning the other planets, while Mercury will be completely engulfed already at the end of the MS, Venus might survive; %however, it should not escape from its fate in the initial phase of the RGB in which the outer planets will experience  increases in the size of their %orbits of the order of $1.2-7.5$ AU. %%Finally, let us recall that in order to have more accurate results concerning the deep-future evolution of the Earth-Sun system,  also the action of %%the other planets of the solar system should be, in principle, accounted for in numerical simulations including the Sun's mass-loss as well.  As a suggestion to other researchers, it would be very important to complement my analytical two-body calculation by performing  simultaneous long-term numerical integrations of the equations of motion of all the major bodies of the solar system by including a mass-loss term in the dynamical force models as well to see if the N-body interactions in presence of such an effect  may substantially change the picture outlined here. It would be important especially in the RGB phase in which the inner regions of the solar system should dramatically change."
    },
    "0806/0806.0679_arXiv.txt": {
        "abstract": "We consider the basic physical properties of matter forming a thin accretion disc in the static and spherically symmetric space-time metric of the vacuum $f(R)$ modified gravity models. The Lagrangian of the generalized gravity theory is also obtained in a parametric form, and the conditions of the viability of the model are discussed. The exact Schwarzschild type solution of the gravitational field equations in the $f(R)$ gravity contains a linearly increasing term, as well as a logarithmic correction, as compared to the standard Schwarzschild solution of general relativity, and it depends on four arbitrary integration constants. The energy flux and the emission spectrum from the accretion disk around the $f(R)$ gravity black holes are obtained, and they are compared to the general relativistic case. Particular signatures can appear in the electromagnetic spectrum, thus leading to the possibility of directly testing modified gravity models by using astrophysical observations of the emission spectra from accretion disks. ",
        "introduction": "Several recent astrophysical observations \\cite{Ri98} have provided the astonishing result that around $95$--$96\\%$ of the content of the Universe is in the form of dark matter $+$ dark energy, with only about $4$--$5\\%$ being represented by baryonic matter. More intriguing, around $70\\%$ of the energy-density may be in the form of what is called ``dark energy'', and is responsible for the acceleration of the distant type Ia supernovae \\cite {PeRa03}. Hence, today's models of astrophysics and cosmology face two severe problems, that can be summarized as the dark energy problem, and the dark matter problem, respectively. Although in recent years many different suggestions have been proposed to overcome these issues, a satisfactory answer has yet to be obtained.  A very promising way to explain these major problems is to assume that at large scales the Einstein gravity model of general relativity breaks down, and a more general action describes the gravitational field. Theoretical models in which the standard Einstein-Hilbert action is replaced by an arbitrary function of the Ricci scalar $R$, first proposed in \\cite{Bu70}, have been extensively investigated lately. The only restriction imposed on the function $f$ is to be analytic, that is, it must possess a Taylor series expansion about any point. Cosmic acceleration can be explained by $f(R)$ gravity \\cite{Carroll:2003wy}, and the conditions of viable cosmological models have been derived in \\cite{viablemodels}. In the context of the Solar System regime, severe weak field constraints seem to rule out most of the models proposed so far \\cite{solartests,Olmo07}, although viable models do exist \\cite{Hu:2007nk,solartests2,Sawicki:2007tf,Amendola:2007nt}. The possibility that the galactic dynamic of massive test particles can be understood without the need for dark matter was also considered in the framework of $\\ f(R)$ gravity models \\cite {Cap2,Borowiec:2006qr,Mar1,Boehmer:2007kx,Bohmer:2007fh}, and connections with MOND and the Pioneer anomaly further explored by considering an explicit coupling of an arbitrary function of $R$ with the matter Lagrangian density \\cite{Bertolami:2007gv,Bertolami:2007vu}. For a recent review of the $f(R)$ modified gravity models see \\cite{rev}.  The study of the static spherically symmetric vacuum solutions of the gravitational field theories is fundamental for the physical understanding and interpretation of the model. In particular, the vacuum solutions provide the theoretical basis for the Solar System testing of the theories, and for the description of the motion of the test particles around massive bodies. It was shown that for a large class of models, including e.g.~the $f(R)=R-\\mu ^{4}/R$ model, the Schwarzschild-de Sitter metric is an exact solution of the field equations.  Solutions in the presence of a perfect fluid were analyzed in \\cite{Multamaki:2006ym}. Other approaches in searching for exact spherically symmetric solutions of $f(R)$ theories of gravity were studied in \\cite {Capozziello:2007wc}, respectively.  Several exact vacuum static and spherically symmetric solutions of the gravitational field equations in $f(R)$ gravity were obtained in \\cite{Multamaki:2006zb}. The set of the modified Einstein's field equations were reduced to a single, third order differential equation, and it was shown how one can construct exact solutions in different $f(R)$ models. In particular, a Schwarzschild type solution of the field equation was constructed. This solution, containing a term linearly increasing with the radial coordinate, as well as a logarithmic term, depends on four new arbitrary integration constants, and reduces to the standard general relativistic case by imposing the zero value to one integration constant, and by appropriately choosing the numerical values of the other constants. However, even that this choice of constants allows the model to pass the solar system tests, we expect that at large distances, or in the presence of strong gravitational fields, the geometry of the space-time in $f(R)$ modified gravity models is different from that of standard general relativity. Therefore, it is important to find a method that could allow one to observationally distinguish, and test in an astrophysical setting, the possible deviations from Einstein's theory. One such possibility is the study of accretion disks around compact objects.  Most of the astrophysical bodies grow substantially in mass via accretion. Recent observations suggest that around almost all of the active galactic nuclei (AGN's), or black hole candidates, there exist gas clouds surrounding the central compact object, together with an associated accretion disc, on a variety of scales from a tenth of a parsec to a few hundred parsecs \\cite{UrPa95}. These gas clouds are assumed to form a geometrically and optically thick torus (or warped disc), which absorbs most of the ultraviolet radiation and the soft X-rays. The gas exists in either the molecular or the atomic phase. The most powerful evidence for the existence of super massive black holes comes from the VLBI imaging of molecular $\\mathrm{H_{2}O}$ masers in the active galaxy NGC 4258 \\cite {Miyo95}. This imaging, produced by Doppler shift measurements assuming Keplerian motion of the masering source, has allowed a quite accurate estimation of the central mass, which has been found to be a $3.6\\times 10^{7}M_{\\odot }$ super massive dark object, within $0.13$ parsecs. Hence, important astrophysical information can be obtained from the observation of the motion of the gas streams in the gravitational field of compact objects.  The mass accretion around rotating black holes was studied in general relativity for the first time in \\cite{NoTh73}.  By using an equatorial approximation to the stationary and axisymmetric space-time of rotating black holes, steady-state thin disk models were constructed, extending the theory of non-relativistic accretion \\cite{ShSu73}. In these models hydrodynamical equilibrium is maintained by efficient cooling mechanisms via radiation transport, and the accreting matter has a Keplerian rotation. The radiation emitted by the disk surface was also studied under the assumption that black body radiation would emerge from the disk in thermodynamical equilibrium. The radiation properties of the thin accretion disks were further analyzed  in \\cite{PaTh74} and in \\cite{Th74}, where the effects of the photon capture by the hole on the spin evolution were presented as well. In these works the efficiency with which black holes convert rest mass into outgoing radiation in the accretion process was also computed.  Later on, the emissivity properties of the accretion disks were investigated for exotic central objects, such as non-rotating or rotating quark, boson or fermion stars  \\cite{Bom,To02,YuNaRe04}. The radiation power per unit area, the temperature of the disk and the spectrum of the emitted radiation were given, and compared with the case of a Schwarzschild black hole of an equal mass.  It is the purpose of the present paper to study the thin accretion disk models applied for black holes in $f(R)$ modified gravity models, and carry out an analysis of the properties of the radiation emerging from the surface of the disk.  The present paper is organized as follows. The $f(R)$ gravity generalization of the Schwarzschild type solution of general relativity is obtained in Section II. In Section III we review the formalism of the thin disk accretion onto compact objects. The basic properties of matter forming a thin accretion disc in the space-time metric of the $f(R)$ modified gravity models are considered in Section IV. We discuss and conclude our results in Section V. In the present paper we use a system of units so that $c=G=\\hbar =k_{B}=1$, where $k_{B}$ is Boltzmann's constant. ",
        "conclusions": "In the present paper we have considered the basic physical properties of matter forming a thin accretion disc in the Schwarzschild type vacuum space-time metric of the $% f(R)$ modified gravity models. The physical parameters of the disc - effective potential, flux and emission spectrum profiles - have been explicitly obtained for several values of the parameters characterizing the vacuum solution of the generalized field equations. All the astrophysical quantities, related to the observable properties of the accretion disc, can be obtained from the black hole metric. Due to the differences in the space-time structure, the modified $f(R)$ gravity black holes present some very important differences with respect to the disc properties as compared to the standard general relativistic Schwarzschild case.  The determination of the accretion rate for an astrophysical object can give a strong evidence for the existence of a surface of the object. A model in which Sgr A*, the $3.7\\times 10^{6}M_{\\odot }$ super massive black hole candidate at the Galactic center, may be a compact object with a thermally emitting surface was considered in \\cite{BrNa06}. For very compact surfaces within the photon orbit, the thermal assumption is likely to be a good approximation because of the large number of rays that are strongly gravitationally lensed back onto the surface. Given the very low quiescent luminosity of Sgr A* in the near-infrared, the existence of a hard surface, even in the limit in which the radius approaches the horizon, places a severe constraint on the steady mass accretion rate onto the source, ${\\dot{M% }}\\leq 10^{-12}M_{\\odot }$ yr$^{-1}$. This limit is well below the minimum accretion rate needed to power the observed submillimeter luminosity of Sgr A*, ${\\dot{M}}\\geq 10^{-10}M_{\\odot }$ yr$^{-1}$.  Thus, from the determination of the accretion rate it follows that Sgr A* does not have a surface, that is, it must have an event horizon. Therefore the study of the accretion processes by compact objects is a powerful indicator of their physical nature. Since, as one can see from Table~\\ref{Efficiency}, the conversion efficiency in the case of the $f(R)$ vacuum solutions is different as compared to the general relativistic case, the determination of this parameter could discriminate, at least in principle, between the different gravity theories, and constrain the parameters of the model."
    },
    "0806/0806.0486_arXiv.txt": {
        "abstract": "{It is generally presumed that the outflows from a YSO are directed close to its rotation axis (i.e. along its angular momentum vector and orthogonal to any attendant accretion disc). Many YSOs are formed from dense prestellar cores embedded in filaments, and therefore the relative orientations of outflows and filaments may place a useful constraint on the dynamics of core formation.} {{\\bf We explore this possibility, from the viewpoint of what it may tell us about the angular momentum delivered to a core forming in a filament. We stress that we are not here addressing the issue of the relationship of filaments and outflows to the prevailing magnetic field direction, although this is evidently also an interesting issue.}} {We use data from the literature and the SCUBA archive to estimate the projected angles between 45 observed outflows and the filaments which appear to contain their driving sources. The distribution of these angles is then compared with model predictions, so as to obtain a statistical constraint on the distribution of intrinsic angles between outflows and filaments.} {Using the limited data available, and neglecting any possible selection effects or correlations between nearby outflows, we infer that the observed outflows have a tendency to be orthogonal to the filaments that contain their driving sources. Specfically, in the cases where the directions of the filaments and outflows are well defined, we infer statistically that $72\\,\\%$ of outflows are within $45^{\\rm o}$ of being orthogonal to the filament, and only $28\\,\\%$ are within $45^{\\rm o}$ of being parallel to the filament.} {This suggests that the prestellar cores which spawned the YSOs driving the observed  outflows had angular momenta which were approximately orthogonal to the filaments out of which the cores formed. We briefly discuss the implications of this for two proposed core formation mechanisms.} ",
        "introduction": "The substructure within a star-forming molecular cloud often appears to be filamentary, and the dense prestellar cores out of which protostars condense are often embedded within filaments (e.g. Schneider \\& Elmegreen 1979; Hatchell et al. 2005; Johnstone \\& Bally 2006; Kirk, Johnstone \\& Tafalla 2007; Nutter \\& Ward-Thompson 2007; Kirk, Ward-Thompson \\& Nutter 2007; Muench et al. 2007; Goldsmith et al. 2008; Narayanan et al. 2008). The implication is that the growth of prestellar cores is fed mainly by material flowing in along filaments{\\bf, and -- although Hatchell et al. (2005) caution against it -- this interpretation is supported by many numerical simulations, (e.g. Passot, V\\'azquez-Semadeni \\& Pouquet 1995; Nagai, Inutsuka \\& Miyama 1998; Padoan \\& Nordlund 1999; Balsara, Ward-Thompson \\& Crutcher 2001; Klein, Fisher \\& McKee 2001; Gammie et al. 2003; Li, P. et al. 2004; Li, Z. \\& Nakamura 2004; Nakamura \\& Li 2005; Oishi \\& Mac Low 2006; Ciolek \\& Basu 2006; Kudoh et al. 2007; Kudoh \\& Basu 2008; Hennebelle et al. 2008; Offner, Klein \\& McKee 2008).} The highly anisotropic inflow {\\bf from a filament onto a core} may then have consequences for the dynamics of core collapse and fragmentation. In particular, the net angular momentum of the inflowing material will strongly influence the orientation of binary orbits and circumstellar accretion discs in the small-$N$ subcluster of protostars forming in a prestellar core.  During the embedded Class 0 and Class I phases of protostellar evolution (e.g. Di Francesco et al. 2007; Ward-Thompson et al. 2007), the orientations of circumstellar accretion discs can be inferred from the directions of the outflows which they drive. An outflow is presumed to be driven by torsional MHD waves propagating along magnetic field lines anchored in a circumstellar disc (e.g. Pudritz \\& Norman 1986). Therefore the outflow should be approximately parallel to the rotation axis of the circumstellar disc, i.e. along the angular momentum vector of the material forming the parental core.  It follows that the orientation of an outflow, relative to the filament in which its driving source has been born, may contain important information on the dynamics of core formation. Specifically it constrains the relationship between the inflow forming the core (if we assume that this inflow is concentrated along the filament) and the angular momentum which this material brings with it (if we assume that this angular momentum is oriented along the direction of the outflow). This information can be used to discriminate between different mechanisms for core formation (e.g. Whitworth et al. 1995; Banerjee, Pudritz \\& Anderson 2006; Banerjee \\& Pudritz 2007).  The present paper is organised as follows. In Section 2, we present the distribution of projected angles, $\\gamma$, between observed filaments and the outflows from YSOs that appear to be embedded in the filaments. We show that the distribution of projected angles is compatible with the distribution expected if outflows are approximately orthogonal to filaments; relative to the orthogonal direction, the intrinsic directions of outflows have a standard deviation $\\sigma_{\\psi}\\sim 45^{\\rm o}$, which implies that $72\\,\\%$ of outflows are within $45^{\\rm o}$ of being orthogonal to the filament which contains their driving source. In Section 3, we discuss two mechanisms for core formation, (i) gravitational fragmentation of a shock-compressed layer (e.g. Whitworth et al. 1995), and (ii) gravo-turbulent fragmentation (e.g. {\\bf Klein, Fisher \\& McKee 2001;} Banerjee, Pudritz \\& Anderson 2006; Banerjee \\& Pudritz 2007{\\bf ; Offner, Klein \\& McKee 2008}). We {\\bf infer} that the first mechanism is more likely to produce the observed distribution{\\bf , but this is not a very strong inference, and more observational data are required to make it more robust}. We also discuss the statistical significance of the result, and the possible influence of selection effects. The data used are presented in Appendix A. The distribution of projected angles expected is derived in Appendix B. ",
        "conclusions": "We have analysed the limited number of protostellar systems in local star formation regions where a YSO embedded in a discernable filament drives an outflow. In these cases, the distribution of angles between filaments and outflows appears to imply that outflows are usually approximately orthogonal to the filaments which contain (and presumably have fed) the cores which have spawned their driving YSOs. Given the chaotic and impulsive nature of multiple star formation in a prestellar core, one should expect some variance in the relative orientation of filaments and outflows. Specifically, it appears that $72\\,\\%$ of outflows are within $45^{\\rm o}$ of orthogonal to the filament which contains their driving YSO. This inference is not very robust, statistically, but, if confirmed, it would favour a model in which filaments were formed by gravitational fragmentation of shock compressed layers, as against one in which filaments were formed from shear instabilities in colliding hydrodynamic or magneto-hydrodynamic flows. Further observations of outflows from filaments are needed to validate this conclusion."
    },
    "0806/0806.4618.txt": {
        "abstract": "{The luminosity-metallicity relation is one of the fundamental constraints in the study of galaxy evolution; yet none of the relations available today has been universally accepted by the community. } {The present work is a first step to collect homogeneous abundances and near-infrared (NIR) luminosities for a sample of dwarf irregular (dIrr) galaxies, located in nearby groups. The use of NIR luminosities is intended to provide a better proxy to mass than the blue luminosities commonly used in the literature; in addition, selecting group members reduces the impact of uncertain distances. Accurate abundances are derived to assess the galaxy metallicity. } {Optical spectra are collected for \\ion{H}{ii} regions in the dIrrs, allowing the determination of oxygen abundances by means of the temperature-sensitive method. For each dIrr galaxy $H$-band imaging is performed and the total magnitudes are measured via surface photometry.} {This high-quality database allows us to build a well-defined luminosity-metallicity relation in the range $-13\\geq M_{H}\\geq-20$. The scatter around its linear fit is reduced to $\\RMS$~dex, the lowest of all relations currently available. There might exist a difference between the relation for dIrrs and the relation for giant galaxies, although a firm conclusion should await direct abundance determinations for a significant sample of massive galaxies.} {This new dataset provides an improved luminosity-metallicity relation, based on a standard NIR band, for dwarf star-forming galaxies.  The relation can now be compared with some confidence to the predictions of models of galaxy evolution. Exciting follow-ups of this work are (a) exploring groups with higher densities, (b) exploring nearby galaxy clusters to probe environmental effects on the luminosity-metallicity relation, and (c) deriving direct oxygen abundances in the central regions of star-forming giant galaxies, to settle the question of a possible dichotomy between the chemical evolution of dwarfs and that of massive galaxies. ",
        "introduction": "According to standard scenarios of galactic chemical evolution, a luminosity-metallicity (\\lz) relation is established through galactic winds -- induced by supernova explosions -- removing the interstellar medium (ISM) before it has been totally converted into stars (see the seminal work of Larson \\cite{larson74}; and also Tinsley \\& Larson \\cite{tinsley_larson79}, Dekel \\& Silk \\cite{dekel_silk86}, Lynden-Bell \\cite{lynden-bell92}, among others). As this process is more efficient in low-mass (low escape velocity) galaxies, less metals should be trapped in dwarf galaxies. Indeed, a well-defined \\lz\\ relation is observed in spheroidal galaxies (e.g. Caldwell et al. \\cite{caldwell_etal92}), mostly field galaxies, in which old stellar populations dominate, and for which the chemical evolution has stopped. Do we expect a similar behaviour in the case of dwarf irregular galaxies (dIrr)? The situation is more complex. First, dIrrs are still active (in the sense of star formation), exhibiting substantial gas fractions and a broad range of star-formation rates (SFR). So, one would expect to encounter in Irrs, ISMs of different chemical maturities and rapidly evolving luminosities: all factors which would loosen any established relation between mass and metallicity. Second, the available sample of dIrrs is made of galaxies belonging to groups rather than distributed in the field. It is well known that the evolution of a galaxy gas content depends on its environment, and that gas stripping is boosted through galaxy interactions in dense groups and cluster sub-clumps. So, while one can see reasons why the terminal chemical evolution in dIrrs might depend on the galactic mass -- similarly to the sample of gas-free spheroidal galaxies --, there are a number of extra processes blurring the situation. Hence, it is unclear whether dIrr galaxies should exhibit a luminosity-metallicity relation.  Under such a scenario dwarf galaxies lose substantial fractions of their mass in the course of their evolution, and should be major contributors to the chemical evolution of the inter-galactic medium (IGM). In particular, mass loss through galactic winds should occur in dIrr galaxies: their study contributes to probe the fraction of the IGM enrichment which is due to dwarf galaxies in general {(e.g., Garnett \\cite{garnett02})}.  Losses of enriched gas in dwarf galaxies have indeed been observed: Martin et al. (\\cite{martin_etal03}) detected a galactic wind of $\\sim6\\times10^{6}\\mathcal{M}{}_{\\odot}$ from NGC~1569, and for the first time, they could measure its metallicity. The authors concluded at $3\\times10^{4}\\mathcal{M}_{\\odot}$ of oxygen in the wind, almost as much as in the disk of the dwarf itself. Hence, the cosmic chemical evolution and the existence of an \\lz\\ relation appear to be closely inter-related: this was an additional motivation to pursue the study presented here.  The case of another dIrr, SagDIG, further illustrates this point. In one of our previous studies (Saviane et al. \\cite{sagdig}) we measured a very low oxygen content ($7.26\\leq12+\\log{\\textrm{(O/H)}}\\leq7.50$), a result which is not compatible with a closed-box evolution. According to the model, such a low abundance would imply a large gas mass fraction $\\mu=m_{{\\rm gas}}/(m_{{\\rm gas}}+m_{{\\rm stars}})\\approx0.97$, whereas the observed gas mass fraction is only $\\mu\\approx0.86$. It is easy to compute that $\\approx1.5\\times10^{6}\\mathcal{M}_{\\odot}$ of gas are missing, and, since this mass is smaller than that of the NGC~1569 galactic wind, it is plausible to conclude that SagDIG lost some of its gas into the IGM.  {On the other hand, a cautionary remark should be added, since a general consensus on the role of galactic winds in the evolution of dwarf galaxies and the IGM has not been reached yet. For example by means of chemo-photometric models Calura \\& Matteucci (\\cite{calura_matteucci06}) conclude that dIrr galaxies play a negligible role in the enrichment of the IGM, and the numerical models of Silich and Tenorio-Tagle (\\cite{silich_tagle98}) show that that galactic winds never reach the escape velocity of these dwarfs. And in the case of dwarf spheroidal (dSph) galaxies, the 3D hydrodynamic simulations of Marcolini et al. (\\cite{marcolini_etal06}) show that no galactic winds develop in these objects.}  In our attempt to place SagDIG on previously established \\lz\\ relations for dIrr galaxies, we realized that the existence of such a relation was controversial. Some studies had concluded at very well-defined correlations (Skillman et al. \\cite{skh89}; Richer \\& McCall \\cite{rm95}; Pilyugin \\cite{pilyugin01b}), others have found only mild relations with substantial scatter (e.g. Skillman et al. \\cite{skillman_scl_hii}), or no correlation at al (Hidalgo G\u00e1mez \\& Olofsson \\cite{anamaria}; Hunter \\& Hoffman \\cite{hunter_hoffman99}). Perhaps the major source of confusion in these investigations comes from the ill-defined samples used in the analysis. Abundance data are taken from different sources (with spectra of variable quality and processed through different reduction and analysis techniques); apparent luminosities and distances are taken from catalogs, meaning that they are largely approximate. Moreover, although the \\lz\\ relation is expected to depend on the environment, this is rarely taken into account in the analysis. Another main source of uncertainty is the common use of the blue absolute magnitude as a tracer of the mass. Optical luminosities can in fact be extremely misleading in Irr galaxies, because of their star-bursting activity (e.g. Tosi et al. \\cite{tosi_etal92} or Schmidt et al. \\cite{schmidt_etal95}). Already, Bruzual \\& Charlot (\\cite{bruzual_charlot83}) showed that at $400\\,{\\rm nm}$ a $1$~Myr old burst is $\\sim$ three orders of magnitude brighter than an underlying old ($15$~Gyr) stellar population of comparable mass. In other words, a dwarf galaxy hosting a recent starburst could be as luminous, in the blue, as a galaxy orders of magnitudes more massive but lacking a recent star formation episode. In comparison, a recent starburst is at most $10$ times brighter in the near-infrared (NIR) than its underlying old population of same mass. Therefore, the NIR window being more stable with regard to star formation episodes is more appropriate for probing the galaxy basic properties, such as its mass.  In order to {address these issues}, we embarked on a medium-term project aimed at gathering nebular oxygen abundances and NIR luminosities for a sample of dIrrs belonging to the three nearest groups of galaxies. In this way we can test the existence of an \\lz\\ relation in well-defined environments (characterized by the group density), for which the scatter in apparent distance modulus is low, and for which an homogeneous set of abundances can be obtained.  Since we started the project, a few investigations have been published, in relation with the context presented above. The first \\lz\\ relation using NIR luminosities (aside the work by Saviane et al. \\cite{sidney}) is that by Salzer et al. (\\cite{salzer05}; hereafter S05). Their sample is dominated by massive galaxies, and their relation was derived using 2MASS data for the luminosities, and proprietary spectra from the KPNO International Spectroscopic Survey (KISS) for the metallicities. The 2MASS survey having a relatively shallow limiting magnitude, for a few additional dwarf galaxies NIR photometry was supplemented by the authors. At the other extreme, the sample assembled by Mendes de Oliveira et al. (\\cite{mendes-de-oliveira_etal06}) is entirely made of dwarf galaxies. They gathered $K$-band luminosities and metallicities -- obtained through the direct method -- for $29$ dIrrs. The NIR data are from 2MASS or from Vaduvescu et al. (\\cite{vaduvescu_etal05}), and metallicities have been picked up from a variety of sources. Finally Lee et al. (\\cite{lee_etal06}) have used the \\emph{Spitzer} Infrared Array Camera to compute $4.5$~\\micron\\ luminosities for $\\sim30$ nearby dIrrs, the distance of which have been derived using standard candles. With oxygen abundances collected from the literature, they {constructed a} 4.5~\\micron\\ \\lz\\ relation.  Although these studies have certainly improved the situation, they do not represent the ideal case yet. S05 abundances are derived through a new but indirect method, and the sample is very scarce in dwarf galaxies. Yet, S05 is the only study which includes $H$-band photometry, so a comparison with their \\lz\\ relation is carried out later in this paper, in Sect.~\\ref{sec:dichot}. As extensively discussed in Sect.~\\ref{subsec:Comparison-with-other}, the use of heterogeneous data from the literature may produce a different \\lz\\ relation than the one obtained with an homogeneous dataset, a likely consequence of distance uncertainties. Therefore, we anticipate that the relation by Mendes de Oliveira et al. (\\cite{mendes-de-oliveira_etal06}) will need to be revised once a better controlled sample, with $K$-band photometry, becomes available. Finally, the use of a non-standard passband does not allow an easy comparison of the \\lz\\ relation by Lee et al. (\\cite{lee_etal06}) to other \\lz\\ relations. For example they need to make several successive assumptions in order to convert their luminosities into masses, and then compare the mass-metallicity relation to the SDSS one. For all these reasons, we believe that so far our approach stands as the one with the smallest number of uncertainties and the broadest applicability.  The paper is structured as described below. In Sect.~2 we explain how the targets were selected, and give a brief account of the data reduction techniques -- described in more detail in the appendices --. The computation of chemical abundances is described in Sect.~3, where we also compare our results with previous abundance determinations. The resulting NIR luminosity-metallicity relation for dIrr galaxies is presented in Sect.~4. It is discussed in Sect.~5, with a possible explanation of its origin presented in Sect.~5.1 and its comparison with a relation based on literature data in Sect.~5.2. In Sect.~5.3 we discuss the possible dwarf vs. giant galaxy dichotomy, through a comparison of our results to those by Salzer et al. (\\cite{salzer05}). Finally a summary and the conclusions of this study are provided in Sect.~6. ",
        "conclusions": "We have measured oxygen abundances in \\ion{H}{ii} regions located inside a number of dIrr galaxies, belonging to the nearby Sculptor and M81 groups. The abundances were measured with the temperature-sensitive (direct) method, based on the ratio of the auroral line {[}\\ion{O}{iii}]$\\lambda4363$ to the nebular lines {[}\\ion{O}{iii}]$\\lambda\\lambda4959,5007$, and were complemented by direct abundances from the literature for two additional galaxies. The weak {[}\\ion{O}{iii}]$\\lambda4363$ line could be measured in \\ion{H}{ii} regions belonging to five galaxies of the Sculptor group and two galaxies of the M81 group. Metallicity gradients were detected in ESO~245-G005 and DDO~42, confirming earlier findings, so only the central highest-metallicity regions were considered for the derivation of the luminosity-metallicity relation. This forced us to discard our DDO~42 measurement, and adopt a literature value. Our fiducial sample then includes six dIrr galaxies of the Sculptor group and two galaxies of the M81 group.  For these eight galaxies we have obtained deep NIR, $H$-band, imaging which allowed us to perform surface photometry and compute their total luminosities. Thanks to the availability of distances with $\\leq15\\%$ errors, the galaxies could be placed in the $12+\\log{\\rm (O/H)}$ vs. $M_{H}$ diagram, revealing a clear \\lz\\ relation with a small $0.11$~dex scatter around the average trend. The scatter is smaller than that of relations obtained at optical wavelengths (e.g. $0.161$~dex using $B$-band data, Lee et al. \\cite{lee_etal06}), and is comparable to the one obtained by Lee et al. (\\cite{lee_etal06}) with their MIR {[}4.5\\micron]-band data ($0.12$~dex). Assuming the existence of a fundamental mass-metallicity relation, the improved definition of the \\lz\\ relation at NIR luminosities must be due to the fact that it better traces the underlying relation with mass, since the NIR mass-to-light ratio is more sensitive to the \\emph{integrated} star formation history of a galaxy. On the contrary, the blue mass-to-light ratio is more sensitive to the instantaneous SFR, which has a large galaxy-to-galaxy scatter (Saviane et al. \\cite{sidney}; Salzer et al. \\cite{salzer05}).  Our work and that of Lee et al. (\\cite{lee_etal06}) are the ones that managed to obtain the best defined \\lz\\ relation, and as an additional advantage, our standard NIR band allows an easy comparison to other relations obtained in independent studies. Indeed we compared our relation to that of Salzer et al. (\\cite{salzer05}), who are the only authors providing an $H$-band relation, at the same time extending it to giant galaxies. Unfortunately a direct comparison with our accurate abundances is not possible, since the abundances of S05 are obtained with the so-called strong-line method (indirect or empirical method). To be consistent with S05 the abundances of our dwarf sample were also computed with the empirical method, and although the new \\lz\\ relation has more scatter than the one using direct abundances, in this way our dwarf galaxies can be placed in a same graph together with the galaxies analyzed by S05. The slope and zero-point of the relation for our dwarfs, using indirect abundances, are different than those of the relation obtained with direct abundances, and, when compared to S05 giants, a possible break-point appears at $M_{H}\\approx-20$. While this could suggest a different gas-consumption mechanism for dwarf and giant galaxies, a solid conclusion cannot be established at this stage. Indeed, the version of the empirical method employed by S05 is different from that classically used, and so a slight inconsistency between our abundances and those of S05 may still exist. The best way for giving the final word on this possible giant-dwarf dichotomy would be to obtain direct abundances for a good number of the KISS galaxies. And precisely, such a project has been started recently by some of us (see Saviane et al. \\cite{saviane_etal07}).  One of the motivations for the current study was to remedy the lack of homogeneous abundances and luminosities in the literature, but still it was interesting to make the experiment of assembling an \\lz\\ relation based on literature data, and see whether our approach really turned out to be superior. We concluded that indeed mixing data from a variety of sources builds an \\lz\\ relation affected by large uncertainties, and very different from the one obtained with a controlled sample and controlled methods. Inferences based on literature data are thus to be taken as qualitative at best, and should be discarded, if possible.  Limiting ourselves to our sample of dIrrs, we attempted to explain the \\lz\\ relation assuming that the chemical evolution of these objects is similar to that of a closed-box model with an effective yield which is $1/3$ its value in the solar vicinity (Skillman et al. \\cite{skillman_scl_hii}, Pilyugin et al. \\cite{pilyugin_etal_04}). If this approximation holds true, and if we assume that the chemical evolution has started at the same time irrespective of mass, then the conclusion is that more massive galaxies have faster chemical evolutions (gas-consumption rates). Alternatively, more massive galaxies had more time to build up their metal content.  Due to a number of reasons (mainly adverse weather conditions) we could assemble a large database only for the Sculptor group of galaxies, while only two galaxies of the M81 group enter our \\lz\\ relation. Although these two dwarfs seem to follow the same relation as that defined by the Sculptor group galaxies, more M81 dIrrs need to be measured, before one can understand whether the higher density in that group is influencing the chemical evolution of its members.  {Indeed in general terms our \\lz\\ relation is based on a small number of galaxies, therefore more objects need to be added to Fig.~\\ref{fig:lz} in order to confirm our \\lz\\ relation with a larger sample. For example looking at the number of dIrr presented in Lee et al. (\\cite{lee_etal06}), we might be able to add $\\sim20$ galaxies to our database. And 10m-class facilities will have to be used to reach this goal in a reasonable time.} In the short term we plan to keep collecting data for the M81 group, and start a similar study for the Centaurus~A group. A natural extension of the project would be to investigate clusters of galaxies, but this is a difficult task from the observational point of view, since the two nearest galaxy clusters are Virgo and Fornax, at $22$ and $24$~Mpc respectively (Ferguson \\& Sandage \\cite{ferguson_sandage90}; hereafter FS90). A more viable alternative would be to observe the Leo group of galaxies. At a distance of $18.2$~Mpc (FS90), this group is the fourth nearest one, and has a relatively high velocity dispersion of $250\\,{\\rm km\\, sec}^{-1}$. Its properties (density, number of galaxies, dwarf-to-giant galaxy ratio, etc.) are intermediate between those of nearby groups and those of clusters of galaxies (see e.g. Figure~1 and following in Ferguson \\& Sandage \\cite{ferguson_sandage91}). It would then allow us to begin the exploration of a cluster-like environment, with NIR cameras at 10m-class telescopes."
    },
    "0806/0806.2026_arXiv.txt": {
        "abstract": "{Very high dynamical range coronagraphs targeting direct exo-planet detection \\rm ($\\rm{10^{9}\\sim10^{10}}$ \\rm contrast) at small angular separation (few $\\lambda/D$ units) usually require an input wavefront quality on the order of ten thousandths of wavelength RMS.} {We propose a novel method based on a pre-optics setup that behaves partly as a low-efficiency coronagraph, and partly as a high-sensitivity wavefront \\rm aberration \\rm compensator (phase and amplitude). The combination of the two effects results in a highly accurate corrected wavefront.} {First, an (intensity-) unbalanced nulling interferometer (UNI) performs a rejection of part of the wavefront electric field. Then the recombined output wavefront has its input \\rm aberrations \\rm magnified. Because of the unbalanced recombination scheme, \\rm aberrations \\rm can be free of phase singular points (zeros) and can therefore be compensated by a downstream phase and amplitude correction (PAC) adaptive optics system, using two deformable mirrors.} {In the image plane, the central star's peak intensity and the noise level of its speckled halo are reduced by the UNI-PAC combination: the output-corrected wavefront \\rm aberrations \\rm can be interpreted as an improved compensation of the initial (eventually already corrected) incident wavefront \\rm aberrations. \\rm } {The important conclusion is that \\rm not all \\rm the elements in the optical setup using UNI-PAC \\rm need \\rm to reach the $\\lambda/10000$ rms surface error quality.} ",
        "introduction": "Optical coronagraphy in space is one of the most useful methods to achieve high dynamic range observations for the direct detection of extra-solar planets (e.g., \\cite{Coul05}; \\cite{Aime06}). A coronagraph can reduce a central star intensity and its diffracted halo light around the star where exo-planets would appear. Several coronagraph designs using  advanced focal plane masks (\\cite{Guyo99}; \\cite{Kuch02}; \\cite{Baba02}; \\cite{Riau03}) have been reported which reduce the light energy, sometimes to zero in theory, inside a re-imaged pupil plane called a Lyot stop, and also at a final image. Some \\rm techniques \\rm of nulling interferometry, which has been mainly considered for mid-infrared long-baseline interferometers (\\cite{Menn05}), \\rm has been applied \\rm to optical coronagraph in a single telescope using rotation-shearing interferometers (\\cite{Baud00}; \\cite{Tavr05}) or lateral-shearing interferometers with overlapped or separated sub-apertures (\\cite{Shao04}; \\cite{Nish05}). Pupil function modification can reduce only the halo intensity by using shaped or apodized pupils (\\cite{Kasd05}; \\cite{Gali05}). Wavefront phase control by an adaptive optics (AO) system can reduce the halo intensity of a limited area, called ``dark hole'' (\\cite{malbet95}). Multiple-stage coronagraphs (\\cite{Aime04}; \\cite{Toll05}), and combinations of these methods, or pre-optics schemes have been proposed to achieve a higher dynamic range (\\cite{Nish05}; \\cite{NishM06}; \\cite{Abe06}).  In any case, the required dynamic range of ${\\rm 10^{9}\\sim10^{10}}$ \\rm for direct detection of earth-like planets by optical coronagraphs can be achieved only with a very high quality wavefront of $\\lambda$/10000 rms and an intensity uniformity of 1/1000 rms (\\cite{Kuch02}; \\cite{Lowman}). Indeed, wavefront \\rm aberrations \\rm throughout the optical train produce a so-called speckle halo noise in the image plane (e.g., \\cite{Bord06}) that prevents direct detection of the planet if these speckles are too bright. Manufacturing or polishing accurate mirrors would be a good direction for small mirrors, while it would not be easy for a telescope primary to reach the required surface accuracy. Then an AO system would be used to achieve the required wavefront accuracy where both measurements and corrections of the wavefront are important. Giant or young planets brighter than $10^{6}$ contrast are the target of ground-based telescopes in the infrared wavelength where wavefront correction by an AO system is also a key technology \\rm for use against \\rm atmospheric turbulence (e.g., \\cite{Maki06}; \\cite{Sera06}; \\cite{Tamu06}).  In an AO system, a deformable mirror (DM) is controlled by a signal from a wavefront sensor (WFS). Recently, control accuracy of a DM has been shown at a level of $\\lambda$/10000 rms (e.g., \\cite{Brow03}; \\cite{Evan05}), while some commercial wavefront sensors have been reported to achieve absolute wavefront measurement accuracy of $\\lambda$/1000 rms by conventional Shack-Hartmann method (e.g., HASO HP 26 from Imagine Optic). The same goes for interferometric sensing, where repeatability on the order of $\\lambda$/10000\\,rms is commonly achieved (e.g., Zygo VeriFire AT). In the high dynamic range coronagraph regime, however, a non-common path error problem, i.e., differences between optics of the WFS and the main path for the star and the planet images, should be considered in the pupil plane WFS such as the Shack-Hartmann sensor, the interferometer, or a curvature sensor (e.g., \\cite{Guyo06}).  Solutions for the non-common path error problem are focal plane WFSs (\\cite{Codo04}; \\cite{Bord06}; \\cite{Give06}) or speckle nulling, \\rm which features \\rm focal plane speckle measurement and iterative wavefront control. A dynamic range of $10^{9}$ within the dark hole has been achieved after a few \\rm thousand iterations \\rm by an AO system (\\cite{Bala06}) where the wavefront correction \\rm approaches \\rm the requirement.  These focal-plane sensing approaches, however, meet other problems. When \\rm the speckle intensity level gets closer to the planet intensity, long exposure times are required to obtain \\rm better S/N measurements of the residual wavefront \\rm aberrations. In addition, \\rm the wavefront stability of the whole telescope system must be guaranteed during \\rm this time lapse. \\rm A multiple stage coronagraph reducing speckles with AOs at every stage (\\cite{Toll05}) faces the same problem. In the case of multiple stages \\rm with approaches \\rm using nulling coronagraphs or nulling interferometers, no wavefront sensing has been considered before (e.g. \\cite{Nish05}; \\cite{NishM06}). But the same wavefront sensing problems would exist, i.e. the intensity in the pupil plane (at the Lyot stop) can be zero or very low, exhibiting a random \\rm intensity pattern \\rm where phase singularity at zero intensity points makes wavefront measurements and corrections difficult. Thus in the very high dynamic range optics for direct detection of exo-planets, maintaining wavefront quality remains an issue as well as the development of coronagraphs.  We propose a pre-optics scheme for coronagraphs, using a combination of an unbalanced nulling interferometer (UNI) and a two-DM AO system (Fig.\\ref{fig_optics}) for Phase and Amplitude Correction (PAC). Similarities can be found with previous studies, i.e. pre-optics by a balanced nulling interferometer (\\cite{Nish05}; \\cite{NishM06}), coronagraphic pre-optics with non-zero amplitude distribution in the pupil after an interference and complex amplitude correction (\\cite{Abe06}), and a multiple-stage wavefront correction by commercial quality AO systems (\\cite{Toll05}).  The pre-optics concept using a balanced nulling interferometer in front of other coronagraphs was shown to be effective in absorbing part of the required dynamic range (\\cite{Nish05}; \\cite{NishM06}) but could not maintain sufficiently high wavefront \\rm aberrations \\rm after the interference to be corrected again by a secondary AO system. As a consequence, it puts some very drastic constraints on the precision and accuracy of the upstream AO system. Abe et al. (2006) showed that a pre-optics, that uses a pseudo-coronagraphic stage and a complex amplitude filter at the Lyot plane, can produce a flat and almost uniform wavefront, in spite of central obscuration and spider-arm patterns (i.e. large amplitude \\rm aberrations \\rm at the entrance pupil). We extended this concept in order to facilitate the phase and amplitude measurement/correction after a nulling stage, which is usually the bottleneck of high contrast coronagraphs. In the proposed concept, the UNI aims at simultaneously performing the absorption of part of the dynamic range, and the magnification of wavefront \\rm aberrations \\rm without phase singularities for later compensation with an AO system (i.e. the PAC), so that the star's speckled halo is effectively attenuated.  \\rm This paper introduces the principle of the UNI-PAC concept, \\rm  a novel method that \\rm avoids the non-common path error problem and the low intensity or phase singularity problem. \\rm The overview, formalization, and simulation of the concept \\rm are \\rm addressed in Sect.\\,\\ref{sect:overview}, Sect.\\,\\ref{sect:formalism}, and Sect.\\,\\ref{sect:simulation}, respectively, followed by a discussion in Sect.\\,\\ref{sect:discussion}. Conclusions \\rm are \\rm drawn in Sect.\\,\\ref{sect:conclusion}.  \\begin{figure} \\centering\\includegraphics[width=\\columnwidth]{f1.eps} \\caption{Schematic of a possible optical layout for the UNI-PAC method (unbalanced nulling interferometer followed by a phase and amplitude correction adaptive optics). The various optical planes are those referred to in the text.} \\label{fig_optics} \\end{figure} ",
        "conclusions": "\\label{sect:conclusion} We have presented the principle of a pre-optics concept for precise wavefront \\rm aberration \\rm reduction in front of a coronagraph in a terrestrial-planet finding telescope. It consists of a combination of an unbalanced nulling interferometer and a two-deformable mirror AO system with a conventional pupil plane wavefront sensor, \\rm where it does not meet the low intensity problem for wavefront measurements. \\rm This method reduces both the source intensity as a nulling coronagraph and the (speckle halo) noise intensity produced by wavefront \\rm aberrations. \\rm  \\rm The wavefront magnification phenomenon at the unbalanced nulling interference makes it possible to correct the wavefront precisely beyond the capabilities of employed AO systems with non-common path errors. In reality, by using the UNI-PAC, the specification for all of the coronagraph optics can be relaxed which is probably key to developing a cost-effective exo-planet detection system, although the performance of the whole system should be investigated for each case. \\rm  \\rm A candidate instrument where the UNI-PAC system can be used would be a modified Michelson interferometer in a collimated beam of an off-axis telescope (\\cite{Shao04}). \\rm We also expect the \\rm present \\rm concept to be applicable with several other coronagraphic concepts instead of the UNI stage \\rm by a Michelson beam combiner, \\rm i.e., with more general contrast-reduction setups (see for example \\cite{Abe06}), but we leave this discussion for a future study.  Some other considerations, such as achromaticity, resolved star nulling, planet image quality, and AO performance limitations should be investigated."
    },
    "0806/0806.1739_arXiv.txt": {
        "abstract": "We develop a general physical model for how galactic disks survive and/or are destroyed in mergers and interactions. Based on simple dynamical arguments, we show that gas primarily loses angular momentum to internal torques in a merger, induced by the gravity of the secondary. Gas within some characteristic radius, determined by the efficiency of this angular momentum loss (itself a function of the orbital parameters, mass ratio, and gas fraction of the merging galaxies), will quickly lose angular momentum to the stars sharing the perturbed host disk, fall to the center and be consumed in a starburst. We use a similar analysis to determine where violent relaxation of the pre-merger stellar disks is efficient on final coalescence. Our model describes both the dissipational and dissipationless components of the merger, and allows us to predict, for a given arbitrary encounter, the stellar and gas content of the material that will survive (without significant angular momentum loss or violent relaxation) to re-form a disk in the merger remnant, versus being dissipationlessly violently relaxed or dissipationally losing angular momentum and forming a compact central starburst. We test these predictions with a large library of hydrodynamic merger simulations, and show that they agree well (with small scatter) with the properties of simulated merger remnants as a function of merger mass ratio, orbital parameters, and gas distributions, in simulations which span a wide range of parameter space in these properties as well as prescriptions for gas physics, stellar and AGN feedback, halo and initial disk structural properties, redshift, and galaxy masses. We show that, in an immediate (short-term) sense, the amount of stellar or gaseous disk that survives or re-forms following a given interaction can be understood purely in terms of simple, well-understood gravitational physics, independent of the details of the ISM gas physics or stellar and AGN feedback. This allows us to demonstrate and quantify how these physics are in fact important, in an indirect sense, to enable disks to survive mergers, by lowering star formation efficiencies in low mass systems (allowing them to retain large gas fractions) and distributing the gas to large radii. The efficiency of disk destruction in mergers is a strong function of gas content -- our model allows us to explicitly predict and demonstrate how, in sufficiently gas rich mergers (with quite general orbital parameters), even 1:1 mass-ratio mergers can yield disk-dominated remnants, and more realistic 1:3-1:4 mass-ratio major mergers can yield systems with $<20\\%$ of their mass in bulges. We discuss a number of implications of this modeling for the abundance and morphology of bulges as a function of mass and redshift, and provide simple prescriptions for the implementation of our results in analytic or semi-analytic models of galaxy formation. ",
        "introduction": "\\label{sec:intro}  In the now established ``concordance'' $\\Lambda$CDM cosmology, structure grows hierarchically \\citep[e.g.][]{whiterees78}, making mergers and interactions between galaxies an essential and inescapable process in galaxy formation. Indeed, mergers are widely believed to be responsible for the morphologies of spheroids \\citep[bulges in disks and elliptical galaxies;][]{toomre77}, and observations find recent merger remnants in considerable abundance in the local universe \\citep{schweizer82,LakeDressler86,Doyon94,ShierFischer98,James99, Genzel01,tacconi:ulirgs.sb.profiles,dasyra:mass.ratio.conditions,dasyra:pg.qso.dynamics, rj:profiles,rothberg.joseph:kinematics} as well as e.g.\\ faint shells and tidal features common around apparently ``normal'' galaxies \\citep{malin80,malin83,schweizer80, schweizerseitzer92,schweizer96}, which are thought to be signatures of galaxy collisions \\citep[e.g.][]{hernquistquinn88,hernquist.spergel.92}.  From both theoretical grounds \\citep[][and references therein]{ostrikertremaine75,maller:sph.merger.rates, fakhouri:halo.merger.rates,stewart:mw.minor.accretion} and observations \\citep[e.g.][]{lin:merger.fraction,barton:triggered.sf, woods:tidal.triggering, woods:minor.mergers} it appears that ``minor'' mergers of mass ratios $\\lesssim 1:10$ are ubiquitous (there are almost no galaxies without mergers of at least this mass ratio in the last few Gyr), and moreover a large fraction ($\\sim1/2$) of the $\\sim L_{\\ast}$ galaxy population is observed and expected to have experienced a ``major'' merger (mass ratio $\\lesssim1:3$) since $z\\sim2-3$ \\citep{lotz:merger.fraction,bell:merger.fraction, bridge:merger.fractions,lin:mergers.by.type,kartaltepe:pair.fractions}. With increasing redshift, kinematic and morphological indications of recent, violent disturbance in disk-dominated galaxies appear more frequent \\citep{hammer:obs.disks.w.mergers, flores:tf.evolution,puech:highz.vsigma.disks,puech:tf.evol}.   Far from there not being enough mergers to explain the abundance of bulges and ellipticals, this has led to the concern that there may be far too {\\em many} mergers to explain the survival and abundance of galactic disks in the context of our present understanding of galaxy formation. \\citet{toomre72} were among the first to point out that mergers are capable of dramatically altering the morphologies of disks, transforming them into elliptical galaxies. Although their neglect of the importance of dissipational star formation and gas dynamics in the mergers led to some controversy \\citep[e.g.][]{ostriker80,carlberg:phase.space,gunn87,kormendy:dissipation}, it is now increasingly well-established that major mergers between spiral galaxies (similar to those observed locally and at $z\\lesssim2-3$) with gas fractions comparable to those observed yields remnants in good agreement with essentially all observed properties of low and intermediate-mass local elliptical galaxies \\citep[e.g.\\ morphologies, shapes, sizes, kinematics, densities, colors, black hole properties, fundamental scaling relations, stellar populations, and halo gas;][]{hernquist.89, barnes.hernquist.91,barneshernquist96, hernquist:phasespace,mihos:gradients,mihos:starbursts.96, dimatteo:msigma,naab:gas,jesseit:kinematics, cox:xray.gas,cox:kinematics,robertson:fp,springel:red.galaxies,burkert:anisotropy, hopkins:clustering, hopkins:cusps.ell,hopkins:cores,hopkins:cusps.fp,hopkins:groups.ell,hopkins:cusps.mergers}.  Many intermediate and low-luminosity ``cusp'' ellipticals (encompassing $\\sim80-90\\%$ of the mass density in ellipticals) contain significant embedded disks \\citep[perhaps all such ellipticals, given projection effects; see][]{ferrarese:type12,lauer:centers}, and they form a continuous sequence with most S0 galaxies, known to have prominent stellar (and even gaseous) disks \\citep{kormendy:spheroidal1, bender:ell.kinematics,ferrarese:type12,kormendy94:review,lauer:95, faber:ell.centers,kormendy99,ferrarese:profiles,emsellem:sauron.rotation}. Indeed, the existence of embedded disks in simulated merger remnants is critical to matching the properties described above.  A wide variety of observations including stellar populations and star formation histories \\citep[e.g.][]{bender89,trager:ages, mcdermid:sauron.profiles} and kinematic and structural analysis of recent merger remnants \\citep{schweizer83,schweizer83:review,schweizerseitzer92, schweizer:ngc34.disk,hibbard.yun:excess.light,rj:profiles} demonstrate that most of these disks are not accreted in the standard cosmological fashion after the spheroid forms -- they must somehow survive the merger or form very quickly thereafter from gas already in and around the galaxies. Therefore, despite the destruction of a large portion of a stellar disk in major mergers, {\\em some} disk must survive mergers, and the amount that does so is a critical component determining many of the photometric and kinematic properties of even bulge-dominated and elliptical galaxies.  Moreover, ``minor'' mergers -- at least those with mass ratios $\\lesssim10:1$ (below which the difference between ``merger'' and accretion becomes increasingly blurred) -- are not generally believed to entirely destroy disks, but they are almost an order of magnitude more frequent than major mergers and as such may pose a more severe a problem for disk survival. In the $\\Lambda$CDM cosmology, and from observed satellite fractions, it is unlikely than any disk (let alone a large fraction of disk galaxies) with a significant stellar age has survived $\\sim5-10$\\,Gyr without experiencing a merger of mass ratio $10:1$ or larger. Simulations \\citep{quinn.84,quinn86:dynfric.on.sats,quinn93.minor.mergers, hernquist.mihos:minor.mergers, walker:disk.fragility.minor.merger,velazquezwhite:disk.heating, naab:minor.mergers, bournaud:minor.mergers,younger:antitruncated.disks,younger:minor.mergers} and analytic arguments \\citep{ostrikertremaine75,tothostriker:disk.heating, sellwood:resonant.disk.thickening} suggest that gas-poor minor mergers can convert a considerable fraction of a stellar disk into bulge and cause significant perturbation (``puffing up'' via dynamical heating) to the disk. The observed coldness of galactic disks suggests that this may be a severe problem: \\citet{tothostriker:disk.heating} argued that large disks such as that in the Milky Way could not have undergone a merger of mass ratio $\\lesssim10:1$ in the last $\\sim10\\,$Gyr. More recently e.g.\\ \\citet{stewart:mw.minor.accretion} and \\citet{hammer:mw.no.mergers} emphasized that the tension between these constraints and the expectation in CDM models that a number of such mergers should occur implies either a deficit in our understanding of hierarchical disk formation or a challenge to the concordance cosmological model.  Given the successes of the $\\Lambda$CDM model on large scales, and the increasing observational confirmation that disks do undergo (and therefore must somehow survive) a large number of mergers, it is likely that the problem lies in our (still relatively poor) understanding of disk galaxy formation. This has led to a great deal of focus on the problem of forming realistic disks in a cosmological context, with many different attempts and debate on the missing elements necessary to produce disks in simulations. Various groups have argued that self-consistent treatment of gas physics and star formation along with implementation of feedback of different kinds is necessary, along with greatly improved numerical resolution \\citep{weil98:cooling.suppression.key.to.disks, sommerlarsen99:disk.sne.fb,sommerlarsen03:disk.sne.fb, thackercouchman00,thackercouchman01,abadi03:disk.structure, governato04:resolution.fx,governato:disk.formation, robertson:cosmological.disk.formation, okamoto:feedback.vs.disk.morphology,scannapieco:fb.disk.sims}, in order to enable disks to survive their expected violent merger histories without completely losing angular momentum and transforming into systems that are too compact and have too much bulge mass (relative to real observed disks) by $z=0$.  It has been known for some time \\citep[see e.g.][]{hernquist:kinematic.subsystems, barneshernquist96} that (even without any feedback) some fraction of the gas in even a major merger of two disks can survive and form new, embedded disks in the remnant -- i.e.\\ despite the problems outlined above, disks are not necessarily completely destroyed in mergers. However, early studies of this were restricted to cases with low gas content ($f_{\\rm gas}\\lesssim10\\%$ in the progenitor disks), most of which was rapidly consumed in star formation, yielding small remnant disks in strongly bulge-dominated remnants. In seminal work, \\citet{springel:spiral.in.merger} and \\citet{robertson:disk.formation} showed that, in idealized merger simulations with significant stellar feedback to allow the stable evolution of extremely gas rich disks ($f_{\\rm gas}\\sim1$), even a major merger can produce a disk-dominated remnant. This has since been confirmed in fully cosmological simulations \\citep{governato:disk.formation}. Together with other recent investigations (see references above), these works have led to the growing consensus that a combination of strong stellar feedback and large gas content is essential to the survival of disk galaxies.  A large number of open questions remain, however. How, exactly, does feedback allow disks to survive mergers? What are the most important physics? Does it require fine-tuning of feedback prescriptions? How might things vary as a function of galaxy mass, redshift, gas content, merger orbits, and environment? Fundamentally, should this be expected for typical cosmological circumstances, or are these cases pathological?  The ambiguity largely owes to the fact that there is no deep physical understanding of how disks survive or re-form after mergers and interactions. It has only just become possible to conduct simulations with the requisite large gas fractions, and thus far theoretical explanations have largely been restricted to phenomenological analysis, with continued efforts to improve resolution and sub-resolution prescriptions. Moreover, without a full model for how disks behave in interactions, these simulations cannot be placed into the broader context of the emergence of the entire Hubble sequence (for example asking the question, are the disks in lenticulars and embedded disks in ellipticals survivors of their pre-merger disks? Are they re-accreted? What determines how large they are? What is the key physics that gives rise to realistic embedded disks, leading to bulge-dominated galaxies with kinematic and photometric properties similar to those in the real universe?) or within a fully cosmological context.  The resolution requirements for full models of disk formation are severe -- limiting any attempt to properly simulate a cosmological box and still achieve the resolution necessary to reliably model a disk population -- and so models of the population of disks, largely semi-analytic, are forced to adopt simplified and un-tested prescriptions for the behavior of disks in mergers. This, in turn, has led to other well-known problems in modeling disk populations (even where prescriptions can ensure no artificial angular momentum losses); even when the cumulative (morphology-independent) galaxy mass function is correctly predicted at the low-mass end, semi-analytic models widely overproduce the relative abundance of low-mass spheroids and underproduce disks \\citep[even when satellites, which have other associated model uncertainties, are removed from consideration; see e.g.][]{somerville:sam, somerville:new.sam,croton:sam, bower:sam,delucia:sam}. Lacking a proper, physically motivated understanding of how low-mass or gas-rich disks may or may not survive mergers, attempts to address this problem in the models have been purely phenomenological and involve arbitrary prescriptions \\citep[see e.g.][]{koda:disk.survival.prescriptions}.  Motivated by these concerns, in this paper we develop a physical, dynamical model for how disks survive and are destroyed in mergers and interactions. We show that, in an immediate (short-term) sense, the amount of stellar or gaseous disk that survives or re-forms following a given interaction can be understood purely in terms of simple, well-understood gravitational physics. Knowing these physics, we develop an analytic model that allows us to accurately predict how much of a given pre-merger stellar and cold gas disk will survive a merger, as a function of the merger mass ratio, orbital parameters, pre-merger cold gas fraction, and mass distribution of the gas and stars. We compare these predictions to the results of a large library of hundreds of hydrodynamic simulations of galaxy mergers and interactions, spanning a wide parameter space in these properties as well as prescriptions for gas physics, stellar and AGN feedback, halo and initial disk structural properties, redshift, and absolute galaxy masses. Our numerical experiments confirm that the analytic scalings accurately describe the behavior and bulge formation/disk destruction in mergers over the entire dynamic range surveyed, and confirm that the parameters not explicitly included in our model do not systematically affect either the mean predictions or the scatter of simulations about those predictions. This allows us to understand the mean behavior of systems with different orbits and mass ratios, as well as why systems with large gas fractions can form little bulge in even major mergers.  This is possible because gas, in mergers, primarily loses angular momentum to internal gravitational torques (from the stars in the same disk) owing to asymmetries in the galaxy induced by the merger. Hydrodynamic torques and the direct torquing of the secondary are second-order effects, and very inefficient. Once gas is drained of angular momentum, there is little alternative but for it to fall to the center of the galaxy and form stars, regardless of the details of the prescriptions for star formation and feedback (these may change things at the $\\sim10-20\\%$ level by blowing out some of the gas, but they cannot fundamentally alter the fact that cold gas with no angular momentum will be largely unable to form any sort of disk, or the fact that a galaxy's worth of gas compressed to high densities and small radii will inevitably form a large mass in stars). But if the systems are sufficiently gas-rich, then there is little stellar material sharing the disk to torque on the gas in the interaction, and little or no angular momentum is lost.  Feedback can dramatically alter the ability of a disk to survive in a cosmological sense: by allowing galaxies to retain large gas fractions (as opposed to no-feedback scenarios, in which cold gas in a disk is usually quickly converted into stars), they are more gas-rich when they undergo interactions, allowing them to avoid angular momentum loss for the reason above. Moreover, we show that in detail (owing to the resonant structure of interactions), it is really gas within a certain radius of the stellar disk that is drained of angular momentum. The commonly-invoked stellar wind feedback then enables cosmological disk survival in a second fashion: by redistributing gas out to large radii, it prevents angular momentum loss and allows rapid re-formation of disks after a merger. Independent of any tuning, our model allows us to quantify the disks expected as a function of interactions of arbitrary properties, and to physically, explicitly quantify what the requirements are for feedback, in a cosmological scenario, to enable disk survival.  In \\S~\\ref{sec:sims} we describe our library of gas-rich merger simulations, which we use to test our physical model for disk destruction and survival. In \\S~\\ref{sec:id} we demonstrate the existence of genuine disks in remnants of even major mergers and briefly consider their properties, and compare methods to separate the disks and bulges in merger remnants. In \\S~\\ref{sec:form.major} we consider the question of how these disks form in and survive mergers: we identify the key components of any merger remnant in \\S~\\ref{sec:form.major:components}, highlighting that these disks originate from a combination of undestroyed pre-merger stellar disks and gas which avoids angular momentum loss in the merger. In \\S~\\ref{sec:form.major:angloss} we discuss how, in detail, that angular momentum loss proceeds. We use this, in \\S~\\ref{sec:model.overview}, to build a physical model for how angular momentum loss proceeds in mergers and predict the surviving disk content of merger remnants: we model and test how this depends on the gas content of the pre-merger disks (\\S~\\ref{sec:model.gas}), and the orbital parameters (\\S~\\ref{sec:model.orbit}) and mass ratio (\\S~\\ref{sec:model.massratio}) of the encounter. We generalize to first passage and fly-by encounters (\\S~\\ref{sec:model.flyby}) and demonstrate that (for otherwise fixed conditions at the time of an encounter) our conclusions are purely dynamical, independent of feedback physics or details in our treatment of e.g.\\ star formation and the ISM gas physics (\\S~\\ref{sec:model.feedback}), although we use our model to determine exactly how these choices can have dramatic {\\em indirect} consequences for disk survival (by altering the state of systems leading into a merger). We discuss some exceptions and pathological cases in \\S~\\ref{sec:model.exceptions}, and relate our results to the long-term secular evolution of barred systems in \\S~\\ref{sec:model.secular}. In \\S~\\ref{sec:prescriptions}, we outline how these results can and should be applied in analytic and semi-analytic models of galaxy formation, and give appropriate prescriptions derived from our numerical experiments. Finally, we summarize our results and discuss some of their cosmological implications and applications to other models and observations in \\S~\\ref{sec:discussion}.  Throughout, we assume a $\\Omega_{\\rm M}=0.3$, $\\Omega_{\\Lambda}=0.7$, $H_{0}=70\\,{\\rm km\\,s^{-1}\\,Mpc^{-1}}$ cosmology, but this has little effect on our conclusions.  \\breaker ",
        "conclusions": "\\label{sec:discussion}  We have derived a general physical model for how disks survive and/or are destroyed in mergers and interactions. Our model describes both the dissipational and dissipationless components of the merger, and allows us to predict, for a given arbitrary encounter, the stellar and gas content of the system that will be dissipationlessly violently relaxed, dissipationally lose angular momentum and form a compact central starburst, or survive (without significant angular momentum loss or violent relaxation) to re-form a disk. We show that, in an immediate (short-term) sense, the amount of stellar or gaseous disk that survives or re-forms following a given interaction can be understood purely in terms of simple, well-understood gravitational physics. Knowing these physics, our model allows us to accurately predict the behavior in full hydrodynamic numerical simulations across as a function of the merger mass ratio, orbital parameters, pre-merger cold gas fraction, and mass distribution of the gas and stars, in simulations which span a wide range of parameter space in these properties as well as prescriptions for gas physics, stellar and AGN feedback, halo and initial disk structural properties, redshift, and absolute galaxy masses.  The fact that we can understand the complex, nonlinear behavior in mergers with this analytic model, and moreover that (for given conditions at the time of merger) our results are independent of the details of prescriptions for gas physics, star formation, and feedback, owes to the fact that the processes that strip angular momentum from gas disks and violently relax stellar disks are fundamentally {\\em dynamical}.  Gas, in mergers, primarily loses angular momentum to internal gravitational torques (from the stars in the same disk) owing to asymmetries in the galaxy induced by the merger (on the close passages and final coalescence of the secondary, during which phase the potential also rapidly changes, scattering and violently relaxing the central stellar populations of the stellar disk).\\footnote{We note again that although we have described these asymmetries as ``bars'' or ``bar-like'' at certain points in this paper, there are a number of properties of the non-axisymmetric distortions induced in mergers (discussed in \\S~\\ref{sec:model.orbit} and \\S~\\ref{sec:model.secular}) that make them -- at least over the short relaxation timescale of the merger -- dynamically distinct from traditional bar instabilities in isolated systems.} Hydrodynamic torques and the direct torquing of the secondary are second-order effects, and inefficient for all but pathological orbits.  Once gas is efficiently drained of angular momentum, there is little alternative but for it to fall to the center of the galaxy and form stars, regardless of the details of the prescriptions for star formation and feedback -- we show that even strong supernova-driven winds (with mass loading efficiencies several times the star formation rate and wind mass-loading velocities well above the halo escape velocity) do not significantly effect our conclusions. Such processes, after all, can blow out some of the gas, but they cannot fundamentally alter the fact that cold gas with no angular momentum will be largely unable to form any sort of disk, or the fact that a galaxy's worth of gas compressed to high densities and small radii will inevitably form a large mass in stars.  For these reasons, many processes and details that are important cosmologically (systematically changing e.g.\\ the pre-merger disk gas fractions) -- in some sense setting the initial conditions for our idealized study of what happens in mergers -- do not alter the basic dynamical behavior within the mergers themselves, and therefore do not change our conclusions.  \\begin{figure*} \\centering \\scaleup \\plotone{f16.ps} \\caption{Summary of our comparison between simulations and analytic model for the mass of disks in merger remnants as a function of appropriate orbital parameters, merger mass ratio, and pre-merger cold gas content. We plot our model prediction versus the simulation remnant disk fraction for all $\\sim400$ full hydrodynamic merger simulations considered in this paper (shown in both a linear and logarithmic scale). Symbols encode some of the parameter studies we consider: orbital parameters, galaxy masses, initial merger redshift, choice of feedback prescription, merger mass ratio, and presence or absence of black holes, as labeled. For each subset of simulations, we sample a wide range in initial and pre-merger gas fractions $\\fgas=0-1$. Solid line is a one-to-one relation. In all cases, our predictions agree well with the simulations, with no systematic offsets owing to any of the parameters we have varied. At high $f_{\\rm disk}$, our predictions are accurate to an absolute uncertainty $\\sim0.05-0.10$ in $f_{\\rm disk}$. At low $f_{\\rm disk}\\lesssim0.1$, our predictions are accurate to a factor $\\sim2-3$ (down to $f_{\\rm disk}\\lesssim1\\%$, where it is difficult to reliably identify disks in the remnant). \\label{fig:summary}} \\end{figure*}  Figure~\\ref{fig:summary} summarizes our results for the ensemble of our simulations. We compare the fraction of the baryonic galaxy mass in the merger remnant that is in a surviving post-merger disk to that predicted by our simple model scalings, and find good agreement over the entire range in disk and bulge mass fractions sampled, with surprisingly small scatter given the complexity of behavior in mergers. We highlight several of the parameter studies, showing that -- for fixed mass ratio, orbital parameters, and gas content {\\em at the time of the final merger}, none of these choices systematically affect our predictions (note that these are not the only parameters varied -- the complete list is discussed in \\S~\\ref{sec:sims}, but it is representative). That is not to say they cannot affect them indirectly, by e.g.\\ altering how much gas is available at the time of merger -- but it emphasizes that the processes we model and use to form our predictions, the processes that dominate violent relaxation and the loss of angular momentum in gas in mergers, are fundamentally dynamical.  This allows us to make robust, accurate physical predictions independent of the (considerable) uncertainty in feedback physics and sub-resolution physics of the ISM. Regardless of how those physics alter the ``initial'' conditions, they do not change basic dynamical processes, and so do not introduce significant uncertainties in our model.  In turn, this means that we can use our model to understand just why and how feedback is important for the cosmological survival of disks. Why, in short, have various works \\citep[see e.g.][]{springel:spiral.in.merger,robertson:disk.formation, governato:disk.formation} concluded that strong feedback is essential for enabling disk survival in mergers? Our results show that it is not that feedback somehow makes the disk more robust to the dynamical torques within the merger, in any instantaneous sense. These torques, at least within the critical radii where the gravitational perturbation from the merger is large and in resonance, are sufficiently strong that any reasonable feedback prescription is a dynamically negligible restoring force. Rather, feedback has two important effects that fundamentally alter the conditions in the merger: first, it allows the galaxy to retain much higher gas content going into the merger. Without feedback from e.g.\\ star formation and supernovae contributing to heating and pressurizing the ISM and redistributing gas spatially, isolated gas-rich disks may be unstable to fragmentation. Even if fragmentation is avoided, it is well-known that star formation in simulations proceeds efficiently under these conditions. This would leave the disks essentially pure stars \\citep[even for idealized simulations beginning with $\\sim100\\%$ gas disks; see e.g.][]{springel:models} by the time of the merger, which guarantees that a major merger will inevitably violently relax the stars (this is a simple collisionless mixing process, and under such circumstances is inescapable). With large gas fractions, however, the system relies on stripping angular momentum from the gas to form new bulge stars, which in turn relies on internal torques from induced asymmetries in the stellar disk. If the gas fractions are sufficiently large, there is little stellar disk to do any such torquing, and the gas survives largely intact.  Second, feedback from supernovae and stellar winds moves the gas to large radii, where it does not feel significant torques from the merger. Again, recall that the most efficient torquing is driven by the internal stellar disk of the galaxy, and as such is most efficient at torquing gas within small radii (this can be thought of as analogous to the well-known co-rotation condition for isolated disk bars). If star formation-driven feedback has blown much of the gas to large radii, then there is little gas inside the radius where torques can efficiently strip angular momentum, yielding little induced starburst and largely preserving the gas disk at large radii.  Not only can we qualitatively identify these requirements for feedback processes, but we can more precisely use our model to set quantitative limits on how much gas must be retained and/or the radii it must be redistributed to in order to enable disk survival under various conditions. This also clearly implies that disks must be able to avoid fragmentation and strong local gravitational instabilities when they achieve these gas fractions. This provides a valuable constraint for feedback models -- how those models affect star formation efficiencies, the ``blowout'' of gas, and the local hydrodynamic state (effective equations of state and phase structure) of ISM gas -- and should be useful for calibrating their (still largely phenomenological) implementations in both numerical and semi-analytic models of galaxy formation.  Our predictions are also of interest in any cosmological model for the emergence of the Hubble sequence, since they apply not just to disk-dominated galaxies but to small disks in bulge-dominated systems. We give a number of simple prescriptions for application of our conclusions to analytic and semi-analytic models of galaxy formation, which can be used to predict the distribution of bulge to disk ratios in cosmological ensembles. But even without reference to a full such model, a number of interesting consequences are immediately apparent.  First, it is a well-known problem that theoretical models systematically overpredict the abundance and mass fractions of bulges in (especially) low-mass galaxies. This is true even in e.g.\\ semi-analytic models, which are not bound by resolution requirements and can adopt a variety of prescriptions for behavior in mergers. However, it is also well-established observationally that disk gas fractions tend to be very high in this regime, with large populations of gas-dominated disks at $M_{\\ast}\\ll 10^{10}\\,M_{\\sun}$ \\citep{belldejong:tf,kannappan:gfs,mcgaugh:tf}. Our models predict that bulge formation should, therefore, be strongly suppressed in precisely the regime required by observations. For e.g.\\ disks with $M_{\\ast}<10^{9}\\,M_{\\sun}$ where observations suggest typical gas fractions $\\sim60-80\\%$, our results show that even a 1:1 major merger would typically yield a remnant with only $\\sim30\\%$ bulge by mass -- let alone a more typical 1:3-1:4 mass-ratio merger, which should yield a remnant with $<20\\%$ bulge. That is not to say that it is impossible to form a bulge-dominated system at these masses, but it should be much more difficult than at high masses, requiring either unusually gas-poor systems, violent merger histories, or rarer merging orbits that are more efficient at destroying disks. Our conclusions therefore have dramatic implications for the abundance of bulges and typical morphologies and bulge-to-disk ratios at low galaxy masses and in gas-rich systems. Low-mass systems, when a proper dynamical model of bulge formation in mergers is considered, should have lower bulge-to-disk ratios -- by factors of several, at least -- than have been assumed and modeled in previous theoretical models. Whether this alone is sufficient to resolve the discrepancies with the observations remains to be seen, but it is clearly of fundamental importance that future generations of models incorporate this scaling.  Second, the importance of this suppression owing to gas content in disks will be even more significant at high redshifts. Observations suggest \\citep[see e.g.][]{erb:lbg.gasmasses} that by $z\\sim2$, even systems with masses near $\\sim L_{\\ast}$ ($M_{\\ast}\\sim 10^{10}-10^{11}\\,M_{\\sun}$) may have gas fractions as high as $\\fgas\\sim0.6$. In this regime, the same argument as above should apply, dramatically suppressing the ability of mergers to destroy disks. Moreover, since most of the mass density is near $L_{\\ast}$, this can change not just the behavior in a specific mass regime but significantly suppress the global mass density of spheroids, modifying the predicted redshift history of bulge formation. (Note that this will not change when {\\em stars} form by very much, so it has little or no effect on e.g.\\ the ages of $z=0$ spheroids).  This redshift evolution may also explain the solution to a fundamental problem in reconciling observed disk populations with CDM cosmologies. Integrated far enough back in time, every galaxy is expected to have experienced a significant amount of major merging. In extreme cases, the mass of the system when it had its last such merger may be so small that it would not be noticed today, but in general, it does not require going far back in redshift (to perhaps $z\\sim2-4$ before almost every $z=0$ galaxy should have had such a merger). How, then, can the abundance of systems with relatively small (or even no) visible bulges be explained? Our conclusions here highlight at least part of the answer: as you go back in time, the gas fractions of systems are also higher, nearing unity. So even though, integrating sufficiently far in time, every system has experienced major mergers, it is also true that the systems were increasingly gas-rich, and therefore that the impact of those mergers was more and more suppressed. Only mergers at later times, below certain gas fraction thresholds, will typically destroy disks.  Third, to the extent that bulge formation is suppressed at increasing redshifts, the existence of an $M_{\\rm BH}-M_{\\rm bulge}$ relation \\citep[e.g.][]{magorrian} implies that black hole growth should also be suppressed. Indeed, bulge formation is suppressed specifically because gas cannot efficiently lose angular momentum in mergers if the systems are gas-dominated -- if the gas cannot lose angular momentum efficiently, then it certainly cannot efficiently be accreted by the nuclear black hole. Since this pertains to gas on the scales of galactic disks, it is probably not relevant for the formation of ``seed'' black holes at very high redshift, but it will in general inhibit the growth of black holes owing to early merging activity. At the same time, of course, higher gas fractions in general imply increasing fuel supplies for black hole growth, so the effects are not entirely clear, and more detailed models are needed to see how this impacts the history of black hole growth and quasar luminosity functions. Nevertheless, this may in part explain why, above $z\\sim2$ (where, for the argument above, these effects become important for the global mass density of spheroids), the global rate of black hole growth (i.e.\\ total quasar luminosity density) appears to decline much more rapidly with increasing redshift than the star formation rate density \\citep[compare e.g.][]{hopkinsbeacom:sfh, hopkins:groups.qso,hopkins:bol.qlf}.  Fourth, our models imply that a large fraction of bulges and disks survive mergers together, rather than being formed entirely separately. It is often assumed that classical bulges -- being similar to small ellipticals in most of their properties -- were formed initially in major mergers, as entirely bulge-dominated systems, and then accreted new gaseous and stellar disks at later times. Although nothing in our modeling would prevent this from happening, our analytic and simulation results generically lead to the expectation that a large (perhaps even dominant) fraction of the bulge population did {\\em not} form in this manner, but rather formed {\\em in situ} from minor mergers or less efficient major mergers (in e.g.\\ very gas-rich systems). Observations tracing the evolution of disk components, kinematics, and morphology in the last $\\sim10\\,$Gyr increasingly suggest that such co-formation or disk regeneration scenario is common \\citep[see e.g.][and references therein]{hammer:obs.disks.w.mergers, conselice:tf.evolution,flores:tf.evolution,puech:tf.evol}. In short, a system with a mass fraction $\\sim0.1-0.2$ in a bulge could be the remnant of an early, violent major merger (when the system was $\\sim0.1$ times its present mass) with a re-accreted disk, or could be the remnant of a typical (low to intermediate gas fraction) 1:10-1:5 mass ratio minor merger, or could even be the remnant of a gas-rich major merger (mass ratio $\\lesssim1:3$, if $f_{\\rm gas}$ is sufficiently large).  Based on a simple comparison of typical merger histories, we would actually expect that the minor merger mechanism should be most common, but all may be non-negligible. Fundamentally, the physics forming the bulge (torquing the gas within some radius owing to internal asymmetries and violently relaxing stars within a corresponding radius) are the same in all three cases, and moreover other indicators such as their stellar populations will be quite similar \\citep[in all cases, the bulge will appear old: this is both because the central stars in even present-day disks are much older than those at more typical radii, and because in any case star formation will cease within the bulge itself, as opposed to the ongoing star formation in the disk, and stellar population age estimates are primarily sensitive to the amount of recent or ongoing star formation; see e.g.][]{trager:ages}. This is also not to say that mergers are the only means of producing bulges. Secular evolution of e.g.\\ barred disks probably represents an increasingly important channel for bulge evolution in later-type and more gas-rich systems \\citep[see e.g.][]{christodoulou:bar.crit.1,sheth:bar.frac.evol, mayer:lsb.disk.bars,debattista:pseudobulges.a,jogee:bar.frac.evol, kormendy.kennicutt:pseudobulge.review,marinova:bar.frac.vs.freq}, and may even be related (albeit through longer timescales of ``isolated,'' post-merger evolution and different physics) to initial bar formation or ``triggering'' in mergers. More detailed theoretical work and analysis of cosmological simulations is needed to develop observational probes that can distinguish between these histories.  Further work is specifically needed to investigate the processes at work in minor mergers with mass ratios $\\sim$1:10 ($\\mu\\sim0.05-0.1$), which cosmological simulations suggest are an important contributor to the growth of disks, especially in later-type systems \\citep{maller:sph.merger.rates, fakhouri:halo.merger.rates,stewart:mw.minor.accretion}. In more minor mergers $\\mu \\ll 0.1$, the secondaries are sufficiently small and dynamical friction times sufficiently long that the disk is unlikely to feel significant external perturbations. More major mergers $\\mu \\gtrsim 0.1$, the cases of interest here, induce sufficiently large responses in the disk and evolve sufficiently rapidly that they can be considered ``merger-dominated'' for the reasons in \\S~\\ref{sec:model.orbit} \\&\\ \\ref{sec:model.secular}. But in the intermediate regime, internal amplification of instabilities in a traditional secular fashion may occur on a timescale comparable to or shorter than the evolution of the secondary orbit, potentially leading to a more complex interplay between the two. It is not entirely clear whether such a system would remain ``locked'' to the driven perturbation, or function as a purely secular system (merely initially driven by the presence of the secondary), or some nonlinear combination of both. A more detailed comparison of the relevant timescales for these processes and their relation to e.g.\\ cosmological triggering of bars and large-scale non-axisymmetric modes in disks will be the subject of future study (in preparation).   Our results are also of direct interest to models of spheroid formation in ellipticals and S0 galaxies. As discussed in \\S~\\ref{sec:intro}, it is increasingly clear that embedded sub-components -- constituting surviving gaseous and stellar disks -- are both ubiquitously observed and critical for theoretical models to match the detailed kinematics and isophotal shapes of observed systems \\citep{naab:gas, cox:xray.gas,cox:kinematics,robertson:fp,jesseit:kinematics, hopkins:cusps.ell,hopkins:cusps.fp}. We have developed a model that allows us to make specific predictions for how disks survive mergers, including both the survival of some amount of the pre-merger stellar disks and the post-merger re-formation of disks and rotationally supported components from gas that survives the merger without losing most of its angular momentum.  Figure~\\ref{fig:summary} shows that we can extend these predictions with reasonable accuracy to surviving rotational systems containing as little as $\\sim 1\\%$ of the remnant stellar mass, comparable to small central subcomponents and subtle features giving rise to e.g.\\ slightly disky isophotal shapes \\citep[see e.g.][]{ferrarese:type12,lauer:centers, mcdermid:sauron.profiles}. Owing to the combination of resolution requirements and desire to understand the fundamental physics involved, most theoretical studies of these detailed properties of ellipticals have been limited to idealized studies of individual mergers. Our results allow these to be placed in a more global context of cosmological models and merger histories. Moreover, our models allow the existence of such features (or lack thereof) to be translated into robust constraints on the possible merger histories and gas-richness of spheroid-forming mergers. Further, \\citet{hopkins:cusps.ell,hopkins:cusps.fp}, studied how the dissipational starburst components arising in gas-rich mergers are critical to explaining the observed properties and scaling relations of ellipticals, and how these components can both be extracted from and related to observed elliptical surface brightness profiles. Because both the starburst and surviving disks arise from gas in mergers, the combination of constraints from the central stellar populations, studied therein, with constraints on the survival and/or loss of gas angular momentum in mergers studied here, should be able to break some of the degeneracies in e.g.\\ pre-merger gas fractions and merger histories in order to enable new constraints and understanding of spheroid merger histories, and new tests of models for spheroid formation in gas-rich mergers.  These points relate to a number of potentially testable predictions of our models. These include the in situ formation of bulges from various types of mergers, and possible associated stellar population signatures, the presence of embedded disks in ellipticals, and how their sizes and mass fractions scale with e.g.\\ the masses and formation times of ellipticals (and how this relates to gas fractions and stellar populations in observed disks). In general, for similar merger histories, the increasing prevalence of later type galaxies (S0's and S0a's) at lower masses where disks are characteristically more gas rich is a natural consequence of our predictions here, and it is straightforward to convert our predicted scalings into detailed predictions for the abundance and mass fractions of disks given some simplified merger histories. To the extent that these processes also give rise to disk heating and/or increasing velocity dispersions in disks, or changing kinematics in both disks and bulges, then there should be corresponding relationships between galaxy shapes, kinematics, and bulge-to-disk ratios along the Hubble sequence. We investigate these possible correlations and tests in subsequent papers (in preparation).  Altogether, our results here elucidate the relevant physics important for both dissipational and dissipationless bulge formation in mergers. They support a new paradigm in which to view bulge and disk formation: gas-richness is not simply a ``tweak'' to existing models of bulge formation and disk destruction in mergers. Rather, if disks are sufficiently gas rich, the qualitative character of mergers is different, with inefficient angular momentum loss giving rise to disk-dominated remnants. This process is not inherently governed by poorly-understood feedback physics (although such feedback may be critical for establishing the conditions necessary in the first place), but rather by well-understood gravitational physics, and as such is robust and fundamentally inescapable. Aspects of galaxy populations such as the continuum of relative bulge and disk mass ratios are not simply consequences of e.g.\\ different amounts of accretion, but can arise owing to the continuum in efficiencies of disk destruction as a function of merger mass ratios, orbital parameters, and gas content. The relative (lack of) abundance of bulges at low galaxy masses and high redshift is a basic consequence of the dynamics of how gas loses angular momentum in mergers, even for similar merger histories. In short, the baryonic physics of mergers ensures that, despite the near self-similarity of the physics and merger histories of their host halos, disk and bulge formation are not a self-similar process, influenced dramatically (well out of proportion to the absolute cold gas mass fractions) by the gas-richness of the baryonic systems."
    },
    "0806/0806.4570_arXiv.txt": {
        "abstract": "Relativistic jets originating from Supermassive Black Holes (SBHs) can have a considerable impact on the Interstellar/Intergalactic Medium (ISM/IGM) within which they propagate. Here we study the interaction which a relativistic jet, and the cocoon associated with its penetration into the ISM, has on the evolution of a dense cloud, placed very near the cocoon's path, by analyzing a series of high-resolution numerical simulations, and studying the dependence on jet input power, between $P_{jet} = 10^{41}-10^{47}\\, {\\rm erg/sec}$.  The density Probability Distribution Function (PDF) within the cocoon can be described in terms of two distinct components, which are also spatially distinct: a low- and a high-density component. The former is associated with the shocked gas within the internal region of the cocoon, while the latter is associated with the outer, shocked region of the cocoon itself. The PDF of the post-shocked region is well approximated by a modified lognormal distribution, for all values of $P_{jet}$. During the \\textit{active} phase, when the jet is fed by the AGN, the cloud is subject both to compression and stripping, which tend to increase its density and diminish its total mass. When the jet is switched off (i.e. during the \\textit{passive} phase) the shocked cloud cools further and tends to become more filamentary, under the action of a back-flow which develops within the cocoon.\\\\ We study the evolution of the star formation rate within the cloud, assuming this is determined by a Schmidt-Kennicutt law, and we analyze the different physical factors which have an impact on the star formation rate. We show that, although the star formation rate can occasionally increase, on time scales  of the order of $10^{5}-10^{6}$ yrs, the star formation rate will be inhibited and the cloud fragments. The cooling time of the environment within which the cloud is embedded is however very long: thus, star formation from the fragmented cloud remains strongly inhibited. ",
        "introduction": "One of the most intriguing research areas in contemporary extragalactic astrophysics involves the study of the interplay between nuclear Black Holes (BHs) in galaxies, the (relativistic) jets which they can produce, and the Interstellar/Intergalactic Medium (ISM/IGM) within which they propagate. Observation and modeling of the propagation of jets within the ISM is an important part of this effort.\\\\ Since the seminal works of \\citet{1974MNRAS.166..513S} and \\citet{1991MNRAS.250..581F}, much effort has been dedicated to the study of the propagation of a jet into the Interstellar Medium (hereafter ISM), and to its consequences for the detectability of the jet. Relatively less attention has been paid to the impact of the jet and the cocoon generated by it on an inhomogeneous ISM. \\citet{2005MNRAS.359..781S} have performed numerical simulations of the interaction of a cocoon with a set of clouds embedded within a diffuse ISM, mostly paying attention to the evolution of the jet's morphology. More recently, \\citet[][hereafter KA07]{2007MNRAS.376..465K} have also studied the turbulence induced by the interaction of jets with cold clouds embedded in an Interstellar/Intergalactic medium. In their simulations, they show that the jet is Kelvin-Helmholtz unstable, and can shear a cold cloud embedded in the ISM/IGM, thus inducing turbulence. Their simulation setup is however different from ours, mainly because they focus their attention on a small portion of the jet to study in detail the interaction with IGM clouds and the resulting turbulence. However, as we see in our simulations, turbulence in the cocoon arises naturally due to the general dynamical expansion of the cocoon itself. \\citet{2007ApJS..173...37S,2007Ap&SS.311..293S} have taken into account the presence of large-scale density gradients in the ISM distribution, and the ability of a jet to emerge out of a galactic disc. They notice that the jet can eventually percolate through the inhomogeneous ISM, and emerges with a very disturbed morphology.\\\\ All the papers quoted made use of fixed-mesh numerical codes: in this paper, we use an Adaptive Mesh Refinement (AMR) code to follow in detail the evolution of turbulence within the cocoon produced by the jet during its propagation within the ISM/IGM. Allowing for a high refinement level, we can resolve small turbulent eddies, and study their statistical properties, and how they affect the evolution of a cloud embedded within the cocoon. We focus our attention on star formation within the cloud, and on the evolution of the density field within the cocoon, as our main aim is to understand how, within the turbulent cocoon produced by typical AGN jets at moderately high redshifts (z$\\approx 1-2$), this turbulence can modify the ongoing star formation within the cloud.\\\\ The paper is organized as follows. In section~\\ref{comp_fram} we briefly describe the numerical and computational set-up. In section~\\ref{comp_fram}, we present the numerical techniques and physical ingredients to simulate the jet and the ISM. We continue in section~\\ref{in_conf} by presenting the initial configuration and the physical parameters space spanned by the simulation, and we also discuss issues of  numerical resolution. In section~\\ref{ev_cocoon} we describe the results of the simulations and the evolution of the cocoon. We consider the evolution during both the \\textit{active} phase, while the jet is present, and the subsequent \\textit{passive} phase. Then, in section~\\ref{ev_pdf}, we discuss the density probability distribution function (PDF) of the matter within the cocoon, and we find that it is well described by a modified lognormal function. In section~\\ref{jet_cloud} we then study the evolution of the Star Formation rate (SFR) within the cloud. In section~\\ref{sec_disc} we discuss the implications of these simulations, comparing our results with those of similar papers. Conclusions are presented in the final section.\\\\ In the following, we adopt cgs units, unless otherwise explicitly stated, but we adopt kpc for lengths. The underlying cosmological parameters are taken from the 3-year WMAP best fit $\\Lambda$CDM model \\citep{2006NewAR..50..850B}: ${\\rm H}_{0} = 74^{+3}_{-3}$ Km/sec/Mpc, $\\Omega_{m} = 0.234\\pm 0.035$, $\\Omega_{b}h^{2} = 0.0223\\pm 0.0008$. The unit of time is taken to be the Hubble time for this model, i.e.: $t_{0} = 1.35\\times 10^{10}$ years. $G$, \\kb and \\mpr denote the gravitational constant, the Boltzmann constant and the proton mass, respectively.\\footnote{ For the interested reader, we have put some animations of these simulations on the website: http://web.ct.astro.it/cosmoct/web\\_group/research.html.}\\\\  \\section[]{Computational framework}\\label{comp_fram} To perform the simulations described here, we have used FLASH v. 2.5 \\citep{2000ApJS..131..273F}, a parallel, Adaptive Mesh Refinement code which implements a second order, shock-capturing PPM solver. FLASH's modular structure allows the inclusion of physical effects like external heating, radiative cooling and thermal conduction (among others). The user is permitted considerable freedom in the specification of the refinement criteria, which can be customized to reach very high spatial and temporal resolutions in selected regions.\\\\ \\noindent The main purpose of this work is to study in detail the interaction of a relativistic jet and the cocoon which it generates with pre-existing clouds, and how it can affect star formation within the cloud itself. Ideally, a full 3D simulation should have a spatial resolution high enough to resolve {\\em both} the turbulent motions within the cocoon and the thermodynamic structure of the cloud, until the end of the simulation. The former task is important when one realizes that, in addition to the direct interaction with the jet, the cloud is also significantly affected by the random interactions with the turbulent eddies present within the cocoon. The computational requirements imposed by this task, however, are prohibitive, when one considers that the smallest turbulent cells to be resolved should have a size comparable to that of the cloud (10 $h^{-1} pc$ here), and the computational box is between 2 and 4 $\\times 10^{2}$ times larger. For this reason, we have restricted ourselves to 2D simulations, where this resolution can easily be reached.\\\\ \\noindent We have included radiative cooling, described by a standard cooling function with  half solar metallicity \\citep{1993ApJS...88..253S}, extended to high temperatures (T $> 10^{7}$ K, see appendix~\\ref{append_2}). Gravity is also included, while thermal conduction and magnetic fields are not. The former could possibly be relevant for the evolution of the cocoon, although on timescales longer than those considered here \\citep{2007MNRAS.376..465K}. Regarding magnetic fields, the evidence is that, if present within the diffuse IGM, their magnitude is not larger than a few microgauss, thus making the IGM a high-$\\beta$ plasma, and the magnetic field would then not significantly affect the global dynamical evolution.\\\\ In the simulations the jet is modeled as a one-component fluid, with a density  \\nj which is a fixed ratio $\\epsilon_{j}$ of the environmental density $\\rho_{env}$. In order to suppress the growth of numerical instabilities at the jet/IGM injection interface, we adopt a steep, but continuous and differentiable transverse velocity and density profile previously adopted in simulations of jet propagation \\citep{1994A&A...283..655B,2004A&A...427..415P, 2005A&A...443..863P}: \\be v_{x,j} = \\frac{V_{j}}{\\cosh\\left\\{\\left(y-y_{j}\\right)^{\\alpha_{j}}\\right\\}} \\label{eq:vinjet} \\ee \\be n_{j} = n_{env} - \\frac{\\left(n_{env} - n_{j}\\right)}{\\cosh\\left\\{\\left(y-y_{j}\\right)^{\\alpha_{j}}\\right\\}} \\label{eq:rhojet} \\ee where: $\\alpha_{j}=10$ is an exponent which determines the steepness of the injection profile and $n_{j}, n_{env}$ denote the jet and environment number densities. This initial profile is highly sheared, and peaks around $y_{j}$, with most of the thrust $n_{jet}v_{jet}^{2}$ concentrated around the center of the profile.   The presence of a highly  sheared injection profile forces the code to refine the grid structure, populating the injection region with subgrid meshes, thus preventing the formation of numerical contact instabilities. In all our runs, the injection point of the jet is chosen to lie at the midpoint of the left boundary. \\section[]{Simulations}\\label{in_conf} Our simulations are characterized by seven parameters. Two of these characterize the diffuse ISM: $n_{env},\\, T_{env}$, two more the cloud: $r_{cl},  M_{cl}$, and finally three characterize the jet: $n_{j},\\, y_{j},\\, V_{j}$. As the SFR is mostly determined by the strength of the shock within the cloud, as previous models seem to suggest \\citep{1994ApJ...420..213K}, we have decided to mostly vary the jet's input power $P_{jet} = 0.5 A_{j}m_{H}n_{j}V_{j}^{3}$ which, together with the density contrast, determines the timescale of the cloud's disruption. We keep most of the other parameters fixed: ISM density and temperature at 1 e$^{-} {\\rm cm}^{-3}$ and $10^{7}$ K, respectively, typical of the hot ISM in the central parts of a massive elliptical at high redshift ($z\\approx 1$). Only in one run we have increased the size of the box, in order to check that the main features of the evolution do not depend on boundary conditions. All runs except one have been performed on a relatively small box (4 $h^{-1}$ kpc), while in run H2 we have used a box twice as large.\\\\ \\noindent In Table 1 we summarize the main parameters of the different runs.  The ratio between jet and environment density, {\\rm n$_{j}$/{\\rm n}$_{env}$} is fixed to $2\\times 10^{-2}$ in runs using small boxes, and decreased to a slightly lower value ($10^{-2}$) in run H2. The injection region has a width of 100 $h^{-1}$ pc in runs in the small box, and twice as much in run H2. The parameter $V_{j}$ in eq.~\\ref{eq:vinjet}, is computed once $\\rho_{j}, d_{j}$ and $P_{jet}$ are assigned. Finally, we have chosen {\\em open} boundary conditions, so the gas is free to flow out of the simulation box. The implication of this is that gas is not allowed to re-enter the box, so circulation motions on scales larger than the simulation box cannot be reproduced. \\\\ \\noindent  \\begin{figure} \\centering \\includegraphics[scale=0.35,angle=0]{fig1.eps} \\caption{Initial configuration for the runs. The plot is a magnification of the actual simulation box, showing a small region of the larger simulation box. The small, dense cloud lying along the path of the jet, placed at (0.6, 1.92) (in $\\, {\\rm kpc}\\, h^{-1}$).} \\label{fig1} \\end{figure} \\subsection[]{Initial configuration}\\label{sec:in_conf} We place a small, dense cloud at a position slightly offset from the jet's propagation direction (see Fig.~\\ref{fig1} and Table~\\ref{tab:cl}). In order to keep the same spatial resolution, the radius of the cloud is doubled in run H2, where the simulation box is larger: consequently also the initial mass of the cloud is larger. \\begin{table*} \\centering \\begin{minipage}{100mm} \\caption{Parameters of the simulation runs. Columns from left to right are as follows: run label, simulation box size, background density (runs H1-3) or central density (runs NFW), jet/background density ratio, jet power. } \\begin{tabular}{lcccc} \\hline run & {\\rm L}$_{box}$ (kpc$\\, h^{-1}$) & n$_{env}$ (cm$^{-3}$) & n$_{jet}$/n$_{env}$ & {\\rm W}$_{jet}$ (\\rm{erg/sec}) \\\\ \\hline H0 & 4 & 1 & $0.02$ & $4\\times 10^{40}$ \\\\ H1 & 4 & 1 & $0.02$ & $8.61\\times10^{41}$ \\\\ H2 & 8 & 1 & $0.01$ & $1.34\\times10^{44}$ \\\\ H3 & 4 & 1 & $0.02$ & $10^{45}$ \\\\ H4 & 4 & 1 & $0.02$ & $10^{46}$ \\\\ \\hline \\end{tabular} \\end{minipage} \\label{tab:cl} \\end{table*} We have chosen to study the effect of the jet on a cloud located very near to its path, because this allows us to check to what extent star formation is affected by the jet under the most extreme conditions. In a forthcoming paper we will study a more realistic setup, where we distribute a set of clouds with a realistic mass spectrum, and study how star formation changes according to the relative position within the cocoon associated with the jet.\\\\ \\begin{table*} \\centering \\begin{minipage}{100mm} \\caption{Cloud parameters. All distances are expressed in ${\\rm kpc} \\, h^{-1}$. From left to right, columns are as follows: simulation box size, x and y coordinates, cloud's radius and mass, the latter in ${\\rm M}_{\\sun}$.} \\begin{tabular}{ccccc} \\hline {\\rm L}$_{box}$ (kpc$\\, h^{-1}$) & x & y & ${\\rm r}_{cl}$ & ${\\rm M}_{cl}$ \\\\ \\hline 4 & 0.6 & 1.92 & 10 & $1.65\\times 10^{8}$\\\\ 8 & 1.2 & 3.84 & 20 & $1.30\\times 10^{9}$\\\\ \\hline \\end{tabular} \\end{minipage} \\end{table*} \\noindent Most of the runs were evolved up to $\\approx 10^{7}\\,$ yrs.,while the jet was active and supplying energy to the cocoon. Only in one run (H4) the jet was switched off after $10^{7}\\,$ yrs., and the further evolution of the system was followed until the cloud was completely destroyed.  \\subsubsection[]{Model of embedded cloud} We have chosen a model for the structure of the embedded clouds suggested by the numerical simulations performed by \\citet{2005ApJ...630..689B}, because the physical ingredients of these simulations are likely to be representative of the physical processes present in the ISM/IGM,  These simulations have been devised to provide a reasonable model for clouds embedded within the IGM. Different sources of heating (UV-background, the wind and radiative flux from the central QSO itself) provide a significant energy input which can promote the formation of pressure-confined clouds through thermal instability. \\citet{2005ApJ...630..689B} have shown that, for typical IGM density and temperatures, similar to those considered in the present paper, the cooling time is a small fraction of the dynamical time, and the ISM is prone to the development of small, pressure-confined clouds. These are almost isothermal, with temperatures near the lower extreme of the cooling function ($T_{cl}\\approx 10^{4}$ K). In the mass range $10^{2.5}\\leq M_{cl} \\leq 10^{7} M_{\\sun}$ they find a relationship between \\rtr and the total mass $M_{cl}$: \\be r_{t} = \\lambda M_{cl,4}^{\\beta} \\label{eq:cleq:2} \\ee where we have defined: $ M_{cl,4} = M_{cl}/10^{4} {\\rm M}_{\\sun}$. If $r_{t}$ is measured in pc, we obtain: $\\lambda = 28.87, \\beta = 0.28\\pm 0.04$. The upper limit of this mass range corresponds to the Jeans mass for these clouds, implying that they are \\textit{not self-gravitating}.\\\\ As we show in Appendix~\\ref{append_1}, a reasonable model for these clouds is the \\textit{Truncated Isothermal Sphere} \\citep[TIS][]{1999MNRAS.307..203S}. Given the cloud's mass \\mtot, the final parameters of the configuration will depend on the background ISM thermodynamic state, and we assume that the latter is described by an unperturbed ideal equation of state with density and temperature $\\rho_{cl}$ and $T_{cl}$, respectively, as appropriate to a high temperature, low density, fully ionized plasma \\citep[e.g.][]{1987smh..book.....P}. We denote the solution of the TIS equation by $F(\\zeta)$, so that the cloud density can be written as: $\\rho_{cl} = \\rho_{0}F(\\zeta)$, $\\rho_{0}$ being the normalization factor and $\\zeta = r/r_{0}$ a normalized distance. The cloud's structure is entirely specified by assigning $r_{0}, \\rho_{0}$ and the truncation distance $r_{t}$. Once the latter is determined by the radius-mass relation found by \\citet{1999MNRAS.307..203S}, the two remaining factors are determined by imposing  pressure equilibrium with the IGM, as shown in Appendix~\\ref{append_1}. \\section[]{Star formation} Although our maximum spatial resolution could allow us to resolve subparsec scales, we cannot follow star formation in any detail, because this would require the inclusion of more physics (for instance a very detailed treatment of molecular cooling) and temporal resolution a few orders of magnitude higher than the one we have adopted. Instead, we assume that the cloud is converting gas into stars at a rate determined by the Schmidt-Kennicutt law \\citep{1963ApJ...137..758S,1959ApJ...129..243S,1998ApJ...498..541K}: $\\dot{\\Sigma} = A\\Sigma^{n}$, with: $A = (2.5 \\pm 0.17)\\times 10^{-4}\\, {\\rm M}_{\\sun}\\, yr^{-1} \\,{\\rm kpc}^{-2}\\,$, $n = 1.4 \\pm 0.15$. At any time, we assume that star formation proceeds only within those regions of the simulation volume where the following criteria are satisfied: \\begin{enumerate} \\item  Mass is larger than the Bonnor-Ebert mass; \\item Temperature is less than a prescribed upper cutoff, i.e. we assume that star formation is sharply inhibited in regions having $T \\geq T_{c} = 1.2\\times 10^{4} {\\rm K}$. \\end{enumerate} The Bonnor-Ebert mass \\citep{1955ZA.....37..217E,1956MNRAS.116..351B} is defined as the largest mass which a pressure confined, gravitating cloud can reach before becoming unstable: \\be M_{be} = 1.18 \\frac{c_{s}^{4}}{\\sqrt{G^{3}p_{ext}}} = 23.55 \\frac{T_{4}^{2}}{p_{ext}^{1/2}} \\, M_{\\sun}\\label{eq:sf:1} \\ee where $p_{ext}$ is the pressure at the surface of the cloud, temperature is expressed in units of $10^{4}$ K, and we have assumed that an isothermal equation of state applies, so that the sound speed is given by: $c_{s} = (k_{B}T/m_{p})^{1/2}$. The temperature cutoff is often adopted in star formation models to exclude regions where the UV flux would be too high to permit significant star formation. Our criteria to select the star forming region are very similar to those adopted by \\citet{2008MNRAS.383.1210S}.\\\\ From now on, all characteristic properties of the cloud such as its mass or size will always be referred to the {\\em star forming region} as just defined.  \\subsection{Numerical resolution} As our main goal is that of studying in detail the evolution of star formation within the cloud, we need sufficiently high spatial resolution during each run. The outer parts of the cloud are stripped due to interaction with the jet, and it is in principle difficult to predict how much the cloud's size and mass will be reduced during the evolution. For this reason, we have applied a refinement criterion which will enable us to get high resolution even during the late phases, when the star forming region is considerably reduced. \\begin{figure} \\centering \\includegraphics[scale=1.0,angle=0]{fig1a.eps} \\includegraphics[scale=1.0,angle=0]{fig1b.eps} \\caption{Computational block distribution for two different steps of run H1, taken at $t = 2.6\\times 10^{5},\\, upper$ and $5.3248\\times 10^{5}, \\,lower$ (time in yrs).  Each little square represents a FLASH {\\em block}, and each of these contains 64 mesh cells (not reproduced in the figures). Note that the size of the region is 50 $h^{-1}$ pc, a small fraction of $L_{box}$, and the blocks represented are at the highest refinement level. } \\label{fig1res} \\end{figure} The spatial resolution in FLASH is determined by three parameters: the number of blocks in the initial decomposition, which in 2D simulations is given by: $N_{blx}\\times N_{bly}$, the number of mesh cells within each block, $nb_{x}\\times nb_{y}$, and the maximum allowed refinement level $l_{r}$. In the present work we have chosen: $N_{blx}=N_{bly}=20$, $n_{bx}=n_{by}=8$, and $l_{r}=6$. Thus, the smallest spatially resolved scale (down to mesh level) along each direction is given by: $\\Delta_{x} = \\Delta_{y} = \\left({\\rm L}_{box}/N_{blx}\\right)/\\left(2^{(l_{r}-1)}n_{bx}\\right)$, which, for ${\\rm L}_{box}=4 h^{-1}$ kpc gives: $\\Delta_{x} = 0.78125 \\, h^{-1}{\\rm pc}$. Note that the maximum refinement level around the cloud is reached already at the start of the simulation, because of the refinement criterion adopted. The approximate number of mesh cells contained within a cloud of radius $R_{cl}$,  $n_{m}$, can be estimated as: $n_{m} = int\\left(\\pi R_{cl}{2}/{\\Delta A_{m}}\\right)$, where $\\Delta A_{m} = \\Delta_{x}\\Delta_{y}$ is the area of a single mesh. In our case initially $R_{cl} \\simeq 10 h^{-1}$ pc, so we get: $n_{m} \\approx 514$. A typical block distribution is shown in Figure~\\ref{fig1res}.\\\\ By default, FLASH refines the grid at those points where one of the components of the second spatial derivative of some user-selected quantities, normalized to the square of the spatial gradient, exceeds some pre-established threshold value. The quantities we check for refinement are density, pressure and temperature. This default criterion is sufficient to resolve the highly compressed regions on the scale of the clouds we are interested in, so we do not add any additional, customized refinement criterion.  \\section[]{Evolution of the cocoon}\\label{ev_cocoon} The injection of energy into the ISM/IGM can engender turbulence, mostly because the jet is supersonic at and near the injection point. Most previous work aimed at describing the general structure and evolution of the cocoon, however, has paid more attention to the global dynamics of the cocoon. Turbulence can have a significant impact on the evolution of the embedded clouds, and for this reason we study it in more detail in the next sections.  \\subsection[]{Active phase: evolution during jet injection}\\label{sec:evol:1} Soon after the jet enters the ISM/IGM a cavity forms, and the gas which has been swept out piles up into a transition layer. In Figure~\\ref{fig_jet_h4_bw} we show an output of one of the runs at a rather advanced stage. We can easily distinguish an internal low-density, high temperature {\\textit cocoon} and a \\textit{shocked ambient gas} region, externally bounded by a tangential discontinuity with the outer ISM/IGM. \\begin{figure*} \\centering \\includegraphics[scale=0.8,angle=90]{fig3_c_v2.eps} \\caption{Density and velocity distribution at $t=8.1\\times 10^{4}$ yrs. for run H4. The logarithmic density scale ranges from $10^{-3}$ to 10$^{3}$ e$^{-}$  cm$^{-3}$. The transition region between the cocoon and the external medium is threaded by a series of transonic shock waves (\\emph{shocked ambient gas layer}). The high-density enhancement within the cocoon originates out of the stripped material from a cloud initially set near (0.6,1.92) $h^{-1}$ kpc, which has been shocked by the jet.} \\label{fig_jet_h4_bw} \\end{figure*} \\noindent One of the most interesting features we find is that these regions are also \\textit{dynamically} very different. The region containing the shocked gas is threaded by a series of weak transonic shocks, and it has on average an expansion motion, while in the cocoon a large-scale circulation parallel and opposite to the main stream of the jet develops, originating from gas reflected away from the region near the hot spot. This circulation produces shear motions which then decay into weak turbulence within the cocoon. Pre-existing clouds embedded within the ISM/IGM are heavily affected by this turbulence. \\begin{figure*} \\centering \\includegraphics[scale=0.8,angle=90]{fig4_c_v2.eps} \\caption{Density and pressure distribution for the same output of Figure~\\ref{fig_jet_h4_bw}. We show two pressure contours, corresponding to 5.77$\\times 10^{9}$ (white) and 2.88$\\times 10^{10}$ (dark grey) K$\\cdot {\\rm cm}^{-3}$. } \\label{fig_jet_h4_bw_pres} \\end{figure*} \\noindent The typical velocities of these shearing motions are large  but, due to the very high temperatures within the cocoon (T$\\approx 10^{9}-10^{11}\\,$ K) , they are only moderately transonic. \\\\ As one can notice by inspecting Figure~\\ref{fig_jet_h4_bw_pres}, the pressure within the cocoon can reach high values, because the temperatures are on average very high. Also the region around the terminal part of the jet, near the hot-spot, is subjected to high pressures, which increase steadily until the cocoon's expansion is halted by the ram pressure. We notice that the highest values of pressure are attained within the {\\em shocked ambient gas region}, mostly driven by the higher density, up to 3 orders of magnitude larger than the average density in the cocoon. The Mach number is higher near the jet, particularly near the injection point, but in the overall region (cocoon and {\\em shocked gas layer}) it never reaches values larger than $\\mathcal{M} \\approx 3.-3.5$. The dynamical evolution of the hot spot, i.e. the region between the tip of the jet and the terminal part of the cocoon, is quite interesting. Here, however, we will concentrate on those features more directly related to the interaction with clouds, leaving to further work a more detailed analysis of the features of interest relevant to the modeling of the radio jet. \\subsection[]{Passive evolution} All runs were stopped when the jet was still active, except for run H4, which was continued for about $5\\times 10^{6}\\,$ yrs. after the jet was switched off. This time was chosen to be $t \\approx 10^{7}\\,$ yrs., within the range of a typical duty cycle for the AGN. \\begin{figure*} \\centering \\includegraphics[scale=0.7,angle=0]{fig5_c_v2.eps} \\caption{Density field of run H4 at $\\Delta t=2.45\\times 10^{5}$ yrs after the jet has been shut off (i.e. at a time $t=1.00245\\times 10^{7}$ yrs since the beginning of the simulation). The cocoon and the tangential discontinuity are now dynamically very different, and their boundary is marked by a zero velocity contour.} \\label{fig_ev_sfr1} \\end{figure*} The cocoon expands up to a maximum radius determined by the jet power and the environmental ram pressure. When the duty cycle of the AGN is completed, the jet injected power decreases very rapidly, and the cocoon is not anymore fed by the jet. The evolution is mostly driven by inertial motions: a snapshot is shown in Figure~\\ref{fig_ev_sfr1}. The typical velocities are still quite high within the cocoon, but the expansion of the shocked ambient gas region has already been slowed down by the ISM/IGM ram pressure, and it is now decelerating, while slowly expanding. Notice that the cloud is threaded by arc-like internal shocks within the cloud (Figure~\\ref{fig_ev_sfr4}, {\\em left}), previously noticed in similar simulations of shock-cloud interactions \\citep{2006ApJS..164..477N}. \\\\ \\begin{figure*} \\centering \\includegraphics[scale=0.7,angle=90]{fig6_c_v2.eps} \\caption{Density evolution at two different epochs after the jet has been switched off. The system evolves towards homogeneity and equilibrium on a temporal scale mostly determined by the decay of the inertial motions. Time is measured in units of Hubble time $t_{0}$, so the plot on the left is at an epoch $t=2.835\\times 10^{5}\\,$ yrs., and the one on the right plot for $t = 1.107\\times 10^{6}\\,$ yrs.} \\label{fig_ev_sfr2} \\end{figure*} \\noindent \\noindent The further evolution of the cloud during the passive phase shows some features not emphasized in previous papers. A continuous, decelerating flow now takes place within the cocoon, primarily driven by the inertial motions, because at the typical densities of the cocoon ($n_{e} \\approx 10^{-4} - 10^{-2}\\, {\\rm cm}^{-3}\\,$) the gravitational pull is less than the inertial force. A zero total velocity line now separates the cocoon from the external ISM. In the latter the decelerating, outward directed expansion motion changes to a wind as the plasma leaks out of the simulation box and the density decreases. A similar motion, but directed in the opposite direction, drives the gas past the left vertical boundary, i.e. towards the AGN and the galaxy bulge, increasing the density within this region. As a consequence of the generalized density decrease, the cooling time increases even further, although it was already much larger than the free fall time-scale.\\\\ \\noindent This back-flow within the cocoon has some interesting consequences for the embedded clouds. During the active phase, the cloud was already subject to a significant compression from the turbulent motions, and, as a consequence, its density had also risen. Its temperature was however typically higher than $5\\times 10^{5}-10^{6}\\, {\\rm K}$, a level maintained mostly by the turbulent dissipation. However, during the passive phase, the cloud is embedded within a continuous decelerating stream, which is now more laminar than turbulent (Figure~\\ref{fig_ev_sfr2}). This flow continues to compress the cloud, which cools efficiently: its average density also increases, and its temperature decreases, a trend which can be clearly seen on the right hand side of figure~\\ref{fig_ev_sfr4}. On a longer time scale the cloud eventually is completely stripped, due to KH instabilities. \\begin{figure*} \\centering \\includegraphics[scale=0.7,angle=90]{fig7_c_v2.eps} \\caption{Evolution of density and temperature in the compressed cloud, during the active phase ({\\em left}), during the which the jet is feeding energy into the cocoon, and some time after it has been switched off ({\\em right}). The contours superimposed on the density maps trace regions of equal temperature, and correspond (from the inner- to the outermost contour) to the following temperatures: ${\\rm T} = 10^{4}, 10^{7} {\\rm and} 10^{9}\\, $K.} \\label{fig_ev_sfr4} \\end{figure*} Our simulations do not have the spatial and temporal resolution to further follow the fragmentation of this cloud, which will be treated in a separate paper. The overall action of the cloud's compression during the \\textit{active} jet phase, and of the cooling and further compression during the passive phase, can result into an occasional enhancement of the region of the cloud where favorable conditions for star formation are possible, as can be seen from figure~\\ref{fig_ev_sfr4}. The gas stripped away by KH instabilities from the cloud tends to form filamentary, high density structures, where temperatures can reach high values (${\\rm T} \\sim 10^{4} - 4\\times 10^{5}\\, $K). Star formation within these filaments, due to these high temperatures, is thus inhibited, and the cloud eventually is completely destroyed. ",
        "conclusions": "\\label{sec_conc} The simulations we have performed in this work have as a target the propagation of the jet into the inhomogeneous ISM of an host galaxy. We have restricted our interest to the central region of a model elliptical galaxy ($r < 4\\,$ $h^{-1}$ kpc), and we have studied the propagation of the jet and its interaction with a single cloud. The adoption of an AMR code like FLASH has proved crucial for studying, in the same simulation, both the large-scale properties of the jet-ISM interaction and the interaction with a small cloud. We have modeled star formation within a small, dense cloud assuming that it follows the Schmidt-Kennicutt law, and we have studied how the SFR is modified by the impact of the jet/cocoon. In general, we find that the SFR could be occasionally enhanced, but in the long run the cloud is stripped by KH instabilities and its SFR decreases. After the jet has been switched off, a laminar back-flow flow develops which continues to compress and strip the cloud, until the latter loses a significant fraction of its mass. Due to the very long cooling time of the cocoon, the cloud is embedded in a very hot, low-density medium, and this eventually results in suppression of star formation. We note that our results are preliminary being restricted to simulations in 2D for a single cloud. In further work we will extend our study to full 3D simulations and to an inhomogeneous protogalaxy."
    },
    "0806/0806.1952_arXiv.txt": {
        "abstract": "The radio source \\sng\\ has recently been identified as the youngest known Galactic supernova remnant, with a putative age of $\\sim$100 years.  We present a radio light curve for \\sng\\ based on 25 epochs of observation with the Molonglo Observatory Synthesis Telescope, spanning 20 years from 1988 to 2007.  These observations are all at the same frequency (843 MHz) and comparable resolutions ($43\\arcsec \\times 91\\arcsec$ or $43\\arcsec \\times 95\\arcsec$) and cover one fifth of the estimated lifetime of the supernova remnant.  We find that the flux density has increased at a rate of $1.22 \\pm {0.24 \\atop 0.16}$ per cent yr$^{-1}$ over the last two decades, suggesting that \\sng\\ is undergoing a period of magnetic field amplification. ",
        "introduction": "The well known deficit of young supernova remnants (SNRs) in our Galaxy has motivated many searches for these objects \\citep[for example;][]{green84,gray94d,misanovic02}. Predictions based on extragalactic supernova (SN) rates suggest that there should be around 40 SNRs younger than 2000 years in our Galaxy \\citep{cappellaro03}, and yet less than 10 have been identified \\citep{green04}. The recent identification of \\sng\\ as a young SNR with age $\\sim100$ years  by \\cite{reynolds08} and \\cite{green08}, adds to this set, and makes it potentially the youngest known Galactic supernova remnant.  Most Galactic SNRs detected at radio wavelengths are detected well after the initial supernova event ($> 10^3 - 10^4$ years). Cassiopeia A \\citep[SN~1681$\\pm$19;][]{fesen06} is the youngest SNR of known age, as $\\sim330$ years. The early stages of radio supernova (RSN) evolution have been studied in a number of bright extragalactic sources; detailed studies include, SN~1979C in M100 \\citep{weiler86}, SN~1980K in NGC~6946 \\citep{weiler86,weiler92}, SN~1993J in NGC~3031 \\citep{weiler07} and SN~1987A in the Large Magellanic Cloud \\citep{ball01,manchester02}. However, the oldest RSN of known age is SN~1923A \\citep{eck98}, which leaves a critical gap in our knowledge of supernova evolution at intermediate ages of $\\sim50-300$ years. Hence the age estimate of $\\sim100$ years for \\sng\\ makes it a useful probe of this period in SNR evolution.  \\sng\\ was first identified as a potential young SNR based on its small angular size of 1.2\\arcmin\\ \\citep{green84}. Recent Chandra observations by \\cite{reynolds08} showed that the remnant had expanded by $16\\pm3$ per cent between earlier VLA observations in 1985 and the Chandra observations in 2007. This was confirmed by VLA follow-up observations by \\cite{green08} which gave an expansion of $15$ per cent over 23 years (or an expansion rate of $0.65$ per cent yr$^{-1}$). This led to an age estimate of $\\sim100$ years (or at most $150$ years) for \\sng. \\citet{green08} also use archival data from a variety of instruments (over frequencies 332 to 5000 MHz) to suggest that \\sng\\ has brightened by a rate of $\\sim 2$ per cent per year over the same period.  We present 20 years of radio observations from the Molonglo Observatory Synthesis Telescope \\cite[MOST;][]{mills81,robertson91}. These observations were carried out at constant frequency (843 MHz) and comparable resolutions ($43\\arcsec\\times91\\arcsec$ or $43\\arcsec\\times95\\arcsec$). They show that \\sng\\ has been increasing in brightness by a rate of $1.22\\pm {0.24 \\atop 0.16}$ per cent yr$^{-1}$ over this period. \\changes{Our new estimate has the advantage that the measurements were taken with the same instrument, whereas the \\citet{green08} estimate was based on observations from a range of instruments, compiled from the literature.} ",
        "conclusions": "\\label{s_discuss} Only a fraction of SNe have detectable radio emission. Observations of these RSNe \\citep[see, for example;][]{eck02} show that radio emission caused by the interaction between the shock and the circumstellar medium of the progenitor star is first detected within days to months of the initial explosion. As the shock travels into regions of decreasing opacity, the radio emission brightens, peaking between days to years after the explosion. After reaching their maximum brightness, RSNe then show a power law decrease in flux density with time \\citep{weiler02}. This is demonstrated by the light curves of SN~1923A, SN~1950B, and others, shown in Fig.~\\ref{f_lumcmp}. The radio behaviour of RSNe is explained well by the models of \\citet{chevalier82,chevalier98}.  There are even fewer Galactic SNRs of known age. Cassiopeia A, the youngest SNR of known age at $\\sim330$ years, Kepler (SN~1604) and Tycho (SN~1572) are shown in Fig.~\\ref{f_lumcmp}. The radio emission we detect from SNRs is due to the interaction between the shock and the interstellar medium. Hence the radio emission from a young SNR will increase once it has swept up enough of the surrounding interstellar medium. The timescale of this brightening is predicted to be $\\sim100$ years \\citep{gull73,cowsik84}.  Between the youngest SNR of known age (Cas A) and the oldest RSNe of known age (SN~1923A), there is a gap in our observational evidence which makes it hard to probe the period after the fading of the radio emission from the RSN and before the SNR switches on. \\citet{eck02} conducted a search for radio emission from SNe of a range of known ages, but did not detect any of the SNe that occurred prior to SN~1923 (for example, SN~1885A and SN~1909A). If the estimated age of $100-120$ years for \\sng\\ is correct, then our measurements help constrain the flux density evolution in this intermediate time range.  To explain the observed expansion rate measured by \\cite{reynolds08} and the new light curve presented here, we need to reconsider some of the standard assumptions made when predicting the luminosity evolution of SNRs. The radio luminosity $L_\\nu$ of a synchrotron source at a given frequency $\\nu$ is a function of the energy spectrum of the ultrarelativistic electrons, the magnetic field $B$ present in the source, and the volume of the source $V$, as given in \\cite{longair94} \\[ L_\\nu = A(\\alpha)V\\kappa B^{1+\\alpha}\\nu^{-\\alpha} \\; , \\] where $\\alpha$ is the spectral index, $A(\\alpha)$ is a constant, and $\\kappa$ is defined in terms of the electron energy spectrum per unit volume \\[ N(E)  dE = \\kappa E^{-(2\\alpha + 1)} dE \\; . \\] Hence for a given frequency \\[ L_\\nu \\propto V \\kappa B^{1+\\alpha} \\; . \\] If we make the assumptions that the expansion of the remnant is adiabatic, and that magnetic flux freezing is applicable, then the magnetic field strength decreases as $B\\propto r^{-2}$. Following \\cite{longair94} we can derive that $V\\kappa \\propto r^{-2\\alpha}$ and so \\[ L_\\nu \\propto r^{-2(2\\alpha+1)} \\; . \\] Now the time dependence of the luminosity is proportional to the time dependence of the source radius, since the spectral index $\\alpha$ should not change during adiabatic expansion. Expressing the change in radius with time as a power law with expansion parameter $m$, $r\\propto t^m$, then \\[ L_\\nu \\propto t^{-2m(2\\alpha + 1)} \\; . \\] During an ideal free expansion phase $r\\propto t$, and in the Sedov phase $r\\propto t^{0.4}$. Using the spectral index of $\\alpha = 0.63$ given by \\cite{reynolds08} we would expect the luminosity to change as $L_\\nu\\propto t^{-4.5}$ in free expansion, or $L_\\nu\\propto t^{-1.8}$ in Sedov expansion.  If we fit the measured expansion rate given by \\cite{reynolds08} with a power law (assuming an age of 100 years in 2008), we get $r\\propto t^{0.55}$, implying $L_\\nu\\propto t^{-2.5}$ which is intermediate between free expansion and the Sedov phase. Using our new light curve data we get $L_\\nu\\propto t$. Clearly, to produce the observed increase in luminosity, either the magnetic field strength or the energy density of the relativistic electrons (or both) must be increasing with time, rather than decreasing as assumed above. \\changes{ A similar conclusion was drawn by \\citet{green08} }. This scenario is supported by what we know about older SNRs --- extrapolating backwards from the observed magnetic field strength in Cas A, the assumption of magnetic flux freezing would result in an implausible magnetic field strength in the progenitor star \\citep{longair94}. Hence at some period in its evolution, the magnetic field strength must have undergone a process of amplification. Simulations by \\citet{jun99} show that a possible mechanism for this amplification is Rayleigh-Taylor instability in the interaction region between the SNR shock and a surrounding cloud of interstellar gas. This happens on timescales of $\\sim100-300$ years after the inital explosion. Our observations suggest that \\sng\\ has entered such a phase, in which $\\kappa B$ is growing."
    },
    "0806/0806.0603_arXiv.txt": {
        "abstract": "In writing a covariant effective action for single field inflation, one is allowed to add a Gauss-Bonnet and axion-type curvature couplings. These couplings represent modifications of gravity, and are the unique higher-curvature terms that lead to second order equations of motion in four dimensions. In this paper we study the observational consequences of such couplings for models with large non-gaussianities. Our focus is on the Gauss-Bonnet term.  In particular, we study an effective action where the scalar Lagrangian is a general function of the inflaton and its first derivative.   We show that, for large non-gaussianities,  one can write $f_{NL}$ in terms of only three parameters.  The shape of $f_{NL}$  is also studied, and we find that it is very similar to that of k-inflation.  We show that the Gauss-Bonnet  term enhances the production of gravitational waves, and allows a smaller speed of sound for scalar perturbations. This, in turn, can lead to larger non-gaussianities which can be constrained by observations. Using current WMAP limits on $f_{NL}$ and the tensor/scalar ratio, we put constraints on all parameters. As an example, we show that for DBI inflation, the Gauss-Bonnet coupling leads to an interesting observational window with both large $f_{NL}$ and a large amplitude of gravitational waves. Finally, we show that the Gauss-Bonnet coupling admits a de-Sitter phase with a relativistic dispersion relation for scalar perturbations. ",
        "introduction": "Cosmology has entered an era of unprecedented progress. High precision measurements of the cosmological parameters have led to a coherent picture of the history of our universe that seems to favor the inflationary paradigm \\cite{Guth}. Moreover, a future detection of large non-gaussianity in the cosmic microwave background (CMB) would falsify the simplest inflationary scenario, namely, single field slow-roll inflation \\cite{Komatsu, Komatsu1, nongauss}.  On the theoretical side, there has been great activity in trying  to produce large non-gaussianities in single and multiple-field inflationary models.  For  single field inflation, large non-gaussianities are easiest to produce in models with a small speed of sound (see e.g. \\cite{Kach, largefnl}).  On a parallel set of developments, there has been recent interest in developing a systematic effective field theory of single field inflation \\cite{Senatore, Weinberg}. In ref. \\cite{Senatore}, such approach was applied directly to the Lagrangian describing the perturbations around the inflationary solution. The effective action can be viewed as an expansion  in powers of $(g^{00} +1)$ and the extrinsic curvature $K_{a b}$ of the constant time hypersurfaces. Such approach is quite general, and provides a straightforward way of calculating all CMB observables directly from the effective action for the fluctuations.   On the other hand, one would like to understand how the various terms in the effective action for the fluctuations relate to the effective action of the inflaton itself. A method to build such an effective action was introduced by Weinberg in \\cite{Weinberg}.  In this approach, one considers all marginal and irrelevant operators involving the inflaton and the metric. Among these terms, there are higher curvature invariants coupled to the inflaton. Generically, such terms will contain higher time derivatives on the fields which need to be eliminated using the first order equations of motion. Otherwise, one would be propagating more degrees of freedoms than intended. Weinberg showed that, after such eliminations and to leading order in the derivative expansion, the resulting action for the inflaton takes the familiar k-inflation type form plus two extra couplings between the inflaton and the Weyl tensor.  Such extra couplings can be written instead in terms of the Gauss-Bonnet tensor and an axion-type coupling. In this way, one has an effective action that leads to second order equations of motion explicitly.   One can then ask if it is possible to re-sum such an expansion. Moreover, one would like to write down a general {\\it local} action for the inflaton coupled to gravity that leads to second order equations of motion for all fields. By abuse of notation we will call such re-summation a ``UV completion\" of the effective theory. An advantage of having such action is that one can then build directly the low energy effective action for the fluctuations as in \\cite{Senatore}, have a clear physical interpretation of the various couplings, and a better assessment of their relative importance. If  we also insist in preserving general covariance in the UV completion, it is not hard to see that an obvious candidate for such action is \\footnote{We cannot claim that this is the most general action that propagates only the spin two graviton and a scalar degree of freedom.  Nevertheless, most alternative  constructions will either break general covariance or involve another scalar field, even if it is not dynamical (see e.g. \\cite{cuscuton}).} \\be \\label{S} S =  \\int \\sqrt{-g}\\left[ \\frac{1}{2}R + P(X,\\phi) + V_1(\\phi) E_4 +  V_2(\\phi) \\epsilon^{abcd} R_{ab}^{\\;\\;\\; ef} R_{cdef}\\right]\\;.\\ee where $E_4$ is the Gauss-Bonnet combination, \\be E_4 = R_{abcd}R^{abcd} - 4 R_{ab} R^{ab} + R^2\\;,\\ee and $X = -\\frac{1}{2} \\nabla_a \\phi \\nabla^a \\phi$ is the kinetic term. In writing (\\ref{S}) we assume that we work in the Einstein frame. The second term in the action is the familiar k-inflation type \\cite{kinf}.  The last two couplings in the effective action would be topological invariants in four dimensions if both potentials $V_1,V_2$ were constant. This is the reason why they lead to second order equations of motion for general $V_i$. Note that if these potentials depended on $X$,  one would end up with equations of motion depending on more than two time derivatives.  The last two terms in the action (\\ref{S}) represent modifications of Einstein's gravity, and as such, they have been studied numerous times (see e.g. \\cite{modified}). Moreover, such couplings are known to arise in string theory \\cite{threshold}. The last term in (\\ref{S}) is of the axion-type and it was studied long ago in \\cite{hwang1}, where it was shown that such coupling does not affect the evolution of scalar fluctuations to quadratic order. We have verified that this is still true to cubic order. Therefore, we discard this coupling in what follows \\footnote{Note, however, that the axion-type coupling does affect the gravitational waves by giving an extra helicity dependence to the tensor power spectrum \\cite{Weinberg, hwang1}.}. The Gauss-Bonnet term, on the other hand, has been studied many times in the context of Dark Energy (e.g. \\cite{modified}).  For other studies in the context of early cosmology see \\cite{Sato}.   This term does contribute to scalar fluctuations and it is the main focus of this work.     The purpose of this paper is to study the observational signatures of the Gauss-Bonnet coupling in the context of inflation. Moreover, our main interest will be in models will large non-gaussianities.  In the following, we compute the non-gaussianity parameter $f_{NL}$ using the action (\\ref{S}) in the limit of a small speed of sound. Moreover, we perform such calculation to leading order in the slow roll parameters, but to all orders in the ``strength\" of the Gauss-Bonnet coupling, defined as \\be \\label{epsilon1}g \\equiv 8 V_1'(\\phi) H \\dot\\phi\\;,\\ee where $H$ is the Hubble parameter. In making the calculations, we assume that the parameter $g$ is slowly varying in time. We find that for large-nongaussianities,  $f_{NL}$ can be written in terms of only three parameters. We study the shape of of $f_{NL}$ as a function of these parameters.   We find that the shape of $f_{NL}$ is always very close to that of k-inflation, even in the limit $g \\rightarrow \\infty$.  We discuss in which cases deformations from this shape might be observed.  The spectrum of gravitational waves is also studied, and we find an enhancement due to the Gauss-Bonnet term.  Using WMAP limits on the equilateral  $f_{NL}$ and the tensor/scalar ratio, we  put constraints on the different parameters.   For the particular case of DBI inflation, we show that one is left with only a two-parameter family of $f_{NL}$. In this case, one can put a more precise constraint on the Gauss-Bonnet coupling: \\bes g_\\text{DBI} \\lesssim 3\\;.\\ees We also discuss implications of the Gauss-Bonnet coupling on the Lyth bound in the context of DBI inflation \\cite{Lyth}.  An interesting aspect of the Gauss-Bonnet coupling is that scalar perturbations are non-trivial even in a de-Sitter background, just like the do in the Ghost Condensate \\cite{ghost}.  However, we find that in our case, quadratic scalar fluctuations have the familiar relativistic dispersion relation $\\omega \\sim k$ instead of the non-relativistic one $\\omega \\sim k^2$ of the Ghost Condensate. This matches perfectly with the new de-Sitter limit found in \\cite{Senatore} using the effective action of the fluctuations. Therefore, the Gauss-Bonnet coupling is precisely the modification of gravity that leads to such limit.   The paper is organized as follows. In section II we study the equations of motion for the background solution and the quadratic scalar fluctuations. We also point out the different limits used in the calculations. In section III we calculate $f_{NL}$ and study its shape.  In section IV we calculate the gravitational wave spectrum. We discuss the various constraints on the parameter space.  In section V we study the case of DBI inflation. Finally, we close with some final comments and future directions in section VI. ",
        "conclusions": "In this article we have studied WMAP constraints on modifications of gravity due to a Gauss-Bonnet coupling in single-field inflation. This is the most general modification of gravity that leads to second order equations of motion, and that also affects the spectrum of scalar fluctuations. We showed that in the slow roll limit, and for a very general class of models with action (\\ref{S}), a large $f_{NL}$ can be written in terms of three parameters: $c_s$, $\\lambda$ and $g$ the Gauss-Bonnet coupling.  We found that the Gauss-Bonnet term has little effect on the shape of non-gaussianities, but it amplifies the tensor power spectrum. Thus, given current limits on the tensor/scalar ratio $r$, we found that large values of the Gauss-Bonnet coupling would require small values of the speed of sound, and hence large non-gaussianities. Using current WMAP limits on $r$ and $f_{NL}$ we were able to constraint the parameter space of such models.  To give better constraints on $g$ we studied a particular model: DBI inflation. In this case we obtained a precise bound on this coupling, Eq. (\\ref{bound}). Moreover, we saw that in this model, a non-zero Gauss-Bonnet couplings leads to an interesting observational window with both large non-gaussianities and a large amplitude of gravitational waves. We also studied how the conditions for the smallness of the slow roll parameters translate to constraints on the scalar potentials. Possible violations of the Lyth bound were also studied. We found that the bound can be violated for large values of $g$. However, for DBI inflation one has an observational restriction of $g \\lesssim 3$, and so the bound is roughly $\\Delta \\phi \\gtrsim 0.5$. Nevertheless, we found that for an AdS throat the fractional change in the scalar field is very small: $\\Delta\\phi/\\phi \\ll 1$. It would be interesting to study the higher curvature corrections to the DBI action, to see if one can realize a Gauss-Bonnet driven inflation in a controlled way. Some of these  corrections were derived in \\cite{Bachas}.   Another interesting aspect of the Gauss-Bonnet term is that scalar fluctuations can exist in a de-Sitter background. Moreover, they have a relativistic dispersion relation, unlike Ghost Inflation \\cite{ghost}. This matches with the new de-Sitter limit found in \\cite{Senatore} using the effective action of the scalar fluctuations.  One might wonder if there are other de-Sitter limits of inflation. In \\cite{Senatore}, it was argued that the answer is negative as any such limits will not make sense as an effective field theory for the fluctuations. However, the authors of \\cite{Senatore} only considered models with one scalar degree of freedom. If we ask about other modifications of gravity, it is well known that  one needs to add more degrees of freedom to the theory. It would be interesting to put constraints in other types of modified gravity using large non-gaussianities."
    },
    "0806/0806.3436_arXiv.txt": {
        "abstract": "Numerical simulations with self-similar initial and boundary conditions provide a link between theoretical and numerical investigations of jet dynamics. We perform axisymmetric resistive magnetohydrodynamic (MHD) simulations for a generalised solution of the Blandford \\& Payne type, and compare them with the corresponding analytical and numerical ideal-MHD solutions. We disentangle the effects of the numerical and physical diffusivity. The latter could occur in outflows above an accretion disk, being transferred from the underlying disk into the disk corona by MHD turbulence (anomalous turbulent diffusivity), or as a result of ambipolar diffusion in partially ionized flows. We conclude that while the classical magnetic Reynolds number $R_{\\rm m}$ measures the importance of resistive effects in the induction equation, a new introduced number, $\\rbeta=(\\beta/2)R_{\\rm m}$ with $\\beta$ the plasma beta, measures the importance of the resistive effects in the energy equation. Thus, in magnetised jets with $\\beta<2$, when $\\rbeta \\la 1$  resistive effects are non-negligible and affect mostly the energy equation. The presented simulations indeed show that for a range of magnetic diffusivities corresponding to $\\rbeta \\ga 1$ the flow remains close to the ideal-MHD self-similar solution. ",
        "introduction": "Collimated outflows of plasma observed to emerge from the vicinity of a wide spectrum of cosmic objects are still a challenge for observational and theoretical astrophysics. These outflows play a key role in the transport of angular momentum and energy of the accreted gas facilitating thus, for example, star formation. Nevertheless, when new observations put more and more severe constraints on models, these seem to be still too rudimentary to provide sophisticated answers.  The starting point of the modeling of jets are the ideal  MHD equations, which can be solved analytically by assuming axisymmetry, time-independence and the self-similarity ansatz. Analytical models of ideal MHD disk winds (Blandford \\& Payne 1982, re-visited in Vlahakis et al. 2000; hereafter V00), provide not only the first insight into the physics of such outflows but equally important they can be used as a test bed of more sophisticated  simulations of the resistive MHD system via various numerical codes. In Vlahakis \\& Tsinganos (1998) general classes of self-consistent ideal-MHD solutions have been constructed. Two sets of exact MHD outflow models have been found: meridionally and radially self-similar ones. Previously known studies were recognised to belong in this more general classification of all available analytical models. In particular, the V00 study remedied the physically unacceptable feature of the Blandford \\& Payne (1982) terminal wind solution which was not causally disconnected from the disk (see also Ferreira \\& Casse (2004) when a resistive disk is included).  Among the basic problems which still remained to be solved however were the common deficiency that all radially self-similar models had, namely, a cut-off of the solution at small cylindrical radii and also at some finite height above the disk where they were unphysical. The reason for such behaviour is a strong Lorentz force close to the system's axis. The invalid analytical solution very close to the axis has been corrected numerically. Also, a search in the numerical simulations for solutions at larger distances from the disk has been performed, in Gracia et al. (2006; hereafter GVT06), where ideal-MHD numerical simulations with the V00 solution as initial condition have been performed using the NIRVANA code (version 2.0, Ziegler 1998). These results have been verified also by using the PLUTO code (Mignone et al. 2007) in the ideal MHD simulations by Matsakos et al. (2008).  \\begin{figure} \\hspace{1.3cm}\\includegraphics[width=6.cm, height=1.cm]{initial_cb.eps} \\includegraphics[width=7.5cm, height=12cm]{initial.eps} \\caption{Initial conditions for our numerical simulations. The solid lines represent logarithmically spaced isocontours of density. It is also shown in colour grading. The dashed lines show poloidal fieldlines or flowlines. The dotted lines indicate, from top to bottom, the position of the fast-magnetosonic ({\\em small dots}), the Alfv\\'en, and the slow-magnetosonic surface ({\\em large dots}), respectively. } \\label{pocet} \\end{figure}  The next step in the exploitation of the available analytical solutions has been the investigation of the dynamical connection of the outflow to the underlying disk (see e.g., K\\\"{o}nigl, 1989; Wardle \\& K\\\"{o}nigl 1993; Li 1995; Ferreira 1997; Casse \\& Keppens 2004; Zanni et al. 2007), providing some understanding of the formation of jets from the accretion disk. At the same time, since the disk is naturally resistive, the need emerged to go beyond the ideal MHD regime. With the magnetic field included, the effects of the magnetic resistivity, both numerical and physical, had to be addressed, discriminated and analysed.  In fully ionized disks an anomalous turbulent diffusivity has to be invoked to allow for accretion of matter crossing a large-scale magnetic field. This anomalous turbulent diffusivity may be present in the outflow as well, at least at distances close to the disk. In partially ionized disks the physical conductivity is properly described by a tensor, consisting of three distinct parts corresponding to the ambipolar diffusion, the Hall effect, and the Ohmic dissipation (e.g. Wardle \\& Ng, 1999; Salmeron et al., 2007). The dominant mechanism in the outflow above the disk is most likely the ambipolar diffusion (e.g., Sano \\& Stone 2002; Kunz \\& Balbus 2004; Wardle 2007), and can be appropriately described only by a multifluid MHD. Nevertheless, in both cases (turbulent or ambipolar diffusion), a scalar conductivity can capture the basic characteristics of the breakdown of ideal MHD related to the magnetic field diffusion and the resistive heating.  The effects of the resistive heating in the formation and acceleration of jets from resistive disks or tori, have been studied in Kuwabara et al. (2000, 2005), concluding that Joule heating is not playing an essential role in jet formation. However, in these studies only low values (lower than the critical value that we define below) of resistivity were examined. Resistive effects have been also studied in Fendt \\& \\v{C}emelji\\'{c} (2002; hereafter FC02). In this case, however, the energy equation was not solved and a polytropic equation of state has been assumed instead, such that the effects of finite resistivity have not been directly incorporated in the energetics of the problem. Safier (1993a) working in the ambipolar diffusion regime of outflows associated with young stars, found that the diffusive term in the energy equation cannot be neglected. The heating of the gas could have significant observational consequences (see e.g. Safier 1993b; Martin 1996b; Cabrit et al., 1999; Garcia et al., 2001a,b; O'Brien et al., 2003; Shang et al., 2004). Safier's (1993a) work also shows why the cooling term in the energy equation and the ionization balance equation need to be considered in a full investigation. He was able to uncover a strong feedback mechanism between the gas temperature and the ionization fraction (which scales inversely with the ambipolar diffusion heating rate). The result of the heating depends also on the geometry of the flow which strongly affects the adiabatic cooling. In a spherical outflow this cooling effectively counters the Joule dissipation heating (Ruden et al., 1999), while in a disk-driven jet with a small streamline divergence the adiabatic cooling is relatively unimportant (at least initially). Interestingly, Joule dissipation can play a role (although generally not a dominant one) also in magnetically guided accretion problems (e.g., Martin 1996a).  Numerical resistivity is implicitly present in any numerical simulation, and its various effects need to be identified and studied in detail. This is one of the main aims of this paper. The other one is to investigate the effects of small physical resistivity, examine the stability of the resistive jet solutions and define when the resistivity is \"small\" and when it becomes \"large\". The novel approach followed here is that the analytical solution is a stationary reference solution, something which was lacking in the previous investigations of resistive MHD flows, e.g. FC02.  The structure of the paper is as follows. The analytical expressions and their modification in the setup for numerical simulations are first presented in Sec. 2. Then the resistive-MHD solutions (Sec. 4) are systematically compared to the ideal-MHD ones of Sec. 3. In Sec. 5 we introduce an extension of the magnetic Reynolds number which quantifies the transition from an ideal-like behaviour of the low diffusivity solutions for values of the diffusivity $\\eta$ below a critical value $\\eta_{crit}$, to a transient and erratic behaviour of the solutions for $\\eta > \\eta_{crit}$. The main possible consequences for astrophysical outflows are briefly discussed. A summary of the main results is given in the last Sec. 6.  \\section[]{Problem setup}\\label{sec2} \\subsection{Governing equations} The resistive-MHD equations solved by the NIRVANA code are, in SI units: \\beqa \\frac{\\partial \\rho}{\\partial t} + \\nabla \\cdot (\\rho \\vec{V})=0 \\,,\\\\ \\rho\\left[ \\frac{\\partial\\vec{V} }{\\partial t}+ \\left( \\vec{V }\\cdot \\nabla\\right) \\vec{V} \\right] + \\nabla p + \\rho\\nabla \\Phi - \\frac{ \\nabla \\times \\vec{B}}{\\mu_0} \\times \\vec{B} = 0 \\label{mom2} \\,, \\\\ \\frac{\\partial\\vec{B} }{\\partial t}- \\nabla \\times \\left( \\vec{V} \\times \\vec{B}-\\eta \\nabla \\times \\vec{B} \\right)= 0 \\,, \\label{faraday}\\\\ \\rho \\left[ {\\frac{\\partial e}{\\partial t}} + \\left(\\vec{V} \\cdot \\nabla\\right)e \\right] + p(\\nabla \\cdot\\vec{V} ) - \\frac{\\eta}{\\mu_0} \\left( \\nabla \\times \\vec{B} \\right)^2= 0 \\,, \\label{enn}\\\\ \\nabla \\cdot \\vec{B}=0 \\,, \\eeqa where $\\vec{V}$ is the flow velocity, $\\vec{B}$ is the magnetic field, $(\\rho, P)$ are the gas density and pressure, and $\\Phi=-{\\cal GM}/r$ is the gravitational potential of the central mass ${\\cal M}$. The internal energy (per unit mass) is related to the pressure and density by \\beq e=\\frac{1}{\\gamma-1} \\ \\frac{p}{\\rho} \\,, \\eeq where $\\gamma$ is the effective polytropic index. The magnetic diffusivity $\\eta$ is assumed constant and is related to the resistivity $\\rho_c=\\mu_0\\eta$, where $\\mu_0$ is the permeability of vacuum.  \\subsection{Initial and boundary conditions} For the initial and boundary conditions of the simulations the self-similar solution of V00 is used. The assumptions of steady-state, axisymmetry and radial self-similarity result in the following expressions for the physical quantities in spherical ($r$, $\\theta$, $\\phi$) and cylindrical ($Z=r\\cos\\theta$, $R=r\\sin\\theta$, $\\phi$) coordinates: \\beqa \\frac{\\rho}{\\rho_0}=\\alpha^{x-3/2}\\frac{1}{M^2}\\,, \\label{rhoss}\\\\ \\frac{p}{p_0}=\\alpha^{x-2}\\frac{1}{M^{2\\gamma}}\\,,\\\\ \\frac{\\vec{B}_p}{B_0}=-\\alpha^{\\frac{x}{2}-1}\\frac{1}{G^2}\\frac{ \\sin\\theta}{\\cos(\\psi+\\theta)} \\left(\\cos\\psi \\hat{R}+\\sin\\psi \\hat{Z}\\right) \\,, \\label{bpss}\\\\ \\frac{\\vec{V}_p}{V_0}=-\\alpha^{-1/4}\\frac{M^2}{G^2}\\frac{ \\sin\\theta}{\\cos(\\psi+\\theta)} \\left(\\cos\\psi \\hat{R}+\\sin\\psi \\hat{Z}\\right) \\,, \\\\ \\frac{{B}_\\phi}{B_0}=-\\lambda\\alpha^{\\frac{x}{2}-1}\\frac{1-G^2}{G(1-M^2)}\\,,\\\\ \\frac{{V}_\\phi}{{V}_0}=\\lambda\\alpha^{-\\frac{1}{4}}\\frac{G^2-M^2}{G(1-M^2)} \\,, \\label{vphiss} \\eeqa where $\\displaystyle \\alpha=\\frac{R^2}{R_0^2 G^2}$, $(M\\,, G\\,, \\psi)$ are functions of $\\theta$, and \\beq \\displaystyle{ V_0=\\frac{1}{\\kappa} \\sqrt{\\frac{ \\cal G M }{R_0}} \\,, \\quad \\rho_0=\\frac{B_0^2}{\\mu_0 V_0^2} \\,, \\quad p_0=\\mu\\frac{B_0^2}{2\\mu_0} \\,. } \\label{norms} \\eeq Here we decomposed vector quantities in poloidal (index $p$) and toroidal (index $\\phi$) components. We set the solution parameters to $(x\\,, \\lambda^2, \\mu\\,, \\kappa\\,, \\gamma)= (0.75\\,, 136.9\\,, 2.99\\,, 2\\,, 1.05)$, as in the V00 solution.  The diffusivity $\\eta$ (which is assumed constant throughout the domain) is normalised as $\\eta = \\etahat \\ V_0 R_0=\\etahat \\sqrt{{\\cal GM}R_0}/\\kappa$, with $\\etahat$ dimensionless. \\begin{figure*} \\includegraphics[width=5.5cm,height=9cm]{density_early.eps} \\includegraphics[width=5.5cm,height=9cm]{density_intermediate.eps} \\includegraphics[width=5.5cm,height=9cm]{density_final.eps} \\caption{Ideal-MHD simulations before, during and after the relaxation, for the resolution of $R\\times Z = (128\\times 256)$ grid cells $=([0,50]\\times [6,100])R_0$. Lines denote thirty logarithmically spaced isocontours of density. The times of simulations are few thousands, few ten thousands and larger than few hundred thousands of the Courant time steps, {\\em Left} to the {\\em Right} panel, respectively. } \\label{idealdens} \\end{figure*} The self-similar solution breaks down near the rotation axis. This becomes evident from the fact that all physical quantities are proportional to a power of the function $1/\\alpha$, which is divergent on the axis ($R=0$), see Eqs. (\\ref{rhoss}-\\ref{vphiss}). In addition, the analytical solution of V00 is not provided for $\\theta$ smaller than 0.025 rad, measured from the axis.  To perform numerical simulations in a computational box with the symmetry axis included, we need to modify/extrapolate the analytical solution. Near the axis, we extrapolated the missing analytical solutions for the tabulated functions $G$, $M$ and $\\psi$, as described in GVT06. A similar result can be achieved with less involved extrapolation, as shown in Matsakos et al. (2008). Modification of the functions G and M means also that the pressure/energy is modified near the axis.  For the magnetic field in the vicinity of the symmetry axis there is an additional problem. With the extrapolated functions $G$ and $\\psi$ the magnetic field as given by Eq.~(\\ref{bpss}) is not divergence-free. This leads to the need for suitable modification of the initial magnetic field. A simple modification is to compute the $B_Z$ component from the $\\hat{Z}$ component of the self-similar expression \\beq \\vec{B}_p=\\frac{B_0 R_0^2}{x} \\nabla \\times \\left( \\alpha^{x/2}\\frac{\\hat{\\phi}}{R} \\right) \\,, \\eeq and subsequently the radial component $B_R$ by solving the $\\nabla\\cdot \\vec B=0$ with boundary condition $B_R(R=0)=0$. \\begin{figure*} \\includegraphics[width=5.5cm,height=8cm]{num_surfaces} \\includegraphics[width=5.5cm,height=8cm]{num_fieldlines} \\includegraphics[width=5.5cm,height=8cm]{num_integrallines} \\caption{Illustration of the effect of {\\em numerical} resistivity, i.e. grid resolution. {\\em Left:} The slow-magnetosonic, Alfv\\'enic, and fast-magnetosonic critical surfaces (from bottom to top). Different line types represent the final states of simulations with resolution $128\\times 256$ ({\\em dotted/red}), $256\\times 512$ ({\\em dot-dashed/blue}), and $512\\times 1024$ ({\\em dashed/green}). The {\\em solid/black} lines show the initial-state critical surfaces in the high resolution reference simulation. {\\em Middle:} The shapes of two different magnetic flux surfaces, i.e. poloidal magnetic field lines, for the same resolutions and line types as described above. The solid/black lines represent again the initial state of the high resolution reference simulation. {\\em Right:} Same as in the {\\em Middle} panel, but for the energy $(E)$ integral lines. } \\label{num_lines} \\end{figure*} \\begin{figure*} \\includegraphics[height=8cm]{num_Q_on_E} \\includegraphics[height=8cm]{num_Q_on_A} \\caption{ Illustration of the effect of {\\em numerical} resistivity, i.e. grid resolution, on the alignment of the MHD integrals with the magnetic flux surfaces. {\\em Left:} As an example for the MHD integrals, the entropy $Q$, normalised to its value at large distances, is plotted along the inner total energy integral line shown in Fig.~\\ref{num_lines}, for the initial state ({\\em solid/black}) of the high resolution reference solution, and the final states at resolutions $128\\times 256$ ({\\em dotted/red}), $256\\times 512$ ({\\em dot-dashed/blue}), and $512\\times 1024$ ({\\em dashed/green}). {\\em Right:} Same as in the {\\em Left} panel but here $Q$ is plotted along a magnetic field line (instead of the energy integral line).} \\label{num_integrals} \\end{figure*}  \\begin{figure*} \\includegraphics[height=8cm]{num_Eterms_on_E} \\includegraphics[height=8cm]{num_Eterms_on_A} \\caption{Illustration of the effect of {\\em numerical} resistivity, i.e. grid resolution, on the contributions to the energy integral $E$. {\\em Left:} Split-down of the energy contributions along the two integral lines shown in Fig.~\\ref{num_lines}. The upper set of curves corresponds to the inner integral line, the lower set of curves to the outer integral line. Colours code the resolution as in Fig.~\\ref{num_lines}. The different line types represent energy $E$ ({\\em solid}), kinetic energy ({\\em dotted}), enthalpy ({\\em dot-dashed}) and Poynting ({\\em dashed}), respectively. The gravitational energy is not shown (it is orders of magnitude smaller). {\\em Right} Same as {\\em Left} panel, but plotted along the field line instead of the integral lines. } \\label{num_eterms} \\end{figure*} The poloidal velocity field should, in result, also be modified. Initially $\\vec{V}_p\\parallel \\vec{B}_p$, as demanded for steady ideal-MHD flow. Therefore, we compute the new poloidal velocity field, maintaining the magnitude of velocity, but correcting the direction so that $\\vec{V}_p\\parallel \\vec{B}_p$ holds.  The modified initial magnetic field is presented in Fig.~\\ref{pocet}, together with the modified density isocontours and positions of the critical surfaces.  For boundary conditions we use symmetry conditions on the rotation axis, and outflow conditions on the outer $R$ and $Z$ boundaries. On the lower boundary $Z=Z_{\\rm min}$ (we take $Z_{\\rm min}=6 R_0$ in all simulations) we fix the values for six physical quantities, namely density, three velocity components, the azimuthal and one poloidal magnetic field component (the other one is given from $\\nabla\\cdot B=0$). The pressure/internal energy boundary condition is kept fixed only in the super-sonic region (where the poloidal velocity is larger than the speed of sound), otherwise it is extrapolated from the flow onto the boundary.  More specifically, in our simulations the $Z$-component of the magnetic field in the first ghost cell is set from the analytical solution, and the radial component is obtained from the divergence-free condition. Therefore, the magnetic flux along the boundary is fixed at all times. In the second ghost cell the $Z$-component of the magnetic field is extrapolated linearly, to allow for change in the field line shape, when the radial component is again obtained from $\\nabla\\cdot B=0$.  In our computations we used various resolutions and sizes of the computational domain. Here we present the results for the resolution $R\\times Z=(256\\times512)$ grid cells $=([0,50]\\times [6,100])R_0$, in the uniform grid. Results comply with the solutions for one fourth, one half and double of this resolution, which we also computed and presented, when needed for direct comparison.  \\subsection{Ideal-MHD integrals} It can be shown, that steady, axisymmetric, ideal-MHD polytropic flows conserve five physical quantities along the poloidal magnetic field lines (Tsinganos 1982). These so called {\\em integrals} are the mass-to-magnetic-flux ratio $\\Psi_A$, the field angular velocity $\\Omega$, the total angular momentum-to-mass flux ratio $L$, the entropy $Q$, and the total energy-to-mass flux ratio $E$ (henceforth we call the latter integral energy for brevity). These integrals are given as \\beq \\Psi_A = \\frac{4 \\pi \\rho V_p}{B_p} \\,, \\eeq \\beq \\Omega = \\frac{V_\\phi}{R} - \\frac{B_\\phi}{B_p} \\frac{V_p}{R}\\,, \\eeq \\beq L = R V_\\phi - \\frac{R B_\\phi B_p}{\\mu_0 \\rho V_p} \\,, \\eeq \\beq Q = p/\\rho^\\gamma \\,, \\eeq \\beq \\label{energy_int} E = \\frac{V^2}{2} + \\frac{\\gamma}{\\gamma-1} \\frac{P}{\\rho} + \\frac{ B_\\phi \\left( B_\\phi V_p - B_p V_\\phi \\right)}{\\mu_0 \\rho V_p} -\\frac{\\cal GM}{r} \\,. \\eeq The various contributions of the energy $E$ correspond to the various terms on the right hand-side of Eq.~(\\ref{energy_int}). From left to right, the kinetic, enthalpy, Poynting, and gravity terms can be recognised.  The degree of alignment of the lines on the poloidal plane where the above quantities are constant together with the poloidal magnetic field lines can be used as a test on how close to a steady-state is the final result of a simulation. ",
        "conclusions": "In this paper we presented resistive numerical simulations of outflows with radially self-similar initial conditions. The analytical solution of V00 has been modified, following GVT06, in a way that it provided a consistent setup for simulations which aimed to extend the failing analytical solution in the close vicinity of the axis, and large distances from the disk. These ideal MHD solutions have been confirmed also by Matsakos et al. (2008), by using the PLUTO code, which is using different numerical methods than the NIRVANA code, used in our simulations.  From the outset, it is not obvious at all that the resistive MHD solutions for such a problem should stay close to the ideal MHD solutions. However, we find that the MHD solution changes smoothly, with a continuous trend for the physical variables, as the resistivity increases. Resistive solutions also reach a well defined stationary state. This in itself is already an interesting and valuable result of the present study. The topology of the moderately resistive MHD solutions turns out to be similar to the ideal MHD case.  This is the case until some critical magnetic diffusivity is reached, when the solutions become increasingly nonconservative in energies and fluxes. The critical transition can be measured through a new dimensionless quantity $\\rbeta$, which measures the influence of resistive effects in the energy equation.  Energy and flux considerations in the present paper are of unprecedented exactness, when it goes for stationarity of the compared runs, which can help to reach some conclusions on resistive MHD behaviour of jets in the astrophysical context. Especially it could be of value for the treatment of the disk corona nearby the disk, where the resistivity of the disk is probably transported up to some height above the disk. The result that the solutions follow the stable {\\em trends}, confirms intuitive expectation.  However, the existence of the critical magnetic diffusivity, now illustrated clearly for the first time in comparison with the analytical solution of the closely related ideal-MHD problem, sets a limit for such intuitive reasoning.  A general conclusion for the resistive MHD simulations of jets is that they are similar to the ideal MHD solutions, for a finite range of the parameter of magnetic diffusivity. In this range, they reach a well defined stationary state. This also extends to the {\\em numerical resistivity} implicitly present in codes. In this respect our result confirms that in numerical simulations with reasonable resolution, the result should not differ significantly from the ideal MHD solution. Departure of the solution from the ideal-MHD regime seems to occur, at least for simple, smooth initial and boundary conditions, only for larger values of the magnetic diffusivity, a few orders of magnitude above the level of numerical magnetic diffusivity. This regime of our solutions will be investigated in more detail in a following study.  Here we presented an application of our results to the case of jets associated with young stellar objects. Evidently, these results are scalable to various other astrophysical cases, as well."
    },
    "0806/0806.1575_arXiv.txt": {
        "abstract": "{ We present the X-ray properties of a massive cluster of galaxies (RXCJ2228.6+2036 at $z=0.421$) using {\\it XMM-Newton} data. The X-ray mass modeling is based on the temperature and density distributions of the intracluster medium derived using a deprojection method. We found that RXCJ2228.6+2036 is a hot cluster ($T_{500}=8.92^{+1.78}_{-1.32}$ keV) showing a cooling flow rate of $12.0^{+56.0}_{-12.0}$ M$_{\\odot}$yr$^{-1}$ based on spectral fitting within the cooling flow radius ($r_{cool}=147\\pm10$ kpc). The total cluster mass is $M_{500}=(1.19\\pm0.35)\\times10^{15}$ M$_{\\odot}$ and the mean gas mass fraction is $f_{gas}=0.165\\pm0.045$ at $r_{500}=1.61\\pm0.16$ Mpc. We discussed the PSF-correction effect on the spectral analysis and found that for the annular width we chose the PSF-corrected temperatures are consistent with those without PSF-correction. We observed a remarkable agreement between X-ray and SZ results, which is of prime importance for the future SZ survey. RXCJ2228.6+2036 obeys the empirical scaling relations found in general massive galaxy clusters (e.g. $S$--$T$, $M$--$T$, $L$--$T$ and $M$--$Y$) after accounting for self-similar evolution. ",
        "introduction": "The gravitational growth of fluctuations in the matter density distribution can be traced by the evolution of the galaxy cluster mass function (e.g. Schuecker et al. 2003). The hot and distant clusters are at the upper end of the mass distribution, thus they can be used to probe the cosmic evolution of large-scale structure and are therefore fundamental probes for cosmology. But to date still very few hot and distant clusters are known. Therefore, it is important to study such clusters in detail, especially in X-ray.  RXCJ2228.6+2036 is one of the distant ({\\it z} = 0.421) and X-ray luminous clusters of galaxies in the northern sky. It is suspected to be massive and hot, and was well recognized as an extended X-ray source in the ROSAT All-Sky Survey, included in both the NORAS galaxy cluster survey (B\\\"{o}hringer et al. 2000) and the ROSAT Brightest Cluster Sample (Ebeling et al. 2000).  The first combined SZ versus X-ray analysis for RXCJ2228.5+2036 is based on the SZ data from the Nobeyama Radio Observatory (NRO) 45 m radio telescope and the X-ray data from ROSAT/HRI. It shows that RXCJ2228.6+2036 is a hot and massive cluster with $T = 10.4 \\pm 1.8$ keV, $M_{tot}(r<R_v=r_{178}=2.9$ Mpc) = $(1.8\\pm0.4)\\times10^{15}$ M$_{\\odot}$, and with a gas mass fraction of $f_{gas} = 0.22 \\pm 0.06$ (Pointecouteau et al.  2002). Recently, LaRoque et al. (2006) performed the {\\it Chandra} X-ray versus OVRO/BIMA interferometric SZ effect measurements for the same cluster, giving $T = 8.43^{+0.78}_{-0.71}$ keV, $f_{gas} = 0.138 \\pm 0.009$ from the X-ray data, and $f_{gas} = 0.188^{+0.035}_{-0.031}$ from the SZ data at $r_{2500}$. RXCJ2228.6+2036, as one of the clusters in the sample of X-ray luminous galaxy clusters with both X-ray ({\\it Chandra}) and SZ observations in Morandi et al. (2007), has a temperature of $T = 6.86^{+0.89}_{-0.71}$ keV and a total mass of $M_{tot}=(4.90\\pm4.35)\\times10^{14}$ M$_{\\odot}$ at $r_{500}=1033 \\pm 464$ kpc. However, the above results are all based on the mass modeling under the assumption of isothermality of the ICM. The {\\it XMM-Newton} EPIC instruments have both high spatial and spectral resolutions and a large field of view, and are therefore suitable for a spatially resolved spectral analysis. We make use of {\\it XMM-Newton} observations to carry on a detailed study of RXCJ2228.6+2036 based on the X-ray mass modeling using a spatially resolved radial temperature distribution and perform a detailed X-ray versus SZ comparison.  The structure of this paper is as follows: Sect. 2 describes the data, background subtraction method and spectral deprojection technique. Sect. 3 presents the spectral measurements using different models to derive the radial temperature profile, cooling time and mass deposition rate. In Sect. 4 we show the radial electron density profile and X-ray mass modeling. In Sect. 5 we discuss the impact of the PSF correction on the spectral analysis, and compare RXCJ2228.6+2036 to the SZ measurements and the empirical scaling relations for massive galaxy clusters. We draw our conclusion in Sect. 6.  Throughout this paper, unless explicitly stated otherwise, we use the 0.5-10 keV energy band in our spectral analysis. The cosmological model used is $H_0$ = $70$ km s$^{-1}${Mpc}$^{-1}$, $\\Omega_m$ = 0.3, and $\\Omega_{\\Lambda}$ = 0.7, in which 1$^{\\prime}$ corresponds to 332.7 kpc at the distance of RXCJ2228.6+2036. ",
        "conclusions": "We presented a detailed analysis of the {\\it XMM-Newton} observations of the distant galaxy cluster RXCJ2228.6+2036 ($z=0.421$) using our deprojection technique. Through the spectral fitting we derived the deprojected temperature profile $T(r)$. Weighted by normalizations, we derived a mean temperature within $r_{500}$, $T_{500}=8.92^{+1.78}_{-1.32}$ keV, which confirms within the error bars the previous results by Pointecouteau et al. (2002) and LaRoque et al. (2006).  Then we calculated the cooling time of this cluster and obtained a cooling radius of $147\\pm10$ kpc. Fitted by a cooling flow model with an isothermal Mekal component, we derived the mass deposition rate $\\sim$ $12.0^{+56.0}_{-12.0}$ M$_{\\odot}$yr$^{-1}$ within $r_{cool}$.  Using the radial density profile $n_e(r)$ and radial temperature profile $T(r)$, we obtained the mass distribution of RXCJ2228.6+2036. At $r_{500}=1.61\\pm0.16$ Mpc, the total mass is $M_{500}=(1.19\\pm0.35)\\times10^{15}$ M$_{\\odot}$, in agreement with the results of Pointecouteau et al. (2002), derived from a combined SZ/X-ray spatial analysis, and the gas mass fraction is $f_{gas}=0.165\\pm0.045$.  We discussed the PSF-correction effect on the spectral analysis and found that the PSF-corrected temperatures are consistent with those without PSF correction.  We found a remarkable agreement within the error bars between our X-ray results and the SZ measurements in Pointecouteau et al. (2002), which is of prime importance for the future SZ survey. The X-ray total mass and X-ray observables for RXCJ2228.6+2036 closely obey the empirical scaling relations found in general massive galaxy clusters, e.g. the $S$--$T$, $M$--$T$, $L$--$T$ and $M$--$Y$ relations, after accounting for self-similar evolution."
    },
    "0806/0806.3093_arXiv.txt": {
        "abstract": "Standard thermodynamical results of ideal Bose gases are used to study the possible formation of a cosmological Bose-Einstein condensate in Scalar Field Dark Matter models; the main hypothesis is that the boson particles were in thermal equilibrium in the early Universe. It is then shown that the only relevant case needs the presence of both particles and anti-particles, and that it corresponds to models in which the bosonic particle is very light. Contrary to common wisdom, the condensate should be a relativistic phenomenon. Some cosmological implications are discussed in turn. ",
        "introduction": "Scalar Field Dark Matter (SFDM) models\\cite{Sin:1992bg,Sahni:1999qe,Matos:2000ng,Hu:2000ke,Arbey:2001qi}, in which the dark matter (DM) particle is a spin-$0$ boson, are becoming a serious alternative to the Cold DM paradigm that invokes the existence of weakly interacting massive particles (WIMP's)\\cite{Taoso:2007qk}.  The WIMP hypothesis has been widely studied in many proposed extensions of the Standard Model of Particle Physics, and there are high expectations for its detection in the near future\\cite{Bertone:2004pz}. Support for the WIMP hypothesis comes from the comparison between high-resolution N-body simulations and the large scale structure we observe in the Universe. The resulting DM distribution seems to be in good agreement with observations\\cite{DelPopolo:2008mr}, from clusters of galaxies down to the scales of the galaxies themselves.  Good agreement means that there are fine details for which the WIMP scenario does not have a definite answer. Those details refer to the internal DM density profile in galaxies, the low number of galaxy satellites, etc., see for instance\\cite{Tasitsiomi:2002hi}. It is not clear if the observed discrepancies are due to corrections to the WIMP hypothesis that have not been taken into account yet, but the possibility exists that the DM particle is something rather different to a WIMP.  It is here where SFDM models enter into play. The usual approximation in cosmological studies so far is to describe SFDM particles by means of a scalar field $\\phi$ minimally coupled to gravity that is endowed with a scalar field potential $V(\\phi)$. The dynamics of any SFDM model depends upon the particular form of the scalar potential, but it is well known that a massive potential as simple as $V(\\phi) = m^2 \\phi^2 /2$, where $m$ is the mass parameter of the SFDM particles, is enough to successfully describe DM in a cosmological setting up to linear order in the formation of large scale structure\\cite{Hwang:1996xd,Matos:2000ss,Matos:2008ag}. The key point here is that the scalar field oscillations around the minimum of the aforementioned quadratic potential make the SFDM particles behaves as cold DM candidates.  More difficult tasks are those of non-linear structure formation, though some progress in that direction has been recently reported\\cite{Woo:2008nn}. Other studies show that the scalar field $\\phi$ can form gravitationally stable structures and that some of them may explain the particular features observed in the rotation curves of galaxies\\cite{Sin:1992bg,Arbey:2001qi,Matos:2007zza}. Nonetheless, the full formation of cosmological structure still is an open question in SFDM models.  The mass $m$ is a free parameter that has to be constrained by cosmological observations. The values considered in the literature cover an ample range of masses, from very heavy values of the order of $10^{13} \\, \\rm{GeV}$ to values as light as $10^{-23} \\, \\rm {eV}$\\cite{Sin:1992bg,Sahni:1999qe,Matos:2000ss,Arbey:2001qi,Liddle:2006qz}. It is usually agreed that very heavy SFDM particles would be indistinguishable from WIMP's in many respects, and the reason is that cosmological bosonic features show up only if the Compton length of the scalar field $\\lambda_C \\equiv m^{-1}$ is of the order of galactic dimensions. This is usually the case in galactic studies, in which very light values of the scalar field mass are ubiquitous.  On the other hand, the existence of SFDM particles, because of its bosonic nature, opens the possibility for the formation of a cosmologically relevant Bose-Einstein condensate (BEC)\\cite{Sin:1992bg,Bernstein:1990kf,Parker:1991jg,Ferrer:2004xj,Matos:2008ag}. Actually, the scalar field $\\phi$ mentioned before is supposed to represent the wave function that describes the evolution of a cosmological BEC, see for instance\\cite{Parker:1991jg}, and the scalar potential $V(\\phi)$ encloses, in general, all the possible self-interaction terms that may be present among the SFDM particles. Therefore, the scalar field description of DM and the formation of a BEC are closely related, but their relation is yet to be fully understood\\cite{Ferrer:2004xj}.  In this paper, I study the conditions for the formation of a cosmological SFDM BEC using known results of the thermodynamics of ideal Bose gases. I shall show that the formation of a BEC is not a generic process in the history of the Universe, but one that only appears for a relativistic SFDM gas composed of very light massive particles.  To set up the general framework, I describe first the general assumptions that will be present throughout the calculations below. First of all, the SFDM particles are supposed to have been in thermal equilibrium with other matter components in the early universe, at a time when they were relativistic. That means the temperature $T_i$ of the early thermal bath was such that $m \\ll T_i \\simeq 10^9 \\, \\rm{GeV}$, as this is the usual initial value assumed for the temperature at the beginning of the so-called Hot Big Bang.  Moreover, the SFDM particles will also decouple from the thermal bath while still relativistic. That is, the temperature at decoupling $T_D$ is much larger than the mass of the particles, $m \\ll T_D < T_i$. Because we expect the Universe to expand adiabatically, the temperature of the SFDM gas will always be very similar to the temperature of the Universe itself.  Another assumption is that SFDM particles may or may not become non-relativistic at some point in the evolution of the Universe. This will allow us to study a ample range of masses for the SFDM particles, including mass values that are well below the present temperature of the Universe, $m < T_0 \\simeq 10^{-4} \\, \\rm{eV}$.  We shall find that, for all practical purposes, the number density of SFDM particles at any given time suffices to decide about the formation of a BEC. Other thermodynamical quantities will be needed at some point, but only to clarify some features about the expansion of the Universe. All calculations below are given in units for which $\\hbar = c = k_B = 1$.  A brief description of the paper is as follows. In Sec.~\\ref{sec:stand-therm-hist} I revise the formation of a BEC using standard formulae for Bose ideal gases; the same is done in Sec.~\\ref{sec:thermal-history-sfdm} but with formulae that include the coexistence of particles and anti-particles. Finally, Sec.~\\ref{sec:conclusions-} is devoted to conclusions. ",
        "conclusions": "} I have described above the main criteria that can lead to the formation of a cosmological BEC within the SFDM hypothesis. The different formulae presented must give a correct description of the SFDM thermodynamics as long as there are no sudden critical changes in the evolution of the Universe that may require a separate study\\cite{Kolb:1990vq}. At the same order of approximation, the effects of the expansion of the Universe are taken into account in the adiabatic behavior of the thermodynamical quantities.  First of all, it is remarkable that the assumption of thermal equilibrium of the SFDM particles imposes tight constraints in the formation of a cosmological BEC. In general, the condensed phase is unreachable because of the adiabatic expansion of the universe.  The existence of a BEC is possible if the presence of anti-particles is explicitly taken into account. The discussion leaded us to Eq.~(\\ref{eq:tphic}), a simple formula that gives the threshold temperature $T_{\\phi,c}$ for the formation of a BEC in terms of basic physical quantities, the mass $m$ of the SFDM particles and their net charge contribution, see $\\eta_\\phi$ in Eq.~(\\ref{eq:etaphi}), to the total matter contents of the universe.  On the other hand, Eq.~(\\ref{eq:tphic}) points out that for massive enough SFDM particles, the cosmological BEC dissapears at late times; the reason lies in the fact that the SFDM gas is diluted by the expansion of the Universe. Actually, we can easily estimate the mass values for which a SFDM BEC exists up to the present time. By taking $T_{\\phi,c} = T_0$ in Eq.~(\\ref{eq:tphic}), we find that a BEC always exists for $m < 10^{-14} \\, \\rm{eV}$. It is exactly this latter case in which a cosmological scalar field $\\phi$ can be considered the right description of SFDM.  For particles with $m > 10^{-14}\\, \\rm{eV}$, the SFDM gas is still relativistic at the temperature at which the BEC disappears, $m \\ll T_{\\phi,c}$, see Eq.~(\\ref{eq:tphic}). After that, the total charge density is entirely accounted by the the excited states, $q_\\phi = \\eta_\\phi T^3_\\phi = \\mu(T_\\phi) T_\\phi /3$, from which we find that the functional form of the chemical potential is the expected adiabatic one, $\\mu(T_\\phi) = 3 \\eta_\\phi T_\\phi$\\cite{Bernstein:1990kf,Haber:1981fg}.  If $m > 1 \\rm{eV}$, the SFDM gas should also become non-relativistic before the present time. But because its number density continues to be that of an ultra-relativistic species, no BEC will be further form at low temperatures. A similar result arises for very large masses so that the SFDM particles decoupled while non-relativistic; the formation of a BEC is not allowed by the adiabatic expansion of the universe\\footnote{Of course, the formation of a BEC should be accomplished, in both cases, in the limit $T=0$ for which the excited states are all depopulated.}.  It is interesting that the existence of a BEC in a cosmological setting requires both high temperatures and relativistic particles. This is contrary to usual expectations according to which the formation of a BEC is considered a low-temperature and non-relativistic phenomenon (see for instance\\cite{Nishiyama:2004ju,Matos:2008ag}).  I should mention here that there is also the possibility that the SFDM particles were produced through a \\emph{non-thermal} process in the early universe (like the typical cases of the infaton and axion fields\\cite{Kolb:1990vq}), so that the appearance of a condensed phase could not be prevented. Such cases, however, are beyond the purposes of this paper.  The overall conclusion of this work is then that the formation of a relativistic BEC is the only possibility for SFDM models if the scalar particles were in thermal equilibrium in the early Universe."
    },
    "0806/0806.3746_arXiv.txt": {
        "abstract": "\\PRE{\\vspace*{.3in}} WIMPless dark matter provides a framework in which dark matter particles with a wide range of masses naturally have the correct thermal relic density.  We show that WIMPless dark matter with mass around 2-10 GeV can explain the annual modulation observed by the DAMA experiment without violating the constraints of other dark matter searches.  This explanation implies distinctive and promising signals for other direct detection experiments, GLAST, and the LHC. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0806/0806.0144_arXiv.txt": {
        "abstract": "Magnetic reconnection plays a critical role in many astrophysical processes where high energy emission is observed, e.g. particle acceleration, relativistic accretion powered outflows, pulsar winds and probably in dissipation of Poynting flux in GRBs. The magnetic field acts as a reservoir of energy and can dissipate its energy to thermal and kinetic energy via the tearing mode instability. We have performed 3d nonlinear MHD simulations of the tearing mode instability in a current sheet. Results from a temporal stability analysis in both the linear regime and weakly nonlinear (Rutherford) regime are compared to the numerical simulations. We observe magnetic island formation, island merging and oscillation once the instability has saturated. The growth in the linear regime is exponential in agreement with linear theory. In the second, Rutherford regime the island width grows linearly with time. We find that thermal energy produced in the current sheet strongly dominates the kinetic energy. Finally preliminary analysis indicates a P(k) ~ 4.8 power law for the power spectral density which suggests that the tearing mode vortices play a role in setting up an energy cascade. ",
        "introduction": "Magnetic reconnection plays a critical role in many astrophysical processes, e.g. particle acceleration \\cite{Phan}, accretion disks \\cite{Hawley} and solar flares \\cite{Shibata}. It is also important in laboratory fusion devices such as tokamaks. Magnetic reconnection is a topological change in the field which violates the frozen-flux condition of ideal magnetohydrodynamics (MHD). If a magnetic field can leak across the plasma it can reach a lower energy state - in the case of a current sheet it can undergo ``tearing'' into filaments or magnetic islands. A current layer of thickness $a$ may dissipate on timescales shorter than the resistive timescale $a^2/\\eta$ due to the tearing mode instability (hereafter TMI). The TMI was first discovered in tokamaks and stellerators and has been extensively studied since the pioneering works \\cite{Furth} and \\cite{Rutherford} (see also \\cite{White}, \\cite{PriestForbes}, \\cite{Biskamp}). In this paper we present numerical simulations of the dissipation of current sheets and the formation of magnetic islands due to the tearing mode instability.  In Section \\ref{Analysis} we remind the reader of the predictions of the linear and weakly nonlinear analyses of Furth et al. \\cite{Furth} and Rutherford \\cite{Rutherford}.  In Section \\ref{Method} we present our numerical method. In Section \\ref{Results} we compare the results with linear and weakly nonlinear theory.  In Section \\ref{Turbulence} we present power spectra derived from 3d simulations.  In Section \\ref{Discussion} we discuss the pitfalls of simulations where reconnection is not properly tracked and the implications for reconnection-driven turbulence. ",
        "conclusions": "\\label{Discussion}    \\subsection{Consequences for driven-turbulence simulations}  Our results show that even without either a driven turbulence mechanism, or an initially turbulent velocity spectrum it is possible for the vortices generated by the TMI to mimic the power spectrum seen in simulations of turbulence. This can have serious consequences for MHD simulations where the resistivity is not explicitly constrained (i.e. numerical resistivity plays a role), since there will be inevitably be some numerical-reconnection driven vortices present in the simulation. However using an explicit resistivity, the contribution from reconnection can be estimated using the formulae in \\cite{Furth,Rutherford}.  Finally, we propose as a benchmark for resistive MHD codes the tearing mode instability test, which can be compared with both analytical results for both the linear and nonlinear regimes, as well as with laboratory experiments. The test may also prove useful as a way of quantifying the effects of numerical resistivity in an MHD code."
    },
    "0806/0806.0372_arXiv.txt": {
        "abstract": "We investigate the implications of our measurement of the Lyman-$\\alpha$ forest opacity at redshifts $2\\leq z\\leq4.2$ from a sample of 86 high-resolution quasar spectra for the evolution of the cosmic ultraviolet luminosity density and its sources. The derived hydrogen photoionization rate $\\Gamma$ is remarkably flat over this redshift range, implying an increasing comoving ionizing emissivity with redshift. Because the quasar luminosity function is strongly peaked near $z\\sim2$, star-forming galaxies likely dominate the ionizing emissivity at $z\\gtrsim3$. Our measurement argues against a star formation rate density declining beyond $z\\sim3$, in contrast with existing state-of-the-art determinations of the cosmic star formation history from direct galaxy counts. Stellar emission from galaxies therefore likely reionized the Universe. ",
        "introduction": "\\label{introduction} The opacity of the Lyman-$\\alpha$ (\\Lya) forest is set by a competition between hydrogen photoionizations and recombinations \\citep[][]{1965ApJ...142.1633G} and can thus serve as a direct probe of the photonization rate \\cite[e.g.][]{1997ApJ...489....7R}. The hydrogen photoionization rate $\\Gamma$~is a particularly valuable quantity as it is an integral over all sources of ultraviolet (UV) radiation in the Universe, \\begin{equation} \\Gamma(z) = 4\\pi \\int_{\\nu_{\\rm HI}}^{\\infty} \\frac{d\\nu}{h \\nu} J_{\\nu}(z) \\sigma(\\nu), \\end{equation} where $J_{\\nu}$ is the angle-averaged specific intensity of the background, $\\sigma(\\nu)$ is the photoionization cross section of hydrogen, and the integral is from the Lyman limit to infinity. As such, it bears a signature of cosmic stellar and quasistellar activity that is not subject to the completeness issues to which direct source counts are prone. Moreover, unlike the redshifted radiation backgrounds observed on Earth, the \\Lya~forest is a \\emph{local} probe of the high-redshift UV radiation, as only sources at approximately the same redshift contribute to $\\Gamma$~at any point in the forest \\citep[e.g.,][]{1996ApJ...461...20H}. In addition to being a powerful probe of galaxy formation and evolution and a fundamental ingredient of cosmological simulations \\citep[e.g.,][]{1992MNRAS.256P..43E}, identifying the sources that contribute most to the UV background is key to our understanding of the reionization history of the Universe.  In this \\emph{Letter}, we derive the photoionization rate implied by our measurement of the \\Lya~forest opacity at $2\\leq z\\leq4.2$ from a sample of 86 high-resolution quasar spectra \\citep[][]{taueffmeas}, for the first time consistently analyzing such a large data set (corrected for both continuum bias and metal absorption) over this redshift interval. We discuss the implications of its flatness over this redshift range for the relative contribution of quasars and star-forming galaxies to the high-redshift cosmic UV background. Throughout, we assume a \\emph{WMAP5} cosmology \\citep[][]{2008arXiv0803.0547K}. The full details of our analysis, as well as supporting arguments, are presented elsewhere \\citep[][]{taueffimp}.  \\begin{figure*}[ht] \\begin{center} \\includegraphics[width=1\\textwidth]{f1.eps} \\end{center} \\caption{(A) Photoionization rate $\\Gamma$ inferred from our \\Lya~effective optical depth measurement (black squares). The contribution from quasars calculated using the \\cite{2007ApJ...654..731H} quasar luminosity function is shown by the short dashes. (B) Comoving UV specific emissivity at 1500~\\AA~obtained by integrating galaxy UV luminosity functions. The green points are from \\cite{2006ApJ...642..653S} (Keck Deep Fields), the cyan points from \\cite{2008ApJS..175...48R} (Keck LBG), the red points from \\cite{2006ApJ...653..988Y} (Subaru Deep Field), the magenta points from \\cite{2007ApJ...670..928B} (Hubble Ultra Deep Field and other deep \\emph{HST} fields), and the blue point is based on the \\cite{1999ApJ...519....1S} $z\\sim4$ LBGs, with the characteristic magnitude and faint-end slope set to the $z\\sim3$ values of \\cite{2008ApJS..175...48R}. See the text for caveats about the error bars shown. The black points show the emissivity implied by our \\Lya~forest measurement, where the normalization was set to match the LF-derived values. (C) Comoving star formation rate density implied by the UV emissivity (see text for details). Same color scheme as in (B). The long dashed curve shows the best fit of \\cite{2006ApJ...651..142H} to the star formation history.} \\label{gamma etc} \\end{figure*} ",
        "conclusions": ""
    },
    "0806/0806.0234_arXiv.txt": {
        "abstract": "In this paper we describe a focal plane array (FPA) prototype, based on Vivaldi elements, developed for the Westerbork Synthesis Radio Telescope (WSRT) to increase its instantaneous field of view by a factor 25 and double its current bandwidth. This prototype is the first step in a project that has the ambition to equip most of the WSRT antennas with FPAs to improve the survey speed of the telescope. Examples of scientific applications are surveys of the northern sky in polarised continuum and HI emission, and efficient searches for pulsars and transients. ",
        "introduction": "Present day synthesis radio telescopes have limited survey capabilities because of field of view restrictions. On single dish radio telescopes this restriction has been alleviated by constructing multiple feeds which provide several beams on the sky, such as the succesful Parkes multi-beam system which surveyed the entire southern sky in HI (Staveley-Smith et al. 1995, 1996), and the Arecibo multi-beam system AlfAlfa (Giovanelli et al. 2005). An alternative method to form multiple beams on the sky is to employ phased array technology in the focal plane of a radio telescope. The advantage of using this technique is that there is great flexibility in forming beams on the sky. Especially in telescopes with small f/D ratios, it is the only way to form beams on the sky which are closely packed, touching at their half power points. In addition it is much easier to optimise the illumination of the telescope dish for maximum gain or minimum sidelobes.  APERTIF (``APERture Tile In Focus'') is such a system that is being developed for the Westerbork Synthesis Radio Telescope (WSRT). In this paper we will briefly describe the technical details of the focal plane array system under development, report first results with a prototype and give a brief overview of a few scientific projects which will come within reach when the WSRT is equipped with FPA feeds. ",
        "conclusions": ""
    },
    "0806/0806.3237_arXiv.txt": {
        "abstract": "We present the third installment of HI sources extracted from the Arecibo Legacy Fast ALFA extragalactic survey.  This dataset continues the work of the Virgo ALFALFA catalog.  The catalogs and spectra published here consist of data obtained during the 2005 and 2006 observing sessions of the survey.  The catalog consists of \\ncat~HI detections within the range $11^h36^m <$ R.A.(J2000) $< 13^h52^m$ and $+08^{\\circ} <$ Dec.(J2000) $< +12^{\\circ}$, and $cz_{\\odot} < 18000$ \\kms. The catalog entries are identified with optical counterparts where possible through the examination of digitized optical images.  The catalog detections can be classified into three categories: (a) detections of high reliability with S/N $>$ 6.5; (b) high velocity clouds in the Milky Way or its periphery; and (c) signals of lower S/N which coincide spatially with an optical object and known redshift. 75\\% of the sources are newly published HI detections.  Of particular note is a complex of HI clouds projected between M87 and M49 that do not coincide with any optical counterparts.  Candidate objects without optical counterparts are few. The median redshift for this sample is 6500 \\kms~ and the $cz$ distribution exhibits the local large scale structure consisting of Virgo and the background void and the A1367-Coma supercluster regime at $cz_{\\odot} \\sim$7000 \\kms.  Position corrections for telescope pointing errors are applied to the dataset by comparing ALFALFA continuum centroid with those cataloged in the NRAO VLA Sky Survey.  The uncorrected positional accuracy averages 27\\arcsec ~(21\\arcsec ~median) for all sources with S/N $>$ 6.5 and is of order $\\sim$21\\arcsec ~(16\\arcsec ~median) for signals with S/N $>$ 12. Uncertainties in distances toward the Virgo cluster can affect the calculated HI mass distribution. ",
        "introduction": "\\label{obs}  ALFALFA utilizes a fixed azimuth drift mode observation scheme with the 7-element ALFA multi-beam receiver system.  Each dual polarization feed has a beam that is $3.3$\\arcmin~$\\times$~ $3.8$\\arcmin~in size.  Beam maps of each feed can be found in Paper I.  The ALFA array is rotated so that parallel feed tracks of constant J2000 declination are spaced at 1.05\\arcmin.  Scans are obtained with the telescope parked along the meridian--small corrections in zenith angle are applied between scans such that the beams are tracking along the constant epoch J2000 Declination.  Fourteen simultaneous spectra are obtained at a sampling rate of 1 Hz in R.A. and scans are composed of 600 one--second records each sweeping 10 minutes of Right Ascension.  The backend correlator setup uses a 100 MHz bandwidth centered at 1385 MHz.  Raw scans have 4096 channels, giving a spectral resolution of 24.4 kHz ($\\delta v\\sim 5.2$~\\kms~ at $cz\\sim$0). Other survey details about the observing and technical modes can be found in Paper I.  Raw data scans are processed offline for each observing session. Initial processing by team members includes flagging of radio interference, bandpass and flux calibration, data quality assessment, and identifying strong continuum sources.  Observing runs are scheduled to fill in target area tiles that are composed of $2.4^{\\circ} \\times 2.4^{\\circ}$ areas of sky. Once scheduled target area tiles are completed, the level I data scans are combined into regularly gridded 3-D data cubes, each 2.4 $\\times$ 2.4 degrees in spatial extent, and 4~$\\times$~1024 channels in velocity space. These data cubes contain spectral data header information, coordinates, both polarizations, continuum data, map weightings, and scan makeup. An automated matched filter algorithm is used for signal extraction (Saintonge 2007). Candidate signals are later examined by eye, and integrated spectral profiles are created.  Flux measurements, velocities, and widths are cataloged for publication and noted as possible followup candidates. The resulting data cubes are approximately 4~$\\times$~380 MB in size.  \\subsection{Continuum maps}  In addition to the spectral line data, continuum maps are also created using data from the level I data scans, described in Giovanelli \\etal (2007). Channels are flagged through the bandpass calibration process and/or the manual flagging process mentioned previously. A background total power continuum value is also computed for all time series records for all records and channels that have not been flagged, excluding point sources; the continuum contribution from these point sources is also stored for the creation of the continuum maps. Sources are detected with an automated peak finding algorithm, which then searches a database of NVSS (Condon \\etal~1998) sources and looks for matches, comparing both fluxes and positions.  Final positions are fit with a 2D Gaussian and stored along with peak flux measurements. These measurements are used to correct for positions offsets as described in section 5, and continuum source variability studies.  \\subsection{Data Access}  As the ALFALFA collaboration includes many worldwide members, rapid dissemination of the data products is of the utmost importance to the survey's success. The catalog results presented in this paper will be added to the existing ALFALFA archive at {\\it http://arecibo.tc.cornell.edu/hiarchive/alfalfa}/. The site provides web services using protocols from the  U.S. National Virtual Observatory\\footnote{This research has made use of data obtained from or software provided by the US National Virtual Observatory(NVO), which is sponsored by the National Science Foundation.}.  The measurements and spectra will join the fast growing archive of HI measurements (Springob \\etal~2005; Giovanelli \\etal~2007; Saintonge \\etal~2008). An ongoing development effort focuses on the long--term public delivery of the 3--D ALFALFA data set through web--based access tools. At this time, delivery of the 3-D data is made possible through the observing team itself, by direct contact to R.G. or M.P.H. A major challenge is data volume: after regridding the 3-D cubes covering the 34 individual ``grids'' from which the current catalog each occupy 50 GB. Allowing access and manipulation of the gridded data publicly will require the development of web services and server applications.  In this work, we present a catalog of HI sources extracted from the ALFALFA grids covering a region stretching from $+8^\\circ$ to $12^\\circ$ in Dec. and from $11^h36^m$ to $13^h52^m$ in R.A. For reference to our database, denomination of the grids from which sources in this catalog were extracted are 1140+09 to 1348+09 and 1140+11 to 1348+11, in steps of 8$^m$ in R.A. The solid angle subtended by this region is $\\sim$132 \\sqd, which is $\\sim1.9$\\% of the sky to be ultimately surveyed by ALFALFA. The coverage of the region is complete by the target goals of ALFALFA, i.e. the region has been sampled by two separate passes with the ALFA array in drift mode. ",
        "conclusions": ""
    },
    "0806/0806.1232_arXiv.txt": {
        "abstract": "We show that using either the method of Page~\\& Carrera or the well-known $1/V_{\\mathrm{a}}$ method to construct the binned luminosity function (LF) of a flux limited sample of Active Galactic Nuclei (AGN) can produce an artificial flattening (or steepening in the case of negative evolution) of the binned LF for bins intersected by the flux cutoff of the sample. This effect is more pronounced for samples with steep and strongly evolving parent LFs but is still present even for non-evolving LFs. As a result of this distortion of the true LF, fitting a model LF to binned data may lead to errors in the estimation of the parameters and may even prompt the erroneous use of broken power law functions. We compute the expected positions of apparent breaks in the binned LF. We show that these spurious breaks in the binned LFs can be avoided if the binning is done in the flux--redshift plane instead of the typically used luminosity--redshift plane. Binning in the flux--redshift plane can be used in conjunction with the binning in the luminosity--redshift plane to test for real breaks in the binned LFs and to identify the features that are the result of binning biases. We illustrate this effect for most typical forms of luminosity dependence and redshift evolution and show how the proposed method helps address this problem. We also apply this method to the MOJAVE AGN sample and show that it eliminates an apparent break in the binned LF. ",
        "introduction": "The luminosity function (LF) of active galactic nuclei (AGN) and its redshift dependence is highly useful for studying the cosmological evolution of AGNs and for statistical tests of AGN unification models. Typically the binned LF is obtained using the standard $1/V_{\\max}$ method \\citep[][]{Schmidt68,Felten76} or its generalization for combined samples $1/V_{\\mathrm{a}}$ \\citep[][]{Avni80}. The evolution of the LF is estimated using the $\\left\\langle V/V_{\\max} \\right\\rangle$ method \\citep[][]{Schmidt68}. Sometimes a model LF is derived by fitting a model LF to the binned LF \\citep[see, e.g.,][]{Padovani07}. However, the binned LF obtained using the $1/V_{\\mathrm{a}}$ method suffers from several biases \\citep[see][and references therein]{LaFranca97,Wisotzki98} and for this reason many prefer fitting the model LF to the unbinned data with, e.g., maximum likelihood techniques. In this case the  binned LF, however, is often used before performing the fit to observe the overall behavior of the luminosity function, e.g., to decide its evolution type and to decide whether or not a double power law LF should be fitted to the data instead of a single power law. It can also be used after the fit to visualize the goodness of fit.  \\citet{PC00} described a different method of computing the binned LF that, according to the authors, gives better estimates of the LF for lower luminosity bins that are close to the flux limit of the survey using an improved (over $1/V_{\\mathrm{a}}$) estimator. However, their method does not eliminate all the biases associated with undersampled bins that are near the flux cutoff of the survey \\citep[see, e.g.,][]{Miyaji01}. In a recent paper \\citep{Cara08} we derived the radio luminosity function of the MOJAVE survey of core-selected AGN \\citep{LH05}. While a single power law provided a good fit to the unbinned data, the binned LF, computed using the method of \\citet{PC00}, was flatter at lower luminosities, and suggested a double power law LF \\citep[see, e.g., ][]{Arsh06}. After an analysis of the problem \\citep[see][]{Cara08} we concluded that the flattening of the binned LF was due to the fact that the values of the LF computed over small chunks of undersampled (because of the flux cutoff) bins are not good approximations of the value of the LF at the centers of the bins, particularly when the LF is a steep and strongly evolving function across the bin.  There are many other examples in the literature where binned luminosity functions have been extended well below the flux cutoff of the survey. For example, \\citet[][Fig. 10]{Dunlop90} used the binned LF to compare the prediction of their pure luminosity evolution (PLE) models with the data. The plot displays a clear flattening of the LF for bins below the flux cutoff. Because of the multiple flux limits in their data, this flattening happens gradually and, therefore, the binned LFs in that plot are better described as being ``bent'' than ``broken''.  A somewhat different example of using bins below the flux cutoff comes from \\citet{Urry95}. In Figure 14 the authors present the \\emph{local} LFs of the FR~II galaxies, steep spectrum radio quasars (SSRQ) and flat spectrum radio quasars (FSRQ). In this case, because the luminosities have been de-evolved, the LFs have been constructed essentially using one large redshift bin with a modified (de-evolved) flux limit. Therefore, almost all the luminosity bins presented in Figure 14 of \\citet{Urry95} lie below the survey limit (however, the precise number of bins below the limit is difficult to determine because the redshift bin width is not presented). At very low luminosities, particularly at luminosities where the survey limit starts cutting into the number of objects in a bin, the number densities will start decreasing because the incomplete bins, due to the shape of the flux cutoff, will preferentially sample a different part of the LF than would be the case for a complete bin. We will show in Section~\\ref{sec:Method} that the position of the dominant break will depend on the minimum redshift of the bin in the case of a non-evolving luminosity function.  A final example is that of the binned luminosity function of the quarter-Jansky sample from \\citet{Wall05}. The flattening of the LF of this sample at low luminosities produces the illusion of the luminosity evolution, which was recognized by the authors as such. The authors attributed the flattening of the LF at lower luminosities to an intrinsic spread in radio spectral indices.  Since this flattening (or steepening in the case of negative evolution) is caused by undersampling of the bins cut by the survey limit, one solution to eliminate these spurious breaks is to stop binning as soon as the lowest luminosity bin reaches the survey limit. This, however, can lead to the exclusion of data points lying between this last complete bin and the survey limit. \\citet{LaFranca97,Miyaji01} proposed to correct for biases in the undersampled bins by prorating them for the missing parts using a model LF determined by independent means. This approach, however, has the disadvantage of depending on an assumed model. \\citet{Efstathiou88} have developed a non-parametric ``step-wise maximum likelihood'' method that allows a binned LF to be fit directly to the unbinned data. While these methods their own advantages and disadvantages, we have chosen to focus in this paper on the method of Page~\\& Carrera. We note that any method that computes the values of the LF over incomplete bins must, in one way or another, use some sort of extrapolation to the parts of the bins lying below the flux limit.  In this paper we show that the biases associated with undersampled bins can produce artificial breaks in the derived luminosity functions, particularly in the case of strong evolution. As evolution of the LF is typically judged by the shifts of features in the LF, these artificial breaks, when interpreted as real, can produce the illusion of luminosity evolution. We show that the flattening (or breaking) of the binned LF for bins close to the flux cutoff of the survey can be minimized if the binning is performed in the flux--redshift plane instead of the luminosity--redshift plane as it is usually done. The improvement is most apparent at lower redshifts where, due to the steepness of the luminosity cutoff in the $\\log L-z$ space, many bins are undersampled, and therefore give strongly biased estimates of the luminosity function. While our analysis of the distortion of the shape of the binned LFs is performed using the method of \\citet{PC00} for constructing binned luminosity functions, this analysis equally applies to the $1/V_{\\mathrm{a}}$ method as well. The binned LFs constructed using the $1/V_{\\mathrm{a}}$ method will exhibit larger deviations from the true LF than the method of \\citet{PC00} at extremely low luminosities ($L<L_{\\mathrm{break,3}}$, see equation~\\ref{eqn:PosBreak3}). However, in the case of non-evolving LFs, the $1/V_{\\mathrm{a}}$ method, unlike the method of Page~\\& Carrera, will not distort the shape of the LF for $L_{\\mathrm{break,3}}<L<L_{\\mathrm{break,1}}$ because it assigns a variable weight to each object in the bin and can correctly handle the flux limit down to lower luminosities~($L_{\\mathrm{break,3}}$). Regardless of the method used, any determination of the LF below $L_{\\mathrm{break,3}}$ is pure extrapolation and should be avoided.  The outline of the paper is as follows: In \\S~\\ref{sec:Method} we describe both binning methods and discuss the main advantages of the binning in the flux--redshift plane over the traditional luminosity--redshift plane binning method. In \\S~\\ref{sec:MCsim} we apply these two binning methods to several model luminosity functions and show that the LF binned in the flux--redshift plane is a better approximation to the model LF. We construct the binned LF of the MOJAVE sample using both methods in \\S~\\ref{sec:MOJAVELF}. We summarize our findings and present our conclusions in \\S~\\ref{sec:conclusions}.  Throughout this paper we assume a cosmology with $\\Omega_m=0.3$, $\\Omega_{\\Lambda}=0.7$, $\\Omega_{r}=0$ and $H_0=70\\,\\mbox{km}\\,\\mbox{s}^{-1}\\,\\mbox{Mpc}^{-1}$. All luminosities are quoted as monochromatic luminosities at specific frequency $\\nu$.  We also adopt the following convention for the spectral index, $\\alpharad$: $S_{\\nu}\\propto\\nu^{\\alpharad}$ and assume that $\\alpharad=0$ for all model LFs and plots in this paper. ",
        "conclusions": "\\label{sec:conclusions}  The binned LFs constructed in the traditional $\\log L-z$ plane show significant distortions in bins that are undersampled due to luminosity limits of the survey. In the presence of strong positive evolution across these bins, the derived luminosity function will flatten for luminosities below the break luminosities given by equations (\\ref{eqn:PosBreak1}), (\\ref{eqn:PosBreak2}) and (\\ref{eqn:PosBreak3}), producing artificial breaks in the binned LFs. In the case of negative evolution the effect is opposite: the luminosity functions steepen. For single power law LFs such breaks can create the illusion of luminosity evolution. The space density at which these spurious breaks occur varies with the redshift bin.  In order to avoid the possibility of producing spurious breaks in binned LFs, the binning in the $\\log L-z$ plane should not extend below luminosities given by $L_{\\mathrm{break,1}}$ from equation~(\\ref{eqn:PosBreak1}) and binning below $L_{\\mathrm{break,3}}$ should be avoided at all costs. However, this may mean that some sources are excluded from the analysis. Alternatively, we propose that the binned LFs be constructed in the $\\log S-z$ plane using complete bins only. This method has the following advantages:  \\begin{enumerate} \\item Sampling in the $\\log S-z$ plane can be done using complete bins without the loss of information (i.e., exclusion of sources with $L<L_{\\mathrm{break,1}}$) that would happen if only complete bins were used in the $\\log L-z$ plane.  \\item Using complete bins eliminates the appearance of spurious breaks or steepening of the LFs.  \\item In the same redshift interval the LFs binned in the $\\log S-z$ plane can sample the parent LF down to lower luminosities than in the case of $\\log L-z$ binning as long as $D^2_{\\mathrm{L}}(z)(1+z)^{-(1+\\alpharad)}$ is an increasing function of redshift (see equation~\\ref{eqn:LcLcDif}).  \\item The previous advantage comes at the expense of blurring the real features in the binned LF.   \\item Binning in the $\\log S-z$ plane should be used in conjunction with binning in the $\\log L-z$ plane (for $L>L_{\\mathrm{break,1}}$) to test for real breaks in the binned LFs and to identify the features that are the result of binning biases, since the biases associated with binning in the two planes are expected to be different. This is especially important because the binning in $\\log S-z$ may blur some weak real features in the binned LF.  \\end{enumerate}"
    },
    "0806/0806.3147_arXiv.txt": {
        "abstract": "s{EDELWEISS is a direct dark matter search situated in the low radioactivity environment of the Modane Underground Laboratory. The experiment uses Ge detectors at very low temperature in order to identify eventual rare nuclear recoils induced by elastic scattering of WIMPs from our Galactic halo. The commissioning of the second phase of the experiment, involving more than 7 kg of Ge, has been completed in 2007. Two new type of detectors with active rejection of events due to surface contamination have demonstrated the performances required to achieve the physics goal of the present phase.} ",
        "introduction": "First indications for the existence of Dark Matter were already found in the 1930s~\\cite{zw}. By now there is strong evidence~\\cite{gbdhjs} to believe that a large fraction (more than 80\\%) of all matter in the Universe is Dark (interacts very weakly with electromagnetic radiation, if at all) and that this Dark Matter is predominantly non-baryonic. Weakly Interacting Massive Particles (WIMPs) are one of the leading candidates for Dark Matter. WIMPs are stable particles which arise in several extensions of the Standard Model of electroweak interactions~\\cite{hhnos}. Typically they are presumed to have masses between few tens and few hundreds of GeV/c$^{2}$ and a scattering cross  section with a nucleon below 10$^{-6}$ pb.  The EDELWEISS experiment (Exp{\\'e}rience pour D{\\'e}tecter les WIMPs en Site Souterrain) is dedicated to the direct detection of WIMPs. The direct detection principle consists in the measurement  of the energy released by nuclear recoils produced in an ordinary matter target by the elastic collision of a WIMP from the Galactic halo. The main challenge is the expected extremely low event rate ($\\leq$ 1evt/kg/year) due to the very small interaction cross section of WIMP with the ordinary matter. An other constraint is the relatively small deposited energy ($\\leq$ 100keV). ",
        "conclusions": "The EDELWEISS-II setup has been validated with calibration and background runs. Energy resolution and discrimination capabilities close to those of EDELWEISS-I have been measured for Ge/NTD detectors. The validation of Ge/NbSi detectors with new aluminum electrodes is in progress and Ge/NTD detectors with interdigitized electrode scheme have shown promising results in surface laboratory.  Low background physics runs will be taken with 28 detectors setup with the aim to reach sensitivity to WIMP-nucleon cross-section of 10$^{-7}$~pb for a WIMP mass of 100~GeV in 2008. Additional detectors will be added in the two coming years to enhance progressively the sensitivity to few 10$^{-9}$~pb thanks to active surface rejection."
    },
    "0806/0806.0622_arXiv.txt": {
        "abstract": "We have studied the dissolution of initially mass segregated and unsegregated star clusters due to two-body relaxation in external tidal fields, using Aarseth's collisional $N$-body code NBODY4 on GRAPE6 special-purpose computers. When extrapolating results of initially not mass segregated models to globular clusters, we obtain a correlation between the time until destruction and the slope of the mass function, in the sense that globular clusters which are closer to dissolution are more strongly depleted in low-mass stars. This correlation fits observed mass functions of most globular clusters. The mass functions of several globular clusters are however more strongly depleted in low-mass stars than suggested by these models. Such strongly depleted mass functions can be explained if globular clusters started initially mass segregated. Primordial mass segregation also explains the correlation between the slope of the stellar mass function and the cluster concentration which was recently discovered by De Marchi et al.\\ (2007). In this case, it is possible that all globular clusters started with a mass function similar to that seen in young open clusters in the present-day universe, at least for stars below $m=0.8$ M$_\\odot$. This argues for a near universality of the mass function for different star formation environments and metallicities in the range $-2 < \\mbox{[Fe/H]} < 0$.  We finally describe a novel algorithm which can initialise stationary mass segregated clusters with arbitrary density profile and amount of mass segregation. ",
        "introduction": "In recent years, stellar mass functions have been obtained for an increasing number of globular clusters by deep {\\it HST} and {\\it VLT} measurements (see DeMarchi et al.\\ 2007 and references therein). These observations have shown that there is a considerable spread in the present-day mass functions of individual clusters, and that a number of star clusters are strongly depleted in low-mass stars. If one expresses the mass function of a cluster as a power-law\\footnote{Note that \\citet{dpp07} used $dN/dm \\sim m^{\\alpha}$ in their paper.} by $dN/dm \\sim m^{-\\alpha}$, where $N$ is the number of stars per unit mass $m$, the observed slopes range from between $\\alpha=1.9$ to $\\alpha=-0.9$ for stars with masses in the range $0.3 < m < 0.8$ M$_\\odot$. For clusters where information from different radii is available, the data point to a global decrease of the number of low-mass stars in the clusters, rather than a local effect due to mass segregation.  The depletion of low-mass stars can in principle be understood by mass segregation and the preferential loss of low-mass stars as a result of the dynamical evolution of star clusters. Indeed, using direct $N$-body simulations, \\citet{bm03} found a correlation between the observed and expected slopes for the then available sample of star clusters. However, for a number of clusters, the difference between theoretical and expected slope is far too large to be explained just by observational uncertainties.  This is emphasised by \\citet{dpp07}, who found a correlation between the mass function slope $\\alpha$ and the concentration parameter $c=\\log_{10}(r_t/r_c)$ for globular clusters, where $r_t$ and $r_c$ are the tidal and core radius of the cluster as determined from the projected light density profile. The correlation found by \\citet{dpp07} is in the sense that clusters with small values of $c$ are depleted in low-mass stars, while clusters with large values of $c$ have mass functions still rising towards small masses. Since simulations show that mass segregation and the preferential loss of low-mass stars should only happen after a cluster has gone into core-collapse, and since core-collapse is connected to the shrinkage of the core size $r_c$, the observed correlation is the exact opposite of the theoretically expected one.  One possible interpretation of this finding would be that star clusters that formed more concentrated have a bottom-heavy IMF, which would be a challenge to star formation theories and dispose the universality of the IMF. However this conclusion needs to be tested by taking into account the stellar-dynamical evolution of the clusters.  In the present paper we compare the observational results with theoretical predictions by \\citet{bm03} (BM03), who have performed a large parameter study of initially not mass segregated multi-mass clusters evolving under the combined influence of relaxation, stellar evolution and an external tidal field. We also report results of new simulations of multi-mass star clusters which start initially mass segregated. Initial mass segregation is expected to occur in star clusters as a result of star formation feedback in dense gas clouds \\citep{mur96} and due to competitive gas accretion and mutual mergers between protostars \\citep{bon02}. Numerous studies have also found observational evidence for it in young star clusters of the Milky Way and the Magellanic Clouds \\citep{bon98,getal04,cdz07}, so that it is certainly possible that globular clusters started mass segregated.  The paper is organised as follows: In \\S2 we compare the results of simulations of non-mass-segregated clusters done by \\citet{bm03} with the observed mass function slopes of globular clusters. In \\S3 we describe the numerical simulations of star clusters with primordial mass segregation and in \\S4 we compare the results of these runs with the observations. We briefly summarize in \\S5. ",
        "conclusions": ""
    },
    "0806/0806.3767_arXiv.txt": {
        "abstract": "We determine the density and mass distribution of dark matter within our Solar System.  We explore the three-body interactions between dark matter particles, the Sun, and the planets to compute the amount of dark matter gravitationally captured over the lifetime of the Solar System.  We provide an analytical framework for performing these calculations and detail our numerical simulations accordingly.  We find that the local density of dark matter is enhanced by between three and five orders of magnitude over the background halo density, dependent on the radial distance from the Sun. This has profound implications for terrestrial direct dark matter detection searches.  We also discuss our results in the context of gravitational signatures, including existing constraints, and find that dark matter captured in this fashion is not responsible for the Pioneer anomaly. We conclude that dark matter appears to, overall, play a much more important role in our Solar System than previously thought. ",
        "introduction": "Introduction}  Accounting for about $23\\%$ of the energy density of the Universe \\cite{Komatsu:2008hk}, dark matter is an integral part of our surroundings. It dominates the more familiar, baryonic components, which comprises only  $4.4\\%$ of the Universe, on the largest scales. Achieving an understanding of this perplexing dark element is one of the most compelling unsolved problems in modern astrophysics.  The astrophysical evidence for the existence of dark matter is overwhelming, as observations of the cosmic microwave background \\cite{Komatsu:2008hk}, the power spectrum of the Universe \\cite{2dFSDSS}, and colliding galaxy clusters \\cite{Clowe:2006} all point towards the same conclusion.  With an understanding that a dark, pressureless, fluid-like component permeates the Universe, astrophysicists have successfully simulated the large-scale processes of structure formation in the context of a $\\Lambda$CDM Universe \\cite{Monstrosity}.  More recently, some attention has been given to the dark matter substructure formed on subgalactic scales, down to scales of order $10^{-2} \\,$pc \\cite{lilCDM}.  However, relatively little consideration has been given to the distribution of dark matter within our own Solar System.  Yet, dark matter may prove to be profoundly important in our Solar System for both its additional gravitational effects on planets and other orbiting bodies \\cite{Monstrosity2,Gron:1996,Frere:2008} as well as the motions of spacecraft \\cite{Anderson:1998,Adler:2008}.  Furthermore, a knowledge of the density and velocity of dark matter particles is particularly important for terrestrial direct detection experiments \\cite{Aprile:2005ww}.  In this paper, we model the Solar System and the dark matter that it encounters in order to quantify how much dark matter we expect the Solar System to have captured over its lifetime.  Through favorable three-body gravitational interactions between a dark matter particle, the Sun, and any of the planets, a non-zero and possibly significant fraction of the dark matter passing through the Solar System will become gravitationally bound to it.  The remainder of this paper is focused on solving this problem, and is laid out as follows: section 2 details the model chosen for the Solar System and galactic properties, and provides analytic details of our calculations.  Section 3 sets forth the computations undertaken to successfully determine the probability of binding dark matter to the Solar System.  Section 4 presents our results, including the density and mass distributions of dark matter with respect to distance from the Sun. We recommend that anyone not interested in the details of our calculations skip directly to section 4.  Finally, section 5 concludes with a discussion of the conclusions reached from our analysis, detailing significant implications for dark matter detection and presenting a comparison of our results with the current experimental and observational constraints. ",
        "conclusions": ""
    },
    "0806/0806.1886_arXiv.txt": {
        "abstract": "This is the third paper in a series analyzing X-ray emission from the hot interstellar medium in a sample of 54 normal elliptical galaxies observed by Chandra. We focus on a subset of 36 galaxies with sufficient signal to compute radial temperature profiles.  We distinguish four qualitatively different types of profile: positive gradient (outwardly rising), negative gradients (falling), quasi-isothermal (flat) and hybrid (falling at small radii and rising at larger radii). We measure the mean logarithmic temperature gradients in two radial regions: from 0--2 $J$-band effective radii $R_J$ (excluding the central point source), and from 2--$4R_J$. We find the outer gradient to be uncorrelated with intrinsic host galaxy properties, but strongly influenced by the environment: galaxies in low-density environments tend to show negative outer gradients, while those in high-density environments show positive outer gradients, suggesting the influence of circumgalactic hot gas. The inner temperature gradient, however, is unaffected by the environment but strongly correlated with intrinsic host galaxy characteristics: negative inner gradients are more common for smaller, optically faint, low radio-luminosity galaxies, whereas positive gradients are found in bright galaxies with stronger radio sources. There is no evidence for bimodality in the distribution of inner or outer gradients. We propose three scenarios to explain the inner temperature gradients: (1) Weak AGN heat the ISM locally, while higher-luminosity AGN heat the system globally through jets inflating cavities at larger radii; (2) The onset of negative inner gradients indicates a declining importance of AGN heating relative to other sources, such as compressional heating or supernovae; (3) The variety of temperature profiles are snapshots of different stages of a time-dependent flow, cyclically reversing the temperature gradient over time. ",
        "introduction": "In the first two papers of this series \\citep[][hereafter Paper I and Paper II]{DiehlGallery,DiehlAGN}, we conducted a comprehensive morphological analysis of the hot gas in normal elliptical galaxies. In Paper I, we introduced a technique to separate the hot gas emission from the contamination of unresolved point sources. We applied this technique to a {\\it Chandra} archive sample of 54 elliptical galaxies and presented a gallery of adaptively binned gas-only images, which were photon-flux calibrated and background corrected. We used these gas maps to derive isophotal ellipticity profiles and conducted a systematic morphological analysis. A comparison between optical and X-ray ellipticities measured in the inner, stellar mass dominated regions shows no correlation, contrary to what would be expected if the gas were in perfect hydrostatic equilibrium. We modeled the expected correlation under various assumptions, and concluded that these systems, in general, are at best only approximately hydrostatic. Moreover, the gas morphologies almost always look disturbed. In Paper II, we introduced a quantitative measure of morphological asymmetry, and found it to be tightly correlated with radio continuum power and nuclear X-ray luminosity. We also found the gas morphology to be influenced, to a comparable degree, by the ambient intergalactic medium. But the AGN--morphology correlation forms a continuous trend down to the lowest detectable AGN luminosities, indicating the importance of AGN feedback, even in rather X-ray faint elliptical galaxies.  In this third paper, we address the question of whether the central AGN is merely redistributing the gas, or heating it as well. We produce radial temperature profiles and find that they fall into a variety of distinct types. In particular, we confirm that negative (outwardly falling) temperature gradients \\citep{HumphreyDarkmatter,FukazawaMassprofiles, RandallXMMNGC4649, PonmanNGC6482} are present, and relatively common, in low-luminosity systems.  Outwardly {\\em rising\\/} (positive-gradient) temperature profiles, nearly ubiquitous in galaxy clusters, X-ray groups, and massive ellipticals, are usually understood as being the result of efficient radiative cooling in the dense central regions. Accreting gas at large radii can additionally shock-heat itself, amplifying the positive gradient. This interpretation is supported by the short central cooling times observed for galaxies, groups, and clusters, which can drop well below $100\\,{\\rm Myr}$. However, cooling times are equally short in galaxies with negative temperature gradients, i.e. with a warm center \\citep{HumphreyDarkmatter}. Several solutions have been proposed to explain these counter-intuitive objects.  \\citet{FukazawaMassprofiles} suggest that the gradients are a function of environment, with outwardly rising (positive) gradients caused by the hotter ambient intracluster or intragroup gas surrounding the galaxy. Galaxies with negative gradients should instead be isolated.  \\citet{HumphreyDarkmatter}, on the other hand, propose a bimodal distribution of temperature gradients, and suggest that the total mass of the system is the decisive factor for the sign of the gradient. The division between their two groups happens at a virial mass of $\\sim 10^{13}M_\\sun$, implying a distinction between normal galaxies and groups. They hypothesize further that the temperature gradients could be related to a significant difference in the galaxies' evolutionary histories.  \\citet{HumphreyDarkmatter} also discuss the role of compressive (gravitational) heating, noting that during a slow inflow of relatively cool gas ($<1-2\\kev$) the energy gain would exceed the radiative losses. For the inflow of hotter baryons, radiative cooling would dominate and one would observe a positive temperature gradient instead. However, they find no reason for hot baryons to be specific to systems above their break mass $\\sim 10^{13}M_\\sun$, and suspect the environment as a fuel source instead.  \\citet{PonmanNGC6482} observe a negative temperature gradient in the fossil group candidate NGC~6482 and argue along the same lines. They model NGC~6482's temperature profile successfully with a steady-state cooling flow with a reasonable cooling rate of $\\dot M = 2M_\\sun\\,\\yr^{-1}$, and adopt it as their preferred solution. They also estimate that type Ia supernovae (SN) may be responsible for balancing about $1/3$ of the radiative losses in this galaxy. They find the contribution from type II SN to be insignificant and argue against AGN feedback on grounds of the very relaxed appearance of NGC~6482.  In this paper, we show that the distribution of temperature gradients is {\\em not\\/} bimodal. We further show that the temperature gradients within the inner 2 optical effective radii are {\\em not\\/} strongly influenced by the environment. Instead, we find evidence that these inner gradients owe their origins either to the specific nature of low-luminosity AGN feedback or to a declining importance of AGN relative to compressive heating or supernovae.  In \\S\\ref{s5.dataanalysis}, we summarize our analyses and results from Papers I and II, and describe the methodology to derive radial temperature profiles. We then discuss the various types of temperature profiles seen in our sample in \\S\\ref{s5.temperature}. For a quantitative analysis, we split the radial range into two regions: the inner region extending out to 2 effective radii and an outer region between $2-4$ effective radii. We fit and analyze the gradients in these two regions separately, and demonstrate that the inner gradient is determined by galaxy properties, while the outer gradient is strongly influenced by the presence of neighboring galaxies and/or a hot ambient medium. In \\S\\ref{s5.discussion}, we discuss the implications of our findings for cooling flows, SN heating, and AGN feedback, before we briefly summarize in \\S\\ref{s5.conclusions}. ",
        "conclusions": "We have reported on the shape of temperature profiles in 36 normal elliptical galaxies. These profiles show a variety of different profile types: purely positive gradients, purely negative gradients, quasi-isothermal and even hybrid profiles. To understand this complexity, we derive mean temperature gradients for an inner region within $2\\RJ$, excluding the central point source, and an outer region between $2-4\\RJ$. We find that the outer temperature gradient is independent of intrinsic galaxy properties, but a strong function of environment, such that positive outer temperature gradients are restricted to cluster and group environments. This suggests that the outer gradients are caused by interaction with hotter ambient gas, whereas galaxies with negative outer gradients are in less dense environments and lack this intergalactic gas reservoir.  The inner temperature gradient, on the other hand, is completely independent of the environmental influence. Instead, we find that it is correlated with a number of intrinsic galaxy properties; in decreasing order of significance, the $20\\cm$ radio luminosity, the central velocity dispersion, the absolute $K$ magnitude, and the $6\\cm$ radio luminosity.  The data cannot rule out the idea that negative gradients can be produced by compressional heating in low-temperature systems, during a slow cooling inflow in a steep gravitational potential. SN feedback may also provide sufficient energy to offset cooling in X-ray faint galaxies, but we find no direct evidence that SN heating dominates.  Our preferred feedback model involves the central AGN. The inner temperature gradient is most strongly correlated with radio luminositiy and central velocity dispersion, which may be interpreted as a surrogate for black hole mass \\citep{TremaineBHsigma}. The nature of these correlations is such that weak AGN hosts show negative temperature gradients, whereas more luminous AGN exclusively live in positive gradient systems. Thus, we propose three scenarios, to explain the observed features. (1) Weak AGN distribute their heat locally, whereas luminous AGN heat the gas more globally with their extended jets. (2) The onset of negative gradients marks the point where AGN heating becomes unimportant, and compressional heating or SN feedback becomes dominant. (3) A cyclic model in which the AGN drives an outflow, which shuts the AGN activity off until the flow reverses itself, fuels the black hole and starts another cycle.  These findings are in agreement with the results from Paper I, which showed that precise hydrostatic equilibrium does not hold for the hot gas in elliptical galaxies, and established the prevalence of disturbances in the X-ray gas morphology. The results of Paper II indicate that the central AGN probably causes these disturbances. Combining these results with the connection between the temperature structure and the radio luminosity of the system produces a strong argument for the general importance of AGN feedback in nearly all normal elliptical galaxies."
    },
    "0806/0806.1589_arXiv.txt": {
        "abstract": "We present optical spectropolarimetric observations of Type Ic supernova (SN) 2007gr with Subaru telescope at 21 days after the maximum brightness ($\\sim 37$ days after the explosion). Non-zero polarization as high as $\\sim 3$ \\% is observed at the absorption feature of \\ion{Ca}{ii} IR triplet. The polarization of the continuum light is $\\sim 0.5$\\% if we estimate the interstellar polarization (ISP) component assuming that the continuum polarization has a single polarization angle. This suggests that the axis ratio of the SN photosphere projected to the sky is different from unity by $\\sim 10$\\%. The polarization angle at the \\ion{Ca}{ii} absorption is almost aligned to that of the continuum light. These features may be understood by the model where a bipolar explosion with an oblate photosphere is viewed from the slightly off-axis direction and explosively synthesized Ca near the polar region obscures the light originated around the minor axis of the SN photosphere. Given the uncertainty of the ISP, however, the polarization data could also be interpreted by the model with an almost spherically symmetric photosphere and a clumpy \\ion{Ca}{ii} distribution. ",
        "introduction": "\\label{sec:intro}  The explosion mechanism of core-collapse supernovae (SNe) is still under debate. Recent theoretical studies show that the effects causing non-spherical explosion, such as magnetic field, rotation, and several kinds of instabilities, are important for the successful explosion (\\eg Blondin et al. 2003; Kotake et al. 2004; Buras et al. 2006; Burrows et al. 2006).  In these circumstances, observational constraints on the SN asymmetry are important. Such observations include the direct imaging of Galactic young supernova remnants (\\eg Hwang et al. 2004) and SN 1987A (Wang et al. 2002), although the number of such objects is limited. The observations of extragalactic, point-source SNe also give clues to the explosion geometry. For example, the shape of emission lines in optical spectra at $\\gsim$ 1 year after the explosion can be used to study the multi-dimensional structure in the innermost part of core-collapse SNe (Mazzali et al. 2005; Maeda et al. 2008; Modjaz et al. 2008).  The most direct way to study the asymmetry of extragalactic, point-source SNe at early phases is polarimetry. Since the polarized lights scattered by electrons in the ejecta are completely cancelled out in the spherically symmetric case, the detection of polarization undoubtedly indicates the deviation from spherical symmetry (Shapiro \\& Sutherland 1982; H\\\"oflich 1991). With spectropolarimetry, the polarization across the P-Cygni profile can give the information on the distribution of elements. Spectropolarimetric studies have clarified the asymmetric nature of core-collapse SNe in detail (\\eg Cropper et al. 1988; Trammell, Hines \\& Wheeler 1993; Wang et al. 2001; Leonard et al. 2001, 2006; Kawabata et al. 2002).  In this paper, we present the spectropolarimetric observation of SN 2007gr. SN 2007gr was discovered in NGC 1058 (Li et al. 2007). Thanks to the short distance to this host galaxy \\footnote{Crockett et al. (2008) adopted 10.6 Mpc for the distance to the SN (Schmidt et al. 1992; Terry et al. 2002; Pilyugin et al. 2004) while Valenti et al. (2008) adopted 9.3 Mpc (Silbermann et al. 1996).}, SN 2007gr reached 12.7 mag in the R band at maximum brightness (Valenti et al. 2008), making the spectropolarimetric observation possible. Photometric and spectroscopic properties of SN 2007gr were shown by Valenti et al. (2008). They classified this SN as Type Ic due to the absence of the H and He lines. The progenitor of Type Ic SNe is thought to have lost its H and He layers before the explosion (\\eg Wheeler et al. 1987; Nomoto et al. 1994). Non-detection in the pre-explosion {\\it Hubble Space Telescope} image at the SN position is consistent with a Wolf-Rayet star progenitor (Crockett et al. 2008). The maximum brightness of SN 2007gr is similar to that of normal SNe, being powered by the decay of $\\sim$ 0.07$-$0.1 $\\Msun$ \\Nifs\\ (Valenti et al. 2008).  We present in \\S \\ref{sec:obs} the observation and data reduction of SN 2007gr. Results of our spectropolarimetric observation are shown in \\S \\ref{sec:res}, where the interstellar polarization (ISP) is also discussed. We study the multi-dimensional explosion geometry of SN 2007gr in \\S \\ref{sec:str} and summarize conclusions in \\S \\ref{sec:con}. ",
        "conclusions": "\\label{sec:con}  We have presented optical spectropolarimetric observations of Type Ic SN 2007gr at 21 days after maximum brightness ($\\sim 37$ days after the explosion). A strong polarization feature is observed across the \\ion{Ca}{ii} IR triplet line, which is as high as $\\sim 3$ \\% independently of the choice of ISP.  If we assume that the SN photosphere has a single polarization angle, the intrinsic polarization of the SN continuum is $\\sim 0.5$\\%. This suggests that the axis ratio of the photosphere projected on the sky is different from unity by $\\sim 10$\\%. The polarization angle of the continuum is $\\sim 60^{\\circ}$, which is similar to that of the strong polarization in the \\ion{Ca}{ii} line ($\\sim 50^{\\circ}$). The polarization of the continuum and the \\ion{Ca}{ii} line in SN 2007gr are similar to those found in broad-line SN Ic 2002ap after maximum but not similar to SN 2002ap around maximum.  The spectropolarimetric features of SN 2007gr can be explained by a bipolar explosion viewed from the slightly off-axis direction (A2 in Fig. \\ref{fig:model}). The photosphere has an oblate shape in the equatorial plane. The explosively synthesized \\ion{Ca}{ii} is distributed near the polar region, and it partially obscure the photosphere. The distributions of \\ion{O}{i} and \\ion{Na}{i} follow the photosphere, being different from \\ion{Ca}{ii}. Given the uncertainty of ISP, however, the polarization data could also be interpreted by almost spherically symmetric photosphere and the aspherical \\ion{Ca}{ii} distribution (B in Fig. \\ref{fig:model})."
    },
    "0806/0806.4824_arXiv.txt": {
        "abstract": "Narrow Line Seyfert 1 galaxies (NLS1) are very interesting objects which display peculiar properties when compared to their broad line analogues (BLS1). Although well studied in many wavebands, their behaviour at $>$10 keV  is poorly studied and yet important to discriminate between models invoked to explain the complexity observed in the X-ray band. Here we present for the first time high energy observations (17-100 keV) of five NLS1 galaxies (3 bona fide and 2 candidates) detected by  \\emph{INTEGRAL}/IBIS and provide for all of them a broad band spectral analysis using data obtained by \\emph{Swift}/XRT below 10 keV. The combined \\emph{INTEGRAL} spectrum is found to be steeper ($\\Gamma$=2.6$\\pm$0.3) than those of classical Seyfert 1 objects. This is due to a high  energy cutoff, which is required in some individual fits as in the average broad band spectrum. The location of this high energy cutoff  is at lower energies (E$\\leq$60 keV) than typically seen in classical type 1 AGNs; a reflection component may also be present but its value (R$<$0.8) is compatible with those seen in standard Seyfert 1s. We do not detect a soft excess in individual objects but only in their cumulative spectrum. Our results suggest a lower plasma temperature for the accreting plasma  which combined to the high accretion rates (close to the Eddington rate) point to different nuclear conditions in broad and narrow line Seyfert 1 galaxies, likely related to different evolutionary stages. ",
        "introduction": "Narrow Line Seyfert 1s (hereafter NLS1) are a sub-class of Seyfert 1 galaxies that display very peculiar and interesting properties. At  optical wavelength they differ from Broad Line Seyfert 1 (BLS1) galaxies for having: (1) the full width at half-maximum (FWHM) of the H$_\\beta$ line lower than 2000 km s$^{-1}$ (but see section 2); (2) the permitted lines only slightly broader than the forbidden lines; (3) the [OIII] $\\lambda$5007/H$_\\beta$ ratio $<$3; and (4) unusually strong FeII and other high ionisation emission line complexes (Osterbrock \\& Pogge 1985).  A popular explanation of the differences in the optical properties across the Seyfert population is that NLS1s have smaller black hole masses than their broad line analogues. Given that NLS1s have comparable luminosity to that of the BLS1s (Pounds, Done \\& Osborne 1995), they must be emitting at higher fractions of their Eddington luminosity; hence, higher fractional accretion rates ($\\dot{m}$ = $\\dot{M}$/$\\dot{M}_{Edd}$) are also required. Alternative explanations for the difference observed across the Seyfert population include the possibility that NLS1s have more distant  broad-line regions (BLR) and, hence, smaller Keplerian line widths (Wandel 1997) or that NLS1s are observed preferentially close to face on. The former possibility is discounted since NLS1s and BLS1 have BLRs of comparable sizes (Kaspi et al. 2000; Peterson et al. 2000); the latter is disfavoured by the analysis of Boroson \\& Green (1992) and the fact that the inner regions of BLS1s also appear to be observed close to face-on (Nandra et al. 1997). The first and most plausible scenario suggests that black holes in NLS1 have not yet been fed enough to become massive and are in a rapidly growing phase (Mathur 2000); if NLS1 are indeed in an early phase of black hole evolution, then they are key objects for studying formation and evolution of Active Galactic Nuclei (AGNs). Their detailed studies over many wavebands in search of similarities/differences with their broad line counterparts can enable us to reveal the formation mechanisms and processes of central black holes in the local Universe and help understanding  QSO formation and evolution.  The waveband where NLS1 show the most marked differences compared to broad line AGNs is the X-ray band. In this waveband the best known properties of NLS1 are the presence of a soft excess in most but not all objects (Turner, George \\& Netzer, 1999), an unusually strong X-ray variability (Boller et al. 1997; Leigly 1999) and 2-10 keV continuum slopes  steeper ($\\Gamma$=2.19) than in normal Seyfert 1s ($\\Gamma$=1.73, Veron-Cetty, Veron and Goncalves 2001). With the higher sensitivity of current X-ray telescopes came the discovery of a sharp spectral drop at around 7 keV in a number of NLS1 (see Gallo 2006 and references therein). The origin of this behaviour is still debated, but two models have been put forward to explain this spectral complexity: partial covering (Tanaka et al. 2004) and reflection/light bending (Fabian et al 2002, Miniutti \\& Fabian 2004). In the first case the drop is produced by absorption of the continuum by dense material which partly obscures the source; in the second case it is simply the blue wing of the relativistically broadened iron line. Despite being successful in explaining most of the 2-10 keV spectral complexity observed in NLS1, the first model does not describe the nature of the primary continuum source nor provides physically acceptable parameters for the soft excess component usually fitted with a black body model. Reflection of the power law continuum source off the cold accretion disc describes more adequately all the features seen in the X-ray spectra of NLS1. Clearly broad band data extending above 10 keV are crucial to discriminate between these two competing models; unfortunately these types of observations are rare for NLS1, with only a few broad band spectra being measured by \\emph{BeppoSAX} (Comastri et al 2001, Dadina 2007). In these few cases, data above 10 keV did not provide conclusive results. Now, the number of known NLS1 bright above 10 keV is increasing thanks to the \\emph{INTEGRAL} and \\emph{Swift} all sky surveys (Bird et al. 2007, Tueller et al. 2007) thus providing the base for a first study of their high energy (and consequently also broad-band) spectral behaviour.  Here we present \\emph{INTEGRAL} results for a set of 5 newly discovered NLS1 (3 bona-fide and 2 candidates); our data  suggest a steeper high energy spectral shape with respect to broad line Seyfert galaxies and the preference for  a lower cut-off energy ($\\leq$60 keV). We do not detect soft excess emission in individual objects but only in their cumulative spectrum, which also provides an upper limit on the average reflection parameter (R$<$0.8). All this observational evidence makes these objects interesting in their own right even if not all of them will be confirmed as NLS1. ",
        "conclusions": "The 5 AGN discussed in this work have optical characteristics which indicate that 3 are bona fide NLS1 while 2 are only candidates and as such provide the first sample (albeit small) of such objects selected above 10 keV. Analysis of their high energy spectra reveals for the first time a different behaviour above 10 keV with respect to their broad line analogues and in particular a steeper ($\\Gamma$=2.6) power law slope than the canonical $\\Gamma$ value of $\\sim$ 2 associated to standard Seyferts.  Based on our analysis, the steeper primary continuum is likely due to a lower  energy cut-off ( E$_{c}$ $\\le$ 60 keV) than typically observed in BLS1 (Malizia et al. 2003) while the presence of a  Compton reflection component is not strongly required by the data and in any case the value found (R$<$0.8) is compatible with those generally observed in broad line AGNs. This observational evidence makes the present sample apart from the standard type 1 AGN discovered by IBIS even if not all will be confirmed NLS1.  Modelling of classical Seyfert spectra ascribe the power law to the inverse Compton scattering of soft photons off hot electrons. Variations to this baseline model depend on the energy distribution of these electrons and their location in relation to the accretion disc, often a hot corona above the disk. Within this model the power law photon index and cut-off energy are directly related to the temperature and optical depth of the Comptonising hot plasma so that precise information on these two parameters provide clues for our understanding of the primary emission mechanism, geometry of the source and physics of the plasma near the central powerhouse. If as suggested,  NLS1 galaxies are evolutionarily young objects, powered by the accretion of gas onto central black holes that are lower in mass and accreting at a higher rate than their broad-line counterparts, measurements of a different high energy continuum may have strong implications on the cosmic evolution of Supermassive Black Holes. Using appropriate relations (Petrucci et al. 2001) it is possible to derive from the values of $\\Gamma$ and E$_c$, estimates of the temperature kTe and optical depth $\\tau$ of the Comptonising hot plasma; in the case of NLS1 galaxies we obtain that kTe$\\sim$ 10-20 keV and $\\tau \\gg$1, i.e. a much lower temperature than typically found in BLS1 but similar plasma thickness. It can therefore be argued that the Comptonising hot plasma of NLS1 did not have enough time to reach the condition of the older, more massive and lower rate accreting black holes found in BLS1. This maybe related to the higher accretion rates associated to NLS1, which  provide the conditions for a much higher  cooling of the emitting plasma and therefore a lower temperature. In fact, the accretion rates estimated for  our sources are  indeed close the Eddington rate.  Clearly, optical spectra with higher signal to noise ratio are desirable  (and encoraged) to firmly assess the NLS1 classification in particular in the case of IGR J14552-5133 and IGR J16185-5928. Furthermore, to better characterise our sources, we need higher quality X-ray data as available from XMM observations and possibly simultaneous broad band coverage as it is possible  with \\emph{Suzaku}. We hope in the near future to be able to obtain both these type of measurements and put stronger constraint on the cut-off energy and reflection components in individual objects. It is also important to properly classify all AGN detected by \\emph{INTEGRAL}/IBIS and \\emph{SWIFT}/BAT in search for new gamma-ray selected NLS1, which can be tested against  our findings.        \\begin{table} \\begin{center} \\centerline{Table 5: IBIS/ISGRI + Swift/XRT  AVERAGE SPECTRUM} \\renewcommand{\\footnoterule}{} \\begin{tabular}{l c c c c} \\hline\\hline & \\texttt{mo po}         & \\texttt{mo bb po}      & \\texttt{mo bb cutoffpl} & \\texttt{mo bb bknpo}    \\\\ \\hline N$_{H}^{\\dagger}$& 0.19$^{+0.01}_{-0.01}$ & 0.27$^{+0.04}_{-0.06}$ & 0.21$^{+0.02}_{-0.02}$ & 0.23$^{+0.05}_{-0.01}$ \\\\ $\\Gamma$         & 1.82$^{+0.04}_{-0.04}$ & 1.97$^{+0.05}_{-0.04}$ & 1.82$^{+0.06}_{-0.06}$ & 1.90$^{+0.05}_{-0.05}$ \\\\ kT               &    -                   & 0.07$^{+0.03}_{-0.01}$ & 0.05$^{+0.03}_{-0.02}$ & 0.06$^{+0.03}_{-0.01}$ \\\\ E$_{cutoff}$     &    -                   &     -                  & 38$^{+17}_{-10}$       &            -           \\\\ E$_{break}$      &    -                   &     -                  &     -                  & 27$^{+7}_{-13}$        \\\\ $\\Gamma_{2}$     &    -                   &     -                  &     -                  & 3.02$^{+0.70}_{-0.43}$ \\\\ C                & 1.09$^{+0.17}_{-0.14}$ & 1.43$^{+0.22}_{-0.08}$ & 2.33$^{+0.53}_{-0.45}$ & 1.60$^{+0.27}_{-0.24}$ \\\\ $\\chi^{2}$/dof   & 458/336                & 438/334                & 397/333                & 400/332                \\\\ \\hline \\hline \\end{tabular} $\\dagger$ in units of 10 $\\times$ 10$^{22}$ cm$^{-2}$\\\\ \\texttt{po}: power law \\texttt{bb po}: blackbody + power law\\\\ \\texttt{bb cutoffpl}: blackbody + cut-off power law\\\\ \\texttt{bb bknpo}: blackbody + broken powe law \\end{center} \\end{table}"
    },
    "0806/0806.1406_arXiv.txt": {
        "abstract": "In the recent literature there is circumstantial evidence that the viscosity of the intracluster medium may not be too far from the Spitzer value. In this letter, we present two-dimensional hydrodynamical simulations of ram pressure stripping of disc galaxies in a viscous intracluster medium (ICM). The values of viscosity explored range between 0.1 and 1.0 times the Spitzer value. We find that viscosity affects the appearance and the dimensions of the galactic wakes but has very little effect on the evolution of the gas mass of the galaxy. ",
        "introduction": "In clusters, galaxies can lose some or all of their gas by ram pressure stripping (RPS) due to their motion through the intracluster medium (ICM). Both, analytical estimates (\\citealt{gunn72}) and (hydro)dynamical simulations (e.g.~\\citealt{abadi99}, \\citealt{quilis01}, \\citealt{schulz01}, \\citealt{vollmer01}, \\citealt{marcolini03}, \\citealt{roediger06}, \\citealt{roediger07}) show that RPS can remove a significant amount of gas from galaxies, and is thus important for the evolution of galaxies and the ICM.  In addition to gas loss by ram pressure pushing, galaxies also suffer gas loss by continuous (sometimes also called turbulent or viscous) stripping (e.g.~\\citealt{nulsen82,quilis00,schulz01,roediger05,roediger06}).  \\citet{nulsen82} has studied the effects of transport processes and turbulence on the flow of gas past a galaxy and has found that they could produce more stripping of gas than ram pressure alone.  For turbulent stripping, he found a mass stripping rate of \\begin{equation} \\dot{M}_{\\rm turb}\\sim \\pi r^2 \\rho_{\\rm ICM} v \\sim 7\\, r_{10}^2 n_{-3}v_3 M_{\\odot} {\\rm yr}^{-1} , \\end{equation} where $r=10r_{10}$ kpc, $\\rho_{\\rm ICM}=10^{-3}n_{-3}m_{\\rm H}$ g cm$^{-3}$ and $v=1000 v_3$ km s$^{-1}$. This is similar to what is found in the simulations by \\citet{roediger06, roediger07} who find rates of about $\\sim 1 M_{\\odot} {\\rm yr}^{-1}$.  Turbulent stripping is mainly caused by the Kelvin-Helmholtz instability, which is suppressed by viscosity that stabilises modes with wavelengths smaller than $r/{\\rm Re}$ (e.g. \\citealt{betchov67}).  There exist only few detailed studies on the wakes of stripped galaxies. Using the hydrodynamical adaptive mesh refinement code FLASH, we have studied galactic ram pressure tails in a constant ICM wind (\\citealt{roediger06wakes}) as well as for galaxies on realistic cluster orbits (\\citealt{roediger08}). The minimum tail width of about 20 to 30 kpc is found near the galaxy. With increasing distance to the galaxy, the tail flares to widths of 30 to 80 kpc at a distance of $\\sim 100\\Kpc$ behind the galaxy. Other hydrodynamical simulations using either grid codes (e.g.~\\citealt{quilis01}, \\citealt{marcolini03}) or smoothed particle hydrodynamics (e.g.~\\citealt{abadi99}, \\citealt{schulz01}) do not focus on the tails. However, according to the snapshots provided in these papers, the gas tails are similar to the ones in our simulations.  Using a sticky-particle code, Vollmer et al.~(e.g.~1999, 2000, 2001a, 2001b, 2003, 2004a, 2005, 2006\\nocite{vollmer99,vollmer00, vollmer01,vollmer01a,vollmer03,vollmer04,vollmer05,vollmer06}) aim at reproducing the RPS history of individual galaxies by comparing simulations and observations. However, these simulations concentrate on the gas distribution close to the galaxy.  Recently, \\citet{oosterloo05} presented observations of a $\\sim 120$ kpc long and $<25\\Kpc$ wide tail of H~I gas associated with NGC~4388, and also suggest that this tail is due to ram-pressure stripping of this galaxy, either in the ICM of the Virgo cluster or in the halo of the nearby elliptical galaxy M86. However, the tail of NGC~4388 is narrower than the tails in the simulations presented in \\citet{roediger06, roediger08}, although it seems to be flaring in a similar fashion as found in the simulation. Also the X-ray and H$\\alpha$ tails observed by \\citet{sun06,sun07,yagi07} are much narrower than simulated ram pressure tails: While being only $\\sim 7\\Kpc$ wide, they reach lengths of $\\sim 70\\Kpc$ and show hardly any flaring. This difference may be caused by the microphysics of the ICM.  The viscosity of the intracluster medium has been discussed before. Based on observations of the Perseus cluster it has been suggested by \\cite{fabian03a}, \\cite{fabian03b} that viscosity may play an important role in dissipating energy injected by the central AGN.  Circumstantial evidence for the presence of significant ICM viscosity is also provided by an examination of the morphology of H$\\alpha$ filaments in the Perseus cluster. Several of the filaments appear to trace well-defined arcs which argues against the presence of strong turbulence in the ICM core, possibly resulting from the action of viscosity. This idea has been tested in numerical simulations by \\citet{ruszkowski04, ruszkowski04a} and \\citet{reynolds05}.     In the case of a fully ionized and unmagnetised, thermal plasma, the relevant coefficient of viscosity is given by \\citet{braginskii58} and \\cite{spitzer62} as $\\mu \\approx 6.0\\times 10^{-17}(\\ln\\Lambda/37)^{-1}T^{5/2}$ g cm$^{-1}$ s$^{-1}$, where $T$ is the temperature of the plasma measured in Kelvin and $\\ln \\Lambda$ is the Coulomb logarithm. It results from the cumulative effect of weak Coulomb collisions. The mean free path of such interactions scales with $T^2$. As the sound speed scales with $T^{1/2}$, the viscosity is proportional to $T^{5/2}$.  It is customary to measure the importance of viscosity through the Reynolds number, Re $=vl/\\nu$, where $v$ and $l$ are characteristic velocities and length-scales of the system and $\\nu=\\mu/\\rho$ the kinematic viscosity, with $\\rho$ being the fluid density.  For our case  \\begin{equation} {\\rm Re} \\sim 26\\, v_3 r_{10}n_{-3}f_v^{-1}\\left (\\frac{kT}{5 {\\rm keV}} \\right )^{-2.5} , \\end{equation} which indicates that viscosity may play a role unless the viscosity is strongly suppressed, i.e. the suppression factor, $f_{v}$, is sufficiently small.  This suppression results from the cluster magnetic fields. Even weak fields lead to a tiny proton gyroradius, which results in a very efficient suppression (factor of $\\sim 10^{23}$ for typical ICM conditions, \\citealt{spitzer62}) of the local viscosity perpendicular to the magnetic field.  Magnetic fields in the ICM are certainly tangled or even chaotic (\\citealt{clarke04,ensslin05}) and will lead to a reduced macroscopic viscosity. The degree of reduction, however, is unknown. Due to the same mechanism, also thermal conduction in the ICM is suppressed perpendicular to the magnetic field lines. In a recent study, \\citet{narayan01} found an effective macroscopic thermal conductivity that is a factor of $f_{v} \\sim 10^{-2}-0.2$ lower than the unmagnetised value. Similar arguments may apply to the viscosity. From studies of ICM turbulence in the Coma cluster, \\citet{schuecker04} derive an upper limit on the kinematic viscosity of the ICM of $\\sim 3\\cdot 10^{29} \\mathrm{cm}^2 \\mathrm{s}^{-1}$. For typical ICM densities and temperatures, this corresponds to a viscosity suppression factor $f_{v}$ around 0.1. Thus, we might expect the viscosity to be suppressed by some factor between $10^{-2}$ and unity.  The magnetic fields of the ICM (e.g.~\\citealt{clarke04,ensslin05}) may themselves influence the appearance of ram pressure stripped galaxies. Given that the thermal pressure in the ICM dominates over the magnetic pressure, the magnetic fields should be frozen-in and follow the ICM flow. Thus, they should be generally parallel to the galaxy's tail. Such a magnetic field structure could attenuate thermal conduction between the stripped gas and the hot ICM, and it could suppress e.g.~Kelvin-Helmholtz and maybe even Rayley-Taylor instabilities in the tail. However, it is unclear how well the magnetic fields will be aligned with the ICM-ISM interfaces and how strong this suppression will be. The influence of the magnetic fields will be studied in a forthcoming paper.  Here, using adaptive-mesh, hydrodynamical simulations, we investigate the effect of a macroscopic viscosity on the stripping of gas from a galaxy. ",
        "conclusions": "Clearly, we have to be careful with quantitative predictions especially concerning the morphology of the wake because our simulations are two-dimensional. However, in previous work we have shown that mass loss rates are the same in 2D and 3D simulations. Moreover, in the non-viscous cases, the wake structures in 2D and 3D were very similar. Thus, our main conclusions should be robust: \\begin{itemize} \\item The mass loss from the gas discs is hardly influenced by viscosity. \\item With increasing viscosity, the wake shows less stucture and turbulence, but larger clumps. \\end{itemize}  A detailed investigation of the fate of the stripped gas requires several additions to our simplified model: a 3D treatment to allow for other inclinations than face-on, prescriptions for heating, cooling and thermal conduction. These processes do not only influence the temperature in the stripped gas, but also determine in which wavebands the stripped gas will be observable. For the fate of the stripped gas, viscosity may make a difference as it may affect the thermal history of the stripped gas. Gradients of density and temperature are smeared out in the wake in the presence of viscosity. This will affect the efficiency of heat conduction and evaporation of clumps in the galactic wakes. It will also affect the efficiency with which material in the wake can form stars as observed by \\citet{sun07}. Comparisons between models and observations will reveal which processes are the dominant ones in shaping galactic tails."
    },
    "0806/0806.2405_arXiv.txt": {
        "abstract": "% Observing massive galaxies at various redshifts is one of the most straightforward and direct approaches towards understanding galaxy formation.  There is now largely a consensus that the massive galaxy (M$_{*} >$ \\mass) population is fully formed by $z \\sim 1$, based on mass and luminosity functions.  However, we argue that the latest data can only rule out number and mass density evolution of a factor of $> 2-3$ at $z < 1.5$.  We furthermore show that the star formation history of M$_{*} >$ \\mass galaxies reveals that 40$\\pm$5\\% of galaxies with M$_{*} >$ \\mass at $z \\sim 1$ are undergoing star formation that effectively doubles their stellar mass between $z = 0.4 - 1.4$. These massive galaxies also undergo 0.9$^{+0.7}_{-0.5}$ major mergers during this same time period. ",
        "introduction": "Understanding when and how massive galaxies form is one of the most outstanding problems in cosmology and galaxy formation.   Galaxies are predicted in Cold Dark Matter based models of structure formation to form gradually with time through the merging of smaller systems (e.g,. Cole et al. 2001).  While there is some evidence for this process, in terms of galaxies (e.g., Patton et al. 2002; Conselice et al. 2003; Bridge et al. 2007), many details are still lacking. Alternatively, massive galaxies, which are mostly ellipticals in today's universe (e.g., Conselice 2006a), may have formed very rapidly in  `monolithic' collapses.  A such, massive galaxies are the dominant test-bed for galaxy models, and understanding their evolution observationally is an important goal.   Most massive galaxies at $z < 1$ are red and passively evolving, yet star formation and merging activity have been seen in ellipticals from $z \\sim 0$ to $z \\sim 1$ (Stanford et al. 2004; Lin et al. 2004; Teplitz et al. 2006; Conselice et al. 2007a,b) - thus it is unclear when or how massive galaxies finally assembled. We must study this process at high redshift since the ages of stars in nearby galaxies cannot reveal their entire formation history, as the assembly of mass is likely decoupled from the formation of stars (e.g., Conselice 2006b; Trujillo et al. 2007).  Observational evidence suggests that passively evolving massive galaxies exist at $z \\sim 1$, and likely at even early times, at $z > 2$ (Daddi et al. 2004; Saracco et al.  2005; Bundy et al. 2006; Conselice et al. 2007a). Recent claims also exist for the establishment of the full massive galaxy population by $z \\sim 1$ (e.g., Drory et al. 2005; Bundy et al. 2006). However, what is not yet clear is if number densities measured in these surveys are able to rule out evolution at $z < 1$ after considering uncertainties in measuring stellar masses, number densities, and cosmic variance.   \\setcounter{figure}{0} \\begin{figure}[!h] \\plotone{conselice_fig1.eps} \\caption{Left panel: the evolution in the number densities for galaxies of various masses between $z \\sim 0.4 - 1.4$.  Right panel: the stellar mass density evolution as a function of galaxy mass at the same redshift intervals.  The points at $z \\sim 0$ are taken from Cole et al. (2001).  The error bars listed on both the number and mass densities reflect uncertainties from stellar mass errors, as well as cosmic variance, and counting statistics (Conselice et al. 2007b).  The solid and dashed lines show predictions of these quantities from the Millennium simulation.} \\end{figure}  On the other hand, at $z > 1.5$ it appears that there are significantly fewer massive galaxies than at $z < 1.5$ (e.g., Drory et al. 2005). Observationally,  the most massive galaxies at $z > 1.5$ are undergoing major mergers, which are able to construct the stellar masses of massive galaxies rapidly (Conselice 2006b; Conselice et al. 2008a).   The situation at $z < 1.5$ is not as clear, with observations inconclusive on whether there is evolution in the massive galaxy population at $z < 1$ (e.g., Brown et al. 2006). We argue here that by using a stellar mass selected sample of galaxies the number and mass densities of massive galaxies are consistent, within their errors, with no evolution at $z < 1.5$, yet there is observable evolution when examining in detail the physical processes occurring within these galaxies. Up to 50\\% of galaxies with M$_{*} >$ \\mass are undergoing star formation at $z \\sim 1$, and roughly one major merger occurs within these systems at $z < 1.4$. We assume throughout a standard cosmology of H$_{0} = 70$ km s$^{-1}$ Mpc$^{-1}$, and $\\Omega_{\\rm m} = 1 - \\Omega_{\\lambda} = 0.3$. ",
        "conclusions": "While massive galaxies can be identified up to $z \\sim 3$ using deep NIR imaging, the evolution of these galaxies is difficult to measure within a factor of 2-3 using simply luminosity or mass functions. By examining the change in the mass function, and through structural parameters and MIPS 24 $\\mu$m and [OII] line emission we show that massive galaxies are still undergoing some evolution, with at least as much as a factor of two increase in stellar mass at $z < 1.4$ from star formation and merging."
    },
    "0806/0806.0293_arXiv.txt": {
        "abstract": "Gravitational waves potentially represent our only direct probe of the universe when it was less than one second old. In particular, first-order phase transitions in the early universe can generate a stochastic background of gravitational waves which may be detectable today. We briefly summarize the physical sources of gravitational radiation from phase transitions and present semi-analytic expressions for the resulting gravitational wave spectra from three distinct realistic sources: bubble collisions, turbulent plasma motions, and inverse-cascade helical magnetohydrodynamic turbulence. Using phenomenological parameters to describe phase transition properties, we determine the region of parameter space for which gravitational waves can be detected by the proposed Laser Interferometer Space Antenna. The electroweak phase transition is detectable for a wide range of parameters. ",
        "introduction": "Since gravitational waves propagate freely through the universe after being generated, their detection provides a powerful test of the very early universe. Various mechanisms that generate such gravitational waves have been discussed: quantum fluctuations during \\cite{inflation} and shortly after inflation \\cite{end-inflation}; bubble wall motion and collisions during phase transitions \\cite{bubble,kos1,ktw92a,kt93,kos2}; cosmic strings and defects, including primordial black holes \\cite{strings}; cosmological magnetic fields \\cite{magnet,cdk04,cd06}, and plasma turbulence \\cite{kmk02,dolgov,kgr05,gkk07}. This paper focuses on gravitational waves generated during phase transitions and study the possibility of direct detection by the planned Laser Interferometer Space Antenna (LISA),  Near-future data from the Large Hadron Collider will for the first time have the ability to probe in detail physics at the electroweak energy scale. This physics determines the nature of the electroweak phase transition in the early universe, when the primordial plasma went from an electroweak-symmetric state to a broken state with distinct electromagnetic and weak interactions. The phase transition took place when the primordial plasma had a temperature on the order of 1 TeV. Intriguingly, the Hubble frequency $H_*$ at this epoch, redshifted to today, falls in the lower end of the detection range for LISA, which ranges from $10^{-4}$ to $10^{-1}$ Hz \\cite{lisa}. If gravitational waves are generated during the electroweak phase transition, their characteristic frequency will be related to this Hubble frequency: remarkably, gravitational waves may provide us with an alternate route to probing electroweak physics.  Having gravitational waves at detectable frequencies is not sufficient for detecting the electroweak phase transitions: the source of gravitational waves must be sufficiently strong to produce radiation with a detectably large amplitude today \\cite{detection,nicolis,gs06,grojean}. If the phase transition is first order, the latent heat of the phase transition partly is transferred into kinetic energy of the walls of expanding bubbles of the broken phase. If these expanding bubbles contain a large enough amount of energy, they will produce a gravitational wave background with a detectably large amplitude today \\cite{bubble}. Models of the electroweak phase transition based on standard-model particle physics do not produce observable gravitational wave signals since they are not first order for allowed values of the Higgs mass \\cite{stdmodelweak}, but common extensions of the standard model, including supersymmetry and extra dimensions, can produce much stronger phase transitions \\cite{strong}.  Bubbles of broken phase in a first-order phase transition expand and percolate to convert the entire universe to the low-energy phase. The kinetic energy in the bubble walls eventually thermalizes, but prior to that the bubbles act to stir the primordial plasma, plausibly generating Kolmogoroff turbulence cascading from the turbulence scale to a much smaller scale where the kinetic energy turns into thermal energy via viscous heating. The turbulent eddy motions can also be a potent source of gravitational waves if the energy input is large enough \\cite{kmk02}. If a small magnetic seed field with nonzero helicity is generated at the phase transition \\cite{helicity,ste08}, magnetohydrodynamic effects can generate an inverse cascade, transferring energy to scales substantially larger than the stirring scale \\cite{gkk07} and resulting in a detectably large signal \\cite{kgr08}, and additionally can give the radiation a non-zero circular polarization \\cite{kgr05} which is also directly detectable \\cite{seto}. So gravitational wave generation from an early universe phase transition can be decomposed into three distinct sources: expanding bubbles of the broken phase, hydrodynamic turbulence stirred by the colliding bubbles, and an inverse cascade due to the amplification of seed magnetic fields.  Previous papers have computed the gravitational wave spectra for each of these sources individually. Here we consider their combined spectrum. The detailed shape and amplitude of the spectrum depends on fundamental properties of the phase transition: its energy scale, bubble nucleation rate, latent heat, efficiency of converting latent heat into plasma motions, and mean helicity of seed magnetic fields. We summarize how the total power spectrum depends on these parameters, then evaluate the regions of parameter space for which the relic gravitational radiation from the phase transition is detectable with LISA, displaying the results as contour plots in the parameter space. While the bubble collisions produce detectable gravitational radiation for a very limited range of parameters, adding the radiation from turbulence and an MHD inverse cascade greatly widens the range of detectable phase transitions.  In the following Section,  we give a brief overview of the physics of phase transitions and the corresponding phenomenological parameters. In Sec.~III, we express the gravitational wave spectra in terms of these phenomenological parameters, and show how the total gravity wave spectrum varies with the parameters. Sec.~IV compares these spectra with the projected LISA sensitivity curve for stochastic backgrounds, displaying detectability regions in the space of phase transition parameters. ",
        "conclusions": ""
    },
    "0806/0806.2891_arXiv.txt": {
        "abstract": "I shall briefly review the observational properties of radio-emitting jets from Galactic X-ray binaries in relation with other wavebands: infrared, optical and X-ray. Special attention is paid to recent results obtained with the Spitzer Space Telescope on quiescent black holes as well as ultra-compact neutron star X-ray binaries. ",
        "introduction": "Traditionally, the key observational aspect of X-ray binary jets lies in their synchrotron radio emission. Such jets appear to come in two types: milliarcsec-scale, continuous jets with flat radio-mm spectra, and arcsec-scale optically thin jets resolved into discrete plasmons moving away from the binary core (Fender 2006).  These are commonly referred to as {\\it steady} and {\\it transient} jets, respectively. Over the last few years, however, observations at higher wavelengths have added substantially to our knowledge in this field. In the following, I shall review the properties of X-ray binary jets in response to global changes in the accretion flow, focusing on recent progress made in the mid-IR thanks to the unprecedented sensitivity of the Spitzer Space Telescope. ",
        "conclusions": ""
    },
    "0806/0806.3758.txt": {
        "abstract": "Some analyses of recent cosmic microwave background (CMB) data have provided hints that there are deviations from Gaussianity in the WMAP CMB temperature fluctuations. Given the far reaching consequences of such a non-Gaussianity for our understanding of the physics of the early universe, it is important to employ alternative indicators in order to determine whether the reported non-Gaussianity is of cosmological origin, and/or extract further information that may be helpful for identifying its causes. We propose two new non-Gaussianity indicators, based on skewness and kurtosis of large-angle patches of CMB maps, which provide a measure of departure from Gaussianity on large angular scales. A distinctive feature of these indicators is that they provide sky maps of non-Gaussianity of the CMB temperature data, thus allowing a possible additional window into their origins. Using these indicators, we find no significant deviation from Gaussianity in the three and five-year WMAP ILC map with \\emph{KQ75} mask, while the ILC unmasked map exhibit deviation from Gaussianity, quantifying therefore the  WMAP team recommendation to employ the new mask \\emph{KQ75} for tests of Gaussianity. We also use our indicators to test for Gaussianity the single frequency foreground unremoved WMAP three and five-year maps, and show that the K and Ka maps exhibit clear indication of deviation from Gaussianity even with the \\emph{KQ75} mask. We show that our findings are robust with respect to the details of the method. ",
        "introduction": "%%%%%%%%%%%%%%%%%%%%%%%%%%  Within the standard approach to cosmological modelling the suggestion that the Universe underwent a brief period of rapid acceleration expansion% ~\\cite{Inflation-1st-refs} before the epoch of primordial nucleosynthesis has become an essential building block of the standard cosmological model. Besides solving the so-called flatness, horizon and monopole problems, such a period of cosmological inflation  provides a mechanism for the production of the primordial density fluctuations, which seeded the observed cosmic microwave background (CMB) anisotropies and the formation of large-scale structure in the Universe.  There are more than one hundred inflationary models (see, e.g., the review articles Refs.~\\cite{Inflation-reviews}), among which the simple ones are based on a slowly-rolling single scalar field. An important prediction of a number of these simple models is that they can generate only tiny non-Gaussianity, which should be undetectable in the Wilkinson Microwave Anisotropy Probe (WMAP) CMB data% ~\\cite{Gauss_Single-field}. There are, however, a large class of inflationary models that can generate non-Gaussianity at a level detectable by WMAP~\\cite{Non-standard-models}. These scenarios comprise models based upon a wide range of mechanisms, including special features of the inflation potential, multiple scalar fields, non-canonical kinetic terms, and  non-adiabatic fluctuations %among others (see the review Ref.~\\cite{Bartolo2004} and references therein). Thus, the detection of non-Gaussianity in CMB data may potentially be useful to discriminate inflationary models and shed light on the physics of the early universe.  In the statistical analyses by using one, three and five-year% ~\\cite{wmap1,wmap3,wmap5,Gauss-consist} CMB measurements  along with some different statistical tools, the WMAP team have found that the CMB data are consistent with Gaussianity. However, some recent analyses  have provided clear hints that there are significant deviations from Gaussianity in the WMAP data. Clearly the study of detectable non-Gaussianities in the WMAP data must take into account that they may have non-cosmological origins as, for example, unsubtracted contamination from galactic diffuse foreground emission~\\cite{Chiang-et-al2003,Naselsky-et-al2005} and unconsidered point sources~\\cite{Eriksen-et-al2004}. % They may also have cosmic topology origin~\\cite{} If they turn out to have a cosmological origin, however, this could have far-reaching consequences on our description of the Universe, particularly on the inflationary picture.  In view of this, a great deal of effort has recently gone into verifying the existence of such non-Gaussianity by employing several different statistical signatures of non-Gaussianity in its various forms (see, e.g., Refs.~\\cite{Some_non-Gauss-refs} and related Refs.~\\cite{Non-Gauss_related}). Apart from revealing the existence of non-Gaussianity in CMB data, different statistical tools are sensitive to different systematics and may be useful in determining their origins. In addition, different indicators can in principle provide information about the multiple types of non-Gaussianity that may be present in CMB data. It is therefore important to test the data for deviations from Gaussianity by using a range of different statistical tools to identify any non-Gaussian signals on the CMB sky.  In this paper, we propose new large-angle non-Gaussianity indicators, based on skewness and kurtosis of large-angle spherical-shaped patches of CMB maps, which provide a measure of departure from Gaussianity on large angular scales. A distinctive feature of these indicators is that they provide sky maps (directional information) of non-Gaussianity of the CMB temperature fluctuations, thus allowing a possible additional window into their causes. Using these indicators, we find no significant deviation from Gaussianity in the WMAP three and five-year ILC \\emph{KQ75} masked map, but ILC unmasked map exhibit deviation from Gaussianity. On the other hand, our indicators reveal deviations from Gaussianity of different degree in the three and five-year single frequency K and Ka \\emph{KQ75} masked maps, which is consistent with the fact that even these masked maps are still foreground contaminated at some level.  The structure of the paper is as follows. In Sec.~\\ref{Indicators} we introduce our non-Gaussianity indicators. Section~\\ref{Indic_vs_data} contains the results of applying our statistical indicators to the three-year and five-year WMAP data, and f\\/inally in Sec.~\\ref{Conclusions} we present the summary of our main results and conclusions.  %%----------------------------------------------- %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% ",
        "conclusions": "\\label{Conclusions} %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% %%-------------------------------------------  We have proposed two new indicators for measuring large-angle directional deviations from Gaussianity in the CMB data, and have used them to search for the large-angle (low $\\ell$) deviation from Gaussianity in the three and five-year foreground reduced ILC and the five single frequency (K, Ka, Q, V, W) \\emph{KQ75} masked maps. Our directional indicators enable us to construct skewness $S$ and kurtosis $K$ maps (as, e.g., Fig.~\\ref{Fig1} and Fig.~\\ref{Fig2}), making it possible to examine the presence and significance of possible large-angle non-Gaussianity in the WMAP CMB temperature fluctuations maps with and without \\emph{KQ75} mask.  To obtain a more quantitative measure of non-Gaussianity we have studied the low $\\ell$ angular power spectrum of the $S$ and $K$ maps generated from the (three and five-year) WMAP ILC and frequency (unremoved foreground) maps with and without \\emph{KQ75} mask. For the full-sky ILC maps we found deviation from Gaussianity for the low $\\ell$, while for the ILC masked maps we found that the low $\\ell$ ($\\,\\ell=1,\\,\\,\\cdots,10\\,$) components are not significantly different from corresponding components of the expected power spectrum calculated from $S$ and $K$  maps obtained from $1\\,000$ Monte Carlo CMB maps generated by considering the Gaussian random hypothesis based on the concordance model~\\cite{wmap5}. Actually, we have found that the values of the multipoles $S_{\\ell}$ and $K_{\\ell}$ (for $\\,\\ell=1,\\,\\,\\cdots,10\\,$) are not statistically significant, i.e. they are within the $95\\%$ values of $S_{\\ell}$'s and $K_{\\ell}$'s of the MC randomly scrambled maps] (Fig.~\\ref{Fig6} and Tables~\\ref{Skew-deviation} and~\\ref{Kurt-deviation}). As regards the frequency maps, we found clear indication of deviation from Gaussianity in the three and five-year frequency \\emph{KQ75} masked maps: K and Ka, which is expected and consistent with the fact that even these masked maps present some level of foreground contamination away from the masked region (see Fig.~\\ref{With_without}). The deviation for the  Q, V, and W masked maps are within the $95\\%$ expected values from MC randomly simulated maps.  To have an overall quantitative assessment of the power spectra $S_{\\ell}$ and $K_{\\ell}$ calculated from  K, Ka, Q, V, W and ILC maps we have performed a $\\chi^2$ test to determine the goodness of fit for low $\\ell$ multipole values as compared to the expected multipoles values from the Gaussian MC maps. In this way we have obtained numbers that collectively quantify the extent to which these \\emph{KQ75} masked maps are consistent with Gaussianity.  The results of our statistical analyses indicate that the current CMB temperature fluctuations ILC three and five-year masked data are consistent with Gaussianity, in agreement with the WMAP team and other analyses made by using different statistical tools~\\cite{wmap1,wmap3,wmap5}. We have demonstrated that the results of our analyses are robust by showing that the $S-$map and $K-$map do not significantly change with different choices of variables involved in our scheme, so long as the statistical noise is kept under control.  The effects of different foreground-reduced algorithms as detected by our non-Gaussianity indicators for other WMAP three and five-year maps~\\cite{KNC,PPG,OT} is under a careful investigation and signs of non-Gaussianity seems to be present in the maps, which may have a non-cosmological origin as, for example, residual foregrounds, artifacts of the cleaning algorithm or simply a statistical fluke.   Finally, we emphasize that the robustness of our scheme with respect to all considered parameters along with the detection on non-Gaussianity in the single frequency (foreground unremoved) maps seem to indicate that our indicators are well-suited to reliably map deviation from Gaussianity at large angular scales in the CMB data, besides being complementary to the other approaches in the literature.  %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%"
    },
    "0806/0806.0546_arXiv.txt": {
        "abstract": "We present the second data  release of the Radial Velocity Experiment (RAVE), an ambitious  spectroscopic survey to measure radial  velocities and stellar atmosphere parameters  (temperature, metallicity, surface gravity, and rotational velocity) of up to one million  stars using the 6dF  multi-object spectrograph on  the 1.2-m UK Schmidt  Telescope  of the  Anglo-Australian  Observatory  (AAO).  The  RAVE program  started  in  2003,  obtaining  medium  resolution  spectra  (median R=7,500) in  the Ca-triplet  region ($\\lambda\\lambda$ 8,410--8,795  \\AA) for southern hemisphere stars drawn from the Tycho-2 and SuperCOSMOS catalogues, in the magnitude range  $9\\!<\\!I\\!<\\!12$.  Following the first data release \\citep{dr1} the current release doubles the sample of published radial velocities, now containing 51,829 radial velocities for 49,327 individual stars observed on 141 nights between April 11 2003 and March 31 2005. Comparison with external data sets shows that the new data collected since April 3 2004 show a standard deviation of 1.3~\\kms, about twice better than for the first data release. For the first time this data release contains values of stellar parameters from 22,407 spectra of 21,121 individual stars. They were derived by a penalized $\\chi^2$ method using an extensive grid of synthetic spectra calculated from the latest version of Kurucz stellar atmosphere models. From comparison with external data sets, our conservative estimates of errors of the stellar parameters for a spectrum with an average signal to noise ratio of $\\sim\\!\\!40$ are 400~K in temperature, 0.5~dex in gravity, and 0.2 dex in metallicity. We note however that, for all three stellar parameters, the internal errors estimated from repeat RAVE observations of 822 stars are at least a factor 2 smaller. We demonstrate that the results show no systematic offsets if compared to values derived from photometry or complementary spectroscopic analyses. The data release includes proper motions from Starnet2, Tycho2, and UCAC2 catalogs and photometric measurements from Tycho-2 USNO-B, DENIS  and 2MASS.  The data release can be accessed via the RAVE webpage: http://www.rave-survey.org and through CDS. ",
        "introduction": "\\label{s:introduction}  This paper presents the second data release from the Radial Velocity Experiment (RAVE), an ambitious spectroscopic survey of the southern sky which has already observed over 200,000 stars away from the plane of the Milky Way ($|b| > 25\\deg$) and with apparent magnitudes $9<I_\\mathrm{DENIS}<13$. The paper follows the first data release, described in \\citet{dr1}, hereafter Paper~I. It doubles the number of published radial velocities. For the first time it also uses spectroscopic analysis to provide information on values of stellar parameters: temperature, gravity, and metallicity. Note that the latter in general differs from iron abundance, because metallicity is the proportion of matter made up of all chemical elements other than hydrogen and helium in the stellar atmosphere. Stellar parameters are given for the majority of the newly published stars. This information is supplemented by additional data from the literature: stellar position, proper motion, and photometric measurements from DENIS, 2MASS and Tycho surveys.  Scientific uses of such a data set were described in \\citet{steinmetz2003}. They include the identification and study of the current structure of the Galaxy and of remnants of its formation, recent accretion events, as well as discovery of individual peculiar objects and spectroscopic binary stars. Kinematic  information derived from the RAVE dataset has been used \\citep{escape_speed} to constrain the Galactic escape speed at the Solar radius to $v_\\mathrm{esc} = 536^{+58}_{-44} \\ \\kms$ (90 percent confidence). The fact that $v_\\mathrm{esc}^2$ is significantly greater than $2 v_\\mathrm{circ}^2$ (where $v_\\mathrm{circ} = 220~\\kms$ is the local circular velocity) is a model-independent confirmation that there must be a significant amount of mass exterior to the Solar circle, i.e.\\ it convincingly demonstrates the presence of a dark halo in the Galaxy. A model-dependent estimate yields the virial mass of the Galaxy of $1.31^{+0.97}_{-0.49} \\times 10^{12}$~M$_\\odot$ and the virial radius of $297^{+60}_{-44}$~kpc (90 per cent confidence). \\citet{veltz} discussed kinematics towards the Galactic poles and identified discontinuities that separate thin disk, thick disk and a hotter component. \\citet{seabroke2008} searched for in-falling stellar streams on to the local Milky Way disc and found that it is devoid of any vertically coherent streams containing hundreds of stars. The passage of the disrupting Sagittarius dwarf galaxy leading tidal stream through the Solar neighborhood is therefore ruled out. Additional ongoing studies have been listed in Paper~I.  The structure of this paper is as follows: Section 2 is a description of the observations, which is followed by a section on data reduction and processing. Data quality is discussed in Section 4, with a particular emphasis on a comparison of the derived values of stellar parameters with results from an analysis of external data sets. Section 5 is a presentation of the data product, followed by concluding remarks on the results in the context of current large spectroscopic surveys. ",
        "conclusions": "This second data release reports radial velocities of 51,829 spectra of 49,327 different stars, randomly selected in the magnitude range of $9 \\simlt I \\simlt 12$ and located more than $25^\\mathrm{o}$ away from the Galactic plane (except for a few test observations). It covers an area of $\\sim 7,200$ square degrees. These numbers approximately double the sample reported in Paper~I. Moreover, this data release is the first to include values of stellar parameters as determined from stellar spectra. We report temperature, gravity and metallicity for 21,121 normal stars, all observed after the first data release. Stars with a high rotational velocity or peculiar type (e.g.\\ binary stars and emission stars) will be discussed separately.  Radial velocities for stars new in this data release are more accurate than before, with typical errors between 1.3 and 1.7~\\kms. These values are confirmed both by repeated observations and by external datasets and have only a weak dependence on the S/N ratio. We used five separate external datasets to check values of stellar parameters derived from RAVE spectra. These included observations with different instruments at different resolving powers and in different wavelength regimes, as well as data from the literature. The uncertainty of stellar parameter values strongly depends on stellar type. Despite considerable effort our calibration observations do not cover (yet) the entire parameter space. We plan to improve on this using dedicated calibration observations with at least 4 telescopes. For this data release we had to resort to extensive simulations which are however tuned by calibration observations. A typical RAVE star has an uncertainty of 400~K in temperature, 0.5~dex in gravity, and 0.2~dex in metallicity. The error depends on the signal to noise ratio and can be $\\simgt 2$ times better/worse for stars at extremes of the noise range. Repeated observations show that these error estimates are rather conservative, possibly due to intrinsic variability of the observed stars and/or non-negligible errors of reference values from the calibration datasets.  Future data releases will follow on an approximately yearly basis. They will benefit from our considerable and ongoing effort to obtain calibration datasets using other telescopes and similar or complimentary observing techniques. Notably we expect that SkyMapper \\citep{skymapper}, an all-southern-sky survey just starting at the Siding Spring Observatory, will provide accurate photometry and temporal variability information for all RAVE stars.  RAVE is planned to observe up to a million spectra of stars away from the Galactic plane. It represents an unprecedented sample of stellar kinematics and physical properties in the range of magnitudes probing scales between the very local surveys (Geneva Copenhagen Survey and \\citet{famaey05}) and more distant ones (SDSSII/SEGUE), complementing the planned AAOmega efforts closer to the Galactic plane. So it helps to complete our picture of the Milky Way, paving the way for the next decade endeavors, like Gaia."
    },
    "0806/0806.4545_arXiv.txt": {
        "abstract": "{We report CO observations towards MS0735.6+7421 a distant cooling core galaxy, and towards M87, the nearest cooling core in the center of the Virgo cluster.  Both galaxies contain radio cavities that are thought to be responsible for the heating that can regulate or stop the cooling of the surrounding gas. In this feedback process, there could still be some gas cooling along filaments, along the borders of the radio cavities.  Molecular gas is known to exist in clusters with cooling cores, in long and thin filaments that can be formed behind the rising bubbles inflated by the central AGN.  CO emission was searched for at several locations along the radio lobes of those two galaxies, but only upper limits were found. These correspond to cold gas mass limits of a few 10$^9$M$_\\odot$ for each pointing in MS0735.6+7421, and a few 10$^6$M$_\\odot$ in M87.  This non detection means that either the cooling is strongly reduced by the AGN feedback or that the gas is cooling in very localized places like thin filaments, possibly diluted in the large beam for MS0735.6+7421. For M87, the AGN heating appears to have stopped the cooling completely. ",
        "introduction": "Giant ellipticals in galaxy cluster centers live in a complex environment. In some cluster cores the intra-cluster medium is known to cool down, but how much of this gas is able to fuel the central galaxy is still not clear.  The hot gas is no longer thought to be flowing directly towards the central galaxy.  The regulation of the intra-cluster medium (ICM) cooling involves some re-heating mechanisms like jet feedback from a central AGN and it becomes observable by the formation of X-ray emitting gas cavities. This feedback mechanism produces mass outflows, weak shocks and heating \\citep{McNa2007}, but also cooling, in long and thin filaments.  Indeed, some very cold gas (10-100 K) has been observed in CO emission in many cooling core systems \\citep{Edge01, Salo03}.  High resolution imaging of CO showed the existence of molecular filaments in Abell 1795 \\citep{Salo04} that are confined to the borders of the radio cavities. These filaments extend even farther outward in NGC~1275 \\citep{Salo08}. The cold gas formation may result from the interaction of the AGN with the surrounding hot gas. The filaments are thought to be cooling filaments that can fuel the giant elliptical with cooled gas.  More examples of this phenomenon need to be observed to better understand the physics of the formation of the cold gas. In this work, we report a search towards two cooling core systems: MS0735.6+7421 and M87 in the Virgo cluster. These two objects are known to be the sites of very strong feedback. The cooling in M0735.6+7421 very likely occurs only in the very central region and M87 is a poor cooling core with almost no cooling, so we do not expect to find much cooled gas.   The giant galaxy MS0735.6+7421 is at a distance of 716 Mpc (1$''$ correponds to 3.47 kpc). In spite of its large distance, its giant X-ray cavities are so large, inflated by radio lobes of size 550 kpc, that they are resolved by the beam of the IRAM-30m telescope.  It is then possible to probe the cold gas distribution and compare it with the X-ray, H$\\alpha$ and radio maps.  The galaxy MS0735.6+7421 hosts a radio source 4C+74.13. According to \\citep{McNa05} the AGN activity could have re-heated sufficiently the ICM to stop the cooling around the central galaxy. Figure \\ref{over} shows that the radio lobe emission of 4C+74.13 fills the X-ray cavities of the intra-cluster medium. These cavities may have displaced and compressed the hot gas surrounding the central cluster galaxy, as predicted by feedback models of cooling flows. Negative feedback models as described by \\citet{Revaz08} predict that the hot X-ray gas when compressed may eventually cool very efficiently down to very low temperatures, forming cold gas filaments visible in CO(1--0) emission. So in this galaxy, we searched for molecular gas emission not only in the central region but also far away from the X-ray cooling core region, along the AGN radio lobes.   M87 is the most nearby X-ray bright cooling core. It lies at a distance of 18 Mpc.  X-ray mapping around M87 reveals a complex structure for the cooling gas, perturbed by several bubbles (Young et al 2002).  The X-ray-cool component detected in the radio lobes is in blobs that are connected to the radio buoyant bubbles \\citep{Ghizzardi2004}.  In a deeper XMM-Newton observation of M87, \\citet{Simio2007} find evidence that cold, metal-rich gas is being transported outwards also outside the X-ray arms, possibly through bubble-induced mixing.  Many X-ray features are better be understood if there are repetitive AGN outbursts.  Depressions appear as the remnants of earlier outbursts, while shocks are seen as brightenings within the prominent X-ray arms.  As in Perseus, shocks may be the most significant channel for AGN energy input into the cooling-flow atmospheres, and outbursts every 3$\\times$10$^7$ yr are sufficient to quench the cooling \\citep{Forman2005}.    \\begin{figure}[htbp] \\centering \\epsfig{file=262fig1.eps, width=7.5cm}% \\caption{{A 250 $\\times$ 250$''$ image of the MS0735.6+7421 cluster center \\citep{McNa05}, one arcsec is 3.5 kpc. The 1.4 GHz radio contours of the jets emitted by 4C+74.13 are superposed. The cavities in the X-ray emission coincide with the radio lobe positions (200 kpc in size at the cluster redshift). The circles represent the 30m telescope beam at 3mm, showing the 7 regions we observed in CO(1--0).  }} \\label{over} \\end{figure} ",
        "conclusions": "We observed 7 different positions along the radio lobes of MS0735.6+7421 and 4 positions for M87, in CO(1--0) with the 30m telescope. We do not detect cold molecular gas in any of the observed regions.  Our upper limits put significant constraints on the cold gas masses. We find M$_{\\rm gas}\\,\\le$ a few 10$^9$M$_\\odot$ in MS0735.6+7421 and M$_{\\rm gas}\\,\\le$ a few 10$^6$M$_\\odot$ in M87. The absence of cold gas could be explained if strong feedback heats sufficiently the intra-cluster medium and prevents any cooling. Nevertheless, as suggested by the recent results of high resolution CO mapping of NGC~1275 \\citep{Salo08}, if the cooling happens only in very thin and elongated filaments, we may have missed such emission in MS0735.6+7421. The spatial and spectral dilution of a collection of CO filaments would make the detection difficult in a single beam of 28$''$ (97 kpc) at this distance. More sensitive observations, with higher spatial resolution are needed to determine whether this is the case.  \\begin{figure}[htbp] \\centering \\includegraphics[width=8.1cm, angle=0]{262fig3.ps} \\caption{MS0735.6+7421 - CO(1--0) spectra obtained at all the positions observed as indicated at upper right in each diagram and as overlaid on Fig. \\ref{over}. The channel width is 50.6 km\\,s$^{-1}$. See Table \\ref{table-upper} for a summary. The Y-axis is in T$_{\\rm mb}$ (mK).} \\label{detection} \\end{figure}  \\begin{figure}[htbp] \\centering \\includegraphics[width=8.1cm, angle=0]{262fig4.ps} \\caption{M87 - CO(1--0) and CO(2-1) spectra obtained at all the positions observed as indicated at upper right in each diagram and as overlaid on Fig. \\ref{over-m87}. The channel width is 41.8 km\\,s$^{-1}$. See Table \\ref{table-upper-m87} for a summary. The Y-axis is in T$_{\\rm mb}$ (mK).} \\label{detection} \\end{figure}"
    },
    "0806/0806.1967_arXiv.txt": {
        "abstract": "{} {While weak lensing cannot resolve cluster cores and strong lensing is almost insensitive to density profiles outside the scale radius, combinations of both effects promise to constrain density profiles of galaxy clusters well, and thus to allow testing of the CDM expectation on dark-matter halo density profiles.} {We develop an algorithm further that we had recently proposed for this purpose. It recovers a lensing potential optimally reproducing observations of both strong and weak-lensing effects by combining high resolution in cluster cores with the larger-scale information from weak lensing. The main extensions concern the accommodation of mild non-linearity in  inner iterations, the progressive increase in resolution in outer iterations, and the introduction of a suitable regularisation term. The linearity of the method is essentially preserved.} {We demonstrate the success of the algorithm with both idealised and realistic simulated data, showing that the simulated lensing mass distribution and its density profile are well reproduced. We then apply it to weak and strong lensing data of the cluster MS~2137 and obtain a parameter-free solution which is in good qualitative agreement with earlier parametric studies.}{} ",
        "introduction": "\\label{Introduction}  The mass distribution in dark-matter halos and the level of substructure in them are among the central predictions of the CDM paradigm for cosmic structure formation. The density profile should asymptotically fall off $\\propto r^{-3}$ at large radii $r$ and flatten considerably within a radial scale $r_\\mathrm{s}$ \\citep{Navarro1996}.The mass distribution should be richly substructured by sublumps of matter with a differential mass function approximated by a power law, $d n/d M\\propto M^{\\alpha}$ with a slope slighly shallower than $\\alpha=-2$ \\citep{Madau2008,Springel2008}.  Galaxy clusters should be weakly influenced by baryonic physics, thus their density profiles and mass distributions outside the cooling radius should well reflect those expected for dark matter. Do they? Although tentative answers exist, showing that estimated density profiles do at least not contradict the CDM expectation, accurate constraints are still missing. Due to its insensitivity to the physical state of the matter, gravitational lensing is perhaps the most promising tool for determining matter distributions.  Weak lensing lacks the resolution necessary to constrain the density profile in cluster centres, while strong lensing is confined to the innermost cluster cores. In combination, they may be able to test the CDM predictions on density profiles well.  Several methods have been suggested to combine weak and strong cluster lensing \\citep{Bradav2005,Cacciato2006,Diego2007}. Among them is our own algorithm aiming at the lensing potential. It is based on minimising a $\\chi^2$ function comparing observed shear measurements with suitable second derivatives of the potential. Expressing the derivatives in terms of finite differences leads to a system of linear equations whose direct inversion yields the solution.  We extend our earlier work in several ways. First, we no longer use the lowest-order approximation in which measured ellipticities estimate the shear, and introduce the reduced shear instead. The non-linearity accommodated in this way can be resolved into an iterative scheme using linear inversion in each step. Second, we wrap the algorithm into an outer iteration loop in which the grid resolution is progressively enhanced. While this step introduces correlations between adjacent pixels that have to be dealt with, it prepares the insertion of the strong-lensing constraints available in cluster cores. Third, we introduce a regularisation term for the two purposes of avoiding overfitting and smoothly joining the strong- and weak-lensing solutions. Finally, to account for the additional computational time we enabled the code to run on parallel machines.  We investigate the performance of our algorithm using two sets of synthetic data, one idealised and one realistic, before we proceed to apply it to the well-known strong-lensing cluster MS~2137, for which we obtain a high-resolution, parameter-free reconstruction.  A brief summary of the lensing notation in Sect.~\\ref{lensingformalism} is followed by an outline of the method in Sect.~\\ref{outlineofthemethod} and a description of its implementation in Sect.~\\ref{implementation}. We present the results in Sect.~\\ref{results} and conclude in Sect.~\\ref{conclusions}. Details of the algorithm are given in Appendix~\\ref{linearisationingridspace}. ",
        "conclusions": "\\label{conclusions}  We have extended the algorithm proposed by \\citet{Cacciato2006} which reconstructs the lensing potential of galaxy clusters on a two-level grid, combining weak and strong-lensing data. Our extensions concern three aspects: First, the shear is replaced by the reduced shear to improve the reconstruction in the mildly non-linear regime. The non-linearity introduced in this way is solved by an inner iteration loop which preserves the linear minimisation of the $\\chi^2$ function. Second, the spatial resolution of the potential grid is gradually increased in an outer iteration loop. While this step causes correlations between neighbouring pixels, which need to be dealt with using a non-diagonal covariance matrix, it prepares for the introduction of a highly resolved grid covering the strong-lensing observations in the cluster core. Third, we add a regularisation term to the $\\chi^2$ function, avoiding overfitting of noise and allowing a smooth transition from the inner, high-resolution to the outer, coarse-resolution grid. \\\\ These extensions of the algorithm, especially the introduction of the full $\\chi^{2}$-function and the increased number of iterations made it necessary to drastically increase the numerical performance of the reconstruction routines. This was done with fast, numerical multiplication schemes and the parallelisation of the code to run on MPI machines. Applications of this algorithm to synthetic data show that the known, simulated cluster lenses are accurately reproduced, although intrinsic ellipticities, measurement and shot noise inevitably lead to a considerable noise level of the reconstruction in the case of realistic data. Still, the quantitative analysis of our convergence maps shows a very good agreement with the expected result. An application to the galaxy cluster MS~2137 qualitatively confirms the results of earlier, parametric studies that show an almost axially symmetric, presumably relaxed mass distribution. Although there is still room for potentially important further improvements, our algorithm should now be ready for reliable recoveries of lensing potentials, density distributions and density profiles of real galaxy clusters."
    },
    "0806/0806.1226_arXiv.txt": {
        "abstract": "The {\\it Swift} satellite early x-ray data shows a very steep decay in most of the Gamma-Ray Bursts light curves. This decay is either produced by the rapidly declining continuation of the central engine activity or by some left-over radiation starting right after the central engine shuts off. The latter scenario consists of the emission from an ``ember'' that cools via adiabatic expansion and, if the jet angle is larger than the inverse of the source Lorentz factor, the large angle emission.  In this work, we calculate the temporal and spectral properties of the emission from such a cooling ember, providing a new treatment for the micro-physics of the adiabatic expansion.  We use the adiabatic invariance of $p_{\\perp}^2/B$ ($p_{\\perp}$ is the component of the electrons' momentum normal to the magnetic field, $B$) to calculate the electrons' Lorentz factor during the adiabatic expansion; the electron momentum becomes more and more aligned with the local magnetic field as the expansion develops. We compare the theoretical expectations of the adiabatic expansion (and the large angle emission) with the current observations of the early x-ray data and find that only $\\sim 20\\%$ of our sample of 107 bursts is potentially consistent with this model. This leads us to believe that, for most bursts, the central engine does not turn off completely during the steep decay of the x-ray light curve; therefore, this phase is produced by the continued rapidly declining activity of the central engine. ",
        "introduction": "The central engine of the GRBs (Gamma-Ray Bursts - see Piran 2005 for a review) is hidden to direct observations and its workings are largely unknown. The only information that we currently have about the GRB is obtained from its electromagnetic radiation.  We have to look for signatures in the radiation mechanism to understand how Nature produces these outbursts.  {\\it Swift} has provided very early x-ray data that shows that for most bursts there is a very steep decay lasting for about 10 minutes (Tagliaferri et al. 2005,  Nousek et al. 2006 - see Zhang et al. 2006 and references therein for possible physical explanations). These observations suggest that the rapidly declining x-ray light curve (LC) and the burst are produced by the same source, because the x-ray LC, when extrapolated backwards in time, matches the gamma-ray LC (O'Brien et al. 2006). Therefore, a natural question arises: Is the early x-ray data really just a rapidly declining continuation of the central engine activity, originally seen in the gamma-ray band, but now seen at lower energy? Or does the central engine switch off abruptly and the early x-ray data doesn't reflect the activity of the central engine?  If the central engine completely shuts off when the gamma-ray photons' flux falls below the gamma-ray detector sensitivity, then the emission from a  cooling ``ember'' would be responsible for the early x-ray steep decay.  This source - which had just produced the gamma-ray emission - would be cooling by adiabatic expansion (AE). In this paper, we study the flux properties of a ``hot'' shell that undergoes AE and cools.  The goal is to determine if the observed x-ray steep decay data is consistent with the AE scenario.  If so, then the central engine did shut off abruptly right after the gamma-ray emission ceased.  On the other hand, if it isn't, then the data is reflecting the rapidly declining activity of the central engine.  The reason for this is that any other process that does not invoke central engine activity to explain the early x-ray data has problems explaining the smooth temporal connection observed in the LC between the prompt emission and the early x-ray data.  AE has been studied before to predict the long-wavelength afterglow from GRBs (M\\'esz\\'aros \\& Rees 1997) and also, the optical flashes from internal and reverse shocks (M\\'esz\\'aros \\& Rees 1999, Sari \\& Piran 1999a).  In this work, we describe the evolution of a collisionless plasma due to AE and show that this is in general different from AE of an ideal gas.  We first present the micro-physics of the AE for a collisionless plasma in \\S2. Then, we use it to calculate the flux properties of a source undergoing AE in \\S3.  We look at what the current observations tell us in \\S4, and we put them in the context of the central engine in \\S5.  We summarize our results and give our conclusions in \\S6. ",
        "conclusions": "The fastest way that the observed flux can decline after the central engine shuts off is set by the Large Angle Emission (LAE) and the Adiabatic Expansion cooling (depending on the value of $\\theta_j$).  The early x-ray steep decay shown in most of {\\it Swift} bursts has been attributed to LAE.  In this paper, we consider both AE and LAE for the very early x-ray data.  LAE and AE both start with the assumption that the central engine shuts off abruptly.  Only $\\sim 20\\%$ of our sample of 107 bursts is possibly consistent with either LAE or AE, thereby suggesting that the observed early x-ray steep decay for a large fraction of bursts might be produced by the rapidly declining continuation of the central engine activity.  The component of the electron's momentum parallel to the magnetic field is unconstrained by the adiabatic invariance. This component would probably cool via Inverse Compton scattering with synchrotron photons.  Another possibility is that the electrons are scattered by small scale fluctuations in the magnetic field, which would effectively couple the parallel and perpendicular components of their momentum, resulting in an adiabatic cooling similar to that in the ideal gas case, \\S2 (personal communication, Granot).  Any process that does not rely on the central engine activity to explain the observed early x-ray steep decay (i.e. RS, FS) has problems explaining the smooth temporal connection observed in the LC between the prompt emission and the early x-ray steep decay."
    },
    "0806/0806.4635_arXiv.txt": {
        "abstract": "We have used the Japanese VLBI array VERA to perform high-precision astrometry of an H$_2$O maser source in the Galactic star-forming region NGC~281~West, which has been considered to be part of a 300-pc superbubble. We successfully detected a trigonometric parallax of 0.355$\\pm$0.030~mas, corresponding to a source distance of 2.82$\\pm$0.24~kpc. Our direct distance determination of NGC~281 has resolved the large distance discrepancy between previous photometric and kinematic studies; likely NGC~281 is in the far side of the Perseus spiral arm. The source distance as well as the absolute proper motions were used to demonstrate the 3D structure and expansion of the NGC~281 superbubble, $\\sim$650~pc in size parallel to the Galactic disk and with a shape slightly elongated along the disk or spherical, but not vertically elongated, indicating the superbubble expansion may be confined to the disk. We estimate the expansion velocity of the superbubble as $\\sim 20$~km~s$^{-1}$ both perpendicular to and parallel to the Galactic disk with a consistent timescale of $\\sim 20$~Myr. ",
        "introduction": "A paradigm shift has been occurring in recent years in the observational study of the interstellar medium (ISM) in the Galaxy and in external galaxies, where the ISM is now recognized as being of fundamental importance not only in the life cycle of stars but also in galactic evolution. The present understanding of the ISM draws a dynamic and violent picture of the gas in the vast interstellar space, rather than a static one envisioned previously. Recent observations with higher angular resolution have revealed the existence of dynamic and diverse gas structures in the Galaxy that cannot be explained by the traditional evolutionary cycle of the ISM, requiring new insights toward a comprehensive understanding of the ISM (e.g., Reynolds 1997, 2002; Dickey 2001). Extreme activity of the ISM is indicated by large-scale circular or arc-like structures (up to a few kiloparsecs in scale) called superbubbles or supershells, many of which were identified earlier by Heiles (1979, 1984) based on H\\emissiontype{I} 21-cm line surveys by Weaver and Williams (1973) and Heiles and Habing (1974). Kiss, Mo\\'or, and T\\'oth (2004) later identified 145 large-scale loop- or arc-like structures in the 2nd Galactic quadrant (i.e.\\ for galactic longitudes $90^\\circ < l < 180^\\circ$) on {\\it IRAS}-based far-infrared maps. Recent results for the 1st, 3rd, and 4th Galactic quadrants by K\\\"onyves \\etal\\ (2007) increased the number of identified loops to 462. Similar objects are also found in nearby galaxies, including the Local Group spirals Andromeda (M31; Brinks \\& Bajaja 1986) and M33 (Deul \\& den Hartog 1990), Local Group irregulars including the Large Magellanic Cloud (LMC; Chu \\etal\\ 1993, 1995; Kim \\etal\\ 1998, 1999), the Small Magellanic Cloud (SMC; Staveley-Smith \\etal\\ 1997; Stanimirovi\\'c \\etal\\ 1999) and IC~10 (Shostak \\& Skillman 1989), large spiral galaxies such as M101 (Kamphuis \\etal\\ 1991) and NGC~891 (Howk \\& Savage 1997, 2000), and dwarf galaxies such as M82 (Matsushita \\etal\\ 2000; Matsumoto \\etal\\ 2001), IC~2574, DDO~47 (Walter \\& Brinks 1999, 2001) and Holmberg~II (HoII; Puche \\etal\\ 1992). \\begin{figure*}[t] \\begin{center} \\FigureFile(86mm,110mm){figure1.eps} \\end{center} \\caption{A wide-field optical image of the star-forming region NGC~281 taken with the WIYN 0.9-meter at Kitt Peak National Observatory (courtesy of T.A. Rector/University of Alaska Anchorage and WIYN/AURA/NSF). The orientation of the image is north up and east left. The red nebulosity of the H\\emissiontype{II} region NGC~281 is obscured in its southwestern quadrant by the adjoining molecular cloud NGC~281~West, where the H$_2$O maser source is located.}\\label{fig:optical} \\end{figure*}   The most widely accepted theory for the origin of superbubbles suggests that supernova explosions and strong stellar winds from OB associations in a galactic disk blow the surrounding gas up into the halo, carving out extensive cavities in the ISM (see Tenorio-Tagle \\& Bodenheimer 1988). Superbubbles are therefore believed to play a vital role in disk-halo interactions, serving as a conduit for matter and energy from the disk into the halo in the form of galactic {`}fountains{'} (Shapiro \\& Field 1976) or {`}chimneys{'} (Norman \\& Ikeuchi 1989).   In the Galaxy, only a small number of chimneys (or also called {`}worms{'}) or fragmenting superbubbles/supershells (on several hundred to one thousand parsec scales) are known and well-studied: e.g., the Orion-Eridanus superbubble (Cowie \\etal\\ 1979; Reynolds \\& Ogden 1979; Boumis \\etal\\ 2001; Welsh \\etal\\ 2005), the Cygnus superbubble (Cash \\etal\\ 1980), the Stockert chimney (M\\\"uller \\etal\\ 1987; Kundt \\& M\\\"uller 1987; Forbes 2000), the Aquila supershell (Maciejewski \\etal\\ 1996), the Scutum supershell (Callaway \\etal\\ 2000; Savage \\etal\\ 2001), GSH~277+00+36 (McClure-Griffiths \\etal\\ 2000), GSH~242-03+37 (McClure-Griffiths \\etal\\ 2006), the W4 chimney/superbubble (Normandeau \\etal\\ 1996; Reynolds \\etal\\ 2001; Terebey \\etal\\ 2003; Madsen \\etal\\ 2006; West \\etal\\ 2007), the Ophiuchus superbubble (Pidopryhora \\etal\\ 2007), and the NGC~281 superbubble (Megeath \\etal\\ 2002, 2003; Sato \\etal\\ 2007). Superbubbles are also a key target of research for understanding their impact on the history of star formation in the galactic disk, including but not limited to superbubble-triggered star formation (e.g., NGC~281 as explained later in the section).   In either the Galaxy or external galaxies, very little is known about the dynamics and energetics of superbubbles/chimneys or about their influence on the formation of the halo or on the star formation history in the galactic disk. A difficulty in revealing the structure and kinematics of superbubbles stems from the limited information obtainable from two dimensions of sky positions and velocity information only along the line of sight. A direct observation of the three-dimensional (3D) motion and structure of a superbubble has not been easy due to the high astrometric accuracy required. However, Sato \\etal\\ (2007) reported the first VLBI (Very Long Baseline Interferometry) observations that obtained direct evidence for expanding motions of a superbubble by measuring the proper motions.   In Sato \\etal\\ (2007), we observed the H$_2$O maser emission in NGC~281~West over 6 months with VERA (VLBI Exploration of Radio Astrometry), a new Japanese radio telescope array dedicated to phase-referencing VLBI astrometry (e.g., Honma \\etal\\ 2000, 2005, 2007; Kobayashi \\etal\\ 2003; Imai \\etal\\ 2006, 2007; Hirota \\etal\\ 2007, 2008a, 2008b; Nakagawa \\etal\\ 2008; Kim \\etal\\ 2008; Choi \\etal\\ 2008). These observations revealed the systemic motion of the NGC~281 superbubble away from the Galactic plane at a velocity of 20$-$30~km~s$^{-1}$. At its estimated heliocentric distance of 2.2$-$3.5~kpc, NGC~281 most likely originated in the Galactic plane and has been blown out by a superbubble expansion (see $\\S$1.1 and Sato \\etal\\ 2007).   A precise and direct determination of the distance to NGC~281 is desirable and essential for a better understanding of the region, including further description of the 3D structure and motion of the superbubble to reveal the origin, the energetics and the timescale of the superbubble. These are the primary reasons why we decided to complete our VERA observations to directly measure the parallactic distance to NGC~281 via an H$_2$O maser source. The astrometric determination of the distance to NGC~281 is also important for studying early-type (O$-$B) high-mass ($M\\gtrsim10\\Mo$) stars and stellar evolution in combination and in comparison with photometric studies of the OB stars in NGC~281 (e.g., Henning \\etal\\ 1994; Guetter \\& Turner 1997). In this paper, we present the results of our parallax measurements with VERA over 18 months, which enables us to further discuss the 3D structure, expanding motion, the origin and the timescale of the NGC~281 superbubble.  \\subsection{NGC~281 Superbubble} The NGC~281 region ($\\alpha_{2000}=$00$^{\\rm h}$52$^{\\rm m}$, $\\delta_{2000}=+$56$^\\circ$\\timeform{34'} or $l=$123$^\\circ$.07, $b=-$6.$^\\circ$31) provides an excellent laboratory for studying in detail the cycle of the ISM and its impact on both star formation and Galactic disk-halo interaction through superbubbles. At its estimated heliocentric distance of 3~kpc, this region is remarkably located far above ($\\sim 300$~pc) the midplane of the Perseus arm of the Galaxy. Of special interest in the NGC~281 region is a possibility of triggered star formation occurring on two different scales: large-scale ($\\sim$300~pc) supperbubble-triggered formation of the first OB stars (Megeath \\etal\\ 2002, 2003), and the sequential and ongoing triggered star formation ($\\sim$1$-$10~pc) in an adjoining giant molecular cloud (NGC~281~West) through interaction with an H\\emissiontype{II} region (the NGC~281 nebula) excited by the first OB stars (Elmegreen \\& Lada 1978; Megeath \\& Wilson 1997).   Figure~\\ref{fig:optical} is an optical image of the region, which clearly shows the red nebulosity of the H\\emissiontype{II} region NGC~281 (also known as Sharpless 184) of diameter 20~arcmin ($\\sim 17$~pc), in which an early-type cluster called IC~1590 is embedded (Guetter \\& Turner 1997). The brightest member of the cluster IC~1590 is an O-type Trapezium-like system HD~5005 (or also called ADS~719), whose component stars HD~5005ab (unresolved), HD~5005c, and HD~5005d have spectral types of O6.5~V, O8~V, and O9~V, respectively (Walborn 1973; Abt 1986; Guetter \\& Turner 1997).   The southwestern quadrant of the NGC~281 nebula is obscured (as seen in figure~\\ref{fig:optical}) by the adjoining molecular cloud NGC~281~West in front of the nebula. Ongoing star formation in the NGC~281~West cloud is indicated by the H$_2$O maser emission and {\\it IRAS} sources within this cloud near the clumpy interface between the H\\emissiontype{II} region and the cloud. This generation of stars may have been triggered by interactions with the H\\emissiontype{II} region (Elmegreen \\& Lada 1978; Megeath \\& Wilson 1997). The NGC~281 molecular cloud complex that surrounds the H\\emissiontype{II} region (including the NGC~281~West and East clouds) was mapped in CO emission lines by Lee and Jung (2003). The central radial velocity of the NGC~281~West molecular cloud, $V_{\\rm{LSR}}=31$~km~s$^{-1}$ (Lee and Jung 2003), agrees well with that of the H$_2$O maser emission in the cloud (Sato \\etal\\ 2007). On a larger scale, Megeath \\etal\\ (2002, 2003) identified the CO molecular cloud complex as being on an H\\emissiontype{I} loop extending over 300~pc away from the Galactic plane toward decreasing galactic latitude. In an $l$ vs.\\ $V_{\\rm{LSR}}$ diagram (i.e.\\ a plot of galactic longitude vs.\\ observed radial velocity), the CO molecular clouds appear to be part of a broken ring of diameter 270~pc, expanding at a velocity of 22~km~s$^{-1}$ parallel to the Galactic plane (Megeath \\etal\\ 2002, 2003). Megeath \\etal\\ (2002, 2003) suggested that these clouds were formed in a fragmenting superbubble, which triggered the formation of the first OB stars, and that these OB stars then ionized the surrounding gas which subsequently triggered ongoing star formation in the neighboring cloud.   In Sato \\etal\\ (2007), we derived the absolute proper motions of the H$_2$O maser features in NGC~281~West by phase-referencing VLBI observations with VERA over 6 months. We found 10 maser features, highly likely excited by two spatially distinct young stellar objects (YSOs), are systematically moving toward southwest, in agreement with the direction of the proper motion of HD~5005 in the H\\emissiontype{II} region measured with larger uncertainties by {\\it Hipparcos} (Perryman \\etal\\ 1995, 1997). Our results with VERA revealed the systemic motion of the NGC~281 region away from the Galactic plane, most likely having originated in the Galactic plane and been blown out by the expansion of the NGC~281 fragmenting superbubble.   The uncertainties in our velocity estimates are largely due to the large distance uncertainty of NGC~281, despite many measurements, e.g., photometric distances of 2~kpc (Sharpless 1952, 1954), 3.68~kpc (Cruz-Gonz\\'alez \\etal\\ 1974), 3.5~kpc (Henning \\etal\\ 1994) and 2.94$\\pm$0.15~kpc (Guetter \\& Turner 1997) and kinematic distances of 2.2~kpc (Georgelin \\& Georgelin 1976; Roger \\& Pedlar 1981) and 3~kpc (Lee \\& Jung 2003). In Sato \\etal\\ (2007), we adopted a value of 2.9~kpc as the distance to NGC~281, based on a recent estimate of 2.94$\\pm$0.15~kpc derived by Guetter and Turner (1997) through photometry of 279 individual stars in and about the young cluster IC~1590 embedded in the NGC~281 nebula.   The most reliable and direct way to determine the distance ($d$) of an astronomical object is to measure the trigonometric parallax ($\\pi$) of the object due to the Earth's motion around the solar system barycenter, where the distance is given by $d$(pc) $=1/\\pi$(\\timeform{\"}). The astrometric determination of the distance to NGC~281 provides valuable information for studying the formation and evolution of early-type (O$-$B) high-mass ($M\\gtrsim10\\Mo$) stars in comparison with photometric studies and distances, for example, by Guetter \\& Turner (1997).   In the present study, we report on our successful determination of the parallactic distance to NGC~281, which reveals the 3D structure of the NGC~281 fragmenting superbubble, and investigate the origin and timescale of the formation of such a large-scale ($\\sim$650~pc) structure. ",
        "conclusions": "In this paper, we have presented the results of our multi-epoch phase-referencing observations with VERA over 18 months of an H$_2$O maser source in the Galactic star-forming region NGC~281~West, associated with a fragmenting superbubble $\\sim300$~pc above the Perseus spiral arm. The primary results are summarized as follows:  \\begin{enumerate} \\item We detected a trigonometric parallax of 0.355$\\pm$0.030~mas, corresponding to a distance of 2.82$\\pm$0.24~kpc to NGC~281. Our parallactic distance agrees well with the photometric distance of 2.94$\\pm$0.15~kpc derived by Guetter \\& Turner (1997), allowing for improved study of the absolute magnitutudes of high-mass OB stars, and resolved the large distance discrepancy of NGC~281 from previous photometric and kinematic studies.\\\\  \\item We revised the absolute proper motions of the H$_2$O maser features measured in Sato \\etal\\ (2007), and using our new determination of the parallactic distance, we derived more precisely the velocity component of the NGC~281 region $v_b\\simeq 18$~km~s$^{-1}$ perpendicular to the Galactic plane. This yields a timescale of $t_b\\sim 17$~Myr and kinectic energy of the region to be $3.0\\times10^{51}$~ergs in the direction perpendicular to the plane.\\\\  \\item We demonstrated the 3D structure and expansion of the NGC~281 superbubbles in comparison of our parallactic distance to NGC~281 with the parallactic distance derived by Moellenbrock \\etal\\ (2007) for another H$_2$O maser source, IRAS~00420+5530 in the superbubble. Our new parallactic distance revealed the structure of the superbubble $\\sim$650 pc in size parallel to the Galactic disk and with a shape slightly elongated along the disk or spherical, but not vertically elongated in the $z$ direction. Therefore, the superbubble expansion may be confined to the disk, possibly due to the magnetic field of the disk. \\\\  \\item We suggested a new possible interpretation of an expanding molecular ring parallel to the Galactic plane in the longitude-velocity diagram by Megeath \\etal\\ (2002, 2003). In either interpretation, the ring center, i.e.\\ the likely origin of the superbubble lies at $l\\simeq 121^\\circ$. The velocity deviation of the superbubble from the Galactic rotation is estimated to be $\\sim 10$~km~s$^{-1}$ inward the Galactic rotation and toward the Sun from our proper motion measurements and from the longitude-velocity diagram, which might be due to velocity jumps over a spiral shock in the Perseus spiral arm.\\\\  \\item We also estimated the velocity component and timescale of the ring expansion to be $v_\\mathrm{ring2}\\simeq 15-20$~km~s$^{-1}$ and $t_\\mathrm{ring2}\\sim 16-21$~Myr, respectively, parallel to the Galactic plane, which are in good agreement with those independently estimated for the direction perpendicular to the Galactic plane, $v_b\\simeq 18$~km~s$^{-1}$ and $t_b\\sim 17$~Myr. The kinetic energy estimate of the region parallel to the plane is $2.1\\times10^{51}$~ergs, and the total kinetic energy of both perpendicular and parallel to the plane is then estimated to be $5.1\\times10^{51}$~ergs, requiring multiple supernovae. \\end{enumerate}  \\medskip \\noindent We are deeply grateful to the referee Dr.~Felix~J.~Lockman for his invaluable comments and suggestions that improved both scientific discussions and presentation of the manuscript. We would like to thank Dr.~Mark~J.~Reid for his tremendous help with carefully reading and improving the manuscript and for his insightful comments on scientific issues. We are very grateful to Prof.~Kohji~Tomisaka for illuminating and insightful discussions. We also greatly appreciate the optical image for figure~1 that was kindly provided by Dr.~Travis~A.~Rector and the National Optical Astronomy Observatory (NOAO). We wish to thank all the support staff at VERA for their efforts and continuous support for our observations. M.~Sato gratefully acknowledges financial support by a Research Fellowship from the Japan Society for the Promotion of Science (JSPS) for Young Scientists. This research has made use of the SIMBAD database, operated at CDS, Strasbourg, France and of NASA's Astrophysics Data System."
    },
    "0806/0806.2822_arXiv.txt": {
        "abstract": "The discovery that the expansion of the Universe is accelerating is the most challenging problem of modern cosmology. In the context of general relativity, there are many dark energy candidates to explain the observed acceleration. In this work we focus our attention on two kinds of simplified Chaplygin gas cosmological accelerating  models recently  proposed in the literature. In the first scenario, the simplified Chaplygin gas works like a Quintessence model while in the second one, it plays the role of a Quartessence (an unification of the dark sector). Firstly, in order to limit the free parameters of both models, we  discuss the age of high redshift objects with special emphasis to the old quasar APM 08279+5255 at $z = 3.91$. The basic finding is that this old high redshift object constrain severely the simplified Chaplygin cosmologies. Secondly, through a joint analysis involving the baryon acoustic oscillations (BAO) and a sample of old high redshift galaxies (OHRGs) we also estimate the value of the Hubble parameter, $H_0$. Our approach suggests that the combination of these two independent phenomena provides an interesting method to constrain the Hubble constant. ",
        "introduction": "\\hspace{0.5cm}  An impressive amount of different astrophysical data are suggesting that the observed Universe can be represented  by an accelerating  spatially flat cosmology driven by an exotic component sometimes called dark energy \\cite{Riess,Sperg07}. In addition to the cosmological constant, the most interesting candidates for describing this dark energy component are: a vacuum decaying energy density, or a time varying $\\Lambda(t)$ \\cite{OzTa87}, the so-called ``X-matter\" \\cite{turner97}, a relic scalar field rolling down its potential \\cite{PR03}, and a Chaplygin Gas \\cite{kamen}. Some recent review articles discussing the history, interpretations, as well as the major difficulties of such candidates have also been published in the last few years \\cite{review}.  On the other hand, some studies have pointed out that the age of old high redshift objects (at moderate and high redshifts) is a powerful physical constraint to cosmological models \\cite{old,Krauss97,HasingerEtAl02,LA,Age}. For old high redshift galaxies (OHRGs), for instance, their ages  can be inferred by measuring the magnitude in different bands and then using stellar evolutionary codes to choose the model that reproduces the observed colors \\cite{old,Krauss97}. In principle, for old high redshift objects, the derived ages can put tight constraints on the physical parameters or even rule out many kinds of dark energy models.  In this article, we discuss some  observational constraints on two different classes of Chaplygin type accelerating cosmologies coming  from  the existence of the  APM 08279+5255, an old quasar at high redshift located  at $z = 3.91$ which has an estimated age of 2-3 Gyr \\cite{HasingerEtAl02}. As reported by  many authors \\cite{LA}, the existence of this object is not compatible with many dark energy models unless the present accepted values  of $\\Omega_M$ and ${H}_0$ be further revised. We also derive new constraints on the Hubble constant $H_0$, by using a joint analysis involving different OHRGs samples and the current SDSS measurements of the baryon acoustic peak \\cite{Eisenst05}. The BAO signature comes out because the cosmological perturbations excite sound waves in the relativistic plasma, thereby producing the acoustic peaks in the early universe. Actually, Eisenstein {\\it et al.} \\cite{Eisenst05} presented the large scale correlation function from the Sloan Digital Sky Survey (SDSS) showing clear evidence for the baryon acoustic peak at $100h^{-1}$ Mpc scale, which is in excellent agreement with the WMAP prediction from CMB data. We stress that the Baryon Acoustic Oscillations (BAO) method is independent of the Hubble constant $H_0$ although being heavily dependent on the value of $\\Omega_M$. In this way, it contributes indirectly to fix the value of $H_0$ since it breaks the degeneracy on the mass density parameter, $\\Omega_M$ \\cite{CLM07}. As we shall see, both set of data (OHRGs sample and BAO signature) provide an interesting method to constrain the Hubble constant. ",
        "conclusions": "It is usually believed that old high-redshift objects may play an important role to the question related to the ultimate fate of the Universe. In a point of fact, their age estimates provide a powerful technique for constraining the free parameters in a given cosmological model \\cite{Krauss97,LA}. In particular, this means that the so-called high-redshift  age crisis is now becoming an important complement to other independent cosmological tests.  In this work we have discussed some constraints in the so-called Simplified Chaplygin Gas (SC-gas) both in the Quintessence and Quartessence versions. Initially, we have investigated the dimensionless age parameter $T_z = H_0 t_z$. In the simplified Quintessence  scenario, by assuming that the quasar has an age of at least a 2 Gyr,  we have obtained  $\\alpha\\geq 0.947$ while for the simplified Quartessence model the limit is $\\alpha>0.965$. We have also studied the constraints on the Hubble constant using a joint analysis involving  different OHRGs samples and the current SDSS measurements of the baryon acoustic peak for both scenarios. In this case, we have found  values of the Hubble constant moderately low, specially, to the  Quintessence case (see Fig. 3 and Table 1).  As a conclusion, we would like to stress that our results point to at least 3 possibilities. The first one is that the simplified Chaplygin type gas models (Quintessence and Quartessence) do not provide a good fit to the Age+BAO data because the best fit predicted for $H_0$ is low. A second possibility is that the estimated ages of the old high redshift galaxies considered here (or even their incubation time) need to be revised, and this may increase the predicted value of the Hubble constant. Finally, as has been advocated by some authors \\cite{Sand06,Blanch}, there exist the possibility that we live in a Universe with Hubble constant smaller than the one predicted by the HST Key Project \\cite{Freedman01}."
    },
    "0806/0806.2649_arXiv.txt": {
        "abstract": "Nailing down the unknown neutrino mixing angle $\\theta_{13}$ is one of the most important goals in current lepton physics. In this context, we perform a global analysis of neutrino oscillation data, focusing on $\\theta_{13}$, and including recent results  [{\\em Neutrino 2008, Proceedings of the XXIII International Conference on Neutrino Physics and Astrophysics, Christchurch, New Zealand, 2008\\/} (unpublished)]. We discuss two converging hints of $\\theta_{13}>0$, each at the level of $\\sim\\!\\! 1\\sigma$: an older one coming from atmospheric neutrino data, and a newer one coming from the combination of solar and long-baseline reactor neutrino data. Their combination provides the global estimate $$\\sin^2\\theta_{13}=0.016\\pm 0.010 \\ (1\\sigma)\\,,$$ implying a preference for $\\theta_{13}>0$ with non-negligible statistical significance ($\\sim\\!\\!90\\%$ C.L.). We discuss possible refinements of the experimental data analyses, which might sharpen such intriguing indication. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0806/0806.0199_arXiv.txt": {
        "abstract": "A two-dimensional (2D) shock-rest-frame model for particle simulations is developed. Then full kinetic dynamics of a perpendicular collisionless shock is examined by means of a 2D full particle simulation. We found that in the 2D simulation there are fewer  nonthermal electrons due to surfing acceleration which was seen in the previous 1D simulations of a high Mach number perpendicular shock in a low-beta and weakly magnetized plasma. This is because the particle motion along the ambient magnetic field disturbs the formation of coherent electrostatic solitary structures which is necessary for electron surfing acceleration. ",
        "introduction": "Observations of shell-type supernova remnants (SNRs) provide us evidence for electron acceleration at SNR shocks, and in some cases the maximum energy of the electrons exceeds $\\sim$10~TeV \\citep{koyama1995,bamba2003,bamba2005}. The most plausible mechanism of the acceleration is the diffusive shock acceleration (DSA), which explains the broadband power-law spectrum with an index around 2 \\citep{drury1983,blandford1987}. Before the DSA phase in which electrons cross the shock front many times, they have to be pre-accelerated by an unknown injection mechanism. This ``injection problem'' is still unresolved.  One of the possible injection mechanisms is the shock surfing acceleration \\citep{Shimada_2000,Hoshino_2002}. \\citet{Shimada_2000} performed a one-dimensional (1D) electromagnetic full particle simulation and found the formation of large-amplitude electrostatic solitary structures during the cyclic reformation of a high Mach number perpendicular shock in a low-beta and weakly magnetized plasma. The coherent solitary structures trap electrons in their electrostatic potential well, resulting in significant surfing acceleration of electrons. The electron surfing has been regarded as an efficient direct acceleration mechanism at SNR shocks and studied by many authors \\citep[e.g.,][]{McClements_2001,Hoshino_2002, Schmitz_2002a,Schmitz_2002b,Amano_2007}. However, these previous works are based on 1D simulations, in which configuration of shock magnetic fields cannot be modified. Since charged particles can easily drift along magnetic fields, it is important to take into account the modification of shock magnetic fields and the free motion of particles along shock magnetic fields. Hence, multi-dimensional simulations are necessary.   In the present study we perform a 2D full particle simulation of a high Mach number and low-beta perpendicular collisionless shock. The preliminary result shows that the electron surfing acceleration is suppressed in a 2D system. ",
        "conclusions": "We developed a 2D shock-rest-frame model for full particle simulations. Then a 2D simulation of a low-beta and high Mach number perpendicular shock was performed, which shows that electrons are thermalized in the direction parallel to the ambient magnetic field in contrast to the previous 1D simulations. The heating of electrons along the ambient magnetic field is due to electromagnetic whistler mode waves excited by temperature anisotropy. In addition, the current-driven instability due to reflected ions does not lead to formation of coherent electrostatic solitary structures because of the free motion of electrons along the ambient magnetic field. As a result of these consequences, the generation of nonthermal electrons via electron surfing acceleration due to electrostatic solitary structures is suppressed in the 2D system. This process is independent of shock parameters such as $\\omega_{pe}/\\omega_{ce}$, $\\beta$, and $M_A$. Even with the realistic parameters, there would be less surfing acceleration in purely perpendicular shocks, because there is no guiding magnetic field to stabilize coherent electrostatic structures. However, the surfing acceleration might be effective in an oblique shock in which a guiding magnetic field could allow the formation of large-amplitude electrostatic structures.   It is noted that \\citet{Ohira_2007} have also shown that electron surfing acceleration becomes absent by a simplified 2D simulation with different shock geometry. That is, their simulation domain has been taken in a plane perpendicular to the ambient magnetic field (which corresponds to the $x$-$z$ plane in the present model). Thus, electrostatic solitary structures generated by the current-driven instability lose coherency by a quite different mechanism. However, both 2D simulations with different axes of coordinate have indicated that the electron surfing acceleration is suppressed more than in 1D simulations. Hence, another unknown injection mechanism for electron acceleration might be necessary."
    },
    "0806/0806.0937_arXiv.txt": {
        "abstract": "Baryon acoustic oscillations (BAOs) and weak lensing (WL) are complementary probes of cosmology. We explore the distance and growth factor measurements from photometric BAO and WL techniques and investigate the roles of the distance and growth factor in constraining dark energy. We find for WL that the growth factor has a great impact on dark energy constraints but is much less powerful than the distance. Dark energy constraints from WL are concentrated in considerably fewer distance eigenmodes than those from BAO, with the largest contributions from modes that are sensitive to the absolute distance. Both techniques have some well determined distance eigenmodes that are not very sensitive to the dark energy equation of state parameters $w_0$ and $w_a$, suggesting that they can accommodate additional parameters for dark energy and for the control of systematic uncertainties. A joint analysis of BAO and WL is far more powerful than either technique alone, and the resulting constraints on the distance and growth factor will be useful for distinguishing dark energy and modified gravity models. The Large Synoptic Survey Telescope (LSST) will yield both WL and angular BAO over a sample of several billion galaxies. \\chgd{Joint LSST BAO and WL can yield $0.5\\%$ level precision on ten comoving distances evenly spaced in $\\log(1+z)$ between redshift $0.3$ and $3$ with cosmic microwave background priors from {\\it Planck}. In addition, since the angular diameter distance, which directly affects the observables, is linked to the comoving distance solely by the curvature radius in the Friedmann-Robertson-Walker metric solution, LSST can achieve a pure metric constraint of 0.017 on the mean curvature parameter $\\Omega_k$ of the universe simultaneously with the constraints on the comoving distances.} ",
        "introduction": "\\label{sec:intr}  The luminosity distance--redshift relation measured from type Ia supernovae (SNe) has provided a vital piece of evidence for an accelerated cosmic expansion in a Friedman-Robertson-Walker universe \\citep{riess98,perlmutter99a}. Such an acceleration suggests that either the universe is dominated by a smooth dark energy component or we need an alternative framework of gravity to account for the SN data \\citep[e.g.,][]{dvali00,deffayet01}. Although different answers have vastly different theoretical implications, they can only be distinguished indirectly, for example, through measurements of the cosmic distance scale as a function of redshift and the amplitude of the density fluctuations as a function of redshift, i.e., the growth factor--redshift relation. One may further derive a phenomenological dark energy equation of state (EOS) as a function of redshift, $w(z)$, from the distances and growth factors and contrast it with predictions from theory.  Besides SNe, baryon acoustic oscillations (BAOs) in the galaxy power spectrum \\citep*{eisenstein98,cooray01b,blake03,hu03b,linder03b} and weak lensing \\citep[WL,][]{hu99a,mellier99,bartelmann01,huterer02, refregier03,song04} can also measure the distance \\citep*{seo03,song05,knox06b,zhan06c}. The sound horizon at the last scattering surface determines the BAO scale \\citep{peebles70,bond84,holtzman89}, which is fixed in comoving space, so that it can be used as a standard ruler to measure the comoving angular diameter distance $D_{A}(z)$ and, with spectroscopy, Hubble parameter $H(z)$. With precise calibration of the sound horizon from cosmic microwave background (CMB) observations, the BAO technique is able to measure the absolute distance, whereas the SN technique measures the relative distance because of the degeneracy between the SN intrinsic luminosity and the Hubble constant $H_0$. The WL technique does not depend on power spectrum features; its ability to measure distance originates from its geometric lensing kernel and the sensitivity of the shear power spectrum amplitude to the absolute distance \\citep[e.g.,][hereafter ZK06b]{zhan06e}.  The WL shear signal is a direct indicator of the gravitational potential of all matter, luminous or not, so the WL technique has the advantage of being able to measure the growth factor of the large-scale structure directly. For the BAO technique\\footnote{Although the BAO features are not related to the growth factor, the BAO technique, being a power spectrum or correlation function analysis, can potentially measure the growth factor.}, the growth factor is largely degenerate with the galaxy clustering bias, though it could be determined from the shape of the galaxy power spectrum \\emph{if} one knew precisely how the amplitude of the linear matter power spectrum affects the slope of the nonlinear matter power spectrum \\emph{and} how the galaxy bias behaves in the mildly nonlinear regime (see Section~\\ref{sec:dgcon}).  The relative importance of the distance and growth factor in constraining the dark energy EOS with WL has been discussed by \\citet{abazajian03}, \\citet[][hereafter SB05]{simpson05}, \\citet*[][hereafter ZHS05]{zhang05}, and \\citet{knox06b}, but their results are apparently inconsistent: SB05 and ZHS05 argue that the distance and growth factor are equally powerful whereas \\citet{abazajian03} and  \\citet{knox06b} show that the growth factor is much less powerful than the distance. This inconsistency is found to be caused by the parameter space of each analysis: with no other parameters, the distance and growth factor place comparable constraints on the dark energy EOS parameters, but the growth factor becomes less powerful when more parameters are allowed to float (ZK06b). Because the WL measurements of the distance and growth factor are entangled, ambiguity arises on the exact meaning of ``using only the distance (or growth) information'' without a reconstruction of distances and growth factors (actually, their errors). In this paper, we give a detailed account of estimating the distance and growth errors and their roles in constraining dark energy.  Throughout this paper, we assume a low-density cold dark matter (CDM) universe with the following parameters: the dark energy equation-of-state parameters $w_0$ and $w_a$ as defined by $w(z) = w_0 + w_az(1+z)^{-1}$ \\citep{chevallier01}, the matter density $\\omega_{m}$, the baryon density $\\omega_{b}$, the angular size of the sound horizon at the last scattering surface $\\theta_{s}$, the equivalent matter fraction of curvature $\\Omega_k$, the optical depth to scattering by electrons in the reionized inter-galactic medium, $\\tau$, the primordial helium mass fraction $Y_{p}$, the spectral index $n_{s}$ and running $\\alpha_{s}$ of the primordial scalar perturbation power spectrum, and the normalization of the primordial curvature power spectrum $\\Delta_R^2$ at $k = 0.05\\,\\mbox{Mpc}^{-1}$. For reasons stated in Section~\\ref{sec:dgpar}, only a subset of these parameters are retained when distance and growth parameters are present. The fiducial model is taken from the 3-year {\\it WMAP} data \\citep{spergel07}: ($w_0$, $w_a$, $\\omega_{m}$, $\\omega_{b}$, $\\theta_{s}$, $\\Omega_k$, $\\tau$, $Y_{p}$, $n_{s}$, $\\alpha_{s}$, $\\Delta_R^2$) = (-1, 0, 0.127, 0.0223, 0.596$^\\circ$, 0, 0.09, 0.24, 0.951, 0, $2.0 \\times 10^{-9}$). The reduced Hubble constant $h = 0.73$ and the present equivalent matter fraction of dark energy $\\Omega_{X} = 0.76$ are implicit in this parametrization, meaning that either one of them can replace $\\theta_{s}$ or any parameter that affects $\\theta_{s}$.  The rest of the paper is organized as follows. Section~\\ref{sec:wb} provides a brief introduction to the power spectrum analysis and error forecast with BAO and WL. For the investigation that follows, we assume a photometric redshift (\\phz{}) survey based on the Large Synoptic Survey Telescope\\footnote{See \\url{http://www.lsst.org}.} \\citep[LSST,][]{ivezic08} and consider only \\phz{} BAO and WL tomography. We compare BAO and WL constraints on the distance and growth factor in Section~\\ref{sec:dgch} and examine the roles of distance and growth for studying dark energy in Section~\\ref{sec:pde}. The conclusions are drawn in Section~\\ref{sec:con}. Since our main goal is to understand the intrinsic properties of the distance and growth factor measured from the two techniques and their applications in dark energy studies, we neglect systematics in Sections~\\ref{sec:dgch} and \\ref{sec:pde}. To be more realistic, we include in Section~\\ref{sec:con} a brief discussion of the constraints on the distance and growth factor from a joint analysis of LSST BAO and WL with conservative estimates of systematic uncertainties in \\phz{}s and power spectra. For more detailed accounts of various systematics, see \\citet*{huterer05a,white05,ma06,huterer06,jain06, zhan06d,eisenstein07b,guzik07,zentner08}. ",
        "conclusions": "\\label{sec:con}  We have disentangled the roles of the distance and growth factor in constraining the dark energy EOS and given a detailed comparison between the distance measurements from BAO and WL as well as their constraints on $w_0$ and $w_a$. We find that if the growth parameters are marginalized without being projected onto $w_0$ and $w_a$ constraints, the EP of LSST WL will increase by a factor of $\\sim 8$. However, if the distance parameters are marginalized instead, the degradation will be two orders of magnitude. This clearly shows that the growth factor is important to WL, but the distance is far more important. In contrast, the BAO EP increases by only 10\\%, when the growth parameters are marginalized. It is also interesting that LSST BAO and WL achieve similar EPs when the growth parameters are marginalized.  The reconstruction of a continuous function depends on how the function is represented by discrete parameters. To explore the intrinsic properties of the reconstructed distances from BAO and WL, we assign a large number of distance and growth parameters and examine various aspects of the distance eigenmodes. This exercise leads to the finding that the WL distance eigenmodes are more complementary to each other than the BAO ones. As such, 5 WL eigenmodes can provide most of the LSST WL distance constraints on $w_0$ and $w_a$, whereas many more are needed for BAO. A more useful result is that both LSST BAO and WL have some well determined distance eigenmodes that are not very sensitive to $w_0$ and $w_a$. These modes may be able to constrain additional dark energy parameters and some systematic uncertainties.  \\begin{figure} \\centering \\epsscale{1} \\plotone{f8} \\caption{Same as Fig.~\\ref{fig:dgcon}, but for distance (solid line and open triangles) and growth factor (dotted line and open circles) constraints from the joint analysis of LSST BAO and WL with a conservative level of systematic uncertainties in the photo-z error distribution and additive and multiplicative errors in the shear and galaxy power spectra (see text for more details). The lines represent the results with other parameters floating, and the symbols are with CMB priors from \\emph{Planck}. The joint constraints on distance are relatively insensitive to the assumed systematics. \\label{fig:dges}} \\end{figure}  The insight gained from the results in Sections~\\ref{sec:dgch} and \\ref{sec:pde} in the absence of various systematics is valuable for understanding the techniques. It is also of great interest for future surveys to have more realistic error estimates that include those systematics. To this end, we present in Fig.~\\ref{fig:dges} constraints on the distance (solid line and open triangles) and growth factor (dotted line and open circles) from the joint analysis of LSST BAO and WL with a conservative level of systematic uncertainties in the photo-z error distribution and additive and multiplicative errors in the shear and galaxy power spectra. The lines represent the results with other parameters floating, and the symbols are with CMB priors from \\emph{Planck}. By comparing with Fig.~\\ref{fig:dgcon}, one sees that the joint constraints on distance are relatively insensitive to the assumed systematics whereas those on the growth factor are affected more.  The results in Fig.~\\ref{fig:dges} are marginalized over additional 30 \\phz{} bias parameters $\\delta z_i$, 30 \\phz{} rms parameters $\\sigma_{z,i}$, 10 shear additive error parameters $A_j^\\gamma$, 10 shear multiplicative error parameters $f_j^\\gamma$, and 20 galaxy additive noise parameters $A_k^\\mathrm{g}$. \\chgd{ We extend the additive and multiplicative shear power spectrum errors in \\citet{huterer06} to include the galaxy power spectrum errors: \\begin{eqnarray} \\nonumber (\\bit{C}_\\ell^{XY})_{ij} &=& (1 + \\delta_{X\\gamma}^{\\rm K} f_i^X + \\delta_{Y\\gamma}^{\\rm K} f_j^Y) P_{ij}^{XY}(\\ell) + \\\\ && \\delta_{XY}^{\\rm K} \\left[\\delta_{ij}^{\\rm K} X_{\\rm rms}^2 \\bar{n}_i^{-1} + \\rho^X A_i^X A_j^Y \\left(\\frac{\\ell}{\\ell_*^X}\\right)^{\\eta^X}\\right], \\nonumber \\end{eqnarray} where $\\rho^X$ determines how strongly the additive errors of two different bins are correlated, and $\\eta^X$ and $\\ell_*^X$ account for the scale dependence of the additive errors. Note that the multiplicative error of galaxy number density is degenerate with the galaxy clustering bias and is hence absorbed by $b_i$. Below the levels of systematics future surveys aim to achieve, the most important aspect of the (shear) additive error is its amplitude \\citep{huterer06}, so we simply fix $\\rho^X = 1$ and $\\eta^X = 0$. For more comprehensive accounts of the above systematic uncertainties, see \\citet{huterer06, jain06, ma06, zhan06d}.}  We have applied priors $\\sigma_P(\\delta z) = 2^{-1/2} \\sigma_P(\\sigma_z) = 0.01(1+z)$ per \\phz{} parameter interval ($\\Delta z = 0.17$) in Fig.~\\ref{fig:dges}. \\chgd{For \\emph{Gaussian} \\phz{} errors with $\\sigma_z = 0.05(1+z)$, these priors correspond to a calibration requirement of 25 spectra per redshift interval of 0.17.} The constraints do not change appreciably even if one tightens the priors to $0.003(1+z)$, because the \\phz{} sample--\\phz{} sample cross correlations can calibrate the uncertainties in $\\delta z$ and $\\sigma_z$ to $10^{-3}$ level \\citep{zhan06d}. Similar results are obtained per redshift bin with spectroscopic sample--\\phz{} sample cross correlations \\citep{newman08}. For the shear multiplicative error, \\citet{massey07a} shows that current methods consistently achieve better than $2\\%$ precision. Allowing for another 5 to 10 years of development, we are hopeful that the multiplicative errors can be controlled to half a percent, i.e., $\\sigma_P(f_i^\\gamma) = 0.005$. \\chgd{For the shear additive error, \\citet{paulin-henriksson08} finds for a 15,000 deg$^2$ ground survey that it will be well under control over relevant scales if the point spread function (PSF) is calibrated over 100 stars. LSST will be able to use on average 2 -- 3 stars per square arcminute to calibrate the PSF in each exposure, so the shear additive error will not have a significant impact on scales larger than $\\sim 5$ arcminute (or $\\ell \\lesssim 2000$) even for a single exposure. This is supported by a study with Subaru data \\citep{wittman05}, which shows that the correlation of the residual stellar shear after the PSF correction using only 0.9 star per square arcminute (out of $\\sim 8$ arcmin$^{-2}$ available) is well below the cosmic shear signal at arcminute scale for a single 10-second exposure and that it is roughly inversely proportional to the number of exposures. } In addition to deep exposures in \\emph{u}, \\emph{g}, and \\emph{y}, LSST plans to take 400 exposures per sky field in each of its \\emph{riz} filters (these will be used for WL as well as \\phz{}). Thus it is reasonable to project that the additive shear error will be sub-dominant to the statistical errors on relevant scales. Full simulations of LSST performance are in progress. To be conservative, we assume that $(A_i^\\gamma)^2 = 10^{-9}$, which is roughly the amplitude of a $z \\gtrsim 2$ shear power spectrum at $\\ell \\sim 1000$. Finally, we infer from the Sloan Digital Sky Survey galaxy angular power spectrum \\citep{tegmark02} that the additive galaxy power spectrum error due to extinction, photometry calibration, and seeing will be at $(A_i^\\mathrm{g})^2 = 10^{-8}$ level.  With the assumed systematic uncertainties, we find that the joint analysis of LSST BAO and WL can achieve $\\sim 1.5\\%$ precision on 11 distances from $z_2 = 0.29$ to $z_{12} = 3.6$ with weak priors on cosmological parameters. Strong priors from \\emph{Planck} will reduce the errors to $\\sim 0.5\\%$ level. The corresponding errors on the growth factors are $\\sim 3.5\\%$ from $z_1 = 0.14$ to $z_{11} = 3.1$ without \\emph{Planck} and $\\sim 2\\%$ with \\emph{Planck}. These estimates are obtained without assuming a specific model for dark energy or modified gravity, so they are fairly model independent.  Finally, we mention that the distance and growth factor reconstruction automatically provides a pure metric measurement of the mean curvature of the universe. The joint LSST BAO and WL analysis can achieve $\\sigma(\\Omega_k) = 0.017$ with the above mentioned systematic uncertainties and weak priors on the other cosmological parameters. Though less impressive than other model-dependent forecasts of $\\sigma(\\Omega_k) \\lesssim 10^{-3}$ \\citep{knox06c,zhan06d}, this estimate does not assume any specific dependence of the distance on cosmological parameters."
    },
    "0806/0806.4808_arXiv.txt": {
        "abstract": "We investigate the BPS configuration of the multi D-vortices produced from the D2${\\bar {\\rm D}}$2 system. Based on the DBI-type action with a Gaussian-type runaway potential for a complex tachyon field, the BPS limit is achieved when the tachyon profile is thin. The solution states randomly-distributed $n$ static D-vortices with zero interaction. With the obtained BPS configuration, we derive the relativistic Lagrangian which describes the dynamics of free massive D-vortices. We also discuss the 90${}^{\\circ}$ and 180${}^{\\circ}$ scattering of two identical D-vortices, and present its implications on the reconnection in the dynamics of cosmic superstrings. ",
        "introduction": "The Bogomolnyi-Prasad-Sommerfield (BPS) limit of solitons~\\cite{Bogomolny:1975de} has been attractive in its nature and tractability of dynamics. The BPS limit is characterized by a minimum-energy configuration of a static soliton which is a solution to the first-order Bogomolnyi equation. The dynamics of multi-BPS solitons is characterized by no interaction. The BPS solitons do not exert any forces to nearby solitons, so they move freely until they collide.  Recently D- and DF-strings have attracted attention as new candidates of cosmic superstrings which produce Y-junctions~\\cite{Copeland:2003bj,Dvali:2003zj}. Concerning string cosmology, such superstrings generated through the decay of D3${\\bar {\\rm D}}$3 system are intriguing since the separated D${\\bar {\\rm D}}$ provides a natural setting for an inflationary epoch~\\cite{Dvali:1998pa} and can be employed in a string cosmological model with fluxes and moduli stabilization~\\cite{Kachru:2003sx}. Once D-strings are produced, understanding their dynamics~\\cite{Jackson:2004zg,Copeland:2006eh} is important in order to see the distinctive evolution between cosmic strings and cosmic superstrings.  In this work, we consider D-vortices which are produced in the coincidence limit of D2${\\bar {\\rm D}}$2. (The model is equivalent to parallel D-strings produced from D3${\\bar {\\rm D}}$3.) We consider an effective field theory of a complex tachyon field described by a Dirac-Born-Infeld (DBI) type action, which reflects the instability of D${\\bar {\\rm D}}$ system~\\cite{Sen:2003tm,Garousi:2004rd}. In the context of effective field theory, D0-branes from D2${\\bar {\\rm D}}$2 have been obtained as D-vortex configurations in (1+2)-dimensions~\\cite{Jones:2002si,Kim:2005tw,Cho:2005wp}.  In this article, we present a precise process obtaining the BPS configuration for multi D-vortices and their moduli-space dynamics based on Refs.~\\cite{Kim:2006xi,Cho:2007jf}. The BPS limit is possibly obtained with a Gaussian-type potential for the tachyon, and the BPS configuration expresses infinitely thin D-vortices randomly-distributed on a plane~\\cite{Kim:2006xi}. We show also that such a configuration reproduces the BPS sum rule and descent relation correctly. With the obtained BPS solution, we derive a relativistic Lagrangian which describes correctly the dynamics of the multi BPS D-vortices~\\cite{Cho:2007jf}. The BPS D-vortices behave as free point-like particles with mass identified with the D0-brane tension when they are separated. We show that the scattering of two identical BPS D-vortices exhibits the 90${}^{\\circ}$ scattering and the 180${}^{\\circ}$ (equivalently 0${}^{\\circ}$) scattering, which corresponds individually to the reconnection and the passing through of D-strings when it is applied to the dynamics of cosmic superstrings.  The article is organized as following. In Section~II, we review the Nielsen-Olesen vortex in the Abelian-Higgs model and present a derivation of its BPS bound. In Section~III, we present precise BPS conditions and obtain the BPS configuration of multi D-vortices. In Section~IV, we derive the classical moduli-space dynamics of these BPS D-vortices. We conclude in Section~V. ",
        "conclusions": "In this work, we obtained the BPS solution to the multi D-vortices produced in the coincidence limit of D2${\\bar {\\rm D}}$2. The model is described by a DBI-type action with a complex tachyon field. With a gaussian type-potential, we showed that the BPS limit is achieved by an infinitely thin tachyon profile. The obtained BPS configuration correctly reproduces the BPS sum rule and descent relation which result in the identification of BPS D-vortices as BPS D-branes of codimension two.  For multi BPS D-vortices, it is expected that the thin vortices are point-like particles moving freely when they are separated. Assuming that the time-dependence of the system is encoded only in the vortex positions in dynamics, we integrated the Lagrangian of the DBI action to obtain a simple Lagrangian $L^{(n)}({\\bf x}_{p},{\\dot {\\bf x}}_{p}) = -{\\cal T}_{0}\\sum_{p=1}^{n}\\sqrt{1-{\\dot {\\bf x}}_{p}^{2}}$ which describes $n$ free relativistic point particles with the mass given by the D0-brane tension.  While the scattering of Nielsen-Olesen vortices exhibits only the 90${}^{\\circ}$ scattering, the scattering of BPS D-vortices exhibits also the $0^{\\circ}$/180${}^{\\circ}$ scattering. When this scattering picture is applied to the interaction of cosmic superstrings, it implies that two identical BPS D-strings can pass through each other ($0^{\\circ}$/180${}^{\\circ}$ scattering) as well as can reconnect (90${}^{\\circ}$ scattering). This is one of the key differences of cosmic superstrings from usual cosmic strings."
    },
    "0806/0806.1046_arXiv.txt": {
        "abstract": "Any hint of non-gaussianity in the cosmological initial conditions will provide us with a unique window into the physics of early universe. We show that the impact of a small local primordial non-gaussianity (generated on super-horizon scales) on the statistics of collapsed objects (such as galaxies or clusters) can be approximated by using slightly modified, but {\\it gaussian}, initial conditions, which we describe through simple analytic expressions. Given that numerical simulations with gaussian initial conditions are relatively well-studied, this equivalence provides us with a simple tool to predict signatures of primordial non-gaussianity in the statistics of collapsed objects. In particular, we describe the predictions for non-gaussian mass function, and also confirm the recent discovery of a non-local bias on large scales \\cite{Dalal:2007cu,Matarrese:2008nc}, as a signature of primordial non-gaussianity. We then study the potential of galaxy surveys to constrain non-gaussianity using their auto-correlation and cross-correlation with the CMB (due to the Integrated Sachs-Wolfe effect), as a function of survey characteristics, and predict that they will eventually yield an accuracy of $\\Delta f_{NL} \\sim 0.1$ and $3$ respectively, which will be better than or competitive with (but independent of) the best predicted constraints from the CMB. Interestingly, the cross-correlation of CMB and NVSS galaxy survey already shows a hint of a large local primordial non-gaussianity: $f_{NL} = 236 \\pm 127$. ",
        "introduction": "The standard cosmological model has been tremendously successful in explaining a large range of cosmological observations, ranging from cosmic microwave background (CMB) anisotropies, to the galaxy redshift surveys, supernovae Ia, baryonic acoustic oscillations and the Lyman-$\\alpha$ forest in the quasar spectra. However, our clues into the nature of the initial conditions of the Universe, as well as its underlying physics, remain scarce.  The theory of cosmological slow-roll inflation remains the most successful and widely studied model for the cosmological initial conditions. However, despite the large inflow of cosmological observations over the past few years, it remains practically impossible to distinguish a large class of inflationary scenarios. Moreover, given the limited amount of observables accessible to a today's observer, it is not clear whether any possible future observations will be able to falsify the inflationary scenario. One exception to this dilemma would be a possible detection of non-gaussianity  of cosmological initial conditions, which is predicted to be unobservably small in standard slow-roll inflationary scenarios (e.g. \\cite{Maldacena:2002vr}). While a large primordial non-gaussianity can still be accommodated in a general inflationary scenario through introducing fine-tuned secondary light fields, higher derivative interactions, or non-adiabatic initial conditions, it could also be interpreted as a smoking gun for alternative early universe models, such as the new ekpyrotic scenario \\cite{Buchbinder:2007ad}, which naturally predicts a large primordial non-gaussianity \\cite{Buchbinder:2007at}.  The best current observational constraints on the primordial non-gaussianity comes from the statistics of the all-sky CMB maps surveyed by the Wilkinson Microwave Anisotropy Probe (WMAP). Using the 3 year data release of the WMAP CMB maps, Yadav and Wandelt \\cite{Yadav:2007yy} claimed a detection of an $f_{NL}$- type local primordial non-gaussianity (see Eq.~(\\ref{fnl})), finding $f_{NL} = 87 \\pm 30$. This is inconsistent with the expectation from single-field slow-roll inflationary models that predict a much smaller value $|f_{NL}| \\lesssim 0.1$, and might be seen as evidence for new ekpyrotic scenario which predicts $f_{NL} \\sim 10-100$ \\cite{Buchbinder:2007at}. However, the more recent data release of 5 year WMAP maps shows a smaller value of $f_{NL} = 51 \\pm 30$ \\cite{Komatsu:2008hk}, which is now consistent with the minimal inflationary expectation at the 2$\\sigma$ level. It is thus clear that an accurate determination of the $f_{NL}$ parameter can have important implications for the early universe models. The next generation of all-sky CMB surveys, namely the Planck satellite, is expected to have enough sensitivity to reduce the error on $f_{NL}$ to $\\sim 3$ (e.g. \\cite{Serra:2008wc}). However, CMB secondary anisotropies (such as Sunyaev-Zel'dovich effect) are expected to contaminate the CMB at this level, and thus prevent a more accurate measurement.  Can we expect complementary (and competitive) constraints from the study of large scale structure at low redshifts? This question was traditionally addressed by looking at the mass function of galaxy clusters, which was expected to be a sensitive probe of non-gaussian initial conditions (e.g. see \\cite{2000ApJ...530...80W,LoVerde:2007ri,Dalal:2007cu} and references therein). Unfortunately, the predicted constraints are not competitive with the CMB, unless $f_{NL}$ has a  strong scale-dependence. Recently, Sefusatti and Komatsu \\cite{Sefusatti:2007ih} argued that the analysis of the bispectrum of upcoming large-volume galaxy surveys can yield competitive constraints with the CMB, but their method relies on the accuracy of a local bias model for the distribution of galaxies on small scales. Interestingly, Dalal et al. \\cite{Dalal:2007cu} recently predicted the presence of an anomalous large scale power in galaxy distribution as a possible smoking gun for local primordial non-gaussianity (also see \\cite{Matarrese:2008nc,McDonald}). This effect would only be seen on very large scales, and thus cannot be confused with any other causal astrophysical effect (although Galactic contamination or large-angle survey systematics could give rise to a similar spurious power). Shortly before the submission of this paper, Slosar et al. \\cite{Slosar:2008hx} published a comprehensive study of this effect in the present large-scale cosmological galaxy surveys, leading to a significant constraint: $f_{NL} = 28 \\pm 25$. We will briefly comment on this result in Sec. \\ref{conclude}.  In this paper, we present a fully analytic approximation that could recast the statistics of collapsed objects (i.e. galaxies, clusters, etc.) with local primordial non-gaussian initial conditions, in terms of the same statistics with gaussian initial conditions but a modified power spectrum. In Sec.\\ref{linear}, we introduce the approximation that recasts the non-gaussian linear density field in terms of a linear combination of gaussian fields, for collapsing objects of a given mass. Sec.\\ref{mass_function} applies this result to the elliptical collapse framework, and compares the resulting mass function with numerical simulations. We then discuss the implications for the large scale clustering of galaxies (due to a non-local bias correction) in Sec.\\ref{clustering}, and show that it could yield competitive (or even superior) constraints ($\\Delta f_{NL} \\sim 0.1$), in comparison with the best upcoming CMB constraints. Finally, Sec.\\ref{isw} shows that cross-correlation of CMB with galaxy surveys, due to its sensitivity to the large scale galaxy bias, can similarly put strong constraints on the primordial non-gaussianity which are also competitive with the CMB ($\\Delta f_{NL} \\sim 3$), and possibly less sensitive to the galaxy survey systematics. Current cross-correlation data between the WMAP CMB maps and the NVSS radio galaxy survey already show a $\\sim 2\\sigma$ hint for such non-gaussian signal.  Unless mentioned otherwise, we use the WMAP5 best fit $\\Lambda$CDM flat cosmology \\cite{Komatsu:2008hk} with $\\Omega_m = 0.279$, $\\Omega_b = 0.0462$, $n_s = 0.96$, $h =0.701$ throughout. For all our numerical calculations, we use the Eisentein \\& Hu analytic approximation \\cite{Eisenstein:1997ik} for the matter transfer function that includes baryonic features, although we show the simpler BBKS, dark matter only, analytic fit \\cite{Bardeen:1985tr} in Eqs.~(\\ref{bbks}) \\& (\\ref{epsilon}), in order to illustrate the asymptotic structure of the transfer function. Finally, all the errors on $f_{NL}$ predicted or quoted in this paper are at 68\\% (1$\\sigma$) confidence level. ",
        "conclusions": "In this paper, we have studied the impact of of a local primordial non-gaussianity (expected from super-horizon evolution in multi-field inflationary models or the new ekpyrotic scenario) on the statistics of collapsed objects in the late universe. We developed an approximation that describes the linear non-gaussian density field as a local bilinear combination of gaussian fields. Applying this to the elliptical collapse framework enabled us to find an analogue gaussian linear density field, which should produce the same predictions as the non-gaussian system for collapsed objects of a given mass, assuming small primordial non-gaussianity. For halo mass function, this prescription leads to an effective analytic correction to $\\sigma(M)$, the amplitude of linear fluctuations on mass scale $M$, which can be readily applied to Press-Schechter-like analytic approximations for the mass function. We found this to give roughly consistent results with direct simulations with non-gaussian initial conditions. We then focused on the impact of non-gaussianity on the clustering of galaxies, and confirmed the recent discovery of a non-local bias correction associated with local non-gaussian initial conditions \\cite{Dalal:2007cu,Matarrese:2008nc}. We found that auto-correlation of galaxy surveys can ultimately lead to a precision $\\Delta f_{NL} \\sim 0.1$ which is an order of magnitude better than the best anticipated CMB constraints. Cross-correlation of CMB and galaxy surveys (through the ISW effect), even though suffering from a low signal-to-noise due to primary CMB contamination, is less prone to survey systematics, and can also yields $\\Delta f_{NL} \\sim 3$, which is again competitive with future CMB constraints. Interestingly, the cross-correlation of CMB and the NVSS radio galaxy survey already provides a $2\\sigma$ hint for a large primordial non-gaussianity $f_{NL} = 236 \\pm 127$.  Finally, we should comment on the future developments of the studies of primordial non-gaussianity. While five more years of WMAP observations can only marginally improve the current constraints to $\\Delta f_{NL} \\sim 15-20$, the most anticipated observational improvement is the first data release of Planck satellite, expected in 2012. After careful accounting for CMB secondary anisotropies, as well as the 2nd order corrections to the CMB Boltzmann solvers, Planck is expected to achieve $\\Delta f_{NL} \\sim 3$. However, as we demonstrated in this paper, competitive or even superior constraints can come from the study of clustering of large scale galaxy surveys. For example, Large Synoptic Survey Telescope (LSST) is expected to survey half of the sky down to magnitude $r\\sim 27.5$ finding $\\sim 10^{10}$ galaxies out to $z_{max} \\sim 2$. Using Fig.~(\\ref{dfnl}), LSST can yield a precision of $\\Delta f_{NL} \\sim 2-3$ through its galaxy auto-correlation function. The cross-correlation of LSST with the CMB sky could independently yield $\\Delta f_{NL} \\sim 7$. Similar proposed or underway deep, large area surveys such as Pan-STARRS, ADEPT, CIP, etc. could potentially lead to comparable constraints. However, one important point to note is that {\\it precise spectroscopic redshifts are not required for non-gaussianity determination, as most of the non-gaussian signal comes from the largest scales of the survey}.  Shortly before the submission of this paper, Slosar et al. \\cite{Slosar:2008hx} published a comprehensive study of this effect in the present large-scale cosmological galaxy surveys, leading to a significant constraint: $f_{NL} = 28 \\pm 25$. Nearly all the statistical weight of this constraint is based on the lack of anomalous power in the lowest l-bin of the auto-power spectrum of a high redshift sample of $\\sim 10^5$ SDSS photometric QSO's, which cover $15\\%$ of sky and have an effective mean bias of $\\sim 2.75$. Plugging these numbers into the error estimates of Sec. \\ref{clustering} (Fig. \\ref{dfnl}), we would predict an error of about a factor of $2$ worse than what was found in \\cite{Slosar:2008hx}. As we pointed out in Sec. \\ref{clustering}, the actual error on $f_{NL}$ will be sensitive to the survey geometry, and so this disagreement between our rough error prediction and Slosar et al.'s findings may not be an immediate cause for alarm. Nevertheless, it is curious that despite the lack of any redshift information (which was used in our error prediction), as well as possible contaminations/survey systematics, Slosar et al. could achieve an accuracy in $f_{NL}$ comparable or better than an optimistic error prediction. A further concern is that the correlated Poisson noise (typically expected from non-linear clustering, e.g. \\cite{McDonald:2006mx}) is not marginalized over in their analysis.  On a different note, Slosar et al. point out that a degeneracy between $f_{NL}$ and the redshift and bias distribution of NVSS survey (which are both poorly known), hinders any accurate determination of $f_{NL}$ from NVSS data, or its correlation with the CMB. While, at face value, this might seem inconsistent with our $\\sim 2\\sigma$ evidence for non-gaussianity from the correlation of NVSS with WMAP, we should note that both redshift and bias distributions may eventually be fixed, either from theoretical modeling, or direct observations, which can then realize the full statistical power of cross-correlation constraints.  On the theoretical front, we are still far from a good theoretical understanding of non-linear structures from non-gaussian initial conditions. While the present work provides a bridge between this problem, and the huge body of work on non-linear/collapsed structures from gaussian initial conditions (such as the semi-analytic frameworks, the halo model, as well the N-body/hydro numerical simulations), it does suffer from some simplifying assumptions. The first is the ellipsoidal collapse scenario \\cite{Sheth:1999su, Sheth:2001dp}. Although for large masses and $f_{NL}\\epsilon(M) \\ll 1$ (i.e. near-gaussian initial conditions), the collapse should be well-described as spherical (leading to a universal spherical collapse threshold), for smaller masses, the threshold depends on the ellipticity of the proto-halo, which is a random number with a distribution set by the statistics of the initial conditions. The current analytic form (Eq.~\\ref{ecol}) is a fit to random realizations of gaussian fields with $\\Lambda$CDM power spectrum, and thus there is no reason to exclude corrections linear in $f_{NL}$, if the linear density field obeys non-gaussian statistics. Therefore, one clear direction for improvement is to investigate/extend the ellipsoidal collapse framework for non-gaussian initial conditions.  The second simplification (and probably the most non-trivial part of this work) is the local approximation of Eq.~(\\ref{deltax}), which was instrumental in all the subsequent analysis/predictions. While the consistency of our results for mass function and clustering with other numerical and analytic studies \\cite{Dalal:2007cu,Matarrese:2008nc} verifies the accuracy of this approximation, Eq.~(\\ref{deltax}) could also be viewed as a {\\it convenient parametrization for the primordial non-gaussianity}. To see this, note that for an $f_{NL}$-type non-gaussianity (Eq.~\\ref{fnl}), the bispectrum of $\\Phi$, $B(k_1,k_2,k_3)$,  is written in terms of the power spectrum of the Gaussian field $\\Phi_{pG}$, $P_{pG}(k)$: \\bea B(k_1,k_2,k_3) = -2f_{NL}P_{pG}(k_1)P_{pG}(k_2)P_{pG}(k_3)\\times\\nonumber\\\\ \\left[P^{-1}_{pG}(k_1)+P^{-1}_{pG}(k_2)+P^{-1}_{pG}(k_3)\\right] \\label{bispectrum}, \\eea where $P_{pG}(k) \\propto k^{-3}$ for scale-invariant spectra. While our local approximation does not exactly preserve this form, different observables are mostly sensitive to a limited region in the $k_1,k_2,k_3$ space, where Eq.~(\\ref{deltax}) may provide a simple parametrization. For example, it is reasonable to think that measurements of mass function of haloes of mass $M$ would only be sensitive to non-gaussianity around $k_1 \\sim k_2 \\sim k_3 \\sim \\pi/R(M)$ (see the line preceding Eq.~\\ref{epsilon}). Therefore, constraints on $f_{NL}$ using mass function could be roughly translated into constraints on the bispectrum for {\\it near-equilateral triangles} of the side $\\sim \\pi/R(M)$.  A more interesting case is the clustering (or the power spectrum) of collapsed haloes (e.g. galaxies) that we studied in Secs. \\ref{clustering}-\\ref{isw}. In this case, we are dealing with a hierarchy of scales: $k_3 \\ll k_1,k_2$, where $k_1 \\ll 0.01~ \\mpc^{-1}$ is the scale over which we measure the galaxy power spectrum, and as we showed pre-dominantly corresponds to the modes of the size of the survey. On the other hand, $k_1,k_2 \\sim \\pi/R(M) \\gtrsim 0.1 ~\\mpc^{-1}$ correspond to the comoving size of the collapsing regions. Therefore, observations of the galaxy power spectrum provides an exquisite measure of the bispectrum for {\\it squeezed triangles}. In particular, the non-gaussian bias correction $\\Delta b_{NG}(k_3) \\propto k_3^{-2}$ found in Sec. \\ref{clustering} (as well as in  \\cite{Dalal:2007cu,Matarrese:2008nc}) are specific to the local form of non-gaussianity (Eq.~\\ref{bispectrum}) which is only produced through super-horizon evolution, where $B(k_1,k_2,k_3) \\propto k_3^{-3}$ as $k_3 \\rightarrow 0$. In contrast, for non-gaussianities produced through sub-horizon evolution (such as in slow-roll single field inflation, DBI, ghost inflation and k-inflation), $B(k_1,k_2,k_3) \\rightarrow $ const. or 0 in this limit \\cite{Babich:2004gb}. Therefore, by way of comparison,  the scale-dependent bias $\\Delta b_{NG}(k_3) \\rightarrow 0$ (and so is unobservable) in the long wavelength limit, for models that produce equilateral primordial non-gaussianity.  We thus close this paper by the exciting (although, by now, familiar) premise of the cosmological microscope as a probe of the fundamental physics of the early universe. In particular, that an anomalously large clustering of galaxies on large scales will be a unique smoking gun for non-linear super-horizon evolution (and thus large primordial isocurvature modes) during the early stages of cosmogenesis.  We would like to thank Neal Dalal for providing us with the mass functions from their non-gaussian simulations \\cite{Dalal:2007cu}. We also would like to thank Neal Dalal, Chris Hirata, Shirley Ho, Marilena LoVerde, Pat McDonald, Bob Scherrer, Uros Seljak, Anze Slosar, David Spergel, Michael Strauss, and James Taylor for useful comments and illuminating discussions.  This research was supported by Perimeter Institute for Theoretical Physics.  Research at Perimeter Institute is supported by the Government of Canada through Industry Canada and by the Province of Ontario through the Ministry of Research \\& Innovation."
    },
    "0806/0806.1877_arXiv.txt": {
        "abstract": "{Spatially resolved observations of circumstellar disks at millimeter wavelengths allow detailed comparison with theoretical models for the radial and vertical distribution of the material.} {We investigate the physical structure of the gas component of the disk around the pre-main-sequence star \\object{HD169142} and test the disk model derived from the spectral energy distribution.} {The $^{13}$CO and C$^{18}$O $J$=2--1 line emission is observed from the disk with $1\\farcs4$ resolution using the Submillimeter Array. We adopt the disk physical structure derived from a model which fits the spectral energy distribution of HD169142. We obtain the full three-dimensional information on the CO emission with the aid of a molecular excitation and radiative transfer code. This information is used for the analysis of our observations and previous $^{12}$CO $J$=2--1 and 1.3~mm continuum data.} {The spatially resolved $^{13}$CO and C$^{18}$O emission shows a Keplerian velocity pattern % The disk is seen at an inclination close to 13$^{\\circ}$ from face-on. We conclude that the regions traced by different CO isotopologues % are distinct in terms of their vertical location within the disk, their temperature and their column densities. With the given disk structure, we find that freeze-out is not efficient enough to remove a significant amount of CO from gas phase. Both observed lines match the model prediction both in flux and in the spatial structure of the emission. Therefore we use our data to derive the $^{13}$CO and C$^{18}$O mass and consequently the $^{12}$CO mass using standard isotopic ratios.} {We constrain the total disk gas mass to (0.6-3.0)$\\times$10$^{-2}$~M$_{\\odot}$. Adopting a maximum dust opacity of 2~cm$^2$g$^{-1}_{\\rm dust}$ we derive a minimum dust mass of 2.16$\\times$10$^{-4}$~M$_{\\odot}$ from the fit to the 1.3~mm data. % Comparison of the derived gas and dust mass shows that the gas to dust mass ratio of 100 is only possible under the assumption of a % dust opacity of 2~cm$^{2}$/g$^{-1}$ and % $^{12}$CO abundance of 10$^{-4}$ with respect to H$_2$. However, our data are also compatible with a gas to dust ratio of 25, with a dust opacity of 1~cm$^{2}$/g$^{-1}$ and $^{12}$CO abundance of 2$\\times$10$^{-4}$.}  {} ",
        "introduction": "\\begin{figure}[!htp] \\includegraphics[angle=0,width=8cm]{olay_6spec.ps} \\caption {Spectra of  C$^{18}$O J$=$2--1 (bottom), $^{13}$CO J$=$2--1 (middle) and $^{12}$CO $J$=2--1 line (top, from \\citet{raman}) summed over the central 4$\\arcsec\\times$4$\\arcsec$ toward HD169142. The $^{13}$CO and $^{12}$CO fluxes are shifted vertically by 2 and 5~Jy respectively. % Black lines show the observed spectra and the red lines show the model fit found in Section 4.3.1 (for $^{12}$CO the model from \\citet{raman} was used).} \\label{spec} \\end{figure}     Although the presence of molecular gas in disks around intermediate mass pre-main-sequence (Herbig Ae) stars was reported a decade ago \\citep{mannings}, the research in this field has focused more on their less massive counterparts (T Tauri stars). The disks around T Tauri stars have masses ranging from 0.001 to 0.1~M$_{\\odot}$ \\citep{beckwith}, usually derived from millimeter continuum fluxes assuming a gas to dust mass ratio of 100, as in molecular clouds. Their outer radii are constrained by molecular line observations and are typically a few hundred AU \\citep[][and references therein]{simon,thi}. Due to the low luminosity of the central star (0.5 to 1~L$_{\\odot}$), these disks are relatively cold (less than 20~K beyond 100~AU from the star) causing a significant depletion of the CO in the midplane of the outer disk. On the other hand, the Herbig Ae stars are about ten times more luminous than T Tauri stars, and consequently their disks are warmer. This allows the easiest to detect and the most commonly used gas tracer, carbon monoxide, to remain in the gas phase even in the disk midplane. Observations of CO and its isotopologues toward Herbig Ae stars are therefore expected to be more powerful probes of the full disk structure. Only a few Herbig Ae disks have been studied thoroughly via spatially resolved observations of molecular line emission that includes the optically thin CO isotopologues:  AB Aur \\citep{pietu2}, MWC480 \\citep{pietu1}, HD163296 \\citep{isella}. ",
        "conclusions": "\\citet{raman} show that the model that fits the SED of HD169142, can also be used to fit the structure and intensity of the resolved $^{12}$CO emission. They note that some weak residual emission is still present after subtracting the model from the observational data, that may correspond to real substructure within the disk. We analyse $^{13}$CO and C$^{18}$O data, which are expected to probe different depths in the disk due to their lower opacities.  \\begin{figure*}[!htp] \\centering \\includegraphics[width=17cm]{13co21mapsidl1.ps} \\caption {-- {\\it{Top panel}:} Channel maps of the observed $^{13}$CO $J$=2--1 emission at a range of velocities over which the line is detected. The lower left corner of this panel shows the size and position angle of the synthesized beam. -- {\\it{Middle panel}:} $^{13}$CO $J$=2--1 chanel maps from the best-fit disk model from Section 4.3.1. --  {\\it{Lower panel}:} Channel maps showing the residual emission after subtraction of the best-fit model from the data. The contour levels are -1,1,2,3,4$\\times$360~mJy~beam$^{-1}$ (2 sigma) in all panels.} \\label{13maps} \\end{figure*}  \\begin{figure*} \\centering \\includegraphics[width=17cm]{c18o21mapsidl1.ps} \\caption {-- {\\it{Top panel}:} Channel maps of the observed C$^{18}$O $J$=2--1 emission at a range of velocities over which the line is detected. The lower left corner of this panel shows the size and position angle of the synthesized beam. -- {\\it{Middle panel}:} C$^{18}$O $J$=2--1 channel maps from the best-fit disk model from Section 4.3.1. --  {\\it{Lower panel}:} Channel maps showing the residual emission after subtraction of the best-fit model from the data. The contour levels are -1,1,2,3,4,5$\\times$360~mJy~beam$^{-1}$ (2 sigma) in all panels.} \\label{18maps} \\end{figure*}  \\subsection{Adopted disk model}  \\citet{dentHD} fit the spectral energy distribution (SED) of HD169142 using an accretion disk model from the \\citet{dalessio} database with an accretion rate of 10$^{-8}$~M$_{\\odot}$yr$^{-1}$, an outer radius of 300~AU, and a 30$^{\\circ}$ inclination. The age of the central star in this model is 10~Myr. They adopt an A2 spectral type for the central star because it provides a slightly better fit than the A6 spectral type model does to the SED in the range of 2-200~$\\mu$m dominated by the warm dust emission. However, there is no significant difference between the SEDs corresponding to these two models. Beyond 44~AU from the star, the model is qualitatively similar to the model of the outer disk of HD169142 described in \\citet{grady}. The detailed model structure, including plots of surface density distribution, temperature at different scale heights and optical thickness of the disk, is provided in the online database of accretion disk models \\citep{dalessio}\\footnote{www.astrosmo.unam.mx/~dalessio/}.  In order to match their resolved submillimeter observations of CO, \\citet{raman} modify the model used in \\citet{dentHD} to have a disk radius of 235~AU. Assuming the mass of the central star to be 2~M$_{\\odot}$, they derive an inclination of 13$\\pm$1$^{\\circ}$ from face-on and use it in the model. These modifications are not expected to affect the quality of the SED fit since the SED of HD169142 alone does not constrain well the outer radius and inclination of the disk. They show that this SED model can be used to fit the structure and intensity of the resolved $^{12}$CO emission. However, the SED does not provide reliable constraints of the disk dust mass, which in the model is set to 2$\\times$10$^{-4}$~M$_{\\odot}$. This parameter depends only on the optically thin thermal emission from the disk in the millimeter wavelength region.  The density in the model is a parameter that can be scaled by a factor of few, without significantly affecting the computed disk temperature structure. The resulting mass only affects the millimeter flux in the modeled SED, i.e., the 1.3~mm data point. In section~4.2 we re-analyse the 1.3~mm data from \\citet{raman} using a radiative transfer code to calculate the emission from the model and better constrain the dust mass of the disk.  In the analysis of our $^{13}$CO and C$^{18}$O line emission (section~4.3) we use this improved disk model, which already encompasses a range of observational constraints. We assume that the gas follows the dust distribution prescribed by the model and that their temperatures are the same. The micro turbulent FWHM of the lines is set to 0.16~km~s$^{-1}$ throughout the disk. One conclusion that we can already draw from the disk model is the fraction of disk mass where freeze-out of CO is likely to be efficient, which is believed to happen at temperatures below 20~K. The mass weighted temperature in the model is 35~K, which indicates that the bulk of the disk material is at temperatures above 20~K. The temperature falls below 20~K only in the midplane region of the disk beyond 150~AU from the star and there the abundance of the gas phase CO and its isotopologues is expected to be heavily decreased due to freeze-out of these molecules onto dust grains. This region contains only 8\\% of disk mass and therefore any mass estimate based on $^{13}$CO and C$^{18}$O emission may represent 92-100\\% of the true disk mass.  In all previously published work on HD169142, the disk gas mass was inferred indirectly from the derived disk dust mass, adopting a standard gas to dust ratio. In our analysis of the $^{13}$CO and C$^{18}$O emission we argue that it is reasonable to assume a standard $^{12}$CO abundance in the case of HD169142, because freeze-out affects less than 8\\% of the mass in this disk. Therefore we investigate the constraints that can be obtained from our molecular line data on the disk gas mass.  \\subsection{Dust continuum emission}  From the observed 1.3~mm continuum flux of 169$\\pm$5~mJy, \\citet{raman} derive a dust mass of 2$\\times$10$^{-4}$~M$_{\\odot}$ assuming an emissivity $\\kappa_{\\nu}^{dust}$=2~cm$^2$g$\\mathrm{^{-1}_{dust}}$ and a single temperature of 30~K. We reanalyse the 1.3~mm emission using the disk temperature and density structure from the model described in section 4.1, where we vary the dust density (i.e., the dust mass) to fit the observations. A large uncertainty on the mass determination from the 1.3~mm continuum flux is due to the dust opacity, which is not well determined in circumstellar disks. \\citet{ossenkopf} study the effects of dust coagulation, and ice coverage of grains on their opacity in protostellar cores and suggest 1.3~mm opacity of 1~cm$^2$g$\\mathrm{^{-1}_{dust}}$ for very dense (n$>$10$^{7}$~cm$^{-3}$) regions, and up to 5 times more if the grains are not covered in ice mantles. The latter is thus only valid for regions above water ice desorption temperature (80-100~K, \\citet{fraser}). In their study of the dust opacity in circumstellar disks \\citet{draine} find that astrosilicate and pirolysed cellulose at 600~$^{\\circ}$C are materials representative of the dust properties which may be expected in circumstellar disks. They explore a large range of grain size distributions and find that these materials have mm-wavelength opacities which are close to the observational constraints on dust opacity from extinction studies of the diffuse interstellar medium presented in \\citet{weingartner}. The resulting opacity could be anywhere between 0.1 and 2~cm$^2$g$\\mathrm{^{-1}_{dust}}$ at the wavelength of 1.3~mm, depending on the adopted grain size. We use this information to calculate the minimum amount of dust needed to produce the observed flux by adopting the opacity (and emissivity) of 2~cm$^2$g$\\mathrm{^{-1}_{dust}}$. The SED modelling of the disk emission at shorter wavelengths is not affected by our assumptions of dust opacity at 1.3~mm. The dust thermal emission is dominated by the cold disk midplane at large radii, likely to contain settled and/or grown dust particles whose optical properties may differ from those of the small dust at disk surface, responsible for the optically thick near-infrared and infrared emission of the disk.   We do the full modeling of the interferometric visibilities using the radiative transfer code of \\citet*{hogerheijde} (RATRAN). We perform a $\\chi^2$ minimisation to fit the 1.3~mm visibilities by varying the disk dust mass. The best match is shown in Fig.~\\ref{uvamp}, top-left panel, resulting in an estimate of the dust mass of the disk, given by \\begin{equation} \\mathrm{M_{dust}= 2.16\\times10^{-2}\\frac{2~cm^2g{^{-1}_{dust}}}{\\kappa_{1.3~mm}}\\times M_{\\odot}}. \\end{equation} Assuming that the adopted dust opacity of 2~cm$^2$g$\\mathrm{^{-1}_{dust}}$ is the maximum value, our estimate presents the lower limit on the dust mass in HD169142: $\\mathrm{M_{dust}\\geq2.16\\times10^{-4}~M_{\\odot}}$. A slightly higher maximum dust opacity of 7~cm$^2$g$\\mathrm{^{-1}_{dust}}$ is obtained for pyrolised cellulose at 800~$^{\\circ}$C but this value exceeds greatly the observational contraints from \\citet{weingartner} and although not impossible, is unlikely in disks \\citep{draine}. However we cannot exclude the possibility that other dust properties like cristallinity, porosity and shape, which are not well understood, may enhance dust submillimeter opacity beyond 2~cm$^2$g$\\mathrm{^{-1}_{dust}}$ by a factor of a few.  \\subsection{Molecular line emission}  Due to their low abundance relative to $^{12}$CO, the $^{13}$CO and C$^{18}$O molecules emit lines that are comparatively less optically thick. As a result of this, they saturate at different heights in the disk allowing us to use them as a probe of the physical conditions in different disk layers and investigate the disk vertical structure \\citep{dartois}. Furthermore, the C$^{18}$O line emission is often optically thin in the outer disk region, where the bulk of disk mass is located, which makes it a valuable probe of the disk gas mass.  The commonly employed route for the analysis of the CO isotopologue and dust continuum emission from circumstellar disks is to convert the continuum flux to disk mass adopting a certain dust opacity, usually 0.02~cm$^2$g$^{-1}_{\\rm gas}$ \\citep{beckwith}, under the assumption of the canonical gas to dust mass ratio of 100. In this way, the molecular line emission is fit for abundance, where any discrepancy with respect to the canonical abundances and isotopic ratios of observed molecules is explained by a depletion factor, indicating the fraction of the CO gas affected by the freeze-out onto dust grains or selective photodissociation. This approach is used in \\citet{raman} where the analysis of the $^{12}$CO J$=$2--1 line emission and 1.3~mm continuum from HD169142 is done for a disk mass of 2$\\times$10$^{-2}$~M$_{\\odot}$ and $^{12}$CO was found to be depleted by a factor four, having an abundance of 2.5$\\times$10$^{-5}$. This corresponds to  6$\\times$10$^{-6}$~M$_{\\odot}$ of $^{12}$CO molecules present in gas phase in the disk. In the present work, we argue that the $^{12}$CO abundance is not significantly affected by freeze-out or selective photodissociation and we chose to follow a different route, in which we derive a rough gas mass estimate from the molecular line observations taking into account all related uncertainties.   In order to interpret our $^{13}$CO and C$^{18}$O observations, we adopt the disk model described in detail in section 4.1. We set the dust mass in the model to 2.16$\\times$10$^{-4}$~M$_{\\odot}$; the value derived as the minimum dust mass for HD169142 in section 4.2. We assume that the molecular gas follows the dust temperature and density structure and we vary the $^{13}$CO and C$^{18}$O gas masses to best fit the observations. We calculate the $^{13}$CO and C$^{18}$O $J$=2--1 line emission from the model using the RATRAN molecular excitation and radiative transfer code. The model visibilities are generated with the miriad task uvmodel, and the comparison to the observed visibilities is done directly. We use assumptions about the molecular abundances of the $^{12}$CO and isotopic ratios to make a conversion to the total gas mass (1.2$\\times$M$_{\\rm H_2}$). We explore a range of values of the gas to dust mass ratio f$_{g/d}$.  \\begin{figure*} \\begin{center} \\parbox{15cm}{ \\it } \\begin{minipage}{1.0\\textwidth} \\includegraphics[width=8cm]{uvampzz.ps} \\includegraphics[width=8cm]{12visxz.ps} \\end{minipage} \\parbox{15cm}{ \\it } \\begin{minipage}{1.0\\textwidth} \\includegraphics[width=8cm]{13visxz.ps} \\includegraphics[width=8cm]{18visxz.ps} \\end{minipage} \\caption{{\\it{Top-left}:} Vector-averaged continuum flux as a function of projected baseline length (black dots). The error bars represent the variance within each annular average. The zero-signal expectation value is shown by the dashed histogram. The red dotted line shows the best-fit model for a dust emissivity of 2.0~cm$^2$g$\\mathrm{^{-1}_{dust}}$ with the disk mass of 2.16$\\times$10$^{-4}$~M$_{\\odot}$. {\\it{Top-right}:} Same for the $^{12}$CO emission from \\citet{raman}, integrated over the width of $\\approx$2.7~km~s$^{-1}$ (10 spectral channels) centered on each line. The red dotted line shows the model presented by \\citet{raman}. {\\it{Bottom-left and bottom-right}:} Same for our $^{13}$CO and C$^{18}$O J$=$2--1 data, respectively. The red dotted lines show the best-fit model from Section 4.3.1. }\\label{uvamp} \\end{center} \\end{figure*}  \\subsubsection{Disk structure probed by $^{13}$CO and C$^{18}$O}  \\begin{figure*} \\centering \\includegraphics[width=17cm]{bastadensbasta.ps} \\caption{Density structure of the best-fit disk model shown as colour scale. The disk is viewed by us from almost straight above. The three dotted lines show the heights above which 90\\% of the integrated line emission of $^{12}$CO, $^{13}$CO and C$^{18}$O $J$=2--1 originates. For the C$^{18}$O emission this line lies below the disk midplane.} \\label{denstemp} \\end{figure*}   We find fits to the observed visibilities of $^{13}$CO and C$^{18}$O line for M$\\mathrm{_{^{13}CO}}$=2.9$\\times$10$^{-7}$~M$_{\\odot}$ and M$\\mathrm{_{C^{18}O}}$=4.6$\\times$10$^{-8}$~M$_{\\odot}$, shown in Fig.~\\ref{uvamp}, bottom panels. For comparison, the visibilities of the $^{12}$CO J$=$2--1 line are also shown in Fig.~\\ref{uvamp}, top-right panel, with the best fit obtained by \\citet{raman} for M$\\mathrm{_{^{12}CO}}$=6$\\times$10$^{-6}$~M$_{\\odot}$ contained in the gas phase. In order to test our fit further, we invert and deconvolve the visibilities using the miriad reduction package and compare the observed (top panels) and modeled (middle panels) velocity channel maps in Figs.~\\ref{13maps} and \\ref{18maps}. The structure and intensity of the emission is found to be well matched. Figures ~\\ref{13maps} and ~\\ref{18maps} (lower panels) show the residual emission resulting from subtraction of the modeled from the observed data. For $^{13}$CO the residuals appear randomly distributed and do not exceed the noise level. Significant residuals are seen in C$^{18}$O at 7.4~km~s$^{-1}$ about 0$\\farcs$5 north from the star. The C$^{18}$O data (Fig. ~\\ref{18maps}) have a marginally significant localised excess $1\\farcs0$ north of the star seen at 3~$\\sigma$ level in the residual emission. This excess emission corresponds to the redshifted peak of the spectral line at v$_{\\rm LSR}$=7.4~km~s$^{-1}$. It is marginally resolved at $1\\farcs5\\times2\\farcs5$ and contains about 20\\% of C$^{18}$O line flux. Similar, but not spatially coincident residual emission is also seen in $^{13}$CO (Figure ~\\ref{13maps}, lower panel) near the noise levels, but not in $^{12}$CO. A cause of this feature may lie in disk midplane asymmetry. Any connection with the reported companion at 9$\\farcs$3 is unlikely, but a hypothesized body within the inner disk is an interesting prospect \\citep{grady}. Nevertheless, an instrumental or calibration artefact may have affected the 7.4~km~s$^{-1}$ channel resulting in a slight increase of flux in that channel in the lower sideband, only noticeable in the weak C$^{18}$O line.  In this section we will use our $^{13}$CO and C$^{18}$O results to make an estimate of the total gas mass M$_{gas}$ in the disk, related to the $^{12}$CO gas mass M$_{^{12}CO}$ through \\begin{equation} \\mathrm{M_{gas}=1.2\\times\\frac{m_{H_2}}{m_{^{12}CO}}\\times\\frac{1}{[^{12}CO]}\\times{M_{^{12}CO}}}. \\end{equation}  where $\\mathrm{m_{H_2}}$ and $\\mathrm{m_{^{12}CO}}$ are the masses of the H$_2$ and $^{12}$CO molecule, and [$^{12}$CO] is the abundance of $^{12}$CO with respect to H$_2$. The factor of 1.2 accounts for one fifth of the gas contained in helium. We can use the derived masses of the $^{13}$CO and C$^{18}$O gas to calculate the mass of the $^{12}$CO gas. We adopt the isotopic ratios of $^{12}$C/$^{13}$C=77$\\pm$7 and $^{16}$O/$^{18}$O=560$\\pm$25 from \\citet{wilson}. \\begin{equation} \\mathrm{M_{^{12}CO}=\\frac{m_{^{12}CO}}{m_{C^{18}O}}\\times\\frac{[^{16}O]}{[^{18}O]}\\times{M_{C^{18}O}}=(2.4\\pm1.0)\\times10^{-5}M_{\\odot}} \\end{equation} \\begin{equation} \\mathrm{M_{^{12}CO}=\\frac{m_{^{12}CO}}{m_{^{13}CO}}\\times\\frac{[^{12}C]}{[^{13}C]}\\times{M_{^{13}CO}}=(2.2\\pm0.6)\\times10^{-5}M_{\\odot}} \\end{equation} The values obtained here are given with errors which are dominated by the rms of our data and also include the errors on isotopic ratios. If we take into account that 8\\% of the $^{12}$CO and its isotopologues in HD169142 may be depleted onto dust grains and therefore not contributing to the emission observed, then the full $^{12}$CO mass range is (1.4-3.7)$\\times$10$^{-5}$~M$_{\\odot}$. The fact that both datasets roughly agree on the $^{12}$CO mass provides an additional argument for the validity of the model. However, the fit to the $^{12}$CO J$=$2--1 data done in \\citet{raman} corresponds to a mass of gas phase $^{12}$CO of 6$\\times$10$^{-6}$~M$_{\\odot}$ (mentioned in section 4.3), two to six times lower than the range derived here. The $^{12}$CO emission is optically thick and therefore it does not trace the entire disk but originates from a layer much higher in the disk than either the $^{13}$CO or C$^{18}$O lines \\citep{dartois}. The discrepancy between the $^{12}$CO result and that of the isotopologues indicates that the adopted disk model may not provide the best description in these high layers. We will discuss this in more detail in section 4.5 where the effects of temperature and turbulence on the line profile are discussed.   Our model calculations allow us to identify which regions of the disk contribute most to the observed lines. We follow the line of sight integration step by step from the observer into the disk and find the location where 90\\% of the integrated intensity is reached for a face on orientation of the disk. We show these locations in Fig.~\\ref{denstemp} in a cross-section of the disk showing the underlying density structure. We choose to use the intensity integrated over a range of wavelengths in which the line is emitted in order to better represent the region corresponding to the physical quantity that is observed, i.e., the line flux. The use of surfaces with equal opacity (often used $\\tau$=3) may be misleading, because the opacity at the line centre is not necessarily consistent with the opacity in the line wings, and the amount of flux still produced in the line wings beyond the $\\tau$=3 surface may contribute significantly to the line flux in some species more than in the others. The calculation was done for $^{12}$CO, $^{13}$CO and C$^{18}$O J$=$2--1 line. The discontinuous shape of the surface above which 90\\% of the integrated intensity in these lines is emitted is due to the discrete cells in the grid of our model that are used to sample the density and temperature structure. As expected, the less optically thick species trace regions deeper in the disk with C$^{18}$O tracing the disk midplane while $^{12}$CO traces the disk surface layer. We analyse the physical properties of the emitting regions in the outer disk, beyond 100~AU from the star (scales resolved by our SMA observations). For the two outermost radial positions in our model grid, we calculate the column of $^{12}$CO gas contained in the region where 90\\% of the emission is coming from, which is located above the vertical position shown in Fig.~\\ref{denstemp}, for each molecular line. These values are shown in Table~\\ref{table}\\footnote{The calculations for $^{12}$CO J$=$2--1 are done for the model presented in \\citet{raman}}. Also, the average temperature of these gas columns is calculated, taking into account the temperature weighted by the column density (i.e., mass) for each cell located above the vertical positions from Fig.~\\ref{denstemp}. The C$^{18}$O J$=$2--1 line generally traces three times as much CO gas as traced by the $^{13}$CO J$=$2--1 line in the outer disk while the $^{12}$CO J$=$2--1 line traces only a small fraction of the gas, a few percent of the column traced by C$^{18}$O. It is clearly seen that the C$^{18}$O traces deep into the disk, to the height of several AU at the far side of the midplane, showing that it is very sensitive to the total gas mass. The 90\\% of the C$^{18}$O J$=$2--1 line emission traces around 60\\% of the total gas column.  \\begin{table*} \\caption{The table shows the column of $^{12}$CO traced by 90\\% of the emission of each line, and the corresponding average temperature, at two radial intervals beyond 100~AU. These values are representative of the regions shown in Fig.~\\ref{denstemp}. The column densities of $^{12}$CO contained above the vertical locations z from Fig.~\\ref{denstemp} are shown. The average temperature values shown are column-weighted (i.e. mass-weighted) values over all cells which are located above the reported height z. } \\label{table} \\centering \\begin{tabular}{c | c c | c c c c | c c} \\hline\\hline Molecular & z(AU) & z(AU) & N$_{\\rm CO}$(cm$^{-2}$) & Fraction Of & N$_{\\rm CO}$(cm$^{-2}$) & Fraction Of & T$_{\\rm ave}$ (K) & T$_{\\rm ave}$ (K)\\\\ Line & 100-149~AU & 149-235~AU & 100-149~AU & Total N$_{\\rm CO}$  & 149-235~AU & Total N$_{\\rm CO}$ & 100-149~AU & 149-235~AU\\\\ \\hline $^{12}$CO J$=$2--1 & 12.7 & 39.0 & 6.5$\\times$10$^{17}$ & 7\\% & 6.6$\\times$10$^{16}$ & 1\\% &33.3 &38.6\\\\ $^{13}$CO J$=$2--1 & 7.6 & 13.0 & 4.2$\\times$10$^{18}$ & 17\\% & 2.6$\\times$10$^{18}$ & 18\\% & 29.1 & 25.6\\\\ C$^{18}$O J$=$2--1 & -2.5 & -4.3 & 1.5$\\times$10$^{19}$ & 62\\% & 8.9$\\times$10$^{18}$ & 62\\% & 25.7 & 21.8\\\\ \\hline \\end{tabular} \\end{table*}      In section 4.1 we already concluded from the disk model that the CO freeze-out could only be efficient in the cold outer disk midplane region which contains 8\\% of the total disk mass. This very small fraction is due to the luminosity of the Herbig Ae star HD169142 which is higher than the typical luminosity of T Tauri stars. In the latter, the circumstellar disks are colder and the CO freeze-out presents a major obstacle to determine gas mass based on observations of CO or its isotopologues. As a caveat, we note that \\citet{pietu3} find somewhat lower dust temperature in the disk around the Herbig Ae star MWC 480 (10-25~K at 100~AU). Since we have no indications to the contrary, we adopt the dust temperature that SED disk modeling suggests. If observed with a sufficiently high resolution and sensitivity, disks where CO freeze-out is efficient are expected to have smaller radii derived from CO isotopologues compared to the outer radius derived from the $^{12}$CO and from the dust emission. Although our observations are limited by the beam size (1$\\farcs$4$\\approx$200~AU) we have tested a range of different outer radii and found that, based on our molecular line data, the outer radius should not be very different from 200~AU for either of the observed lines. We also do not find any significant discrepancy between the molecular line data and dust continuum that could indicate different outer radii for gas and dust.  Another factor that can influence the $^{12}$CO abundance is photodissociation. The effect of photodissociation via the stellar UV radiation on CO abundance in circumstellar disks is studied in \\citet{vanzadelhoff} and \\citet{vanzadelhoffpaper}. They find that the CO abundance decrease due to photodissociation is significant only in the upper disk layers and that the CO column density is not much affected by it. \\citet{jonkheid} show that the stellar UV field is only able to heat 10\\% of the gas mass in the disk to considerably high temperatures, in the surface disk layer. The effect of photodissociation on the CO content of HD169142 is expected to be small, due to the significant dust mass of the disk. Since we adopt a maximum emissivity and therefore minimal dust mass, the conclusion that photoionization can be neglected is robust. In addition, the C$^{18}$O emission that we use to derive the gas mass, arises predominantly from the outer disk midplane region that contains the bulk of the disk mass and does not receive any direct UV illumination from the central star. This minimises further the effects of photodissociation on our results. Therefore, we find it appropriate to use the $^{12}$CO mass derived from our best fit to the C$^{18}$O and $^{13}$CO data to make an estimate of the total gas mass.  The abundance of $^{12}$CO is a key factor in this conversion and the values found in the literature for the interstellar medium range from 9.5$\\times$10$^{-4}$ \\citep{frerking} to 2.0$\\times$10$^{-4}$ \\citep{lacy}. If we adopt the $^{12}$CO abundance of 10$^{-4}$ with respect to H$_2$, we obtain $\\mathrm{M_{\\rm gas}=1.2\\times M_{\\rm H_2}=(2.0^{+1.0}_{-0.8})\\times10^{-2}M_{\\odot}.}$ The error estimate includes the error due to the rms of our C$^{18}$O data (41\\%) and to the isotopic ratio of [$^{16}$O]/[$^{18}$O] (4\\%). The upper margin is extended to include the 8\\% uncertainty in the gas mass estimate due to freeze-out. The factor 1.2 takes into account the gas mass contained in helium, the second most important gas mass constituent after molecular hydrogen. Adopting a $^{12}$CO abundance of 2.0$\\times$10$^{-4}$ with respect to H$_2$ results in M$_{\\rm gas}$ as low as 6$\\times$10$^{-3}$~M$_{\\odot}$. Therefore, due to the uncertainty largely dominated by the $^{12}$CO abundance, we can only conclude that the gas mass of the disk around HD169142 falls in the range of (0.6-3.0)$\\times$10$^{-2}$~M$_{\\odot}$.      \\subsection{Gas to dust ratio}  \\begin{figure} \\includegraphics[width=8cm,angle=0]{3x.ps} \\caption{Comparison of the observed vector-averaged $^{12}$CO ({\\it{top panel}}), $^{13}$CO ({\\it{middle panel}}) and C$^{18}$O ({\\it{lower panel}}) line fluxes (black) to four models with different gas to dust ratios (red lines). The error bars represent the variance within each annular average. The models assume the dust mass of 2.16$\\times$10$^{-4}$~M$_{\\odot}$ and gas to dust mass ratio f$_{\\rm g/d}$=100,50,33,25 marked respectively by the full, dashed, dashed-dotted and dotted red lines in each panel. The line fluxes are integrated over the width of $\\approx$2.7~km~s$^{-1}$ (10 spectral channels) centered on each line. }\\label{uvamp2} \\end{figure}  In order to investigate the gas to dust ratio, we ran several models in which we assume M$\\mathrm{_{\\rm dust}}$=2.16$\\times$10$^{-4}$~M$_{\\odot}$, [CO]=10$^{-4}$, $^{12}$C/$^{13}$C=77, $^{16}$O/$^{18}$O=560 and gas to dust mass ratios of 100, 50, 33 and 25. From the comparison of the synthetic visibilities of these models with the observations (Fig.~\\ref{uvamp2}, middle and lower panel) we find that a good fit for $^{13}$CO and C$^{18}$O data is obtained for f$_{\\rm g/d}$=100. Under the assumptions used above, lower values of f$_{\\rm g/d}$ fail to provide a fit to the data. The $^{12}$CO data (Fig.~\\ref{uvamp2}, top panel) shows a discrepancy, already discussed in section 4.3.1, with respect to the $^{13}$CO and C$^{18}$O data in the sense that the model we use overestimates the $^{12}$CO emission. This suggests that the model assumptions like the temperature, micro turbulence and $^{12}$CO abundance used for the upper disk layers where the $^{12}$CO emission arises from should be reconsidered. From the gas and dust masses derived in sections 4.2 and 4.3 we see that a gas to dust mass ratio of 100:1 only holds for special circumstances considered here: a maximum dust emissivity of 2~cm$^2$g$\\mathrm{^{-1}_{\\rm dust}}$ at 1.3~mm and a $^{12}$CO abundance of 10$^{-4}$ with respect to H$_2$, or lower. For equally likely values of 1~cm$^2$g$\\mathrm{^{-1}_{\\rm dust}}$ for the dust emissivity and 2$\\times$10$^{-4}$ for the $^{12}$CO abundance, a gas to dust mass ratio of 25 would be inferred, suggesting that the disk would have lost significant amounts of gas with respect to dust. In fact, gas loss is expected in circumstellar disks due to photoevaporation of gas by the stellar radiation. The typical timescale for this process of a few Myr \\citep{hollenbach,font} is similar to the estimated age of HD169142 of 6$^{+6}_{-3}$~Myr \\citep{grady}.  \\subsection{Micro turbulence}  As mentioned in the previous section, the overestimate of the $^{12}$CO emission by the model which fits well the $^{13}$CO and C$^{18}$O $J$=2--1 line emission, may be due to model assumptions used for the upper disk layers, e.g., temperature, $^{12}$CO relative abundance, turbulence.   In section 4.1 we mention that a disk model with a less luminous A6 spectral type star also reproduces the SED of HD169142 well \\citep{dentHD}. However, the temperature in that model is less than 10\\% lower than the temperature used in our model. This is insufficient to provide a significant decrease in the $^{12}$CO line flux.  A fit to all observed lines simultaneously by varying the isotopic ratios, yields an unusualy low value of both the $^{12}$C/$^{13}$C and the $^{16}$O/$^{18}$O isotopic ratio: 21 and 171, respectivelly. These values are far from the observationally derived ratios in the interstellar medium and this scenario would be highly unlikely. An over-abundance of $^{13}$CO was measured at large radial distances from the star (several hundreds of AU) in the disks DM Tau, MWC 480 and LkCa 15 \\citep{pietu1} In general, disks are often found to have $^{13}$CO and C$^{18}$O abundances depleted by large factors, mainly in disks around T Tauri stars \\citep{dutrey1,dutrey2}.  In the above analysis, the turbulent line FWHM was set to 0.16~km~s$^{-1}$ throughout the disk. In a more realistic scenario, the turbulence in the cold and dense midplane differs from the turbulence in the warmer and less dense surface layers. We investigate the effect of micro turbulence on the modeled line emission and find that a simultaneous match is possible in our model, as shown in Fig.~\\ref{turb}. This fit is obtained for a gas to dust ratio of 33, i.e., a disk gas mass of 7$\\times$10$^{-3}$~M$_{\\odot}$, if a higher value of turbulence is assumed for the disk midplane (turbulent FWHM of 0.33~km~s$^{-1}$) and no turbulence for the upper disk layers traced by the $^{12}$CO emission (0.00~km~s$^{-1}$, only thermal broadening). For the $^{12}$CO J$=$2--1 line, that is fully optically thick at all radii in our model, the turbulent line broadening is degenerate with temperature and density in its effect on the line flux and profile. This is because the turbulence additionally broadens the thermally broadened line profile and in this way the optical thickness decreases allowing for more gas column to be traced. On the other hand, by decreasing the turbulence, at a given temperature, the same $^{12}$CO data can be fit with comparably more M$_{\\rm CO}$. The effect of the increase in turbulence on the $^{13}$CO and C$^{18}$O lines is present but less pronounced, due to their lower optical thickness. In particular, the C$^{18}$O $J$=2--1 line emission is only marginally optically thick. The effect of turbulence on the modeled C$^{18}$O $J$=2--1 line emission is below the rms of our observations. It is not possible to obtain simultaneous fits to all data if the gas to dust ratio is higher than 33 (disk mass higher than 7$\\times$10$^{-3}$~M$_{\\odot}$). A lower gas to dust ratio (disk mass lower than 7$\\times$10$^{-3}$~M$_{\\odot}$) would require the turbulent FWHM of the lines to be larger than 0.33~km~s$^{-1}$ in disk midplane.   \\begin{figure} \\includegraphics[width=8cm,angle=0]{3xvt.ps} \\caption{Comparison of the observed vector-averaged $^{12}$CO ({\\it{top panel}}), $^{13}$CO ({\\it{middle panel}}) and C$^{18}$O ({\\it{lower panel}}) line fluxes (black) to a model which assumes the total disk mass of 7$\\times$10$^{-3}$~M$_{\\odot}$, gas to dust mass ratio f$_{\\rm g/d}$=33, no microturbulent line broadening in the upper disk layers from which the $^{12}$CO emission originates (only thermal broadening), and a microturbulent broadening of 0.33~km~s$^{-1}$ FWHM in the layers from which the $^{13}$CO and C$^{18}$O emission arises.}\\label{turb} \\end{figure}"
    },
    "0806/0806.4380_arXiv.txt": {
        "abstract": "We present VULCAN/2D multi-group flux-limited-diffusion radiation hydrodynamics simulations of binary neutron star (BNS) mergers, using the Shen equation of state, covering $\\sgreat$100\\,ms, and starting from azimuthal-averaged 2D slices obtained from 3D SPH simulations of Rosswog \\& Price for 1.4\\,\\mo (baryonic) neutron stars with no initial spins, co-rotating spins, and counter-rotating spins. Snapshots are post-processed at 10\\,ms intervals with a multi-angle neutrino-transport solver. We find polar-enhanced neutrino luminosities, dominated by $\\bar{\\nu}_e$ and ``$\\nu_\\mu$'' neutrinos at peak, although $\\nu_e$ emission may be stronger at late times. We obtain typical peak neutrino energies for $\\nu_e$, $\\bar{\\nu}_e$, and ``$\\nu_\\mu$'' of $\\sim$12, $\\sim$16, and $\\sim$22\\,MeV. The super-massive neutron star (SMNS) formed from the merger has a cooling timescale of $\\sles$1\\,s. Charge-current neutrino reactions lead to the formation of a thermally-driven bipolar wind with <$\\dot{M}$>$\\sim$10$^{-3}$\\,\\mo s$^{-1}$, baryon-loading the polar regions, and preventing any production of a $\\gamma$-ray burst prior to black-hole formation. The large budget of rotational free energy suggests magneto-rotational effects could produce a much greater polar mass loss. We estimate that $\\sles$10$^{-4}$\\,\\mo of material with electron fraction in the range 0.1-0.2 become unbound during this SMNS phase as a result of neutrino heating. We present a new formalism to compute the $\\nu_i\\bar{\\nu}_i$ annihilation rate based on moments of the neutrino specific intensity computed with our multi-angle solver. Cumulative annihilation rates, which decay as $t^{-1.8}$, decrease over our 100\\,ms window from a few $\\times$10$^{50}$ to $\\sim$10$^{49}$\\,erg\\,s$^{-1}$, equivalent to a few $\\times$10$^{54}$ to $\\sim$10$^{53}$ $e^-e^+$ pairs per second. ",
        "introduction": "Coalescing neutron stars are one of the primary progenitor candidates for short-duration (i.e., $\\sles$\\,2\\,s) $\\gamma$-ray bursts (GRBs; \\citealt{paczynski:86,goodman:86,eichler:89,narayan:92,mochkovitch:93}). In-spiral of the two neutron star components occurs due to energy and angular momentum loss through gravitational radiation \\citep{taylor:82,weisberg:05}, which is emitted as the system evolves towards coalescence. Modern $\\gamma$-ray and X-ray satellites have considerably improved our understanding of short-hard GRBs \\citep{nakar:07}, and have more strongly associated them with binary neutron star (BNS) merger events \\citep{lee:07}. Gravitational wave detectors such as LIGO \\citep{abramovici:92} will most likely improve our understanding of these events in the coming decade.  The engine powering such short-hard GRBs and their associated high-Lorentz-factor ejecta is thought to be linked to mass accretion from a torus onto a central black hole, formed when the super-massive neutron star (SMNS) born from the coalescence eventually collapses. In recent years, models of GRB production from a high-mass, magnetar-like, neutron star have also received some attention \\citep{usov:92,usov:94,kluzniak:98,rosswog:03c,dai:06,metzger:07,dessart:07}. We will address the likelyhood that the SMNS produces a GRB prior to black-hole formation in \\S~\\ref{sect:dynamics}. A fraction of the gravitational energy in the disk around this black hole is radiated as neutrinos which power a disk wind. A fraction of these neutrinos and antineutrinos annihilate into $e^-e^+$. Magnetic processes may, through their role in angular momentum transport, be decisive in setting the time scale between merger and collapse and, in addition, they may extract rotational energy from the central black hole \\citep{blandford:77}. The ejecta may be confined by the magnetic field morphology or by the intense neutrino-driven baryon-loaded wind thought to accompany coalescence and black hole mass accretion \\citep{levinson:00,rosswog:03b,aloy:05}.  Numerical simulations of BNS mergers have a rich history, and have been considerably refined over the years, moving from 2D to 3D; from Newtonian to post-Newtonian, conformally flat, and finally to full General Relativity (GR); from simple polytropic to more sophisticated equations of state (EOS); from the neglect of neutrinos to approximate neutrino-trapping schemes. BNS merger studies yield predictions for 1) their associated gravitational wave signals \\citep{ruffert:96,calder:02,faber:02,shibata:02,shibata:03,shibata:05,oechslin:07a,anderson:08a}, 2) the production of short-hard GRBs \\citep{ruffert:97,ruffert:99,rosswog:03b,rosswog:03c,rosswog:05,janka:06,birkl:07}, and 3) the pollution of the environment by r-process nuclei \\citep{lattimer:74,lattimer:76,symbalisty:82,eichler:89,davies:94,freiburghaus:99,rosswog:99a,surman:08}.  Different modeling ingredients and approaches have been employed. 3D Newtonian simulations without neutrino transport were performed by \\cite{oohara:89,oohara:90}, \\cite{nakamura:89,nakamura:91}, \\cite{zhuge:94,zhuge:96}, and \\cite{rosswog:00}. \\cite{ruffert:96,ruffert:97} included neutrinos using a grid-based code, and this was followed by \\cite{rosswog:02a} and \\cite{rosswog:03a} using a 3D SPH code. Recently, \\cite{setiawan:04,setiawan:06} applied such a neutrino-leakage scheme in their study of torus-disk mass accretion around a black hole formed in a BNS merger, and addressed neutrino emission and annihilation. Magnetic fields in BNS merger evolution were introduced in 3D Newtonian simulations by \\cite{price:06} and \\cite{rosswog:07a}.  In parallel, there have been improvements to the Newtonian approach to incorporate the effects of GR (although largely neglecting microphysics), which become important as the two neutron stars come closer and eventually merge. Post-Newtonian simulations were performed by \\cite{oohara:92}, \\cite{faber:00,faber:02}, \\cite{faber:01}, \\cite{ayal:01}, and \\cite{calder:02}. The conformally-flat approximation was introduced in \\cite{wilson:96} and followed by \\cite{oechslin:02}, \\cite{faber:04}, \\cite{oechslin:06,oechslin:07a}, and \\cite{oechslin:07b}. Full 2D GR simulations, sometimes including magnetic fields,  have been performed for super-massive neutron stars (SMNS), possibly resulting from BNS mergers, by \\cite{baumgarte:00}, \\cite{morrison:04}, \\cite{duez:04,duez:06}, and \\cite{shibata:06}, while full 3D GR simulations of the merger were carried out by \\cite{shibata:02,shibata:03,shibata:05,shibata:07}, \\cite{anderson:08b}, and \\cite{baiotti:08}.  All these simulations have been conducted at various levels of sophistication for the thermodynamic properties of matter, ranging from simplistic and not so physically-consistent polytropic EOSs, to those employing a detailed microphysical representation of nuclear matter at an arbitrary temperature \\citep{lseos:91,shen:98a,shen:98b}. Investigations performed with such detailed EOSs have been conducted by, e.g., \\cite{ruffert:96,ruffert:97}, \\cite{rosswog:99a}, and \\cite{rosswog:02a}, and the dependency of BNS merger properties on the adopted EOS has been discussed by \\cite{oechslin:07a} and \\cite{oechslin:07b}. Variations of a few times 10\\% in the maximum allowed mass are seen, depending on the compressiblity of nuclear matter, but recent observations may suggest that neutron stars with a gravitational mass in excess of 2\\,\\mo do exist \\citep{nice:05,freire:07}, supporting a stiff EOS for nuclear matter, such as the Shen EOS we employ here. In Fig.~\\ref{fig_BH}, we show, as a function of the central density, the baryonic and gravitational neutron star masses that obtain for our implementation of the Shen EOS. The maximum occurs at $\\sim$1.3$\\times$10$^{15}$\\,g\\,cm$^{-3}$, corresponding to a gravitational (baryonic) mass of 2.32\\,\\msun\\, (2.77\\,\\msun).\\footnote{Note that the values given in \\citealt{dessart:08} are not exact.} Moreover, in the context of differentially rotating (and possibly magnetized) SMNSs, the maximum mass that can be supported by a given EOS may be increased by up to $\\sim$50\\% compared to the equivalent non-rotating object \\citep{baumgarte:00,morrison:04,duez:04,duez:06,shibata:06,sumiyoshi:07,kiuchi:08}. Ultimately, understanding the mechanisms and timescales for angular momentum to be redistributed to lead to solid-body rotation is, therefore, important. One possible agency of redistribution is the magneto-rotational instability \\citep{bh:91,akiyama:03,pessah:06}. The implications are non-trivial because the efficiency of angular-momentum transport determines in part whether a black hole forms promptly, or after a short or a long delay. It also determines how much mass can be accreted, i.e., whether there is $\\sim$0.01 or $\\sim$0.1\\,\\mo available in the torus surrounding the black hole after it has formed, and what the timescale is over which such accretion can take place to power relativistic ejecta. Although indirect, the relevance to short-hard GRBs and their properties is obvious, and perhaps central. In fact, this issue is also germane to the production of collapsars and long-duration GRBs \\citep{dessart:08}.  \\begin{figure} \\plotone{f1.eps} \\caption{Gravitational (solid line) and baryonic (dotted line) masses, as well as the corresponding neutron star radius (dashed line), versus central density derived for our Shen EOS and the Oppenheimer-Volkov equation (one form of the general relativistic equation of hydrostatic equilibrium), and assuming a non-rotating neutron star characterized by a uniform electron fraction of 0.1 and a uniform temperature of 1\\,MeV. Here, the transition to a black hole will occur when the central density in the neutron star reaches $\\sim$1.3$\\times$10$^{15}$\\,g\\,cm$^{-3}$, corresponding to a gravitational (baryonic) mass of 2.32\\,\\msun\\, (2.77\\,\\msun) and a neutron star radius of $\\sim$13\\,km. \\label{fig_BH}} \\end{figure}  In this work, we perform 2D multi-group, flux-limited-diffusion (MGFLD), radiation hydrodynamics of merged BNSs using the Shen EOS. The main goal of this work is to document in detail the neutrino signatures from such merger events. Our approach is to solve the neutrino transport problem self-consistently (although with a flux-limiter and assuming diffusion), in combination with the dynamics of the system (but assuming axisymmetry), over $\\sgreat$100\\,ms. At selected times, we post-process such models with the multi-angle, $S_n$, neutrino transport algorithm of \\cite{livne:04} and described recently in \\cite{ott:08}. In particular, our investigation allows for the first time the computation of the neutrino-antineutrino annihilation rate from a solution of the transport equation, rather than using an approximate leakage scheme. Our investigation applies strictly to the neutron-star phase of such BNS mergers.  This paper is organized as follows. In \\S~\\ref{sect:model}, we present the initial models we employ in our work, based on 3D SPH simulations of BNS mergers with different spin configurations (\\S~\\ref{sect:prog}). In \\S~\\ref{sect:vulcan}, we present the characteristics and procedures we employ to evolve such initial conditions with the MGFLD radiation hydrodynamics code VULCAN/2D. In \\S~\\ref{sect:results}, we present our results for the neutrino signatures (\\S~\\ref{sect:neutrinos}) and the dynamics of BNS mergers (\\S~\\ref{sect:dynamics}). In \\S~\\ref{sect:annihilate}, we address the neutrino-antineutrino annihilation rate in our three BNS merger models. We first present results employing the approach used so far \\citep{ruffert:96,ruffert:97,rosswog:03a,setiawan:04,setiawan:06,birkl:07}, based on a leakage scheme for the neutrino-flavor dependent opacities and emissivities and a summation of all paired grid cells contributing at any given location. In \\S~\\ref{annihil:sn}, we then present a new formalism which uses moments of the neutrino specific intensity computed with a multi-angle, $S_n$, scheme. Finally, we present our conclusions in \\S~\\ref{sect:conclusion} and discuss the most striking implications for the powering of short-duration GRBs. For completeness, in the appendix, we apply our formalism to the computation of the annihilation rate in the context of slow- and fast-rotating single protoneutron stars (PNSs).  \\begin{figure*} \\epsscale{0.35} \\plotone{f2a.eps} \\plotone{f2b.eps} \\plotone{f2c.eps} \\plotone{f2d.eps} \\plotone{f2e.eps} \\plotone{f2f.eps} \\caption{ Colormap of the temperature (MeV; log scale) at 0 ({\\it top row}; inner 150$\\times$150\\,km$^2$ region) and 100\\,ms ({\\it bottom row}; inner 250$\\times$250\\,km$^2$ region) after the start of the simulation. In the bottom panels, velocity vectors are added, with a saturation length at 15\\% the width of the display and corresponding to a velocity of 30,000\\,\\kms. From left to right columns, we show the no-spin, co-rotating, and counter-rotating models. We also overplot white density contours for every decade, starting at a maximum of 10$^{14}$\\,g\\,cm$^{-3}$. Note that the minimum of the colorbar changes between the top and the bottom rows. } \\label{fig_temp_evol} \\end{figure*} ",
        "conclusions": "\\label{sect:conclusion}   We have presented multi-group flux-limited-diffusion radiation hydrodynamics simulations of binary neutron star mergers, starting from azimuthal-averaged 2D slices based on the 3D SPH simulations produced with the MAGMA code \\citep{rosswog:07a}, for neutron star components with initially no spins, co-rotating spins, and counter-rotating spins. The main virtues of this work are 1) the solution of the radiation transport problem for $\\nu_e$, $\\bar{\\nu}_e$, and ``$\\nu_\\mu$'' neutrinos for eight energy groups, coupled to the hydrodynamical evolution of such mergers over a typical timescale of $\\sgreat$100\\,ms after the two components come into contact, and 2) the first quantitative assessment of baryon-pollution by the neutrino-driven wind produced by such SMNSs. Interestingly, our results on neutrino signatures from such merger evens confirm the broad adequacy of previous work that avoided solving the radiation transport problem by designing a neutrino-trapping, or leakage, scheme \\citep{ruffert:96,ruffert:97,rosswog:03a}. The only noticeable discrepancy is the much stronger ``$\\nu_\\mu$'' luminosities we predict with VULCAN/2D, something we associate with nucleon-nucleon bremsstrahlung not currently accounted for in the above leakage schemes.  \\begin{figure} \\epsscale{1.2} \\plotone{f19.eps} \\caption{ Variation of the log of the cumulative annihilation rate $\\int^{\\theta_{\\rm max}} dV Q^+(\\nu\\bar{\\nu})$ with polar angle $\\theta_{\\rm max}$ in the BNS merger models initially with no spin ({\\it solid}), co-rotating spins ({\\it dotted}), and counter-rotating spins ({\\it dashed}), and shown here at a time of 10 (black) and 60\\,ms (red) after the start of the VULCAN/2D simulation. In other words, $\\theta_{\\rm max}$ is the half-opening angle of the cone that defines the integration volume for the quantity plotted along the ordinate axis. We use the $S_n$ results computed with 8 $\\vartheta$-angles and accounting only for a density cut for the sites of emission and deposition. Energy deposited in the cooling region is not subtracted off, yielding values larger by a factor of two compared with the corresponding values given in Table \\ref{tab:edep}. } \\label{fig:edep_vs_theta} \\end{figure}  At 10\\,ms intervals and for each BNS merger model, we select a sequence of VULCAN/2D snapshots computed with the MGFLD solver and post-process them with the multi-angle $S_n$ solver, relaxing the radiation variables, but freezing the fluid variables. Based on a knowledge of the energy-dependent, species-dependent, and angle-dependent neutrino specific intensity $I_\\nu(\\varepsilon_\\nu,\\vec{n})$, we compute the neutrino-antineutrino annihilation rate using a new formalism that incorporates various moments of $I_\\nu(\\varepsilon_\\nu,\\vec{n})$. We find that the total annihilation rate computed with the $S_n$ solution and our new formalism is larger, but at most by a factor of a few, compared to the results based on a combination of a leakage scheme \\citep{ruffert:96} and paired-cell summation \\citep{ruffert:97}. With density cuts alone, the annihilation rates based on the $S_n$ solution increase by a factor of about two. In our simulations, we find that all neutrino luminosities decrease after peak, resulting in a decrease in the annihilation rate with time, i.e., $\\propto t^{-1.8}$. We find a cumulative total rate that decreases over 100\\,ms from a few $\\times$10$^{50}$ to $\\sim$10$^{49}$\\,erg\\,s$^{-1}$. Adopting an average energy of $8kT\\sim$40\\,MeV for the annihilating $\\nu_e\\bar{\\nu}_e$ \\citep{brt:06}, the number of $e^-e^+$ pairs produced varies over that time span from a few $\\times$10$^{54}$ down to $\\sim$10$^{53}$\\,s$^{-1}$.  We show in Fig.~\\ref{fig:edep_vs_theta} the cumulative annihilation rate computed with the $S_n$ solution and our new formalism at 10 (black) and 60\\,ms (red) after the start of the VULCAN/2D simulations. Only $\\sim$1\\% of the total is deposited along the rotation axis and in a cone with an half-opening angle of $\\sim$10$^{\\circ}$, yielding a rate of $\\sim$10$^{49}$\\,erg\\,s$^{-1}$ at 10\\,ms but down to $\\sim$10$^{48}$\\,erg\\,s$^{-1}$ at 60\\,ms. Integrating over a tenth of a second, we obtain a total energy deposition of $\\sles$10$^{48}$\\,erg. This result is about an order of magnitude smaller than obtained by \\cite{setiawan:04,setiawan:06}, although in their approach, neutrino emission is considerably boosted by shear-heating in the differentially-rotating disk. Their annihilation rates are stronger and weaken slightly more gradually with time, i.e., $\\propto t^{-1.5}$. In the VULCAN/2D simulations presented here, we have no physical viscosity to include shear heating, nor is there adequate resolution to simulate the magneto-rotational instability. Our temperatures and neutrino emissions are therefore underestimates of what would occur in Nature. With viscous dissipation as an energy source, we anticipate that our results might yield more attractive powers for the generation of a relativistic $e^-e^+$-pair jet, an issue we defer to a future study.  Baryon loading of relativistic ejecta is a recognized problem in the production of a GRB, but this is the first time the neutrino-driven wind, the baryon polluter, has been simulated dynamically. All three BNS merger simulations we performed reveal the birth of a thermally-driven wind through neutrino energy deposition in a gain layer at the surface, and primarily at high latitudes. The wind blows mostly along the polar funnel, contained in a cone with an opening angle of $\\sim$20$^{\\circ}$ within $\\sles$500\\,km, but widening at larger distances to $\\sim$90$^{\\circ}$. The electron fraction of the associated ejecta is close to 0.5, along the pole, but is $\\sim$0.1-0.2 along all lower latitudes. Low-$Y_{\\rm e}$ material at near-equatorial latitudes and large distances has a radial velocity below the local escape speed, and since local pressure gradients are not negligible, it is unclear whether it will eventually escape. Overall, it appears that if this pre-black-hole phase lasts for 100 milliseconds, $\\sles$10$^{-4}$\\,\\mo of ``r-process material'' will feed the interstellar medium through this neutrino-driven wind. Moreover, such baryonic pollution along the polar direction may affect any subsequent relativistic ejecta, although the wind from the SMNS may last for only a short time, perhaps 100\\,ms. Travelling at a tenth the speed of light, it will reach only out to 0.01 light-second, thus much shorter than the duration of short GRBs. This initial wind phase cannot provide the sustained confinement needed for the potential relativistic ejecta. Although we focus on the phase prior to black hole formation, a neutrino-driven wind may blow after the black hole forms too, but this time from the surrounding hot and dense torus, and driven by charge-current neutrino absorption at the disk surface. Importantly, because of the centrifugal barrier felt by particles coming in from sizable orbits in the disk, these injected baryons should remain away from the polar region, while neutrino-antineutrino annihilation would continue to contribute along the rotation axis of the SMNS, but now in an essentially baryon-free environment. This provides an attractive mechanism for confinement, since neutrino emission would yield both the power source for the relativistic ejecta (caused by neutrino-antineutrino annihilation) and the confined neutrino-driven wind (caused by charge-current reactions). In the future, we will investigate the neutrino signatures that obtain in the context where the compact object is a black hole, and only the surrounding torus material radiates neutrinos. Such a disk wind would offer a natural confining ingredient for any relativistic ejecta propelled along the rotation axis and powered by annihilation of neutrinos and antineutrinos radiated from the disk, as proposed by \\cite{mochkovitch:93}.  The huge amount of free energy of rotation, on the order of 10$^{52}$\\,erg in the supermassive neutron star and its torus disk, together with millisecond orbital periods, suggest that the magneto-rotational instability could play a major role in redistributing angular momentum, leading to solid-body rotation, and increasing the magnetic pressure in the corresponding layers by orders of magnitude. Magneto-rotational effects should be strong and would lead to a considerably stronger neutrino-driven wind, as in the AIC of white dwarfs \\citep{dessart:07}, and delay the formation of a black hole. Similarly, \\cite{dessart:08} found that in collapsar candidates \\citep{woosley:06,yoon:06,meynet:07}, the large amount of free energy of rotation in their iron core at the time of collapse can lead to a magnetically-driven explosion, associated with very large mass loss rates which may compete with the mass accretion rate from the torus disk and jeopardize the formation of a black hole. By contrast with such collapsar models, supermassive neutron stars formed from merger events are unambiguously endowed with a large rotational energy budget (stored in the orbital motion), of which a large fraction is differential and may be tapped. Hence, the phase prior to black hole formation should be modeled at very high resolution and in combination with neutrino transport to address these issues. The maximum neutron star mass allowed by the EOS will ultimately determine how much mass accretion can take place from the torus-disk, and it thus represents a central question for short-duration GRBs. Presently, there is a lack of consensus that leaves this issue largely unsettled."
    },
    "0806/0806.2179_arXiv.txt": {
        "abstract": "Cosmological shock waves are induced during hierarchical formation of large-scale structure in the universe. Like most astrophysical shocks, they are collisionless, since they form in the tenuous intergalactic medium through electromagnetic viscosities. The gravitational energy released during structure formation is transferred by these shocks to the intergalactic gas as heat, cosmic-rays, turbulence, and magnetic fields. Here we briefly  describe the properties and consequences of the shock waves in the context of the large-scale structure of the universe. ",
        "introduction": "Shock waves are ubiquitous in astrophysical environments; from solar winds to the largest scale of the universe \\citep{rkhj03}. In the current paradigm of the cold dark matter (CDM) cosmology, the large-scale structure of the universe forms through hierarchical clustering of matter. Deepening of gravitational potential wells causes gas to move supersonically. Cosmological shocks form when the gas accretes onto clusters, filaments, and sheets, or as a consequence of the chaotic flow motions of the gas inside the nonlinear structures. The gravitational energy released during the formation of large-scale structure in the universe is transferred by these shocks to the intergalactic medium (IGM).  Cosmological shocks are collisionless shocks which form in a tenuous plasma via collective electromagnetic interactions between baryonic particles and electromagnetic fields \\citep{quest88}. They play key roles in governing the nature of the IGM through the following processes: in addition to the entropy generation, cosmic-rays (CRs) are produced via diffusive shock acceleration (DSA) \\citep{bell78,bo78}, magnetic fields are generated via the Biermann battery mechanism \\citep{kcor97} and Weibel instability \\citep{msk06} and also amplified by streaming CRs \\citep{bell04}, and vorticity is generated at curved shocks \\citep{binn74}.  \\begin{figure*}[t] \\vspace{0.2cm} \\begin{center}\\includegraphics[scale=0.72]{f1.eps}\\end{center} \\vspace{0cm} \\caption{Two-dimensional images showing x-ray emissivity (top left), locations of shocks with color-coded shock speed $V_s$ (top right), perpendicular component of vorticity (bottom left), and magnitude of vorticity (bottom right) in the region of $(25\\ h^{-1}{\\rm Mpc})^2$ around a galaxy cluster at present ($z=0$). Color codes $V_s$ from 15 (green) to 1,800 km s$^{-1}$ (red).} \\end{figure*}  Cosmological shocks in the intergalactic space have been studied in details using various hydrodynamic simulations for the cold dark matter cosmology with cosmological constant ($\\Lambda$CDM) \\citep{rkhj03,psej06,krco08}. In this contribution, we describe the properties of cosmological shocks and their implications for the intergalactic plasma from a simulation using a PM/Eulerian hydrodynamic cosmology code \\citep{rokc93} with the following parameters: $\\Omega_{BM}=0.043$, $\\Omega_{DM}=0.227$, and $\\Omega_{\\Lambda}=0.73$, $h \\equiv H_0$/(100 km/s/Mpc) = 0.7, and $\\sigma_8 = 0.8$. A cubic region of comoving size 100 $h^{-1}$ Mpc was simulated with $1024^3$ grid zones for gas and gravity and $512^3$ particles for dark matter, allowing a uniform spatial resolution of $\\Delta l = 97.7h^{-1}$ kpc. The simulation is adiabatic in the sense that it does not include radiative cooling, galaxy/star formation, feedbacks from galaxies/stars, and reionization of the IGM. A temperature floor was set to be the temperature of cosmic background radiation. ",
        "conclusions": "Shocks are inevitable consequences of the formation of large-scale structure in the universe. They heat gas, accelerate cosmic-ray particles, produce vorticity and turbulence, and generate and amplify magnetic fields in the IGM. By applying detailed models of the DSA and the turbulence dynamo to the data of a cosmological hydrodynamic simulation of a concordance $\\Lambda$CDM universe, we have made the following quantitative estimates: $\\varepsilon_{\\rm CRp} / \\varepsilon_{\\rm therm} \\sim 0.4$ in the IGM, $\\left<\\varepsilon_{\\rm turb} /\\varepsilon_{\\rm therm}\\right>_{\\rm mass} = 0.1 - 0.3$ in the ICM, $\\left<\\varepsilon_{\\rm turb}/ \\varepsilon_{\\rm therm}\\right> \\sim 0.5 - 2$ in filaments and sheets, $B \\ga 1 \\mu$G inside clusters/groups, $B \\sim 0.1 \\mu$G around clusters/groups, and $B \\sim 10$ nG in filaments at present. Our results suggest that the non-thermal components can be energetically significant  in the intergalactic plasma of large-scale structure, as in the interstellar plasma inside Our Galaxy."
    },
    "0806/0806.2981_arXiv.txt": {
        "abstract": "The H.E.S.S. array of Cherenkov telescopes has performed, from 2004 to 2007, a survey of the inner Galactic plane at photon energies above 100 GeV. About 400 hours of data have been accumulated in the region between -30 and +60 degrees in Galactic longitude, and between -3 and +3 degrees in Galactic latitude. Assuming that \\DM\\, is composed of Weakly Interacting Massive Particles, we calculate here the H.E.S.S. sensitivity map for \\DM\\, annihilations, and derive the first experimental constraints on the % $^{\\backprime\\backprime}$mini-spikes$^{\\prime\\prime}$ scenario, in which a gamma-ray signal arises from \\DM\\, annihilation around Intermediate Mass Black Holes. The data exclude the proposed scenario at a 90\\% confidence level for \\DM\\, particles with velocity-weighted annihilation cross section $\\sigma v$ above $\\rm 10^{-28}\\,cm^3s^{-1}$ and mass between 800 GeV and 10 TeV. ",
        "introduction": "A substantial body of cosmological and astrophysical measurements suggests that $\\sim$22\\% of the Universe is composed of non-baryonic \\DM\\, (DM), e.g.~\\cite{spergel}, commonly assumed to be in the form of Weakly Interacting Massive Particles (WIMPs) arising in extensions of the Standard Model of Particle Physics (for recent reviews see~\\cite{bergstrom0,bertone0}). The lightest neutralino arising in supersymmetric (SUSY) extensions of the Standard Model~\\cite{jungman}, and the first excitation of the Kaluza-Klein bosons, $\\tilde{B}\\rm ^{(1)}$, in universal extra dimension (UED) theories~\\cite{appelquist,cheng,servant} are among the most widely discussed \\DM\\, candidates.   Although accelerator and direct \\DM\\, searches may well provide useful hints on the nature of \\DM, a correct identification is likely to require the combination of different techniques, including the indirect searches, based on the detection of the annihilation products of \\DM\\, particles. The annihilation rate being proportional to the square of the \\DM\\, density integrated along the line of sight, regions with enhanced \\DM\\, density are primary targets of indirect searches. Among them are the Galactic halo~\\cite{silk}, external galaxies~\\cite{baltz0}, galaxy clusters~\\cite{fornengo}, substructures~\\cite{calcaneo,tasitsiomi,stoehr,khousiappas0,baltz1,pieri0,koushiappas,diemand,Pieri:2007ir}, and the Galactic Center (GC)~\\cite{aharonian0,zaharijas,mambrini,profumo,cesarini}.   The GC, in particular, has attracted  significant interest. The distribution of \\DM\\, at the GC is actually highly uncertain, due to lack of resolution in N-body simulations, and to the many astrophysical effects that further complicate the situation, such as the presence of the supermassive black hole coincident with Sgr A*, gravitational scattering of \\DM\\, off the stellar cusp, and \\DM\\, annihilation~\\cite{bertone1}.  Dwarf spheroidal galaxies in the Local Group~\\cite{mateo} have also been considered  as targets for gamma-ray detection since they seem to represent dark matter dominated regions~\\cite{evans,profumo1,sanchez,bergstrom,strigari,aharonian1}. More recently, the Sloan Digital Sky Survey (SDSS) revealed the existence of new satellites~\\cite{willman,irwin,walsh} offering appealing features for \\DM\\, searches. Radial velocity dispersion of stars in Galaxy satellites usually  implies large mass-to-luminosity ratio. Nevertheless, the lack of accurate measurements on the velocity dispersion may induce large systematic effects on the parameters used for the modelling of the \\DM\\, halo. Even in case of accurate kinematic data for some of the Galaxy satellites, only faint annihilation signals are expected for smooth dark matter halos due to their distance of $\\sim$100 kpc.  Mini-spikes around Intermediate Mass Black Holes (IMBHs) have been recently proposed as promising targets for indirect \\DM\\, detection~\\cite{zhao}. Accurate predictions, in the context of well-defined astrophysical scenarios, have been derived for the distribution and luminosity of these objects~\\cite{bertone34}. Mini-spikes might in fact be detected as bright point-like sources in gamma-rays~\\cite{bertone3} and neutrinos~\\cite{bertone4}, and the prospects for detection with satellites with large field of view such as the upcoming GLAST experiment~\\cite{glast} to be launched in 2008 appear particularly promising.  Current Imaging Atmospheric Cherenkov Telescopes (IACTs) can also effectively search for these objects. The H.E.S.S. (High Energy Stereoscopic System) experiment has already surveyed a significant part of the Galactic plane. With the combination of a large field of view, very good angular resolution and off-axis performance, H.E.S.S. has reached the sensitivity to accurately map the Galactic plane in scan-based observations.  In this paper, H.E.S.S. data are used to derive for the first time experimental exclusion limits on the \\DM\\, annihilation signals within the context of the mini-spike scenario. The paper is organized as follows. Section II is devoted to the mini-spike scenario and the gamma-ray flux expected from \\DM\\, annihilations in mini-spikes. In section III, we present the H.E.S.S. data from the Galactic plane survey and compute the H.E.S.S. flux sensitivity map to \\DM\\, annihilations in the Galactic plane region. Exclusion limits are then derived for \\DM\\, annihilation from mini-spikes. Finally, section IV is devoted to the discussion of the results obtained in this study. ",
        "conclusions": "H.E.S.S. observations (2004-2006) of the Galactic plane allowed to discover more than 20 VHE sources~\\cite{sp2}. Some of them have been identified owing to their counterparts at other wavelengths, but almost half of the sources have no obvious counterpart and are still unidentified~\\cite{darksources}. For these sources, an accurate reconstruction of their energy spectra has been carried out. All spectra were consistent with a pure power-law of spectral indices between 2.0 and 2.5, spanning up to two orders of magnitude in energy above the energy threshold, as shown in Fig.~8 of Ref.~\\cite{darksources}. None of them exhibits an energy cut-off, characteristic of \\DM\\, annihilation spectra, in the energy range from $\\sim$ 100 GeV up to 10 TeV. Furthermore, the detailed study of their morphology~\\cite{darksources} show that all the sources have an intrinsic spatial extension greater than $\\sim$5', while mini-spikes are expected to be point-like sources, since the bulk of the gamma-ray emission comes from a region of size $\\approx 10^{-3}$ pc  at a typical distance from GC of 10 kpc. In the survey region discussed here, only three point-like gamma-ray sources are detected: HESS J1826-148, HESS J1747-281 and HESS J1745-290. They are identified with well-known objects: LS 5039, G0.9+0.1 and the Galactic Center, respectively. H.E.S.S. has detected so far no IMBH candidate within the survey range.   Based on the absence of plausible IMBH candidates in the H.E.S.S. data, we can derive constraints on one scenario for neutralino or LKP dark matter annihilations. These constraints are shown as upper limits on the annihilation cross section (Figures~\\ref{fig:exclusionlimit_pMSSM} and~\\ref{fig:exclusionlimit_KK_AMSB}) but are actually constraints on the entire gamma-ray production scenario, assuming the mass function of scenario B and a coreless halo model. In the survey range, the expected mean number of IMBHs is 4.3, within the assumptions of scenario B, as shown in Fig.~\\ref{fig:Distribution_IMBH_SurveyRange}. The number of IMBHs in the whole halo depends on the assumed cosmological parameters in the formation scenario~\\cite{Koushiappas:2005qz}, in particular the reionization redshift. The value of 4.3 {\\it unmerged} IMBHs is for a reionization redshift of 16. The redshift of $10.8\\pm1.4$ determined from latest WMAP data~\\cite{Komatsu:2008hk} favors a larger number of IMBHs since they had less time for merger events to occur~\\cite{Koushiappas:2005qz}. \\begin{figure}[!ht] \\mbox{\\hspace{-0.cm}\\includegraphics[scale=0.45]{fig5.eps}} \\caption{Distribution of the number of IMBHs for the scenario B in the survey range corresponding to a field of view between $\\pm$3$^{\\circ}$ in Galactic latitude and from -30$^{\\circ}$ to 60$^{\\circ}$ in Galactic longitude. The distribution is obtained from 200 stochastic realizations of Milky Way-like halos (see Sec. II. A for details). The mean of the distribution is 4.3 with an rms of 2.3. The distribution is well fitted to a Poisson distribution with mean 4.4 (black dashed line).} \\label{fig:Distribution_IMBH_SurveyRange} \\end{figure} The distribution of the IMBH number has an intrinsic rms of 2.3. The probability to find no IMBH in the survey range is 2.5\\%. The simulations of the Milky Way halo show that the IMBH number distribution in the sky survey range is well fitted with a Poisson distribution with mean 4.4.  For a given gamma-ray flux sensitivity $\\Phi$, we define the probability density function of black holes yielding an integrated gamma-ray flux larger than $\\Phi$, $d^2N_{BH}(>\\Phi)/dbdl$, from the IMBH catalog used in~\\cite{bertone3} by : \\begin{eqnarray} \\label{eqn:nbhdetected} \\lefteqn{\\,\\,\\,\\,\\,\\,\\,N_{\\rm BH}({\\rm detected},>\\Phi)\\,=\\,} \\nonumber & & \\\\ & & \\int_{b=-3^{\\circ}}^{b=+3^{\\circ}}\\int_{l=-30^{\\circ}}^{l=+60^{\\circ}} \\frac{d^2N_{\\rm BH}(>\\Phi)}{dbdl}\\,dbdl\\,. \\end{eqnarray} $N_{BH}({\\rm detected},>\\Phi)$ denotes the expected number of black holes yielding a gamma-ray flux larger than $\\Phi$ in the survey, $\\Phi$ corresponding to the H.E.S.S. sensitivity. The total number of black holes in the survey is computed by integrating the probability density function $d^2N_{\\rm BH}(>\\Phi)/dbdl$ over the latitude and longitude ranges of the Galactic plane survey. Since no H.E.S.S. source is a plausible IMBH candidate, we thus calculate for each \\DM\\, particle mass the limit at 90\\% C.L. on the velocity-weighted cross section $\\sigma v$, assuming a Poisson distribution.   Fig.~\\ref{fig:exclusionlimit_pMSSM} shows the exclusion limit at the 90\\% C.L. on $\\sigma v$ as a function of the neutralino mass $m_{\\rm DM}$. The neutralino is assumed to annihilate into $b\\bar{b}$ and $\\tau^+\\tau^-$ with 100\\% BR, respectively. \\begin{figure}[!hb] \\mbox{\\hspace{-0.cm}\\includegraphics[scale=0.45]{fig6.eps}} \\caption{Constraints on the IMBH gamma-ray production scenario for different neutralino parameters, shown as upper limits on the annihilation cross section $\\sigma v$ as a function of the mass of the \\DM\\, particle $m_{\\rm DM}$, but with a number of implicit assumptions about the IMBH initial mass function and halo profile (see text for details). For scenario B, the probability of having no observable halos in our Galaxy is 10\\% from Poisson statistics making these limits essentially 90\\% confidence level exclusion limits for this one particular (albeit optimistic) scenario (grey shaded area). The \\DM\\, particle is assumed to be a neutralino annihilating into $b\\bar{b}$ pairs or $\\tau^+\\tau^-$ pairs to encompass the softest and hardest annihilation spectra. The limit is derived from the H.E.S.S. flux sensitivity in the Galactic plane survey within the mini-spike scenario. SUSY models from the pMSSM (black points) are plotted together with those satisfying the WMAP constraints on the \\DM\\, particle relic density (magenta points).} \\label{fig:exclusionlimit_pMSSM} \\end{figure} Below $\\sim$ 2 TeV, the upper contour of the gray shaded area is given by the $b\\bar{b}$ assumption and yields the conservative exclusion limit. Above 2 TeV, the upper contour corresponds to the $\\tau^+\\tau^-$ annihilation spectrum. The maximum sensitivity of H.E.S.S. to DM annihilation for the $b\\bar{b}$ channel is achieved for masses of order 5 TeV as expected from Fig.~\\ref{fig:Integrated_Flux}. For neutralino masses in the TeV energy range, we obtain limits on one mini-spike scenario (scenario B). The limits on $\\sigma v$ are at the level of $10^{-28}$ $\\rm cm^{-3}s^{-1}$ for the $b\\bar{b}$ channel. A rapid decrease in sensitivity is observed for WIMP masses less than $\\sim$~800 GeV. This corresponds to the threshold effect seen in Fig.~\\ref{fig:Integrated_Flux}. \\begin{table}[!ht] \\caption{\\label{tab:parameterspace} pMSSM parameter space randomly scanned to generate SUSY models. A set of parameters corresponds to a specific pMSSM model.} \\begin{ruledtabular} \\begin{tabular}{ccc} Parameter&Minimum& Maximum\\\\ \\hline \\\\ $m_0$ &1 TeV& 30 TeV\\\\ $M_2$ &1 TeV &50 TeV\\\\ $\\mu$ &1 TeV& 50 TeV\\\\ $m_A$ &1 TeV&50 TeV\\\\ $A_{t,b}$ &-300 GeV&+300 GeV\\\\ $tan \\beta$ &3& 60\\\\ \\end{tabular} \\end{ruledtabular} \\end{table} Predictions for SUSY models are generated using the DarkSUSY code~\\cite{darksusy} in a phenomenological MSSM (pMSSM) framework characterized by the following independent parameters: the common scalar mass $m_0$, the gaugino mass parameter $M_2$, the higgsino mass parameter $\\mu$, the CP-odd Higgs mass $M_A$, the trilinear couplings A$_{t,b}$, and the Higgs vacuum expectation value ratio $tan \\beta$. The set of parameters for a given model is randomly chosen in the pMSSM parameter space encompassing a large class of pMSSM models, as described in Table~\\ref{tab:parameterspace}. Models providing a neutralino thermal relic density, $\\Omega_{\\rm DM}h^2$, in the range [0.08,0.12] are overlaid to account for the cold dark matter density inferred from the measurements of the CMB anisotropies of the Wilkinson Microwave Anisotropy Probe (WMAP) satellite. Some models in the mass range 0.8 - 6 TeV can be excluded. \\begin{figure}[ht] \\mbox{\\hspace{-0.cm}\\includegraphics[scale=0.45]{fig7.eps}} \\caption{90\\% C.L. exclusion limits on the annihilation cross section $\\sigma v$ as a function of the mass of the \\DM\\, particle $m_{\\rm DM}$ for a Kaluza-Klein boson $\\tilde{B}^{(1)}$ (solid blue line) and a pure wino $\\tilde{W}$ in AMSB (solid red line). The limit is derived from the H.E.S.S. flux sensitivity in the Galactic plane survey within the mini-spike scenario. % Kaluza-Klein models in UED scenario (blue dashed line) and AMSB models (red dotted line) are plotted together with those satisfying the WMAP constraints on the \\DM\\, particle relic density (pink dashed line and green dotted line, respectively).} \\label{fig:exclusionlimit_KK_AMSB} \\end{figure}  Fig.~\\ref{fig:exclusionlimit_KK_AMSB} shows the 90\\% C.L. exclusion limit on $\\sigma v$ as a function of the DM particle mass $m_{\\rm DM}$ in the case of AMSB and Kaluza-Klein scenarios. In both cases, the BRs of each annihilation channel entering in the calculation of the total annihilation differential spectra are uniquely determined as discussed in section II. In the AMSB scenario, the neutralino is considered here to be a pure wino annihilating with 100\\% BR into the $W^+W^-$ channel. The predictions for $\\sigma v$ are parameterized using the results of~\\cite{moroi} where it is derived in the case of the latter assumption on the annihilation scheme. The AMSB models fulfilling in addition the WMAP constraints are overlaid. Models yielding neutralino masses between 900 GeV and 6 TeV are excluded including those having a neutralino thermal relic density compatible with the WMAP measurements. Predictions on the annihilation cross section for pairs of $\\tilde{B}^{(1)}$ from UED theories~\\cite{sigmavKK} are also plotted as well as those satisfying the WMAP constraints. In this scenario, $\\tilde{B}^{(1)}$ masses in the range from 0.8 to 6 TeV can be excluded."
    },
    "0806/0806.2723_arXiv.txt": {
        "abstract": "We measure the ellipticity of isolated clusters of galaxies in the Sloan Digital Sky Survey (SDSS) using gravitational lensing. We stack the clusters, rotating so that the major axes of the ellipses determined by the positions of cluster member galaxies % are aligned. We exclude the signal from the central 0.5 $h^{-1}$ Mpc to avoid problems with stacking alignment and cluster member contamination. We fit an elliptical NFW profile and find a projected, two-dimensional axis ratio % for the dark matter of $f = b/a = 0.48^{+0.14}_{-0.09} \\ (1 \\sigma)$, and rule out $f=1$ at 99.6 per cent confidence thus ruling out a spherical halo. We find that the ellipticity of the cluster galaxy distribution is consistent with being equal to the dark matter ellipticity. The results are similar if we change the isolation criterion by 50 per cent in either direction. ",
        "introduction": "Cosmological simulations can be used to predict many different statistics of the mass distribution in the Universe. The most commonly employed statistic is the two-point correlation function, or its Fourier counterpart the power spectrum. Three-point statistics are much harder to predict and measure, and higher orders are rarely discussed. A more popular statistic is the number of peaks in the mass distribution, as given by the number of clusters of galaxies. The dark matter power spectrum and the number of clusters of galaxies are often cited as among the best ways to constrain the properties of dark energy~\\citep[e.g.][]{detf}.  Uncertainties on cosmological parameters are decreased when two measurements have different parameter degeneracies and are often referred to as `complementary'. % In this paper we consider a statistic which % may offer % complementary constraints on cosmology: the \\emph{shapes} of peaks in the mass distribution, as probed by the ellipticity of galaxy cluster dark matter halos. In addition this may place important constraints on modifications to the law of gravity since we may compare the results from both dark and light matter, as also tested by studying the dark and light matter distributions in the bullet cluster~\\citep{clowebgmrjz06}.  Predictions of cluster ellipticities come mostly from numerical simulations \\citep{west89, detheije95, jing02, floor03, ho04, flores05, rahman06}. The ellipticity is expected to depend on cosmological parameters \\citep{evrard93, splinter97, buote97, suwa03, rahman04} and to evolve with redshift \\citep{kasun05, allgood06}, an evolution which itself might depend on cosmology \\citep{hopkins05, ho06}. The distribution of sub-halos within a cluster halo is found to be an indicator of the overall halo ellipticity \\citep{bode07} but is slightly less elliptical.  The ellipticity of the brightest cluster galaxy and the ellipticity of the distribution of cluster member galaxies is much easier to observe than any other cluster ellipticity measure. This has been studied by a number of authors \\citep{west90, plionis91, rhee91, strazzullo05}. Most recently, \\cite{wang07} studied groups of galaxies in SDSS and found an alignment between the brightest cluster galaxy (BCG) and the distribution of galaxies which was strongest in the most massive groups and between red BCGs and red group member galaxies.  In this paper we focus on measuring the dark matter ellipticity directly using gravitational lensing. We also compare this ellipticity to the ellipticity of the cluster member galaxy distribution, to see how reliable a tracer of ellipticity the light is, and therefore whether it can be used by itself for cosmological studies.  Gravitational lensing has been used very successfully to measure the mass and profile of clusters of galaxies by many authors. The work most relevant to our study is that of \\cite{sheldon01,sheldonI07} and \\cite{sheldonIII07} who stack the lensing signal from many clusters of galaxies to find an average signal. \\cite{nat2000} proposed that when stacking the shear signal from many halos the stack could be made while retaining information about the major axis of the halo, as observed from the distribution of light. The stacked shear map should then provide a constraint on the ellipticity of the dark matter halo, if indeed the mass and light were aligned. They considered lensing by galaxies but we apply the same technique to clusters of galaxies here.  The ellipticity of the dark matter distribution has been studied observationally for galaxy sized halos using gravitational lensing using data from the Red-Sequence Cluster Survey % \\citep{hoekstrayg04}, the Sloan Digital Sky Survey \\citep{mandelbaumea06} and the Canada-France-Hawaii Telescope Legacy Survey \\citep{parker07}. % This has proved to be extremely difficult and we present here, for the first time, results from stacking galaxy cluster halos. \\cite{cypriano04} has previously made measurements of the cluster dark matter ellipticity using gravitational lensing measurements from individual clusters. They found a good agreement between the dark matter halo orientation and the orientation of the brightest cluster galaxy.  When calculating angular diameter distances and the mean density of the Universe we assume a flat cosmology with $\\Omega_m = 0.3$.  The NFW halo profile has a very weak dependence on the fluctuation amplitude. We assume $\\sigma_8=0.8$ for this calculation.  The structure of this paper is as follows. In Sec.~\\ref{sec:data} we describe our data set and operations we have performed on the data such as rotating and stacking and removal of neighboring clusters.  In Sec.~\\ref{sec:modelling} we use two theoretical models - a Singular Isothermal Ellipsoid (SIE) and a Navarro-Frenk-White \\citep[NFW, see ][]{nfw} model - of the mass and light distribution, and look at correction factors from the redshift distribution and cluster decontamination. Our results, and some interpretations of these, are presented in Sec.~\\ref{sec:results}.  Conclusions are summarized and discussed in Sec.~\\ref{sec:conclusions}. ",
        "conclusions": "\\label{sec:conclusions} We have used galaxy clusters from the catalogue of \\cite{bencat} and shear maps as used in \\cite{sheldon04} to investigate the light and dark matter ellipticities of galaxy cluster halos.  We rotate the selected clusters so that their major axes are aligned, and stack them according to the method which is described in \\cite{nat2000} for galaxy-galaxy lensing. This is the first time that this method has been applied to cluster lensing.  Through the pattern imprinted on the shear maps by intervening massive cluster halos, we have detected a dark matter ellipticity of these halos at a 99.6\\% level, with an axis ratio of $f_{\\rm DM} = 0.48^{+0.14}_{-0.09}$ from a joint $\\chi^2$-analysis for an NFW model, using 4281 clusters between $0.10 < z \\le 0.30$. We have corrected for dilution of the shear signal by cluster members left in the shear galaxy catalogues but have masked out the central areas where this correction factor grows too large.  The light matter distribution of the clusters, as traced by the number density of individual cluster members, is also clearly elliptical, with a joint axis ratio $f_{\\rm LM} = 0.60^{+0.004}_{-0.005}$ for the NFW model.  Using the shear catalogue alone gives very similar results to using the shear catalogue concatenated with the cluster member catalogue, which means that we are essentially comparing like with like when comparing dark and light map ellipticities. This is because both the light matter (shear catalogue) and the dark matter have been stacked in the same way.   We find that the ellipticity of the dark matter distribution is consistent with the ellipticity of the galaxy number density distribution. Our result is limited by the uncertainty in the ellipticity of the dark matter distribution. The results for the NFW and SIE agree within the errors, but any differences could be attributed to a changing ellipticity with radius.  We have not attempted to measure this effect since the uncertainties are too large.  Shear maps from neighboring clusters will influence each other, making the pattern more elliptical or less, depending on the redshift of the neighbors. In order to reduce this effect as much as possible, we use only the cluster with the highest number of members ($N_{\\rm{gals}}^{r200}$) when two (or more) clusters are closer together than $5 h^{-1}$ Mpc, as seen on the sky. We find that increasing or decreasing the minimum distance to a close neighbor by 50 \\% does not significantly affect our main results.  We have also simulated the effect that a possible misalignment between the light and dark matter could have, and concluded that, according to our results, the light and dark matter must be misaligned by less than $\\sim 50^\\circ$.  Our measurement is very insensitive to overall calibration biases in shear measurements, since these are degenerate with cluster mass, rather than cluster ellipticity. It is unlikely that biases in shear measurement would vary with angle around the cluster: residual point spread function anisotropies would be oriented at random with respect to the cluster major axis and would cancel out on stacking; the cluster member light may leak into the postage stamps used to measure shears of background galaxies, but we remove the central region where the number of confirmed cluster members is significant.   Possibly the biggest potential systematic is intrinsic alignment of cluster members pointing at the cluster center \\citep{ciottid94,kuhlendm07,pereirabg08,knebe08,faltenbacherea08} which would be a problem for us because we have not made a significant attempt to remove cluster members from our analysis and our background catalogue is not particularly deep. A thorough assessment of this effect is beyond the scope of the current work.  To first order this would make our observed shear maps less elliptical since the contamination by cluster members is greatest along the cluster major axis, and thus the strong gravitational shears expected along this axis will be partially cancelled out by the cluster members which have the opposite ellipticity since they tend to point at the cluster center.  However a full assessment of this effect would also have to take into account the variation in intrinsic alignment with respect to the cluster major axis, which appears to be more complicated \\citep{kuhlendm07} and possibly weakens the degree of cancellation preferentially along the cluster major axis.  In \\cite{hopkins05}, the authors use a large-scale, high-resolution $N$-body simulation to predict cluster ellipticities and alignments in a $\\Lambda$CDM universe. They find an ellipticity \\be \\langle \\epsilon \\rangle = 1 - \\langle f \\rangle = 0.33 + 0.05z. \\label{eq:hopkins} \\ee for the redshift range $0<z<3$. This redshift evolution is negligible for the redshift range considered here, and due to the large uncertainties we do not try to detect any trends with redshift. For a redshift in the middle of the redshift range of our cluster sample ($z=0.2$), this formula yields an axis ratio of $f = 0.66$ which is in good agreement with our result of $f_{\\rm DM}=0.48^{+0.14}_{-0.09}$.  \\cite{ho06} have used numerical simulations of cluster formation with the aim of investigating the possibility of using cluster ellipticities as a cosmological probe. They find that the mean ellipticity of high mass clusters is approximated by \\be \\bar{e}(z) = 0.245 \\ \\Bigl[ 1 - 0.256 \\ \\frac{\\sigma_{8}(z)}{0.9} + 0.00246 \\ \\frac{\\Omega_m}{0.3} \\Bigr]. \\ee Using $\\sigma_{8}(z) = 0.8$ and $\\Omega_{m}=0.3$, this gives an ellipticity of $e \\sim 0.2$, or $f \\sim 0.8$. This result is more circular than our main result. However, to make a proper comparison of our result with theory it would be necessary to make a theoretical prediction that takes into account our observation method, especially the overlaps in the shear field due to close neighbors and the impact of selecting the most isolated clusters.   \\cite{hoekstrayg04} report on a first weak-lensing detection of the flattening of galaxy dark matter halos, using data from the Red-Sequence Cluster Survey. They find an average galaxy dark matter halo ellipticity of $\\langle \\epsilon \\rangle = 0.33^{+0.07}_{-0.09}$. They also find a detection that dark matter halos are rounder than the light map. In a later work, \\cite{mandelbaumea06} did not find a definite detection of this effect in the larger SDSS dataset. However, \\cite{parker07}, measuring the ratio of tangential shear in regions close to the semi-minor versus that close to the semi-major axes of the lens, find some evidence of a halo ellipticity of $\\sim 0.3$ using early data from the Canada-France-Hawaii Telescope Legacy Survey (CFHTLS). % \\cite{mandelbaumea06} commented that the stronger signal in clusters of galaxies means that there is more chance of making a detection of ellipticity in this higher mass data set. % Therefore we have confirmed the detection of dark matter halo ellipticity, extending the measurement to cluster scales. We find no evidence for different ellipticities for the light and dark matter distribution on cluster scales.  \\cite{mandelbaumea06} focus on measuring the ratio of (dark matter) halo ellipticity to galaxy (light) ellipticity $\\epsilon_h/\\epsilon_g$ where $\\epsilon =(a^2 - b^2)/(a^2 + b^2)$. They find $\\epsilon_h/\\epsilon_g=0.60 \\pm 0.38$ (68 per cent confidence) for elliptical galaxies, which is most comparable with our result of $f_{\\rm DM}/f_{\\rm LM} = 0.91^{+0.19}_{-0.25}$ (68 per cent confidence) where $f=b/a$. Repeating our calculation using the different ellipticity parameters we find $\\epsilon_h/\\epsilon_g = 1.37^{+0.35}_{-0.26}$. It is expected that our value is greater than unity, since we find that the dark matter is more elliptical (lower axis ratio) than the light matter. However it is not significantly larger. The result of \\cite{mandelbaumea06} is the opposite side of unity, but the difference is not significant and we note that our result is for clusters and that of \\cite{mandelbaumea06} is for galaxies.  The anisotropy of the lensing signal around individual galaxies, if definitely detected, has been said to pose a serious problem for alternative theories of gravity \\citep{hoekstrayg04}.  For galaxy clusters it is more complicated, as the dominant source of baryons is the intracluster gas. Lensing by clusters has been found to pose a problem for Modified Newtonian Dynamics (MOND), as there seems to be a need to include a dark matter component \\citep[]{sanders03, takahashi07}.  Because the gas - the dominant baryonic component of clusters - is collisional, we can suppose that on cluster scales it will be less elliptical than the light. We may end up concluding that either (i) the gas distribution is elliptical (which we would not expect) or (ii) our result is inconsistent with MOND. However, further study and simulations of this is clearly needed."
    },
    "0806/0806.1824.txt": {
        "abstract": "{ % Context heading {\\bb We present new polarisation  observations at 13~cm, acquired using the Westerbork Synthesis Radio Telescope (WSRT), of 65 sources,} from the B3-VLA {\\bb sample} of Compact Steep-Spectrum sources. These new data are combined with {\\bb our} VLA polarisation data, at 3.6, 6 and, 21~cm, presented in a previous paper. Due to the multi-channel frequency capabilities of the WSRT, these new 13~cm observations enable {\\bb a more reliable determination of integrated Rotation Measures, and of depolarisation behaviour with wavelength. The new data are inconsistent with the  depolarisation models that we used earlier, and we propose an alternative model which seems to work properly. We also revise our previous model for the external Faraday screen, and its dependence on the source redshift.}  % Aims heading %{Aims } % Methods heading %{Methods} % Results heding %{Results} % Conclusions heading %{Conclusion} ",
        "introduction": " ",
        "conclusions": "\\label {concl}  We have observed 65 radio sources from the B3-VLA CSS sample (\\citeauthor*{Fanti01}) using the WSRT at 13~cm, to study the source polarisation properties. At the WSRT resolution, the radio sources are all unresolved, and we can therefore only discuss their global properties. The new 13-cm data, combined with earlier low-resolution VLA data at 3.6, 6, and 21~cm, have improved the determination of Rotation Measures ($RM$) and the Rotation Measure Dispersions ($\\sigma_{\\rm RM}$). This has allowed their properties to be defined more carefully as a function of both redshift and Linear Size, and the characteristics of the surrounding ISM to be modelled.  The main results of the paper are the following:  \\begin{description} \\item[{\\mdseries 1)}] Radio sources $\\le 5$~kpc are mostly unpolarised at 13~cm. This result is similar to results found at 20~cm \\citep{Cotton03} and  at 3.6 and 6~cm (\\citeauthor*{Fanti04}), and the ``critical size'' that we find is intermediate between those found earlier.  \\item[{\\mdseries 2)}] The 13~cm polarisation angles have led to a revision of the earlier $RM$s (\\citeauthor*{Fanti04}) for $\\approx 30\\%$ of the sources.  \\item[{\\mdseries 3)}] The integrated fractional polarisation $m_{\\lambda}$, as a function of wavelength shows in general a decrease between 3.6 and 13~cm, followed by a flattening between 13 and 20~cm. This is the major observational result of the present paper. At variance with the conclusions of ~\\citeauthor*{Fanti04}, the ~\\citet{Burn66} and ~\\citet{Tribble91} models are not adequate to describe this behaviour. However, an empirical variant of the ~\\citet{Burn66} model, which introduces  a ``partial coverage'' by the depolarising curtain, appears to reproduce the data well. The adopted formula allows the Rotation Measure Dispersion, $\\sigma_{\\rm RM}$, and the fraction, $f_{\\rm c}$, of source covered by the depolarising curtain, to be determined.  \\item[{\\mdseries 4)}] For a {\\bb minority of sources}, $m_{\\lambda}$ shows an irregular, possibly oscillatory, behaviour with $\\lambda$. We propose that these sources contain sub-components with different Rotation Measures that produce beats {\\bb with} $\\lambda$ in the integrated fractional polarisation.  \\item[{\\mdseries 5)}]  $\\sigma_{\\rm RM}$ shows a clear dependence on redshift (up to $z \\approx 1.8$),  and on projected Linear Size. %(see \\ref{...}).   \\item[{\\mdseries 6)}] We have analysed these dependences using a model similar to that presented in ~\\citeauthor*{Fanti04}, in which a depolarising Faraday curtain, with a King-like distribution of the magneto-ionic medium, is produced either by a smooth medium with an irregular magnetic field on small scales, or by  a clumpy magnetised medium (NL region). The parameters derived in either case  differ from those we obtained in ~\\citeauthor*{Fanti04}, because of the use of the new 13~cm data and of the new model adopted {\\bb for} the depolarisation behaviour with  $\\lambda$. In both models the core radius of the curtain has {\\bb to be $\\le 0.5$~kpc.}  In the smooth medium model, the required central density, at $z = 0$, is in the range of that found using X-ray observations in the centre of early-type galaxies. Beyond the core radius, it decreases according to $r^{-\\eta}$, where $1.5 \\le \\eta \\le 2.0$. The central magnetic field, at $z = 0$, is $\\approx 150$~$\\mu$G, and the magnetic field energy  density is close to the thermal energy.  In  the clumpy medium model, {\\bb identified with} the NLR, we assumed pressure equilibrum between the clouds and the diffuse medium in which they are embedded. The required magnetic field, inside the clouds at $z =0$, is $\\approx 6~\\mu$G, and the magnetic energy is a negligible ($\\approx 10^{-4}$) fraction of the thermal energy. Beyond the core radius, the density of the smooth medium, which confines the clouds, should decline as $r^{-\\eta}$, with  $1.2 \\le \\eta \\le 1.6$. The product $B_{\\rm NL,0} \\times n_{\\rm NL,0}$  is a strong function ($\\propto (1+z)^4$) of redshift up to $z \\approx 1.8$.  The $NL~ model$ requires more parameters than the {\\it smooth medium model}, and these parameters need tuning to avoid the ``NL covering  problem'' discussed in Sect.~\\ref{farprop}. For these reasons, we favour the {\\it smooth medium model}.  \\item[{\\mdseries 7)}] The ISM density profiles required by the depolarisation models  are consistent with the expectations of the evolutionary models of young radio sources.  \\item[{\\mdseries 8)}]  The partial coverage, that we introduced to describe the long wavelength behaviour of $m_{\\lambda}$ vs $\\lambda^2$, is probably due to {\\bb orientations effects} of the source with respect to the line of sight. For sources that are almost orthogonal to the line of sight, the two lobes are depolarised similarly  by a spherically-symmetric medium, and $m_{\\lambda}$ is fitted by the Burn model. However, when the source is at large angles with respect to the line of sight, the far lobe is more depolarised causing the rapid decrease of $m_{\\lambda}$ at short wavelengths, while the {\\bb nearby one} is less depolarised giving rise to a  lower decrement of $m_{\\lambda}$ with wavelength. The total fractional polarisation $m_{\\lambda}$ then appears to be almost constant  at long wavelengths.  \\item[{\\mdseries 9)}]  The total source-frame Rotation Measures are generally too large to be the residual of random rotations across the source due to a totally irregular magnetic field, and require a large-scale ordered magnetic-field component. The Rotation Measures show hints of  a correlation with both redshift and Linear Size in a way similar to the Rotation Measure dispersions. If these correlations are real, the  ordered magnetic-field component would have  parameters similar to those of the random field component. \\end{description}"
    },
    "0806/0806.4275_arXiv.txt": {
        "abstract": "We present the derivation of two-integral distribution functions for the first four members of the family of generalized Kalnajs discs, recently obtained by \\cite{GR}, and which represent a family of axially symmetric galaxy models with finite radius and well behaved surface mass density. In order to do this we employ several approaches that have been developed starting from the potential-density pair and, essentially using the method introduced by \\cite{KAL2}, we obtain some distribution functions that depend on the Jacobi's integral. Now, as this method demands that the mass density can be properly expressed as a function of the gravitational potential, we can do this only for the first four discs of the family. We also find another kind of distribution functions by starting with the even part of the previous distribution functions and using the maximum entropy principle in order to find the odd part and so a new distribution function, as it was pointed out by \\cite{DEJ}. The result is a wide variety of equilibrium states corresponding to several self-consistent finite flat galaxy models. ",
        "introduction": "The problem of finding self-consistent stellar models for galaxies is of wide interest in astrophysics. Usually, once the potential-density pair (PDP) is formulated as a model for a galaxy, the next step is to find the corresponding distribution function (DF). This is one of the fundamental quantities in galactic dynamics specifying the distribution of the stars in the phase-space of positions and velocities. Although the DF can generally not be measured directly, there are some observationally accesible quantities that are closed related to the DF: the projected density and the light-of-sight velocity, provided by photometric and kinematic observations, are examples of DF's moments. Thus, the formulation of a PDP with its corresponding equilibrium DFs establish a self-consistent stellar model that can be corroborated by astronomical observations.  On the other hand, a fact that is usually assumed in astrophysics is that the main part of the mass of a typical spiral galaxy is concentrated in a thin disc (\\cite{BT}). Accordingly, the study of the gravitational potential generated by an idealized thin disc is a problem of great astrophysical relevance and so, through the years, different approaches has been used to obtain the PDP for such kind of thin disc models (see \\cite{KUZ} and Toomre (\\citeyear{T1,T2}), as examples). Now, a simple method to obtain the PDP of thin discs of finite radius was developed by \\cite{HUNT}, the simplest example of disc obtained by this method being the Kalnajs disc (\\cite{KAL1}).  In a previous paper (\\cite{GR}) we use the Hunter method in order to obtain an infinite family of axially symmetric finite thin discs, characterized by a well-behaved surface density, whose first member is precisely the well-known Kalnajs disc (\\cite{KAL1}) and, more recently, a detailed study of the kinematics characterizing such discs was be made by \\cite{RLG}. So, in order to obtain  self-consistent stellar models for this family of generalized Kalnajs discs, we will consider at the present paper the derivation of some two-integral DFs for the first four members of the family.  Now, as is stated by Jeans's theorem, an equilibrium DF is a function of the isolating integrals of motion that are conserved in each orbit and, as it has been shown, it is possible to find such kind of DFs for PDPs such that there is a certain relationship between the mass density and the gravitational potential. The simplest case of physical interest corresponds to spherically symmetric PDPs, which are  described by isotropic DFs that depends on the total energy $E$. Indeed, as was be shown by \\cite{EDD}, it is possible to obtain this kind of isotropic  DFs by first expressing the density as a function of the potential and then solving an Abel integral equation.  Recently, \\cite{JO} found that a similar procedure can be performed in the axially symmetric case, where the equilibrium DF depends on the energy $E$ and the angular momentum about the axis of symmetry $L_{z}$, i.e. the two classical integrals of motion. They developed a formalism that essentially combines both the Eddington formulae and the \\cite{FRICKE} expansion in order to obtain the DF's even part, starting from a density that can be expressed as a function of the radial coordinate and the gravitational potential. Once such even part is determined, the DF's odd part can be obtained by introducing some reasonable assumptions about the mean circular velocity or using the maximum entropy principle.  On the other hand, another method appropriated to find DFs depending only of the Jacobi's integral, $E_r = E - \\Omega L_z$, for axially symmetric flat galaxy models was introduced by \\cite{KAL2} for the case of disc-like systems and basically consists in to express the DF as a derivative of the surface mass density with respect to the gravitational potential. Such method does not demands solving an integral equation, but instead is necessary properly to express the mass density as a function of the gravitational potential, a procedure that is only possible in some cases.  In this paper we will use both of the mentioned approaches in order to obtain equilibrium DFs for some generalized Kalnajs discs. Accordingly, the paper is organized as follows. First, in Section \\ref{sec:finteq}, we present the fundamental aspects of the two methods that we will use in order to obtain the equilibrium DFs. Then, in Section \\ref{sec:dfkal}, we present a summary of the main aspects of the generalized Kalnajs discs, at Section \\ref{sec:kal}, and then, in Sections \\ref{sec:dfm1} -- \\ref{sec:dfm4}, we derive the DFs for the first four members of the family. Finally, in Section \\ref{sec:conc}, we summarize our main results. ",
        "conclusions": "\\label{sec:conc}  We presented the derivation of two-integral equilibrium DFs for some members of the family of generalized Kalnajs discs previously obtained by \\cite{GR}. Such two-integral DFs were obtained, esentially, by expresing them as functionals of the Jacobi's integral, as it was sketched in the formalism developed by \\cite{KAL2}. Now, since such formalism demands that the surface mass density can be written as a potential-dependent function, the above procedure can only be implemented for the first four members of the family, the discs with $m=1,2,3,4$.  Indeed, the procedure requires that the expression given the relative potential $\\Psi_r$ as a function of the spheroidal variable $\\eta$ can be analytically inverted in order to express the surface mass density $\\Sigma$ as a function of the relative potential. So,  we can do this in a simple way for the discs with $m = 1$ to $4$. However, when $m > 4$ we must to solve an equation of grade larger than four, whose analytical solution do not exists.  For the first two members of the family, the discs with $m=1,2$, we also use the method introduced recently by \\cite{JO} in order to find the even part of the DF and then, by introducing the maximum entropy principle, we can determines the full DF. This procedure was also used for the other three discs, starting from the Kalnajs method, so defining another class of two-integral DFs. Such kind of DFs describes stellar systems with a preferred rotational state, characterized by the parameter $\\alpha$.  This paper can be considered as a natural complement of the work previously presented by \\cite{GR} and \\cite{RLG}, where the PDP formulation and the kinematics, respectively, of the generalized Kalnajs discs were analyzed. Now, by the construction of the corresponding two-integral DFs, the first four members of this family can be considered as a set of self-consistent stellar models for axially symmetric galaxies."
    },
    "0806/0806.1666_arXiv.txt": {
        "abstract": "XPOL, the first correlation polarimeter at a large millimeter telescope, uses a flexible digital correlator to measure all four Stokes parameters simultaneously, i.e. the total power $I$, the linear polarization components $Q$ and $U$, and the circular polarization $V$. The versatility of the backend provides adequate bandwidth for efficient  continuum observations as well as sufficient spectral resolution (40 kHz) for observations of narrow lines. We demonstrate that the polarimetry specific calibrations are handled with sufficient precision, in particular the relative phase between the Observatory's two orthogonally linearly polarized receivers.  The many facets of instrumental polarization are studied at 3mm wavelength in all Stokes parameters: on--axis with point sources and off--axis with beam maps. Stokes $Q$ which is measured as the power difference between the receivers is affected by instrumental polarization at the 1.5\\% level. Stokes $U$ and $V$ which are measured as cross correlations are very little affected (maximum sidelobes 0.6\\% ($U$) and 0.3\\% ($V$)). These levels critically depend on the precision of the receiver alignment. They reach these minimum levels set by small ellipticities of the feed horns when alignment is optimum ($\\simlt0.3''$).  A second critical prerequisite for low polarization sidelobes turned out to be the correct orientation of the polarization splitter grid. Its cross polarization properties are modeled in detail.  XPOL observations are therefore limited only by receiver noise in Stokes $U$ and $V$ even for extended sources. Systematic effects set in at the 1.5\\% level in observations of Stokes $Q$. With proper precautions, this limitation can be overcome for point sources. Stokes $Q$ observations of extended sources are the most difficult with XPOL. ",
        "introduction": "Spectropolarimetric observations at millimeter wavelengths are of considerable astrophysical interest. Several molecular transitions which are bright in star forming regions are polarized due to the magnetic fields associated with most stages of star formation \\citep{Vallee}. Lines may either be circularly polarized due to the Zeeman effect or linearly polarized through the Goldreich--Kylafis effect, or both. Line polarization is also known in late stages of stellar evolution where  dense circumstellar shells often give rise to maser emission in millimeter lines \\citep{Menten}. At least in the case of SiO, the transitions may have very complex polarization characteristics due to magnetic fields and maser propagation.  The millimeter continuum may also be polarized in Galactic sources due to emission from non-spherical paramagnetic dust grains.  The continuum radiation of extragalactic sources, notably the synchrotron emission from active galactic nuclei, is also often polarized.  Its detection at millimeter wavelengths is of interest, since the lower optical depth allows us to look deeper into the sources than is possible at longer wavelengths and since these wavelengths are less affected by interstellar scattering.  The ideal polarimeter should therefore be capable of observing lines as well as continuum, and all four Stokes parameters should be measured. Given the usual weakness of the polarization signal in the millimeter range, a large telescope is advantageous if compact sources are to be measured.  The few spectral polarization observations made previously at the IRAM 30m telescope measured  % circular polarization \\citep{Dick1,Dick2,TM99}. Linear polarization measurements were attempted in the continuum with a bolometer \\citep{Lemke}.    % Many more observations of linear polarization were made at the James--Clark--Maxwell telescope \\citep[see][and references therein]{Jane}.  In October 1999 when we began developing XPOL, the IRAM 30m telescope had started a profound upgrade of its receivers.  At the end of this transformation, the telescope was equipped with 4 single beam, dual polarization heterodyne  receivers which cover most of the 75 to 270 GHz frequency range. All receivers have linearly polarized feed horns. Here we concentrate on work at 3mm wavelength with receivers A\\,100 (vertically polarized) and B100 (horizontally polarized). The four Stokes parameters  are derived from the IF signals from these receivers (the Intermediate Frequecy is at 1.5 GHz) and refer to the right handed Nasmyth coordinate system $\\mathcal{K_N}$ (Fig.~\\ref{f:schema}) in which these receivers are stationary. In this configuration, the Stokes parameters relate to the time--averaged horizontal, $E_x$, and vertical, $E_y$, electric fields and their relative phase $\\delta$  \\citep[chap.~10]{BW}: \\begin{eqnarray} I  =   \\langle E_x^2\\rangle + \\langle E_y^2\\rangle \\label{e:Stokes} \\\\ Q  =   \\langle E_x^2\\rangle - \\langle E_y^2\\rangle  \\\\ U  =   2\\ \\langle E_x E_y\\,{\\rm cos}\\,\\delta\\rangle   \\\\ V  =   2\\ \\langle E_x E_y\\,{\\rm sin}\\,\\delta\\rangle \\end{eqnarray} This description of the Stokes parameters in terms of electric fields lends itself most naturally to our case where heterodyne receivers amplify the incoming electromagnetic fields. Stokes $I$ and $Q$ are thus obtained from power measurements, whereas Stokes $U$ and $V$ are measured as correlations. The fractional linear, \\pL, and circular, \\pC, polarization and the polarization angle, $\\tau$, are then derived in the usual way.  Correlation polarimeters are well known in radio astronomy \\citep[chapter 4.7]{RW}. Heiles and collaborators \\citep{Heiles-I,Heiles-All,Heiles-Mueller} give a very detailed description of this measurement technique as used at the Arecibo Observatory. To our knowledge, XPOL at the IRAM 30m telescope and its analog predecessor IFPOL \\citep{IFpol} are the first such correlation polarimeters used at a large millimeter telescope. ",
        "conclusions": "XPOL, the first correlation polarimeter on a large single dish millimeter telescope, is a versatile instrument, capable of measuring line and continuum in all four Stokes parameters simultaneously. The setup makes use of the Observatory's 3mm (single pixel) heterodyne receivers which are housed in separate dewars on mechanically independent mounts. Our study shows that these receivers and their arrangement, which were not designed with polarization observations in mind, can nevertheless be used efficiently for many types of polarization observations.  In particular, all issues related to calibration were solved satisfactorily. The instrumental phase is precisely measured ($\\simlt1$\\dgr) for all spectral channels individually, and it is very stable in time. A new method is applied for precisely measuring the decorrelation losses.  The many facets of instrumental polarization, in practice the factors limiting the precision of XPOL observations, were studied at 3mm wavelength for all Stokes parameters in much detail. The most information can be drawn from maps of the Stokes beams. We showed that the pattern and strength of the polarization sidelobes very sensitively depends on the precision of the alignment between the receivers. Only when a precision of $\\simlt0.3''$ was reached, did we see the 4--lobed pattern expected in a Cassegrain telescope with its symmetry about the radio axis. This symmetry does not appear to be destroyed by additional reflections on the flat mirrors of our Nasmyth focus. The 4--lobed structure of the polarization sidelobes and their maxima are then determined by slight differences in the electric parameter (polarization, beam width, etc.) of the feedhorns.  The practical limits for XPOL observations of point sources can be inferred from the values in Stokes $Q,\\ U$ and $V$ at the center of the Stokes maps. These values are compatible with the ``false polarization'' measured on strong and unpolarized point sources. We conclude that Stokes $Q$, which is measured as a power difference, is most affected by instrumental polarization, whereas Stokes $U$ and $V$, which are measured as  cross correlations, are very little affected. The lower instrumental polarization of $U$ and $V$ is due to the insensitivity  of correlations to fluctuations of receiver gain and atmospheric emission. An additional disadvantage arises for $Q$ under an unstable atmosphere where the wobbling period may be too long for complete cancellation of the emission fluctuations.   In summary, XPOL observations of Stokes $U$ and $V$ are possible down to the $\\sim0.5$\\% level before systematic effects set in, both for point--like and extended sources. Instrumental polarization limits Stokes $Q$ observations at the 1.5\\% level. With proper precaution, this limitation can be overcome for point sources. Measurements of the linear polarization of extended sources are the most difficult observations with XPOL.   The design of a new receiver optimized for polarization  would have to pay special attention to the feed, as its cross polarization constitutes the ultimate limit for systematic measurement errors. The alignment of the two feed horns, the next largest source of systematic error, appears easier if the feed horns (and the subsequent radiofrequency components) are housed in the same dewar, or even better, only one feed is used for the two polarizations. It therefore appears quite feasible on a big single dish millimeter telescope for a carefully designed polarization receiver to keep instrumental polarization well below the 0.5\\% level."
    },
    "0806/0806.1358_arXiv.txt": {
        "abstract": " ",
        "introduction": "Cataclysmic variables (CVs) consist of a low-mass star filling its Roche lobe and transferring mass to its white-dwarf (WD) companion, a procedure resulting in the formation of an accretion disc. Due to their unpredictability and diversity leading to various types of CVs, these stars are since long    the subject of many studies.  Our aim is to better understand the properties of these systems by detecting periodicities and by analysing the various physical processes that occur, such as outbursts, superhumps, eclipses, flickering, etc. Our data will be useful for studying these CVs in the future, when a system possibly is in a different brightness state. In this study we present time-resolved photometric data of nine poorly-studied CVs, dwarf novae (DN) and nova-likes (NLs) in particular. A deeper study of five more CVs, was published by \\citet{pap06}. ",
        "conclusions": ""
    },
    "0806/0806.4261_arXiv.txt": {
        "abstract": "{Magnetic fields are thought to be one of the possible mechanisms responsible for shaping the generally spherical outflow of evolved stars into often aspherical planetary nebulae. However, direct measurements of magnetic fields during the transition to the planetary nebula phase are rare.}{The aim of this project is to expand the number of magnetic field measurements of stars in the (proto-)planetary nebula phase and find if the magnetic field strength is sufficient to affect the stellar outflow.}{We used Very Long Baseline Array observations to measure the circular polarization due to the Zeeman splitting of 22~GHz \\water masers in the envelope of the proto-planetary nebula candidate star IRAS~19296+2227 and the planetary nebula K3-35.}{A strong magnetic field of $B_{\\rm ||}=-135\\pm28$ is detected in the \\water maser region of the proto-planetary nebula candidate IRAS 19296+2227. The \\water masers of K3-35 are too weak to detect circular polarization although we do present the measurements of weak linear polarization in those masers.}{The field measured in the masers of IRAS 19296+2227 is dynamically important and, if it is representative of the large scale field, is an important factor in driving the stellar mass loss and shaping the stellar outflow. ",
        "introduction": "At the end of their evolution, a majority of stars go through a period of high mass loss that is an important source for replenishing interstellar space with processed materials. During the asymptotic giant branch (AGB) phase, this mass loss produces circumstellar envelopes which are found to undergo a major modification during the rapid transition from AGB star to Planetary Nebula (PN). The standard assumption is that the initial slow AGB mass loss in a short time changes into a fast superwind generating shocks and accelerating the surrounding envelope \\citep{Kwok78}. However, a large fraction of PNe have asymmetric shapes, with the majority of the young PNe being bipolar. Thus, at some point during the evolution to a PNe the AGB stars must undergo a process in which the spherically symmetric outflow is altered to produce a-spherical PNe morphologies. It has been shown that the energy contained in the outflows of young bipolar PNe is often orders of magnitude larger than can be provided by radiation pressure \\citep{Bujarrabal01}. The source of this energy has been argued to be magnetic fields, binary or disk interaction or a combination of these \\citep[see e.g.][and references therein]{Balick02, Frank07}.  Maser polarization observations are the predominant source of information about the role of magnetic fields during the late stages of stellar evolution. Most observations have focused on the masers in the envelopes of AGB stars, as OH, H$_2$O and SiO masers are fairly common in these sources, and have revealed strong magnetic fields throughout the entire envelope \\citep[e.g.][]{Etoka04, Vlemmings05, Herpin06}. Because they are a short-lived ($\\approx1000$~yr) transition phase between AGB star and PNe, proto-planetary nebulae (PPNe) are fairly rare. However, OH maser observations reveal similar strength magnetic fields as in their progenitor AGB stars \\citep[e.g][]{Bains03, Bains04}. Additionally, a very small fraction of the PPNe maser stars show highly collimated H$_2$O maser jets. Currently 11 of these so-called 'water-fountain' sources have been identified, and they are likely the progenitor of bipolar PNe \\citep[][and references therein]{Imai07}. Polarization observations of the \\water masers of W43A have recently revealed that the maser jet is magnetically collimated \\citep{Vlemmings06c}.  Here we present the detection of a strong magnetic field in the \\water maser region of the circumstellar envelope of the PPNe candidate IRAS~19296+2227, only the second such detection in an object beyond the AGB phase. Additionally, we shortly describe the measurement of linear polarization of the \\water masers of the PNe K3-35 \\citep{Miranda01}, which was observed together with IRAS~19296+2227. The PNe K3-35 is only one of 4 PNe towards which \\water masers have been detected \\citep{Suarez07}. IRAS~19296+2227 is a source with both \\water and OH maser emission and IRAS colors indicative of an AGB star with strong mass loss. Its near kinematic distance is estimated by \\citet{Engels96} to be $3.2\\pm1.3$~kpc, which implies an infrared luminosity of $5000~L_\\cdot$. Although this luminosity is consistent with that of an AGB star, its maser properties are atypical since both the OH and \\water spectra are narrow compared with the maser spectra of OH/IR stars. However, its maser properties do resemble those of a number of PPNe candidates. Therefore, \\citet{Engels96} concluded that IRAS~19296+2227 should also be considered a PPNe candidate.  \\begin{figure*}[htf] \\resizebox{\\hsize}{!}{\\includegraphics{fig1.eps}} \\hfill \\caption[specs]{Observed VLBA spectra as measured in the image plane of the \\water masers of IRAS~19296+2227 (left) and K3-35 (right).} \\label{Fig:specs} \\end{figure*} ",
        "conclusions": "\\label{concl} We have measured a magnetic field of $B_{\\rm ||}=-135\\pm28$~mG in the \\water maser region of the PPNe candidate star IRAS 19296+2227. The field strength is of the same order as the previously measured field strengths in evolved stellar envelopes. This implies that there is no significant evolution of the magnetic field in the transition from AGB star to early PNe. The magnetic field, if representative of the large scale field, is strong enough to play an important role in the dynamics of the stellar outflow and possibly also in driving the mass-loss. However a direct interpolation of the measured field strength to the stellar atmosphere is impossible.  While no linearly polarized \\water maser emission is detected in the envelope of IRAS~19296+2227, $\\sim 1.6$\\% fractional linear polarization is found for the strongest maser feature of the PNe K3-35, implying that the masers aree at least partially saturated. With sufficiently sensitive observations, the measurement of linear polarization of the \\water masers of PNe will enable us to probe the magnetic field morphology early in the PNe phase. The circular polarization measurement and lack of linear polarization suggests the \\water masers of IRAS~19296+2227 are unsaturated."
    },
    "0806/0806.3394_arXiv.txt": {
        "abstract": "We show that the mechanical energy injection rate generated as the intra-cluster medium (ICM) flows around cold clouds may be sufficient to power the optical and near infra-red emission of nebulae observed in the central regions of a sample of seven galaxy clusters.  The energy injection rate is extremely sensitive to the velocity difference between the ICM and cold clouds, which may help to explain why optical and infra-red luminosity is often larger than expected in systems containing AGNs. We also find that mass recycling is likely to be important for the dynamics of the ICM. This effect will be strongest in the central regions of clusters where there is more than enough cold gas for its evaporation to contribute significantly to the density of the hot phase. ",
        "introduction": "Line-emitting nebulae surround approximately a third of all Brightest Cluster Galaxies (BCGs) \\citep[][]{crawf}. The occurance of these optical nebulae appears to have some correspondence to the short radiative cooling time of the cluster. Nearby examples, such as NGC1275, in Perseus, NGC4696 in Centaurus, and A1795, show extended filamentary structures up to 50 kpc from the central galaxy \\citep[see][and references therein]{hatch07}. Many of these BCGs also contain reservoirs of $10^{8}-10^{11.5}\\,{\\rm M_{\\odot}}$ of molecular hydrogen \\citep[e.g.][]{edge01}.  One of the prevalent emission lines in these nebulae is [NII]$\\lambda 6584$. Since [NII] is collisionally excited by thermal electrons, this line measures the heating rate in the gas, whereas the H$\\alpha$ line measures the ionisation rate \\citep[][]{don91}. A large heating rate compared to the ionisation rate leads to large [NII]/H$\\alpha$ and [OII]/H$\\alpha$ ratios. To date, the power source remains unknown, although numerous heating processes have been investigated. These include ionisation by the central active galactic nucleus (AGN), ionisation by hot young stars, heating by X-rays from the ICM, conduction of heat from the ICM \\citep[][]{don00}, and turbulent mixing layers \\citep[][]{crawf92}. Cosmic rays, preferentially diffusing along the magnetic field lines trailing behind rising bubbles, could possibly drive the excitation in those filaments that are located in bubble wakes \\citep[][]{rusz07}. However, there are problems associated with each possiblity listed above \\citep[see][]{hatch}. As a consequence, the nature of the power source(s) remains unknown.  Energy dissipation due to drag has, so far, not been considered as a potential source of energy for the filaments. In principle, dissipation by drag could be responsible for at least the heating of the optically emitting gas, and thus the [NII] emission, if not the H$\\alpha$ emission.  Recent numerical simulations suggest that the filamentary structures can be formed by an outflow stripping material from cold clouds \\citep[][]{pope08}. These filaments only become long if the wind is sufficiently fast. The kinematics and morphology also suggest that the filaments may be many tens of Myr old. Dissipation due to drag would necessarily occur all along the filament and would provide a spatially distributed source of energy that could power the optical emission.  So far the presence of vast quantities of cold material has been ignored in most numerical simulations of the diffuse gas in galaxy clusters. This is partly due to the difficulty of implementing a cold phase in such simulations, but also because the importance of its presence has not been hitherto established. The effect of a cold phase may not be entirely negligible. In fact, the presence of significant quantities of cold gas might alter the flow dynamics in the central regions of galaxy clusters. This could occur through the addition of mass, stripped from cold clouds, or energy and momentum transfer associated with the mass loss from the clouds.  In this article we estimate mass transfer parameters for a sample of galaxy clusters: Virgo, Perseus, Hydra, A2597, A2199, A1795 and A478. We also investigate drag as a possible power source for the optical emission of these nebulae. In section 2 we discuss the constraints on cloud properties. In section 3 we concentrate on the significance of mass transfer between the phases. Section 4 focusses on mechanical energy as a power source for the optical emission, and we summarise in section 5. ",
        "conclusions": "The effect of mass injection into a diffuse flow in a gravitational potential was investigated by \\cite{pittard04} using 1D hydrodynamical simulations. Regularly spaced shock-like structures were obtained under some conditions, which could provide an alternative explanation for the structures generally believed to be AGN-driven shocks. Other simulations revealed a differential luminosity function which resembled those deduced from observations \\citep[e.g.][]{peterson03}. However, along with these successes are a number of areas for concern. Firstly, the amount of material that was added was not related to the mass deposition rate (expected from cooling), and eventually dominated the total gas mass in the system. However, this is less of a problem if the flows are fountain-like. Secondly, the resulting X-ray luminosity was far too centrally peaked. This may be related to the form of the gravitational potential used, and it would be worth revisting the model, to see what the effect of a realistic cluster potential, and realistic mass transfer rates, would have. It may also be necessary to include thermal conduction as well, to see whether these structures would be smoothed out by the same process that evaporates the clouds.  Our analysis in the current work indicates that the rate of mechanical energy injection is comparable with the optical luminosity in the sample. Since the [NII] emission is a tracer of the heating rate of the nebular gas, it seems possible that the observed emission is indeed powered by mechanical energy injection. Furthermore, given that the energy dissipation rate is proportional to $(u-v)^{3}$ it is not surprising that optical emission occurs frequently in cool-core clusters, since AGN are frequently found in such environments. That is not to say that radiation from the AGN is responsible for the emission, but that AGN-induced bulk motions in the ICM are likely to cause significant interactions with the cold material in the central regions of clusters which may be evident through the optical emission. It may also be revealing that optical emission in clusters with AGN is generally more luminous than would be predicted. It is also possible that a component of the ${\\rm H}_\\alpha$ is due to the same process, though photoionisation also plays a role.  As we have already stated, it is also possible that mechanical energy also powers the NIR emission. Again, prominent NIR emission seems to occur in systems with powerful AGN. As a result, it is possible that the presence of cold material may automatically lead to physical processes which power the observed optical and infra-red emission. In this picture the injection of energy will necessarily be relatively constant and distributed, in accordance with observational constraints.   This work suggests that there are plausible theoretical reasons which could explain the rough proportionality between the total optical luminosity and the total mass of cold material observed by \\cite{edge01}. The total cold mass is likely to depend on the heating/cooling balance in the cluster centres, thus providing us important information on the deviation from equilibrium between these processes.  Our numerical simulations \\citep[][]{pope08} indicate that, if formed by outflows, filaments can only be formed if there is a sufficiently fast wind ($>{\\rm few} \\times 10^{7}\\,{\\rm cm\\,s^{-1}}$). Without this wind, any cold, dense material stripped from a cloud would sink towards the cluster centre with the cloud, rather than forming a long tail behind the cloud. Thus, a slower wind may lead to amorphous optical emission, without filaments. Consequently, the morphology of the optically emitting region is likely to be strongly dependent on the magnitude of the relative motion between the ICM and the cold phase.  To summarise, we have two main points:  1) There is enough cold gas within the central 30 kpc, or so, of most galaxy clusters for its evaporation to contribute significantly to the density of the hot gas in this region. Therefore, to obtain a more complete understanding of the behaviour of the ICM this effect should be included in future numerical simulations.  2) In general, it seems that the mechanical energy or momentum injection will not affect the thermal structure of the hot phase, but it may account for, at least, the [NII] component of the optical emission of the filaments and the NIR emission. Under exceptional circumstances this process could also alter the flow dynamics of the ICM and account for some heating in the central regions of clusters."
    },
    "0806/0806.3441_arXiv.txt": {
        "abstract": "We use the COMPLETE Survey's observations of the Perseus star-forming region to assess and intercompare three methods for measuring column density in molecular clouds: near-infrared extinction mapping; thermal emission mapping in the far-IR; and mapping the intensity of CO isotopologues. Overall, the structures shown by all three tracers are morphologically similar, but important differences exist amongst the tracers.  We find that the dust-based measures (near-IR extinction and thermal emission) give similar, {\\it log-normal}, distributions for the full ($\\sim 20$ pc-scale) Perseus region, once careful calibration corrections are made.   We also compare dust- and gas-based column density distributions for physically-meaningful sub-regions of Perseus, and we find significant variations in the distributions for those (smaller, $\\sim$few pc-scale) regions. Even though we have used \\co\\ data to estimate excitation temperatures, and we have corrected for opacity, the \\thco\\ maps seem unable to give column distributions that consistently resemble those from dust measures.  We have edited out the effects of the shell around the B-star HD 278942 from the column-density distribution comparisons. In that shell's interior and in the parts where it overlaps the molecular cloud, there appears to be a dearth of \\thco, which is likely due either to \\thco\\ not yet having had time to form in this young structure, and/or destruction of \\thco\\ in the molecular cloud by the HD 278942's wind and/or radiation.  We conclude that the use of either dust or gas measures of column density without extreme attention to calibration (e.g. of thermal emission zero-levels) and artifacts (e.g. the shell) is more perilous than even experts might normally admit.  And, the use of \\thco\\ data to trace total column density in detail, even after proper calibration, is unavoidably limited in utility due to threshold, depletion, and opacity effects.  If one's main aim is to map column density (rather than temperature or kinematics), then dust extinction seems the best probe, up to a limiting extinction caused by a dearth of sufficient background sources.  Linear fits amongst all three tracers' estimates of column density are given, allowing us to quantify the inherent uncertainties in using one tracer, in comparison with the others. ",
        "introduction": "In this paper, it is our goal to use data from the COMPLETE Survey of Star-Forming Regions\\footnote{All of the data from the COMPLETE (\\textbf{CO}ordinated \\textbf{M}olecular \\textbf{P}robe \\textbf{L}ine \\textbf{E}xtinction \\textbf{T}hermal \\textbf{E}mission) Survey are available online at http://www.cfa.harvard.edu/COMPLETE.} to assess and intercompare three methods for measuring column density in molecular clouds: near-infrared extinction mapping; thermal emission mapping in the far-IR; and mapping the intensity of CO isotopologues.  We wish we could offer a snapshot of {\\it volume} density, as numerical modelers can, but volume density cannot be mapped from our vantage point on Earth without making very model-dependent assumptions.  We discuss the position-position-{\\it velocity} distribution of material in other papers \\citep[e.g.][]{Erik:2007, dendro}: here we focus solely on the column density distribution.  Our aim is to offer insight into the biases of particular techniques, and to provide the best estimates to date of ``true\" column density distributions in molecular clouds suitable for comparison with current and future numerical simulations.  We focus our discussion on the Perseus molecular cloud complex, which covers nearly 10 square degrees on the sky \\citep[$\\sim200$~pc$^2$ at 250 pc; ][]{Enoch06,Hirota:2008}. A companion paper to this one \\citep{Pineda07} includes an historical perspective on the use of dust extinction and emission and of CO isotopologue mapping to study column density in molecular clouds, and it offers an in-depth look at abundance variations in the cloud. ",
        "conclusions": ""
    },
    "0806/0806.1734_arXiv.txt": {
        "abstract": "We present new spectroscopic and photometric observations of the HAT-P-1 planetary system. Spectra obtained during three transits exhibit the Rossiter-McLaughlin effect, allowing us to measure the angle between the sky projections of the stellar spin axis and orbit normal, $\\lambda = 3\\fdg 7 \\pm 2\\fdg 1$. The small value of $\\lambda$ for this and other systems suggests that the dominant planet migration mechanism preserves spin-orbit alignment. Using two new transit light curves, we refine the transit ephemeris and reduce the uncertainty in the orbital period by an order of magnitude. We find a upper limit on the orbital eccentricity of 0.067, with 99\\% confidence, by combining our new radial-velocity measurements with those obtained previously. ",
        "introduction": "Prior to 1995, it was expected that Jovian planets around other stars would inhabit wide, circular orbits similar to the Solar System gas giants. It was therefore a surprise when the first exoplanet was discovered with a minimum mass of 0.468~\\mjup\\ and a semimajor axis of only 0.05~AU \\citep{mayor95}. Since then, 85 ``hot Jupiters''---Jovian planets with periods $\\leq$10~days---have been detected around Sun-like stars \\citep{butler06,torres08}. It is unlikely that these planets formed \\emph{in situ}\\, due to the low surface densities and high temperatures of the inner regions of circumstellar disks \\citep{lin96}. A more likely scenario is that these massive planets formed at a distance of several astronomical units, and then migrated inward to their current locations.  Theories for the inward migration of planets can be divided into two broad categories. The first category involves tidal interactions between the planet and a remaining gaseous disk \\citep{lin96,moorhead08}. The second category involves few-body gravitational dynamics, such as planet--planet scattering \\citep{rasio96,chat07}, dynamical relaxation \\citep{pap01,adams03}, and Kozai cycles accompanied by tidal friction \\citep{holman97,fab07,wu07,nagasawa08}. One possible way to distinguish between these categories is to examine the present-day alignment between the stellar rotation axis and the planetary orbital axis. Assuming that these axes were initially well aligned, disk-planet tidal interactions would preserve this close alignment \\citep{ward94}, while the second category of theories would at least occasionally result in large misalignments. For example, \\citet{adams03} predict a final inclination distribution for dynamically relaxed planetary systems that peaks near $20^\\circ$ and extends to $85^\\circ$. Likewise, \\citet{fab07} and \\citet{wu07} simulated systems of planets with randomly aligned outer companions and found that the Kozai interaction resulted in a wide distribution of final orbital inclinations for the inner planet, with retrograde orbits ($\\lambda > 90^\\circ$) not uncommon. Similar results were found by \\citet{nagasawa08}, for the case in which Kozai oscillations are caused by an outer planet, rather than a companion star.  Spin-orbit alignment can be measured by taking advantage of the Rossiter--McLaughlin (RM) effect that occurs during a planetary transit. As the planet blocks portions of the rotating stellar surface, the star's rotational broadening kernel becomes asymmetric and its spectrum appears to be anomalously Doppler-shifted. The RM effect has previously been observed and modeled for eight transiting planetary systems \\citep{queloz00,winn05,winn06b,winn07a,wolf07, narita07,narita08,bouchy08,lo08,winn08}. In this work, we add HAT-P-1 to this sample.  HAT-P-1 (ADS\\,16402B) is a member of a G0V/G0V visual binary and harbors a short--period, Jovian planet. The transits of HAT-P-1b were discoverd by \\citet{bakos07} as part of the Hungarian-made Automated Telescope Network (HATNet). The planet has a 4.465~day orbital period, a mass of 0.53~\\mjup, a radius $R_P = 1.20$~\\rjup\\ \\citep{bakos07, winn07b}. We have monitored HAT-P-1 using precise radial velocity (RV) and photometric measurements made both in and out of transit in order to measure the RM effect and improve the precision with which the system's orbital parameters are known. In the following section we describe our observations and data reduction procedures. In~\\S \\ref{model} we present the transit model that we fit to our observations, and in \\S~\\ref{results} we present our results, and we conclude in \\S~\\ref{discussion} with a brief discussion. ",
        "conclusions": "\\label{discussion}  We have obtained high--precision photometric and spectroscopic measurements of the star HAT-P-1. Our in--transit spectroscopic observations clearly show the anomalous Doppler shift due to the Rossiter--McLaughlin effect, and we find that the angle between the sky projections of the stellar spin axis and the orbit normal is $3\\fdg7\\pm 2\\fdg1$. Additional Doppler measurements made during non-transit orbital phases allow us to constrain the orbital eccentricity to $e < 0.067$ with 99\\% confidence. We measured the transit times from two new light curves and refined the orbital period by nearly an order of magnitude.  The HAT-P-1 system is an interesting case for planetary migration theories because it is known to have a stellar companion in a wide orbit \\citep{bakos07}, a key ingredient for the Kozai mechanism. It is also suggestive that the radius of the planet HAT-P-1b is on the high end of theoretical expectations \\citep{bakos07,winn07b}, which may a relic of tidal energy dissipation \\citep{fab07}. Had the planetary system exhibited a large spin--orbit misalignment, it would have provided evidence for a scenario in which HAT-P-1b migrated to its current orbit as a result of Kozai oscillations, coupled with tidal dissipation within the planetary interior \\citep{fab07}. However, the small value of $\\lambda$ does not necessarily rule out the Kozai migration scenario, as small spin--orbit angles are not excluded by the simulations \\citep{fab07, wu07, nagasawa08}. But taken together with the 7 other planetary systems with small values of $\\lambda$, it seems likely that the dominant migration mechanism responsible for hot Jupiters preserves spin--orbit alignment.  Even if gravitational few-body mechanisms such as Kozai cycles do not represent the dominant migration channel for the formation of hot Jupiters, these mechanisms may be nonetheless be responsible for configurations of some of close--in planets. A prime example is the orbit of HD\\,17156b, for which \\citet{narita08} reported a 2.5$\\sigma$ detection of a large misalignment ($\\lambda = 62^\\circ \\pm 25^\\circ$). Since small values of $\\lambda$ are not excluded by any of the existing theories of planet migration, misaligned systems like HD\\,17156 provide the most important tests of the various migration mechanisms. With wide--field transit surveys discovering new planets at an accelerating pace, the sample of measured spin--orbit angles will soon be large enough and precise enough to directly confront the theory."
    },
    "0806/0806.1787_arXiv.txt": {
        "abstract": "Using a large sample of 26623 quasars with redshifts in the range $1.5 \\le z\\le 5.1$ with \\civ $\\lambda$1549 \\AA\\ emission line in Fifth Data Release of the Sloan Digital Sky Survey (SDSS), we investigate the cosmological evolution of the Baldwin Effect, i.e. the relation between the equivalent width (EW) of the \\civ emission line and continuum luminosity. We confirm the earlier result that there exists a strong correlation between the \\civ EW and the continuum luminosity, and we find that, up to $z\\approx 5$, the slope of the Baldwin Effect seems to have no effect of cosmological evolution. A sub-sample of 13960 quasars with broad \\civ $\\lambda$1549 \\AA\\ emission line from SDSS is used to explore the origin of the Baldwin Effect. We find that \\civ EW have a strong correlation with the mass of supermassive black hole (SMBH), and a weak correlation with the Eddington ratio, $\\lb/\\ledd$. This suggests that the SMBH mass is probably the primary drive for the Baldwin Effect. ",
        "introduction": "The discovery of an anti-correlation between the equivalent width (EW) of \\civ $\\lambda$1549 emission line and the continuum luminosity measured at 1450\\AA\\ in quasar rest frame, known as the Baldwin Effect, was first made by Baldwin (1977) (e.g. see a review by Shields 2006). This correlation was explored in Active Galactic Nuclei (AGNs) for other broad emission lines, such as Ly$\\alpha$, \\ciii $ \\lambda$1909, \\heiv $\\lambda$1640, \\mgii  $\\lambda$ 2800, \\siiv $\\lambda$1400, \\ovi $\\lambda$1034, \\nv $\\lambda$1240, H$\\beta$ $\\lambda$4861; for narrow emission lines, such as \\oiii$\\lambda 5007$, \\oii $\\lambda 3727$, [N{\\sc v}] $\\lambda$3426, [N{\\sc iii}] $\\lambda$3869; even for the Fe K$\\alpha$ emission line in X-ray spectrum (e.g. Wampler et al. 1984; Baldwin et al. 1989; Netzer et al. 1992; Borson \\& Green 1992; Iwasawa \\& Taniguchi 1993; Laor et al. 1995; Francis et al. 1995; Green 1996; Nandra et al. 1997; Croom et al. 2002; Dietrich et al. 2002; Baskin \\& Laor 2004; Warner et al. 2004; Page et al. 2004; Zhou et al. 2005; Netzer et al. 2006). Over the past 30 years, a significant amount of effort has been expended to confirm this effect and explore its origin and evolution.  The slope of line EW versus the luminosity for different ions depends on their ionization energy (e.g. Dietrich et al. 2002). However we are not sure whether the slope of the Baldwin Effect evolves with redshift.  Several interpretations about the origin of Baldwin effect have been proposed. It is not clear whether the Baldwin effect is driven by the underlying physical parameters, such as the redshift, Eigenvector 1 of Boroson \\& Green (1992), the Eddington ratio (i.e. the ratio of the bolometric luminosity to the Eddington luminosity), the mass of supermassive black hole (SMBH), the gas metallically (e.g. Dietrich et al. 2002; Shang et al. 2003; Sulentic et al. 2007; Baskin \\& Laor, 2004; Bachev et al. 2004; Warner et al. 2004). Baskin \\& Laor (2004) found a strong correlation between the \\civ EW and the Eddington ratio, and suggested that the Eddington ratio is the primary physical parameters driving the Baldwin effect. However, Shang et al. (2003) used the method of Spectral Principal Component (SPC) to find there is no correlation between the Baldwin effect and the Eddington ratio or the mass as underlying physical parameters.  With the reverberation mapping result, we may calculate the virial SMBH mass from the broad lines (e.g., \\hb, \\ha, \\mgii, \\civ) (e.g. Kaspi et al. 2000; Bian \\& Zhao 2004; Vestergaard \\& Peterson 2006; Salviander et al. 2007). It provides the possibility to explore these underlying physical parameters for the origin of \\civ Baldwin effect.  In the past, most of the work was made by using relatively small samples. For a sample of 20 quasars, Baldwin (1977) first discovered this effect. Wills et al. (1999) used a complete sample of 22 QSOs from the Bright Quasar Survey (LBQS) to discuss the relation between the Baldwin effect and the Eigenvector 1. For a sample of 993 quasars from the Large Bright Quasar Survey (LBQS), Green, Forster \\& Kuraszkiewicz (2001) studied the relations of \\civ EW with the continuum luminosity and redshift \\footnote {Their luminosity at 2500 \\AA is derived from the $B_{\\rm J}$ photometric magnitude, since the LBQS spectra are not spectrophotometric}. Dietrich et al. (2002) used a large sample of 744 type 1 AGNs ($0< z< 5$) from International Ultraviolet Explorer (IUE) and Hubble Space Telescope (HST) to discuss the correlations of continuum and emission-line properties in the rest-frame ultraviolet and optical spectral ranges.  Here we use a large sample of 26623 quasars with \\civ emission line from the Fifth Data Release (DR5, Adelmann et al. 2007) catalog of the Sloan Digital Sky Survey (SDSS) to discuss \\civ Baldwin Effect and its evolutionary effect, and we also use a large sample of 13960 quasars with broad \\civ emission line to discuss the underlying physical parameters of the Baldwin Effect. The sample is described in \\S 2, the results and the analysis are given in \\S 3, and the conclusions are presented in \\S 4. ",
        "conclusions": "Baldwin Effect has been explored for a long time and we are still doing our efforts to explore it. In our paper, we discuss the Baldwin Effect and its evolutionary effect from a large sample of 26623 quasars ($1.5 \\le z\\le 5.1$) from DR5. And also, using a sample of 13960 quasars with broad \\civ emission line from DR5, we explore the physical parameters that drive the \\civ Baldwin Effect. The main conclusions can be summarized as follows: (1) Baldwin Effect exists in the large sample of 26623 quasars from DR5 of SDSS; (2) According to the slopes in our four redshift-bins, up to $z\\approx 5$, there's no evolutionary effect for the slope of Baldwin Effect, however, these slopes are relatively steeper than the slope for the total sample of 26623 quasars; (3) By the 13960 quasars with SMBH masses estimate, we find that the relation between \\civ EW and SMBH mass is stronger than that between \\civ EW and $\\lb/\\ledd $ (from $R=-0.344$ to $R=-0.128$), SMBH mass seems to be a better underlying physical parameter."
    },
    "0806/0806.1608_arXiv.txt": {
        "abstract": "{} {We study the dependence of turbulent transport coefficients, such as the components of the $\\alpha$ tensor ($\\alpha_{ij}$) and the turbulent magnetic diffusivity tensor ($\\eta_{ij}$), on shear and magnetic Reynolds number in the presence of helical forcing.} {We use three-dimensional direct numerical simulations with periodic boundary conditions and measure the turbulent transport coefficients using the kinematic test field method. In all cases the magnetic Prandtl number is taken as unity.} {We find that with increasing shear the diagonal components of $\\alpha_{ij}$ quench, whereas those of $\\eta_{ij}$ increase. The antisymmetric parts of both tensors increase with increasing shear. We also propose a simple expression for the turbulent pumping velocity (or $\\gamma$ effect). This pumping velocity is proportional to the kinetic helicity of the turbulence and the vorticity of the mean flow. For negative helicity, i.e.\\ for a positive trace of $\\alpha_{ij}$, it points in the direction of the mean vorticity, i.e.\\ perpendicular to the plane of the shear flow. Our simulations support this expression for low shear and magnetic Reynolds number. The transport coefficients depend on the wavenumber of the mean flow in a Lorentzian fashion, just as for non-shearing turbulence. } {} ",
        "introduction": "Understanding the origin of non-uniform large-scale magnetic fields in stars, galaxies, and accretion discs continues to pose important challenges. Such fields are commonly thought to be the result of dynamo action that converts the kinetic energy of turbulent motions and large-scale shear into magnetic energy. The usual framework for studying such dynamo actions is mean field electrodynamics \\citep[e.g.][]{KR80}. Over the years, however, the applicability of this framework has been questioned \\citep{P81,VC92}. In particular, an important debate in this connection revolves around the role played by magnetic helicity \\citep{GD94}. Magnetic helicity is conserved for ideal (non-dissipative) magnetohydrodynamics (MHD) and also in the limit of large magnetic Reynolds numbers. This conservation places severe constraints on the growth of the mean magnetic field and may regulate the quenching of the $\\alpha$ effect as the magnetic Reynolds number ($\\Rm$) increases~\\citep{BS05}. Under certain circumstances (no magnetic helicity fluxes), $\\alpha$ quenching becomes more severe with $\\alpha$ proportional to $\\Rm^{-1}$. It has been suggested that this can be alleviated by expelling magnetic helicity through open boundaries \\citep{BF00,KMRS00}, possibly mediated by shear \\citep{VC01,SB04,BS04}. Furthermore, shear itself is an important ingredient in MHD dynamos in solar and stellar settings as, for example, in the solar tachocline. Hence, it is important to understand how shear affects the turbulent transport coefficients, in particular the components of the $\\alpha_{ij}$ and $\\eta_{ij}$ tensors.  Several studies have recently looked at various aspects of this problem \\citep[see, e.g.,][]{RK03,RK04,Bra05,RK06, RS06,LK08,BRRK08}. These works employ different tools and make different assumptions and are often applicable to limited regions of the parameter space. As a consequence, care must be taken in comparing these results. For example, using semi-analytical tools, which treat the nonlinear Lorentz force feedback perturbatively, \\cite{LK08} have found that shear can reduce  $\\alpha$ in helically forced turbulence. This is analogous to the ``rotational quenching'' of turbulent transport coefficients with increasing Coriolis or inverse Rossby numbers \\citep{KRP94,PRK96}. There are also other related cases in which the presence of shear enhances the growth rate of the dynamo. For example; using direct numerical simulations of the MHD equations in the presence of shear and {\\em non-helical} forcing, \\cite{YHSKRICM07,YHRSKRCM08} have found large-scale dynamos whose growth rate increases linearly with shear. Such scaling has also been found for $\\alpha$--shear dynamos where the $\\alpha$ effect is due to stratified convection with shear \\citep{KKB08}. Furthermore, using the kinematic test field method (described below), \\cite{BRRK08} have studied the dynamo coefficients in the presence of shear, but in the absence of helicity, and they find that Gaussian fluctuations of $\\alpha$ about zero are strong enough to drive an incoherent $\\alpha$--shear dynamo~\\citep{VB97,P07}. The significance of the incoherent $\\alpha$ effect has also been stressed by \\cite{HP08}, although their system may have also had a net $\\alpha$ effect. They dismissed this on the grounds that for an imposed uniform magnetic field $\\alpha$ is very small. However, this result disagrees with recent calculations of $\\alpha$ using the test field method \\citep{KKB08b}.  In this paper we use three-dimensional direct numerical simulations of the kinematic test field equations with {\\em helical} forcing in order to study the dependence of turbulent transport coefficients on shear and magnetic Reynolds number. In Section 2 we give a brief account of our model. Section 3 contains our results and finally we conclude in Section 4. ",
        "conclusions": "We have studied the effects of varying shear and magnetic Reynolds number on the turbulent transport coefficients in the presence of helicity in the kinematic limit. We have shown that for fixed $\\Rm$, $\\alpha$ is reduced (quenched) with increasing shear. Despite the differences in the assumptions made, this quenching is qualitatively similar to the recent results obtained by \\cite{LK08}. To the best of our knowledge this is the first numerical study to show quenching of $\\alpha$ as a function of shear in helical turbulence. We find that $\\eta_{\\rm t}$ increases with increasing shear in the range of $\\Rm$ values that we have considered here. A similar behaviour for $\\eta_{\\rm t} $ was also seen in \\cite{BRRK08} where the forcing was non-helical.  We also compute the off--diagonal components of the $\\alpha_{ij}$ and $\\eta_{ij}$ tensors. The antisymmetric part of $\\alpha_{ij}$ corresponds to a turbulent pumping velocity $\\bm{\\gamma}$ in the direction perpendicular to the plane of the shear flow. It shows a roughly linear increase with shear for small magnetic Reynolds number. We propose simple expressions for $\\gamma$ and $\\delta$ in Eq.~(\\ref{eqgda}) which show reasonable agreement with our numerical results for small $\\mbox{Sh}$ and $\\Rm$. Our expression shows that for negative helicity, $\\bm{\\gamma}$ points in the direction of the vorticity of the mean flow. Regarding the component $\\eta_{21}$ we find indications that, at least in one or two cases, this component changes sign and becomes negative. This could be of significance in connection with the shear--current effect.  We also find that all the turbulent transport coefficients depend on the wavenumber of the mean flow in a Lorentzian fashion, just as in the case of non-shearing turbulence \\citep{BRS08}. This means that the kinematic values of $\\alpha$ and $\\eta_{\\rm t}$ for $k=k_1$ are close to the values obtained for $k\\to0$. This is not the case for certain non-turbulent flows such as the Galloway-Proctor flow \\citep{Cou08,RB08}. In an earlier paper, \\cite{Cou06} considered only the limiting case $k=0$ for this flow.  Several aspects of the present investigations could be of astrophysical relevance. Turbulence in celestial bodies is helical and exhibits an $\\alpha$ effect. In addition, shear ($S$) can be an important ingredient in that the efficiency of large-scale dynamo action is determined by the product of $\\alpha$ and $S$. However, as $S$ increases, $\\alpha$ itself becomes quenched when $S$ becomes comparable with the inverse turnover time, i.e.\\ $S\\tau=O(1)$. Furthermore, the turbulent diffusivity becomes enhanced, suppressing the dynamo even further. Finally, it is found that the combined action of helicity and shear gives rise to a pumping velocity of mean magnetic field perpendicular to the plane of the shear flow. The existence of such a pumping velocity has not been emphasized before. On the other hand, it is well known that $\\alpha\\Omega$ (or rather $\\alpha S$) dynamos can have travelling wave solutions~\\citep{BBS01}. When the product of $\\alpha$ and $S$ is positive, these waves travel in the positive $z$ direction, which agrees with the direction of pumping. In the near-surface shear layer of the Sun this pumping would therefore support the equatorward migration in that layer.  It is important to understand the quenching of turbulent transport coefficients in the presence of shear beyond the kinematic approximation. In that case one needs to include the induction equation (in addition to the test field equations) and incorporate the resulting Lorentz force in the momentum equation  \\citep{BRRS08}. This would also help in alleviating problems of strong late-time fluctuations arising from the vorticity dynamo, because the vorticity dynamo tends to be suppressed by magnetic fields of equipartition strengths~\\citep{KB08}. Similarly, given their potential importance in allowing the escape of magnetic helicity, the effects of open boundary conditions also needs to be considered. These questions are under study and will be reported elsewhere."
    },
    "0806/0806.0324_arXiv.txt": {
        "abstract": "Photon breeding in relativistic jets involves multiplication of high-energy photons propagating from the jet to the external environment and back with the conversion into electron-positron pairs. The exponential growth of the energy density of these photons is a super-critical process powered by the bulk energy of the jet. The efficient deceleration of the jet outer layers creates a structured jet morphology with the fast spine and slow sheath. In initially fast and  high-power jets even the spine can be decelerated efficiently leading to very high radiative efficiencies of conversion of the jet bulk energy into radiation. The decelerating, structured jets have angular distribution of radiation significantly  broader than that predicted by a simple  blob model with a constant Lorentz factor. This reconciles the discrepancy between the high Doppler factors determined by the fits to the spectra of TeV blazars and the low apparent velocities observed at VLBI scales as well as the low jet Lorentz factors required by  the observed statistics and luminosity ratio of Fanaroff-Riley I  radio galaxies and BL Lac  objects. Photon breeding produces a population of high-energy leptons  in agreement with the constraints on the electron injection function required by spectral fits of the TeV blazars. Relativistic pairs created outside the jet and emitting gamma-rays by inverse Compton process might explain the relatively high level of the TeV emission from the misaligned jet in the radio galaxies. The  mechanism reproduces basic spectral features observed in blazars including the blazar sequence (shift of the spectral peaks towards lower energies with increasing luminosity). The mechanism is very robust and can operate in various environments characterised by the high photon density. ",
        "introduction": "A general consensus exists on the radiative processes responsible for the blazar emission: synchrotron, synchrotron self-Compton (SSC) and external radiation Compton (ERC) mechanisms. However,  neither  the distance from the central black hole to  the gamma-ray emitting region, nor the mechanisms of the energy dissipation and electron acceleration in relativistic jets are well understood. The internal shocks model \\cite{spa01}  as well as magnetic  reconnection models \\cite{lut03} have been discussed as possible scenarios for  the energy dissipation mechanism.  The arguments based on the gamma-ray transparency and the observed short time-scale variability  constrain the gamma-ray emitting region in blazars to distances of about $\\sim10^{17}$ cm \\cite{s94}. This distance is  often associated with the broad-emission line region (BLR). In internal shock model, there is no physically sound way of explaining this distance. In addition,  the internal shocks are rather  inefficient radiators, unless huge fluctuations of the Lorentz factors are involved \\cite{amb00}. Very rapid TeV variability detected from PKS 2155--304 \\cite{aha07}  and Mrk 501 \\cite{alb07} requires small distances and/or small emission region size and makes the internal shock models questionable. The  reconnection models are not yet at the level to predict the distance or efficiency.  Recently, a novel photon breeding  mechanism for dissipation of the bulk energy of relativistic jets was suggested \\cite{SP06,SP08}. The mechanism  involves multiplication of high-energy photons propagating from the jet to the external environment and back with the conversion into electron-positron pairs. The exponential growth of the energy density in these photons is a super-critical process  \\cite{SP08b} powered by the jet bulk energy. The general scheme of the converter mechanism, where conversion of charged particles to neutral ones and back plays an important role in particle acceleration, was first discussed independently in \\cite{der03,st03}. The numerical studies of the  operation of the photon breeding mechanism in ultrarelativistic shocks \\cite{st03}, showed  the high radiative efficiency of the bulk energy dissipation.   A similar problem for relativistic jets in active galactic nuclei was investigated by us \\cite{SP06,SP08}  using a ballistic model of the jet and  a detailed treatment of particle propagation and interactions. We showed that a large fraction of the jet bulk energy can be converted into gamma-rays, if the following  conditions are met: jet Lorentz factor $\\Gammaj \\gtrsim 5$--10,  the presence of the ambient  soft photon field  and a transversal or a chaotic magnetic field in the jet and the ambient medium.   In this paper we discuss the observational predictions of the photon breeding mechanism and compare them to the available data on blazars and radio galaxies. In Section \\ref{sec:sites}  we discuss the possible sites, where photon breeding can operate. In Section \\ref{sec:jet} we describe the character of the jet radiation in the photon breeding regime, particularly, the angular distributions of the emission components. We also show that it naturally resolves the Doppler factor crisis.  In Section \\ref{sec:spectr} we compare numerical simulations to the observed broad-band blazar spectra as well as to the electron  distributions obtained from the spectral fits. ",
        "conclusions": "Photon breeding mechanism  can operate at various sites characterised by sufficiently  high photon density such as in the vicinity of the accretion disc,  the broad-emission line region, at parsec scale in the infrared radiation field of hot dust, at kpc scale in the stellar radiation field and even at 100 kpc scale using cosmic-microwave radiation as a target.  As a result of operation of photon breeding a decelerating, structured jets is produced.  The angular distribution of radiation is much broader than that predicted by a blob model moving with a  constant Lorentz factor.  The broad emission pattern reconciles the discrepancy between the high Doppler factors determined by the fits to the spectra of TeV blazars and the low apparent velocities observed at VLBI scales  as well as low jet Lorentz factors required by the statistics and luminosity ratio of FR I  radio galaxies and BL Lac  objects. The inverse Compton emission from the pairs produced outside the jet predicts a high level of gamma-ray emission from  the misaligned jets in radio galaxies.  Photon breeding effectively   accelerates  predominantly high-energy  leptons  in agreement with the constraints on the electron injection function determined by the spectral fits to TeV blazars.  The  mechanism also reproduces basic spectral features observed in blazars including  the blazar sequence."
    },
    "0806/0806.0865_arXiv.txt": {
        "abstract": "{ In a companion paper we studied the detailed evolution of stellar collision products that occurred in an $N$-body simulation of the old open cluster M67 and compared our detailed models to simple prescriptions. In this paper we extend this work by studying the evolution of the collision products in open clusters as a function of mass and age of the progenitor stars.  We calculated a grid of head-on collisions covering the section of parameter space relevant for collisions in open clusters. We create detailed models of the merger remnants using an entropy-sorting algorithm and follow their subsequent evolution during the initial contraction phase, through the main sequence and up to the giant branch with our detailed stellar evolution code. We compare the location of our models in a colour-magnitude diagram to the observed blue straggler population of the old open clusters M67 and NGC 188 and find that they cover the observed blue straggler region of both clusters. For M67, collisions need to have taken place recently. Differences between the evolution tracks of the collision products and normal main sequence stars can be understood quantitatively using a simple analytic model. We present an analytic recipe that can be used in an $N$-body code to transform a precomputed evolution track for a normal star into an evolution track for a collision product. }  \\date{Today} ",
        "introduction": "Star clusters are important laboratories for a wide range of astrophysical processes. It has become clear that the evolution of a star cluster is driven by the complex interplay between stellar dynamics and stellar and binary evolution \\citep{1999A&A...348..117P,article:hurley_first_m67}. Physical collisions between stars in the dense cluster core play a pivotal role in this interaction. The products of stellar collisions between main-sequence stars potentially stand out as blue stragglers in a colour-magnitude diagram. The blue straggler population of a cluster can therefore be used to study its dynamical history. Since stellar collision products generally have a very different thermal and chemical structure than normal main-sequence stars \\citep{lombardi_collisions}, detailed calculations of their evolution are required \\citep{article:sills_evcolprod1,sills_evcolprod2}.  In a previous paper \\citep{paperI} we have developed an efficient procedure to import detailed models of stellar collision products into a fully implicit stellar evolution code and to evolve the remnants well beyond the main sequence. Our work is similar to that of \\citet{article:sills_evcolprod1} but our code is faster and much more robust, allowing for the first time a systematic study of stellar merger remnants and an exploration of the parameter space. \\citet{paperI} studied the evolution of collision remnants that occurred in the $N$-body simulation of M67 by \\citet{article:hurley_m67} and compared these with the evolution tracks predicted by the prescription of \\citet{article:hurley_bse} as well as tracks of fully mixed versions of the remnants. We found that our merger remnants have shorter main-sequence lifetimes than predicted by either the \\citet{article:hurley_bse} prescription or the fully mixed models. Our models are also brighter than normal stars of the same mass, but not as blue as fully mixed models.  In this work we explore the parameter space for collisions between low-mass main sequence stars relevant for blue straggler formation in old open clusters (such as M67) and place the findings of \\citet{paperI} in context. \\refereebf{A large grid of stellar collisions covering this parameter space was calculated by \\citet{2005MNRAS.358.1133F}. They mainly focussed on collisions with high impact velocity suitable for the galactic centre and their simulations had insufficient resolution to resolve mixing due to the collision. Our work complements theirs by investigating the mixing as well as the long term evolution of the merger remnants.} We describe the detailed structure and evolution of the collision products and discuss their dependence on the collision parameters (masses and time of collision) in \\S\\ref{sec:properties_and_tools}. We compare the results of our models with observations of blue stragglers in the open clusters M67 and NGC188 in \\S\\ref{sec:comparison}. Finally, we show how the global properties of the collision products (luminosity, radius and lifetime) can be described by a simple analytic prescription that can be included in a parametric model of stellar collisions and in $N$-body simulations (\\S\\ref{sec:recipe}). We discuss our results in \\S\\ref{sec:discussion}. ",
        "conclusions": "\\label{sec:discussion}  We have calculated a grid of stellar evolution models for the remnants of collisions between main sequence stars. These models improve on the models in the literature by using a better model for the initial structure of the merger remnant \\citep{article:bailyn_pinsonneault} or by studying a larger portion of the $t, q, M$ parameter space \\citep{article:sills_evcolprod1}.  We have compared the position of our models in the Hertzsprung-Russell diagram with the observed blue straggler regions in the open clusters M67 and NGC 188 and found that our models can populate the blue straggler region, provided that the blue stragglers in M67 are relatively `young' and those in NGC 188 are relatively `old'.  We have used our results to formulate a recipe that can be used to transform a normal evolution track, which can be obtained from fitting formulae \\citep{article:hurley_sse} or by interpolation on a table \\citep[\\emph{e.g.}][]{article:pols_evmodels,padova_tracks_2002}. This can be used to improve the treatment of collisions in $N$-body calculations and population modelling.  Our models do not include the effect of rotation. Rotation can modify our results in two ways: collision products need to lose angular momentum and therefore mass to reach hydrostatic equilibrium and rapid rotation can induce mixing. It is clear that for a collision product to survive angular momentum must be lost but the mechanism by which this occurs is unclear. Magnetic fields may be involved, but their role is not well understood. Rotational mixing can enrich the stellar envelope in helium, increasing the luminosity and extending the stellar lifetime. By ignoring rotation we have assumed that an efficient mechanism to remove the angular momentum operates that spins the star down sufficiently that rotational mixing is not important for the long-term evolution. The fact that blue stragglers in M67 are observed to be rotating slowly \\citep{mathys_bss_m67} can be interpreted as an indication that such an efficient mechanism for removing angular momentum does operate in reality.  Although our models can fill the observed blue straggler region in M67 and NGC 188, this does not directly give information about the expected distribution of blue stragglers in the colour magnitude diagram because the models need to be weighted with the probability for each collision. This requires a dynamical (or statistical, \\emph{e.g.} Monte-Carlo) model for the cluster evolution.  There are several directions for future research that could be explored. From the side of cluster modelling it is interesting to study the effect of our predicted corrections to the evolution tracks on the number of expected blue stragglers and on the blue straggler luminosity function. On the topic of detailed modelling of merger remnant evolution, including rotation and rotational mixing \\citep[extending the work of][]{ sills_evcolprod2,sills_evcolprod3} is a logical next step. However, this needs to be coupled with a mechanism for angular momentum loss from the merger remnant, which requires a better understanding of mass loss from rapidly rotating low mass stars, possibly involving magnetic fields as well. A simple equilibrium dynamo may not be adequate to model the magnetic field of a merger remnant, which might be very different from the magnetic field of a normal star following the collision. Magnetohydrodynamic models of stellar mergers and stellar evolution models including magnetic fields are also relevant in this light and are a logical next step."
    },
    "0806/0806.2098_arXiv.txt": {
        "abstract": "The rotational spectra of singly substituted $^{13}$C and $^{15}$N isotopic species of vinyl cyanide have been studied in natural abundances between 64 and 351~GHz. In combination with previous results, greatly improved spectroscopic parameters have been obtained which in turn helped to identify transitions of the $^{13}$C species for the first time in space through a molecular line survey of the extremely line-rich interstellar source Sagittarius B2(N) in the 3~mm region with some additional observations at 2~mm. The $^{13}$C species are detected in two compact ($\\sim 2.3''$), hot (170~K) cores with a column density of $\\sim 3.8 \\times 10^{16}$ and $1.1 \\times 10^{16}$ cm$^{-2}$, respectively. In the main source, the so-called ``Large Molecule Heimat'', we derive an abundance of $2.9 \\times 10^{-9}$ for each $^{13}$C species relative to H$_2$. An isotopic ratio $^{12}$C/$^{13}$C of 21 has been measured. Based on a comparison to the column densities measured for the $^{13}$C species of ethyl cyanide also detected in this survey, it is suggested that the two hot cores of Sgr~B2(N) are in different evolutionary stages. Supplementary laboratory data for the main isotopic species recorded between 92 and 342~GHz permitted an improvement of its spectroscopic parameters as well. ",
        "introduction": "\\label{intro}  Vinyl cyanide, (VyCN for short), also known as acrylonitrile or propenenitrile, H$_2$C=CH$-$C$\\equiv$N or shorter C$_2$H$_3$CN, is an important molecule in space. It was detected in the massive star-forming region Sagittarius B2 (Sgr B2 for short) by Gardner and Winnewisser as early as 1975 \\cite{VyCN_det_space}. Massive star-forming regions are comparatively dense and hot, thus, transitions of VyCN in its $v_{11} = 1$ and  $v_{15} = 1$ excited vibrational states have been detected in Orion KL \\cite{Orion-KL_325-360} and in Sgr B2(N) \\cite{VyCN-vib}. Matthews and Sears reported the first detection of this molecule in a dark cloud (TMC-1) \\cite{VyCN_TMC-1}; and very recently, Ag\\'undez et al. reported its detection in the circumstellar envelope of the late-type star IRC~+10216 (CW~Leo) \\cite{VyCN_IRC+10216}.  VyCN had also been among the molecules whose rotational spectra were studied rather early by Wilcox et al. in 1954 \\cite{VyCN_Wilcox_1954} and by Costain and Stoicheff in 1959 \\cite{VyCN_Costain_1959}. The latter also investigated the singly-substituted $^{13}$C species as well as C$_2$H$_3$C$^{15}$N and H$_2$CCDCN. More recently, additional reports have been published \\cite{VyCN_Gerry_1973,VyCN_Demaison_1973,VyCN_Sutter_1985,VyCN_Cazzoli_1988,VyCN_Demaison_1994,VyCN_Baskakov_1996,VyCN_Colmont_1997}. Its dipole moment components \\cite{VyCN_Wilcox_1954,VyCN_Sutter_1985} as well as increasingly accurate $^{14}$N hyperfine structure parameters \\cite{VyCN_Costain_1959,VyCN_Sutter_1985,VyCN_Baskakov_1996,VyCN_Colmont_1997} have been obtained. A total of 14 isotopic species have been investigated in Ref~\\cite{VyCN_Colmont_1997}, including all stable singly substituted ones, all multiply deuterated ones except D$_2$CCHCN, and the singly substituted $^{13}$C species of H$_2$CCDCN. The rotational constants were used to derive structural parameters according to several structural models. Finally, a combined fit of these previous data was reported by Thorwirth et al. in their work on the related vinyl acetylene molecule \\cite{VA_2004}.  We have carried out millimeter line observations of Sgr B2(N) with the IRAM 30\\,m telescope. These observations yielded a variety of interesting results, e.g. the first detection of aminoacetonitrile, a possible precursor of the simplest aminoacide glycine, which was secured by accompanying interferometric observations \\cite{AAN-det}. In the course of the ongoing analyses of the observational data supplementary laboratory spectroscopic investigations have been carried out to facilitate these and future analyses.  In the present contribution we report on laboratory rotational spectroscopic investigations of the singly substituted isotopic species of VyCN containing $^{13}$C or $^{15}$N as well as supplementary data on the main isotopic species. In addition, we present the first detection of $^{13}$C containing isotopologues of VyCN in space. ",
        "conclusions": "\\label{concl}  Very accurate transition frequencies have been obtained for five isotopic species of VyCN in their ground vibrational states within the millimeter-wave domain. The resulting spectroscopic parameters are accurate enough to identify the minor isotopic species in the interstellar medium or circumstellar envelopes up to about 700~GHz. This corresponds approximately with the foreseen upper frequency limit of ALMA of 720~GHz which may even be extended up to around 950~GHz later. On the other hand, none of the high frequency molecular line surveys of Orion~KL carried out with comparatively large ($\\sim$10~m) single dish telescopes covering the frequency regions $455 - 507$~GHz \\cite{Orion-KL_455-507}, $607 - 725$~GHz \\cite{Orion-KL_607-725}, and $795 - 903$~GHz \\cite{Orion-KL_795-903} revealed any features of the main isotopic species of VyCN while even vibrationally excited VyCN was observed in the $325 - 360$~GHz survey \\cite{Orion-KL_325-360}. Therefore, it is not clear at present if the minor isotopic species of VyCN will be observable with ALMA in its highest frequency bands; however the expected very high sensitivity and spatial resolution may make this fairly likely. Because of previous laboratory data~\\cite{VyCN_Demaison_1994}, predictions for the main isotopic species should be reliable up to at least 1~THz. The current data significantly improve the predictions for the weaker lines. Extensive predictions of the rotational spectra will be available in the CDMS~\\cite{CDMS1,CDMS2} for all these species. These predictions will be a contribution to minimize the so-called ``weed problem'', that means that features of abundant molecules in space, such as VyCN, may prevent weaker features of desired molecules to be identified unambiguously unless sufficiently weak features of the more abundant species are known accurately enough. The ``weed problem'' is of particular concern for radiotelescope arrays such as the SMA or, even more so, for ALMA.   With the accurate transition frequencies obtained in the present work, we could identify for the first time in space the three $^{13}$C species of VyCN in our 3\\,mm molecular line survey of the massive star forming region Sgr~B2(N). The emission is detected in two compact ($\\sim 2.3''$) hot cores with a temperature of 170~K, a column density of $\\sim 3.8 \\times 10^{16}$ and $1.1 \\times 10^{16}$ cm$^{-2}$, and a velocity of 63 and 73~km~s$^{-1}$, respectively. In the strongest source, the ``Large Molecule Heimat'' with a velocity of 63~km~s$^{-1}$, a relative abundance of $2.9 \\times 10^{-9}$ has been derived for each $^{13}$C species of VyCN, and an isotopic ratio $^{12}$C/$^{13}$C of 21 has been measured. Based on a comparison to the column densities measured for the $^{13}$C species of ethyl cyanide also detected in this survey, it is suggested that the two hot cores of Sgr~B2(N) are in different evolutionary stages."
    },
    "0806/0806.3460_arXiv.txt": {
        "abstract": "We present the results of a numerical study on the effects of metal enrichment and metal cooling on galaxy formation and cosmic star formation (SF) history using cosmological hydrodynamic simulations.  We find following differences in the simulation with metal cooling when compared to the run without it: (1) the cosmic star formation rate (SFR) is enhanced by about 50 \\& 20\\% at $z=1$ \\& 3, respectively; (2) the gas mass fraction in galaxies is lower; (3) the total baryonic mass function (gas + star) at $z=3$ does not differ significantly, but shows an increase in the number of relatively massive galaxies at $z=1$; (4) the baryonic mass fraction of intergalactic medium (IGM) is reduced at $z<3$ due to more efficient cooling and gas accretion onto galaxies. Our results suggest that the metal cooling enhances the galaxy growth by two different mechanisms: (1) increase of SF efficiency in the local interstellar medium (ISM), and (2) increase of IGM accretion onto galaxies. The former process is effective throughout most of the cosmic history, while the latter is effective only at $z<3$ when the IGM is sufficiently enriched by metals owing to feedback. ",
        "introduction": "\\label{sec:intro}  Gas cooling plays a key role in galaxy formation.  According to the hierarchical structure formation model, galaxies form in gravitationally collapsed dark matter halos.  The dissipative baryonic matter cools by emitting radiation and condenses into halos to form galaxies \\citep{Rees.Ostriker:77, White.Rees:78}. Subsequently stars form through the radiative cooling of interstellar medium \\citep[ISM; e.g.,][]{McKee.Ostriker:07}.  The cooling rate of gas varies with its chemical abundance, because different atoms and molecules have different cooling rates.  The chemical abundance of cosmic gas evolves as a function of time and environment. In the early universe, the primordial gas consists of mainly hydrogen (mass fraction of $X\\simeq 0.74$) and helium (mass fraction of $Y\\simeq 0.25$) with a negligible amount of light metals \\citep[e.g.][]{Burles.etal:01}. Throughout this paper, all the elements heavier than helium are collectively called `metals'.  Most metals are produced inside stars and spread out to the space by galactic winds driven by supernova (SN) explosions. In the metal-enriched gas, the total cooling rate is significantly enhanced relative to that of the primordial gas by the atomic emission lines from recombination of ionised metals.  Metals may influence galaxy formation in two different ways. Firstly, the SF efficiency is increased owing to the shorter gas cooling time. Star formation generally takes place in dusty metal-enriched ISM, therefore, the enhanced cooling rate by metals would boost up the SF efficiency. However, the energy/momentum feedback by SNe may suppress the subsequent star formation after the initial starburst by heating up the ambient gas. This complicates the situation, and it is not clear whether the net effect of metal cooling and SN feedback would be negative or positive. The feedback process is highly nonlinear, therefore, a direct numerical simulation would be a useful tool to explore the effects of feedback and metal cooling.  Secondly, the metals dispersed into the IGM by SN feedback also enhance the cooling of IGM.  This can lead to the increase of IGM accretion onto galaxies, because colder gas can sink into galaxies more easily. Both of the above two effects may affect the galaxy growth and cosmic SF history significantly.  Therefore it is essential to include the effects of metal cooling and chemical enrichment by feedback in the studies of galaxy formation and evolution, and capture the complex two-way interactions between galaxies and IGM. In addition, metal enrichment also affects the equation of state by altering the mean molecular weight of gas, which influences the hydrodynamic calculation of gas thermal state.  Cosmological hydrodynamic simulations are widely used in the studies of galaxy formation and cosmic SF history \\citep[e.g.,][]{Cen:92, Katz.etal:92, Cen.Ostriker:93, Katz.etal:96, Nagamine.etal:00, Pearce.etal:01, Nagamine.etal:01a, Nagamine.etal:01b, Ascasibar.etal:02, Murali.etal:02, Nagamine:02, Weinberg.etal:02, Kawata.Gibson:03, Kravtsov:03, Marri.White:03, Springel.Hernquist:03:S, Governato.etal:04, Robertson.etal:04, Weinberg.etal:04, Nagamine.etal:04, Nagamine.etal:05a, Nagamine.etal:05b, Cen.etal:05, Keres.etal:05, Finlator.etal:06, Hoeft.etal:06, Nagamine.etal:06, Scannapieco.etal:06, Tasker.Bryan:06, Dave.Oppenheimer:07, Finlator.etal:07, Governato.etal:07, Kobayashi.etal:07, Tornatore.etal:07, Finlator.etal:08}. Most of these works included the treatments of radiative cooling by H \\& He, star formation, and SN feedback. However, some of the simulations did not include the effects of metal cooling and/or chemical enrichment by galactic wind. Furthermore, the effects of metal cooling on galaxy growth and cosmic SF history has not been explored in detail using cosmological hydrodynamic simulations and presented in the literature, as it is costly to run large cosmological simulations with and without the effect of metal cooling.  In this paper, we investigate the effects of metal cooling and metal enrichment on galaxy growth and cosmic SF history using a series of cosmological hydrodynamic simulations with and without metal cooling. The aim of this paper is to single out the effects of metal cooling among our simulations.  This work can be regarded as our initial step towards the long-term goal of developing more complete cosmological hydrodynamic code with physically motivated models of star formation and feedback.  The paper is organised as follows. In \\S~\\ref{sec:method}, we describe our simulation method focusing on how we implement the metal cooling. In \\S~\\ref{sec:global}, we study the metal cooling effects on the global properties, such as cosmic SFR, phase space ($\\rho$ vs. $T$) distribution of gas, and the evolution of four phases (hot, warm-hot, diffuse, and condensed) of baryons. We then study the galaxy mass functions (\\S~\\ref{sec:MF}) and gas mass fractions (\\S~\\ref{sec:GasFrac}). Finally, we discuss and summarise our findings in \\S~\\ref{sec:summary}. ",
        "conclusions": "\\label{sec:summary}  Using cosmological hydrodynamic simulations with metal enrichment and metal cooling, we studied their effects on galaxy growth and cosmic SFR. Owing to metal cooling, the SFR density increases about 20\\% at $z=3$ and about 50\\% at $z=1$. Our results suggest that metal cooling enhances the star formation through two different processes: 1) more efficient conversion of local ISM into stars (i.e., increase of local SF efficiency), and 2) the increase of IGM accretion onto galaxies. The former process is in effect essentially at all times at $z\\lesssim 15$, because the local ISM can be instantaneously enriched by SN explosions. This process enhances the SFR, but does not noticeably change the total baryonic mass of galaxies. The latter process, on the other hand, can increase the total baryonic mass of galaxies, as well as enhancing the overall SFR density by supplying more gas for star formation.   This process becomes effective only at lower redshifts ($z\\lesssim 3$), because it takes some time to enrich the IGM up to $Z \\gtrsim 10^{-2}\\Zsun$.  In this paper, we used two different prescriptions for star formation in the runs with metal cooling: 1) the `constant $\\rhoth$' scheme and 2) the `varying $\\rhoth$' scheme. The value of $\\rhoth$ was fixed in the former scheme, while it was modulated according to the metallicity in the latter scheme. Both schemes show similar increase of cosmic SFR density (see Figure~\\ref{fig:sfr_N216L10}), which implies that the overall SF enhancement by metal cooling is not so sensitive to the choice of $\\rhoth$ (but with some differences as we discussed in \\S~\\ref{sec:cos_sfr}).  \\begin{figure*} \\begin{center} \\includegraphics[width=1.9\\columnwidth,angle=0] {kenn_mcool_cut2.eps} \\caption{Comparison of the N216L10 series with the empirical Kennicutt law. \\label{fig:kenni}} \\end{center} \\end{figure*}  The multiphase ISM model for star formation by \\citet{Springel.Hernquist:03} contains two free parameters: $\\rhoth$ and the normalisation of gas consumption time-scale, $t_\\star^0$.  The values of these two parameters were originally chosen to match the empirical Kennicutt-Schmidt law \\citep{Kennicutt:98a,Kennicutt:98b} using simulations of isolated disk galaxies.  Since we did not change the value of $t_\\star^0$ in the runs with metal cooling (as well as $\\rhoth$ in the `mc' runs), it is possible that our simulations may now violate the Kennicutt law.  However, as we show in Figure~\\ref{fig:kenni}, the plots of $\\NHI$ vs. $\\Sigma_{\\rm SFR}$ for the N216L10 and N216L10mc runs are not so different from each other. This can be understood as follows.  As Figure~\\ref{fig:Evolv4}c showed, metal cooling increases the density and lowers the temperature of star-forming gas.  The fundamental scaling relationship between SFR and cold gas density does not need to change when the metal cooling is introduced; the star-forming gas particles would simply slide upward along the Kennicutt law, $\\Sigma_{\\rm SFR} \\propto \\NHI^{1.4}$.  In the case of N216L10mv run, we varied $\\rhoth$ according to the gas metallicity, primarily scattering it to lower values because higher metallicity increases the cooling rate and lowers $\\rhoth$.  Therefore more gas particles become eligible to form stars and $\\Sigma_{\\rm SFR}$ is scattered upward above the Kennicutt law for a given value of $\\NHI$, resulting in a broader deviation from the power-law relationship of the Kennicutt law near $\\log \\NHI \\simeq 20$.  At higher values of $\\NHI$, $\\Sigma_{\\rm SFR}$ is enhanced, but still within the range of the Kennicutt law. This difference can be alleviated by adjusting the SF timescale \\citep[see][]{Springel.Hernquist:03}, however, we try to emphasise the effect of metal cooing on SFR by keeping this parameter the same in all the runs. Further investigations of the \\HI\\ aspect of our simulations is beyond the scope of the present paper, and we will present the results elsewhere.  Upon evaluating the effect of metal cooling, \\citet{Springel.Hernquist:03:S} considered two different gas phases where the cooling become important. One is the diffuse gas in galactic halos, which must radiate its thermal energy in order to collapse onto the high-density ISM, and the other is the multiphase star-forming gas.  They argued in their \\S~5.2 that the cosmic SFR density at $z\\gtrsim 6$ would not be affected by the metal cooling very much, because at high-$z$ cooling is so efficient that the gas in diffuse atmospheres of halos cools nearly instantly, even without any metal cooling. They expected their model results to be largely independent of metal enrichment, because in this regime the evolution of SFR density is driven by the fast gravitational growth of the halo mass function. For the star-forming multiphase ISM, they argued that the parameters of their model (e.g. cold gas evaporation efficiency and gas consumption time-scale) can be adjusted such that the normalisation of the Kennicutt law can be maintained, yielding to first order unaltered model predictions. However, our simulations, in particular the comparison between the N216L10 and N216L10mc run, suggest that the metal cooling can enhance the SFR density even at $z\\gtrsim 6$ by $\\sim 20-30$\\%.  This is because the gas density becomes higher and temperature becomes lower (see Figure~\\ref{fig:Evolv4}c) by responding to the enhanced cooling rate by metals.  In summary, we consider that our study serves to demonstrate the importance of metal enrichment and metal cooling in galaxy formation, beyond our original motivation to perform a simple numerical study. Our study highlights the role of metals at low-$z$ and high-$z$ to enhance the star formation in the Universe, and demonstrates that the metals increase the SF efficiency in the star-forming ISM, as well as enabling more IGM to accrete onto galaxies and fuel star formation. We expect that metal cooling would also alter some of the galaxy properties such as the mass-metallicity relationship \\citep[e.g.,][]{Tremonti.etal:04} and specific SFR \\citep[e.g.,][]{Noeske.etal:07}.  We plan to investigate these issues in the future, as well as alternative models for star formation and feedback."
    },
    "0806/0806.1653_arXiv.txt": {
        "abstract": "Simulations of the formation of Population III (Pop III) stars suggest that they were much more massive than the Pop II and Pop I stars observed today. This is due to the collapse dynamics of metal-free gas, which is regulated by the radiative cooling of molecular hydrogen. We study how the collapse of gas clouds is altered by the addition of metals to the star-forming environment by performing a series of simulations of pre-enriched star formation at various metallicities.  To make a clean comparison to metal-free star formation, we use initial conditions identical to a Pop III star formation simulation, with low ionization and no external radiation other than the cosmic microwave background (CMB). For metallicities below the critical metallicity, $Z_{cr}$, collapse proceeds similarly to the metal-free case, and only massive objects form. For metallicities well above $Z_{cr}$, efficient cooling rapidly lowers the gas temperature to the temperature of the CMB.  The gas is unable to radiatively cool below the CMB temperature, and becomes thermally stable.  For high metallicities, $Z \\ga 10^{-2.5} Z\\subsun$, this occurs early in the evolution of the gas cloud, when the density is still relatively low.  The resulting cloud-cores show little or no fragmentation, and would most likely form massive stars.  If the metallicity is not vastly above $Z_{cr}$, the cloud cools efficiently but does not reach the CMB temperature, and fragmentation into multiple objects occurs.  We conclude that there were three distinct modes of star formation at high redshift ($z \\ga 4$): a `primordial' mode, producing massive stars (10s to 100s $M\\subsun$) at very low metallicities ($Z \\la 10^{-3.75} Z\\subsun$); a CMB-regulated mode, producing moderate mass (10s of $M\\subsun$) stars at high metallicites ($Z \\ga 10^{-2.5} Z\\subsun$ at redshift $z\\sim$15-20); and a low-mass (a few $M\\subsun$) mode existing between those two metallicities. As the universe ages and the CMB temperature decreases, the range of the low mass mode extends to higher metallicities, eventually becoming the only mode of star formation. ",
        "introduction": "Understanding the nature of the first stars in the universe is, in principle, a very straightforward problem to solve.  The initial conditions are defined by well-constrained cosmological parameters (e.g., \\citet{2007ApJS..170..377S,2008arXiv0803.0547K}) and accepted calculations of Big Bang Nucleosynthesis (e.g., \\citet{1993ApJS...85..219S}).  The relative simplicity of the chemistry of primordial gas \\citep{1997NewA....2..181A,1998A&A...335..403G}, along with powerful numerical methods, have made it possible to accurately characterize the formation process of primordial stars from the assembly of their dark matter halos through the end of the optically-thin regime of the collapsing protostar \\citep{2002Sci...295...93A,2002ApJ...564...23B, 2004NewA....9..353B,2006ApJ...652....6Y,2007ApJ...654...66O,2007MNRAS.378..449G}.  Results from theory and numerical simulations suggest that the first stars were 10s to 100s $M\\subsun$ \\citep{2002Sci...295...93A,2003ApJ...589..677O,2004ApJ...603..383T, 2006ApJ...652....6Y,2007ApJ...654...66O}. A direct calculation, however, involves simulating the complex processes of accretion and radiative feedback, a capability currently just beyond the state of the art. In contrast, the challenge of understanding the second generation of stars is significantly more complex.  The initial conditions of the first stars essentially depend only on the principles of cold dark matter cosmology and the properties of molecular hydrogen, but the initial conditions of the second stars require a complete solution to the formation and evolution of their predecessors as well as host of additional chemistry and physical processes.  The chemical composition and physical conditions of second-generation star-forming environments strongly depend on the exact masses of the first stars \\citep{2002ApJ...567..532H, 2005ASPC..336...79M,2006NuPhA.777..424N,2006astro.ph..8028R} and the mechanics of their supernovae \\citep{2005ApJ...630..675K,2007ApJ...670....1G}.  At the heart of the problem regarding forming the second stars is understanding how the star formation process is altered by the introduction of the first metals in the universe, created in the explosions of the first stars. The chemical composition of a gas determines the efficiency with which it can cool radiatively. Cold, metal-free gas ($T < 10^{4}$ K) cools solely through ro-vibrational lines of H$_{2}$, whose lowest-lying transition has an energy equivalent temperature of $\\sim$ 512 K, resulting in a minimum achievable temperature of $\\sim$ 200 K.  The energy levels of H$_{2}$ become thermalized at rather low densities, $n$ $\\sim$ 10$^{3-4}$ cm$^{-3}$, where $n$ is the number density, above which the cooling rate is only proportional to $n$, instead of $n^{2}$.  Numerical simulations have shown that this creates a stalling point in the star formation process, where the free-fall collapse of metal-free gas comes to a halt as the cooling time increases to be significantly above the dynamical time \\citep{2002Sci...295...93A,2002ApJ...564...23B}.  Entry into this `loitering' phase marks the end of hierarchical fragmentation that occurs during the free-fall period, a fragmentation that occurs because the temperature is able to continually decrease with increasing density.  The final fragmentation mass scale is then set by the Jeans mass at the temperature and density corresponding to the point where the gas can no longer get colder with increasing density \\citep{1985MNRAS.214..379L,2005MNRAS.359..211L}. In the metal-free case, this yields a mass scale of $\\sim$ 1000 $M\\subsun$, resulting in the high-mass nature of the first stars.  The addition of metals enhances the cooling rate at low temperatures through fine-structure and molecular transitions, as well as by continuum emission from dust grains.  At low to moderate densities ($n <$ 10$^{9}$ cm$^{-3}$), only metals in the gas phase contribute significantly to the cooling, unless the level of metal enrichment is very high \\citep{2008MNRAS.385.1443S}. In the phases where H$_{2}$ cooling is strong ($n \\la 10^{4}$ cm$^{-3}$, an increase in the cooling due to metals is first seen only at temperatures less than the 200 K temperature minimum that exists in metal-free star formation \\citep{2008MNRAS.385.1443S}.  At abundances of $Z \\la$ 10$^{-4}$ $Z\\subsun$, where metal cooling is only significant for $T <$ 200 K, the low-density evolution of the cloud is identical to the metal-free case \\citep{2000ApJ...534..809O,2001MNRAS.328..969B,2005ApJ...626..627O}. At higher metallicities, gas-phase metal cooling becomes strong enough for the collapsing cloud to bypass the loitering phase and to undergo continued fragmentation \\citep{2001MNRAS.328..969B,2003Natur.425..812B,2006ApJ...643...26S,2007ApJ...661L...5S}. If dust grains are present, their influence becomes important at very high densities ($n >$ 10$^{12}$ cm$^{-3}$).  This has been shown to induce fragmentation for metallicities as low as 10$^{-5.5} Z\\subsun$ \\citep{2005ApJ...626..627O,2006MNRAS.369.1437S,2006ApJ...642L..61T, 2008ApJ...672..757C}.  Calculations by \\citet{2004MNRAS.351.1379S} predict that up to 30\\% of the progenitor mass is converted into dust in a pair-instability supernova, but observations of Type II supernova in the local universe remnants have not returned conclusive evidence of dust \\citep{2004MNRAS.355.1315G,2004Natur.432..596K}. In the context of star formation in the early universe, the arrival at this critical metallicity, $Z_{cr}$, at which fragmentation occurs beyond the capabilities of primordial gas, is predicted to be the point where the universal mode of star formation shifts permanently from the high-mass, solitary mode of the first stars, to the low-mass, multiply-producing mode that is presently observed  \\citep{2003Natur.425..812B,2006ApJ...643...26S, 2007ApJ...661L...5S}.  \\begin{deluxetable*}{lccccccc} \\tablecolumns{8} \\tablewidth{0pt} \\tablecaption{Simulations} \\tablehead{ \\colhead{Run} & \\colhead{IC} & \\colhead{$Z$ (Z$\\subsun$)} & \\colhead{$z_{col}$} & \\colhead{Grids} & \\colhead{Cells ($\\times10^{7}$)} & \\colhead{$n_{max}$ (cm$^{-3}$)} & \\colhead{$\\Delta t_{col}$ (yr)}} \\startdata r1\\_mf       & 1 & 0            & 14.761 & 13790 & 6.49 & 5.96 $\\times 10^{11}$ & -\\\\ r1\\_Z-6      & 1 & 10$^{-6}$     & 14.762 & 13015 & 6.42 & 6.43 $\\times 10^{11}$ & 2.9 $\\times10^{4}$\\\\ r1\\_Z-5      & 1 & 10$^{-5}$     & 14.783 & 12889 & 6.43 & 6.74 $\\times 10^{11}$ & 5.7 $\\times10^{5}$\\\\ r1\\_Z-4.25   & 1 & 10$^{-4.25}$  & 14.809 & 12955 & 6.40 & 6.75 $\\times 10^{11}$ & 1.2 $\\times10^{6}$\\\\ r1\\_Z-4      & 1 & 10$^{-4}$     & 14.830 & 12964 & 6.38 & 6.22 $\\times 10^{11}$ & 1.8 $\\times10^{6}$\\\\ r1\\_Z-3.75   & 1 & 10$^{-3.75}$  & 14.848 & 13003 & 6.38 & 6.97 $\\times 10^{11}$ & 2.2 $\\times10^{6}$\\\\ r1\\_Z-3.5    & 1 & 10$^{-3.5}$   & 14.874 & 12883 & 6.40 & 6.89 $\\times 10^{11}$ & 2.9 $\\times10^{6}$\\\\ r1\\_Z-3.25   & 1 & 10$^{-3.25}$  & 14.936 & 12888 & 6.36 & 3.97 $\\times 10^{11}$ & 4.5 $\\times10^{4}$\\\\ r1\\_Z-3      & 1 & 10$^{-3}$     & 15.073 & 12581 & 6.25 & 4.42 $\\times 10^{11}$ & 7.9 $\\times10^{6}$\\\\ r1\\_Z-2.5    & 1 & 10$^{-2.5}$   & 16.180 & 11187 & 5.69 & 4.70 $\\times 10^{11}$ & 3.3 $\\times10^{7}$\\\\ r1\\_Z-2      & 1 & 10$^{-2}$     & 19.481 & 8693  & 4.57 & 1.56 $\\times 10^{11}$ & 8.8 $\\times10^{7}$\\\\ \\tableline r2\\_mf       & 2 & 0            & 17.409 & 8684  & 4.76 & 7.97 $\\times 10^{11}$ & -\\\\ r2\\_Z-4      & 2 & 10$^{-4}$     & 17.555 & 8509  & 4.74 & 8.39 $\\times 10^{11}$ & 2.5 $\\times10^{6}$\\\\ r2\\_Z-3.5    & 2 & 10$^{-3.5}$   & 17.654 & 8476  & 4.71 & 8.20 $\\times 10^{11}$ & 4.2 $\\times10^{6}$\\\\ r2\\_Z-3      & 2 & 10$^{-3}$     & 17.955 & 8408  & 4.67 & 7.30 $\\times 10^{11}$ & 9.2 $\\times10^{6}$\\\\ r2\\_Z-2.5    & 2 & 10$^{-2.5}$   & 18.537 & 8022  & 4.51 & 4.41 $\\times 10^{11}$ & 1.8 $\\times10^{7}$\\\\ r2\\_Z-2      & 2 & 10$^{-2}$     & 20.441 & 7194  & 4.09 & 1.92 $\\times 10^{11}$ & 4.4 $\\times10^{7}$\\\\ \\tableline r3\\_mf       & 3 & 0            & 23.885 & 7771  & 4.25 & 1.63 $\\times 10^{12}$ & -\\\\ r3\\_Z-4      & 3 & 10$^{-4}$    & 23.966 & 7722  & 4.23 & 1.65 $\\times 10^{12}$ & 6.7 $\\times10^{5}$\\\\ r3\\_Z-3.5    & 3 & 10$^{-3.5}$  & 24.122 & 7640  & 4.21 & 1.42 $\\times 10^{12}$ & 1.9 $\\times10^{6}$\\\\ r3\\_Z-3      & 3 & 10$^{-3}$    & 24.390 & 7366  & 4.17 & 1.29 $\\times 10^{12}$ & 4.1 $\\times10^{6}$\\\\ r3\\_Z-2.5    & 3 & 10$^{-2.5}$  & 24.732 & 7145  & 4.13 & 1.20 $\\times 10^{12}$ & 6.7 $\\times10^{6}$\\\\ r3\\_Z-2      & 3 & 10$^{-2}$    & 25.028 & 7083  & 4.09 & 6.53 $\\times 10^{11}$ & 8.9 $\\times10^{6}$\\\\ r3\\_Z-2noCMB & 3 & 10$^{-2}$    & 25.255 & 7424  & 4.29 & 6.02 $\\times 10^{11}$ & 1.1 $\\times10^{7}$\\\\ \\enddata \\tablecomments{$z_{col}$ is the redshift at the onset of runaway collapse.  The total number of grid cells includes those that are covered by child grids at higher levels of refinement.  $n_{max}$ is the proper maximum baryon number density within the box.  $\\Delta$t$_{col}$ is the time difference to runaway collapse from the metal-free case.} \\label{tab:sims} \\end{deluxetable*}  The protostellar collapse of metal-enriched gas clouds has been studied extensively with one-zone models coupled to large chemical networks \\citep{2000ApJ...534..809O, 2003Natur.422..869S,2005ApJ...626..627O,2006MNRAS.369.1437S}.  These studies have produced insight into the evolution of the density and temperature of collapsing gas clouds with finite metallicities, but one-zone models cannot speak to the actual process of fragmentation, which requires attention to complex cloud geometries that can only be given by fully three-dimensional hydrodynamic simulations \\citep{2007astro.ph..1733L}.  The first of such simulations were carried out by \\citet{2001MNRAS.328..969B}, who included cooling from C, N, O, Fe, S, and Si, but not H$_{2}$, finding that clouds with metallicities, $Z$ $\\ge$ 10$^{-3}$ $Z\\subsun$ are able to fragment to mass scales lower than in the metal-free case.  These simulations had a mass resolution of only 100 $M\\subsun$ and were, therefore, unable to investigate the formation of solar-mass stars.  More recently, \\citet{2007ApJ...661L...5S} have performed a series of high-resolution simulations of metal-enriched gas collapse that were able to follow the evolution of gas fragments to sub-solar mass scales.  They, too, find that gas with metallicities, $Z$ $\\ge$ 10$^{-3}$ $Z\\subsun$ will fragment into multiple clumps, while gas with $Z$ $\\le$ 10$^{-4}$ $Z\\subsun$ will produce only one object.  However, \\citet{2008MNRAS.385.1443S} reported the existence of regions in density and temperature that are thermally unstable in gas with metallicities as low as 10$^{-4}$ $Z\\subsun$.  Fragmentation is traditionally thought to happen when the cooling time is less than the dynamical time, as the gas is able to cool and form perturbations before they can be smoothed out by sound waves.  When this condition, referred to as the fragmentation criterion, is satisfied, fragmentation can also be aided by thermal instabilities \\citep{1965ApJ...142..531F}, where a slight decrease in temperature or increase in density leads to a higher cooling rate, causing differences in temperature between pockets of gas to grow in a runaway fashion. In the adaptive mesh refinement simulations of \\citet{2007ApJ...661L...5S}, grid refinement was performed based on baryon and dark matter overdensities and by ensuring that the Truelove criterion, $l_{J}$ $<$ $\\Delta x$, where $l_{J}$ is the local Jeans length and $\\Delta x$ is the grid cell size, was satisfied \\citep{1997ApJ...489L.179T}.  However, refinement was not performed when the cooling time was less than the hydrodynamic time-step.  While this was not explicitly wrong, since the radiative cooling solver iterates with time-steps that are no larger than 10\\% of the cooling time, it may have artificially suppressed the growth of thermal instabilities that could have formed extra fragments.  We rerun the simulations of \\citet{2007ApJ...661L...5S} with an additional refinement criterion which ensures that the hydrodynamic time-step is always less than the cooling time on the finest level of resolution. In addition, we extend the series of simulations to include lower metallicities and to more carefully examine the metallicity range between 10$^{-4}$ $Z\\subsun$ and 10$^{-3}$ $Z\\subsun$.  We also run similiar simulations with two extra sets of initial conditions to confirm the robustness of the results. We describe the setup of our simulations and the improvements over \\citet{2007ApJ...661L...5S} in \\S\\ref{sec:setup}.  In \\S\\ref{sec:results}, we calculate the gas-phase critical metallicity, evaluate the validity of our assumption of optical thinness, and present the results of the suite of simulations. In \\S\\ref{sec:discussion}, we discuss our results in the context of star formation at high redshift and make the case for an initial mass function (IMF) that evolves over cosmic time.  We also include a discussion of the caveats and limitations of this work.  Finally, we conclude with a brief summary of the main conclusions of this work in \\S\\ref{sec:conclusion}. ",
        "conclusions": "\\label{sec:conclusion}  We have performed a series of high resolution simulations of metal-enriched star formation using cosmological, Pop III style initial conditions, and assuming fully homogeneous metal enrichment. We have shown that our results apply to more than a single star forming region by using 3 different sets of initial conditions with identical cosmological parameters and resolution, but with 3 unique random seeds with which to create the initial perturbations in the density and velocity fields.  From the results of these simulations, the main conclusions of this work are:  1. Fragmentation does not occur when the metallicity is only slightly above $Z_{cr}$, since this only leads to a small delay in the onset of the loitering phase that is brought on by a decrease in the efficiency of H$_{2}$ cooling.  The density at which the temperature begins to increase with increasing density is only marginally higher than in the metal-free case, and therefore does not lead to a signicant lowering of the minimum Jeans mass.  Within our simulations, $Z_{cr}$ is roughly 10$^{-3.9} Z\\subsun$. We find that fragmentation does not occur until the metallicity is roughly 0.5 dex above $Z_{cr}$.  2. Fragmentation is suppressed when the metallicity is high enough such that the gas is able to cool to the temperature of the CMB when the density of the collapsing cloud is still relatively low ($n \\sim 10^{4}$ cm$^{-3}$). The Jeans mass at the density and temperature at which the cloud first reaches the CMB temperature sets the minimum fragmentation mass within the cloud, and as such only massive clumps are able to form.  3. Metal-enriched star formation occurs in three modes that are separated by two metallicity thresholds, $Z_{cr}$ and $Z_{CMB}$.  $Z_{cr}$ is the conventional critical metallicity and $Z_{CMB}$ is the metallicity at which the gas is able to cool to the temperature of the CMB at a given redshift.  At the approximate collapse redshifts of our simulations, $z \\sim 20$, $Z_{CMB}$ is between 10$^{-3}$ and 10$^{-2.5} Z\\subsun$.  The three modes of star formation are: \\begin{itemize} \\item The primordial mode ($Z < Z_{cr}$) - The additional cooling provided by the metals is not enough to significantly alter the thermal structure of the cloud relative to the metal-free case.  No fragmentation occurs and the star will have a mass similar to a Pop III star. \\item The metallicity-regulated mode ($Z_{cr} < Z < Z_{CMB}$) - Metal cooling is high enough to allow the cloud to continue to cool past the metal-free loitering phase, but not high enough to cool it to the temperature of the CMB. The minimum fragmentation mass is set at lower temperatures and higher densities than in the primordial case and fragmentation into multiple objects occurs.  Based on the masses of clumps formed within our simulations, stars forming in this mode could have masses of the order of a few $M\\subsun$ or less. \\item The CMB-regulated mode ($Z > Z_{CMB}$) - Fragmentation is suppressed when the cloud is able to cool to the CMB temperature, as is described in point 2 of the conclusions.  At minimum, stars forming in the CMB-regulated mode will be more massive than stars forming in the metallicity-regulated mode.  $Z_{CMB}$ will increase as the CMB temperature lowers with time.  As such, the masses of stars forming in the CMB-regulated mode will slowly decrease. This mode will vanish altogether when the CMB temperature reaches the observed minimum temperature of nearby molecular clouds ($T \\sim$ 10 K at $z \\sim$ 2.7). \\end{itemize}  4. By demonstrating that the CMB can suppress fragmentation, we have provided a key conceptual piece CMB-IMF hypothesis \\citep{2007ApJ...664L..63T}.  As pointed out by \\citet{2007ApJ...664L..63T}, an IMF that evolves with redshift, producing more massive stars in the past, has consequences that may already be testable by observations."
    },
    "0806/0806.1009_arXiv.txt": {
        "abstract": "{} {We build a catalogue PPM-Extended (PPMX) on the ICRS system which is complete down to a well-defined limiting magnitude and contains the best presently available proper motions to be  suited for kinematical studies in the Galaxy.} {We perform a rigorous weighted least-squares adjustment of individual observations, spread over more than a century, to determine mean positions and proper motions. The stellar content of PPMX is taken from GSC 1.2 supplemented by catalogues like ARIHIP, PPM and Tycho-2 at the bright end. All observations have been weighted according to their individual accuracy.  The catalogue has been screened towards rejecting false entries in the various source catalogues.} {PPM-Extended (PPMX) is a catalogue of 18,088,920 stars containing astrometric and photometric information. Its limiting magnitude is about 15.2 in the GSC photometric system. PPMX consists of three parts: a) a survey complete down to R$_U$ = 12.8 in the magnitude system of UCAC2; b) additional stars of high-precision proper motions, and c) all other stars from GSC 1.2 identified in 2MASS. The typical accuracy of the proper motions is 2mas/y for 66 percent of the survey stars (a) and the high-precision stars (b), and about 10 mas/y for all other stars. PPMX contains photometric information from ASCC-2.5 and 2MASS.} {} ",
        "introduction": "According to IAU Resolution B2 of the XXIIIrd General Assembly (1997), the Hipparcos Catalogue \\citep{1997yCat.1239....0E} is the primary realisation of the International Celestial Reference System (ICRS) at optical wavelengths. Since then, it has been the aim of the astrometric community to extend the Hipparcos system to higher star densities and fainter limiting magnitudes. The first, and most important catalogue is Tycho-2 \\citep{2000A&A...355L..27H}, based on observations of the Tycho experiment onboard the ESA-Hipparcos satellite. The old-epoch observations of Tycho-2 were taken from new reductions \\citep{1998AJ....115.1212U} of the observations of the Astrographic Catalogue. Tycho-2 contains about 2.5 million stars and is 90 percent complete down to $V$ = 11.5. The Tycho B and V magnitudes were transferred to the Johnson system in ASCC-2.5 \\citep{2001KFNT...17..409K}.  In 2004, the UCAC2 catalogue \\citep{2004AJ....127.3043Z} has been published. It is based on a new full-sky survey made with a newly developed astrographic camera in the years from 1998 to 2004. The first epochs are taken from yet unpublished remeasures of the AGK2 plates and scans from the NPM and SPM plates, also known as the Yellow Sky (YS). To a minor extent, old epochs come also from AC and smaller astrometric catalogues of the last century. UCAC2 is not yet complete, the declination zone from about +40 degrees to the north pole is still missing. UCAC2 contains some 48 million stars down to $R_U$ = 16 mag. Proper motion errors are about 1-3 mas/yr for stars to 12th magnitude, and about 4-7 mas/yr for fainter stars to 16th magnitude.  The largest catalogue in the optical regime is USNO B1.0 \\citep{2003AJ....125..984M} with more than one billion objects. However, it should be noted that USNO B1.0 is not in the system of ICRS, it contains relative, not absolute proper motions  \\citep[see][]{2003AJ....125..984M}. A comparison of USNO B1.0 and UCAC2 performed in the present work yielded systematic differences (in areas of square degrees) of up to 5 mas/y in proper motions and up to 0.5 arcseconds at present-day epoch.  For kinematical studies in the Milky Way, more precisely in the closer neighbourhood of the Sun, a catalogue of proper motions in the ICRS system and with a well-defined completeness limit is indispensable. At Astronomisches Rechen-Institut, Heidelberg, we have therefore started an initiative to extend the Tycho-2 system to fainter magnitudes using all the relevant astrometric observations of the last century, and perform a rigorous least-squares adjustment to derive proper motions, instead of determining proper motions from the diffences of positions from two epochs.  The first result of this effort is PPM-Extended (PPMX). This paper summarizes the astrometric and photometric sources, from which PPMX is build together with its properties. An important remark is appropriate at this place. The aim of having a proper motion catalogue of homogeneous accuracy and complete down to a well-defined limiting magnitude in the visual regime (considerably fainter than Tycho-2) is not achievable at present. The main reason for this is the inhomogeneity of old-epoch observations mainly of AC. So, PPMX is a compromise between having a bright completeness limit and proper motions of homogeneous accuracy and a fainter completeness limit with a more inhomogeneous distribution of the accuracies of proper motions.  In the next section we summarise the properties of the source catalogues which enter PPMX, then we briefly describe the construction of the catalogue, and finish with an overview of the properties of PPMX. ",
        "conclusions": ""
    },
    "0806/0806.1523_arXiv.txt": {
        "abstract": "The detection of outgassing water vapor from Enceladus is one of the great breakthroughs of the \\emph{Cassini} mission. The fate of this water once ionized has been widely studied; here we investigate the effects of purely neutral-neutral interactions within the Enceladus torus. We find that, thanks in part to the polar nature of the water molecule, a cold ($\\sim 180$ K) neutral torus would undergo rapid viscous heating and spread to the extent of the observed hydroxyl cloud, before plasma effects become important.  We investigate the physics behind the spreading of the torus, paying particular attention to the competition between heating and rotational line cooling. A steady-state torus model is constructed, and it is demonstrated that the torus will be observable in the millimeter band with the upcoming \\textit{Herschel} satellite. The relative strength of rotational lines could be used to distinguish between physical models for the neutral cloud. ",
        "introduction": "Undoubtedly one of the major discoveries of the \\emph{Cassini} mission was the plume of water vapor and dust emanating from Saturn's moon Enceladus \\citep{hansen.2006,porco.2006,waite.2006}. In addition to opening up many lines of research as to the origin of the vents and their source of heat \\citep[e.g.][]{roberts.2008,meyer.2007,halevy.2008}, the discovery solved a hitherto outstanding mystery of the Saturn system: since 1993 it has been known that there exists a large population of hydroxyl radicals in the inner magnetosphere \\citep{shemansky.1993, jurac.2002}, centered near the orbit of Enceladus at $4 R_S$ ($1 R_S=6.0  \\times 10^9$ cm) and extending to $R \\sim 10 R_S$. In order to populate and maintain the OH torus, it was estimated that $\\dot N_W \\sim 10^{28}$ molecules per second of water needed to be injected into the magnetosphere \\citep{jurac.2005} from a then unknown source. \\cite{hansen.2006} estimate that the release of water from Enceladus is sufficient ($\\dot N_W > 5 \\times 10^{27} {\\rm ~s}^{-1}$) for the small moon to be the source of the observed OH torus.  The ``tiger stripe'' regions from which the water plume emerges \\citep{spitale.2007} possess inferred temperatures as high as 180K \\citep{spencer.2006}. Recently, a hotspot temperature of 180 K was measured, again by the CIRS instrument (NASA news release 2008-050). The gravitational pull of  Enceladus (escape velocity $0.2 {\\rm ~km~s}^{-1}$) does little to slow the outflow of gas at this temperature. Free expansion of water vapor at 180K  from the surface of Enceladus and subsequent intersection of the jet with previously emitted vapor will result initially in a narrow ($\\Delta R \\simeq 0.4 R_S$) torus of vapor with characteristic temperature $\\lesssim 180$K, i.e. velocity dispersion $u \\lesssim 0.5 {\\rm ~km~s}^{-1}$ \\citep{johnson.2006}.  The lifetime of the neutral water is limited by ionization and dissociation processes whose rates are listed in Table \\ref{tab:proc}.  \\cite{johnson.2006} demonstrated that radial spreading of the water torus can occur via one of these processes: charge exchange with corotating ions. In this paper we investigate a different spreading mechanism. From Table \\ref{tab:proc} it is clear that a neutral molecule does not ``feel'' the influence of the magnetospheric plasma or solar UV during a time $\\tau \\gtrsim 6 \\times 10^6$ s after production. We shall demonstrate that  this is sufficient time for the neutral water molecules to interact with \\emph{each other} many times, and for the torus to heat and spread to the extent of the observed OH torus.  The dynamics of moon-generated tori have been considered by \\cite{smyth.1993} in the context of the atomic hydrogen tori of Titan and Triton. \\cite{marconi.2003} and \\cite{decker.1994} investigated the viscous heating of the Triton torus via monte carlo simulations. The Enceladus torus is unique in consisting of nearly pure water vapor \\citep{waite.2006}. The particular properties of water as a polar molecule make the Enceladus water torus a fascinating physical and chemical system. Comets are another astrophysical system in which almost pure water is observed, and much of the work on the spectrum and collisions of water vapor has been done in this context \\citep[e.g.][]{bensch.2004,buffa.2000}. We shall make use of these results for the cometary case. Global simulations of the water torus around Saturn, including all species of neutrals and ions, have been carried out by \\cite{jurac.2002} and \\cite{jurac.2005,jurac.2007}. In these studies, no significant viscous heating was found. We discuss the reasons for this discrepancy in Section \\ref{sec:discuss}.  In Section \\ref{sec:xs} we calculate the cross section for collisions between water molecules, and estimate the collision rate in the water cloud. The nature of viscous evolution at low collision rates is discussed in Section \\ref{sec:stir}, where we also  investigate the impact of collisional inelasticity upon both spreading and heating rates. Section \\ref{sec:cooling} is devoted to the evaluation of the inelasticity of neutral water collisions in the Enceladus torus, and then in Section \\ref{sec:spread} we find the degree of viscous evolution permitted before the neutrals are lost to charge exchange. The steady-state column density profile of the heated torus is estimated in Section \\ref{sec:sstate}. In Section \\ref{sec:discuss} we discuss the consequences of the neutral-neutral interaction and make comparisons with previous studies. The suitability of the upcoming \\textit{Herschel} mission for observing the neutral water is considered in Section \\ref{sec:obs}.   \\begin{table}[htdp] \\caption{Rates of external (non neutral-neutral) processes. Charge exchange and electron impact ionization rates are applicable at radial location of Enceladus ($4 R_S$), assuming ions and neutrals occupy the same volume.  Notes: (a) from \\cite{johnson.2006}; (b) from \\cite{burger.2007}} \\begin{center} \\begin{tabular}{l|r} Process & Rate (s$^{-1}$)\\\\ \\hline Photodissociation & $1.5 \\times 10^{-7}$ (b)\\\\ Photoionization & $7 \\times 10^{-9}$ (b)\\\\ Charge exchange & $1.7 \\times 10^{-7}$ (a)\\\\ Electron impact ionization & $1.6 \\times 10^{-8}$ (b)\\\\ Keplerian orbital & $5 \\times 10^{-5}$~~~~~~\\\\ frequency at Enceladus,  $\\Omega_K$&\\\\ \\end{tabular} \\end{center} \\label{tab:proc} \\end{table}% ",
        "conclusions": "Our results serve to show that significant spreading of the torus can occur via neutral-neutral interactions alone, and thus such interactions, with correct cross sections, should be in future included in the global models such as those of \\cite{jurac.2005} which include all relevant processes.  A further issue that can be addressed is deposition of water on to the outer parts of the rings. \\cite{jurac.2007} estimate using their model that 17\\% of the water from Enceladus is precipitated on to the rings. \\cite{farrell.2008} report measurements of a decrease in plasma density at the edge of the rings, and use this to calculate the rate of absorption of ions. In both of the cases shown in Fig. \\ref{fig:coldens}, about 70 kg~s$^{-1}$ (25\\%) of the total 280 kg~s$^{-1}$ released by Enceladus precipitates on to the rings. This is probably a lower limit, both because other processes also contribute to loss on to the rings; and because our model probably represents a lower limit to the degree of spreading of the torus.  The detailed models of \\cite{jurac.2002,jurac.2005,jurac.2007} do not show significant viscous heating or interactions between neutrals. We suggest that this is because their adopted water-water cross section was that of atomic oxygen \\citep{jurac.2005}, which has radius $r \\simeq 0.6$ \\AA. In fact, the characteristic collisional size of the molecule is substantially larger, at least that of the O$^{2-}$ ion ($r \\simeq 1.4$ \\AA), and is further enhanced by the permanent dipole moment as detailed in the Introduction. Because there are fewer collisions between neutrals before they are impacted by other processes, they do not spread or heat significantly in the models of Jurac et al.  \\cite{marconi.2003} performed monte carlo simulations of the low-collision rate atomic hydrogen torus around Neptune's moon Triton. This approach follows individual particles and so does not suffer the potential inaccuracies of our thin disk model. The comparison between Marconi's results and ours, however, is encouraging. After a similar number of collision times, the torus has spread by only a slightly larger amount than we find in our Section \\ref{sec:sstate}."
    },
    "0806/0806.4306_arXiv.txt": {
        "abstract": "We present values of temperature and spectral index of the galactic diffuse radiation measured at 600 and 820 MHz along a 24 hours right ascension circle at declination $\\delta = +42^{\\circ}$. They have been obtained from a subset of absolute measurements of the sky temperature made with TRIS, an experiment devoted to the measurement of the Cosmic Microwave Background temperature at decimetric-wavelengths with an angular resolution of about $20^{\\circ}$.  Our analysis confirms the preexisting picture of the galactic diffuse emission at decimetric wavelength and improves the accuracy of the measurable quantities. In particular, the signal coming from the halo has a spectral index in the range $2.9-3.1$ above 600 MHz, depending on the sky position. In the disk, at TRIS angular resolution, the free-free emission accounts for the $11\\%$ of the overall signal at 600 MHz and $21\\%$ at 1420 MHz. The polarized component of the galactic emission, evaluated from the survey by Brouw and Spoelstra, affects the observations at TRIS angular resolution by less than $3\\%$ at 820 MHz and less than $2\\%$ at 600 MHz. Within the uncertainties, our determination of the galactic spectral index is practically unaffected by the correction for polarization.  Since the overall error budget of the sky temperatures measured by TRIS at 600 MHz, that is 66 mK(systematic)$+$18 mK (statistical), is definitely smaller than those reported in previous measurements at the same frequency, our data have been used to discuss the zero levels of the sky maps at 150, 408, 820 and 1420 MHz in literature. Concerning the 408 MHz survey, limiting our attention to the patch of sky corresponding to the region observed by TRIS, we suggest a correction of the base-level of $(+3.9\\pm 0.6)$K. ",
        "introduction": "Measurements of brightness temperature and frequency spectrum of the galactic radio emission are of primary importance for astrophysics. There are tight links between the radio signal and the cosmic ray electrons which moving through the interstellar magnetic field produce synchrotron radiation. A detailed knowledge of the galactic radio emission is also important for observational cosmology. The galactic radio emission is a foreground which hampers the search and detection of spectral distortions, anisotropies and residual polarization of the Cosmic Microwave Background (CMB). Here we present and discuss results on the diffuse galactic emission at decimetric wavelengths obtained with TRIS, an experiment dedicated to the search for CMB spectral distortions, which included a subset of observations especially made to extract the galactic signal. The complete experiment and its cosmological results are presented in the two accompanying papers I \\cite[]{zannoni08} and II \\cite[]{gervasi2008a}. TRIS is a system of three absolute radiometers operating at nominal frequencies of 600, 820 and 2500 MHz respectively, installed at Campo Imperatore (Italy, latitude $=+42^{\\circ}$). To minimize the level of radio-interferences at the observing site, the receivers were tuned at 600.5, 817.75 and 2427.75 MHz. The three antennas, wavelength-scaled pyramidal horns with the same beam ($HPBW_H\\times HPBW_E = 23^{\\circ} \\times 18^{\\circ}$) at the three frequencies, were aimed at the zenith with their E-planes tilted $7^\\circ$ westward from the meridian. Along the right ascension, in a direction nearly parallel to the antennas' H-plane, the size of the beam was $23^{\\circ}$. The radiometers have been used in two different operation modes: (1) in drift-scan mode to measure profiles of the variations of the sky temperature versus the right ascension $\\alpha$, along a circle of constant declination $\\delta = + 42^{\\circ}$ and (2) in absolute mode, to get the absolute temperature of the sky at various points along the same circle. This configuration guarantees that the same sky region can be observed with the same beam at different frequencies, so that the comparison of the signals is straightforward. Moreover, every 24 hours a complete profile ({\\it drift scan profile}) of the antenna temperature $T_a$ versus $\\alpha$ is obtained. After subtraction of the local (ground and atmospheric) contributions, data collected at night time at different epochs during a year are combined in a profile of $T_{sky}$, the sky brightness temperature, covering the full right ascension range at $\\delta =+42^{\\circ}$ \\cite[]{zannoni08}. The sky signal is a superposition of isotropic components of extragalactic origin, that is CMB and unresolved extragalactic radio sources (UERS, see \\cite[]{gervasi2008b}), and the anisotropic galactic signal:  \\begin{equation}\\label{totalpower} T_{sky}(\\nu,\\alpha,\\delta)=T_{Gal}(\\nu,\\alpha,\\delta)+T_{uers}(\\nu)+T_{CMB}(\\nu) \\end{equation}  \\noindent where $T_{Gal}\\propto \\nu^{-\\beta_{Gal}}$. The undefined frequency dependence of $T_{CMB}$ which appears in Eq. (\\ref{totalpower}) accounts for possible spectral distortions of the CMB. \\par\\noindent Variations of $T_{sky}$ \\textit{vs} $\\alpha$ are due only to $T_{Gal}$.  Analyzing these variations by means of the T-T plot technique proposed by \\cite{turtle62} we obtain $\\beta_{Gal}^{TT}$. Having $\\beta_{Gal}^{TT}$ it is then possible to disentangle the components of $T_{sky}$ \\cite[]{gervasi2008a} and get absolute values of $T_{Gal}$. An important remark is that the physical interpretation of the values of $\\beta_{Gal}$ found by means of this technique is not always straightforward. In particular, near the galactic plane, the spectral index of the diffuse emission found in this way is really produced by a blend of synchrotron and free-free emission. A possible approach followed by some authors (e.g. \\cite[]{platania1998}) to extract the thermal contribution near the galactic plane, in particular at frequencies higher than 1 GHz, is to use an HII catalog.  The outline of the paper is the following. In Section \\ref{dataset} we will describe the data set we used in our analysis. Then we will show the results obtained using just genuine TRIS data in Section \\ref{trisresults}, while a more general discussion, including also polarization effects, can be found in Section \\ref{generalconsideration}, where our data set is fully exploited. The astrophysical issues, and possible hints for future experiments, will be discussed in Section \\ref{discussion}, followed by a last section summarizing the results (Section \\ref{conclusions}). ",
        "conclusions": "\\label{conclusions}  The main purpose of the study we presented in this paper was to evaluate the galactic diffuse emission at decimetric wavelengths at $\\delta=+42^\\circ$. This work is necessary to extract the temperature of the Cosmic Microwave Background at frequencies close to 1 GHz from TRIS absolute measurements. In fact, distortions of the blackbody CMB spectrum are expected at decimetric wavelengths \\cite[]{gervasi2008a}.  \\noindent We extended however our analysis of the galactic signals and added to our measurements data in literature covering the frequency range 150 MHz - 1420 MHz.  We conclude that at angular resolutions of $\\sim 20^{\\circ}$: i) the steepening of the galactic spectral index from $\\beta_{Gal} \\simeq 2.3\\div 2.5$ to $\\beta_{Gal} \\simeq 2.8 \\div 3.0$ at high galactic latitude is confirmed; ii) the change of slope is maximum between 408 and 600 MHz; iii) this behavior is fully compatible with the features of the Cosmic Ray Electron energy spectrum and the intensity of the Galactic magnetic field; iv) in the disk the signal is a mixture of synchrotron and thermal emission from HII regions, with HII contribution to $T_{Gal}$ variable from 11$\\%$ at 600 MHz to 21$\\%$ at 1420 MHz; v) at this angular resolution the polarization of the synchrotron radiation affects for less of 1$\\%$ the spectral index values.  We have also shown, in section \\ref{Tris3}, how a detailed knowledge of the spectral index can help in recognizing and eventually controlling systematic uncertainties affecting absolute temperature scans.  The impact of the polarization on the knowledge of the galactic spectral index has been evaluated, and, for an experiments like TRIS, with a beam around $\\sim$20$^{\\circ}$ FWHM, we have shown that it is negligible everywhere, especially towards the galactic disk, even if there may be a suggestion that depolarization is less effective in the sky region $5^h\\lesssim \\alpha \\lesssim 8^h$ (at $\\delta=+42^{\\circ}$), going from the anti-center towards the halo in the northern direction (see Figure \\ref{fig5}).  TRIS-600 MHz data are accurate enough to allow a critical discussion of the zero level of the surveys at 150, 408, 820 and 1420 MHz, at least in the Northern Sky where our observations have been done. The results, obtained as described in section \\ref{par_zero_maps} and discussed at the end of the previous section, are summarized in Table \\ref{tab9}. The effect of this new assessment of the zero levels on the evaluation of the spectral index between 408 and 1420 MHz is also considered: a substantial steepening of the local spectral index, with $\\Delta \\beta \\cong +0.2$, is observed towards the regions at high galactic latitude, at TRIS angular resolution. This adjustment is important because these maps, noticeably the 408 MHz map \\cite[]{haslam82}, are frequently used as a benchmark in astrophysical and cosmological studies, e.g. for searching CMB anisotropies (e.g. Planck mission, \\cite[]{tauber2004} and SPOrt mission, \\cite[]{cortiglioni04}) and in planning future observation of the diffuse radiation from space or from ground stations."
    },
    "0806/0806.2185_arXiv.txt": {
        "abstract": "We present visible and near-infrared reflectance spectra and interpreted surface mineralogy for asteroid 10537 (1991 RY16). The spectrum of this object is without precedent amongst the Main Belt asteroids. A unique absorption band centered at 0.63 microns could be attributed to one of several mineralogies. Pronounced 1- and 2-micron absorption bands suggest that the composition of 10537 is a mixture of pyroxenes and olivine and that it originated from a parent body that was partially or fully differentiated. The closest available analog is the large Main Belt asteroid 349 Dembowska but 10537 may be an isolated fragment from a completely eroded parent body. ",
        "introduction": "The spectra of asteroids have revealed many aspects of Solar System history. Differences between the spectra of S- and C-type asteroids are thought to represent a thermal gradient from the primordial solar nebula \\citep{Gradie82}, the spectral slope of S-type asteroids gives clues to space weathering processes \\citep{Nesvorny05}, and V-type asteroids trace the heating and melting of solid bodies during the epoch of planet formation \\citep{Gaffey93b}. The discovery of new features in asteroid spectra provides insight to unstudied processes and mineralogies.  In the early Solar System proto-planetary bodies were heated by the decay of short-lived radioactive isotopes (SLRs) such as $^{26}$Al and $^{60}$Fe \\citep{Goswami05}. Temperatures in bodies that accreted quickly enough to incorporate high abundances of SLRs reached the solidus of silicates and partially melted, producing basaltic melt product \\citep{Hevey06}. Isotopic analyses suggest that the iron and basaltic achondrite meteorites represent at least 60 differentiated parent bodies \\citep{Chabot06,Yamaguchi01} and mineralogical analyses suggest that many asteroid families represent differentiated bodies \\citep{Nathues05,Gaffey84,Mothe05,Sunshine04}, thus basaltic material should be common throughout the Main Belt. However, with the exception of Vesta and the dynamically associated Vestoids \\citep{McCord70,Binzel93}, only three basaltic asteroids have been discovered \\citep{Lazzaro00,Hammergren06,Binzel06,Roig08}.  In the inner Main Belt ($a<2.5$ AU) the basaltic asteroid population is dominated by Vestoids, fragments that originated from the surface of Vesta in an impact event at least 3.5 Gyr ago \\citep{Bottke05a}. The Yarkovsky effect and resonance scattering have made it difficult to map the dynamical boundaries of the Vestoid family \\citep{Carruba05,Carruba07,Nesvorny08}, thus non-Vestoid basaltic asteroids  have yet to be unambiguously identified in the inner Main Belt.  The dynamics and in some cases the mineralogy of basaltic asteroids orbiting beyond the 3:1 mean motion resonance with Jupiter ($a>2.5$ AU) make them unlikely to have originated from the surface of Vesta \\citep{Roig08,Hardersen04,Mich02}. Thus, the three basaltic asteroids in this region [1459 Magnya, 21238 (1995 WV7) and 40521 (1995 RL95)] likely represent independent cases of differentiation.  The paucity of non-Vestoid basaltic material in the Main Belt has motivated a number of recent searches \\citep{Moskovitz08,Roig06,Hammergren06,Binzel06} that utilize the Sloan Digital Sky Survey Moving Object Catalog [SDSS MOC, \\citet{Ivezic02}]  to identify taxonomically unclassified asteroids whose photometric colors indicate basaltic surface material. We present the optical and near-infrared (NIR) spectroscopic follow-up of one such asteroid, 10537 (1991 RY16). The orbital elements of 10537 ($a=2.85$ AU, $e=0.07$ and $i=7.25^\\circ$) place it exterior to the 3:1 resonance, suggesting that it is not a fragment from Vesta. ",
        "conclusions": "We have shown that 10537 is a spectroscopically unique asteroid whose composition may be suggestive of a partially or fully differentiated parent body. Its diameter is between 5 and 15 km, as determined from a plausible range of albedos (0.05-0.4) and an absolute magnitude of $H=12.9$, and therefore is unlikely to represent an intact differentiated body \\citep{Hevey06} that has remained unfragmented since the time of planet formation \\citep{Bottke05b}.  It is unlikely that 10537 is a scattered Vestoid because it orbits exterior to the 3:1 mean motion resonance \\citep{Mich02,Roig08,Nesvorny08}. We numerically integrated the orbit of this asteroid and the four planets, Earth, Mars, Jupiter and Saturn using the N-body integration package MERCURY \\citep{Chambers99}. Our simulations show that the orbit of 10537 is stable over 2 Gyr and presently is not in any mean-motion resonance with other planetary bodies.  10537 is not a member of any major asteroid family or near any other of our non-Vestoid basaltic candidates (Fig. \\ref{fig.orbits}). Interactions with nearby non-linear secular resonances may have aided its orbital migration away from any parent body, particularly in inclination. Two such resonances are the $g+s-g_6-s_6$ and the $g+s-g_5-s_7$ \\citep{Milani90, Milani92}. The former resonance for an eccentricity of 0.08 is shown in Fig. \\ref{fig.orbits}. Accounting for the slightly larger eccentricity of 10537 would move this resonance to higher inclinations and closer to the orbit of 10537, thus increasing its relevance to the dynamical evolution of this body.  It is likely that the Yarkovsky effect also played a role in the orbital evolution of 10537. The maximum Yarkovsky drift in semi-major axis over 4.5 Gyr for a 5-km body (a lower limit for 10537) is $\\sim0.1$ AU \\citep{Bottke01}. This distance rules out the possibility that 10537 originated from the Eunomia or Eos families which may be derived from partially melted or differentiated parent bodies \\citep{Nathues05,Mothe05}. Furthermore, 10537 would have had to pass through the 5:2 or 7:3 mean-motion resonance with Jupiter to reach its current orbit, a migration that would have been slower for a 5-km body \\citep{Bottke01} than the dynamical lifetime of an asteroid in the 7:3 resonance \\citep{Tsiganis03}.  One possible parent body for 10537 is the large ($\\sim 140$ km) Main Belt asteroid 349 Dembowska. Dembowksa is one of only four known R-type asteroids in the Main Belt \\citep{Bus02} and like 10537 has an uncertain surface mineralogy. \\citet{Abell00} suggest that Dembowska has undergone some igneous processing and resembles an acapulcoite or a lodranite meteorite but has no known terrestrial analogue. \\citet{Feierberg80} suggest that it is the cumulate mantle of a fully differentiated body. \\citet{Gaffey93a} suggest it is the silicate residue from the partial melt of a partially differentiated body. It has also been suggested that Dembowska is a space weathered ordinary chondrite \\citep{Hiroi01}.  Regardless of Dembowska's interpreted mineralogy, compositional heterogeneity across its surface \\citep{Abell00} implies that a cratering event could have produced fragments like 10537. The spectral similarity between 10537 and Dembowska is consistent with this interpretation (Fig. \\ref{fig.spec}):  the difference of $\\sim10\\%$ between these spectra is less that the spectral variation amongst the Vestoids. Until its cause is better understood, the 0.63 $\\mu$m feature in the spectrum of 10537 does not argue against a genetic relation to Dembowska. Dembowksa's BAR of 1.135 and BI center of 0.9435 \\citep{Gaffey93a} place it near 10537 in Fig. \\ref{fig.gaffey}. From a purely dynamical stand point, a combination of ejection velocity \\citep{Marchi04}, interactions with secular resonances and the influence of the Yarkovsky effect suggest that 10537 could have migrated to its current orbit from the surface of Dembowska.  In spite of these similarities we note that Dembowska is merely the closest available analog to 10537 and may not be directly related in a petrogenetic sense. Under the standard paradigm of primordial clearing of the asteroid belt \\citep{Bottke05b} most large bodies were disrupted, leaving behind representative fragments. This is a likely explanation for basaltic asteroid Magnya \\citep{Mich02} and may apply to 10537 as well.  Further observations coupled with more robust mineralogical analyses and detailed dynamical simulations will be necessary to better characterize 10537. A spectroscopic survey of asteroids in the dynamical space surrounding 10537 could reveal the presence of other similar bodies and establish or refute any connection to Dembowska. However, irrespective of whether these two asteroids are related, the spectral properties of 10537 indicate that there is more mineralogical diversity in the Main Belt than can be accounted for by traditional taxonomic definitions of asteroid families."
    },
    "0806/0806.1590_arXiv.txt": {
        "abstract": "We study optical radiation of aspherical supernovae (SNe) and present an approach to verify the asphericity of SNe with optical observations of extragalactic SNe. For this purpose, we have developed a multi-dimensional Monte-Carlo radiative transfer code, \\texttt{SAMURAI} \\ (SupernovA MUlti-dimensional RAdIative transfer code). The code can compute the optical light curve and spectra both at early phases ($\\lsim$ 40 days after the explosion) and late phases ($\\sim$ 1 year after the explosion), based on hydrodynamic and nucleosynthetic models. We show that all the optical observations of SN 1998bw (associated with GRB 980425) are consistent with polar-viewed radiation of the aspherical explosion model with kinetic energy $20 \\times 10^{51}$ ergs. Properties of off-axis hypernovae are also discussed briefly. ",
        "introduction": "\\label{sec:introduction}  Although the explosion mechanism of core-collapse supernovae (SNe) is not well understood, there is several observational evidence of non-spherical explosion, obtained by the imaging of {\\it very} nearby SNe, \\eg SN 1987A \\citep{wan02} and Galactic supernova remnants \\citep{hwa04}. Even when the imaging is not possible, the detection of polarization from extragalactic SNe \\citep{wan01, kaw02, leo06} suggests that SNe are not spherical. In addition to these studies, spectroscopy of SNe can also give clues of the structure of SN explosion. For example, emission line profiles in the late time spectra ($t \\gsim$ 1 year, where $t$ is the time after the explosion) reveals the asphericity of the explosion \\citep{maz05}.  It is well established that a special class of Type Ic SNe \\footnote{SNe that do not show H, He, and strong Si absorption in the early time spectra ($t \\lsim 40$ days) are classified as Type Ic \\citep{fil97}.} are associated with the long gamma-ray bursts (GRBs, see \\citep{woo06} and references therein). This class of SNe is thought to be highly energetic, so called hypernovae (here defined as SNe with ejecta kinetic energy $E_{51} = \\KE / 10^{51} {\\rm ergs} > 10$; \\eg \\citep{nom06}). The asphericity of hypernovae is of great interest related to the nature of GRBs.  Since GRBs are induced by relativistic jets, hypernovae are also thought to be aspherical. However, the large kinetic energy of hypernovae is estimated by the analysis under the spherical symmetry. No realistic multi-dimensional explosion models have been verified against the observed early phase spectra.  To study the multi-dimensional nature of the SN explosion through the various observational facts, radiative transfer calculations are required to connect observables and hydrodynamics models. In this paper, radiative transfer in SN ejecta is solved with a multi-dimensional Monte-Carlo radiative transfer code, \\samurai \\ (SupernovA MUlti-dimensional RAdIative transfer code), based on hydrodynamic and nucleosynthetic models of hypernovae. The results are compared with observations of SN 1998bw, and implications for off-axis hypernovae are discussed.   \\begin{figure} \\includegraphics[scale=0.6, angle=270]{f1.eps} \\caption{ Aspherical explosion model A20 with the kinetic energy $E_{51} = \\KE / 10^{51} {\\rm ergs} =20$. {\\it Left}: Density distribution (log ${\\rm g \\ cm^{-3}}$) at t=10 days. The contour also shows the density. {\\it Right}: Mass fraction of \\Nifs. The velocity can be used as spatial coordinate thanks to the homologous expansion ($r \\propto v$). \\label{fig:model}} \\end{figure} ",
        "conclusions": "\\label{sec:discussion}  We have presented the detailed simulations of optical radiation with realistic jet-like hypernova models. The emergent LC and spectra are different for different viewing angles. The spectral properties are determined by the combination of aspherical abundances and anisotropic ionization states. Although the agreement of the spectra is far from perfect, the spectra of the model with mixing are in qualitative agreement with those of SN 1998bw.  The LC study shows that the kinetic energy of an aspherical model that explains SN 1998bw is $E_{51} = 20$, which is less than that of a well-fitting spherical model ($E_{51} = 50$). The early phase spectra can also be explained by the model with $E_{51} = 20$. However, it should be noted that the higher kinetic energy than the canonical SNe ($E_{51} \\sim 1$) is still required.  The simulations enable us to predict the radiation from off-axis hypernovae. The LC viewed off-axis rises more slowly than that of on-axis, and its maximum brightness is fainter (Fig. \\ref{fig:LC}). The spectra viewed off-axis show (1) a slightly lower absorption velocity, (2) stronger peaks around 4000 and 4500 \\AA\\ (narrower absorption of Fe) and (3) a stronger \\ion{Na}{i} $\\lambda$5890 line. However, the spectra still show general appearance of ``hypernovae'' or ``broad-line supernovae''. At later phases, off-axis hypernovae would show double-peaked [\\ion{O}{i}] emission profile \\citep{maz05, mae06b}.    \\begin{theacknowledgments} M.T. is supported through the JSPS (Japan Society for the Promotion of Science) Research Fellowship for Young Scientists. \\end{theacknowledgments}"
    },
    "0806/0806.1073_arXiv.txt": {
        "abstract": "We provide a brief review of the physical processes behind the radiative driving of the winds of OB stars and the Bondi-Hoyle-Lyttleton capture and accretion of a fraction of the stellar wind by a compact object, typically a neutron star, in detached high-mass X-ray binaries (HMXBs). In addition, we describe a program to develop global models of the radiatively-driven photoionized winds and accretion flows of HMXBs, with particular attention to the prototypical system Vela X-1. The models combine XSTAR photoionization calculations, HULLAC emission models appropriate to X-ray photoionized plasmas, improved models of the radiative driving of photoionized winds, FLASH time-dependent adaptive-mesh hydrodynamics calculations, and Monte Carlo radiation transport. We present two- and three-dimensional maps of the density, temperature, velocity, ionization parameter, and emissivity distributions of representative X-ray emission lines, as well as synthetic global Monte Carlo X-ray spectra. Such models help to better constrain the properties of the winds of HMXBs, which bear on such fundamental questions as the long-term evolution of these binaries and the chemical enrichment of the interstellar medium. ",
        "introduction": "As described by Castor, Abbott, and Klein (hereafter CAK) \\cite{cak75}, mass loss in the form of a high-velocity wind is driven from the surface of OB stars by radiation pressure on a multitude of resonance transitions of intermediate charge states of cosmically abundant elements. The wind is characterized by a mass-loss rate $\\Mdot\\approx 10^{-6}$--$10^{-5}~\\Msun~\\yr $ and a velocity profile $v(r)\\approx v_\\infty (1-R_\\star /r)^\\beta $, where $r$ is the distance from the OB star, $\\beta \\approx \\frac{1}{2}$, $v_\\infty\\approx 3 v_{\\rm esc}$ is the terminal velocity, $v_{\\rm esc}=(2GM_\\star /R_\\star )^{1/2}$ is the escape velocity, and $M_\\star $ and $R_\\star $ are respectively the mass and radius of the OB star. As described by Hoyle and Lyttleton \\cite{hoy39}, Bondi \\cite{bon52}, and Shapiro and Lightman \\cite{sha76}, in detached high-mass X-ray binaries (HMXBs), a compact object, typically a neutron star, captures the fraction of the wind of the OB star that passes within the accretion radius $r_{\\rm a} = \\zeta 2GM_{\\rm x}/v_{\\rm rel}^2$ of the compact object, where $\\zeta $ is a constant of order unity, $v_{\\rm rel} = (v_{\\rm w}^2+v_{\\rm x}^2)^{1/2}$ is the velocity of the wind relative to the compact object, $v_{\\rm w}$ is the velocity of the wind at $r=a$, $v_{\\rm x} = 2\\pi a/P_{\\rm orb}$ is the orbital velocity of the compact object, $a=[G(M_\\star +M_{\\rm x} )(P_{\\rm orb}/2\\pi )^2]^{1/3}$ is the binary separation, and $P_{\\rm orb}$ is the binary orbital period. The mass-accretion rate onto the compact object $\\Mdot_{\\rm a} = \\pi r_{\\rm a}^2\\rho _{\\rm w}v_{\\rm rel}$, where $\\rho _{\\rm w} = \\Mdot_{\\rm w}/4\\pi a^2v_{\\rm w}$ is the wind mass density and $\\Mdot_{\\rm w}$ is the wind mass-loss rate, hence $\\Mdot_{\\rm a} = \\frac{1}{4}\\, (r_{\\rm a}/a)^2 (v_{\\rm rel}/v_{\\rm w}) \\Mdot_{\\rm w}\\equiv f \\Mdot_{\\rm w}$. This compact expression expands out to \\begin{equation} \\Mdot_{\\rm a}\\approx 0.012\\, \\zeta^2 \\eta^{-4} (M_{\\rm x}/\\Msun )^2 (M_\\star/\\Msun )^{-8/3} (R_\\star/\\Rsun )^2 (P_{\\rm orb}/{\\rm d})^{-4/3} \\Mdot_{\\rm w}, \\end{equation} where $\\eta\\equiv (v_{\\rm w}/v_{\\rm esc})\\approx 3\\, (1-R_\\star/a)^\\beta$ for an isolated OB star. For the parameters appropriate to the prototypical detached HMXB Vela X-1 (\\S 2),\\footnote{$M_\\star =23.8~\\Msun $, $R_\\star =30~\\Rsun $, $M_{\\rm x}=1.86~\\Msun $, $P_{\\rm orb}=8.964$ d, $v_\\infty = 1700~\\rm km~s^{-1}$, $\\beta=0.8$, and $\\Mdot _{\\rm w}\\approx 10^{-6}~\\Msun~\\yr $, hence $a=3.7\\times 10^{12}$ cm, $r_{\\rm a} = 5.7\\times 10^{10}$ cm, $v_{\\rm x}=300~\\rm km~s^{-1}$, $v_{\\rm w}=880~\\rm km~s^{-1}$, $v_{\\rm rel}=930~\\rm km~s^{-1}$, $v_{\\rm esc}=550~\\rm km~s^{-1}$, and $\\eta = 1.6$.} $\\Mdot_{\\rm a}\\approx 7\\times 10^{-5} \\Mdot_{\\rm w}\\sim 7\\times 10^{-11}~\\Msun~\\yr $. Accretion of this material onto the neutron star powers an X-ray luminosity $L_{\\rm x} = G\\Mdot_{\\rm a}M_{\\rm x}/R_{\\rm x}\\sim 1\\times 10^{36}~{\\rm erg~s}^{-1}$, where $M_{\\rm x}$ and $R_{\\rm x}$ are respectively the mass and radius of the compact object. The resulting X-ray flux photoionizes the wind and reduces its ability to be radiatively driven, both because the higher ionization state of the plasma results in a reduction in the number of resonance transitions, and because the energy of the transitions shifts to shorter wavelengths where the overlap with the stellar continuum is lower. To first order, the lower radiative driving results in a reduced wind velocity $v_{\\rm w}$ near the compact object, which increases the accretion radius $r_{\\rm a}$, which increases the accretion efficiency $f$, which increases the X-ray luminosity $L_{\\rm x}$. In this way, the X-ray emission of HMXBs is the result of a complex interplay between the radiative driving of the wind of the OB star and the photoionization of the wind by the neutron star.  Known since the early days of X-ray astronomy, HMXBs have been extensively studied observationally, theoretically \\cite{hat77,mcc84,ste90}, and computationally \\cite{blo90,blo91,blo94,blo95}. They are excellent targets for X-ray spectroscopic observations because the large covering fraction of the wind and the moderate X-ray luminosities result in large volumes of photoionized plasma that produce strong recombination lines and narrow radiative recombination continua of H- and He-like ions, as well as fluorescent lines from lower charge states.  \\begin{figure} \\includegraphics[scale=0.875]{mauche_fig1.eps} % \\caption{Three views of the {\\it Chandra\\/} HETG Medium Energy Grating count spectrum of Vela X-1 in eclipse, with the strongest emission features labeled: ({\\it a\\/}) the broad 0.7--10 keV bandpass, ({\\it b\\/}) the line-rich 4--15~\\AA \\ bandpass, and ({\\it c\\/}) the 6.1--7.3~\\AA \\ bandpass containing emission lines from H- and He-like Si as well as fluorescence emission features from L-shell Si.} \\end{figure} ",
        "conclusions": ""
    },
    "0806/0806.3076_arXiv.txt": {
        "abstract": "We study a sample of Gamma--Ray Bursts detected by the Swift satellite with known redshift which show a precursor in the {\\it Swift}--BAT light curve.  We analyze the spectra of the precursors and compare them with the time integrated spectra of the prompt emission.  We find neither a correlation between the two slopes nor a tendency for the precursors spectra to be systematically harder or softer than the prompt ones. The energetics of the precursors are large: on average, they are just a factor of a few less energetic (in the source rest frame energy range 15--150 keV) than the entire bursts.  These properties do not depend upon the quiescent time between the end of the precursor and the start of the main event.  These results suggest that what has been called a ``precursor'' is not a phenomenon distinct from the main event, but is tightly connected with it, even if, in some case, the quiescent time intervals can be longer than 100 seconds. ",
        "introduction": "What happens in a Gamma--Ray Burst (GRB) before the main $\\gamma$--ray prompt event is still an open question. Koshut et al. (1995, hereafter K95) searched in the \\ba\\ sample for precursors defined as pulses with a peak intensity lower than that of the main GRB and separated from it by a quiescent phase at least as long as the main burst duration.  They found that a precursor was present in $\\sim 3$\\% of the 995 GRBs detected up to May 1994: their duration appeared weakly correlated with those of the main GRBs and on average shorter than that of the burst.  The spectral properties of the precursors (based on hardness ratios) showed no relation with those of the GRB, being both softer and harder.  Lazzati (2005; L05 hereafter) searched for precursors as weak events {\\it preceding} the \\ba\\ trigger.  He found them in $\\sim$20\\% of the bursts.  These precursors account for only a small fraction of the total GRB counts, and their duration show a weak correlation with that of the burst.  In contrast to those of K95, the precursors studied by L05 are softer than the main events.  K95 also found that the typical precursor--to--burst separation time is $\\Delta t \\simeq$100 s, whereas L05 showed that most precursors were closer to the onset of the main event (with $\\Delta t \\simeq$ 30 s). These timescales are computed in the observer frame since we do not know the redshift for most of the \\ba\\ bursts. This difference could be due to the different definition of precursor--to--burst separation adopted (peak--to--peak separation and interval between the onset times for K95 and L05, respectively).  Both analysis revealed, however, that the minimum separation is of the order of 10 s.  The main limitations of the above studies were: i) the lack of known distances, that prevented to quantify the absolute energy of the precursor event; ii) the poor spectral characterization of the precursor, that was based on low resolution spectra: typically the spectrum was described by either the hardness ratio (K95) or by a fit to a 3--4 channel broad band spectrum (L05).  Theoretical models for precursors can be separated into three classes: the ``fireball precursor'' (Li 2007; Lyutikov \\& Blandford, 2003; Meszaros \\& Rees, 2000; Daigne \\& Mochkovitch, 2002; Ruffini et al. 2001); the ``progenitor precursor'' models (Ramirez--Ruiz, MacFadyen \\& Lazzati, 2002; Lazzati \\& Begelman 2005) and the ``two step engine'' model (Wang \\& Meszaros 2007). In the first class the precursor is associated to the initially trapped fireball radiation being released when transparency is reached.  In the second class, based on the collapsar scenario, the precursor is identified with the interaction of a weakly relativistic jet with the stellar envelope. A strong terminal shock, breaking out of the envelope, is expected to produce transient emission. In both classes of models the precursors emission is predicted to be thermal, characterized by a black--body spectrum. In the third class the collapse of the progenitor star leads to the formation of a neutron star whose emission would be responsible for the precursor; subsequent accretion onto the neutron star causes its collapse onto a black hole, originating the GRB prompt.   Since the launch of the \\sw\\ satellite (Gehrels et al. 2004), redshifts have been measured for a considerable number of GRBs. It is thus possible to investigate the rest frame properties and calculate the energetics of any precursor associated with them. These of course are key physical quantities for the understanding of their origin.  A major issue however is the very definition of ``precursor'', as there is no obvious objective criterion.  For this reason we have considered a ``loose'' operative definition for their ``selection'', allowing ourselves to sub--select, {\\it a--posteriori}, those events sharing some characterizing property. Thus, in our operative approach, a precursor is defined simply as an initial signal for which: \\begin{itemize} \\item the peak flux is smaller than that of the main event in the \\sw--BAT 15--350 keV light curve; \\item the flux falls below the background level before the start of the main event. \\end{itemize} Our definition is quite similar to that adopted by K95 although we do not require the precursor to precede the bursts by a time delay at least as long as the main burst duration and, differently from L05, we do not impose the condition that a precursor did not trigger the detector. Indeed, such a definition might comprise initial pulses with very short time separation from the start of the main event, making them effectively undistinguishable from first pulses of the burst prompt emission. But how long should the temporal separation be to let us distinguish among the two?  Or are there, instead, other (e.g. spectral) properties which can neatly distinguish precursors from main events, independently of the time separation?  And are they really physically different?  These are some of the issues we are trying to investigate.  We adopt a $\\Lambda$CDM cosmology with $\\Omega_\\Lambda=h_0=0.7$ and $\\Omega_{\\rm M}=0.3$. ",
        "conclusions": "Our results point to a clear but puzzling conclusion: the spectra and energetics of the selected initial pulses, being them bona fide precursors or not, are indistinguishable from those of the main event. While this could be not surprising for ``precursors'' which were in fact the initial pulse of the main event, in cases like GRB 060124 and GRB 050820A the precursor precedes the main event by $\\sim$100 s (rest frame time): yet they behave as the rest of the main emission, like ``normal'' initial pulses.  This forces us to re--consider what the very same precursor phenomenon is.  Our finding contrasts with that by L05, who found precursors much fainter and significantly softer than the main event.  However the precursor selection criteria are different and in particular the requirement by L05 that the precursors did not trigger \\ba\\ obviously biased the sample against strong precursors. Our result are instead more consistent with the Koshut et al. (1995) one, whose selection criteria is similar to ours. We therefore cannot exclude that there are two kinds of ``precursors'': one as strong as and spectrally similar to the main event and the other being softer and dimmer. {\\it But -- independent of that -- both can occur $\\sim$100 s before the main event.} Indeed, this long delay is both the most intriguing feature and the main difficulty for all the proposed progenitor interpretation.  As discussed by Wang \\& Meszaros (2007), the progenitor class of models cannot explain delays longer than $\\sim$10 s.  The origin of quiescent times has been discussed by Ramirez--Ruiz, Merloni \\& Rees (2001), who considered the possibility that a temporal modulation in Lorentz factor of ejected shells/relativistic outflow would lead to time dependent emission via dissipation in internal shocks.  ``Fireball'' models predict too short quiescent timescales if the main prompt emission mechanism is internal shocks taking place at typical radii $R= 10^{13}R_{13}$ cm, since $\\Delta t\\sim R/(c\\Gamma^2)\\sim 0.03 R_{\\rm 13}/\\Gamma^2_2$ s. External shocks occurring at $R\\sim 10^{16}$ cm can lead to time delays similar to what observed, but -- in the case of a homogeneous fireball interacting with an homogeneous interstellar medium (ISM) -- this process hardly accounts for fast prompt variability, suppressed by the curvature effect.  More complex external shock scenarios can overcome this problem (e.g. Dermer et al. 1999), but in turn require a strongly clumped ISM.  A second issue emerging from our results concerns the spectral shape of the precursor. The non--thermal appearance of the spectra is not the chief problem, as this may arise as convolution of black-body emission at different temperatures and/or from different locations, consistently with the predicted thermal character. What remains puzzling (or revealing) is that, on average, the power--law fit spectral indices are very similar to those of the main event. The large energetics of the precursors studied here is also difficult to explain within the precursor models proposed so far as, whatever the progenitor nature, they rival the main event energetics.  In the collapsar model, the precursor photons may be produced in a region emerged from the progenitor star. Indeed, heated cocoon material has been proposed as responsible for the precursor (Ramirez-Ruiz, McFadyen \\& Lazzati 2002), but the expected energetics would be low compared to our findings.  This also applies to the scenario proposed by Lazzati \\& Begelman (2005), where the jet opening angle increases in time, so an observer off axis could detect the prompt emission after the precursor, when the jet angle becomes equal to the viewing angle. The ``two steps'' engine model (Wang \\& Meszaros 2007) envisages that the precursor is associated with the cooling phase of the proto--neutron star and the delay time should correspond to the accretion phase which ultimately leads to the collapse of the neutron star to a black hole, when the ``normal'' GRB activity begins.  An alternative possibility is that precursors do not represent any distinct physical process, but are simply a manifestation of the same phenomenon producing the prompt emission, which sometimes does give raise to quiescent intervals between emission peaks. We can put an upper limit on the energy emitted during these quiescent times considering that the BAT sensitivity for a 5$\\sigma$ detection is $2\\times 10^{-10} (dT/20~ {\\rm ks})^{-0.5}$ erg cm$^{-2}$ s$^{-1}$ (Markwardt et al. 2007), where $dT$ is the exposure time.  Using the delay times reported in Tab. 1 we estimate the mean value of the 1$\\sigma$ upper limit to the energy emitted during the quiescence. The ratio of these limits to the precursors energy ranges from 0.012 for GRB~061007 to 0.25 for GRB~060124, with an average of 0.14.  Finally, we applied the same selection criterion, adopted for precursors, to search for emission episodes (``postcursors'') following the main bursts and separated by a quiescent phase. Within the 15 GRBs with precursors, GRB~060210 and GRB~0508020A show two and three pulses after the main burst\\footnote{We recall that for GRB~050820A, \\sw\\ entered the {\\it SAA} during the main event and therefore we could not perform the spectral analysis of the BAT data.}.  For GRB~060210 the two postcursors (separated by 60 and 150 s from the end of the main burst) have spectral indices $-1.76 \\pm 0.28$ and $-1.83\\pm0.39$ and energetics $(7.31 \\pm 2.29)\\times 10^{51}$ erg and $(5.04\\pm2.14)\\times 10^{51}$ erg, respectively. The spectra are softer and the energetics smaller than the main event and the precursor. Since this is the only burst, in our sample, having both a precursor and a postcursor, we cannot draw any strong conclusion. We plan to study spectra and energetics of postcursors by relaxing the condition of having also a precursor in a forthcoming paper."
    },
    "0806/0806.1842_arXiv.txt": {
        "abstract": "In this paper we study the sensitivity of the Large Sky Area Multi-Object Fiber Spectroscopic Telescope (LAMOST) project to the determination of cosmological parameters, employing the Monte Carlo Markov Chains (MCMC) method. For comparison, we first analyze the constraints on cosmological parameters from current observational data, including  WMAP, SDSS and SN Ia. We then simulate the $3$D matter power spectrum data expected from LAMOST, together with the simulated CMB data for PLANCK and the SN Ia from 5-year Supernovae Legacy Survey (SNLS). With the simulated data, we investigate the future improvement on cosmological parameter constraints, emphasizing the role of LAMOST. Our results show the potential of LAMOST in probing for the cosmological parameters, especially in constraining the equation-of-state (EoS) of the dark energy and the neutrino mass. ",
        "introduction": "The measurement of the large scale galaxy clustering has been an important probe in constraining the cosmological models. The large scale structure (LSS) measurements have made remarkable progresses by the observational efforts such as 2dFGRS and Sloan Digital Sky Survey (SDSS), which have provided an accurate measurement of the galaxy power spectrum and given a robust constraint on cosmological parameters \\cite{Tegmark:2006az,Cole:2005sx}.    The LAMOST \\cite{lamost} project is a 4m quasi-meridian reflecting Schmidt telescope laid down on the ground. It has a 5 degree field of view, and may accommodate as many as 4000 optical fibers and the light from 4000 celestial objects will be led into a number of spectrographs simultaneously. Thus the telescope will be the one that possesses the highest spectrum acquiring rate in the world. The spectroscopic survey which contains the information about the radial positions of galaxies, can probe the 3D distribution of galaxies effectively. In this paper, we study the sensitivity of LAMOST to the determination of cosmological parameters with the simulated galaxy power spectrum. In our analysis, we also consider the simulated observations for the future CMB and SN Ia measurements from PLANCK and the 5-year SNLS, which are presumably conducted during the same time period as LAMOST survey. Our results indicate that the LAMOST has the promising potential in probing for the cosmological parameters, especially in constraining on the EoS of the dark energy and the neutrino mass.   The paper is organized as follows: In Section II we describe the method of fitting and the simulation technique. In section III, we present the results and discussions. The last section contains a summary. ",
        "conclusions": "\\begin{table}{\\footnotesize TABLE I. Constraints on cosmological parameters from the current observations and the future simulations.  For the current constraints we have shown the mean values $1\\sigma$ (Mean) and the best fit results together. For the future mocked data we list the standard deviation (SD) of these parameters with fiducial model I (FMI) and fiducial model II (FMII). In order to highlight the contribution from LAMOST, we compare the results with/without LAMOST. \\begin{center}  \\begin{tabular}{|c|cccc|cccc|}   \\hline  &\\multicolumn{2}{c|}{Current for $\\Lambda$CDM}&\\multicolumn{2}{c|}{~Future~(SD with FMI)}&\\multicolumn{2}{c|}{Current for Run w}&\\multicolumn{2}{c|}{~Future~(SD with FMII)}\\\\  \\cline{2-9}  &\\multicolumn{1}{c|}{~Best~Fit~}&\\multicolumn{1}{c|}{~~~~Mean~~~~}&\\multicolumn{1}{c|}{PLANCK$+$ SNLS}&\\multicolumn{1}{c|}{PLANCK$+$ SNLS}&\\multicolumn{1}{c|}{~Best~Fit~}&\\multicolumn{1}{c|}{~~~~Mean~~~~} &\\multicolumn{1}{c|}{PLANCK$+$SNLS}&\\multicolumn{1}{c|}{PLANCK$+$ SNLS} \\\\  &\\multicolumn{1}{c|}{}&\\multicolumn{1}{c|}{}&\\multicolumn{1}{c|}{}&$+$LAMOST&\\multicolumn{1}{c|}{}&\\multicolumn{1}{c|}{}&\\multicolumn{1}{c|}{}&$+$LAMOST\\\\  \\hline  $w_0$&$-1$&$-1$&$0.118$&$0.100$&$-1.16$&$-1.03^{+0.15}_{-0.15}$&$0.0899$&$0.0598$\\\\  \\hline  $w_a$&$0$&$0$&$0.522$&$0.417$&$0.968$&$0.405^{+0.562}_{-0.587}$&$0.332$&$0.162$\\\\  \\hline  $~\\Omega_{\\Lambda}~$&$0.760$&$0.762^{+0.015}_{-0.015}$&$0.0115$&$0.00547$&$0.756$&$0.760^{+0.017}_{-0.018}$&$0.0125$&$0.00460$\\\\  \\hline  $H_0$&$73.1$&$73.3^{+1.6}_{-1.7}$&$1.594$&$0.828$&$70.3$&$71.2^{+2.3}_{-2.3}$&$1.840$&$0.673$\\\\  \\hline  $\\sigma_8$&$0.769$&$0.755\\pm0.031$&$0.0223$&$0.0174$&$0.634$&$0.675\\pm0.068$&$0.0299$&$0.0220$\\\\  \\hline     \\end{tabular} \\end{center}} \\end{table}  In Table I, we present the numerical results of the constraints on the cosmological parameters from the current data and the error forecast from the simulated data. To show the importance of LAMOST, we compare the two sets of results, one from PLANCK $+$ SNLS, the other from  PLANCK $+$ SNLS $+$ LAMOST. As we know, the matter power spectrum is directly related to the horizon size at matter-radiation equality, in turn the matter power spectrum will make accurate measurement of $\\Omega_m h$. On the other hand, there are degeneracies between $\\Omega_m h$ and the other cosmological parameters, {\\it e.g.} $\\Omega_{\\Lambda}$, $H_0$, $w_0$, $w_a$ and so on, hence the tight constraint on $\\Omega_m h$ will be helpful for breaking these degeneracies and improve the constraints on these cosmological parameters. For example, in Table I, one can find the constraints on $\\Omega_{\\Lambda}$ and $H_0$ are tightened a lot by including LAMOST. Also from Figure 1, one can see the constraits on the age of universe is also shrank obviously, this is because the age is directly related to the hubble constant and $\\Omega_m$.   In figure \\ref{fig1}, we plot the $2-$D cross correlation and $1-$D probability distribution of some of the basic cosmological parameters in Eq.(\\ref{para1}) and also some of the reduced parameters. The black solid lines are given by fitting with the simulated PLANCK and SNLS, and the red solid lines are provided by including the simulated LAMOST data. From the comparison, we find that LAMOST have promising potentials in constraining cosmological parameters, such as, the EoS of dark energy, the dark energy density budget $\\Omega_{\\Lambda}$, the age of universe, $\\sigma_8$ and the Hubble constant.  In order to see explicitly the effect of LAMOST on dark energy constraints, in figure \\ref{fig2}, we plot the $2\\sigma$ confidence level contours on $w_0$ and $w_a$. The black solid line is given by the current constraints and the red solid line is given by fitting with the simulated data of PLANCK and SNLS 5 year data with fiducial model II, while the red dashed line is given by including the simulated LAMOST data. This comparison shows clearly LAMOST will contribute significantly in tightening the constraints on the EoS of dark energy. Numerically we find the best fit model with the current data is given by the dynamical quintom model with EoS across -1, however the cosmological constant is within $1\\sigma$ confidence level. The future PLANCK measurement and 5-year SNLS SN Ia will be able to distinguish the cosmological constant from the dynamical model at $2\\sigma$ confidence level, while LAMOST can improve this sensitivity significantly at $3.3\\sigma$. The blue solid lines and blue dashed lines show the comparison between the results with and without LAMOST for fiducial model I.  On the other hand, for the parameter related to the inflation models, such as $n_s$ and $A_s$, the constraints are mainly from PLANCK, as pointed out in our previous paper\\cite{07081111}. Adding in LAMOST can not further tighten the constraints. We have also done another analysis with the additional parameters $\\alpha_s$ and $r$, and obtained the similar conclusion, where $\\alpha_s$ characterizes the running of the primordial power spectrum index and $r$ is the ratio of tensor to scalar perturbations.   Now we study the cosmological constraint on the neutrino mass by adding in a new parameter $f_{\\nu}$ in Eq.(\\ref{para1}). The parameter $f_{\\nu}$ is the dark matter neutrino fraction at present, namely, \\begin{equation} f_{\\nu}\\equiv\\frac{\\rho_{\\nu}}{\\rho_{DM}}=\\frac{\\Sigma m_{\\nu}}{93.105~eV~\\Omega_ch^2}~, \\end{equation} where $\\Sigma m_{\\nu}$ is the sum of the neutrino masses. In this study, the mocked data we use are generated by assuming the massless neutrino {\\it i.e.} $f_{\\nu} = 0$ in the fiducial models. Consequently the constraints on $f_{\\nu}$ should be regarded as the upper limits of the neutrino mass which the future observations will be sensitive to. It is well known that the massive neutrinos modify the shape and amplitude of the matter power spectrum, and also the epoch of matter-radiation equality, angular diameter distance to the last scattering surface. Thus they leave imprints on the observations of CMB and LSS. In Table II, we provide the constraints on neutrino mass from the current observations and the future simulated data. For the current data\\footnote{Usually Lyman-$\\alpha$ data will give stringent constraint on the neutrino mass, however, its systematics are quite unclear currently. In our global analysis to be conservative, we have not included it \\cite{silk,komastue}.}, within the framework of the $\\Lambda$CDM model, we get $\\Sigma m_{\\nu}<0.958~eV~(95\\%)$ which is consistent with the result in Ref.\\cite{Tegmark:2006az}. For the time evolving EoS of dark energy model, this limit is relaxed to $\\Sigma m_{\\nu}<1.59~eV~(95\\%)$, due to the degeneracy between the dark energy parameters and the neutrino mass\\cite{Xia:2006wd,Hannestad:2005gj}.  \\begin{table}% TABLE II. Constraints on neutrino mass from the current observations and the future simulations. We have shown the $2\\sigma$ upper limits. In order to highlight the contribution from LAMOST, we compare the results with/without LAMOST. \\begin{center}  \\begin{tabular}{|c|c|c|c|c|}  \\hline  \\multicolumn{2}{|c|}{Current}&\\multicolumn{3}{c|}{Future}\\\\  \\cline{3-5}  \\multicolumn{2}{|c|}{}&&PLANCK+SNLS&PLANCK+SNLS+LAMOST\\\\  \\hline  $\\Lambda$CDM&$<$0.958eV&FMI&$<$0.957eV&$<$0.377eV\\\\  Run w&$<$1.59eV&FMII&$<$0.915eV&$<$0.346eV \\\\ \\hline      \\hline   \\end{tabular} \\end{center} \\end{table}    With the simulated data, in figure \\ref{fig4}, we illustrate the one dimensional probability distribution of the total neutrino mass $\\sum m_{\\nu}$. The black solid line is given by the current constraints, the red solid line is given by fitting with the simulated PLANCK $+$ SNLS with fiducial model II and the red dashed line is given by including the simulated LAMOST. The blue solid line and blue dashed line are the results obtained with the fiducial model I. Our results show that the LAMOST can provide a more stringent constraint on the neutrino mass. For example, the $2\\sigma$ neutrino mass limit is changed from $0.957 eV$ to $0.377 eV$ by including the simulated LAMOST data with the fiducial model I.    \\begin{figure}[htbp] \\begin{center} \\includegraphics[scale=0.9]{LAMOST_A2A3_tri9.eps} \\vskip-1.3cm \\vspace{10mm}\\caption{one-dimensional distributions and two-dimensional $68\\%$ and $95\\%$ limits on the cosmological parameters. The black solid lines are obtained with the simulated PLANCK $+$ SN Ia and the red solid lines are from PLANCK $+$ SN Ia $+$ LAMOST.\\label{fig1}} \\end{center} \\end{figure}   \\begin{figure}[htbp] \\vspace{3mm} \\begin{center} \\includegraphics[scale=0.5]{DE_2D2.eps} \\vskip-1.3cm \\vspace{10mm}\\caption{2-d joint $68\\%$ and $95\\%$ confidence regions of the parameters $w_0$ and $w_a$ for flat universe. The black solid line is given by the current constraints, the red solid line comes from the simulated data of PLANCK and 5year SNLS with fiducial model II, while the red dashed line is by combining the simulated LAMOST data. The blue solid line and blue dashed line are the results with the fiducial model I with/without LAMOST respectively.\\label{fig2}} \\end{center} \\end{figure}   \\begin{figure}[htbp] \\begin{center} \\includegraphics[scale=0.45]{Neu.eps} \\vskip-1.3cm \\vspace{10mm}\\caption{one-dimensional constraints on the neutrino mass. The black solid line is given by the current data, the red solid line is given by fitting with the simulated PLANCK $+$ SNLS with fiducial model II and the red dashed line is given by including the simulated LAMOST. The blue solid line and blue dashed line are the results obtained with the fiducial model I. \\label{fig4}} \\end{center} \\end{figure}"
    },
    "0806/0806.1245_arXiv.txt": {
        "abstract": "Ekpyrotic and cyclic cosmologies provide theories of the very early and of the very late universe. In these models, the big bang is described as a collision of branes - and thus the big bang is not the beginning of time. Before the big bang, there is an ekpyrotic phase with equation of state $w=\\frac{\\cP}{\\rho}\\gg 1$ (where $\\cP$ is the average pressure and $\\rho$ the average energy density) during which the universe slowly contracts. This phase resolves the standard cosmological puzzles and generates a nearly scale-invariant spectrum of cosmological perturbations containing a significant non-gaussian component. At the same time it produces small-amplitude gravitational waves with a blue spectrum. The dark energy dominating the present-day cosmological evolution is reinterpreted as a small attractive force between our brane and a parallel one. This force eventually induces a new ekpyrotic phase and a new brane collision, leading to the idea of a cyclic universe. This review discusses the detailed properties of these models, their embedding in M-theory and their viability, with an emphasis on open issues and observational signatures. ",
        "introduction": "\\label{section intro}  The evolution of our universe is very well understood from the time of big bang nucleosynthesis until the present. However, if we go beyond these time frontiers, almost nothing is known with a great amount of certainty. We know that the early universe must have been in a very special state (very homogeneous and flat, but with tiny fluctuations in curvature), but we do not know why. The theory of inflation \\cite{Guth:1980zm,Linde:1981mu,Albrecht:1982wi} provides a possible explanation by means of a period of rapid expansion preceding nucleosynthesis, but it has not yet been possible to test inflation conclusively, which is why it is important to keep an open mind. Also, we know that dark energy has come to dominate the energy density of the universe \\cite{Riess:1998cb,Perlmutter:1998np} and that it will determine the expansion of the universe in the near future. If dark energy turns out to be a cosmological constant, which is the theoretically simplest way of modeling it, then the universe will continue to expand at an accelerating pace and become increasingly dilute and cold, making life very difficult, if not impossible. However, it should be borne in mind that we have no fully convincing explanation of dark energy at present.  This review will be concerned with a set of ideas, strongly inspired by string theory, that suggests alternative solutions to the early universe puzzles mentioned above, and considers an alternative fate for the future of our universe. What makes these ideas exciting is that, on the one hand, they provide a theoretical playground for applying string theory to cosmology, and, on the other hand, certain observational signatures predicted from these models are in a range that will be tested by near-future satellite experiments. This provides a certain timeliness to a subject concerned with eras far removed from everyday experience.  \\begin{figure}[t] \\begin{center} \\includegraphics{braneworld.eps} \\caption{\\label{figure braneworld} {\\small The braneworld picture of our universe. Think of a sandwich: the filling is the 5-dimensinonal bulk spacetime, which is bounded by the two pieces of bread a.k.a. the 4-dimensional boundary branes. There is no space ``outside'' of the sandwich, but the branes can be infinite in all directions perpendicular to the line segment. In the M-theory embedding, there are 6 additional internal dimensions at each 5-dimensional spacetime point.}} \\end{center} \\end{figure}   Ekpyrotic\\footnote{The name ekpyrosis can be translated as all-engulfing cosmic fire. In Stoic philosophy, it represents the contractive phase of eternally-recurring destruction and re-creation \\cite{Wiki}.} and cyclic cosmology are based on the braneworld picture of the universe, in which spacetime is effectively 5-dimensional, but with one dimension not extending indefinitely, but being a line segment\\footnote{This setting will be discussed in detail in section \\ref{section fundamental}, along with its motivation from and embedding in heterotic M-theory.}, see Fig. \\ref{figure braneworld}. The endpoints of this line segment (orbifold) are two $(3+1)$-dimensional boundary branes. All matter and forces, except for gravity, are localized on the branes, while gravity can propagate in the whole spacetime. Our universe, as we see it, is identified with one of the boundary branes and, as long as the branes are far apart, can interact with the other brane only via gravity. The ekpyrotic model assumes that there is an attractive force between the two branes, which causes the branes to approach each other. This ekpyrotic phase has the rather non-intuitive property that it flattens the branes to a very high degree. Eventually the two branes collide and move through each other (since there is no space ``outside'' of the boundary branes, it makes no difference whether we say that the branes bounce off each other or move through each other). It is this collision that, from the point of view of someone living on one of the branes, looks like the big bang. The collision is slightly inelastic and produces matter and radiation on the branes, where the standard cosmological evolution now takes place. However, due to quantum fluctuations, the branes are slightly rippled and do not collide everywhere at exactly the same time. In some places, the branes collide slightly earlier, which means that the universe has a little bit more time to expand and cool. In other places, the collision takes place slightly later, and those regions remain a little hotter. This provides a heuristic picture of the way temperature fluctuations are naturally produced within the model. Shortly after the brane collision, the distance between the boundary branes gets almost stabilized, but the branes start attracting each other again very slightly. This very slight attraction acts as quintessence, and is identified with the dark energy observed in the universe. After a long time and as the branes become closer again, they start attracting each other more strongly so that we get a new ekpyrotic phase and eventually a new brane collision with the creation of new matter. In this way, a cyclic model of the universe emerges.  The preceding description of course only provided a very rough outline of the main ideas. All of this will be made precise in the following. However, at this stage it is already clear that conceptually the cyclic universe differs substantially from the standard big bang picture. For one, the big bang is seen as a physical event and not a mysterious moment of creation. As such, it does not represent the beginning of time (note that if quantum gravity is unitary, this point of view seems unavoidable). Thus there was plenty of time before the big bang for the universe to be in causal contact over large regions, and in this way the horizon problem is automatically solved. The long timescales existing before the big bang also imply that other cosmological puzzles don't have to be solved in the extremely short time interval between the big bang and nucleosynthesis. What's more, properties of the very early universe (such as cosmological perturbations) and properties of the late universe (such as dark energy) are described by the same ingredients, namely the motion of branes.    In the description above, the higher-dimensional interpretation was emphasized. This higher-dimensional viewpoint is useful for understanding the origin of the ideas, and provides a concrete realization of them. However, it is important to realize that many of the processes mentioned above can also be discussed purely in a 4-dimensional effective theory. Some processes, such as the ekpyrotic period of slow contraction before the big bang, are quite general and could have a different origin than the brane motion just described. Therefore, the four subsequent sections of this review, which describe the fundamentals of ekpyrotic and cyclic cosmology, are phrased mostly in terms of a 4-dimensional effective theory. Section \\ref{section ekpyrotic} describes the ekpyrotic phase and the way it resolves the standard cosmological puzzles. Section \\ref{section crunch} describes the approach to the big crunch in more detail. The cosmological fluctuations generated by the model are discussed in section \\ref{section perturbations}. This includes a treatment of scalar perturbations with their second-order non-gaussian corrections as well as tensor perturbations. The following section is concerned with the details of the cyclic universe, such as its relationship to the second law of thermodynamics and its global structure. Section \\ref{section fundamental} explains the origin of the 4-dimensional effective theory used in the preceding sections and its relation to string and M-theory. Finally, section \\ref{section outlook} provides a brief conclusion and an outlook to possible future developments.   It should be noted from the outset that ekpyrotic and cyclic cosmology are theories that are still in full development. Certain aspects of the theories have gained or lost in importance over time. Here are two examples: in the original ekpyrotic proposal, the brane collision was between a boundary and a bulk brane, the second boundary brane being a spectator. It was later realized that the model could be simplified by considering a collision of the boundary branes. It was also thought that the dark energy phase was crucial for smoothing and isotropizing the universe. Again, it was later realized that the ekpyrotic phase can achieve this on its own. In this review the subject will be discussed only according to its present understanding. Moreover, every effort will be made to highlight assumptions that are in need of a precise calculation, or that are due for example to limitations in our understanding of string theory. This should hopefully provide the reader with opportunities to contribute to the subject. ",
        "conclusions": "\\label{section outlook}  Postulating a slowly contracting phase of the universe with equation of state $w \\gg 1$ (preceding the standard expanding phase of big bang cosmology) has extremely powerful consequences: it renders the universe flat and homogeneous over large patches, it delays the onset of chaos until the quantum gravity regime is reached, it prepares the universe ideally for a big crunch/big bang transition, and, on top of that, it generates small-scale scalar perturbations with a nearly scale-invariant spectrum. Hence, even if the higher-dimensional picture and the link to string theory seem unappealing to you, at the very least the ekpyrotic phase represents a new mechanism for solving the standard cosmological puzzles and a new way of generating cosmological perturbations. Moreover, these perturbations contain a significant non-gaussian component, so that they can be tested observationally and compared to simple, gaussian, inflationary models. Further into the future, the prediction of a blue small-amplitude gravitational spectrum will also be tested against observations.  But more generally, ekpyrotic and cyclic models present a very broad conceptual framework, in which the big bang is not the beginning of time, but rather a physical event amenable to a testable physical description. Cyclic models allow for evolutionary timescales much larger than the time since the last big bang, and can provide qualitatively new explanations in this way. The embedding of ekpyrotic/cyclic models in string theory provides a geometric interpretation for many cosmological processes, and can eventually provide observational tests of the theory. Many open issues remain, which should be tractable within string theory: in particular, the issue of moduli stabilization requires further work, as does the closely related question of ``initial'' conditions in the presence of many moduli. Also, the all-important cyclic potential has not yet been derived from first principles. Most promising is the recent realization that the brane collision, and the corresponding 4-dimensional big crunch/big bang transition, can perhaps be understood within the context of generalizations of the non-perturbative gauge/gravity correspondence. Much work remains to be done here before a physically realistic case can be treated; but this is a topic of such importance to cosmology that we should see this as an opportunity, rather than a drawback. It would be truly magnificent if we could find observational signatures of higher-dimensional physics, and of a time before the big bang."
    },
    "0806/0806.2841_arXiv.txt": {
        "abstract": "Gravitational lensing provides a unique and powerful probe of the mass distributions of distant galaxies.  Four-image lens systems with fold and cusp configurations have two or three bright images near a critical point. Within the framework of singularity theory, we derive analytic relations that are satisfied for a light source that lies a small but finite distance from the astroid caustic of a four-image lens. Using a perturbative expansion of the image positions, we show that the time delay between the close pair of images in a fold lens scales with the cube of the image separation, with a constant of proportionality that depends on a particular third derivative of the lens potential. We also apply our formalism to cusp lenses, where we develop perturbative expressions for the image positions, magnifications and time delays of the images in a cusp triplet. Some of these results were derived previously for a source asymptotically close to a cusp point, but using a simplified form of the lens equation whose validity may be in doubt for sources that lie at astrophysically relevant distances from the caustic. Along with the work of \\citet{Keeton-fold}, this paper demonstrates that perturbation theory plays an important role in theoretical lensing studies. ",
        "introduction": "\\label{sec:intro}  Gravitational lensing, or the bending of light by gravity, offers an exciting synergy between mathematics and astrophysics. Singularity theory provides a powerful way to describe lensing near critical points \\citep{Schneider_lensing,Petters_book}, which turns out to have important implications for astrophysics and the quest to understand dark matter. Four-image lensed quasars can be broadly classified, according to image geometry, as either folds, cusps or crosses.  We are interested in folds and cusps, which occur when the light source is close to the caustic curve, along which the lensing magnification is infinite (see top row of Fig.~\\ref{fig:u-theta-def}). Fold lenses feature a close pair of bright images, while cusp lenses have a close triplet of bright images (see bottom row of Fig.~\\ref{fig:u-theta-def}). By expanding the gravitational potential of the lens galaxy in a Taylor series, one finds that the image magnifications satisfy simple analytic relations, viz. \\begin{equation} |\\mu_A| - |\\mu_B| \\approx 0 \\quad {\\rm and} \\quad |\\mu_A| - |\\mu_B| + |\\mu_C| \\approx 0 \\end{equation} for fold pairs and cusp triplets, respectively \\citep{Blandford_Narayan, Mao_cusp, Schneider_cusp, Schneider_lensing, zakharov-magsum, Gaudi-fold, Gaudi-cusp, Keeton-cusp, Keeton-fold}. For a fold pair, the two images have opposite parity, hence the negative sign. For a cusp triplet, the middle image (B) has opposite parity from the outer images (A and C). Note that in practice, one works with the image fluxes, which are directly observable, rather than the magnifications, which are not. This leads to the equivalent relations: \\begin{equation} \\label{eq:fluxratdef} \\Rfold \\equiv \\frac{F_A - F_B}{F_A + F_B} \\approx 0 \\quad {\\rm and} \\quad \\Rcusp \\equiv \\frac{F_A - F_B + F_C}{F_A + F_B + F_C} \\approx 0 , \\end{equation} where the image flux $F_i$ is related to the source flux $F_0$ by $F_i = |\\mu_i| F_0$.  Observationally, the fold and cusp relations are violated in several lens systems \\citep[e.g.,][]{Hogg_1422, Falco_0414, Keeton-shear-ellip, Keeton-cusp, Keeton-fold}.  These so-called ``flux-ratio anomalies'' are taken as strong evidence that lens galaxies contain significant small-scale structure \\citep[e.g.,][]{Mao-flux, Metcalf-substruc, Chiba_substruc}.  Since global modifications of the lens potential (modeled with multipole terms, for example) cannot explain the anomalous flux ratios \\citep{Evans-multipole, Congdon-multipole, Yoo-1115, Yoo-rings}, it is generally believed that lens galaxies must contain substructure in the form of mass clumps on a scale $\\ga 10^6\\,M_\\odot$ \\citep{Dobler-finite}.  Indeed, substructure is observed in at least two lens systems, in the form of a dwarf galaxy near the main lens galaxy \\citep{Schechter-0414, McKean-2045}.  The substructure could well be invisible, though; numerical simulations of structure formation in the cold dark matter (CDM) paradigm predict that galaxy dark matter halos are filled with hundreds of dark subhalos \\citep[e.g.,][]{Klypin-missing, Moore-substructure}.  The abundance of flux-ratio anomalies in lensed radio sources implies that lens galaxies contain $\\sim$2\\% (0.6--7\\% at 90\\% confidence) of their mass in substructure \\citep{Dalal-substruc}.  This may be somewhat higher than the amount of CDM substructure \\citep[e.g.,][]{Mao-substructure}, although new higher-resolution simulations are attempting to refine the theoretical predictions and make more direct comparisons with lensing observations \\citep[e.g.,][]{Gao-substructure, Strigari-VL, Diemand-vialactea}. While there is still work to be done along these lines, it is clear that violations of the ideal fold and cusp relations provide important constraints on the small-scale distribution of matter in distant lens galaxies.  \\citet{Keeton-cusp, Keeton-fold} pointed out that the ideal fold and cusp relations only hold for a source asymptotically close to the caustic.  If we want to use flux-ratio anomalies to study small-scale structure in lens galaxies, it is vital that we understand how much $\\Rfold$ and $\\Rcusp$ can deviate from zero for realistic smooth lenses just because the source lies a small but finite distance from the caustic.  To that end, \\citet{Keeton-cusp} used a Taylor series approach to demonstrate that $\\Rcusp = 0$ to lowest order, and used Monte Carlo simulations to suggest that $\\Rcusp \\propto d^2$, where $d$ is the distance between the two most widely separated cusp images.  In order to derive the leading non-vanishing term in $\\Rcusp$ analytically, it would be necessary to extend the Taylor series to a higher order of approximation. This has not been done before, since including higher-order terms substantially complicates the analysis. Instead, previous authors have made assumptions about which terms are important and which terms can be neglected, in order to obtain an analytically manageable problem.  As we shall see, perturbation theory provides a natural way to overcome the difficulties of the Taylor series approach. \\citet{Keeton-fold} used perturbation theory to show that for fold lenses, $\\Rfold$ is proportional to the image separation $d_1$ of the fold pair.  We extend their analysis in several important ways. We derive the leading-order non-vanishing expression for $\\Rcusp$ using perturbation theory. We also show, for the first time, how the combination of singularity theory and perturbation theory can be used to study {\\em time delays} between lensed images for both fold and cusp systems.  Working in analogy to flux ratios, we derive time-delay relations for fold and cusp lenses, which we will apply (in a forthcoming paper) to observed lenses in order to identify ``time-delay anomalies.''  As time delays join flux ratios in probing small-scale structure in lens galaxies \\citep[see][]{Keeton-tdel}, singularity theory and perturbation theory once again provide the rigorous mathematical foundation. ",
        "conclusions": "We have developed a unified, rigorous framework for studying lensing near fold and cusp critical points, which can (in principle) be extended to arbitrary order. We have found that the differential time delay of a fold pair assumes a particularly simple form, depending only on the image separation and the Taylor coefficient $h = \\psi_{222} (\\bmo)/6$. This result is astrophysically relevant, since it is quite accurate even for sources that are not asymptotically close to the caustic. We have also obtained perturbative expressions for the image positions, magnifications and time delays of a cusp triplet. These results rest on the key insight that a source at a given distance $\\epsilon$ from a cusp along the relevant symmetry axis of the caustic can only move a perpendicular distance of $\\epsilon^{3/2}$ in order to remain inside the caustic \\citep{Blandford_Narayan}. We have shown rigorously that the distance dependence of the magnification ratio $\\Rcusp$ conjectured by \\citet{Keeton-cusp} is correct. We have also demonstrated that the leading-order expression for the image positions is given by the relations presented by \\citet{Keeton-cusp}, and have provided the necessary framework for deriving the image positions corresponding to a Taylor expansion of the lens potential at arbitrary truncation order. Finally, we have derived cusp time delays analytically for the first time. Our results provide a rigorous foundation for identifying anomalous flux ratios and time delays in gravitational lens systems, and for using them to study small-scale structure in the mass distributions of distant galaxies."
    },
    "0806/0806.0850_arXiv.txt": {
        "abstract": "Galaxy clusters exhibit regular scaling relations among their bulk properties.  These relations establish vital links between halo mass and cluster observables.  Precision cosmology studies that depend on these links benefit from a better understanding of scatter in the mass-observable scaling relations.  Here we study the role of merger processes in introducing scatter into the \\mtx\\ relation, using a sample of 121 galaxy clusters simulated with radiative cooling and supernova feedback, along with three statistics previously proposed to measure X-ray surface brightness substructure.  These are the centroid variation ($w$), the axial ratio ($\\eta$), and the power ratios ($P_{20}$ and $P_{30}$).  We find that in this set of simulated clusters, each substructure measure is correlated with a cluster's departures $\\delta \\ln \\tx$ and $\\delta \\ln M$ from the mean \\mtx\\ relation, both for emission-weighted temperatures $\\tew$ and for spectroscopic-like temperatures $\\tsl$, in the sense that clusters with more substructure tend to be cooler at a given halo mass.  In all cases, a three-parameter fit to the \\mtx\\ relation that includes substructure information has less scatter than a two-parameter fit to the basic \\mtx\\ relation. ",
        "introduction": "Clusters of galaxies play a critical role in our understanding of the Universe and its history and are potentially powerful tools for conducting precision cosmology.  For example, large cluster surveys can discriminate between cosmological models with different dark-energy equations of state by providing complementary observations of the shape of the cluster mass function, evolution in the number density of clusters with redshift, and bias in the spatial distribution of clusters \\citep{1998ApJ...508..483W, 2002ApJ...577..569L, 2003PhRvD..67h1304H,2004ApJ...613...41M,Voit:2005PhysRevD}. However, this potential to put tight constraints on the properties of dark energy will be realized only if we can accurately measure the masses of clusters and can precisely characterize the scatter in our mass measurements.  Scatter in X-ray cluster properties is directly related to substructure in the intracluster medium.  If clusters were all structurally similar, then there would be a one-to-one relationship between halo mass and any given observable property.  Generally speaking, deviations from a mean mass-observable relationship are attributed to structural differences among clusters.  One kind of structural difference is the presence or absence of a cool core, in which the central cooling time is less than the Hubble time at the cluster's redshift, and the prominence of a cool core is observed to be a source of scatter in scaling relationships \\citep{1994MNRAS.267..779F,1998ApJ...504...27M,2002ApJ...576..601V,2004ApJ...613..811M}. We also expect structural differences to arise from substructure in the dark matter, galaxy, and gas distributions.  For instance, there may be a spread in halo concentration at a given mass, variations in the incidence of gas clumps, differences in the level of AGN feedback, or various effects due to mergers.  All of these deviations can be considered forms of substructure that produce scatter in the mass-observable relations one would like to use for cosmological purposes.  While it may ultimately be possible to constrain the amount of scatter and its evolution with redshift using self-calibration techniques \\citep{2005PhRvD..72d3006L}, such constraints would be improved by prior knowledge about the relationship between scatter and substructure.  Traditionally, the most worrisome form of substructure has been that due to the effects of merger events.  Clusters are often identified as ``relaxed'' or ``unrelaxed'', with the former assumed to be nearly in hydrostatic equilibrium and the latter suspected of being far from equilibrium.  Cosmological simulations of clusters indicate that the truth is somewhere in between.  The cluster population as a whole appears to follow well-defined virial relations with log-normal scatter around the mean, showing that clusters do not cleanly separate into relaxed and unrelaxed systems (e.g.,\\cite{2007astro.ph..2241E}). Even the most relaxed-looking clusters are not quite in hydrostatic equilibrium (e.g.,\\cite{2006ApJ...650..128K}).  Instead of simply being ``relaxed'' or ``unrelaxed,'' clusters occupy a continuum of relaxation levels determined by their recent mass-accretion history.  Quantifying this continuum of relaxation offers opportunities for reducing scatter in the mass-observable relations.  If mergers are indeed responsible for much of the observed scatter around a given scaling relation, then there may be correlations between a cluster's morphology and its degree of deviation from the mean relation.  Once one identifies a morphological parameter that correlates with the degree of deviation, one can construct a new mass-observable relation with less scatter by including the morphological parameter in the relation.  Such an approach would be analogous to the improvement of Type Ia supernovae as distance indicators by using light-curve stretch as a second parameter to indicate the supernova's luminosity \\citep{1993ApJ...413L.105P,1996ApJ...473...88R}.  Here we investigate how merger-related substructure in simulated clusters affects the relationship between a simulated cluster's mass and the temperature of its intracluster medium, building upon \\cite{1995ApJ...452..522B} and \\cite{2006ApJ...639...64O}. \\cite{1995ApJ...452..522B} quantified the morphologies and dynamical states of observed clusters and found structure measures to be an indicator of the dynamical state of a cluster. \\cite{2006ApJ...639...64O} also examined morphological measurements, for both observed and simulated clusters, and found that simulations without cooling showed no correlation between substructure and scaling relation scatter.  In this work we examine substructure for simulated clusters with radiative cooling and focus on the idea that merger processes introduce intrinsic scatter into the \\mtx\\ relationship by displacing clusters in the \\mtx\\ plane away from the mean X-ray temperature $\\langle{\\tx}\\rangle|_{_M}$ at a given mass $M$, either to higher or lower average ICM temperature.  We then adopt a set of statistics \\citep{1995ApJ...452..522B,2006ApJ...639...64O} for quantifying galaxy cluster substructure and merger activity in order to investigate this hypothesis.  Section 2 discusses the \\mtx\\ scaling relationship in our sample of simulated clusters and shows that disrupted-looking clusters in this sample tend to be cooler at a given cluster mass.  In Section 3 we attempt to quantify the relationship between morphology and temperature using four different substructure statistics and compare it to similar studies.  We then show that substructure in these simulated clusters indeed correlates with scatter in the \\mtx\\ relationship and assess the prospects for using that correlation to reduce scatter in the \\mtx\\ plane.  Section 4 summarizes our results. ",
        "conclusions": "Using a sample of galaxy clusters simulated with cooling and feedback, we investigated three substructure statistics and their correlations with temperature and mass offsets from mean scaling relations in the \\mtx\\ plane.  First, we showed that the substructure statistics $w$, $\\eta$, $P_{20}$ and $P_{30}$ all correlate significantly with $\\delta \\ln \\tx$, though with non-negligible scatter.  In all cases this scatter is larger for $\\delta \\ln$ $\\tsl$ than it is for $\\delta \\ln$ $\\tew$.  Next, we considered the possibility that \\mtx\\ scatter is driven by low-mass clusters.  We tested the degree to which scatter can be reduced by filtering out these systems.  This consisted of performing a cut at 2 keV, for which we saw that it yielded a modest improvement in mass estimates.  To see whether incorporating substructure could refine these mass estimates, we first showed that $w$, $\\eta$, $P_{20}$, and $P_{30}$ correlate significantly with the difference $\\delta \\ln M$ between masses predicted from the mean $M(\\tx)$ relation and the true cluster masses, with non-negligible scatter that again is less for $M(\\tew)$ than it is for $M(\\tsl)$. Then we adopted a full three-parameter model, $M$-$\\tx$-$S$, which includes substructure information $S$ estimated using $w$, $\\eta$, $P_{20}$, and $P_{30}$. Scatter about the basic two-parameter $M$-$\\tew$ relation was 0.094.  Including substructure as a third parameter reduced the scatter to 0.072 for centroid variation, 0.084 for axial ratio, 0.081 for $P_{20}$, and 0.084 for $P_{30}$.  Scatter about the basic two-parameter $M$-$\\tsl$ relation was 0.124, and including substructure as a third parameter reduced the scatter to 0.085 for centroid variation, 0.112 for axial ratio, 0.110 for $P_{20}$, and 0.108 for $P_{30}$.  As one last test, and to increase our confidence that our substructure measures are not relying on potentially non-physical core structure in the simulations, we also repeated the comparison of mass-estimates for $\\tew$, with the core regions of the clusters excised.  First, removing the core slightly increased the scatter in \\mtx\\, possibly by making the average temperature more sensitive to structure outside the core.  Second, even with the cores removed the improvement in mass-estimates obtained using substructure information remains.  Based on these results, it appears that centroid variation is the best substructure statistic to use when including a substructure correction in the $M$-$\\tew$ relation.  However, the correlations we have found in this sample of simulated clusters might not hold in samples of real clusters, because relaxed clusters in the real universe tend to have cooler cores than our simulated clusters do."
    },
    "0806/0806.0916_arXiv.txt": {
        "abstract": "The third law of black hole thermodynamics is violated in the test fluid approximation for the process of phantom energy accretion onto a rotating or an electrically charge black hole. The black hole mass is continuously decreasing but the angular momentum or electric charge is remaining constant in this process. As a result, a black hole reaches the extreme state during a finite time with a threat of black hole transformation into the naked singularity and violation of the cosmic censorship conjecture. We demonstrate this by using new analytical solutions for spherically symmetric stationary distribution of a test perfect fluid with an arbitrary equation of state in the Reissner-Nordstr\\\"om metric. Our speculative assumption, however, is that the cosmic cosmic censorship conjecture remains valid even for phantom energy case if one takes into account the back reaction of an accreting fluid onto a near extreme black hole. Some hint for the validity of this hypothesis comes from the specific case of the ultra-hard fluid accretion onto the rotating black hole. In this case the energy density of an accreting fluid diverges at the event horizon of an extreme black hole, thus violating the test fluid approximation. ",
        "introduction": "\\label{intro}  The problem of matter accretion onto the compact objects in the Newtonian gravity was formulated in a self-similar manner by Bondi \\cite{Bondi}. In the framework of General Relativity a steady-state accretion of test gas onto a Schwarzschild black hole was investigated by Michel \\cite{Michel}. The detailed studies of spherically symmetric accretion of different types of fluids onto black holes were further undertaken in a number of works, see e.~g. \\cite{accretion}. In \\cite{bde04} (see also \\cite{accrph} for further discussion) it was shown that accretion of a {\\sl phantom} fluid onto the Schwarzschild black hole results in a diminishing of black hole mass due to a negative flux of energy through the event horizon. Usually in the General Relativity it is assumed that matter has a {\\sl suitable} form. By {\\sl suitable}, one can imply that the stress-energy tensor of matter satisfy particular energy conditions, i.e. the weak energy condition (see for details, e.g. \\cite{Wald97}). Meanwhile the phantom energy violates by definition the weak energy condition $\\rho+p<0$, where $\\rho$ is an energy density and $p$ is a pressure. It is believed that phantoms generically contain ghosts \\cite{ghosts} and thus should be denied as nonphysical. It turns out, however, that it is possible to construct a physically reasonable model of phantom, which is stable in the ultraviolet limit, thus giving no catastrophic instabilities of vacuum \\cite{Rubakov}. Thus the study of phantoms seems to be not fully meaningless from the physical point of view. Violation of weak energy condition by phantom brings unusual consequences: in the cosmological context the Big Rip singularity can be formed \\cite{bigrip}; while the accreting of phantom leads to the diminishing of a black hole mass \\cite{bde04,accrph}.  In this paper, we find a solution for the steady-state accretion of a test perfect fluid with an arbitrary equation of state, $p=p(\\rho)$, onto the charged black hole, when the event horizon for the Reissner-Nordstr\\\"om metric exists, $m^2>e^2$. A similar analysis was performed in \\cite{Mubasher08}. We show that a phantom energy accretion leads to the decreasing of a black hole mass. This result is consistent with \\cite{bde04} for Schwarzschild case. On the other hand, we show that when the Reissner-Nordstr\\\"om metric describes a naked singularity with $m^2<e^2$, a perfect fluid does not accrete at all onto the the naked singularity, instead, a static atmosphere is formed.  As we show below, when neglecting the gravitational back reaction of the accreting fluid on the background metric, the extreme state of electrically charged black hole is reached by phantom energy accretion during the finite time. It is natural to ask then, whether it is possible to convert a Reissner-Nordstr\\\"om black hole into a naked singularity by accretion of phantom fluid. When phantom accretes onto a Schwarzschild black hole, the latter becomes smaller and smaller with time and might completely disappear, i.e. in the Big Rip scenario. However, in the case of phantom energy accretion onto a charged black hole, when the mass becomes smaller than the charge, $m^2<e^2$, we might naively think that black hole transforms to a naked singularity by phantom energy accretion. This would also imply that the third law of black hole thermodynamics breaks down as well \\cite{bch73}. Such a process of a Kerr-Newman black hole transformation into a naked singularity by accretion of phantom energy was first discussed in \\cite{MadridGonz}.  The key conjecture of general relativity is the cosmic censorship by R. Penrose \\cite{Penrose69} prohibiting the appearance of naked singularity in the gravitational collapse of suitable matter. At the same time, the inevitable formation of singularity inside a black hole horizon is guaranteed by singularity theorems \\cite{HawPenr70,hawking73}. It is worthwhile to note that by ``usual means'' it seems to be impossible to make a naked singularity from a black hole. For example, the Kerr-Newman black hole cannot be transformed into a naked singularity by capturing test particles with an electrical charge or orbital angular momentum \\cite{bch73,Wald74} and can be only approached to the extreme state in a limiting process.  We argue, however, that a test fluid approximation is inevitably violated when the Reissner-Nordstr\\\"om black hole or naked singularity is near to the extreme state. If this true, the back reaction of the accreting fluid on the background geometry may prevent black transformation into the naked singularity.  The paper is organized as follows. In Sec.~{\\ref{Sec accretion}} we construct a general formalism for the steady-state spherically symmetric accretion of a test perfect fluid onto the Reissner-Nordstr\\\"om black hole. In Sec.~\\ref{Sec solutions} we apply these results to the the particular examples of perfect fluids, the linear equation of state and the Chaplygin gas. The formation of fluid atmosphere around naked singularities is described in Sec.~\\ref{Sec atmosphere}. The approaching to the extremal case by accretion of phantom and a possible violation of the third law of thermodynamics is studied in Sec.~\\ref{Sec approaching}. We conclude and discuss the obtained results of the paper in Sec.~\\ref{Sec discussion}. ",
        "conclusions": "% \\label{Sec discussion} In this paper we considered the stationary distribution of the test perfect fluid with an arbitrary equation of state, $p=p(\\rho)$, in the Reissner-Nordstr\\\"om metric. Similarly to the well-known case of the stationary accretion of perfect fluid onto the Schwarzschild black hole, the corresponding solution for the accretion exists also in the case of Reissner-Nordstr\\\"om black hole. On the contrary, there is no stationary accretion of the perfect fluid onto the Reissner-Nordstr\\\"om naked singularity, $e>m$. Instead, the static atmosphere of the fluid is formed around the naked singularity. In both cases of the black hole and the naked singularity we found the analytical solution to the problem of the steady state configurations of the perfect fluids with an arbitrary equation of state, $p=p(\\rho)$. As the particular cases, we considered fluid with the linear equation of state, $p=\\alpha(\\rho-\\rho_0)$ and the Chaplygin gas, $p=\\alpha/\\rho$.  When the accreting fluid is phantom, $\\rho+p<0$, the mass of the Reissner-Nordstr\\\"om black hole decreases. This result is in the agreement with the previous findings \\cite{bde04,accrph}. This immediately leads us to the question, whether it is possible to transform the Reissner-Nordstr\\\"om black hole into the naked singularity by by accretion of phantom. Formally it seems so, since the accreting phantom decreases the black hole mass, while the electric charge of the black hole remains the same. Thus, one can expect that at some finite moment of time a black hole will turn into the naked singularity. Indeed, as we have shown in Sec.~\\ref{Sec approaching}, it takes the finite time for the Reissner-Nordstr\\\"om black hole to reach the extreme case. The similar result also holds for the Kerr black hole.  However, this naive picture, taken out in the test fluid approximation, does not seem realized if one takes into account the back reaction of perfect fluid onto the background metric. First of all, in the case of ultra-hard equation of state, $p=\\rho$, the fluid density diverges at the horizon, $r=r_+$, when black hole is approaching to the extremal state, $m\\to e$. Thus, the test fluid approximation breaks down and the results is not applicable. In other cases, when $c_s^2\\neq 1$, the situation is more subtle and beyond the scope of this paper. Our preliminary investigations show, however, that in general case, the test fluid approximation breaks down at the extreme case. We expect that black hole cannot be converted into the naked singularity even by the accretion of phantom energy, and thus the third law of thermodynamics is not violated in this case.  We would like to stress, that although the test fluid approximation seems to break down for the near-extreme state of the black hole/naked singularity, for the far-from-the-extreme state black hole (in particular, for the Schwarzschild solution), the parameters of the perfect fluid and the boundary condition at the infinity can be tuned so, that the test fluid approximation describes well the problem under consideration. In particular, we stress, that the phantom energy accretion indeed leads to decreasing of the black hole mass.  Recently Gao {\\sl et al} \\cite{Gao} contested the validity of this result, claiming that the inclusion of the back reaction would result in the opposite process, namely, the mass of a black hole increases in the process of phantom accretion. The particular solution was presented in \\cite{Gao} to support this point. However, the conclusion of \\cite{Gao} is doubtful. First of all, it is only valid for {\\sl imperfect} fluids, thus making impossible the application of their arguments to the {\\sl perfect} fluids, for which the effect of black hole mass decreasing was found. The perfect fluid does not accrete in the solution of Gao {\\sl et al}, and the heat flux is introduced by hand to support this kind of configuration. Note that the heat flux ingredient is crucial to cook up this solution. Without the heat flux the system must evolve to the more physical configuration, which is likely to be the stationary accretion \\cite{bde04}. Even when the heat flux is introduced, it is unclear whether such an artificial configuration is stable with respect to the small perturbations. Finally, it is worthwhile to note that the temperature of the fluid diverges at the black hole event horizon for the solution by Gao {\\sl et al}.  If the back reaction does not prevent the process of phantom accretion onto a charged black hole or rotating black hole, then it is a way to violate the cosmic censorship. Otherwise, the phantom energy must be totally forbidden on the more fundamental basis, as for instance, a quantum instability.  \\ack We would like to thank V.~Beskin, V.~Lukash, Ya.~Istomin, A.~Vikman and K.~Zybin for useful discussions and C.~Gao for correspondence. The work of EB was supported by the EU FP6 Marie Curie Research and Training Network ``UniverseNet'' (MRTN-CT-2006-035863). The work of other coauthors was supported in part by the Russian Foundation for Basic Research grant 06-02-16342 and the Russian Ministry of Science grant LSS 959.2008.2."
    },
    "0806/0806.3638_arXiv.txt": {
        "abstract": "{}{We are trying to probe conditions in the youngest super star clusters, those still embedded in dense obscuring clouds.}{The hydrogen recombination lines in the radio and infrared can be observed through the obscuration, as the optical and UV lines cannot, and give us the kinematics of the ionized gas.}{The line profiles of the clusters resemble superpositions of the lines of many very young ultra-compact or hyper-compact HII regions.  This can be explained if each OB star is individually embedded in dense material which it is accreting, even as it ionizes. }{We speculate on what this implies for conditions in the clusters.} ",
        "introduction": "Compact, populous star clusters are the main mode of star formation in many starburst galaxies.  These Super Star Clusters (SSCs) contain a few $10^5$ to  a few $10^7$ $M_\\odot$ in radii 3 to 10 pc. (\\citet{TB03}, \\citet{J04}, \\citet{Ts06}, \\citet{GG07},  and references therein.) SSCs  are expected to contribute greatly to the mechanical energy, mass transport and metal enrichment of the starburst. In their mature stages, when they may be measured directly, they are observed to drive strong winds that can create superbubbles and strongly influence the galactic ISM \\citep{Te07}. The evolution of these cluster winds is fairly well understood and the models agree reasonably well with what is observed \\citep{S07}).  But  the very youngest clusters, still deeply embedded in gas and dust clouds, are not so well understood. They cannot be observed in the optical or UV because of the extinction, but only at infrared, millimeter and radio wavelengths. The cluster stars are not directly observable;   only the nebula they excite can be measured.  These cluster sources are called RISN (radio-infrared supernebulae) or UDHII (ultra-dense HII regions) \\citep{TB03, KJ99}.   They contain as many as 10,000 O stars in volumes no larger than a few parsecs \\citep{TB04, TB03, Hea07, Be02}, yet do not have the cluster winds of their more evolved counterparts.  The gas motions in these clusters, as revealed in infrared and radio recombination line shapes, do not resemble O star winds.  In this note we discuss the line shapes and what they imply for the internal structure and evolution of the nebula and cluster. ",
        "conclusions": "We have demonstrated that the recombination line shapes observed in RISN are like those that would be seen from a cluster of UCHII regions. Since the line shapes of UCHII regions depend  on dense obscuring material local to the star, this implies that the cluster are still embedded in, and accreting from, their individual cocoons.  We suggest that in RISN whose lines show wider  velocity  components, some of the stars have removed their  surrounding material and developed faster winds.  Eventually, all the stars will have expelled their natal cocoons and the cluster will appear as a normal SSC."
    },
    "0806/0806.1910_arXiv.txt": {
        "abstract": "We determine an absolute calibration for the MIPS 24$\\mu $m band and recommend adjustments to the published calibrations for 2MASS, IRAC, and IRAS photometry to put them on the same scale. We show that consistent results are obtained by basing the calibration on either an average A0V star spectral energy distribution (SED), or by using the absolutely calibrated SED of the sun in comparison with solar-type stellar photometry (the solar analog method). After the rejection of a small number of stars with anomalous SEDs (or bad measurements), upper limits of $\\sim $ 1.5\\% (rms) are placed on the intrinsic infrared SED variations in both A dwarf and solar-type stars. These types of stars are therefore suitable as general-purpose standard stars in the infrared. We provide absolutely calibrated SEDs for a standard zero magnitude A star and for the sun to allow extending this work to any other infrared photometric system. They allow the recommended calibration to be applied from 1 to 25$\\mu $m with an accuracy of $\\sim$ 2\\%, and with even higher accuracy at specific wavelengths such as 2.2, 10.6, and 24$\\mu $m, near which there are direct measurements. However, we confirm earlier indications that Vega does not behave as a typical A0V star between the visible and the infrared, making it problematic as the defining star for photometric systems. The integration of measurements of the sun with those of solar-type stars also provides an accurate estimate of the solar SED from 1 through 30$\\mu$m, which we show agrees with theoretical models. ",
        "introduction": "For many applications of astronomical photometry, an accurate absolute calibration in physical units and at multiple wavelengths is critical. Previous calibrations have made use of two basic approaches (Rieke et al. 1985). In direct calibrations, measurements are made of celestial sources in ways allowing the signals to be compared directly (if sometimes through a long chain of measurements) with signals from calibrated emitters. For indirect calibrations, direct calibrations at wavelengths well removed from those of interest are interpolated or extrapolated through physical modeling of astronomical sources. For direct calibrations, rigorous error analysis is possible, although there is always a risk of systematic terms that are not captured. Error analysis is far more difficult for indirect calibrations, since the systematic errors in theoretical modeling are often not apparent and there is usually no rigorous way to quantify the errors. Therefore, indirect calibrations must be used with caution until they have been confirmed by other indirect approaches or, better, by direct ones.  Infrared calibrations have often included some aspect of indirect calibration by extrapolating from the high quality visible direct calibrations. There are now high quality direct calibrations in the infrared, so the extrapolations can be tested. In this work, we start with the infrared calibrations, test their internal consistency, and then examine the consequences for extrapolation from them into the visible. This approach is preferred in principle because stellar behavior is relatively simple in the infrared (e.g., small temperature uncertainties have little effect on the shape of a nearly Rayleigh-Jeans spectrum for stars of the temperatures considered here).  The 24$\\mu $m band of the Multiband Imaging Photometer for Spitzer (MIPS) achieves high photometric accuracy - with typical errors an order of magnitude smaller than those previously achieved at similar wavelengths (Rieke et al. 2004; Engelbracht et al. 2007). The Infrared Array Camera (IRAC) also is providing a large body of accurate and homogeneous photometry in the 3 - 8$\\mu$m range (Fazio et al. 2004; Reach et al. 2005). In addition, a new absolute calibration of unprecedented accuracy is now available in the thermal infrared (Price et al. 2004). These advances make it both desirable and feasible to establish a very accurate calibration of the MIPS photometry at this wavelength and to provide guidelines to tie it in consistently with photometry at shorter wavelengths.  Bessell (2005) has reviewed photometry in general, but with only a few comments on the infrared. Price (2004) has reviewed infrared calibrations. His review concentrates on a huge body of work by Cohen, Walker, Witteborn, Price, and other co-workers on this topic. The review states that the results are \"in substantial disagreement with previous direct calibrations,\" thus leaving open the question of possible undetected systematic errors. In addition, the review points out a number of discrepancies between the measured properties of the sun (e.g., Thuillier et al. 2003) and measurements of solar analog stars, which indicate possible issues in the photometric system.  The availability of large and homogeneous sets of data such as the Hipparcos photometric catalog (Perryman et al. 1997) and the 2MASS survey (Cutri et al. 2003) opens new possibilities to probe these issues and to improve calibrations, as well as providing a solid foundation to extend calibrations uniformly over the entire sky. Doing so is the goal of this paper. The paper is organized as follows. In Section 2, we review measurements of the absolute flux from Vega at 2.22 and 10.6$\\mu $m, and the extrapolation of these measurements to the MIPS effective wavelength of 23.675$\\mu $m. In Section 3, we discuss an independent approach to calibration using the solar analog method introduced by Harold Johnson (1965a). In Section 4, we show explicitly the discrepancy in the Vega-based visible calibration and the infrared one. In Section 5, we recommend adjustments to the absolute calibration of other photometric systems to bring them into agreement with the work reported here. The paper is summarized in Section 6. Those not wishing to plow through the details can go to that section for a summary of the useful results, which include recommendations for an absolute calibration accurate to 2\\% or better across the near and mid infrared. ",
        "conclusions": "We have reviewed the calibration of infrared (1 to 25$\\mu $m) photometry. Our most important conclusion is that there is very consistent behavior of solar-type and A-type stars, and that they in turn are closely consistent with virtually all direct calibration measurements and with models of their spectra. Concerns of significant inconsistencies (Price 2004; Bohlin 2007) can therefore be put aside, and we can proceed to develop a procedure for calibration of infrared measurements with assurance that there are unlikely to be serious undetected systematic errors.  We have therefore established a consistent calibration across the near and mid-infrared spectral region (1 to 25$\\mu $m). The foundation of the calibration is the accurate direct measurements near 2.2$\\mu $m and particularly near 10$\\mu $m. The accuracy of the absolute calibration is 2\\% or better across this entire wavelength range. We provide guidelines for applying it to 2MASS, IRAC, and MIPS photometry. Because of the overall agreement among the previous calibrations, the adjustments to apply to them for a fully consistent infrared calibration are small, generally within the stated errors.  After the rejection of a few stars with anomalous SEDs, upper limits of $\\sim $ 1.5\\% (rms) are placed on the intrinsic infrared SED variations in both A dwarf and solar-type stars. These types of star are therefore suitable as general-purpose standard stars, allowing the calibration to be extended readily to other photometric bands and systems. We provide spectral energy distributions of a fiducial A star and of the sun for use in extending the calibration to other systems, or for generating fainter or brighter mid-infrared standards by extrapolation from accurate near-infrared measurements.  The suggested calibration is summarized in a number of tables. Table 1 gives the zero points at 10.6 and 23.675$\\mu$m. Tables C2 and C3 provide a list of accurate measurements of A and solar-type stars at the latter wavelength. Table C1 provides measurements of solar-type stars at 7.872 $\\mu$m. These measurements can be transferred to any wavelength using the spectral energy distributions in Table A1. The listed stars are bright enough that they can be measured at high signal to noise from the ground at least through the 10$\\mu$m atmospheric window, so they provide a direct transfer between the Spitzer calibration and groundbased observations. Zero points for 2MASS photometry are provided in Table 6.  In many cases, however, it is more convenient simply to adjust measurements using alternative calibrations to the consistent scale suggested in this paper. The relevant correction factors for 2MASS K$_S$, IRAC Band 4 and IRAS Bands 1 and 2 are listed in Table 5. Corrections to the other 2MASS and IRAC bands should be similar to those listed.  Previous work has suggested an inconsistency between infrared and visible measurements of Vega, when fitted to a standard A0V star model. The improved accuracy of the infrared calibration, and its confirmation through many approaches, makes it clear that this inconsistency is real. The A0V star model normalized to the infrared calibration and extended to the visible (0.5556$\\mu $m) predicts a flux density 4.2\\% = 0.045 mag brighter than the absolute measurements of Vega near this wavelength. We have independently verified this result by transforming the Vega measurements of Johnson et al. (1966) into the 2MASS system, showing that Vega would be $\\sim$ - 0.036 at 2MASS K$_S$. It is likely that the discrepancy has roots in Vega being a rapidly rotating star seen pole-on, so that its output spectrum is affected by the surface temperature gradient associated with the rapid rotation (Gulliver et al. 1994). We conclude that the V - K$_S$ color of Vega is about +0.045 magnitudes (plus 0.014 magnitudes to account for the excess due to its ring at 2$\\mu$m) compared with the average colors of A0V stars, and that this star is not suitable to define a photometric zero point because of its eccentric SED for its spectral type.  An important feature of our approach is its use of large, homogeneous databases - for example, Hipparcos and 2MASS photometry, the Nstars classification of nearby solar-type stars, and the FEPS and MIPS GTO samples of solar-type and A-type stars observed to identify debris disks. The homogeneity and generally high accuracy of these data support a new approach to calibration. Initially large samples of stars are cleaned of members with anomalous colors (due, e.g., to reddening or photometric errors) and the results of still moderately large samples are then averaged to drive the measurement errors to small values. In addition to providing an accurate calibration that does not depend on a small number of \"ideal\" objects such as Vega was once thought to be, this approach facilitates extending the calibration over the entire sky."
    },
    "0806/0806.4368_arXiv.txt": {
        "abstract": "We investigate models of dark energy with purely kinetic multiple k-essence sources that allow for the crossing of the phantom divide line, without violating the conditions of stability. It is known that with more than one kinetic k-field one can possibly construct dark energy models whose equation of state parameter $\\wx$ crosses $-1$ (the phantom barrier) at recent red-shifts, as indicated by the Supernova Ia and other observational probes. However, such models may suffer from cosmological instabilities, as the effective speed of propagation $\\cx$ of the dark energy density perturbations may become {\\it imaginary} while the $\\wx = -1$ barrier is crossed. Working out the expression for $\\cx$ we show that multiple kinetic k-essence fields do indeed lead to a $\\wx = -1$ crossing dark energy model, satisfying the stability criterion $\\cx^2 \\geq 0$ as well as the condition $\\cx \\leq 1$ (in natural units), which implies that the dark energy is not super-luminal. As a specific example, we construct a phantom barrier crossing model involving three k-fields for which $\\cx$ is a constant, lying between $0$ and $1$. The model fits well with the latest Supernova Ia Union data, and the best fit shows that $\\wx$ crosses $-1$ at red-shift $z \\sim 0.2$, whereas the dark energy density nearly tracks the matter density at higher red-shifts. ",
        "introduction": "}  The late-time accelerated expansion of the universe is now well established through the observations of the luminosity distance -- red-shift relation for type Ia Supernovae (SNIa) \\cite{perl,tonry,riess1,riess2,astier,davis,kowal,rubin}. Measurements of the cosmic microwave background (CMB) temperature fluctuations by the Wilkinson Microwave Anisotropy Probe (WMAP) \\cite{wmap1,wmap2} as well as the large scale red-shift data from the Sloan Digital Sky Survey (SDSS) \\cite{sdss} further indicate that the universe is very nearly spatially flat and consists of about $70 \\%$ of dark energy (DE) which drives the cosmic acceleration. The value of the DE equation of state (EoS) parameter, i.e., the ratio of pressure to energy density of DE, $\\wx = \\px/\\rx$, should be less than $-1/3$ to maintain this acceleration. Although the simplest candidate for DE is a cosmological constant ($\\L$), for which $\\wx = -1$, there are serious theoretical problems associated with it \\cite{lambda}. $\\L$ needs to be extremely {\\it fine tuned} so that the DE at present is small compared to the Planck scale. A fine tuning is also required to keep the ($\\L$-) dark energy density $\\rx$ comparable with the present critical density $\\r_{0c}$, while the acceleration begins only in the recent past --- the {\\it coincidence problem}. From the observational point of view, although there are indications that $\\wx$ is very close to $-1$ at present, combined studies of SNIa, CMB and baryon acoustic oscillation (BAO) peaks do actually point toward a {\\it dynamical} DE with EoS parameter $\\wx$ slowly varying with red-shift $z$. More specifically, the so-called $\\L$CDM model, which consists of $\\L$ and cold dark matter as dominant components, although consistent with the low-$z$ observational data, does not produce sufficiently good fits with the data at relatively higher $z$ ($\\gtrsim 1$). Accordingly, there have been a large number of model-independent and model-dependent DE parameterizations explored in the literature \\cite{cpl,alam,visser,gong,param}. Extensive studies reveal that in parameterizing the DE to produce good fits with the observational data, it is always desirable to have no inherent restriction imposed on $\\wx$ from the underlying theory. For instance, it is not desirable to have $\\wx$ always $\\geq -1$, which is required for the theoretical consistency of many popular DE models, such as quintessence, tachyon, dilaton, etc. For the reviews on various dynamical aspects of DE, see \\cite{de-rev}.  Within the scope of minimally coupled canonical scalar field DE models, such as quintessence \\cite{quin}, the crossing from $\\wx > -1$ to $\\wx < -1$ can never happen and there is a need of more than one scalar field, at least one of which is of {\\it phantom} nature, i.e., carries a wrong sign in front of the kinetic term \\cite{phantom}. For this reason the $\\wx = -1$ line is commonly known as the phantom divide line (PDL). Phantom fields violate the dominant energy condition, and therefore could give rise to classical instabilities \\cite{cht}. Moreover, since the phantom energy density is unbounded from below, the quantum mechanical stability is in jeopardy. The vacuum is unstable against the decay into ghosts and positive energy particles, which couple to gravitons producing further decay into cosmic gamma rays \\cite{cline}. In an expanding Friedmann universe, the extremely repulsive (anti-gravitating) nature of the phantom fields pose more problems as the phantom energy density generically tends to infinity in finite time in future leading to the so-called ``Big Rip\" singularity \\cite{cald}, as well as some other mild singularities. For a general categorization of these singularities see for instance \\cite{noj}.  Apart from the coupled quintessence and phantom models, commonly known as the {\\it quintom} models \\cite{quintom}, there are various other models in which the phantom barrier crossing is shown to be achieved in a consistent way. Notable among these are the scalar-tensor models \\cite{polarski}, brane-world models \\cite{sahni}, string-inspired dilatonic ghost condensate models \\cite{dil}, quantum-corrected Klein-Gordon models with quartic potential \\cite{onemli}, coupled DE models \\cite{coupDE}, modified gravity models \\cite{modgrav}, H-essence (complex scalar) models \\cite{wei}, etc. In this work, we focus on the so-called k-essence models \\cite{k-ess1,k-ess2}, which incorporate a natural generalization of quintessence via a Lagrangian with non-linear dependence on the kinetic term(s) for the scalar field(s). For such models $\\wx$ is not constrained to be $\\geq -1$ and may in principle cross the PDL. However, it has been argued by Vikman \\cite{vik} that in the case of a single field k-essence model, with a Lagrangian of the form $P(\\f,X)$, where $X = -\\fr 1 2 [\\nabla \\f]^2$ is the kinetic term of the scalar field $\\f$, the PDL crossing either leads to instabilities against cosmological perturbations or is realized by a discrete set of phase space trajectories\\footnote{There are however some counter-arguments and counter-examples set in \\cite{odint,andrianov} and it is shown that for a generalized single field k-essence Lagrangian of the form $P(\\f,X) = f(\\f) X - V(\\f)$ stable PDL crossing is indeed possible for some specific functional form of $f(\\f)$ and $V(\\f)$.}. For a multiple k-essence model, involving more than one k-fields $\\fI ~(I=0,\\dots,N-1)$, with kinetic terms $\\xI = -\\fr 1 2 [\\nabla \\fI]^2$, the PDL crossing is shown to be possible even when the Lagrangian is purely a function of the kinetic terms, $P = \\sum \\pI(\\xI)$ \\cite{chim}. However, the stability of such a model against cosmological perturbations have not been explored in the literature.  A major difficulty in dealing with the equations of scalar-type cosmological perturbations for multiple scalar fields is that these equations are extremely coupled in the curvature and isocurvature modes, even when the scalar fields are only minimally coupled to gravity \\cite{hwang,hwangnoh1,kodamasasaki,maartens}. For multiple k-essence, the perturbation equations are complicated further, however there are some specific cases in which the expression for the speed of propagation $\\cx$ of the curvature perturbations could be obtained and the stability analysis could be performed \\cite{langlois,arroja}. For example, when the k-essence Lagrangian is $P({\\mathbf \\F},X)$, where ${\\mathbf \\F} = (\\f_{_0}, \\f_{_1}, \\cdots)$ and $X = - \\fr 1 2 {\\cal G}^{IJ} ({\\mathbf \\F}) \\nabla_\\m \\fI \\nabla^\\m \\fJ$ \\cite{langlois}, the propagation speed is simply given by the well-known result of Garriga and Mukhanov \\cite{gm}, i.e., $\\cx^2 = \\fr{\\pa P}{\\pa X} {\\big /} \\fr{\\pa \\rx}{\\pa X}$. On the other hand, in the multi-field Dirac-Born-Infeld (DBI) inflationary scenario, the propagation speed for the Lagrangian $P = - f^{-1} ({\\mathbf \\F}) \\sqrt{1 - 2 X f ({\\mathbf \\F})} - V ({\\mathbf \\F})$, is given in a homogeneous background by $\\cx = \\sqrt{1 - 2 X f ({\\mathbf \\F})}$ \\cite{arroja}. However, none of these expressions for $\\cx$ will hold for the most general Lagrangian $P = \\sum \\pI(\\fI,\\xI)$, corresponding to the $N$ non-interacting fields ${\\mathbf \\F} = (\\f_{_0}, \\f_{_1}, \\cdots)$. A comprehensive study of the curvature perturbation equations reveals that one can, in principle, obtain the expression of $\\cx$ for the multi k-field Lagrangian $P = \\sum \\pI(\\fI,\\xI)$ \\cite{ss}, leaving aside the isocurvature modes which couple weakly with the curvature mode in the large-scale limit \\cite{hwangnoh1,kodamasasaki}.  In this paper we first investigate whether the multiple kinetic k-essence DE models that allow the PDL crossing is cosmologically {\\it stable}, i.e., the speed of propagation $\\cx$ of the DE density perturbations satisfies the condition $\\cx^2 \\geq 0$. It is also relevant to check whether $\\cx \\leq 1$ (in natural units), so that the DE perturbations do not travel faster than light \\cite{bonvin}. However, the requirement $\\cx \\leq 1$ is not potentially stringent, because it has been argued that causal paradoxes do not necessarily arise even when $\\cx > 1$ \\cite{brun}. We find that in most circumstances (including the model described in \\cite{chim}) where $\\cx$ dynamically varies with time, the PDL crossing is associated with a flip of sign of $\\cx^2$, thus giving rise to violent instability. An appropriate choice for $\\cx^2$ in a stable DE model is therefore a {\\it constant} whose value lies between $0$ and $1$. Imposing this condition on $\\cx^2$ we work out the constraint on the field configuration. As a specific example we try to construct a stable PDL crossing DE model using this constraint and with the least number of model parameters possible, so as to determine these parameters with less uncertainties by fitting with the observational data. While within the scope of two k-fields the stability constraint leads to complicated field solutions, we construct a reasonably simpler model involving three k-fields for which stable PDL crossing is possible.  We constrain the parameters of the model using the most recent data-set released by the Supernova Cosmology Project (SCP) team \\cite{kowal}, viz., the ``Union\" compilation, which consists of $414$ SNIa reduced to $307$ usable data points after various selection cuts. Marginalizing over the Hubble constant $H_0$, we find that the best fit value of the present matter density parameter $\\O_{0m} \\simeq 0.265$, whereas the best fit DE EoS parameter $\\wx$ crosses the PDL at red-shift $z = z_c \\sim 0.2$, fairly in agreement with $z_c$ found in many model-independent DE parameterizations \\cite{pdlc}. The minimized value of the $\\chi^2$ statistic is found to be about $310.4$, which shows an improvement of $\\D\\chi^2 \\sim - 2.7$ over the minimized $\\chi^2$ for the $\\L$CDM model found with the Union data-set \\cite{rubin}. The minimized $\\chi^2$ per degree of freedom for the model discussed in this paper is $1.011$, which is also marginally better than that found in \\cite{davis} for the $\\L$CDM model using the older SNIa data-set complied by Riess {\\it et al} \\cite{riess2}. Additionally, we see that the best fit value of the DE density almost tracks the matter density over sufficiently longer red-shift range in the past until becoming dominant after $z \\simeq 0.4$. Thus the model provides a possible resolution of the coincidence problem. Extrapolation of the best fit $\\wx$ to future epochs shows a tendency of $\\wx$ to turn back towards $-1$ after reaching a minimum value. Thus the phantom regime appears to be transient and future singularities, such as the Big Rip, may be avoided.  The present model also differs from the usual quintom model \\cite{quintom} implemented using two (or possibly more) scalar fields, (at least) one of which has a negative definite kinetic energy density (i.e., phantom), and as such the EoS parameter is less than $-1$, whereas the other field(s) has(have) positive definite kinetic energy (quintessence-type) and EoS parameter(s) greater than $-1$. The negative energy field comes to dominate at late times and is responsible for the PDL crossing. In the three-field k-essence model which we study, two of the k-fields have positive definite energy density, whereas the energy density of the third k-field is negative for a sufficiently longer time in the past and present, but may become positive in distant future, as found by extrapolating the best fit model parameters to future (negative) red-shifts. In that sense the third k-field may not be regarded as equivalent to an eternal phantom, and quantum mechanical instabilities may be avoided. None of the fields have EoS parameter less than $-1$, unlike quintom models. Moreover, the negative energy k-field is not responsible for the super-acceleration ($\\wx < -1$), it is a positive energy k-field that causes the PDL crossing, again in contrast with what happens in quintom models.  This paper is organized as follows: in sec. \\ref{sec:gf} we describe the general framework of multiple k-essence DE models assuming that there are no mutual interactions between the individual k-fields. In sec. \\ref{sec:kks} we consider the simplistic case of purely kinetic multiple k-essence and show that only for more than one dynamical k-fields the PDL crossing is possible, although the field configuration is severely constrained by the stability criterion. We work out this constraint and choose the value of the speed of propagation $\\cx$ of the DE density perturbations to be a constant, in order to have a stable PDL crossing multiple kinetic k-essence DE model. In sec. \\ref{sec:model} we construct a specific multiple k-field model for which $\\cx$ is constant, and $\\wx$ crosses the PDL in recent past. In sec. \\ref{sec:modelfit} we fit this model with the $307$ Union SNIa data \\cite{kowal}, to reconstruct phenomenologically the multiple k-essence Lagrangian. We conclude with a summary and open questions in sec. \\ref{sec:concl}. In the Appendix, we work out the equations of scalar-type cosmological perturbations for a multiple k-essence field configuration so as to find the expression for the propagation speed $\\cx$. ",
        "conclusions": "}  While a PDL crossing DE model involving a single k-essence scalar field suffers from cosmological instabilities \\cite{vik}, a set of multiple k-essence fields are shown to give rise to such PDL crossing, satisfying the conditions for stability even in the simplistic scenario where the k-fields are mutually non-interacting and the individual k-essence Lagrangian densities are purely functions of the kinetic terms. Although there are severe restrictions imposed on the k-field configuration by virtue of the stability criterion, i.e., $\\cx^2 \\geq 0$, and the non-super-luminescence ($\\cx \\leq 1$) of the DE density perturbations, we still find certain field configurations for which the DE EoS parameter can cross the cosmological constant barrier. We have worked out the general equation of constraint on the field configuration, with an appropriate choice of the propagation speed, viz., $\\cx^2 =$ a constant, lying between $0$ and $1$. While a configuration of two kinetic k-fields, satisfying the constraint equation, leads to complicated solutions, we have shown that a three k-field configuration gives rise to a cosmologically stable PDL crossing DE model. For certain choices of some of the parameters, the model is found to fit well with the latest SNIa Union data compiled by the Supernova Cosmology Project team \\cite{kowal}. The value of the red-shift $z = z_c$ at which the DE equation of state parameter $\\wx$ (best fit with the data) crosses the PDL is found to be $\\simeq 0.2$, whereas the best fit present matter density parameter is found to be $\\O_{0m} \\simeq 0.265$. These are in good agreement with the values of $z_c$ and $\\O_{0m}$ found in the likelihood analysis done using other model-dependent or model-independent ansatze in the literature \\cite{gong,param,pdlc}. On the other hand, the value of $z$ at which the universe transits from a decelerating phase to an accelerating phase is found to be $\\sim 0.65$, which also agrees fairly well with other independent studies \\cite{gong,param,pdlc}.  In the model which we have presented here, the value of the square of the speed of propagation of the DE density perturbations, $\\cx^2$, is as low as $1/6$ throughout the entire course of the DE evolution. We have worked out the expression for $\\cx^2$ in a general setting of multiple k-essence fields considering the most general scalar type of metric perturbations in a Friedmann-Robertson-Walker background. Although the perturbation equations have been extremely coupled in the curvature and the isocurvature modes, assuming the coupling to be very weak in the sub-Hubble scales the speed of propagation of the curvature perturbations in the multi-k-field case is found to be the one that generalizes the well-known expression for $\\cx$ for a single k-essence field found in ref. \\cite{gm}.  The present model exhibits some additional features: Firstly, the dark energy density (best fit with the Union data) almost tracks the matter density in distant past until exceeding the later at $z \\simeq 0.4$, thus providing a possible resolution to the coincidence problem. Secondly, the best fit dark energy density, extrapolated to future epochs, does not run away to very high values leading to Big Rip-like singularities even in distant future. Thirdly, the extrapolated best fit of the dark energy equation of state parameter shows a tendency to turn back from a minimum negative value ($< -1$) towards the cosmological constant value ($= -1$) in distant future, thus indicating a transient phantom regime.  We have reconstructed the entire k-essence Lagrangian $P(\\ya,\\yb,\\yc)$, as well as the individual Lagrangian components $\\pI (\\yI)$ in terms of the corresponding field gradients $\\yI = \\dfI$, using the values of the model parameters best fitting the observational data. The reconstructed total Lagrangian, $P$ is found to be a smooth function of the total field gradient $y = \\sum \\yI$, i.e., there are no discontinuities or multi-valuedness in the parametric plot of $P$ versus $y$ for the red-shift range under consideration, $0 \\leq z \\leq 1.75$. All the individual Lagrangian components $\\pI$ as functions of the corresponding $\\yI$ also show the same smooth nature.  The reconstructed Lagrangian $P$ is therefore substantially better than the double-valued Lagrangian found for a single kinetic k-essence field in ref. \\cite{aas05} using popular parameterizations, such as that due to Alam {\\it et al} \\cite{alam} or due to Chevallier-Polarski-Linder \\cite{cpl} and with the old SNIa gold data-set \\cite{riess1}.  A few problems that may be looked upon in the context of the present model are as follows: (i) Can we find a systematic prescription in order to build a stable PDL crossing multiple kinetic k-essence model, instead of choosing the number of k-fields by hand and implementing the value of $\\cx^2$ arbitrarily? (ii) Can we extend the model to sufficiently higher red-shifts so as to match with the CMB and SDSS data? (iii) Is there a way to ascertain how far the isocurvature modes of the density perturbations affect the model at higher red-shifts? (iv) Can we ascertain the status of the future Big Rip singularities, if any, in the context of multi-k-essence, generically, i.e., by not just resorting to particular models? Or, can we generically ascertain whether the phantom regime would be eternal or transient for multiple k-essence? Research works addressing some of these questions are already in progress \\cite{ss} and we hope to report on them in near future.  Finally, it should be mentioned that although we have not specifically looked to construct tracking/scaling dark energy models in the present paper, the possible avoidance of the coincidence problem in our multi-k-essence model is particularly noteworthy. The value of $\\cx$ remains constant and less than the speed of light, unlike in the coincidence resolving single field k-essence models, which are shown to be affected by the super-luminal propagation ($\\cx > 1$) of the density perturbations \\cite{bonvin}. However, whether our purely kinetic model specifically belongs to the class of tracking or scaling k-essence dark energy, is an interesting issue which we hope to circumvent in future. Assisted accelerated solutions are shown to exist generically for the multi-field k-essence models that have scaling solutions \\cite{tsuji}. The nature of such assisted acceleration may also be of some interest for the purely kinetic multi-k-essence, that we have studied here, and is worth investigating.   \\ack{This work is supported by the Natural Sciences and Engineering Research Council of Canada. The authors acknowledge useful discussions with S. Odintsov, A. Dasgupta and especially with S. Shankaranarayanan. The authors thank L. Perivolaropoulos, whose numerical code for the likelihood analysis, using Mathematica, has been modified and used in this paper. The authors are also thankful to the anonymous referee for important remarks.}"
    },
    "0806/0806.1251_arXiv.txt": {
        "abstract": "Using a homotopic family of boundary eigenvalue problems for the mean-field $\\alpha^2$-dynamo with helical turbulence parameter $\\alpha(r)=\\alpha_0+\\gamma\\Delta\\alpha(r)$ and homotopy parameter $\\beta \\in [0,1]$, we show that the underlying network of diabolical points for Dirichlet (idealized, $\\beta=0$) boundary conditions substantially determines the choreography of eigenvalues and thus the character of the dynamo instability for Robin (physically realistic, $\\beta=1$) boundary conditions. In the $(\\alpha_0,\\beta,\\gamma)-$space the Arnold tongues of oscillatory solutions at $\\beta=1$ end up at the diabolical points for $\\beta=0$. In the vicinity of the diabolical points the space orientation of the 3D tongues, which are cones in first-order approximation, is determined by the Krein signature of the modes involved in the diabolical crossings at the apexes of the cones. The Krein space induced geometry of the resonance zones explains the subtleties in finding $\\alpha$-profiles leading to  spectral exceptional points, which are important ingredients in recent theories of polarity reversals of the geomagnetic field. ",
        "introduction": "%  Polarity reversals of the Earth's magnetic field have fascinated geophysicists since their discovery by David and Brunhes \\cite{B1906} a century ago. While the last reversal occurred approximately                           % 780000 years ago, the mean reversal rate (averaged over the last       % few million years) is approximately 4 per Myr.                     % At least two, but very likely  three \\cite{COURTILLOT}             % superchrons have been identified                                   % as ''quiet'' periods of some tens of millions of years showing    % no reversal at all.                                             %  The reality of reversals is quite complex and there            % is little hope to understand all their details within a simple   % model. Recent computer simulations of the geodynamo in             % general and of reversals in particular \\cite{G2002,WICHTOLSON,TAKAHASHI}  % have progressed much since the first fully coupled  3D simulations     % of a reversal by Glatzmaier and Roberts in 1995 \\cite{GLRO1995}.          % Most interestingly, polarity reversals were also                    % observed in one \\cite{BERHANU} of the recent liquid sodium        % dynamo experiments which have flourished during the last decade     % \\cite{RMP,ZAMM}.                                                   %  However, it is important to note                               % that neither in simulations nor in experiments it is possible to accommodate all dimensionless parameters of the geodynamo \\cite{GUBBINS}, and many of them are not even well known \\cite{ROGLA2}. In an interesting attempt to bridge the gap of several orders of magnitude between realistic and numerically achievable parameters, Christensen and Aubert \\cite{CHRAUBERT} were able to identify remarkable  scaling laws for some appropriate non-dimensional numbers.                                             %  The use of appropriate simplified models \\cite{MELBOURNE,HOYNG,RYAN,LUCA} represents another attempt to understand better the basic principle and the typical features of reversals. Most prominent among those features  are the distinct asymmetry (with a slow decay and a fast recovery phase) \\cite{VMG2005}, the clustering property of reversal events \\cite{C2006}, and the appearance of several maxima (at multiples of ~95 kyr) of the residence time distribution which has been explained in terms of a stochastic resonance phenonemon with the Milankovic cycle of the Earth's orbit excentricity \\cite{CONSOLINI,LORITO}.   One of the simplest reversal models which seems capable to explain all those three reversal features in a consistent manner \\cite{EPJB,IP} relies basically on the existence of an exceptional point in the spectrum of the non-self-adjoint dynamo operator, where two real                                        % eigenvalues coalesce and continue as a complex conjugated pair of                % eigenvalues. The importance of the specific interplay between oscillatory and non-oscillatory modes for the reversal mechanism had been early expressed by Yoshimura \\cite{YOSHI}, Sarson and Jones \\cite{SAJO}, and Gubbins and Gibbons \\cite{GUGI}. In the framework of a simple mean-field $\\alpha^2$-dynamo with a spherically symmetric helical turbulence parameter $\\alpha$ it was possible to identify reversals as noise-triggered relaxation                                                   % oscillations in the vicinity of an exceptional point \\cite{SG2005,EPSL,S2007}.                                                       % The key point is that the exceptional point is associated with a nearby local maximum of the                                                           % growth rate situated at a slightly lower magnetic Reynolds number.               % It is the negative slope of the growth rate curve between this local           % maximum and the exceptional point that makes stationary                       % dynamos vulnerable to  noise. Then, the instantaneous                          % eigenvalue is driven towards the exceptional point and beyond into the         % oscillatory branch where the sign change of the dipole polarity happens.       % Therefore, the existence of an exceptional point is an essential ingredient for reversals, although non-linear dynamics and the influence of noise must be invoked for a more detailed understanding of those events.  From the spectral point of view, the reversal phenomenon of the geomagnetic field is strongly linked to other fields of physics, like van-der-Pol like oscillators \\cite{vdP1926}, geometric phases \\cite{MKS2005}, $\\mathcal{PT}$-symmetric quantum mechanics \\cite{Ben2007,Grae2008}, $\\mathcal{PT}$-symmetric optical waveguides \\cite{PT-opt2008}, microwave resonators \\cite{D2003}, and dissipation-induced instabilities \\cite{B56,M2005,K2008}.  A particular problem of all those systems in which exceptional points are involved is a strong sensitivity of the eigenvalues on boundary conditions (BCs). As for the geodynamo, the periodic occurrence of so-called superchrons is usually  attributed to the changing thermal BCs at the core-mantle boundary \\cite{COURTILLOT}, but the growth of the inner core may also play a role \\cite{S2007} by virtue of a spectral resonance phenomenon \\cite{GK2006}.   In this context it is worthwhile to note that                          % important features of dynamos are most easily understandable when     % treated with idealized (i.e. non-physical) boundary conditions.       % This was the case for explaining the famous eigenvalue symmetry        % between dipole and quadrupole modes as it was done by Proctor in 1977   % \\cite{PROCTOR1,PROCTOR2}.                                               %  Standing in this tradition, the present paper is devoted to a          % better understanding of the interplay of BCs, the spectral             % resonance phenomenon and oscillatory regimes in  dynamos.             % ",
        "conclusions": "%  In summary, we have found that the underlying network of DPs and their intersection indices for $\\beta=0$ substantially determine the choreography of eigenvalues for $\\beta=1$ and, in particular, the loci of EPs which are important to explain the reversals of the geomagnetic field. Although this has been exemplified for the unphysical monopole ($l=0$)   % mode of a simplified spherically symmetric $\\alpha^2$ dynamo model,      % the general idea is well generalizable to physical modes and to more      % realistic dynamo models. Work in this direction is in progress.         %"
    },
    "0806/0806.3581_arXiv.txt": {
        "abstract": "s{ Heavy stable charged particles can exist, hidden from us in bound atomlike states. Models with new stable charged leptons and quarks give rise to realistic composite dark matter scenarios. Significant or even dominant component of O-helium (atomlike system of He4 nucleus and heavy -2 charged particle) is inevitable feature of such scenarios. Possible O-helium explanation for the positron excess in the galactic bulge and for the controversy between the positive results of DAMA and negative results of other experiments is proposed.} ",
        "introduction": "The widely shared belief is that the dark matter, corresponding to $25\\%$ of the total cosmological density, is nonbaryonic and consists of new stable particles. One can formulate the set of conditions under which new particles can be considered as candidates to dark matter (see e.g. \\cite{book,Cosmoarcheology,Bled07} for review and reference): they should be stable, saturate the measured dark matter density and decouple from plasma and radiation at least before the beginning of matter dominated stage. The easiest way to satisfy these conditions is to involve neutral weakly interacting particles. However it is not the only particle physics solution for the dark matter problem. As we show here, new stable particles can have electric charge, but escape experimental discovery, because they are hidden in atom-like states maintaining dark matter of the modern Universe.  Recently several elementary particle frames for heavy stable charged particles were proposed:  (a)  A heavy quark of fourth generation \\cite{I,lom,Khlopov:2006dk} accompanied by heavy neutrino \\cite{N}; which can avoid experimental constraints \\cite{Q,Okun} and form composite dark matter species; (b) A Glashow's ``sinister'' heavy tera-quark $U$ and tera-electron $E$, forming a tower of tera-hadronic and tera-atomic bound states with ``tera-helium atoms'' $(UUUEE)$ considered as dominant dark matter \\cite{Glashow,Fargion:2005xz}. (c) AC-leptons, predicted in the extension \\cite{5} of standard model, based on the approach of almost-commutative geometry \\cite{bookAC}, can form evanescent AC-atoms, playing the role of dark matter \\cite{5,FKS,Khlopov:2006dk}. Finally, it was recently shown in \\cite{KK} that an elegant composite dark matter solution is possible in the framework of walking technicolor models (WTC) \\cite{Sannino:2004qp}.  In all these models (see review in \\cite{Khlopov:2006dk,Bled07,Khlopov:2008rp}), the predicted stable charged particles form neutral atom-like states, composing the dark matter of the modern Universe. It offers new solutions for the physical nature of the cosmological dark matter. The main problem for these solutions is to suppress the abundance of positively charged species bound with ordinary electrons, which behave as anomalous isotopes of hydrogen or helium. This problem is unresolvable, if the model predicts stable particles with charge -1, as it is the case for tera-electrons \\cite{Glashow,Fargion:2005xz}.  The possibility of stable doubly charged particles $A^{--}$ and $C^{++}$, revealed in the AC model, offered a candidate for dark matter in the form of elusive (AC)-atoms. In the charge symmetric case, when primordial concentrations of $A^{--}$ and $C^{++}$ are equal, a significant fraction of relic $C^{++}$, which is not bound in (AC)-atoms, is left in the Universe and the suppression of this fraction in terrestrial matter involves new long range interaction between A and C, making them to recombine in (AC)-atoms inside dense matter bodies.  In the asymmetric case, corresponding to excess of -2 charge species, as it was assumed for $(\\bar U \\bar U \\bar U)^{--}$ in the model of stable $U$-quark of a 4th generation, their positively charged partners effectively annihilate in the early Universe. The dark matter is in the form of nuclear interacting O-helium - atom-like bound states of -2 charged particles and primordial helium, formed as soon as $He$ is produced in the Standard Big Bang Nucleosynthesis (SBBN). Such an asymmetric case was realized in \\cite{KK} in the framework of WTC, where it was possible to find a relationship between the excess of negatively charged anti-techni-baryons $(\\bar U \\bar U )^{--}$ and/or technileptons $\\zeta^{--}$ and the baryon asymmetry of the Universe.   It turned out that the necessary condition for these scenarios, avoiding anomalous isotopes overproduction, is absence of stable particles with charge -1, so that stable negatively charged particles $X^{--}$ should only have charge -2.  After it is formed in SBBN, $^4He$ screens the $X^{--}$ charged particles in composite $(^4He^{++}X^{--})$ {\\it O-helium} ``atoms'' \\cite{I}. For different models of $X^{--}$ they are also called ANO-helium \\cite{lom,Khlopov:2006dk}, Ole-helium \\cite{FKS,Khlopov:2006dk} or techni-O-helium \\cite{KK}. We'll call them all O-helium ($OHe$) in our further discussion.  In all these forms of O-helium $X^{--}$ behave either as leptons or as specific \"hadrons\" with strongly suppressed hadronic interaction. Therefore O-helium interaction with matter is determined by nuclear interaction of $He$. These neutral primordial nuclear interacting objects contribute the modern dark matter density and play the role of a nontrivial form of strongly interacting dark matter \\cite{Starkman,McGuire:2001qj}. The active influence of this type of dark matter on nuclear transformations seems to be incompatible with the expected dark matter properties. However, it turns out that the considered scenario is not easily ruled out \\cite{FKS,I,KK} and challenges the experimental search for various forms of O-helium and its charged constituents. O-helium scenario might provide explanation for the observed excess of positrons in the galactic bulge and for the controversy between positive results of dark matter searches in DAMA/NaI (see for review \\cite{Bernabei:2003za}) and DAMA/Libra \\cite{Bernabei:2008yi} experiments and negative results of other experimental groups. ",
        "conclusions": ""
    },
    "0806/0806.1856_arXiv.txt": {
        "abstract": "We discuss the generation of large-scale magnetic fields due to the breaking of the conformal invariance in the electromagnetic field through the $CPT$-even dimension-six Chern-Simons-like effective interaction with a fermion current in inflationary cosmology. It is shown that the magnetic fields on 1Mpc scale with the field strength of $\\sim 10^{-9}$G at the present time can be generated even for the scale of the effective interaction being the Planck scale. ",
        "introduction": "It is observationally known that there exist magnetic fields with the field strength $10^{-7}$--$10^{-6}$G on 10kpc--1Mpc scale in clusters of galaxies as well as $\\sim 10^{-6}$G on $1-10$kpc scale in galaxies of all types and at cosmological distances (for reviews, see~\\cite{reviews}). However, the origin of the cosmic magnetic fields, in particular the large-scale magnetic fields in clusters of galaxies, is not well understood yet. Although the dynamo amplification mechanism~\\cite{EParker} amplifies very weak seed magnetic fields up to $\\sim 10^{-6}$G, its effectiveness in galaxies at high redshifts and clusters of galaxies is still unsatisfactory. Furthermore, it is difficult for astrophysical processes~\\cite{Biermann1, PI}, cosmological phase transitions~\\cite{PT} and primordial density perturbations before or at the epoch of recombination~\\cite{DP} to generate the magnetic fields on megaparsec scales with  sufficient field strengths to account for the magnetic fields observed in galaxies and clusters of galaxies without dynamo amplification mechanism.  The most natural origin of the large-scale magnetic fields is from electromagnetic quantum fluctuations existed at the inflationary stage~\\cite{Turner:1987bw}. This is because inflation naturally produces effects on very large scales, larger than the Hubble horizon, starting from microphysical processes operating on a causally connected volume. In the Friedmann-Robertson-Walker (FRW) spacetime, the metric is conformally flat, while the ordinary Maxwell theory is conformally invariant. It is clear that the conformal invariance must have been broken at the inflationary stage in order that electromagnetic quantum fluctuations can be induced during inflation~\\cite{Parker:1968mv}. Note that this does not apply when the FRW background has nonzero spatial curvature~\\cite{NZC}. Various breaking mechanisms of the conformal invariance in the electromagnetic field, such as non-minimal gravitational coupling~\\cite{Turner:1987bw}, coupling to a scalar field~\\cite{Ratra:1991bn, DE, C-S, pseudoscalar, charged scalar, ScalarED, ScalarED-2, Enqvist:2004yy, DBI}, conformal anomaly induced by quantum effects~\\cite{Dolgov:1993vg}, spontaneous breaking of the Lorentz invariance~\\cite{Bertolami:1998dn}, noncommutative field theories~\\cite{NC}, a preferred minimal length~\\cite{Ashoorioon:2004rs}, and cosmic defects~\\cite{Hollenstein:2007kg}, have been proposed.  Recently, an $CPT$-even dimension-six Chern-Simons-like effective interaction between a fermion current and the electromagnetic field in inflationary cosmology was proposed to induce the cosmological birefringence~\\cite{Geng:2007va, GHN2} as well as  baryon number asymmetry~\\cite{BGH}. The electromagnetic quantum fluctuations can be generated during inflation as the interaction breaks not only the Lorentz invariance but also the conformal invariance of the electromagnetic field. In the present paper, we study the generation of large-scale magnetic fields due to this effective interaction.  This paper is organized as follows. In Sec.\\ II we describe our model and derive equations of motion for the $U(1)$ gauge field. In Sec.\\ III we consider the evolution of the $U(1)$ gauge field and estimate the present strength of the large-scale magnetic fields. Finally, Sec.\\ IV is devoted to a conclusion. We use units of $k_\\mathrm{B} = c = \\hbar = 1$ and adopt Heaviside-Lorentz units of electromagnetism. ",
        "conclusions": "In summary, we have studied the generation of the large-scale magnetic fields due to the breaking of the conformal invariance in the electromagnetic field through the $CPT$-even dimension-six Chern-Simons-like effective interaction with a fermion current in inflationary cosmology. We have found that the magnetic fields on 1Mpc scale with the present amplitude of $\\sim 10^{-9}$G can be generated. This strength is enough to explain the magnetic fields observed in galaxies and clusters of galaxies through only adiabatic compression without requiring any dynamo amplification~\\cite{Turner:1987bw}. If the number density of the fermion $\\psi$ interacting with the electromagnetic field evolves in proportion to $a^{-3}(t)$ during inflation, the scale of the effective interaction can be the Planck scale."
    },
    "0806/0806.1067_arXiv.txt": {
        "abstract": "Existence of metastable quarks of new generation can be embedded into phenomenology of heterotic string together with new long range interaction, which only this new generation possesses. We discuss primordial quark production in the early Universe, their successive cosmological evolution and astrophysical effects, as well as possible production in present or future accelerators. In case of a charge symmetry of 4th generation quarks in Universe, they can be stored in neutral mesons, doubly positively charged baryons, while all the doubly negatively charged \"baryons\" are combined with He-4 into neutral nucleus-size atom-like states. The existence of all these anomalous stable particles may escape present experimental limits, being close to present and future experimental test. Due to the nuclear binding with He-4 primordial lightest baryons of the 4th generation with charge $+1$ can also escape the experimental upper limits on anomalous isotopes of hydrogen, being compatible with upper limits on anomalous lithium. While 4th quark hadrons are rare, their presence may be nearly detectable in cosmic rays, muon and neutrino fluxes and cosmic electromagnetic spectra. In case of charge asymmetry, a nontrivial solution for the problem of dark matter (DM) can be provided by excessive (meta)stable anti-up quarks of 4th generation, bound with He-4 in specific nuclear-interacting form of dark matter. Such candidate to DM is surprisingly close to Warm Dark Matter by its role in large scale structure formation. It catalyzes primordial heavy element production in Big Bang Nucleosynthesis and new types of nuclear transformations around us. ",
        "introduction": "Introduction}  The question about existence of new quarks and/or leptons is among the most important in the modern particle physics. Possibility of existence of new (meta)stable quarks which form new (meta)stable hadrons is of special interest. New stable hadrons can play the role of strongly interacting dark matter \\cite{Dimopoulos:1989hk,Starkman,McGuire:2001qj}. This question is believed to find solution in the framework of future Grand Unified Theory. A strong motivation for existence of new long-living hadrons comes from a possible solution \\cite{Dvali} of the \"doublet-triplet splitting\" problem in supersymmetric GUT models. Phenomenology of string theory offers another motivation for new long lived hadrons.  A natural extension of the Standard model can lead in the heterotic string phenomenology to the prediction of fourth generation of quarks and leptons \\cite{Shibaev,Sakhenhance} with a stable 4th neutrino \\cite{Fargion99,Grossi,Belotsky,BKS}. The comparison between the rank of the unifying group $E_{6}$ ($r=6$) and the rank of the Standard model ($r=4$) can imply the existence of new conserved charges. These charges can be related with (possibly strict) gauge symmetries. New strict gauge U(1) symmetry (similar to U(1) symmetry of electrodynamics) is excluded for known particles but is possible, being ascribed to the fermions of 4th generation only. This provides theoretic motivation for a stability of the lightest fermion of 4th generation, assumed to be neutrino. Under the condition of existence of strictly conserved charge, associated to 4th generation, the lightest 4th generation quark $Q$ (either $U$ or $D$) can decay only to 4th generation leptons owing to GUT-type interactions, what makes it sufficiently long living.  Whatever physical reason was for a stability of new hypothetical particles, it extends potential for testing respective hypothesis due to its implications in cosmology. Especially rich in this sense is a hypothesis on (meta)stable quarks of new family. It defines the goal of current work.  As we will show, in the case when 4th generation possesses strictly conserved $U(1)$-gauge charge (which will be called $y$-charge), 4th generation fermions are the source of new interaction of Coulomb type (which we'll call further $y$-interaction). It can be crucial for viability of model with equal amounts of 4th generation quarks and antiquarks in Universe. The case of cosmological excess of 4th generation antiquarks offers new form of dark matter with a very unusual properties. Owing to strict conservation of $y$-charge, this excess should be compensated by excess of 4th generation neutrinos.  Recent analysis \\cite{Okun} of precision data on the Standard model parameters admits existence of the 4th generation particles, satisfying direct experimental constraints which put lower limit 220 GeV for the mass of lightest quark \\cite{PDG}.  If the lifetime of the lightest 4th generation quark exceeds the age of the Universe, primordial $Q$-quark (and $\\bar Q$-quark) hadrons should be present in the modern matter. If this lifetime is less than the age of the Universe, there should be no primordial 4th generation quarks, but they can be produced in cosmic ray interactions and be present in cosmic ray fluxes. The search for this quark is a challenge for the present and future accelerators.  In the present work we will assume that up-quark of 4th generation ($U$) is lighter than its down-quark ($D$). The opposite assumption is found to be virtually excluded, if $D$ is stable. The reason is that $D$-quarks might form stable hadrons with electric charges $\\pm 1$ ($(DDD)^-$, $(\\bar Du)^+$, $\\anti DDD^+$), which eventually form hydrogen-like atoms (hadron $(DDD)^-$ is combined with $^4He^{++}$ into $+1$ bound state), being strongly constrained in surrounding matter. It will become more clear from consideration of $U$-quark case, presented below.  The following hadron states containing (meta)stable $U$-quarks (U-hadrons) are expected to be (meta)stable and created in early Universe: ``baryons'' $(Uud)^+$, $(UUu)^{++}$, $(UUU)^{++}$; ``antibaryons'' $\\anti UUU^{--}$, $\\anti UUu^{--}$, meson $(\\bar U u)^0$. The absence in the Universe of the states $\\anti Uud$, $(U\\bar u)$ containing light antiquarks are suppressed because of baryon asymmetry. Stability of double and triple $U$ bound states $(UUu)$, $(UUU)$ and $\\anti UUu$, $\\anti UUU$ is provided by the large chromo-Coulomb binding energy ($\\propto \\alpha_{QCD}^2 \\cdot m_Q$) \\cite{Glashow:2005jy,Fargion:2005xz}. Formation of these states in particle interactions at accelerators and in cosmic rays is strongly suppressed, but they can form in early Universe and cosmological analysis of their relics can be of great importance for the search for 4th generation quarks.  We analyze the mechanisms of production of metastable $U$ (and $\\bar U$) hadrons in the early Universe, cosmic rays and accelerators and point the possible signatures of their existence.  We'll show that in case of charge symmetry of U-quarks in Universe, a few conditions play a crucial role for viability of the model. An electromagnetic binding of $\\anti UUU^{--}$ with $^4He^{++}$ into neutral nucleus-size atom-like state (\\OHe) should be accompanied by a nuclear fusion of $(Uud)^+$ and $^4He^{++}$ into lithium-like isotope $[^4He(Uud)]$ in early Universe. The realization of such % a fusion requires a marginal supposition concerning respective cross section. Furthermore, assumption of $U(1)$-gauge nature of the charge, associated to U-quarks, is needed to avoid a problem of overproduction of anomalous isotopes by means of an $y$-annihilation of U-relics ($[^4He(Uud)]$, $(UUu)$, $(UUU)$, $^4He\\anti UUU$, $^4He\\anti UUu$, $(\\bar U u)$). Residual amount of U-hadrons with respect to baryons in this case is estimated to be less than $10^{-10}$ in Universe in toto and less than $10^{-20}$ at the Earth.  A negative sign charge asymmetry of U-quarks in Universe can provide a nontrivial solution for dark matter (DM) problem. For strictly conserved charge such asymmetry in $\\bar U$ implies corresponding asymmetry in leptons of 4th generation. In this case the most of $\\bar U$ in Universe are contained in \\OHe states $[^4He\\anti UUU]$ and minor part of them in mesons $\\bar Uu$. On the other hand the set of direct and indirect effects of relic U-hadrons existence provides the test in cosmic ray and underground experiments which can be decisive for this hypothesis. The main observational effects for asymmetric case do not depend on the existence of $y$-interaction.   The structure of this paper is as the following. Section 2 is devoted to the charge symmetric case of U-quarks. Cosmological evolution of U-quarks in early Universe is considered in subsection 2.1, while in subsection 2.2 the evolution and all possible effects of U-quarks existence in our Galaxy are discussed. The case of charge asymmetry of quarks of 4th generation in Universe is considered in Section 3. Section 4 is devoted to the questions of the search for U-quarks at accelerators. We summarize the results of our present study, developing earlier investigations \\cite{BFK,KPS}, in Conclusion. ",
        "conclusions": "Conclusion}  To conclude, the existence of hidden stable or metastable quark of 4th generation can be compatible with the severe experimental constraints on the abundance of anomalous isotopes in Earths atmosphere and ground and in cosmic rays, even if the lifetime of such quark exceeds the age of the Universe. Though the primordial abundance $f= r_4/r_b$ of hadrons, containing such quark (and antiquark) can be hardly less than $f \\sim 10^{-10}$ in case of charge symmetry, their primordial content can strongly decrease in dense astrophysical objects (in the Earth, in particular) owing to the process of recombination, in which hadron, containing quark, and hadron, containing antiquark, produce unstable charmonium-like quark-antiquark state.  To make such decrease effective, the equal number density of quark- and antiquark-containing hadrons should be preserved. It appeals to a dynamical mechanism, preventing segregation of quark- and antiquark- containing hadrons. Such mechanism, simultaneously providing strict charge symmetry of quarks and antiquarks, naturally arises, if the 4th generation possesses new strictly conserved U(1) gauge ($y$-) charge. Coulomb-like $y$-charge long range force between quarks and antiquarks naturally preserves equal number densities for corresponding hadrons and dynamically supports the condition of $y$-charge neutrality.  It was shown in the present paper that if $U$-quark is the lightest quark of the 4th generation, and the lightest free $U$-hadrons are doubly charged $(UUU)$- and $(UUu)$-baryons and electrically neutral $(U \\bar u)$-meson, the predicted abundance of anomalous helium in Earths atmosphere and ground as well as in cosmic rays is below the existing experimental constraints but can be within the reach for the experimental search in future. To realize this possibility nuclear binding of all the $(Uud)$-baryons with primordial helium is needed, converting potentially dangerous form of anomalous hydrogen into less dangerous anomalous lithium. Then the whole cosmic astrophysics and present history of these relics are puzzling and surprising, but nearly escaping all present bounds.  Searches for anomalous isotopes in cosmic rays and at accelerators were performed during last years. Stable doubly charged $U$ baryons offer challenge for cosmic ray and accelerator experimental search as well as for increase of sensitivity in searches for anomalous helium. In particular, they seem to be of evident interest for cosmic ray experiments, such as PAMELA and AMS02. $+2$ charged $U$ baryons represent the low $Z/A$ anomalous helium component of cosmic rays, whereas $-2$ charged $\\bar U$ baryons look like anomalous antihelium nuclei. In the baryon asymmetrical Universe the predicted amount of primordial single $\\anti Uud$ baryons is exponentially small, whereas their secondary fluxes originated from cosmic ray interaction with the galactic matter are predicted at the level, few order of magnitude below the expected sensitivity of future cosmic ray experiments. The same is true for cosmic ray $+2$ charged $U$ baryons, if $U$-quark lifetime is less than the age of the Universe and primordial $U$ baryons do not survive to the present time.  The models of quark interactions favor isoscalar $(Uud)$ baryon to be the lightest among the 4th generation baryons (provided that $U$ quark is lighter, than $D$ quark, what also may not be the case). If the lightest $U$-hadrons have electric charge $+1$ and survive to the present time, their abundance in Earth would exceed the experimental constraint on anomalous hydrogen. This may be rather general case for the lightest hadrons of the 4th generation. To avoid this problem of anomalous hydrogen overproduction the lightest quark of the 4th generation should have the lifetime, less than the age of the Universe. Another possible solution of this problem, using double and triple $\\bar U$ baryons $\\anti UUU$ and catalysis of $(Uud)$ annihilation in atom-like bound systems $^4He\\anti UUU$ is considered in \\cite{BFK}.  However short-living are these quarks on the cosmological timescale in a very wide range of lifetimes they should behave as stable in accelerator experiments. For example, with an intermediate scale of about 10$^{11} \\GeV$  (as in supersymmetry models \\cite{Benakli}) the expected lifetime of $U$-(or $D$-) quark $\\sim 10^6$ years is much less than the age of the Universe but such quark is practically stable in any collider experiments.  First year operation of the accelerator LHC has good discovery potential for U(D)-quarks with mass up to 1.5 TeV. U-hadrons born at accelerator will distinguish oneself by high $p_t$, low velocity, by effect of a charge flipping during their propagation through the detectors. All these features enable strongly to increase efficiency of event selection from not only background but also from alternative hypothesis.   In the present work we studied effects of 4th generation having restricted our analysis by the processes with 4th generation quarks and antiquarks. However, as we have mentioned in the Introduction in the considered approach absolutely stable neutrino of 4th generation with mass about 50 GeV also bears $y$-charge. The selfconsistent treatment of the cosmological evolution and astrophysical effects of $y$-charge plasma of neutrinos, antineutrinos, quarks and antiquarks of 4th generation in charge symmetric case will be the subject of special studies. An attempt of such a treatment has been undertaken in the case of charge asymmetry, described in this paper.  We believe that a tiny trace of heavy hadrons as anomalous helium and stable neutral $\\OHe$ and mesons\\footnote{Storing these charged and neutral heavy hadrons in the matter might influence its $e/m$ properties, leading to the appearance of apparent fractional charge effect in solid matter \\cite{BFK}. The present sensitivity for such effect in metals ranges from $10^{-22}$ to $10^{-20}$ .} may be hidden at a low level in our Universe ($\\frac{n_U}{n_b} \\sim 10^{-10} - 10^{-9}$) and even at much lower level here in our terrestrial matter a density $\\frac{n_U}{n_b} \\sim 10^{-23}$ in case of charge symmetry. There are good reasons to bound the 4th quark mass below TeV energy. Therefore the mass window and relic density is quite narrow and well defined, open to a final test.  In case of charge asymmetry of 4th generation quarks, a nontrivial solution of the problem of dark matter (DM) can be provided due to neutral $\\OHe$-like U-hadrons states (ANO-helium in case of $y$-interaction existence). Such candidates to DM have many interesting implications in BBN, large scale structure of Universe and physics of DM \\cite{I,lom,KPS,FKS,Majorana,KK,Bled07,DM08}. It should catalyze new types of nuclear transformations, reminding alchemists' dream on the philosopher's stone. It challenges direct search for species of such composite dark matter and its constituents. A very low probability for their existence is strongly compensated by the expectation value of their discovery."
    },
    "0806/0806.2969_arXiv.txt": {
        "abstract": "{ }{We study the dynamical status of the poor, low X--ray luminous galaxy clusters Abell 610, Abell 725, and Abell 796 (at $z=0.1$, 0.09, and 0.16, respectively), containing diffuse radio sources (relic, relic, and possible halo, respectively).} {Our analysis is based on new spectroscopic data obtained at the William Herschel Telescope for 158 galaxies, new photometry obtained at the Isaac Newton Telescope with the addition of data recovered from the Data Release 5 of the Sloan Digital Sky Survey.  We use statistical tools to select 57, 36, and 26 cluster members and to analyze the kinematics of cluster galaxies, as well as to study the 2D cluster structure.}{The low values we compute for the global line--of--sight velocity dispersion of galaxies ($\\sigma_{\\rm V}=420-700$ \\kss) confirm that these clusters are low--mass clusters. Abell 610 shows a lot of evidence of substructure.  It seems to be formed by two structures separated by $\\sim 700$ \\ks in the cluster rest--frame, having comparable $\\sigma_{\\rm V} \\sim 200$ \\ks and likely causing a velocity gradient. The velocity of the brightest cluster member (BCMI; a bright radio source) is very close to the mean velocity of the higher velocity structure.  A third small, low--velocity group hosts the second brightest cluster member (BCMII). The analysis of the 2D galaxy distribution shows a bimodal distribution in the core elongated in the SE--NW direction and likely associated to BCMI and BCMII groups. Abell 725 and Abell 796, which are less sampled, show marginal evidence of substructure in the velocity space. They are elongated in the 2D galaxy distribution.  For both Abell 610 and Abell 725 we shortly discuss the possible connection with the hosted diffuse radio relic.}{Our results show that relic radio sources are likely connected with merger events, but are not limited to massive clusters. About the possible halo source in Abell 796, there is some evidence of a merger event in this non--massive cluster, but a pointed radio observation is necessary to confirm this halo.} ",
        "introduction": "Clusters of galaxies are not recognized as simple relaxed structures, but rather as evolving via merging processes in a hierarchical fashion from poor groups to rich clusters.  A recent aspect of these investigations is the possible connection of cluster mergers with the presence of extended, diffuse radio sources: halos and relics. These radio sources are large (up to $\\sim1$ \\hh), amorphous cluster sources of uncertain origin and generally steep radio spectra (Hanisch \\cite{han82}; see also Giovannini \\& Feretti \\cite{gio02} for a more recent review).  They are rare sources that appear to be associated with very rich clusters that have undergone recent mergers. Therefore, it has been suggested by various authors that cluster halos/relics are related to recent merger activity (e.g., Tribble \\cite{tri93}; Burns et al. \\cite{bur94}; Feretti \\cite{fer99}).  The synchrotron radio emission of halos and relics demonstrates the existence of large scale cluster magnetic fields, of the order of 0.1--1 $\\mu$G, and of widespread relativistic particles of energy density 10$^{-14}$ -- 10$^{-13}$ erg cm$^{-3}$.  The difficulty in explaining radio halos/relics arises from the combination of their large size and the short synchrotron lifetime of relativistic electrons.  The expected diffusion velocity of the electron population is of the order of the Alfven speed ($\\sim$100 \\kss), making it difficult for the electrons to diffuse over a megaparsec--scale region within their radiative lifetime. Therefore, one needs a mechanism by which the relativistic electron population can be transported over large distances in a short time, or a mechanism by which the local electron population is reaccelerated and the local magnetic fields are amplified over an extended region.  The cluster--cluster merger can potentially supply both mechanisms (e.g., Giovannini et al. \\cite{gio93}; Burns et al. \\cite{bur94}; R\\\"ottgering et al. \\cite{rot94}; see also Feretti et al. \\cite{fer02a}; Sarazin \\cite{sar02} for reviews). However, the question is still debated since the diffuse radio sources are quite uncommon and only recently have we been able to study these phenomena on the basis of sufficient statistics (few dozen clusters up to $z\\sim0.3$, e.g., Giovannini et al. \\cite{gio99}; see also Giovannini \\& Feretti \\cite{gio02}; Feretti \\cite{fer05}).  Growing evidence of the connection between diffuse radio emission and cluster merging is based on X--ray data (e.g., B\\\"ohringer \\& Schuecker \\cite{boh02}; Buote \\cite{buo02}). Studies based on a large number of clusters have found a significant relation between the radio and the X--ray surface brightness (Govoni et al. \\cite{gov01a}, \\cite{gov01b}) and between the presence of radio halos/relics and irregular and bimodal X--ray surface brightness distribution (Schuecker et al. \\cite{sch01}).  Optical data are a powerful way to investigate the presence and the dynamics of cluster mergers, too (e.g., Girardi \\& Biviano \\cite{gir02}). The spatial and kinematical analysis of member galaxies allow us to detect and measure the amount of substructure, to identify and analyze possible pre--merging clumps or merger remnants.  This optical information is complementary to X--ray information since galaxies and intracluster medium react on different timescales during a merger (see, e.g., numerical simulations by Roettiger et al. \\cite{roe97}).  Unfortunately, to date optical data are lacking or are poorly exploited to investigate the phenomenon of diffuse radio sources.  The sparse literature contains a few individual clusters (e.g., Barrena et al. \\cite{bar02}; Mercurio et al. \\cite{mer03}; Maurogordato et al. \\cite{mau08}).  In this context, we are conducting an intensive observational and data analysis program to study the internal dynamics of radio clusters by using member galaxies. Clusters already analyzed are: Abell 2219 (Boschin et al. \\cite{bos04}); Abell 2744 (Boschin et al. \\cite{bos06}); Abell 697 (Girardi et al. \\cite{gir06}); Abell 773 (Barrena et al. \\cite{bar07a}); Abell 115 (Barrena et al. \\cite{bar07b}). These are all massive clusters and, indeed, to date analyzed clusters containing a radio halo or a relic source have a large gravitational mass (larger than $0.7$ \\mqui within 2\\hh; see Giovannini \\& Feretti \\cite{gio02}). From the theoretical point of view, a large mass is an expected property in radio clusters since the energy available to accelerate relativistic particles in a merger scales as $\\sim M^2$, as discussed by Buote (\\cite{buo01}).  In this paper, we report our results about three poor Abell clusters: Abell 610, Abell 725, and Abell 796 (hereafter A610, A725, and A796) having Abell richness (i.e., the number of galaxies within 1 Abell radius and magnitude between $m_3$ and $m_{3}+2$; see Abell et al. \\cite{abe89}) of 46, 36, and 56 galaxies respectively. These clusters contain a radio relic or a (possible) radio halo. In particular, A610 exhibits a relic source confirmed by pointed observations (Giovannini \\& Feretti \\cite{gio00}); A725 presents a relic source well visible in survey data (Kempner \\& Sarazin \\cite{kem01}); and A796 possibly hosts a halo source visible in survey data (Kempner \\& Sarazin \\cite{kem01}) which needs to be confirmed. To date, no measure of internal velocity dispersion and/or X--ray temperature have been reported in the literature.  We have recently carried out spectroscopic observations with the William Herschel Telescope giving new redshift data for 158 galaxies in the field of these clusters, as well as photometric observations at the Isaac Newton Telescope for A725. We recover additional photometric and spectroscopic information from the Data Release 5 of the Sloan Digital Sky Survey (SDSS DR5).  The paper is organized as follows. We present the new optical data in Sect.~2. We present the relevant analyses and conclusions in Sects.~3, 4, and 5 for A610, A725, and A796, respectively.  We summarize and discuss our results in Sect.~6.  Unless otherwise stated, we give errors at the 68\\% confidence level (hereafter c.l.).  Throughout the paper, we assume a flat cosmology with $\\Omega_{\\rm m}=0.3$, $\\Omega_{\\Lambda}=0.7$ and $H_0=70$ $h_{70}$ \\ks Mpc$^{-1}$. For this cosmological model, 1\\arcm corresponds to 108, 102, and 163 \\kpc at the A610, A725, and A796 redshift. ",
        "conclusions": ""
    },
    "0806/0806.2158_arXiv.txt": {
        "abstract": "High-redshift submillimetre-bright galaxies identified by blank field surveys at millimetre and submillimetre wavelengths appear in the region of the IRAC colour--colour diagrams previously identified as the domain of luminous active galactic nuclei (AGNs).  Our analysis using a set of empirical and theoretical dusty starburst spectral energy distribution (SED) models shows that power-law continuum sources associated with hot dust heated by young ($\\la 100$ Myr old), extreme starbursts at $z>2$ also occupy the same general area as AGNs in the IRAC colour--colour plots.  A detailed comparison of the IRAC colours and SEDs demonstrates that the two populations are distinct from each other, with submillimetre-bright galaxies having a systematically flatter IRAC spectrum ($\\ga1$ mag bluer in the observed [4.5]--[8.0] colour).  Only about 20\\% of the objects overlap in the colour--colour plots, and this low fraction suggests that submillimetre galaxies powered by a dust-obscured AGN are not common.  The red IR colours of the submillimetre galaxies are distinct from those of the ubiquitous foreground IRAC sources, and we propose a set of IR colour selection criteria for identifying SMG counterparts that can be used even in the absence of radio or \\spitzer\\ MIPS 24 \\micron\\ data. ",
        "introduction": "One of the most exciting developments of the past decade has been the resolution of the cosmic far-infrared background into discrete sources, providing a first glimpse of the rapid build-up of massive galaxies in the early universe long predicted by theory.  Deep, wide blank--field surveys at millimetre (mm) and submillimetre (submm) wavelengths \\citep{smail97,barger98,hughes98,eales99,scott2002,wang04,wang06,greve04,laurent05,mortier05,bertoldi07,scott08} have shown that ultraluminous infrared galaxies (ULIRGs) at $z\\ga 1$ contribute significantly to the observed far-IR background.  Multi--wavelength follow--up studies of these so-called submillimetre galaxies (SMGs) suggest that they are massive, young galaxies seen during the period of rapid stellar mass build-up, with very high specific star formation rates % at $z>1$ \\citep[see review by][]{blain02}.  A major obstacle in understanding the nature of this luminous dusty galaxy population is the limited angular resolution ($>10\\arcsec$) of current instrumentation, which prevents unambiguous identification of their counterparts at other wavelengths.  Deep interferometric radio imaging and \\spitzer\\ 24 \\micron\\ MIPS imaging are shown to be effective for identifying a subset of SMGs at $z\\la3$ \\citep{ivison02,chapman05,pope06}, and our understanding of their nature and evolution is based almost entirely on some 100 SMGs identified this way.  Recent 890 \\micron\\ continuum imaging of the seven brightest AzTEC \\citep{wilson08} 1100 \\micron\\ sources in the COSMOS field \\citep{scoville07} using the Smithsonian Submillimeter Array (SMA) by \\citet{younger07} has shown that there may be a substantial population of higher redshift ($z>3$) SMGs that are extremely faint or undetected at radio and/or 24 \\micron\\ MIPS bands.  This discovery accentuates the need for a new method to identify and investigate the higher redshift SMGs in order to quantify their contribution to the cosmic energy budget and star formation history, and to map their evolution over time.  One common feature among all seven COSMOS AzTEC/SMA sources and the four other  SMGs detected with the SMA by \\citet{iono06}, \\citet{wang07}, and \\citet{younger08} is that they are all detected in the \\spitzer\\ IRAC bands at $\\ga1\\mu$Jy level, raising the exciting possibility that deep IRAC imaging may provide a powerful new tool for identifying and obtaining deeper understanding of the SMG phenomenon.  The IRAC bands cover at $z\\sim3$ the rest-frame optical to near-IR portion of the spectral energy distribution (SED) and offer direct insight to the properties of their stellar component.  A detailed analysis of the multiwavelength properties of the AzTEC/SMA sources is presented elsewhere (Yun et al., in prep.).  In this paper, we examine the rest-frame optical/near-IR properties of a large sample of well-studied SMGs and the utility of the IRAC colour--colour plots for identifying and investigating the nature of the SMG population in general.  \\begin{figure*} \\centering \\includegraphics[width=5.5in]{figure1.eps} \\caption{{\\it Spitzer}/IRAC $S_{5.8}/S_{3.6}$ vs. $S_{8.0}/S_{4.5}$ colour--colour diagram for SMGs and IR QSOs. {\\bf (a)} Empty squares are SMGs securely identified by radio and CO interferometric imaging.  Filled squares and pentagons are SMGs identified as ``starburst'' and ``starburst+AGN'' by \\spitzer\\ IRS spectra, respectively \\citep{menendez07,valiante07,rigby08,pope08}.  IR QSOs with power-law spectrum identified in the FLS field by Lacy et al. and \\citet{martinez08} are shown as stars.  Small dots represent 4000 random field galaxies selected in the COSMOS field \\citep{sanders07}, and the large circle centered near ($-$0.4,$-$0.4) represents the centroid of the IRAC sources in the First Look Survey (FLS) field identified by \\citet{lacy04}. The long-dashed line outlines the region for AGNs proposed by Lacy et al.  The solid line shows the extended region we propose for the identification of the SMG counterparts (Eq.~1).  An error bar corresponding to a typical 10\\% uncertainty in the flux density measurement is shown on the right side. {\\bf (b)} The same {\\it Spitzer}/IRAC colour--colour diagram for IRAC sources found within a 6\\arcsec\\ radius of the nine securely identified mm/submm sources.  Filled circles are the nine SMGs identified by direct submm interferometric measurements obtained using the SMA while the empty circles represent the foreground/interloper IRAC sources. {\\bf (c)} Redshift evolution colour--colour tracks for three different starburst ages based on theoretical starburst SED models by \\citet{efstathiou00}.  Filled and empty squares along the redshift evolution colour tracks mark the redshifts of $z=0.5$, 1, 2, 3, 4, and 5.  The thick solid line represents power-law spectrum sources with $\\alpha$=0.3-1.0 ($S_\\nu \\propto \\nu^{-\\alpha}$). {\\bf (d)} The effects of extinction are demonstrated by the 72 Myr old starburst colour--colour model tracks with total extinction of $A_V=50$ and 200. } \\label{fig:colourcolourplot} \\end{figure*} ",
        "conclusions": "\\label{sec:summary}  High-redshift submillimetre-bright galaxies identified by blank field surveys at millimetre and submillimetre wavelengths appear in the region of the IRAC colour--colour diagrams previously identified as the domain of luminous AGNs by \\citet{lacy04} and \\citet{stern05}.  Rather than interpreting this as a sign that the majority of the SMGs are powered by a luminous AGN, we have shown using empirical and theoretically motivated dusty starburst SED models that their IR colours can be interpreted as those of power-law continuum associated with hot dust heated by young ($\\la 100$ Myr old), extreme {\\it starbursts} at $z>2$.  These SMGs fall along the branch of galaxies extending from the blue photospheric peak toward the red power-law region in the IRAC colour--colour plot by \\citet{lacy04}, and our analysis suggests that this branch is a heterogeneous ensemble of power-law AGNs and dust obscured starbursts whose continuum near-IR luminosity exceeds that of the photospheric emission from cool stars.  In fact, our analysis demonstrates that the popular red IRAC colour selection methods for AGN identification, such as the criteria by Lacy et al. and Stern et al., should be used with caution because of the significant starburst contribution expected.  While there is some overlap between SMGs and AGNs in these IRAC colour--colour plots, SMGs are systematically bluer ($\\langle[4.5]-[8.0]\\rangle\\ga 1$ mag), consistent with 30 to 70 Myr old starbursts observed at redshifts between $z\\sim1$ and $z\\sim5$.  Our examination of the model dusty starburst SEDs of \\citet{efstathiou00} shows that the main physical parameters that determine the observed IRAC colours are starburst age and redshift, while extinction (column density) plays a less important role.  For the models examined in Figures~\\ref{fig:colourcolourplot} and \\ref{fig:colourcolourplot2}, starburst age and redshift are nearly degenerate.  The use of IRAC colours as a redshift indicator was first proposed by \\citet{simpson99} and \\citet{sawicki02}, and various empirical photometric redshift relations have been proposed recently with estimated accuracies of $\\delta z/(1+z)=0.1-0.2$ \\citep[see][]{pope06,wilson08b}.   This photometric redshift method may not work well for galaxies with red power-law continuum and/or weak 1.6 \\micron\\ photospheric feature that are more common at $z\\ga2$.  The degeneracy between the starburst age and redshift found with these theoretical model SEDs suggests that a careful calibration using a large sample of SMGs is needed in order to quantify fully the systematic uncertainty of this method.  A detailed comparison of the IRAC colour--colour plots and SEDs shows that AGNs and SMGs are distinct from each other due to intrinsic differences in their energy source and dust distribution.  SMGs as a group have a flatter SED (bluer by $\\langle[4.5]-[8.0]\\rangle\\ga 1$ mag) in comparison with AGNs.  Only 20\\% of the objects overlap in the colour--colour plots shown in Figures~\\ref{fig:colourcolourplot}~\\&~\\ref{fig:colourcolourplot2}, and this suggests that SMGs powered by an AGN are not common.  In the context of the ULIRG-QSO evolutionary scenario \\citep{sanders88,norman88}, the little overlap between the AGNs and the SMGs population may indicate that the transition period is much shorter than the duration of the SMG or the IR AGN phase.  The red IR colours of the SMGs are distinct from the colours of the majority of the foreground IRAC sources, which are both early and late type galaxies at $z\\la2$.  We show that colour selection criteria similar to those of AGNs proposed by Lacy et al. and Stern et al. can be used to pinpoint the IRAC counterpart to an SMG uniquely or to pare down the candidates to just a few, improving the efficiency of expensive follow-up observations at other wavelengths.  The lack of high-quality multi-wavelength data for a large sample of SMGs is currently the primary limiting factor that prevents a better understanding of the SMG phenomenon.  The IRAC colour selection method discussed here appears to work even in cases when the radio and MIPS 24 \\micron\\ counterparts are too faint to be detected at the present sensitivity, and an unbiased investigation of the entire SMG population may become possible using this new identification method.  There are some important advantages and disadvantages to using the ($S_{5.8}/S_{3.6}$ vs. $S_{8.0}/S_{4.5}$) colour--colour plot (Figure~\\ref{fig:colourcolourplot}) versus the ([5.8]--[8.0] vs. [3.6]--[4.5]) colour--colour plot (Figure~\\ref{fig:colourcolourplot2}) for SMG candidate identification.  The former uses IR colours with a longer baseline in wavelength, making the colour measurements more robust and the model colour tracks better behaved, as shown in Figure~\\ref{fig:colourcolourplot}.  On the other hand, the instrumental sensitivity of the IRAC 3.6 \\micron\\ and 4.5 \\micron\\ bands are a factor of a few better than the 5.8 \\micron\\ and 8.0 \\micron\\ bands, and the 5.8 \\micron\\ and 8.0 \\micron\\ channels will be inoperable during the upcoming warm \\spitzer\\ operation.  The model colour tracks shown in Figure~\\ref{fig:colourcolourplot2} are more irregular because of shorter wavelength baselines in these colours, but the [3.6]--[4.5] colour is the best determined quantity among all IRAC colour combinations.  A far simpler and potentially more robust alternative SMG candidate selection criterion is [3.6]--[4.5] $>-0.2$.  This single colour selection method does only a slightly poorer job of rejecting foreground sources than the full colour section criteria described by Eq.~2.  Given the superior sensitivity of the IRAC 3.6 \\micron\\ and 4.5 \\micron\\ bands and their availability during the warm \\spitzer\\ mission, this simpler colour selection criterion may be the more effective method for identifying SMG counterpart candidates in the longer term.  Alternatively, colours of IRAC to optical or near-IR bands or to MIPS 24 \\micron\\ bands have also been proposed previously for identifying AGN activity \\citep[see][]{huang04,ashby06,webb06}, and similar analysis may also prove fruitful for SMG candidate identification."
    },
    "0806/0806.0531_arXiv.txt": {
        "abstract": "{} {We present mid-IR spectro-interferometry of the Seyfert type 1 nucleus of \\object{NGC 3783}. The dusty circumnuclear environment is spatially resolved and the wavelength dependence of the compact emission is discussed.} {The observations were carried out with the MIDI instrument at the Very Large Telescope Interferometer in the $N$-band. Spectra and visibilities were derived with a spectral resolution of $\\lambda/\\Delta\\lambda\\sim30$ in the wavelength range from 8 to $13\\,\\mu$m. For the interpretation we developed a simple dusty disk model with small and variable covering factor.} {At baselines of 65 and 69\\,m, visibilities in the range of 0.4 to 0.7 were measured. The $N$-band spectra show a monotonic increase of the measured flux with wavelength with no apparent silicate feature around $10\\,\\mu$m. We find that the mid-IR emission from the nucleus can be reproduced by an extended dust disk or torus with a small covering factor of the radiating dust clouds. } {Our mid-IR observations of NGC~3783 are consistent with a clumpy circumnuclear dust environment. {  The interpretation in terms of a dusty torus with low covering factor supports a clumpy version of the unified scheme for AGN}. The inferred sizes and luminosities are in good agreement with dust reverberation sizes and bolometric luminosities from optical and X-ray observations. } ",
        "introduction": "Unified scenarios of AGN suggest that the central continuum source is surrounded by a dusty molecular torus. This torus is thought to be responsible for aspect-angle-dependent obscuration of the central source \\citep[e.g.,][]{Antonucci1993ARAA_31_473}. Mid-infrared (MIR) interferometry of the Seyfert~2 nucleus in \\object{NGC 1068} \\citep{Jaffe2004,Poncelet2006AA.450.483} and the \\object{Circinus galaxy} \\citep{Tristram2007AA.474.837} provided supporting evidence for the dusty torus model. The MIR emission of these nuclei was resolved by MIR interferometry with the MIDI instrument in both cases. The resolved structures are belived to be the IR-bright innermost region of the torus. In Seyfert type 1 sources the situation is not as clear as NIR interferometry data obtained with the Keck interferometer on \\object{NGC 4151} \\citep{Swain2003ApJ.596L.163} have been interpreted as accretion disk emission.  At a redshift of $z=0.0097$ \\citep[][$D = 44$\\,Mpc]{Theureau1998AAS_130_333}, the galaxy NGC~3783 harbors a bright Seyfert 1 core with a black hole of $M_{\\rm BH} = 3.0 \\pm 0.5 \\times 10^7 M_\\odot$ derived from reverberation mapping \\citep{Onken2004ApJ_615_645}. NGC~3783 has been the target of extensive ultraviolet spectral monitoring \\citep[e.g.,][]{Reichert1994ApJ_425_582}. A narrow iron K line is seen in XMM-Newton spectra \\citep{Reeves2004ApJ_602_648} consistent with an origin in Compton-thick matter such as a molecular and dusty torus. The optical spectrum shows blueward asymmetries in forbidden lines and broad lines with up to 3800 km/s in the broadest components \\citep{Evans1988ApJS_67_373}. The  extended narrow line region of NGC~3783 seen with HST \\citep{Schmitt2003ApJS_148_327} appears halo-like in [O{\\sc iii}] with a size of $1\\farcs9$\\, (400 pc). The NIR spectrum of the nucleus \\citep{Reunanen2003MNRAS_343_192} features a broad  Br$\\gamma$ line and three coronal lines are detected. MIR spectra within a beam of $4\\farcs2$\\, were presented by \\citet{Roche1991MNRAS_248_606} and showed an almost featureless spectrum with a $N$-band luminosity of $\\sim 3\\times 10^9 L_\\odot$ in this aperture. VISIR photometry presented by \\citet{Horst2006AA_457L_17} show that NGC~3783 falls onto the MIR vs. hard X-ray correlation for AGN with a rising flux of 0.66 to 0.72\\,Jy across the $N$-band. The UV luminosity is estimated from $L_{\\rm UV} \\approx 6\\,\\lambda L_\\lambda (5100$\\,\\AA) to be $1-2\\times 10^{37}\\,$W. This is consitent with the $L_{\\rm bol} - L_{\\rm X}$-relation \\citep{Marconi2004MNRAS_351_169} which predicts a bolometric luminosity of $3 \\times 10^{37}$\\,W for NGC~3783. The accretion disk around the central black hole is therefore radiating at $\\sim 10\\%$ of the Eddington limit. All these characteristics of a classical type 1 core of suggest the presence of a dusty nuclear torus seen almost face-on. Direct evidence of this torus can only be obtained through high spatial resolution interferometry of the hot dust emission in the near- and mid-IR. In this paper we present the first spectro-interferometric observations of NGC~3783 with MIDI/VLTI. In Sec.\\,2 we describe the observations and data reduction. In Sec.\\,3 we develop a simple model for the circumnuclear dust distribution and discuss the implications for dusty tori in the unification scheme of AGN. Sec.\\,4 gives a short summary. ",
        "conclusions": "\\subsection{General properties} The first and most notable feature of the visibilities in Fig.\\,\\ref{N3783VisC} and the spectra in Fig.\\,\\ref{N3783Spec} is the absence of the $9.7\\,\\mu$m absorption or emission feature of silicate dust expected to be present in the torus. This absence may be due to missing silicates in AGN dust, because of the lower sublimation temperature $T_{\\rm Si} \\sim 1000\\,$K when compared with graphite grains $T_{\\rm C} \\sim 1500\\,$K. But the mid-IR emission is expected to come from $T\\sim 300\\ldots 800\\,$K dust, where both kinds of dust grains should be present. Alternatively, the dust grain size distribution maybe dominated by large grains, for which the silicate feature is less prominent. We have no direct evidence for this in NGC~3783, except for dust-reverberation arguments (see below). The third possibility, and our choice for the models below in Sec.\\,\\ref{Model:Sec}, is a radiative transfer effect: For moderate optical depth in the 8--$13\\,\\mu$m range with $\\tau \\sim 0.5$ and small temperature variations along the line of sight through the dusty torus, the emission spectrum approaches a black body and the feature is suppressed.  Secondly, the flux spectrum $F_\\nu$ in Figs.\\,\\ref{N3783Spec} and \\ref{N3783SpecComp} shows a linear rise with wavelength, in contrast to most radiative transfer simulation of face-on dusty AGN tori \\citep[e.g.,]{Honig2006AA_452_459}. This may indicate the presence of additional, spatially distributed heating sources, which can increase the amount of dust at intermediate temperatures of 150 to 300\\,K. Another possibility is a smaller covering factor of hot dust compared to cooler dust in a disk heated by the central engine. This idea is employed in the models below.  In addition, we find almost constant visibilities between 8--$13\\,\\mu$m suggesting a linear rise of the source size with wavelength. The derived visibilities are smaller than expected for a compact AGN dust torus, for which the NIR and MIR emission is dominated by a small region close to the dust sublimation radius. The diameter of the inner rim of the torus is expected (see Eq.\\,\\ref{Temp:Eq} below) to be $2.1\\,$mas corresponding to $V$ close to unity. A small dust covering factor in the models below can explain the data and imply a large extension of the MIR emission region.  Information on the torus inclination can be infered from spectropolarimetry in the optical 400 to 500\\,nm region, which shows a relatively large polarisation fraction of $\\sim 1$\\% % at P.A.$\\,\\sim 130\\degr$ \\citep[][and Kishimoto priv.comm.]{Lira2007ASPC..373..407L}. The axis of the optical scattering region is therefore inclined with respect to our line of sight suggesting the minor axis of an elongated MIR disk to be close to the P.A.$\\sim120\\degr$ of our observations. The P.A. of the interferometric baselines for the datasets 1\\,\\&\\,2 are very similar, so that we cannot constrain any non-circularity of the MIR emission from our data.  \\subsection{Modeling visibilities and spectra with a dusty disk \\label{Model:Sec}} For the interpretation of the data presented above we employ a simplified version of the clumpy torus model \\citep{Beckert2004AA_426_445,Honig2006AA_452_459}. A torus seen almost face-on ($i \\approx 0\\degr$) appears as a disk composed of dusty clouds. The inclination has almost no observable effect on spectra and colors as long as $i$ is smaller than the half opening angle of the torus. We assume single temperature clouds, where the cloud temperature depends only on the radius in the disk. For simplicity and comparison with other models we use a MRN \\citep{Mathis1977ApJ_217_425} grain size distribution between $0.005\\,\\mu$m and $0.25\\,\\mu$m. The inner radius of the disk or torus is at $R_{\\rm in}$, where we assume a dust temperature of $T_{\\rm in} = 1500\\,$K, which is suggested by NIR colors of type 1 AGN \\citep{Kishimoto2007AA.476.713}. The local intensity $I_\\nu(r) $ emitted by clouds is a modified black body $B(\\nu,T)$ at temperature $T = T_{\\rm in} \\times (r/R_{\\rm in})^{-\\eta}$ with optical depth $\\tau_\\nu$, so that \\begin{equation} I_\\nu(r) = B(\\nu,T(r))\\left(1-e^{-\\tau_\\nu}\\right)\\;. \\end{equation} We are interested in lines of sight through the disk parallel to the disk symmetry axis. For simplicity, we use a constant $\\tau_\\nu$ independent of the radius. The frequency dependence of $\\tau_\\nu$ is derived from the MRN distribution and a mixture of 50\\% silicates and 50\\% graphite grains with the optical properties taken from \\citet[][and references therein]{Weingartner2001ApJ_548_296}\\footnote{Files with tabulated optical constants have been taken from http://www.astro.princeton.edu/$\\sim$draine/dust/dust.diel.html.}. For comparison we fix $\\tau_\\nu$ in the optical $\\tau_V$ at $\\lambda=560\\,$nm. For normal black body emission in radiative equilibrium with the central illuminating UV/optical radiation we expect an exponent $\\eta = 0.5$. This leads to an almost constant flux spectrum $F_\\nu$ when integrated over the disk surface. We employ two modifications (models A and B) to fit the total flux spectrum in Fig.~\\ref{N3783Spec} (dash-dotted lines). Model A-1 \\& A-2 for dataset 1 \\& 2 assumes a different exponent $\\eta =0.35$ (see Table\\,\\ref{table:1}). This is motivated by the equilibrium temperature of graphite grains \\citep[see e.g.,][]{Barvainis1987ApJ_320_537,Kishimoto2007AA.476.713} \\begin{equation}\\label{Temp:Eq} \\frac{T}{1500\\,{\\rm K}} = \\left[\\frac{L}{10^{37}\\,{\\rm W}}\\right]^{0.179} \\left[\\frac{R}{0.13\\,{\\rm pc}}\\right]^{-0.357}\\left[\\frac{\\tilde{a}}{0.05\\,\\mu{\\rm m}}\\right]^{-0.179} \\end{equation} where the grain size  $\\tilde{a} = \\left<a^3\\right>/\\left<a^2\\right>$ is an average size for the grain size distribution. Such a model can explain the spectra in Fig.\\ref{N3783Spec} but not the visibilities as shown in Fig.\\,\\ref{N3783VisC}. We can accommodate the suggested large size of the dust emitting region by introducing a covering factor $f(r)$, for which we assume a radial dependence $f(r) = f_0 (r/R_{\\rm in})^\\xi$. For models A, we use $\\xi =0$ and $f_0$ can be adjusted to get the size of the source. The idea of a low covering factor is a direct consequence of clumpy torus models \\citep[see e.g.,][]{Honig2006AA_452_459}. In this scenario only clouds directly illuminated by the central source contribute significantly to the torus emission and due to the clumpy structure there is a non-zero probability for such clouds at all radii.  The implied AGN luminosity for models A-1 \\& A-2 are much smaller (see Tabel \\ref{table:1}) than derived from optical and X-ray observations. We therefore concentrate in the following on models B-1 \\& B-2 which provide a better fit to the spectra in Fig.\\ref{N3783Spec} and the visibilities in Fig.\\,\\ref{N3783VisC} with a running covering factor. All models exhibit decreasing visibilities with wavelength in Fig.\\,\\ref{N3783VisC} and both models A and B give acceptable fits to the datasets 1 \\& 2. The parameters of the models are summarized in Table\\,\\ref{table:1}. The inner torus radii of the models are determined by the total flux level and the covering factor in the combination $R_{\\rm in}/\\sqrt{f_0}$. The AGN luminosity (assuming isotropic emission) in Table\\,\\ref{table:1} is derived from Eq.\\,\\ref{Temp:Eq}. The larger luminosities for B-1, B-2 and B-3 are a direct consequence of the smaller covering factor at the inner boundary $R_{\\rm in}$. Note that the angular scale for $R_{\\rm in}=0.1\\,$pc % is only 0.47\\,mas. The actual $T=300\\,$K diameter is as large as $70\\,$mas for model B-1 or 60\\,mas for A-1.  \\begin{table} \\caption{Dusty disk/torus model results. The covering factor $f_0$ and the exponents $\\xi$  and $\\eta$ for the covering factor and dust temperature are described in Sec.\\,\\ref{Model:Sec}. $\\tau_V$ is the optical depth along a vertical path through clouds in the disk taken at $560\\,$nm. } % \\label{table:1} % \\begin{center} % \\begin{tabular}{c c c c c c c} % \\hline\\hline % Model & $\\tau_V$ & $f_0$ & $\\xi$ & $\\eta$ & $R_{\\rm in}$ & $L_{\\rm bol}$ \\\\ % & & & & & [pc] & [$10^{37}$\\,W] \\\\ \\hline % A-1 & 6 & 0.04 & 0 &$0.35$ & 0.064 & 0.16 \\\\ A-2 & 12 & 0.06 & 0 & $0.35$ & 0.04 & 0.065 \\\\ B-1 & 4 & $4.0\\,10^{-3}$ & 0.8  & $0.5$ & 0.30 & 3.7 \\\\ % B-2 & 8 & $3.8\\,10^{-3}$  & 0.9 & $0.5$ & 0.21 & 1.7 \\\\ B-3 & 8 & $3.2\\,10^{-3}$  & 0.9 & $0.5$ & 0.23 & 2.2 \\\\ \\hline % \\end{tabular} \\end{center} \\end{table} {  The visibilities derived with the combined total flux spectrum in Fig.~\\ref{N3783SpecComp} are much more consistent with each other. The difference between dataset 1 \\& 2 can be completely attributed to the difference in baseline length. This allows to fit the long wavelength side the NGC~3783 datasets with a single model B-3 which is very similiar to B-2 (see Table~\\ref{table:1}). In the 8.3 to $9.3\\,\\mu$m region the derived $V(\\lambda)$ is substantially lower than suggested by all simple power-law models. NGC~3783 appears to be even larger than the models in this wavelength region.}  \\subsection{Implications for unified models} Models of dusty nuclear tori suggested by the unified scheme of AGN \\citep{Antonucci1993ARAA_31_473} require the dust to be contained in individual clouds or clumps \\citep{Krolik1988ApJ_329_702}. In recent years, models for the cloud size and cloud distribution \\citep{Beckert2004AA_426_445} and for the emission from a clumpy dust distribution \\citep{Nenkova2002ApJ_570L_9,Honig2006AA_452_459} have been developed. The dynamical model based on cloud collisions \\citep{Beckert2004AA_426_445} suggests that a geometrically thick and obscuring torus is expected for high mass accretion rates close to the Eddington limit of the central black hole. From the X-ray and optical/UV luminosity of NGC~3783 an Eddington ratio of 10\\% has been inferred. Our simple model for the mid-IR emission suggests that the 'torus' has a much smaller covering factor than expected for the collisional scenario.  Long term optical/IR photometry of NGC~3783 \\citep[][and references therein]{Glass2004MNRAS_350_1049} suggests a time delay of 190\\,d between $U$- and $L$-band\\footnote{In an earlier paper \\citep{Glass1992MNRAS_256P_23} a time delay of 85\\,d was suggested, which would imply a much smaller physical scale for the response.}. This implies a physical path difference of $0.16\\,$pc or an angular scale of 0.75\\,mas, if the response is in the plane of the torus/disk. Given the uncertainties entering in Eq.\\,\\ref{Temp:Eq} we get an agreement with the B-models within a factor of 2. {  For comparision, the NIR interferometry of NGC~4151 by \\citet{Swain2003ApJ.596L.163} suggests a more compact emission region like an accretion disk or a ring at $R_{\\rm in}$. } A discussion of dust sublimation radii, dust reverberation time delay and luminosity dependence can be found in \\citet{Kishimoto2007AA.476.713}. The inferred AGN luminosity from the B-models (Table \\ref{table:1}) of $L_{\\rm bol} \\approx 3\\,10^{37}\\,$W or from the reverberation size via Eq.\\,\\ref{Temp:Eq} are in nice agreement with the $L_{\\rm bol} - L_{\\rm X}$-relation \\citep{Marconi2004MNRAS_351_169}."
    },
    "0806/0806.2228_arXiv.txt": {
        "abstract": "Radio synthesis imaging is dependent upon deconvolution algorithms to counteract the sparse sampling of the Fourier plane. These deconvolution algorithms find an estimate of the true sky brightness from the necessarily incomplete sampled visibility data. The most widely used radio synthesis deconvolution method is the CLEAN algorithm of \\Hogbom. This algorithm works extremely well for collections of point sources and surprisingly well for extended objects. However, the performance for extended objects can be improved by adopting a multi-scale approach. We describe and demonstrate a conceptually simple and algorithmically straightforward extension to CLEAN that models the sky brightness by the summation of components of emission having different size scales. While previous multiscale algorithms work sequentially on decreasing scale sizes, our algorithm works simultaneously on a range of specified scales. Applications to both real and simulated data sets are given. ",
        "introduction": "Radio synthesis imaging of astronomical sources was revolutionized by the invention of the CLEAN algorithm \\citep{Hogbom1974}. This simple algorithm enabled synthesis imaging of complex objects even with relatively poor Fourier plane coverage, such as occurs with partial earth rotation synthesis or with arrays composed of small numbers of antennas \\citep[see][]{Thompsonetal2001}.  The algorithm is motivated by the observation that in the image formed by simple Fourier inversion of the sampled visibility data (the ``dirty'' image), each point on the sky is represented by a suitably scaled and centered point spread function. Hence one can find the brightest point source by simply performing a cross-correlation of the dirty image with the point spread function. Removing the effects of the brightest point then enables one to find the next brightest, and so on. This simple algorithm works very well for collections of point sources, and surprisingly well for extended objects. Convergence is slow for extended objects (as might be expected since one is trying to identify a potentially large number of pixels one by one), and instabilities occur \\citep{Cornwell1983}. For this reason, other deconvolution algorithms such as the Maximum Entropy Method \\citep{GullDaniell1978,CornwellEvans1985} are often preferred. However, the CLEAN algorithm remains attractive for many purposes so the natural question arises as to how the basic algorithm can be improved in convergence and stability while maintaining the advantages of simplicity and noise behavior.  There are many different deconvolution algorithms in use in various scientific domains but one thread common to many of the best algorithms is the use of the multi-scale approach \\citep{Starcketal2002}. In multi-scale methods, the object to be recovered is modelled as being composed of various different scale sizes. The reconstruction algorithm then has the task of estimating the strengths of the various scales, instead of estimating the strength of pixels. Thus the number of degrees of freedom in the reconstruction can be significantly reduced with a concomittant increase in various measures of performance, such as robustness, stability, and signal to noise. Examples of multi-scale methods are:  \\begin{itemize} \\item {\\bf Multi-Resolution CLEAN}: \\cite{WakkerSchwarz1988} developed a simple strategy for running the CLEAN algorithm emphasizing broad emission first and then finer and finer resolution. The dirty image and point spread function are smoothed and decimated to emphasize the broad emission. The image resulting from CLEANing this dirty image is then used as an initial model for a CLEAN deconvolution of the full resolution image. \\item {\\bf Multi-scale Maximum Entropy}: \\cite{Weir1992}, and \\cite{Bontekoeetal1994} noted that the performance of Maximum Entropy deconvolution could be improved by decomposing the image to be estimated into several channels of different resolutions. A hierarchy of scale sizes is specified and an image reconstructed by estimating pixels in the combined space such that the convolution equation is satisfied. \\item {\\bf Wavelets} Numerous authors have described the virtues of wavelet analysis and its application to deconvolution. The recent textbook by \\cite{Starcketal1998} provides an excellent summary and includes discussion of the connections between Multi-Resolution CLEAN and wavelet analysis. Various authors have described an extension of the Maximum Entropy Method to wavelets as basis functions \\cite{PantinStarck1996,Starcketal2001,Maisingeretal2004}. \\item {\\bf Pixons} \\cite{PuetterPina1993} developed a method for estimating not only pixel strengths but an associated scale size. The combination of strength and scale size, they dubbed a ``pixon''. The Pixon method has been extended considerably, and the original algorithm drastically improved \\citep{Dixonetal1996,PuetterYahil1999}. Performance is extremely good, especially as measured by the statistical whiteness of the residuals. However, there has been no published success in applying the algorithm to synthesis observations because a key assumption, that the PSF is compact, does not hold for Fourier synthesis. \\item {\\bf Adaptive Scale Pixels} \\cite{BhatnagarCornwell2004} developed a method for fitting extended components during a deconvolution process. This has good deconvolution performance but is computationally expensive. \\end{itemize}  Thus multi-scale methods have demonstrated advantages in deconvolution, encouraging further development of multi-scale algorithms. This conclusion is supported by the explosion of interest and technical advances in ``compressive sampling'' \\cite{HauptAndNowak2006a, HauptAndNowak2006b, BobinEtAl2008, BaraniukEtAl2008}. Compressive Sampling (CS) theory shows that under quite general conditions, a sparse signal can be reconstructed from a relatively small number of random projections. This is reassuring for radio astronomers since deconvolution of radio synthesis observations has been the norm for about thirty years.  The recent growth of work in compressive sampling algorithms holds great promise for new and efficient algorithms for deconvolution of radio synthesis observations. However, here we concentrate on a simple extension of the CLEAN algorithm. We describe a conceptually and algorithmically straightforward multi-scale generalization of the CLEAN method that improves convergence and stability (as well as some other properties). This retains the pattern matching motivation of the \\Hogbom\\ algorithm but extends it to encompass extended emission as well as point sources. Unlike the Multi-Resolution CLEAN and Wavelet CLEAN, this algorithm selects among scales considered simultaneously rather than sequentially. We present some background in the next section, our multi-scale algorithm in section 3, some demonstrations and comparisons to other algorithms in section 4, and summarize our work in section 5. ",
        "conclusions": "Multi-Scale CLEAN has been in use for 5-6 years.  It seems to work well in practice. For examples of the application of the algorithm, see \\citep{Subrahmanyanetal2003, Momjianetal2003}.  Besides the obvious efficacy of the algorithm, the major attraction is that it is simple to understand and to implement. The most notable deficiencies are that the scales must be chosen arbitrarily, and an arbitrary bias towards small scales must be introduced. For extended emission, Multi-Scale CLEAN is much faster than \\Hogbom\\ or Clark CLEAN but usually slower than MEM. Multi-Scale CLEAN shows much less bias than MEM at low to moderate signal to noise, and can be applied to images with negative flux."
    },
    "0806/0806.4312_arXiv.txt": {
        "abstract": "We investigate the thermodynamics of the interstellar medium (ISM) and the formation of molecular hydrogen through numerical simulations of spiral galaxies. The model follows the chemical, thermal and dynamical response of the disc to an external spiral potential. Self-gravity and magnetic fields are not included. The calculations demonstrate that gas can cool rapidly when subject to a spiral shock. Molecular clouds in the spiral arms arise through a combination of compression of the ISM by the spiral shock and orbit crowding. These results highlight that local self-gravity is not required to form molecular clouds. Self-shielding provides a sharp transition density, below which gas is essentially atomic, and above which the molecular gas fraction is $>0.001$. The timescale for gas to move between these regimes is very rapid ($\\leq$1 Myr). From this stage, the majority of gas generally takes between 10 to 20 Myr to obtain high H$_{2}$ fractions ($>50$ percent). These are however strict upper limits to the H$_{2}$ formation timescale, since our calculations are unable to resolve turbulent motions on scales smaller than the spiral arm, and do not include self-gravity. True cloud formation timescales are therefore expected to be even shorter.  The mass budget of the disc is dominated by cold gas residing in the spiral arms. Between 50 and 75 percent of this gas is in the atomic phase. When this gas leaves the spiral arm and drops below the self-shielding limit it is heated by the galactic radiation field. Consequently, most of the volume in the interarm regions is filled with warm atomic gas. However, some cold spurs and clumps can survive in interarm regions for periods comparable to the interarm passage timescale. Altogether between 7 and 40 percent of the gas in our disc is molecular, depending on the surface density of the calculation, with approximately 20\\% molecular for a surface density comparable to the solar neighbourhood. ",
        "introduction": "Most stars form in giant molecular clouds (GMCs), and so understanding how GMCs form and disperse is a major goal of star formation research in a Galactic context. GMC formation is likely to be strongly dependent on the nature of the interstellar medium (ISM). Two prominent theories that have previously been advanced are that GMCs form by large scale magnetic and/or gravitational instabilities in thermally or turbulently supported gas \\citep{Elmegreen1979,Balbus1985,Elmegreen1994,Chou2000,Kim2006}, or by the collisional build up of smaller clouds or clumps (e.g. \\citealt{FieldS1965,Levinson1981,Tomisaka1984,Roberts1987,Kwan1987}). The difference between these two scenarios is that in the first the ISM behaves essentially as a fluid, whilst in the latter the behaviour is  more ballistic. The first case resembles the warm, low density component of the ISM, whilst the second appears more consistent with the cold HI or even molecular component which is predominantly situated in clumps. In the majority of galaxies, the ISM may be expected to consist of both of these components, with both instabilities and collisions contributing to GMC formation \\citep{Zhang1999,Dobbs2008}.  A major limitation of many galactic simulations is that the ISM is assumed to be isothermal, either warm \\citep{Kim2002,Slyz2003,Chak2003,Li2005b,Li2006,Shetty2006} cold HI \\citep{Dobbs2006}, or a two phase medium \\citep{Dobbs2007}. In this paper we include a comprehensive treatment of the thermodynamics of the ISM in global simulations, to investigate whether cooling in spiral shocks produces a clumpy, molecular phase even in the absence of self gravity. The only other calculations of heating and cooling in grand design galaxies are presented in the recent paper by \\citet{Wada2008}, who uses an energy equation. Here it is possible to see how the nature of the shock changes with time as the gas changes from a warm smooth distribution to cold clumps. A similar approach has been adopted for simulations without a spiral potential, which show complex thermal and density distributions \\citep{Wada2001,Wada2002,Tasker2006}.  Another key issue related to GMC formation is the determination of the timescale over which these clouds are formed. Although there is no generally accepted definition of the formation timescale, one plausible measure is the time required for the gas to become fully molecular. To address this, simulations ideally need to include the chemistry of H$_2$. The formation of H$_2$ has recently been modelled in simulations of dwarf galaxies \\citep{Pelupessy2006}, but only a very simplistic treatment has been applied to spiral galaxies \\citep{Wada2000,DBP2006}.  Simulations which include a treatment of the thermodynamics of the ISM typically focus on much smaller length scales than those we study here (see, for instance, the recent review by \\citealt{Henne2007}).  Simulations of shock compressed layers show the formation of structure in the cold neutral medium (CNM), and the generation of low levels of turbulence \\citep{Koyama2002,Heitsch2005,Vaz2006,HAudit2007}. Even without self gravity, densities in excess of 1000~cm$^{-3}$ are produced \\citep{Audit2005}. The inclusion of self gravity leads to star formation in the density enhancements from the initial turbulence \\citep{Vaz2007}. In addition to the thermodynamics, \\citet{Glover2007b} also follow the evolution of molecular hydrogen, They find that in a turbulent environment, H$_2$ formation takes place over timescales of a few Myr, suggesting that molecular clouds may form very quickly.  In previous work \\citep{DBP2006}, we showed that molecular clouds can form from cold gas by agglomeration in a spiral shock \\citep{Dobbs2008}. However these isothermal calculations assumed that either cold gas is prevalent in the ISM, or that warm gas cools very quickly in spiral shocks. Here we investigate whether those calculations are plausible, by including cooling and heating of the ISM. We include the thermodynamics described in \\citet{Glover2007a}, which takes into account the main processes responsible for cooling and heating of the ISM, rather than assuming a prescribed relationship between density and temperature. There are essentially two possible outcomes. If cooling is very rapid in the spiral shock, cold clumps are likely to form and the agglomeration of these clumps into GMCs is then a possibility. In this scenario, we are also interested to see whether the gas remains cold in the interarm regions, and whether spurs can be produced. Alternatively, cooling may not be particularly efficient, and little structure is produced in the spiral arms. In this case, another mechanism would be required to form GMCs. ",
        "conclusions": "The simulations presented in this paper indicate that cooling in spiral shocks leads to the formation of molecular clouds in the spiral arms of a galaxy. Self gravity of the gas is not required for this process. Cooling of the gas is very quick as gas enters the shock (almost instantaneous), hence cold gas clumps can agglomerate into larger structures, as described in previous isothermal simulations \\citep{Dobbs2006,DBP2006,Dobbs2008}, within the spiral shock. Larger GMCs form after successive shocks, once both the arm and interarm material has become colder and denser. Again, similar to previous calculations, these clumps are sheared into interarm spurs.  For a disc of surface density comparable to the solar neighbourhood, around 70\\% of the gas lies in the cold regime, and 20\\% is molecular. The transition from atomic to molecular gas is shown to be strongly dependent on the self shielding limit. Spiral shocks compress the gas sufficiently that large amounts of gas exceed the self shielding limit, thus enabling rapid formation of H$_2$. Gas obtains moderate H$_2$ fractions immediately after compression. The maximum H$_2$ fraction obtained then depends on the path of the gas through the shock. Gas which encounters little high density material may obtain a fraction of 0.01 or so for a short period (10 -- 20 Myr) after which photodissociation reduces the H$_2$ fraction. On the other hand gas may encounter dense, cold material, in which case it will undergo a stronger shock, be subject to greater cooling and become predominantly molecular. The time for the transition to fractions of 0.5 or so is typically 10 to 20 Myr, although this is a strict upper limit because our calculations cannot fully resolve internal turbulence and neglect self-gravity of the gas. True cloud formation timescales therefore will be considerably shorter. Without stellar feedback however, this gas tends to remain at high densities, and therefore retain appreciable H$_2$ fractions for substantially longer times ($\\sim40$ Myr). We note the cloud formation time is generally not a well defined quantity, and suggest that determining the time from which gas enters a spiral arm, to the point at which it becomes observable (in CO), may be a more robust measure in future calculations."
    },
    "0806/0806.1701_arXiv.txt": {
        "abstract": "{ The common flare scenario comprises an acceleration site in the corona and particle transport to the chromosphere. Using satellites available to date it has become possible to distinguish between the two processes of acceleration and transport, and study the particle propagation in flare loops in detail, as well as complete comparisons with theoretical predictions.} {We complete a quantitative comparison between flare hard X-ray spectra observed by RHESSI and theoretical predictions. This enables acceleration to be distinguished from transport and the nature of transport effects to be explored. } {Data acquired by the RHESSI satellite were analyzed using full sun spectroscopy as well as imaging spectroscopy methods. Coronal source and footpoint spectra of well observed limb events were analyzed and quantitatively compared to theoretical predictions. New concepts are introduced to existing models to resolve discrepancies between observations and predictions. } {The standard thin-thick target solar flare model cannot explain the observations of all events. In the events presented here, propagation effects in the form of non-collisional energy loss are of importance to explain the observations. We demonstrate that those energy losses can be interpreted in terms of an electric field in the flare loop. One event seems consistent with particle propagation or acceleration in lower than average density in the coronal source.} {We find observational evidence for an electric field in flare loops caused by return currents.} ",
        "introduction": "\\label{Introduction}  Solar flares have been studied in detail both observationally and theoretically ever since their discovery by \\citet{Ca859}. Although increasingly more sophisticated instrumentation provides ever more detailed data, we still lack the basic understanding of many processes at work in a solar flare.  The common flare picture as deduced from hard X-ray (HXR) observations features an HXR source in the corona~\\citep[coronal or loop-top source,][]{Fr71,Hu78}, and two or more HXR sources (footpoints) in the chromosphere \\citep{Hoy81}. These sources are thought to be due to bremsstrahlung emission produced by fast electrons accelerated somewhere above the loop. If we assume that a single particle beam creates both coronal and footpoint emission, the most basic model would involve thin target emission at the top of the coronal loop and thick target emission from the footpoints, which both produce characteristic spectra. \\citet{Wh95} developed a more sophisticated model (intermediate thin-thick target or ITTT model) to fit observations by Yohkoh. They based their model on observations by \\citet{Fe94}, who found high column densities at the loop top, which might act as a thick target below a certain electron energy. In the ITTT model, the shape of the coronal and footpoint non-thermal spectra and the relation between them, observed by Yohkoh, can be explained. The column density in the coronal source determines a critical energy level for the electrons. Electrons that have an energy below this critical energy are stopped in the coronal region. Consequently, the distribution of electron energies measured at the footpoints is depleted in low energy electrons. If the column density is high, the coronal source may act as a thick target to electrons of energies as high as 60 keV, which would leave almost no footpoint emission. Observational evidence for such coronal thick targets were found in RHESSI observations \\citep[eg.][]{Ve04}.  Less extreme cases, flares with one or more footpoints, have frequently been observed by RHESSI. To study the spectral time evolution of individual sources, five well-observed events were analyzed by \\citet{Ba06} who focused on the differences between the spectral indices of coronal and footpoint spectra. They found that in two of those events, the differences at specific times as well as the time-averaged difference was significantly larger than two, ruling out a simple thin-thick target interpretation. In \\citet{Ba07}, the spectra of the five events were compared with the predictions of the ITTT model. The authors exploited the order of magnitude improvement in spectral resolution of RHESSI over the 4-point Yohkoh spectra and showed that most RHESSI observations could not be explained by the ITTT model.  \\citet{Ba07} proposed that by considering non-collisional energy loss inside the loop this inconsistency could be resolved. A possible mechanism that causes non-collisional energy loss is an electric field. Accelerating electrons out of the coronal source region drives a return current to maintain charge neutrality in the whole loop. For finite conductivity, Ohm's law implies that an electric field must be present. The beam electrons lose energy because of work expended in moving inside the electric potential. This produces a change in the shape of the electron spectrum at the footpoints. The formation and evolution of these return currents were studied by various authors \\cite[e.g.][]{Kn77,SS84,La89,Oo90}. \\citet{Zh06} proposed that return currents could explain the high energy break observed in flare HXR spectra. Most studies have, however, been theoretical proposals or numerical simulations, based on standard flare values, that do not attempt to explain or reproduce true solar flare observations.  \\citet{Ba07} compared RHESSI spectra to the ITTT-model of \\citet{Wh95}, demonstrating that the qualitative shape and relations between coronal and footpoint-spectrum often do not agree with the model predictions. In this study, we take an additional step by completing a quantitative analysis of the relation linking coronal and footpoint spectra in the context of the thin-thick target model; we demonstrate that, in some cases, electric fields related to return currents can indeed explain the relation between coronal and footpoint spectra.  In Sect.~\\ref{Theory} we summarize the basic physical concepts applied in the paper. Section~\\ref{evdescription} provides a brief overview of the analyzed events and a description of the spectral analysis. In Sect. \\ref{Method}, we describe our calculation of the energy loss required to reproduce the observed footpoint spectrum, constrained by the coronal emission. Our results are presented in Sect.~\\ref{results}. In Sect. \\ref{retcurrnefield}, we link those results to the concept of return currents. ",
        "conclusions": "\\label{conclusions} The spectral relations between coronal and footpoint HXR-sources provide information about electron transport processes in the coronal loop between the coronal source and the footpoints. Most models neglect these processes in the prediction of the shape and quantitative differences between the source spectra. As shown by \\citet{Ba07}, the observations of some solar flares do not fit the predictions of such models, in particular the intermediate thin-thick target model by \\citet{Wh95}: there is a discrepancy concerning the difference in coronal and footpoint spectral indices, which is expected to be 2.  We have analyzed the two out of five events that display a spectral index difference larger than two in more detail. Such a behavior can be attributed to energy loss during transport that is not proportional to electron energy, but $\\mathcal{E}_{loss}/E$ is larger at low energies. Such an energy loss causes the footpoint spectrum to flatten, which increases the difference in spectral indices. Two loss mechanisms come to mind immediately: Coulomb collisions and an electric potential. Figure 4 demonstrates that the assumption of an electric potential reproduces the observations more accurately.  In one of the two events, there remains a discrepancy between the observed and expected footpoint emission, such that the electron flux at the footpoints is larger than predicted. This flux was estimated from the observed non-thermal HXR (photon) flux and the observed thermal emission of the coronal source. We attribute the discrepancy to propagation or acceleration in low density plasma, which also heats the adjacent high-density regions.  The energy loss can therefore be explained by an electric field in the loop associated to the return current, which builds up as a reaction to the electrons streaming down the loop and the associated beam current. In a steady state ($j_{beam}=j_{ret}$), the return current is unstable to wave growth in one event for all realistic temperature and density parameters in the loop. The kinetic current instability drives a wave turbulence that enhances the electric resistivity by many orders of magnitude. This anomalous resistivity in turn significantly enhances the electric field. In the event of 13 July 2005, the return current may be stable if the loop density is high and the temperature is low, and Spitzer conductivity is applied. Both cases (out of five) present strong evidence for a return current in flares for the first time.  Transport effects by return currents constitute a considerable energy input by Ohmic heating into the loop outside the acceleration region. It may be observable in EUV. Comprehensive MHD modeling including the coronal source, the footpoints, and the region in-between, may be the goal of future theoretical work."
    },
    "0806/0806.4124_arXiv.txt": {
        "abstract": "We present a new, high resolution ($5\\arcsec$ per pixel) \\nir\\ extinction map of the \\nb\\ using data from the UKIDSS-GPS. Using photometry from the $J$, $H$ and $K$-bands we show that the extinction law parameter $\\alpha$ is also highly variable in this region on similar scales to the absolute extinction. We show that only when this extinction law variation is taken into account can the extinction be measured consistently at different wavelengths. ",
        "introduction": "Study of the \\nb\\ (NB) is extremely difficult as it is one of the most highly obscured regions of the Galaxy. The extinction towards the NB is so great that it is almost impossible to observe at visual wavelengths. It is necessary to observe in the \\nir\\ (NIR) to obtain imaging similar to that normally obtained in visual. Previous measures of the NIR extinction towards the NB \\citep{catc90, schu99, dutr03} found the extinction to be highly spatially variable. We present a new mapping of the extinction towards the NB at a much higher resolution than previously possible. ",
        "conclusions": "We have presented a new, highly detailed map of the NIR extinction towards the NB using data from the UKIDSS-GPS. We find both variation in the degree of absolute extinction, and in the extinction law parameter $\\alpha$. We show that the extinction law is not ``universal'' as had previously been thought \\citep{mart90, riek85}, but is highly variable, and that only when this variation is taken into account can absolute extinction be calculated consistently at different wavelengths. We recommend that this variable extinction law be taken into account in all future extinction corrections. We intend to compare our extinction law and absolute extinction maps to dust emission observed with {\\it Spitzer} ({\\it top-right} insert in Figure \\ref{f:extmap}) to better determine the properties of the extincting material and so the distribution of gas and dust in the Bulge."
    },
    "0806/0806.1647_arXiv.txt": {
        "abstract": "{Evidence for a spectral flattening of the supernova remnant (SNR) HB3 (G132.7+1.3) was recently claimed in literature based on previously published total flux density data, and the flattening was further interpreted as the discovery of thermal bremsstrahlung emission in the shell of HB3.} {A spectral flattening has never been observed from any SNR before. Reliable observations of HB3 at frequencies above 3000~MHz are crucial to confirm such a spectral behaviour.} {We extracted 4800~MHz total intensity and polarisation data of HB3 from the Sino-German 6~cm polarisation survey of the Galactic plane made with the Urumqi 25~m telescope, and analysed the spectrum of HB3, together with Effelsberg data at 1408~MHz and 2675~MHz.} {We found an overall spectral index of HB3 of $\\alpha=-0.61$$\\pm$$0.06$ between 1408~MHz and 4800~MHz, similar to the index at lower frequencies. There is no spectral flattening at high frequencies. We detected strong polarised emission from HB3 at 4800 MHz. Our 4800~MHz data show a tangential field orientation in the HB3 shell.} {} ",
        "introduction": "Radio emission from a supernova remnant (SNR) originates from synchrotron emission of relativistic electrons gyrating in the enhanced magnetic fields in its shell or filaments. The observed radio spectrum of a SNR generally follows a power-law ($S_\\nu \\propto \\nu^{\\alpha}$, where $\\alpha$ is the spectral index). A definite case of a spectral break was detected for the SNR S147 \\citep{fr86,xfrh08}.  HB3 is a large ($1\\fdg5\\times2\\degr$) and evolved SNR, with an age of about 30000~yr \\citep{ls06}. It is located at the edge of the HII regions W3/W4 and belongs to a complex of molecular clouds, star-formation and HII regions (W3/W4/W5) \\citep[e.g.,][]{dlp+96} in the Perseus arm at a distance of about 2.2~kpc \\citep{rdlv91,xrzm06}. Although the SNR HB3 partly overlaps with the adjacent HII regions, W3/W4, in the low-resolution radio maps, signs of a real interaction, such as a deformed shell in the high-resolution maps or OH masers in the joint region from the HB3 shock, have not been detected \\citep{kfg+98,kffu06}. Nevertheless, enhanced radio emission has been observed at all radio bands in the southeastern regions of HB3 (towards the lower-left in Fig.~\\ref{fig1}).  Previous estimates of the spectral index of HB3 were made from flux densities determined at frequencies up to 3~GHz, mostly without point source subtraction. The spectral index of HB3 was quoted to have values of $\\alpha=-0.64$$\\pm$$0.01$ \\citep{fdwn95}, $\\alpha=-0.56$$\\pm$$0.03$ \\citep{gre07}, $\\alpha=-0.60$$\\pm$$0.04$ \\citep{lvdr87}, and $\\alpha=-0.66$$\\pm$$0.02$ \\citep{kffu06}. Flux-density estimates of HB3 from low-resolution observations are uncertain, especially at high frequencies, because of the partial overlap with W3/W4 \\citep[see][]{gre07}.   \\begin{figure*} \\centering \\includegraphics[angle=270,width=0.30\\textwidth]{0088fg1a.ps} \\includegraphics[angle=270,width=0.30\\textwidth]{0088fg1b.ps} \\includegraphics[angle=270,width=0.30\\textwidth]{0088fg1c.ps} \\\\[2mm] \\includegraphics[angle=270,width=0.30\\textwidth]{0088fg1d.ps} \\includegraphics[angle=270,width=0.30\\textwidth]{0088fg1e.ps} \\includegraphics[angle=270,width=0.30\\textwidth]{0088fg1f.ps}\\\\ \\caption{HB3 maps at 1408~MHz, 2675~MHz, and 4800~MHz. The top row shows the data after the subtraction of point-like sources in the area of HB3 and $10\\arcmin$ outside. The angular resolution at 1408 MHz, 2675 MHz, and 4800 MHz is $9\\farcm4$, $4\\farcm3$, and $9\\farcm5$, respectively.  The bottom row displays the same data convolved to a common angular resolution of $10\\arcmin$.} \\label{fig1} \\end{figure*}  Recently, \\citet{tl05} obtained images of HB3 at 408~MHz and 1420~MHz from the Canadian Galactic Plane Survey \\citep[CGPS: ][]{tgp+03} with compact sources and diffuse background emission subtracted.  They derived a spectral index for HB3 of $\\alpha = -0.34$$\\pm$$0.25$ between these two frequencies \\citep[][erratum]{tl05,tl06}. This result was confirmed by the TT-plot method.  They also found a surprising result: {\\it a spectral flattening of the overall HB3 spectrum above 1~GHz}, derived from all published flux density data from 38~MHz to 3900~MHz.  \\citet{upl07} fitted the spectrum using a two-component model of synchrotron emission and thermal bremsstrahlung emission. They and concluded that the flattening is evidence for the detection of thermal bremsstrahlung emission produced in the shell of the SNR.  However, there are severe problems in the determination of flux density data of HB3, as discussed by \\citet{gre07}. The uncertainties in the quoted flux density data, especially at high frequencies (1~GHz to 3~GHz), are large, therefore the spectral flattening has been discredited by \\citet{gre07}. Radio observations at even higher frequencies are required to clarify this extraordinary spectral behaviour.  The ongoing {\\it Sino-German $\\lambda$6~cm polarisation survey of the Galactic plane} \\citep{shr+07} has already covered the HB3 area. In this paper, we analyse our $\\lambda$6~cm (4800~MHz) data together with 1408~MHz and 2675~MHz data from the Effelsberg Galactic plane surveys \\citep{rrf97,frrr90} to determine the high-frequency spectrum of HB3. Furthermore, instead of using radio emission of the entire SNR, where its eastern part obviously suffers from contamination or confusion by thermal emission, we limit our study to the part being almost free of distortions. The first detection of polarised emission of HB3 at $\\lambda$6~cm will also be briefly discussed. ",
        "conclusions": "We studied the spectrum of HB3 using maps at 1408~MHz, 2675~MHz, and 4800~MHz for the regions of HB3 where no contamination from W3/W4 is visible.  We obtained a spectral index for HB3 between 1408~MHz and 4800~MHz of $\\alpha = -0.61$$\\pm$$0.06$, which is consistent with values obtained by \\citet{fdwn95}, \\citet{lvdr87}, \\citet{kffu06}. The TT-plot method was used to verify the spectrum. The new 4800~MHz data directly show that the spectrum of HB3 is constant with $\\alpha \\sim -0.6$ from 1~GHz to 5~GHz. There is no evidence for a spectral flattening to higher frequencies discussed earlier by \\citet{tl05}. This conclusion is in entire agreement with \\citet{gre07}. Therefore, the thermal bremsstrahlung emission discussed by \\citet{upl07} does not exist in the HB3 areas not contaminated by thermal emission related to W3/W4.  Polarised radio emission is clearly detected, which indicates that the magnetic field is aligned with the shell of HB3, except for the southern spherical limb, which we suspect to be an independent feature."
    },
    "0806/0806.1471_arXiv.txt": {
        "abstract": "We have searched for solar axions with a detector which consists of a $4\\,{\\rm T}\\times2.3\\,{\\rm m}$ superconducting magnet, PIN-photodiode X-ray detectors, and an altazimuth mount to track the sun. The conversion region is filled with cold helium gas which modifies the axion mass at which coherent conversion occurs. In the past measurements, axion mass from 0 to 0.27eV have been scanned. Since no positive evidence was seen, an upper limit to the axion-photon coupling constant was set to be $\\gagg<6\\hbox{--}10\\times10^{-10}\\rm/GeV$ (95\\% CL) depending on the axion masses. We are now actively preparing for a new stage of the experiment aiming at one to a few eV solar axions. In this mass region, our detector might be able to check parameter regions which are preferable to the axion models. ",
        "introduction": "Axions \\cite{pq1977} that are thermally produced in the core of the sun through the Primakoff process (Fig. \\ref{fig:principle}) are called the solar axions. Sikivie \\cite{sikivie1983} proposed an ingenious experiment to detect such axions. The detection device called axion helioscope is a system of a strong magnet and an X-ray detector, where the solar axions is converted into X-ray photons through the inverse Primakoff process (Fig. \\ref{fig:principle}) in the magnetic field. Conversion is coherently enhanced even for massive axions by filling the conversion region with light gas. If the axion mass $m_a$ is at around a few eV, detection of the solar axions becomes feasible.  We have constructed an axion helioscope with a dedicated superconducting magnet. Cold helium gas was used as the conversion medium, and the highest axion mass was targeted at 2.6\\,eV. We report the current status of this experiment. \\begin{figure} \\begin{minipage}[b]{0.45\\textwidth} \\includegraphics[width=\\textwidth]{fig1.eps} \\caption{The solar axions produced via the Primakoff process in the solar core are, then, converted into X-rays via the reverse process in the magnet.} \\label{fig:principle} \\end{minipage} \\hskip0.1\\textwidth plus 1fil minus 5pt \\begin{minipage}[b]{0.45\\textwidth} \\includegraphics[width=\\textwidth]{fig2.eps} \\caption{The schematic view of the axion helioscope.} \\label{fig:sumico} \\end{minipage} \\end{figure} ",
        "conclusions": "We have searched for solar axions with an axion helioscope, and axion mass from 0 to 0.27eV have been scanned. No evidence of axion signal has been seen so far, implying an exclusion limit to the axion-photon coupling as shown in Fig.\\ \\ref{fig:limit}. In the lower mass region, the above limit has been renewed by CAST. However, we are preparing for the next stage aiming at a higher mass region where our detector might be able to check the parameter regions preferred by theory. In December 2007, a preliminary measurement started at $m_a\\simeq1\\rm\\,eV$. Remaing problems are being resolved one by one. The ultimate density will be reached in 2008, although a full range scan will take years. \\ack The authors thank the ex-director general of KEK, Professor H. Sugawara, for his support in the beginning of the experiment. The authors are also thankful to the support from the Grant-in-Aid for COE research by the Japanese Ministry of Education, Science, Sports and Culture, and the support from the Matsuo Foundation."
    },
    "0806/0806.3474_arXiv.txt": {
        "abstract": "We develop \\textit{information field theory} (IFT) as a means of Bayesian inference on spatially distributed signals, the information fields.  A didactical approach is attempted. Starting from general considerations on the nature of measurements, signals, noise, and their relation to a physical reality, we derive the information Hamiltonian, the source field, propagator, and interaction terms. Free IFT reproduces the well known Wiener-filter theory. Interacting IFT can be diagrammatically expanded, for which we provide the Feynman rules in position-, Fourier-, and spherical harmonics space, and the Boltzmann-Shannon information measure. The theory should be applicable in many fields. However, here, two cosmological signal recovery problems are discussed in their IFT-formulation. 1) Reconstruction of the cosmic large-scale structure matter distribution from discrete galaxy counts in incomplete galaxy surveys within a simple model of galaxy formation. We show that a Gaussian signal, which should resemble the initial density perturbations of the Universe, observed with a strongly non-linear, incomplete and Poissonian-noise affected response, as the processes of structure and galaxy formation and observations provide, can be reconstructed thanks to the virtue of a response-renormalization flow equation. 2) We design a filter to detect local non-linearities in the cosmic microwave background, which are predicted from some Early-Universe inflationary scenarios, and expected due to measurement imperfections. This filter is the optimal Bayes' estimator up to linear order in the non-linearity parameter and can be used even to construct sky maps of non-linearities in the data. ",
        "introduction": "\\subsection{Motivation}  The optimal extraction and restoration of information from data on spatially distributed quantities like the cosmic \\textit{large-scale structure} (LSS) or the \\textit{cosmic microwave background} (CMB) temperature fluctuations in cosmology, but also on many other signals in physics and related fields, is essential for any quantitative, data-driven scientific  inference. The problem of how to design such methods possesses many technical and even conceptual difficulties, which have led to a large number of recipes and methodologies.  Here, we address such problems from a strictly information theoretical point of view. We show, as others have done before, that information theory for distributed quantities leads to a statistical field theory, which we name \\textit{information field theory} (IFT). In contrast to the previous works, which mostly treat such problems on a classical field level, as will be detailed later, here, we take full advantage of the existing field theoretical apparatus to treat interacting and non-classical fields. Thus, we show how to use diagrammatic perturbation theory and renormalization flows in order to construct optimal signal recovering algorithms and to calculate moments of their uncertainties. Non-classicality manifests itself as quantum and statistical fluctuations in quantum and statistical field theory (QFT \\& SFT), and very similarly as  uncertainty in IFT.   The information theoretical perspective on signal inference problems has technical advantages, since it permits to design information-yield optimized algorithms and experimental setups. However, it also provides deeper insight into the mechanisms of knowledge accumulation, its underlying information flows, and its dependence on data models, prior knowledge and assumptions than pure empirical evaluations of ad-hoc algorithms alone could provide.   We therefore hope that our work is of interest for two types of readers. The first are applied scientists, who are mainly interested in the practical aspect of IFT since they are facing  a concrete inverse problem for a spatially distributed quantity, especially but not exclusively in cosmology. The second are more philosophical or theoretically inclined scientists, for whom IFT may serve as a framework to understand and classify many of the existing methods of signal extraction and reception. Since we expect that many interested readers are not very familiar with field theoretical formalisms, we introduce some of its basic mathematical concepts. Due to this anticipated non-uniform readership, not everything in this article might be of everyones interest, and therefore we provide in the following a short overview on the structure and content of the article.    \\subsection{Overview of the work}  The remainder of this introduction section contains a detailed discussion of the previous work on signal inference theory as well as a very brief introduction into the here relevant works on the cosmic LSS and the CMB. The main part of this article falls into two categories: abstract IFT and its application. The concepts of IFT are introduced in Sec. \\ref{sec:concepts}, where Bayesian methodology, the distinction of physical and information fields, the definition of signal response and noise, as well the design of signal spaces are discussed. The basic IFT formalism including the free theory is introduced in Sec. \\ref{sec:basics}, which, according to our judgement, summarizes and unifies the previous knowledge on IFT before this paper. An impatient reader, only interested in applying IFT and not worrying about concepts, may start reading in Sec. \\ref{sec:basics}. From Sec. \\ref{sec:interacting fields} on the new results of this work are presented, starting with the discussion of interacting information fields, their Hamiltonians and Feynman rules, and the Boltzmann-Shannon information measure. The normalisability of sensibly constructed IFTs is shown, as well the classical information field equation is presented there. A step-by-step recipe of how to derive and implement an IFT algorithms is also provided.   Details of the notation can be found, if not defined in the main text, in Appendix \\ref{sec:notation}.  Applications of the theory are provided in the following two sections, which can be skipped by a reader interested only in the general theoretical framework. Although specific inference problems are addressed, they should serve as a blueprint for the tackling of similar problems. In Sec. \\ref{sec:poisson} the problem of the reconstruction of the cosmic matter distribution from galaxy surveys is analyzed in terms of a Poissonain data model. In Sec. \\ref{sec:fnl} we derive an optimal estimator for non-Gaussianity in the CMB, and show how it can be generalized to map potential non-Gaussianities in the CMB sky. Our summary and outlook can be found in Sec. \\ref{sec:summary}.     \\subsection{Previous works}\\label{sec:lit}  The work presented here tries to unify information theory and statistical field theory in order to provide a conceptual framework in which optimal tools for cosmological signal analysis can be derived, as well as for inference problems in other disciplines. Below, we provide very brief introductions into each of the required fields\\footnote{This work has tremendously benefitted in a direct and indirect way from a large number of previous publications in those fields. We, the authors, have to apologize for being unable to give full credit to all relevant former works in those fields for only concentrating on a brief summary of the papers more or less directly influencing this work. This collection is obviously highly biased towards the cosmological literature due to our main scientific interests and expertise, and definitely incomplete.} (information theory, image reconstruction, statistical field theory, cosmological large-scale structure, and cosmic microwave background), for the orientation of non-expert readers. An expert in any of these fields might decide to skip the corresponding sections.   \\subsubsection{Information theory and Bayesian inference}  The fundament of information theory was laid by the work of \\citet{Bayes} on probability theory, in which the celebrated Bayes theorem was presented. The theorem itself (see Eq. \\ref{eq:Bayes}) is a simple rule for conditional probabilities. It only unfolds its power for inference problems if used with belief or knowledge states, described by conditional probabilities.  The advent of modern information theory is probably best dated by the work of Shannon \\citep{Shannon1948,1949mtc..book.....S} % on the concept of information measure, being the negative Boltzmann-entropy, and the work of Jaynes, combining the language of statistical mechanics and Bayes probability theory and applying it to knowledge uncertainties \\citep{1957PhRv..106..620J, 1957PhRv..108..171J, jaynes1963, 1965AmJPh..33..391J, jaynes1968, 1982ieee...70..939J, 2004PhT....57j..76J}. The required numerical evaluation of Bayesian probability integrals suffered often from the curse of high dimensionality. The standard recipe against this, still in massive use today, is importance sampling via Markov-Chain Monte-Carlo Methods (MCMC), following the ideas of \\citet{metroplis}, % \\citet{hastings}, and % \\citet{gibbsamp}, % where the latter authors already had image reconstruction applications in mind. The Hamiltonian MCMC methods \\citep{HMCMC}, in which the phase-space sampling is partly following Hamiltonian dynamics, are also of relevance here. There the Hamiltonian is introduced as the negative logarithm of the probability, as we do in this work.  With such tools, higher dimensional problems, as present in signal restoration, could and can be tackled, however, for the price of getting stochastic uncertainty into the computational results. For a recent review on image restoration MCMC techniques, see \\cite{2005cs........4037M}.  The applications and extensions of these pioneering works are too numerous to be listed here. Good monographs exist and the necessary references can be found there \\switch{% : \\cite{toolsstatinf}, % \\cite{neal1993}, % \\cite{bayeschoice}, % \\cite{bayesdataanal}, and % \\cite{paramest}. % }{% \\cite{toolsstatinf, % neal1993, % bayeschoice, % bayesdataanal, % paramest, % 2008arXiv0803.4089T}. % }   \\subsubsection{Image reconstruction in astronomy and elsewhere}\\label{sec:lit:map}  The problem of image reconstruction from incomplete, noisy data is especially important in astronomy, where the experimental conditions are largely set by the nature of distant objects, weather conditions, etc., all mainly out of the control of the observer, as well as in other disciplines like medicine and geology, with similar limitations to arrange the object of observations for an optimal measurement.  Some of the most prominent methods of image reconstruction, which are based on a Bayesian implementation of an assumed data model, are the Wiener-filter \\citep{1949wiener}, % the Richardson-Lucy algorithm \\citep{1972JOSA...62...55R, 1974AJ.....79..745L}, and the maximum-entropy image restoration \\citep{1972JOSA...62..511F}\\switch{% \\footnote{See also \\cite{1972JOSA...62..511F}, % \\cite{1978Natur.272..686G}, % \\cite{1979MNRAS.187..145S}, % \\cite{1980MNRAS.191...69B}, % \\cite{1983CVGIP..23..113B}, % \\cite{1983iimp.conf..267G}, % \\cite{1984Natur.311..446S}, % \\cite{1984Natur.312..381T}, % \\cite{1984MNRAS.211..111S}, % \\cite{1986JMOp...33..287B}, % \\cite{gull1989}, % \\cite{gullskilling}, and % \\cite{1998mebm.conf....1S}.% }}{(see also \\cite{% 1978Natur.272..686G, % 1979MNRAS.187..145S, % 1980MNRAS.191...69B, % 1983CVGIP..23..113B, % 1983iimp.conf..267G, % 1984Natur.311..446S, % 1984Natur.312..381T, % 1984MNRAS.211..111S, % 1986JMOp...33..287B, % gull1989, % gullskilling,  % 1998mebm.conf....1S}).% }  The Wiener filter can be regarded to be a full Bayesian image inference method in case of Gaussian signal and noise statistics, as we will show in Sect. \\ref{sec:free}. It will be the working horse of the IFT formalism, since the Wiener filter represents the algorithm to construct the exact field theoretical expectation value given the data for an interaction-free information Hamiltonian. The filter can be decomposed into two essential information processing steps, first building the information source by response-over-noise weighting the data, and then propagating this information through the signal space, by applying the so called Wiener variance.  The Richardson-Lucy algorithm is a maximum-likelihood method to reconstruct from Poissonian data and therefore is also of Bayesian origin. This method has usually to be regularized by hand, by truncation of the iterative calculations, against an over-fitting instability due to the missing (or implicitly flat) signal prior. A Gaussian-prior based regularization was recently proposed by \\citet{2008MNRAS.389..497K}, and the implementation of a variant of this is presented here in Sect. \\ref{sec:classicalPoisson}.  Maximum entropy algorithms will not be the topic here, as well as not a number of other existing methods, which are partly within and partly outside the Bayesian framework. They may be found in existing reviews on this topic \\citep[e.g.][]{1986ARA&A..24..127N, imagerest}.   \\subsubsection{Statistical and Bayesian field theory}  The relation of signal reconstruction problems and field theory was discovered independently by several authors. In cosmology, a prominent work in this directions is \\citet{1987ApJ...323L.103B}, % in which the path integral approach was proposed to sample primordial density perturbations with a Gaussian statistics under the constraint of existing information on the large scale structure. The work presented here can be regarded as a non-linear, non-Gaussian extension of this. Many methods from statistics and from statistical mechanics were of course used even earlier, e.g. the usage of moment generating function for cosmic density fields can already be found in \\citet{1985ApJ...289...10F}.  Simultaneously to  Bertschinger's work, \\citet{1987PhRvL..58..741B, 1988PhRvL..61.1512B} % argued that visual perception can be modeled as a field theory for the true image, being distorted by noise and other data transformations, which are summarized by a nuisance field. A probabilistic language was used, but no direct reference to information theory was made, since not the optimal information reconstruction was the aim, but a model for the human visual reception system. However, this work actually triggered our research.  \\citet{1996PhRvL..77.4693B} % applied a field theoretical approach to recover a probability distribution from data. Here, a Bayesian prior was used to regularize the solution, which was set up ad-hoc to enforce smoothness of the reconstruction, obtained from the classical (or saddlepoint, or maximum a posteriori) solution of the problem. However, an ``optimal'' value for the smoothness controlling parameter was derived from the data itself, a topic also addressed by \\citet{2000AJ....120.2163S} and by a follow up publication to ours \\citep{PURE}. \\citet{1996PhRvL..77.4693B} also recognized, as we do, that an IFT can easily be non-local.  Finally, the work of Lemm and coworkers\\switch{% \\footnote{\\citet{1998cond.mat..8039L, 1999physics..11077L, 2000FBS....29...25L, 2000quant.ph..2010L, 2000PhRvL..84.2068L, 2000PhRvL..84.4517L, 2000PhLA..276...19L, 2001AIPC..568..425L, 2001EPJB...20..349L, 2003quant.ph.12191L, 2005EPJB...46...41L}} }{ \\cite{1999physics..12005L, 2000FBS....29...25L, 2000PhRvL..84.2068L, 2000PhRvL..84.4517L, 2000PhLA..276...19L, 2001AIPC..568..425L, 2001EPJB...20..349L, 2005EPJB...46...41L}} established a tight connection between statistical field theory and Bayesian inference, and proposed the term \\textit{Bayesian field theory} (BFT) for this. However, we prefer the term {\\it information field theory} since it puts the emphasis on the relevant object, the information, whereas BFT refers to a method, Bayesian inference. The term \\textit{information field} is rather self-explaining, whereas the meaning of a  \\textit{Bayesian field} is not that obvious.  The applications considered by Lemm concentrate on the reconstruction of probability fields over parameter spaces and quantum mechanical potentials by means of the maximum a posteriori equation. The extensive book summarizing the essential insights of these papers, \\cite{1999physics..12005L}, clearly states the possibility of perturbative expansions of the field theory. However, this is not followed up by these authors probably for reasons of the computational complexity of the required algorithms. In contrast to many of the previous works on IFT, which deal with ad-hoc priors, the publication by \\citet{1998cond.mat..8039L} % is remarkable, since it provides explicit recipes of how to implement a priori information in various circumstances more rigorously.  The mathematical tools required to tackle IFT problems come from SFT and QFT, which have a vast literature. We have specially made use of the books of \\citet{Binney1992}, \\citet{PeskinSchroeder}, and \\citet{2003qftn.book.....Z}.  \\subsubsection{Cosmological large-scale structure}\\label{sec:lit:lss}  Our first IFT example in Sec. \\ref{sec:poisson} is geared towards improving galaxy-survey based cosmography, the reconstruction of the large-scale structure matter distribution. We provide here a short overview on the relevant background and works.  The LSS of the matter distribution of the Universe is traced by the spatial distribution of Galaxies, and therefore well observable. This structure is believed to have emerged from tiny, mostly Gaussian initial density fluctuations of a relative strength of $10^{-5}$ via a self-gravitational instability, partly counteracted by the expansion of the Universe. The initial density fluctuations are believed to be produced during an early inflationary epoch of the Universe, and to carry valuable information about the inflaton, the field which drove inflation, in their $N$-point correlation functions, to be extracted from the observational data.  The onset of the structure formation process is well described by linear perturbation theory and therefore to conserve Gaussianity, however, the later evolution, the structures on smaller scales, and especially the galaxy formation require non-linear descriptions. The observational situation is complicated by the fact that the most important galaxy distance indicator, their redshift, is also sensitive to the galaxy peculiar velocity, which causes the observational data on the three-dimensional LSS to be partially degenerated. There are analytical methods to describe these effects% \\switch{% \\footnote{ Of special interest in this context may be the following works, and the papers they refer to: \\cite{1970A&A.....5...84Z}, % \\cite{1986ApJ...304...15B}, % \\cite{1986ApJ...310L..21M}, which already applies path-integrals,  % \\cite{1980lssu.book.....P}, % \\cite{1987MNRAS.227....1K}, % \\cite{1990ApJ...362....1P}, % \\cite{1992ApJ...390L..61B}, % \\cite{1996ApJ...462...25Z}, % \\cite{hamilton-1998}, % \\cite{1999MNRAS.309..543B}, % \\cite{1999MNRAS.308....1B}, % \\cite{1999ApJ...520...24D}, % \\cite{2002astro.ph..6052Z}, % \\cite{2003MNRAS.341.1311S}, % and \\cite{2004PhRvD..70h3007S}. % }}{% \\footnote{% Of special interest in this context may be \\cite{1986ApJ...310L..21M}, which already applies path-integrals,   % % % % \\cite{1970A&A.....5...84Z, % 1986ApJ...304...15B, % 1980lssu.book.....P, % 1987MNRAS.227....1K, % 1990ApJ...362....1P, % 1992ApJ...390L..61B, % 1996ApJ...462...25Z, % hamilton-1998, % 1999MNRAS.309..543B, % 1999MNRAS.308....1B, % 1999ApJ...520...24D, % 2002astro.ph..6052Z, % 2003MNRAS.341.1311S, % 2004PhRvD..70h3007S}, % % and the papers they refer to.}}, and also extensive work on $N$-body simulations of the structure formation, the latter probably providing us with the most detailed and accurate statistical data on the properties of the matter density field \\citep[e.g][]{2005Natur.435..629S}.    In recent years, it was recognized that the evolution of the cosmic density field and its statistical properties can be addressed with field theoretical methods by virtue of renormalization flow equations. Detailed semi-analytical calculations for the density field time propagator, the two- and three- point correlation functions are now possible due to this, which are expected to play an important role in future approaches to reconstruct the initial fluctuations from the observational data\\switch{% \\footnote{Relevant references are \\cite{2004A&A...421...23V, 2007A&A...476...31V, 2008A&A...484...79V}, \\cite{2006PhRvD..73f3519C}, % \\cite{2006PhRvD..73f3520C}, % \\cite{2006PhRvD..74j3512M, 2006PhRvD..74l9901M, 2007PhRvD..75d3514M}, \\cite{2006ApJ...651..619J, 2008arXiv0805.2632J}, % \\cite{2007JCAP...06..026M}, % \\cite{2007JPhA...40.6849G}, % \\cite{2008MPLA...23...25M}, % \\cite{2008PhRvD..77f3530M}, and \\cite{2008arXiv0806.0971P}. % }}{ \\cite{2004A&A...421...23V, 2007A&A...476...31V, 2008A&A...484...79V, 2006PhRvD..73f3519C, % 2006PhRvD..73f3520C, % 2006PhRvD..74j3512M, 2006PhRvD..74l9901M, 2007PhRvD..75d3514M, 2006ApJ...651..619J, 2008arXiv0805.2632J, % 2007JCAP...06..026M, % 2007JPhA...40.6849G, % 2008MPLA...23...25M, % 2008PhRvD..77f3530M, 2008arXiv0806.0971P}} % .  It was recognized early on, that the primordial density fluctuations can in principle be reconstructed from galaxy observations \\citep{1987ApJ...323L.103B}. This has lead to a large development of various numerical techniques for an optimal reconstruction\\switch{% \\footnote{Some of the main references are: \\cite{1989ApJ...336L...5B, 1991ASPC...15...67B}, % \\cite{1989ApJ...344L..53P}, % \\cite{1990ApJ...364..349D}, % \\cite{1991ASPC...15..111K}, % \\cite{1991ApJ...380L...5H}, % \\cite{1992MNRAS.254..315W}, % \\cite{1992ApJ...391..443N}, % \\cite{1992ApJ...398..169R}, % \\cite{1993ApJ...405..449G}, % \\cite{1993ApJ...415L...5G}, % \\cite{1994A&A...291..697B}, % \\cite{1994ApJ...421L...1N}, % \\cite{1994ASPC...67..171L}, % \\cite{1994ApJ...423L..93L}, % \\cite{1995MNRAS.272..885F}, % \\cite{1995MNRAS.277..933S}, % \\cite{1995ApJ...449..446Z}, % \\cite{1995ApJ...453..533T}, % \\cite{1997MNRAS.285..793C}, % \\cite{1998ApJ...508..440N}, % \\cite{1998ApJ...504..601P}, % \\cite{1998ApJ...503..492S, 1998ApJ...506...64S}, % \\cite{1998ApJ...492..439B}, % \\cite{1999MNRAS.306..491T}, % \\cite{1999ApJ...515..471N}, % \\cite{1999ApJ...520..413Z}, % \\cite{2000ASPC..201..282G}, % \\cite{2000ApJ...544...21G}, % \\cite{2000MNRAS.316..464K}, % \\cite{2001ApJ...550..522B}, % \\cite{2001ApJ...552..413G}, % \\cite{2002Natur.417..260F}, % \\cite{2002MNRAS.331..901Z}, % \\cite{2003MNRAS.346..501B}, % \\cite{2003A&A...406..393M}, % \\cite{2004astro.ph.10063M}, % \\cite{2004MNRAS.349..425B}, % \\cite{2005ApJ...635L.113M}, % and \\cite{2006MNRAS.365..939M}. % }}{ \\cite{1989ApJ...336L...5B, 1991ASPC...15...67B, % 1989ApJ...344L..53P, % 1990ApJ...364..349D, % 1991ASPC...15..111K, % 1991ApJ...380L...5H, % 1992MNRAS.254..315W, % 1992ApJ...391..443N, % 1992ApJ...398..169R, % 1993ApJ...405..449G, % 1993ApJ...415L...5G, % 1994A&A...291..697B, % 1994ApJ...421L...1N, % 1994ASPC...67..171L, % 1994ApJ...423L..93L, % 1995MNRAS.272..885F, % 1995MNRAS.277..933S, % 1995ApJ...449..446Z, % 1995ApJ...453..533T, % 1997MNRAS.285..793C, % 1998ApJ...508..440N, % 1998ApJ...504..601P, % 1998ApJ...503..492S, 1998ApJ...506...64S, % 1998ApJ...492..439B, % 1999MNRAS.306..491T, % 1999ApJ...515..471N, % 1999ApJ...520..413Z, % 2000ASPC..201..282G, % 2000ApJ...544...21G, % 2000MNRAS.316..464K, % 2001ApJ...550..522B, % 2001ApJ...552..413G, % 2002Natur.417..260F, % 2002MNRAS.331..901Z, % 2003MNRAS.346..501B, % 2003A&A...406..393M, % 2004astro.ph.10063M, % 2004MNRAS.349..425B, % 2005ApJ...635L.113M, % 2006MNRAS.365..939M}}% . Many of them are based on a Bayesian approach, since they are implementations and extension of the Wiener filter. However, also other principles are used, like, e.g. the least action approach, or Voronoi tessellation techniques\\switch{% \\footnote{ E.g. \\cite{1991QJRAS..32...85I}, % \\cite{1991MNRAS.250..519I}, % \\cite{1996MNRAS.279..693B}, % \\cite{2000A&A...363L..29S}, % \\cite{2001A&A...368..776R}, % and \\cite{2006ChJAA...6..387Z}. % }.}{ \\cite[e.g.][]{1991QJRAS..32...85I, % 1991MNRAS.250..519I, % 1996MNRAS.279..693B, % 2000A&A...363L..29S, % 2001misk.conf..268V, % 2001A&A...368..776R, % 2006ChJAA...6..387Z}.} % A discussion and classification of the various methods can be found in \\cite{2008MNRAS.389..497K}.  Especially the Wiener filter methods were extensively applied to galaxy survey data\\switch{% \\footnote{Survey based reconstructions of the cosmic matter fields can be found in \\cite{1990ApJ...364..370B}, % \\cite{1991ApJ...372..380Y}, % \\cite{WienerFSL}, % \\cite{1995ApJ...454...15S}, % \\cite{1996ApJ...461L..17B}, % \\cite{1997MNRAS.287..425W}, % \\cite{1997ApJ...474..553Y}, % \\cite{1999AJ....118.1146S}, % \\cite{1999MNRAS.303..179N}, % \\cite{tegmark-1999-518}, % \\cite{2000ApJ...535L...5H}, % \\cite{2001ApJ...550...87G}, % \\cite{2002MNRAS.333..739M}, % \\cite{2004MNRAS.352..939E}, % \\cite{2004ogci.conf....5V}, % \\cite{2005ASPC..329..135H}, % \\cite{2005MNRAS.356.1168P}, % and \\cite{2006MNRAS.373...45E}.} % % }{% \\footnote{Survey based reconstructions of the cosmic matter fields can be found in \\cite{1990ApJ...364..370B, % 1991ApJ...372..380Y, % WienerFSL, % 1995ApJ...454...15S, % 1996ApJ...461L..17B, % 1997MNRAS.287..425W, % 1997ApJ...474..553Y, % 1999AJ....118.1146S, % 1999MNRAS.303..179N, % tegmark-1999-518, % 2000ApJ...535L...5H, % 2001ApJ...550...87G, % 2002MNRAS.333..739M, % 2004MNRAS.352..939E, % 2004ogci.conf....5V, % 2005ASPC..329..135H, % 2005MNRAS.356.1168P, % 2006MNRAS.373...45E}.} % % } and permitted partly to extrapolate the matter distribution into the \\textit{zone of avoidance} behind the galactic disk\\switch{\\footnote{ \\cite{1994ASPC...67..185H}, % \\cite{2000ASPC..218..173Z}, % and \\cite{2000A&ARv..10..211K}.} % and to close the data-gap there, a topic we also address in Sect. \\ref{sec:poisson}. } { and to close the data-gap there, c.f. \\cite{1994ASPC...67..185H, % 2000ASPC..218..173Z, % 2000A&ARv..10..211K}% , a topic we also address in Sect. \\ref{sec:poisson}. }   Another cosmological relevant information field to be extracted from galaxy catalogues is the LSS power spectrum\\switch{% \\footnote{\\citep[E.g.][]{1994MNRAS.267.1020P, 1996ApJ...465...34V, 1997ApJ...486...21Z, 1997PhRvL..79.3806T, 1999ApJ...511....5E}.}}{ \\cite[e.g.][]{1994MNRAS.267.1020P, 1996ApJ...465...34V, 1997ApJ...486...21Z, 1997PhRvL..79.3806T, 1999ApJ...511....5E}% }. This power is also measurable in the CMB, and for a long time the CMB provided the best spectrum normalization \\citep{1992MNRAS.258P...1E, 1995ApJ...441L...9B}.   \\subsubsection{Cosmic Microwave Background}\\label{sec:lit:cmb}  Since our second example deals with the CMB, we give a brief overview on it and on related inference methods.  The CMB reveals the statistical properties of the matter field at a time, when the Universe was about 1100 times smaller in linear size than it is today. The photon-baryon fluid, which decouples at that epoch into neutral Hydrogen and free streaming photons, has responded to the gravitational pull of the then already forming dark matter structures. The photons from that epoch cooled due to the cosmic expansion since then into the CMB radiation we observe today, and carry information on the physical properties of the photon-baryon fluid of that time like density, temperature and velocity. To very high accuracy, the spectrum of the photons from any direction is that of a blackbody, with a mean temperature of $2.7$ Kelvin and fluctuations of the order of $10^{-5}$ Kelvin, imprinted by the primordial gravitational potentials at decoupling.  Therefore, mapping these temperature fluctuations permits precisely to study many cosmological parameters simultaneously, like the amount of dark matter producing the gravitational potentials, the ratio of photons to baryons, balancing the pressure and weight of the fluid, and geometrical and dynamical parameters of space-time itself. The observations are technically challenging, and therefore require sophisticated algorithms to extract the tiny signal of temperature fluctuations against the instrument noise, but also to separate it from other astrophysical foreground emission with the best possible accuracy.  A number of such algorithms were developed\\switch{% \\footnote{E.g.by \\cite{1992issa.proc..391J}, % \\cite{1994ApJ...432L..75B}, % \\cite{1997MNRAS.290..313M}, % \\cite{1997PhRvD..56.4514T}, % \\cite{1997ApJ...480L..87T}, % \\cite{1997ApJ...482..577D}, % \\cite{1998MNRAS.300....1H}, % \\cite{natoli}, % \\cite{2001A&A...374..358D}, % \\cite{maxima}, % \\cite{2004PhRvD..70h3511W}, % \\cite{2004ApJS..155..227E}, % \\cite{2004ApJ...609....1J}, % \\cite{mirage}, % \\cite{madam}, % \\cite{2006ApJS..162..401S}, % \\cite{2007ApJ...656..653L}, % and \\cite{Hinshaw:2008kr}.} % }{ \\cite[e.g.][]{1992issa.proc..391J, % 1994ApJ...432L..75B, % 1997MNRAS.290..313M, % 1997PhRvD..56.4514T, % 1997ApJ...480L..87T, % 1997ApJ...482..577D, % 1998MNRAS.300....1H, % natoli, % 2001A&A...374..358D, % maxima, % 2004PhRvD..70h3511W, % 2004ApJS..155..227E, % 2004ApJ...609....1J, % mirage, % madam, % 2006ApJS..162..401S, % 2007ApJ...656..653L, % Hinshaw:2008kr}}% , which in many cases implement the Wiener filter. Thus, the required numerical tools for an IFT treatment of CMB data are essentially available.  The expected temperature fluctuations spectrum can be calculated from a linear perturbative treatment of the Boltzmann equations of all dynamical active particle species at this epoch, and fast computational implementations exists permitting to predict it for a given set of cosmological parameters. Well known codes for this task are publicly available\\footnote{E.g. \\texttt{cmbfast} (\\switch{\\texttt{http://cmbfast.org, http://ascl.net/cmbfast.html}}{\\url{http://cmbfast.org}, \\url{http://ascl.net/cmbfast.html}}, \\cite{1996ApJ...469..437S}), \\texttt{camb} (\\switch{\\texttt{http://camb.info/}}{\\url{http://camb.info/}}, \\cite{Lewis:1999bs}), and \\texttt{cmbeasy} (\\switch{\\texttt{http://www.cmbeasy.org/}}{\\url{http://www.cmbeasy.org/}}, \\cite{2005JCAP...10..011D}).} and permit to extract information on cosmological parameters from the measured CMB temperature fluctuation spectrum via comparison to their predictions for a given parameter set. It was recognized early on that this should happen in an information theoretically optimal way, and Bayesian methods were therefore adapted in that area well before in other astrophysical disciplines \\citep[e.g.][]{1995ApJ...446...49B, 1997ApJ...480...22T, 1997PhRvD..55.5895T,Nolta:2008ih}.  The initial metric and density fluctuations, from which the CMB fluctuations and the LSS emerged, are believed to be initially seeded by quantum fluctuations of a hypothetical inflaton field, which should have driven an inflationary expansion phase in the very early Universe \\citep{1981PhRvD..23..347G, 1982PhLB..108..389L, 1982PhRvL..48.1220A, 1982PhRvL..49.1110G, 1982PhLB..117..175S, 1983PhRvD..28..679B}. The inflaton-induced fluctuations have a very Gaussian probability distribution, however, some non-Gaussian features seem to be unavoidable in most scenarios and can serve as a fingerprint to discriminate among them \\citep[e.g.][]{2001PhRvD..64h3005H, 2002PhRvD..66j3506B, 2004PhR...402..103B, 2004JCAP...08..009B}. Observational tests on such non-Gaussianities based on the three-point correlation function of the CMB data \\citep[e.g.][]{2002ApJ...566...19K, 2004PhRvD..70h3005B, 2005ApJ...634...14K, 2007ApJ...664..680Y, 2008ApJ...678..578Y} were so far mostly negative, however not sensitive enough to seriously constrain the possible theoretical parameter space of inflationary scenarios, see e.g. \\cite{2003ApJS..148..119K, 2008arXiv0804.0136C}. % Recently, there has been the claim of a detection of such non-Gaussianities by \\citet{2008PhRvL.100r1301Y} and a confirmation of this with better data and improved algorithms is therefore highly desirable. In Sect. \\ref{sec:fnl} we make a proposal for improving the algorithmic side of this challenge. A recent review on the current status of CMB-Gaussianity can be found in \\cite{2008arXiv0805.4157M}. ",
        "conclusions": "\\label{sec:summary}  Starting with fundamental information theoretical considerations about the nature of measurements, signals, noise and their relation to a physical reality given a model of the Universe or the system under consideration, we reformulated the inference problem in the language of  \\textit{information field theory} (IFT). IFT is actually a statistical field theory. The information field is identified with a spatially distributed signal, which can freely be chosen by the scientist according to needs and technical constraints. The mathematical apparatus of field theory permits to deal with the ensemble of all possible field configurations given the data and prior information in a consistent way.  With this conceptual framework, we derived the Hamiltonian of the theory, showed that the free theory reproduces the well known results of Wiener-filter theory, and presented the Feynman-rules for non-linear, interacting Hamiltonians in general, and in particular cases. The latter are information fields over Fourier- and spherical harmonics-spaces for inference problems in $R^n$ and $S^2$, respectively. Our ``philosophical'' considerations permitted to argue why the resulting IFTs are usually well normalized, but often non-local. Since the propagator of the theory is closely related to the Wiener-filter, for which nowadays efficient numerical algorithms exist as image reconstruction and map-making codes, and the information source term is usually a noise weighted version of the data, the necessary computational tools are at hand to convert the diagrammatic expressions into well performing algorithms.  Furthermore, we provided the Boltzmann-Shannon information measure of IFT based on the Helmholtz free energy, thereby highlighting the embedding of IFT in the framework of statistical mechanics.  As examples of the IFT recipe, two concrete IFT problems with cosmological motivation were discussed, which are also thought as blueprints for other inference problems. The first was targeting at the problem of reconstructing the spatially continuous cosmic LSS matter distribution from discrete galaxy counts in incomplete galaxy surveys. The resulting algorithm can also be used for image reconstruction with low-number photon statistics, e.g in low-dose X-ray imaging.  The second example was the design of an optimal method to measure or constrain any possible local non-linearities in the CMB temperature fluctuations. This may serve as a blueprint for statistical monitoring of the linearity of a signal amplifier.  We conclude here with a short outlook on some problems that are accessible to the presented theory.  Many signal inference problems involve the reconstruction of fields without precisely known statistics.  Some coefficients in the IFT-Hamiltonians may only be phenomenological in nature, and therefore have to be derived from the same data used for the reconstruction itself. This more intricate interplay of parameter and information field can also be incorporated into the IFT framework, as we will show with a subsequent work.  For cosmological applications, along the lines started in this work, clearly more realistic data models need to be investigated. For example, to understand the response in galaxy formation to the underlying dark matter distribution in terms of a realistic, statistical model, to be used in constructing the corresponding IFT Hamiltonian for a dark-matter information field, detailed higher-order correlation coefficients have to be distilled from numerical simulations or semi-analytic descriptions. Also the CMB Hamiltonian may benefit from the inclusion of remnants from the CMB foreground subtraction process, permitting to gather more solid evidence on fundamental parameters which are hidden in the CMB fluctuations, like the amplitude of non-Gaussianities.  Furthermore, there exist a number of more or less heuristic algorithms for inverse problems, which have proven to serve well under certain circumstances. Reverse engineering of their implicitly assumed priors and data models may permit to understand better for which conditions they are best suited, as well how to improve them in case these conditions are not exactly met.  Finally, we are very curious to see whether and how the presented framework may be suitable to inference problems in other scientific fields."
    },
    "0806/0806.3538_arXiv.txt": {
        "abstract": "Gott has promulgated a rule for making probabilistic predictions of the future duration of a phenomenon based on the phenomenon's present age [{\\sl Nature\\/} {\\bf 363}, 315 (1993)].  I show that the two usual methods for deriving Gott's rule are flawed.  Nothing licenses indiscriminate use of Gott's rule as a predictor of future duration. It should only be used when the phenomenon in question has no identifiable time scales. ",
        "introduction": "In an article$^1$ published in {\\sl Nature\\/} in 1993 and in subsequent publications$^{\\hbox{\\scriptsize2--5}}$ and the concluding chapter of a book,$^6$ J.~Richard Gott~III has promulgated a formula for making probabilistic predictions of the future duration of a phenomenon based on the phenomenon's present age.  When you observe that a phenomenon has lasted a time $\\tp$, Gott instructs you to predict that the phenomenon will last an additional time $\\tf\\ge Y\\tp$ with probability \\begin{equation} G(\\tf\\ge Y\\tp)={1\\over1+Y}\\;. \\label{eq:Grule} \\end{equation} For example, Gott's rule predicts that a phenomenon has a probability of $1/2$ to survive an additional time at least as long as its past ($Y=1$) and, by the same token, a probability of $1/2$ to end before reaching twice its present age.  In applying his rule to a host of phenomena, Gott usually couches his predictions in terms of a particular 95\\% confidence interval, 95\\% confidence being his standard for a scientific prediction.  According to his rule, the probability that a phenomenon's future duration, $\\tf$, will be between $1/39$ and $39$ times its present age, $\\tp$, is $G(\\tf\\ge\\tp/39)-G(\\tf\\ge39\\tp)=39/40-1/40=0.95$. The flip side of this prediction is that the phenomenon has a 2.5\\% chance to end before reaching 1/39 of its present age and a 2.5\\% chance of lasting longer than 39 times its present age.  Gott bases his formula on a temporal version of the Copernican principle: when you observe the phenomenon's present age, your observation does not occur at a special time.  Here I show, distilling the essence of a previous critical analysis$^7$ of Gott's work, that although the Copernican principle does lead directly to a version of Gott's rule, this version is essentially meaningless and, in particular, does not authorize his predictions for future longevity based on present age.  In published papers and his book, Gott is on record as applying his rule to himself,$^{2,6}$ Christ\\-ianity,$^6$ the former Soviet Union,$^{1,6}$ the Third Reich,$^6$ the United States,$^6$ Canada,$^4$ world leaders,$^{2,4,6}$ Stonehenge,$^4$ the Seven Wonders of the Ancient World,$^6$ the Pantheon,$^6$ the Great Wall of China,$^6$ {\\sl Nature},$^1$ the {\\sl Wall Street Journal},$^6$ {\\sl The New York Times},$^6$ the Berlin Wall,$^{\\hbox{\\scriptsize1--4,6}}$ the Astronomical Society of the Pacific,$^2$ the 44 Broadway and off-Broadway plays open and running on 27~May 1993,$^{\\hbox{\\scriptsize2--4,6}}$ the Thatcher-Major Conservative government in the UK,$^{\\hbox{\\scriptsize2--4,6}}$ Manhattan (New York City),$^6$ the New York Stock Exchange,$^6$ Oxford University,$^6$ the internet,$^6$ Microsoft,$^6$ General Motors,$^6$ the human spaceflight program,$^{\\hbox{\\scriptsize1--6}}$ and {\\it homo sapiens}.$^{\\hbox{\\scriptsize1--6}}$ In all these cases---even the New York plays---Gott uses his rule to make probabilistic predictions for the survival of individual phenomena whose present age is known. For example, given {\\sl Nature\\/}'s 123 years of publishing in 1993, Gott predicted that {\\sl Nature\\/} had a 95\\% chance to continue publishing for a period between 3.15 years (already exceeded) and 4,800 years.$^1$  Most notably, Gott has used the 200,000-year present age of {\\it homo sapiens} to predict that we have a 95\\% chance to go extinct sometime between 5,100 years and 7.8 million years from now.$^{\\hbox{\\scriptsize1--6}}$  Although Gott issues occasional cautionary statements about the applicability of his rule,$^{2,6}$ the list of phenomena to which he has applied the rule indicates that these cautions don't cramp his style much.  Gott's predictions have received attention in the popular media, including a favorable piece by Timothy Ferris in {\\sl The New Yorker},$^8$ which highlighted Gott's predictions for human survival (and which motivated me to write my original paper$^7$), and a recent article by John Tierney in {\\sl The New York Times},$^9$ which focused on the implications for space colonization.  My late 1999 posting of the paper$^7$ that eventually appeared in {\\sl Contemporary Physics\\/} prompted a sympathetic article in {\\sl The New York Times\\/} by James Glanz,$^{10}$ then a science writer at {\\sl The Times}.  Glanz, as a long-suffering fan of the Chicago White Sox, was particularly interested in Gott's prediction, issued in 1996, for the Sox's World Series prospects.  Gott opined that the Sox, having not won a World Series title since 1917, would, with 95\\% confidence, win a Series sometime between 1999 and 5077.  In his 2000 article, Glanz noted that the Sox hadn't yet succeeded, but he was clearly dismayed by the long wait evoked by the mere mention of 5077.  Happily for him and other Sox fans, they did win the Series in 2005.  In 1996 Gott would have predicted a World Series title in 2005 or before with probability~0.10, considerably less than the probability, $1-(29/30)^9=0.26$, that comes from assuming that the Sox had the same chance each year as the other 30 major-league ball clubs.  Gott has given two main derivations of his rule: the argument from the Copernican principle, which he calls the delta-$t$ argument,$^{\\hbox{\\scriptsize1--6}}$ and a Bayesian analysis,$^{2,6,11}$ which he adopted from criticism due to Buch.$^{12}$  Both of these derivations are flawed.$^7$  Here I begin in Sec.~\\ref{sec:deltat} with an analysis of the delta-$t$ argument, because it led Gott to his rule and because he consistently portrays it as the chief justification for his predictions.  I show that the delta-$t$ argument does not lead to any prediction of future duration based on present age.  I then turn in Sec.~\\ref{sec:Bayesian} to the usual Bayesian derivation of Gott's rule, which has greater appeal to most other contributors to the literature on the subject.  I demonstrate that this derivation is simply wrong and sketch the correct Bayesian analysis.  A concluding Sec.~\\ref{sec:conclusion} considers the assumptions that are required to get Gott's predictions.  It should be emphasized at the outset that we will not conclude that Gott's rule is ``wrong,'' but rather that its two primary derivations are wrong.  In science flawed justifications are as bad as---perhaps worse than---being obviously wrong, because they are more pernicious. They can mislead you into using methods that don't apply in your situation and can get you into trouble when you export those methods to other contexts.  Determining the assumptions that underlie whatever you are doing in science is essential, so that you know when to abandon what you are doing in favor of something else.  The purpose of this article is thus to debunk the two primary derivations of Gott's rule and to identify the assumptions that underlie Gott's rule in its predictive form, so that you will know what you are doing should you choose to use it.  The discussion in this paper is couched mainly in terms of a simple graphical representation, which is equivalent to the more formal, Bayesian analysis given in Ref.~7.  Section~\\ref{sec:deltat} on the delta-$t$ argument is phrased almost entirely in terms of the graphical representation.  The results of the Bayesian analysis in Sec.~\\ref{sec:Bayesian} can be understood by referring to the graphical representation, but the Bayesian equations are included for those who prefer to see the details.  The evidence from papers$^{\\hbox{\\scriptsize13--17}}$ that cite Ref.~7 is that its argument and conclusions have not been appreciated and understood. The goal of the present paper, with its graphical mode of presentation, is to rectify that situation. ",
        "conclusions": ""
    },
    "0806/0806.0820_arXiv.txt": {
        "abstract": "{}{We analyze different data of the variation of the fine structure constant obtained with different methods to check their consistency.}{We test consistency using the modified student test and confidence intervals. We split the data sets into smaller intervals.  A criterion for this selection is proposed.}{Results show consistency for reduced intervals for each pair of data sets considered.}{Results are at variance with the ones obtained considering mean values over the whole interval.} ",
        "introduction": "\\label{Intro}  The attempt to unify all fundamental interactions resulted in the development of multidimensional theories like  string-motivated field theories \\citep{Wu86,Maeda88,Barr88,DP94,DPV2002a,DPV2002b}, related brane-world theories \\citep{Youm2001a,Youm2001b,branes03a,branes03b}, and (related or not) Kaluza-Klein theories \\citep{Kaluza,Klein,Weinberg83,GT85,OW97}.  Among these theories, there are some in which the gauge coupling constants may vary over cosmological timescales. On the other hand, theoretical frameworks based on first principles, were developed by different authors \\citep{Bekenstein82,Bekenstein2002,CLV01,BSM02,OP02,BM05} to study the variation of the fine structure constant ($\\alpha$) or the electron mass ($m_e$).  Different versions of the theories mentioned above predict different time behaviors of the fundamental constants. Thus, bounds obtained from astronomical and geophysical data are an important tool to test the validity of these theories.  The experimental research can be grouped into astronomical and local methods. The latter ones include geophysical methods such as the natural nuclear reactor that operated about $1.8\\ 10^9$ years ago in Oklo, Gabon \\citep{DD96,Fujii00}, the analysis of natural long-lived $\\beta$ decayers in geological minerals and meteorites \\citep{Olive04b} and laboratory measurements such as comparisons of rates between clocks with different atomic numbers \\citep{PTM95,Sortais00,Marion03,Bize03,Fischer04,Peik04}. The astronomical methods are based mainly on the analysis of high-redshift quasar absorption systems. The relative magnitude of the fine splitting of resonance lines of alkaline ions is proportional to $\\alpha^2 $. \\citet{Murphy01b} and \\citet{Chand05} have applied this method to SiIV doublet absorption lines systems at different redshifts. An extension of this method was proposed by \\citet{Bahcall04}. These authors use  strong nebular emission lines of O III to constrain the variation of $\\alpha$. Furthermore, this method was improved   by comparing transitions of different species with widely far atomic masses and led to the first results consistent with a time varying fine structure constant for a range of redshifts ($0.5 < z < 3.5$) \\citep{Webb99,Webb01,Murphy03b}. This method is known in the literature as the many multiplet method (MM). However, other recent independent analyses of similar data \\citep{Srianand04,Chand04,Chand06,QRL04} found no variation. On the other hand, the standard MM technique can be revised to avoid the deficiencies pointed out earlier in the literature \\citep{Bahcall04,QRL04}. This improved method called in the literature as revised many multiplet (RMM) method, was applied by \\citet{QRL04} to a homogeneous sample of FeII lines at redshift $z=1.15$.  Another method, to test cosmological variation of $\\alpha$,  from pairs of Fe II lines observed in individual exposures from a high-resolution spectrograph was proposed by \\citet{levshakov05} (this method is known in the literature as SIDAM). The authors found no variation of $\\alpha$ at $z=1.84$ and $z=1.15$ \\citep{Levshakov06}. However, a recent reanalysis of spectrum of the quasar Q1101-264 found variability within $1\\sigma$ \\citep{Levshakov07}. Besides, by comparing  optical and radio redshifts, a bound on $\\alpha ^2g_p\\frac{me}{mp}$ (where $g_p$ is the proton $g$ factor) can be obtained \\citep{CS95,Tzana07}. Furthermore, comparing molecular and radio lines provides a bound on $g_p \\alpha^2$ and the most stringent constraints were obtained by \\citet{Murphy02}. On the other hand, \\citet{darling04} reports bounds on the variation of $\\alpha$ at $z=0.2467$ from the satellite $18$ cm $OH$ conjugate lines. Finally, \\citet{kanekar05} compared the HI and OH main line absorption redshifts of the different components in the $z=0.765$ absorber and the $z=0.685$ lens toward B0218+357 to establish stringent constraints on changes in $F= g_p \\left({\\frac{\\alpha^2 m_e}{m_p}}\\right)^{1.57}$. Besides, the time variation of the gauge coupling constants in the early universe can be constrained using data from the cosmic microwave background (CMB) \\citep{Martins02,Rocha03} and the primordial abundances of light elements \\citep{Ichi02,Nollet}.  In this paper, we would like to focus on the discrepancy between the data on time variation of $\\alpha$ from astronomical observations, following a purely statistical criterion. As described above, different methods are able to constrain the variation of $\\alpha$. However, to achieve statistical consistency for each fixed redshift interval, the reported value should be the same. The usual assumption is that if there is any time variation of $\\alpha$, such variation is constant over all observed redshifts. So, the reported results are means over a range of redshifts. However, theories mentioned above predict different time evolutions of $\\alpha$ yielding different variations for different times \\citep{Okada,Marciano,Bekenstein82}. Therefore, to test the evolution of $\\alpha$ predicted by those kind of theories, not only a mean value over a range of redshifts, but each individual measurement should be considered. However, not all the reported results are consistent. The most important discrepancy in the literature are the results reported by \\citet{Murphy03b} and those reported by \\citet{Chand04}.   It should be also mentioned that there is an important debate in the literature about the reliability of the \\citet{Chand04} results \\citep{Murphy07,Srianand07,Murphy08}. The aim of this paper is to test the consistency of different reported values for the variation of $\\alpha$ for smaller intervals than the reported ones, in each case. In this way, we will be able to define a reduced interval where consistency can be assured and thus the theoretical prediction for the evolution of $\\alpha$ with time within this interval can be tested. The selection of the redshift interval to be tested is not trivial and a method for this selection is also proposed and described in Sects. \\ref{tools} and \\ref{resultados}.   It is also important to note that observational errors are important and can not be ignored in any statistical analysis.    In Sect. \\ref{tools} we describe the statistical tools (which include the observational errors) we use for testing consistency. In Sect. \\ref{resultados}, we apply these tools  to check consistency between the data on varying $\\alpha$. Not all groups of data can be tested using the modified student test because the requirements of the method for selecting the redshift intervals can not be fulfilled by most groups of data.  Therefore, we compute confidence intervals of a group of data and compare them with a single reported value of another author to check consistency in most of the cases. We find that consistency between pairs of data sets  can be assured for a reduced interval in each case.  In Sect. \\ref{discusion}, we discuss our results and write our conclusions. ",
        "conclusions": "\\label{discusion}  In Sect. \\ref{resultados}  we have obtained two consistency redshift intervals (($0.45, 1.14$) and ($2.03, 2.3$)), when comparing the data sets of \\citet{Murphy03b} and \\citet{Chand04}.  Now, we apply the modified student test defined in Sect. \\ref{student} to the two consistency intervals. We obtain $\\lambda^*=0.28$ and $\\lambda^*=0.42$ respectively, which confirms the consistency of these two redshift intervals as analyzed using confidence intervals. In this case, we do not have to worry about the number of samples being equal, since the selection of the redshift interval is already done. It is also interesting to repeat the calculation of confidence intervals for different values of $\\beta$ (or $n$). Table \\ref{table13} shows that  for a higher value of $\\beta$, the consistency interval is larger  than the interval obtained for a lower value of $\\beta$.   We also performed the modified student test for the consistency intervals of table \\ref{table13} and obtain $\\lambda^* > 0.25$. We have also calculated the consistency intervals for the other groups of data for larger values of $\\beta$. We find that there is consistency between the confidence interval and the reported interval for $\\beta=0.05$ and $\\beta=0.1$ in all cases but the comparison of the data set from \\citet{Murphy03b} and \\citet{Tzana07}. For this case, the redshift consistency interval for $\\beta=0.05$ is $(0.238,1.78)(2.17,2.47)$ and for $\\beta=0.1$ we obtain $(0.238,1.79)(2.14,2.47)$. On the other hand, we would like to comment on the comparison between the \\citet{Chand05} and \\citet{Murphy01b} data sets. In this case, both the modified student test and the calculation of confidence intervals were used to analyze consistency with different results. However, it should be stressed that the modified student test is able to test redshift intervals larger than $0.76$, while the length of the confidence intervals lies between $0.26$ and $0.53$.  A   statistical criterion has been given to estimate the appropriate  sample sizes (which determine the width of the redshift intervals). The criterion is based on the probability of not rejecting the null hypothesis (consistency of two experiments compared) when a definite alternative hypothesis is supposed to be true. While being well justified on statistical grounds,  this proposal  involves a subtle point: The choice of a given alternative hypothesis is not provided by a theoretical framework. Rather, our choice is suggested by the results of \\citet{Murphy03b}, which  motivated our analysis. In that work, the non null result for the variation of $\\alpha$ was supported by a departure from the null hypothesis  $\\delta\\sim 4.5 \\, \\sigma = 4.5 \\,S/\\sqrt{n}$ ($S$ the sample variance); our choice for the alternative hypothesis $\\delta\\sim S$ is consistent with this for the employed sample sizes.  We  thus have analyzed the consistency between data sets on varying $\\alpha$, using the statistical tools described in Sect. \\ref{tools}. Usually, results are taken to be means over a range of redshifts. Here, instead, we have analyzed smaller intervals. The criterion for the selection of intervals is based on limits on the probability of type II error (see section \\ref{size}). We have obtained two consistency intervals for the data set of \\citet{Murphy03b} and \\citet{Chand04}, while the mean values and errors of the whole interval show no consistency.  We have also shown that the data of \\citet{Tzana07} and \\citet{Murphy03b} are not consistent when tested over small redshift intervals. We have also obtained consistency intervals for this case, which should be used when testing theories on varying $\\alpha$. For the other data sets, the consistency depends on the desired value of $\\beta$. The result for $\\beta=0.025$ is also at variance with the one obtained calculating the mean values and errors over the whole interval.   \\begin{table}[!ht] \\renewcommand{\\arraystretch}{1.3} \\caption{Consistency intervals obtained for different values of $\\beta$ ($\\lambda=0.025$ for all cases).} \\label{table13} \\begin{center} \\begin{tabular}{ccccc} \\hline \\hline $\\beta$  & $n$ & Consistency Intervals & $n_1$ & $n_2$   \\\\ \\hline $0.1$ &$12$& $(0.45,1.17)(1.86,2.30)$  &$12$ &$69$\\\\ $0.05$ &$15$& $(0.45,1.14) (1.90,2.30)$ &$12$ &$67$\\\\ $0.025$ &$18$& $(0.45,1.14) (1.92,2.30)$ &$11$ &$67$ \\\\ \\hline \\multicolumn{5}{l}{Notes.- $n_1$ is the number of data reported by  }\\\\ \\multicolumn{5}{l}{\\citet{Chand04} that are consistent with $n_2$ data }\\\\ \\multicolumn{5}{l}{from \\citet{Murphy03b}.} \\end{tabular} \\end{center} \\end{table}"
    },
    "0806/0806.3890_arXiv.txt": {
        "abstract": "A class of dynamical dark energy models, dubbed Spinor Quintom, can be constructed by a spinor field $\\psi$ with a nontraditional potential. We find that, if choosing suitable potential, this model is able to allow the equation-of-state to cross the cosmological constant boundary without introducing any ghost fields. In a further investigation, we show that this model is able to mimic a perfect fluid of Chaplygin gas with $p=-c/\\rho$ during the evolution, and also realizes the Quintom scenario with its equation-of-state across $-1$. ",
        "introduction": "The recent data from type Ia supernovae and cosmic microwave background (CMB) radiation and so on \\cite{Perlmutter:1996ds, Riess:2004nr, Seljak:2004xh} have provided strong evidences for a spatially flat and accelerated expanding universe at the present time. In the context of Friedmann-Robertson-Walker (FRW) cosmology, this acceleration is attributed to the domination of a component with negative pressure, called dark energy. So far, The nature of dark energy remains a mystery. Theoretically, the simplest candidate for such a component is a small positive cosmological constant, but it suffers the difficulties associated with the fine tuning and the coincidence problems. So many physicists are attracted by the idea of dynamical dark energy models, such as Quintessence \\cite{Ratra:1987rm}, Phantom \\cite{Caldwell:1999ew}, K-essence \\cite{ArmendarizPicon:2000dh}, Quintom \\cite{Feng:2004ad} and so on (see Refs. \\cite{Copeland:2006wr, Padmanabhan:2007xy} for a review). Usually, dark energy models are constructed by scalar fields which are able to accommodate a rich variety of behaviors phenomenologically. However, there is another possibility that the acceleration of the universe is driven by a classical homogeneous spinor field. Some earlier studies on applications of spinor fields in cosmology are given in Refs. \\cite{taub:1937, Brill:1957fx, Parker:1971pt}. In recent years there are many works on studying spinor fields as gravitational sources in cosmology, for example: see Refs. \\cite{Obukhov:1993fd, ArmendarizPicon:2003qk} for inflation and cyclic universe driven by spinor fields; see Refs. \\cite{Saha:2000nk} for spinor matter in Bianchi Type I spacetime; see Refs. \\cite{Ribas:2005vr, Chimento:2007fx} for a dark energy model with spinor matter; and so on.  Although the recent fits to the data in combination of the 3-year WMAP \\cite{Spergel:2006hy}, the recently released 182 SNIa Gold sample \\cite{Riess:2006fw} and also other cosmological observational data show remarkably the consistence of the cosmological constant, it is worth noting that a class of dynamical models with the equation-of-state (EoS) across $-1$ {\\it Quintom} is mildly favored \\cite{Zhao:2006qg, Wang:2006ts}. In the literature there have been a lot of theoretical studies of Quintom-like models. For example, motivated from string theory, the authors of Ref. \\cite{Cai:2007gs} realized a Quintom scenario by considering the non-perturbative effects of a generalized DBI action. Moreover, a No-Go theorem has been proven to constrain the model building of Quintom \\cite{Xia:2007km} (see Ref. \\cite{Vikman:2004dc} for some earlier considerations), and according to this No-Go theorem there are models which involve higher derivative terms for a single scalar field \\cite{Li:2005fm}, models with vector field \\cite{ArmendarizPicon:2004pm}, making use of an extended theory of gravity \\cite{Cai:2005ie}, non-local string field theory \\cite{Aref'eva:2005fu}, and others (see e.g. \\cite{Cai:2006dm, Quintom_1, Xiong:2007cn}). The similar work applied in scalar-tensor theory has also been studied in Ref. \\cite{Elizalde:2004mq, Boisseau:2000pr}.  Usually, a Quintom model involves a ghost field with its kinetic term to be negative which leads to quantum instability. In this paper we study the dark energy model with spinor matter. Interestingly, we find that this type of model can realize the Quintom scenario with its EoS across the cosmological constant boundary $w=-1$ without introducing a ghost field. Instead, the derivative of its potential with respect to the scalar bilinear $\\bar\\psi\\psi$, which is defined as the mass term, becomes negative when the spinor field lies in the Phantom-like phase. If this model can realize its EoS across $-1$ more than one time, the total EoS of the universe can satisfy $w\\geq -1$ during the whole evolution which is required by the Null Energy Condition (NEC) \\cite{Qiu:2007fd}, and we expect to treat this process as a phase transition merely existing for a short while. Moreover, due to a perfect mathematic property of the spinor field, it is possible to combine the Quintom scenario and the picture of Chaplygin gas with the EoS evolving from $0$ to $-1$ smoothly in Spinor Quintom. In the literature a dark energy model of Chaplygin gas has been proposed to describe a transition from a universe filled with dust-like matter to an accelerated expanding stage, and hence it has been argued that the coincidence problem of dark energy may be alleviated in this model \\cite{Kamenshchik:2001cp}.  Interestingly, in our model we are able to evade the drawbacks of considering the Phantom field, of which the kinetic energy is negative and so is unstable in quantum level. We notice that Ref. \\cite{Kahya:2006hc} has investigated the quantum stability of a Phantom phase of cosmic acceleration and shown that a super-acceleration phase can be obtained by quantum effects. Our model is different from that one since the super-acceleration is realized by the background contribution. However, it is interesting that both the two approaches are stable in quantum level since the first order of perturbation theory can be defined.  This paper is organized as follows. In Section II, we simply review the basic algebra of a spinor field in FRW universe which is minimally coupled to Einstein's gravity. In Section III, we present the condition for the spinor field to realize a Quintom scenario and give some detailed examples. In Section IV, we provide a unified model of Spinor Quintom and a perfect fluid of Chaplygin gas by taking certain potentials. By solving the model numerically, we will study the evolution of its EoS and fraction of energy density. Section V is the conclusion and discussions of our paper. ",
        "conclusions": "The current cosmological observations indicate the possibility that the acceleration of the universe is driven by dark energy with EoS across $-1$, which will challenge the theoretical model building of the dark energy if confirmed further in the future. In this paper we have studied various Quintom scenarios in virtue of a spinor field and proposed a unified model of Spinor Quintom and a generalized Chaplygin gas. As shown in the present work, this model can give rise to the EoS crossing the cosmological constant boundary during the evolution by varying the sign of the term $V'$. Compared with other models with $w$ across $-1$ in the literature, so far the present one is also economical in the sense that it merely involves a single spinor field."
    },
    "0806/0806.1547_arXiv.txt": {
        "abstract": "A proxy for horizontal electric currents in the solar photosphere was suggested. For a set of evolving active regions (ARs) observed with {\\it Solar and Heliospheric Observatory (SOHO)} Michelson Doppler Imager (MDI) in the high resolution mode, the dissipation spectrum, $k^2E(k)$, and the spatial structure of dissipation, i.e., the Stokes dissipation function $\\varepsilon(x,y)$, were calculated from  the observed $B_z$ component of the magnetic field. These functions allowed us to calculate (a part of) the horizontal electric current density in the photosphere. It was shown that as an active region emerges, large-scale horizontal electric currents are gradually generated and determine a bulk of dissipation. When an active region decays, the large-scale horizontal currents decay faster than the small-scale ones. The density of horizontal currents in active regions is in the range of $\\langle j_h \\rangle \\sim (0.008 - 0.028)$ A/m$^2$, that is compatible with the density of vertical currents in active regions. We suggest two possible mechanisms for generation of such horizontal currents in the photosphere. One of them is the drift motions of charged particles in the medium of varying plasma pressure gradient in a horizontal plane at the periphery of a sunspot. Such a drift can produce quasi-circular closed horizontal currents around sunspots. Another possibility could be an existence of horizontal axial current inside a highly twisted horizontal magnetic structure laying in the photosphere along the magnetic neutral line. The horizontal currents may contribute significantly to the dynamics of the photosphere/corona coupling, as well as the estimation of non-potentiality of ARs. ",
        "introduction": "\\subsection {Spatial structures of $B_z$ dissipation}  Because the $B_z$ component is transported in the photosphere in the same way as a passive scalar does (Parker 1979, Petrovay \\& Szakaly 1993), one may apply the Stokes dissipation function (Monin \\& Yaglom 1975) to obtain the spatial structure of dissipation:  \\begin{equation} \\varepsilon(x,y) = \\frac{1}{2} \\left ( \\left ( \\frac {{\\it d}B_z}{{\\it d}x} \\right ) ^2 + \\left ( \\frac {{\\it d}B_z}{{\\it d}y} \\right ) ^2 \\right ), \\label{eps} \\end{equation} where $(x,y)$ is a current point on the horizontal plane and $z$ represents direction normal to the solar surface. In general, the right hand part of Eq. (\\ref{eps}) requires the magnetic diffusion coefficient, $\\eta$, which is unknown. Assuming the quasi-uniform behavior of $\\eta$ over the active region area, we will analyze the spatial distribution of dissipation as it is derived from Eq. (\\ref{eps}) and measured in $Gs^2/Mm^2$ units.  A relationship between the dissipation $\\varepsilon$ and the squared electric current, $\\bf j^2$, can be easily derived from the general formula for the squared curl of the magnetic field vector, ${\\bf B} \\equiv (B_x,B_y,B_z)$: \\begin{eqnarray} \\mu^2 {\\bf j}^2 =(\\nabla \\times {\\bf B})^2  = & \\nonumber \\\\ & \\left ( \\frac {{\\it d}B_z}{{\\it d}y} -  \\frac {{\\it d}B_y}{{\\it d}z} \\right ) ^2 + \\left ( \\frac {{\\it d}B_x}{{\\it d}z} -  \\frac {{\\it d}B_z}{{\\it d}x} \\right ) ^2 + \\left ( \\frac {{\\it d}B_y}{{\\it d}x} -  \\frac {{\\it d}B_x}{{\\it d}y} \\right ) ^2 =   \\label{j2} \\\\ & \\varepsilon(x,y)  + \\mu^2j_z^2 + f(B_x,B_y,{\\it d/dz}). \\nonumber \\end{eqnarray} Thus, dissipation $\\varepsilon(x,y)$ represents a part of the horizontal electric current. The spatial structure of $\\varepsilon(x,y)$ is shown in Figures \\ref{Fig0}, \\ref{Fig2}, and \\ref{Fig3}. Areas of strong dissipation, and, therefore, strong horizontal currents, tend to be concentrated around strong magnetic elements. Below we discuss the distribution of horizontal currents with possible mechanisms for their formation.   \\begin{figure}[!h] \\centerline {\\epsfxsize=4.0truein \\epsffile{Fig_0.eps}} \\caption{\\sf MDI/HR magnetograms ({\\it left}) and corresponding dissipation structures, $\\varepsilon(x,y)$, calculated for NOAA AR 0365 observed on 2003 May 26/01:01 UT. North is to the top, west is to the right. The image size is $ 156 \\times 126$ arcsec. The magnetogram is scaled from -500 G to 500G. The dissipation map is scaled from 0 to $ 10^5$ G$^2$/Mm$^2$.} \\label{Fig0} \\end{figure}  Averaging $\\varepsilon(x,y)$ over the magnetogram area, we obtain the mean value of dissipation $\\langle \\varepsilon \\rangle$ associated with $B_z$ component.  Having the value of $\\langle \\varepsilon \\rangle$ for an active region, we can estimate (a part of) the horizontal electric current by using Eq. (\\ref{j2}): \\begin{equation} \\langle {\\bf j}_h^2\\rangle =\\frac{1}{\\mu^2}\\langle \\varepsilon \\rangle. \\label{jtr} \\end{equation} Here $\\mu=\\mu_0=4\\pi 10^{-7}$ H/m is the vacuum permeability.   \\subsection {Spectrum of dissipation of $B_z$}  For any measured quantity, which characterizes the photospheric magnetized plasma, one can calculate the spectrum of energy, i.e., a function that shows how the energy (per unit volume) is distributed in the wavenumber space, $k=2\\pi/s$, where $s$ is a spatial scale. The energy spectrum, $E(k)$ (often referred also as a power spectrum), is essentially a squared Fourier transform of a measured array (here we consider a one-dimensional spectrum of a two-dimensional array, see Abramenko (2005) for details of calculation). For a turbulent medium, the energy spectrum displays a maximum at large scales (see Figure \\ref{Fig1}, blue curve) followed by a power-law decay at smaller scales. The wavenumber where the maximum is reached, marked as $k_e$ ($e$ stands for {\\it energy}), refers to the scales which bear the maximum energy and where the energy input occurs (in our case, this is the largest magnetic elements on a magnetogram).  The spectrum of {\\it dissipation} can be derived from the energy spectrum (Monin \\& Yaglom 1975) as \\begin{equation} W(k) = k^2E(k). \\label{W} \\end{equation} Usually, for fully developed turbulence at high magnetic Reynolds number (the ratio of the inertial force to the viscous force), the magnetic dissipation becomes significant at high wavenumbers only. As a result, the maximum of the dissipation spectrum, $k_d$ (see Figure \\ref{Fig1}), is located at higher wavenumbers than the maximum, $k_e$, of the energy spectrum. The energy interval (which corresponds to the bulk of energy) and the dissipation interval (which corresponds to the bulk of dissipation) are separated in the wavenumber space: the energy is stored mostly at large spatial scales and dissipated at much smaller scales. The energy cascade is formed from large to small scales, which manifests the presence of the fully developed turbulent state. Turbulent magnetic diffusion seems to be a key mechanism for dissipation in this case.  \\begin{figure}[!h] \\centerline {\\epsfxsize=5.0truein \\epsffile{Fig_1.eps}} \\caption{\\sf MDI/HR magnetograms ({\\it left}) and corresponding energy spectrum, $E(k)$ ({\\it blue curve and right axis}) and dissipation spectrum, $W(k)=k^2E(k)$ ({\\it red curve, left axis}), calculated for a coronal hole area recorded on 2001 June 5/13:07 UT. Noise influence on the spectra was eliminated. North is to the top, west is to the right. The image size is $ 266 \\times 202$ arcsec. The magnetograms are scaled from -150 G to 150G. The arrows $k_e$  and $k_d$ mark the maxima of the $E(k)$ and $W(k)$ spectra, respectively.  The locations of the maxima were derived from 5 point box car averaging ({\\it thin curves}). Peaks of $E(k)$ and $W(k)$ are distinctly separated in the wavenumber space. } \\label{Fig1} \\end{figure}    Sometimes the maximum of dissipation, $k_d$, is located close to $k_e$, so that the energy and dissipation intervals are overlapped. In such a case, the dissipation regime differs from the one described above. Namely, the bulk of dissipation occurs at large scales, and both the ohmic dissipation of large-scale currents and the turbulent diffusion at smaller scales contribute into the dissipation process. Thus, the dissipation spectrum seems to be useful to diagnose the turbulence and the dissipation regime.  Another advantage of the dissipation spectrum arises from its straightforward relation with the mean squared curl of a field $\\bf B$ (Monin \\& Yaglom 1975, Biskamp 1993), in other words, with the electric currents:  \\begin{equation} \\langle {(\\nabla \\times {\\bf B} )}^2 \\rangle = 2 {\\int_0^{\\infty}}  k^2 E(k) dk. \\label{rot} \\end{equation} This quantity gives us the mean value of dissipation per unit volume per unit time. We have in our possession the $B_z$ components of the magnetic field only. When we apply Eq. (\\ref{rot}) for $B_z$ by substituting the energy spectrum of $B_z$ as $E(k)$, we acquire the part of dissipation associated with the $B_z$ component, i.e., the mean value of dissipation per unit area per unit time measured in units of G$^2$/Mm$^2$. We denoted this value as $\\langle W \\rangle$: \\begin{equation} \\langle W \\rangle = 2{\\int_0^{\\infty}}  k^2 E(k) dk. \\label{Wsr} \\end{equation}  One can see that $\\langle W \\rangle$ and $\\langle \\varepsilon \\rangle$ determine the same physical quantity derived, however, from different approaches: from spatial and spectral considerations. This circumstance will allow us to provide two estimations for the mean magnetic energy dissipation. Besides, the two approaches allow us to study different aspects of the horizontal electric currents. The current density can be derived as \\begin{equation} \\langle {\\bf j}_h^2\\rangle =\\frac{1}{\\mu^2}\\langle W \\rangle. \\label{jtrW} \\end{equation} ",
        "conclusions": ""
    },
    "0806/0806.4362_arXiv.txt": {
        "abstract": "Gravitational instabilities in a magnetized Friedman - Robertson - Walker (FRW) Universe, in which the magnetic field was assumed to be too weak to destroy the isotropy of the model, are known and have been studied in the past. Accordingly, it became evident that the external magnetic field disfavors the perturbations' growth, suppressing the corresponding rate by an amount proportional to its strength. However, the spatial isotropy of the FRW Universe is not compatible with the presence of large-scale magnetic fields. Therefore, in this article we use the general-relativistic (GR) version of the (linearized) perturbed magnetohydrodynamic equations with and without resistivity, to discuss a generalized {\\em Jeans criterion} and the potential formation of density condensations within a class of homogeneous and anisotropically expanding, self-gravitating, magnetized fluids in curved space-time. We find that, for a wide variety of anisotropic cosmological models, {\\em gravito-magnetic} instabilities can lead to sub-horizonal, magnetized condensations. In the non-resistive case, the {\\em power spectrum} of the unstable cosmological perturbations suggests that most of the power is concentrated on large scales (small $k$), very close to the horizon. On the other hand, in a {\\em resistive} medium, the critical wave-numbers so obtained, exhibit a delicate dependence on resistivity, resulting in the reduction of the corresponding Jeans lengths to smaller scales (well bellow the horizon) than the non-resistive ones, while increasing the range of cosmological models which admit such an instability. ",
        "introduction": "It is known that, the formation of large-scale structures (such as galaxies, super-clusters of galaxies and/or super-clouds of stars within galaxies) both in cosmological and in astrophysical scale is closely related to the concept of {\\em instability}~\\cite{1}, \\cite {2}.  The first attempt towards a theory of galaxy formation carried out by Sir James Jeans in the early nineteenths~\\cite{3}, \\cite{4}. He assumed that the Universe is filled with a self-gravitating non-relativistic fluid, characterized by its mass-density $(\\rh)$, the corresponding pressure $(p)$ and a velocity field $(\\vec{u})$. As a consequence, the resulting dynamical system is governed by the equation of continuity, together with Euler's and Newton's equations for the matter-content and the gravitational field, respectively~\\cite{5}. However, Jeans' theory does not apply to the formation of galaxies in an expanding Universe, because it involves a static medium. Nevertheless, he arrived at the conclusion that, density fluctuations with wavelengths $\\lm$ greater than a {\\em critical value} $\\lm_J$ will grow, so that the system, eventually, becomes {\\em unstable} (gravitational instability). As a consequence, any isothermal gaseous sphere with a length-scale greater than $\\lm_J$ is gravitationally unstable and is going to contract~\\cite{6}.  The first satisfactory theory of gravitational instabilities in an expanding Universe was proposed by Lifshitz~\\cite{7}. He showed that, in a homogeneous and isotropic Universe filled with a perfect fluid, cosmological perturbations with wave-numbers $(k)$ below the Jeans' one $(k_J)$ grow, not exponentially, but, as a power of the corresponding scale factor, $R (t)$. Among the various scenarios of cosmic fluctuations there is also the case where gravitation and cosmic expansion are essentially irrelevant. In this context, Weinberg's work~\\cite{8} is of particular interest: For the first time, the notion of an {\\em imperfect fluid} was used and the role of {\\em dissipation} in the survival of protogalaxies was taken into account, along the lines of Eckart's formalism~\\cite{9}, which differs from the corresponding of Landau \\& Lifshitz~\\cite{10}. Accordingly, Weinberg found that the protogalactic fluctuations, which behave as ordinary sound waves during the period where their matter-content is much lower than the corresponding Jeans mass $(M << M_J)$, damp in the acoustic phase due to the {\\em viscosity} of the imperfect fluid.  The study of gravitational instabilities in the presence of magnetic fields in an expanding Universe arose as a natural generalization~\\cite{11}. Accordingly, solutions for the evolution of the mass-density and the magnetic field $(H^i)$ perturbations (in the linear regime) have been quested in FRW cosmological models, in which, the magnetic field was assumed to be too weak to destroy the isotropy of the model~\\cite{12} - \\cite{16}.  Although it has been argued that a uniform magnetic field slows down the growth-rate of the unstable modes, several authors have studied the influence of cosmological magnetic fields (either homogeneous or inhomogeneous) to the formation of structures within a cosmic medium (e.g. see~\\cite{17} - \\cite{25}, for an extensive, though incomplete list). In this context, it was shown~\\cite{26} that several magnetohydrodynamic processes like pinch-effects, hose-instabilities, sausage- and kink-instabilities, may also play a significant role to structure formation. These processes correspond to individual harmonics in the eigenfunction solutions of the wave equation for plasma instabilities and although they are different, their time-development is similar.  In the present article, we intend to discuss gravitational instabilities in a homogeneous and anisotropic cosmological model, in which the large-scale anisotropy is due to the presence of a homogeneous magnetic field along the $\\hat{z}$-direction.  Our choice is justified by the fact that, mathematically speaking, the spatial isotropy of the FRW Universe is not compatible with the presence of large-scale magnetic fields. In fact, an anisotropic cosmological model can and should be imposed for the treatment of magnetic fields whose {\\em coherent length} is comparable to the {\\em horizon length}~\\cite{27}. Therefore, although current observations give a strong motivation for the adoption of a FRW model, the effects one may lose by neglecting the large-scale anisotropy induced by the background magnetic field, should be investigated. Not to mention that the anisotropy of the unperturbed model facilitates a closer study of the coupling between magnetism and geometry.  The Paper is organized as follows: In Section II, we summarize the general-relativistic version of magnetohydrodynamics (MHD) and present its application to the Thorne's class of homogeneous and anisotropic cosmological models~\\cite{28}. They correspond to a class of {\\em semi-realistic} solutions to the Einstein-Maxwell equations, adequate to describe a model which departs from isotropy along the radiation epoch~\\cite{29}. In Section III, we perturb the corresponding MHD equations to study the evolution of linear fluctuations propagating along the $\\hat{x}$-direction, i.e. perpendicularly to the magnetic field. An immediate result is the derivation of the corresponding {\\em dispersion relation}, which may reveal the {\\em critical values} of the perturbations' wave-numbers in order to discuss on their potential growth (or decay). We find that, for a wide class of anisotropic cosmological models, {\\em gravito-magnetic} instabilities can lead to sub-horizonal, magnetized condensations. In this case, the {\\em power spectrum} of the unstable cosmological perturbations suggests that most of the power is concentrated on large scales (small $k$), close to the horizon. These results are in contrast to what has been found in the isotropic case, where the magnetic field disfavors large-scale matter aggregation, by reducing the growth-rate of the corresponding density perturbations \\cite{11}, \\cite{13}, \\cite{14}. In Section IV, we consider the relativistic MHD equations within a {\\em resistive} medium. Now, the critical wave-numbers so obtained exhibit a delicate dependence on resistivity, resulting in the reduction of the corresponding Jeans lengths to smaller scales than the non-resistive ones, while increasing the range of cosmological models which admit such an instability. ",
        "conclusions": "Gravitational instabilities in a magnetized FRW Universe, in which the magnetic field was assumed to be too weak to destroy the isotropy of the model, are known and have been studied in the past~\\cite{13}, \\cite{14}, \\cite{16}. However, mathematically speaking, the spatial isotropy of the FRW Universe is not compatible with the presence of large-scale magnetic fields. In fact, an anisotropic cosmological model can and should be imposed for the treatment of magnetic fields whose coherent length is comparable to the horizon length~\\cite{27}. Interesting homogeneous and anisotropic cosmological models of Bianchi Type I with an ambient magnetic field are known~\\cite{28} and have been used to discuss astrophysical implications~\\cite{6}, \\cite{19}, \\cite{25}, \\cite{29}, \\cite{31}.  In this article, we have used the GR version of the MHD equations~\\cite{18} in a homogeneous and anisotropically expanding, self-gravitating, magnetized plasma, to study the gravito-magnetic excitation of cosmological perturbations. Accordingly, we have found that:  \\begin{itemize}  \\item In a non-resistive medium:  \\begin{enumerate}  \\item In contrast to the isotropic case, where the role of the magnetic field (as an agent opposing the perturbations' growth) is clear, there is a wide class of anisotropic, magnetized cosmological models - parameterized by $\\gm$ in the range ${1 \\over 3} \\leq \\gm \\lesssim 0.5$, for which a {\\em gravito-magnetic} instability can lead to sub-horizonal magnetized condensations [clf Eqs (29), (31) and (32)].  \\item We have determined explicitly the {\\em power spectrum} of the unstable cosmological perturbations [Eq (37)], which indicates that most of the power is concentrated on large scales, close to the horizon [clf Eq(40)].  \\end{enumerate}  \\item In a resistive medium:  \\begin{enumerate}  \\item The critical wave-numbers so obtained, exhibit a delicate dependence on resistivity [Eq (52)], resulting in the reduction of the corresponding Jeans lengths to smaller scales (well below the horizon) than the non-resistive ones. In fact, the non-zero resistivity induces convective currents, resulting in the enhancement of the perturbations' energy (frequency) and hence in the reduction of their comoving wavelength.  \\item Furthermore, the non-zero resistivity prolongs the range of the potential values of $\\gm$ (and therefore the number of the corresponding cosmological models) which admit sub-horizonal scales of such an instability (from ${1 \\over 3} \\leq \\gm \\lesssim {1 \\over 2}$ in the non-resistive case, to ${1 \\over 3} \\leq \\gm \\lesssim {2 \\over 3}$ in the fully-resistive one) [clf Eq (54)].  \\end{enumerate}  \\end{itemize}"
    },
    "0806/0806.1771_arXiv.txt": {
        "abstract": "We analyze the orbital and mass evolution of planets that undergo run-away gas accretion by means of two- and three-dimensional hydrodynamic simulations. The disk torque distribution per unit disk mass as a function of radius provides an important diagnostic for the nature of the disk-planet interactions. We first consider torque distributions for nonmigrating planets of fixed mass and show that there is general agreement with the expectations of resonance theory. We then present results of simulations for mass-gaining, migrating planets. For planets with an initial mass of $5$ Earth masses ($\\MEarth$), which are embedded in disks with standard parameters and which undergo run-away gas accretion to one Jupiter mass ($\\MJup$), the torque distributions per unit disk mass are largely unaffected by migration and accretion for a given planet mass. The migration rates for these planets are in agreement with the predictions of the standard theory for planet migration (Type~I and Type~II migration). The planet mass growth occurs through gas capture within the planet's Bondi radius at lower planet masses, the Hill radius at intermediate planet masses, and through reduced accretion at higher planet masses due to gap formation. During run-away mass growth, a planet migrates inwards by only about $20$\\% in radius before achieving a mass of $\\sim 1\\,\\MJup$. For the above models, we find no evidence of fast migration driven by coorbital torques, known as Type~III migration. We do find evidence of Type~III migration for a fixed mass planet of Saturn's mass that is immersed in a cold and massive disk. In this case the planet migration is assumed to begin before gap formation completes. The migration is understood through a model in which the torque is due to an asymmetry in density between trapped gas on the leading side of the planet and ambient gas on the trailing side of the planet. ",
        "introduction": "In the core accretion picture of planet formation \\citep[and references therein]{bodenheimer1986,wuchterl1991b,% pollack1996,hubickyj2005}, a small mass solid core initially rapidly accretes solid material, followed by a slow evolution phase of gas and solid accretion. During this slow evolution phase, the planet is limited in its ability to accrete gas by the thermal heating caused by the impacting solids. Once the planet's gas mass is greater than its solid mass, typically at several Earth masses, the planet undergoes ``run-away'' gas accretion, in which it can accrete whatever mass is provided to it. These processes have been treated by one-dimensional, spherically symmetric structure calculations in the above papers.  On the other hand, multi-dimensional hydrodynamical calculations of a protostellar disk interacting with the planet has revealed various flow properties of the gas, including the gap opening by tidal effects, previously anticipated by one-dimensional disk models \\citep{lin1986}. In addition, planet migration that results from disk-planet interactions has been analyzed by means of such simulations. Good agreement is often, but not always, found between the simulations and the expectations of theory \\citep{rnelson2000,bate2003,gennaro2003b,anelson2003b,li2005,gennaro2006}. These calculations typically do not include the mass evolution of the planet. Usually they apply accretion boundary conditions onto the planet as a means of modelling the run-away gas accretion process. One aim of this paper is to analyze the effects of planet mass growth on migration.  Several controversies remain on the effects of gas. The role of coorbital torques on planet migration, in the subgiant mass range, is not well understood. \\citet[hereafter MP03]{masset2003} suggested on the basis of a model and simulations that a fast mode of migration (sometimes called Type III migration) can occur due to strong coorbital torques. \\citet[hereafter OL06]{ogilvie2006} found support for the concept of coorbital dominated migration under certain conditions. At higher grid resolution under the conditions specified by MP03, simulations by \\citet[hereafter DBL05]{gennaro2005} found that the migration rate was much slower.  Another subject of interest is how planet masses may be limited by a reduction in the gas accretion rate. \\citet{lin1986} proposed such a reduction by tidal torques that open a gap about the orbit of the planet. The value of the highest planet mass achieved in the presence of gap opening is somewhat controversial. Some studies \\citep{lubow1999,bate2003,gennaro2003b} have suggested that the maximum planet mass is about $6$--$10\\,\\MJup$, corresponding to the upper limit of the observed range of extrasolar planets \\citep{marcy2005,butler2006}. This limit suggests that some other process, such as disk dispersal or other self-limiting feedback on planetary accretion, is responsible for the lower masses ($\\sim 1\\,\\MJup$) typically found observationally. Other studies suggest that the tidal limit is $\\sim 1\\,\\MJup$ and therefore no additional process is required to explain the typical masses \\citep[e.g.,][]{dobbs-dixon2007}.  We will address these and other issues in this paper by analyzing the orbital evolution of a mass-gaining planet embedded in a gas disk. In section~\\ref{sec:TorqueDistributions} we analyze the torque distributions for planets of constant mass on fixed circular orbits. In section~\\ref{sec:GrowingMigratingPlanets} we analyze the orbital and mass evolution of migrating planets that undergo run-away mass accretion. Section~\\ref{sec:typeiii} describes a model that appears to exhibit migration that is dominated by coorbital torques, i.e., Type~III migration. Section~\\ref{sec:summary} contains the summary and discussion. ",
        "conclusions": "\\label{sec:summary}  We have analyzed the evolution of migrating planets that undergo run-away gas accretion by means of multi-dimensional numerical simulations. The results agree with the predictions of Type~I and Type~II migration (see Figures~\\ref{fig:a_sn} through \\ref{fig:a_sn_nu}) for a planet of time-varying mass that we obtain from simulations. The set of simulations include cases with disk densities as low as the minimum mass solar nebula value and as high as $5$ times that value, viscosities $\\alpha\\approx 0.004$ and $10$ times that value (also $50$ times that value, $\\alpha\\approx 0.2$, in a test reported in Appendix~\\ref{sec:MVV}), disk temperatures corresponding to $H/r = 0.05$ and a colder case of $H/r=0.04$. The mass accretion rates onto the planet (see Figures~\\ref{fig:mp_sig} and \\ref{fig:mp_hnu}) are in general agreement with previous determinations based on fixed mass planets on fixed circular orbits, when comparing cases with the same planet mass and disk properties for which the growth timescale is longer than the gap opening timescale. Planet mass growth rates can be understood in terms of accretion within the Bondi radius at lower planet masses, accretion within the Hill radius at intermediate masses, and accretion through the gap at higher planet mass (see Fig.~\\ref{fig:tauG_sn_fr}). Mass growth rates typically peak at the intermediate planet masses, a few tenths of a Jupiter mass.  An important diagnostic for the nature of the disk-planet torques is the scaled torque density distribution per unit disk mass as a function of the scaled radial distance from the planet. In the linear regime of the standard theory of disk-planet resonances, for a fixed form of the disk density distribution $d \\ln{\\Sigma}/d\\ln{r}$ and gas properties (sound speed and viscosity), this scaled torque distribution should be universal, independent of planet mass and disk density value. We verified this universality for low mass planets, although the distribution varies somewhat with planet mass for larger mass planets that open gaps (compare low mass cases in Figures~\\ref{fig:dTdM}, \\ref{fig:dTdM_sn_mig}, and \\ref{fig:dTdM_xsn_mig}).  \\begin{figure} \\centering% \\resizebox{\\linewidth}{!}{% \\includegraphics{f20}} \\caption{Orbital migration of a planet undergoing run-away gas accretion. Orbital radius in units of $a_{0}$ ($5.2\\,\\AU$), as a function of time in units of the initial orbital period ($\\approx 12$ years). The initial planet mass is $5\\,\\MEarth$. The initial surface density is $\\Sigma_{p}=3\\times10^{-4}\\,\\densu% \\approx 100\\,\\mathrm{g}\\,\\mathrm{cm}^{-2}$ at the planet's initial orbital radius and $H/r=0.05$. \\textit{Solid curve}: Results from the same simulation as in Figure~\\ref{fig:a_sn} for a migrating, mass-gaining planet. \\textit{Short-dashed curves}: Predictions based on Type~I migration theory, obtained by solving equations~(\\ref{eq:tanakadotaI}) and (\\ref{eq:tanakadotaIsat}), taking into account the time variable mass of the planet (Fig.~\\ref{fig:mp_sig}, \\textit{solid line}) and roughly accounting for the gas depletion near the planet (see text). The upper (lower) dashed curve is for migration with unsaturated (saturated) coorbital torques. } \\label{fig:a_snc} \\end{figure} There is no fundamental distinction between the torques involved in Type~I and Type~II migration. This follows from the near independence of the scaled torque density distribution with planet mass. Previous concepts of Type~II suggested that a planet in a clean gap would migrate inward like a test particle that follows the disk accretion. Here we describe a view for cases where the gap is not completely clear of material. The difference in Type~I and Type~II rates is due to the mass density distribution of the gas that multiplies torque density in determining the torque on the planet. In Type~II migration, the density distribution adjusts so that the net torque on the planet causes it to migrate at approximately the viscous evolution rate of the disk. The transition between the two forms of migration is quite smooth. To illustrate this point, we plot in Figure~\\ref{fig:a_snc} the orbital evolution of the planet through the phase where gap formation sets in. The figure shows that the migration rate can be accounted for by standard Type~I theory, corrected for the gas depletion in the gap region, although there is no unique prescription to do this. In the figure, the theoretical curves (\\textit{dashed lines}) are determined by Type~I migration theory, equations~(\\ref{eq:tanakadotaI}) and (\\ref{eq:tanakadotaIsat}), with the density $\\Sigma$ taken to be the average value in a radial band of half-width $0.15\\,a$ (a typical gap width) centered on the planet.  For a given planet mass, the torque density diagnostic reveals that the distribution is not strongly affected by migration or accretion (Figure~\\ref{fig:dTdM} is quite similar to Figures~\\ref{fig:dTdM_sn_mig} and \\ref{fig:dTdM_xsn_mig} for similar planet masses). In particular, there is no evidence for strong coorbital torques. However, in a certain case, we do find evidence for strong coorbital torques, or Type~III migration. This case has a planet of fixed mass, $0.3 \\MJup$, that is immersed into a cold, smooth disk. The planet is held at a fixed orbit for relatively short time ($\\sim 100$ orbits) before being released, so that gap clearing is incomplete. The torque distribution at the time the planet is released follows the expectations of the standard theory for nonmigrating planets. This can be seen by comparing the curve for the $0.3\\,\\MJup$ case in Figure~\\ref{fig:dTdM} (\\textit{dot-dashed curve}) with the short-dashed curve (time $t=t_{\\mathrm{rls}}$) in Figure~\\ref{fig:dTdM_fast}. (There are differences in the magnitude of the distributions because of differences in gas sound speeds and dimensionalities of the calculations, but the forms of the distributions are similar). At later times, Figure~\\ref{fig:dTdM_fast} reveals a transition to a completely different torque distribution, where coorbital torques play a critical role.  The strong coorbital torque can be understood in terms of an asymmetry in the streamlines between the leading and trailing sides of the planet, in accord with the analytic model of OL06 (see Figures~\\ref{fig:str-nonmig} and \\ref{fig:str-mig}) and also along the lines of \\citet{pawel2004}. This asymmetry causes trapped material to persist on the leading side of the planet which has a different density from the ambient gas that flows on the trailing side (see Fig.~\\ref{fig:sig_B_fast}). The asymmetry gives rise to the coorbital torque. We suggest some criteria for this form of migration (see section~\\ref{sec:typeiii_conditions}). More exploration is needed to test them.  Although we find evidence for Type~III migration, the conditions required appear somewhat artificial, i.e., incomplete gap clearing (nonequilibrium gap) of a cool disk with a planet of fixed mass. It is not clear whether and/or how conditions for Type~III could arise in a more plausible evolution scenario.  We have generally assumed that the planet is able to accrete almost all gas the disk is able to provide (so-called run-away gas accretion). For a disk viscosity $\\nu\\gtrsim 1\\times 10^{-5}\\,\\viscu$, the mass accretion rates are large. The time to build a $1\\,\\MJup$ planet starting with a $5\\,\\MEarth$ planet in a minimum mass solar nebula is shorter than $\\sim 10^5$ years, substantially less than the observationally determined disk lifetimes of $\\sim 10^6$ years \\citep{haisch2001,flaherty2008}. Over this $10^5$ year time interval, the planet has radially migrated inward by only $\\sim 20$\\% of its initial radius. In other words, for these models, the mass doubling timescale for a $\\Mp \\la 1\\,\\MJup$ planet is short compared to the migration timescale and the disk lifetime. This situation stands in strong contrast to the earlier phases of planet formation where the migration timescales are shorter than the planet mass doubling timescales and disk lifetimes \\citep[e.g.,][]{ward1997,hubickyj2005}.  The run-away accretion rates pose some challenges for explaining the mass distribution of planets \\citep{butler2006}. Typical accretion rates in T Tauri stars are $\\sim 1\\times 10^{-8}\\,M_{\\odot}$ per year \\citep{hartmann1998}. For a steady-state unperturbed disk (without a planet), the accretion rate is given by $3 \\pi \\nu \\Sigma$. The initial disks considered in this paper are not in a steady state, but this accretion rate provides a reasonable estimate. For the minimum mass nebula model (Fig.~\\ref{fig:mp_sig}), the kinematic  turbulent viscosity we typically adopt, $\\nu\\sim 10^{-5}\\,\\viscu$, was chosen so that the accretion rate evaluates to about this same value, $\\sim 1\\times 10^{-8}\\,M_{\\odot}$ per year. In the case of a planet embedded in a disk, if there is a comparable accretion rate onto a planet of mass $\\Mp \\la 1 \\MJup$ \\citep[as found by][]{lubow2006}, then the mass doubling timescale for a Jupiter-mass planet is about $10^5$ years, consistent with what we found in the simulations in this paper. But then it is not at all clear why planets would not almost always achieve masses higher than $\\sim 1\\,\\MJup$, in contradiction with the observed mass distribution of extra-solar planets and the case of Saturn. Special timing for disk dispersal could be invoked, but may be artificial.  There are a few possible explanations. A colder disk ($H/r < 0.05$) would experience stronger tidal truncation effects from a Jupiter-mass planet, as $H$ becomes significantly smaller than $\\Rhill\\simeq 0.07\\,a$. This effect could certainly reduce the accretion rate by a large enough factor (say $10$), so that the mass doubling timescale for a $\\sim 1 \\MJup$ planet would become of order $10^6$ years, consistent with the suggestion by \\citet{dobbs-dixon2007}. The issue then is how cold the disk would need to be. The model in this paper with $H/r =0.04$ (Fig.~\\ref{fig:mp_hnu}, \\textit{long-dashed line}) has the same unperturbed overall disk accretion rate as the $H/r =0.05$ case (Figures~\\ref{fig:mp_sig} and \\ref{fig:mp_hnu}, \\textit{solid line}), since $\\nu$ and the initial surface density $\\Sigma(r)$ are the same. The accretion rate onto the planet from the colder disk ($H/r=0.04$) differs from the case with $H/r=0.05$ by only $10$\\%. Model $h$ in \\citet{lubow2006} for a nonmigrating planet of fixed mass, having $H/r=0.03$, has an accretion rate that is $\\sim 25$\\% less than model $b$, which has $H/r=0.05$ and the same unperturbed overall disk accretion rate. Consequently, in this disk thickness range ($0.03\\le H/r \\le 0.05$), we find that the reduction in the accretion rate onto the planet due to cooler disks is not significant. For higher mass planets ($5$--$10\\,\\MJup$), we expect that the accretion rate will be reduced by tidal truncation effects to a level where the planet mass doubling timescale is comparable to the disk lifetime, as found in previous studies of planets on fixed orbits \\citep{lubow1999,bate2003,gennaro2003b}. This effect  may set the upper limit to planet masses.  Another possibility is that there is a feedback effect that limits the gas accretion  rate. Perhaps the heating of the protoplanet envelope by impacting solids continues to later times than is assumed in the standard core accretion model. Depletion of disk solids near the planet occurs in the standard core accretion model, when planet migration is not included. With migration, it is possible that continued accretion of disk solids would occur \\citep[e.g.,][]{alibert2005a}, resulting in continued heating that could limit the gas accretion rate further. It is not clear how well this possibility works, since planetesimals will get trapped into resonances as the planet migrates \\citep{zhou2007}. Having a higher mass solid core is problematic in the case of Jupiter, whose solid core mass is thought to be a small fraction of the total mass (see \\citealp{guillot2005} and discussion in \\citealp{lissauer2007}). It is also possible that winds emanating from the circumplanetary disk within the planet's Hill sphere could reduce the accretion rate onto the planet. Magnetically driven winds are believed to play an important role in the case of young stars \\citep{blandford1982,pudritz1986,shu1994}. There are likely differences in the flow properties from the stellar outflow case \\citep{fendt2003}. However, it is not clear that the winds would be able to expel a large enough fraction (say $90$\\%) of the accreting gas to sufficiently reduce the accretion rate onto the planet."
    },
    "0806/0806.3404_arXiv.txt": {
        "abstract": "{} {Our goal is to measure the nitrogen  abundance in 5 Turn Off (TO) stars of the two Globular Clusters NGC 6397 and NGC 6752, and to compare the cluster abundances with those of field stars of comparable metallicity. In addition we wish to investigate the   variations of nitrogen abundance and its  connections to variations in the abundances of other light elements. } {We determine the nitrogen abundance from the compact band head of the (0-0) vibrational band of the NH $A^3\\Pi-X^3\\Sigma^-$ system at 3360 \\AA, using spectra  of resolution R$\\sim 45000$ obtained with the UVES spectrograph on the ESO Kueyen-VLT 8.2m telescope, analysed with spectrum synthesis based on plane parallel LTE model atmospheres. We apply the same method previously used on field stars, to allow a direct comparison of the results. } {Nitrogen is found to have the same abundance in two of the NGC 6397 stars, in spite of a  difference  of one order of magnitude in oxygen abundance between  them. In a third star of the same cluster, the value is  slightly lower, but compatible with the other two, within the uncertainties. All the stars in NGC 6397 are N-rich with respect to  field objects of similar metallicity. The two stars in NGC 6752  show a difference in nitrogen abundance by over one order of magnitude. The same stars differ in the abundances of other elements such as Na, O and Li,  however   only by a factor $\\sim$3-4.} {NGC 6397 and NGC 6752 are, at present, the only two globular clusters in which N has been measured in TO stars in a way consistent with similar measurements in  field stars. The behaviour of N is  different in the two clusters: no variation is observed NGC 6397, while a large variation is observed in NGC 6752. This is consistent with a picture in which the stars in NGC 6752 have been formed by a mixture of ``pristine'' material and material which has been processed by an early generation of stars, referred to as ``polluters''. The N abundances here reported will help to constrain the properties of the polluters. In the case of NGC 6397  a simple pollution history is probably not viable, since the observed variations in O abundances are not accompanied by corresponding variations in N or Li. } ",
        "introduction": "In spite of the fact that nitrogen is one of the  most abundant elements in the cosmos, only recently it has been possible to produce  large samples of N abundances in metal poor stars, enabling us  to quantitatively study its evolution in the Galaxy (see e.g. Israelian et al. 2004, and references therein). Nitrogen tracks iron ([N/Fe]$\\sim 0$) down to a metallicity of $\\sim -1.0$. The evolution at lower metallicity is less clear; while many stars seem to indicate that N continues to track iron, a few field stars are found to be nitrogen-rich (the prototypes of the class being HD 74000 and HD 160617 \\citealt{BN82}), while others seem to show [N/Fe] ratios below solar. In their investigation of the chemical composition in the lowest metallicity stars \\citet{spite05} suggested that a possible interpretation of the data in the metallicity range $\\rm -4.0 \\le [Fe/H]\\le -2.5$ is a bimodal distribution of nitrogen abundances, the majority of stars at [N/Fe]$\\sim +0.4$ and the rest at [N/Fe]$\\sim -1.0$. Given the complexity of  its nucleosynthesis, iron is not the best reference element to study chemical evolution, as discussed in \\citet{cayrel04}, and an $\\alpha$-chain element such as O or Mg is preferable. The [N/$\\alpha$] ratio begins to decline as one moves from solar to lower metallicities (see e.g. figure 17 of \\citealt{spite05}) and at low metallicities it reaches a constant value, a {\\em plateau} at  [N/$\\alpha$]$\\sim -0.9$, although it could be argued that there are in fact {\\em two} plateaus, the second found around [N/$\\alpha$]$\\sim -1.4$.  From the nucleosynthetic point of view there is only one mode of nitrogen production: it is formed in the course of the CNO cycle, at the expenses of available C and O. Since this requires the pre-existence of nuclei of C and O, one may expect the production of N to increase with increasing abundance of C and  O, i.e. with metallicity. Nitrogen would then be a so-called purely {\\em secondary} element, to distinguish it from {\\em primary} elements, like the $\\alpha$-chain elements, whose yield is independent of metallicity. The ratio of a purely {\\em secondary} element to a primary one should be a monotonically increasing function of metallicity. The fact that [N/$\\alpha$] shows a plateau at low metallicity demonstrates that N cannot be a purely secondary element. We note here that although the concept of {\\em primary} and {\\em secondary} elements is a useful simplification, for the majority of the elements the production history is complex enough that it is inappropriate to label it with either of these terms.   For nitrogen, although the nucleosynthetic process is unique, it is necessary to invoke different sites where the process occurs to explain the two kinds of behaviour. Intermediate mass stars in the AGB phase (IM-AGB stars, Ventura et al. 2002), in which the H-burning shell operates through the CNO cycle, have always been considered major producers of nitrogen. For these stars the necessary C and O nuclei are those already present at the time of the formation of the star and the production mode would be {\\em secondary}. Massive stars are other possible N producers.  In such stars the central H burning takes place  through the CNO cycle, and again the mode of production would be secondary. If, however, the star is rotating, as soon as the central He burning is ignited, fresh C and O produced in the core may be transported through rotational mixing to the H-burning shell. This  would then produce N in a primary fashion, since the necessary parent nuclei are created {\\em in situ} (see Maeder and Meynet (2005) for a complete discussion). In such a scenario N production cannot be purely primary, since some secondary production is unavoidable. A primary mode of production of N may occur also in IM-AGB stars, provided a way of mixing the products of the He burning shell to the H burning shell exists.  The two above-mentioned producers (IM-AGB stars and massive stars) release their nucleosynthetic products, which become available for the next generation of stars, on quite different time-scales. Massive stars will explode as core-collapse supernovae on time scales shorter that 10 Myr, while IM-AGB will release their products through winds or in the planetary nebula phase on a time scale between 300 Myr and 1Gyr, depending on their mass and metallicity.  The pattern of  N abundances in Globular Clusters is somewhat more complex than in field stars. In the first place, like other light elements, N displays sizeable star to star scatter within the cluster. From studies conducted with intermediate band photometry or low resolution spectra it has been clear that many clusters display variations of N abundances \\citep[see][for a review of these older results]{kraft}. For many years such effects were supposed to be caused by mixing in the giant stars, which were the only ones it was possible to analyse \\citep[see][for an interesting idea on the role of rotation]{norris81}. In the recent years, largely thanks to advent of 8m class telescopes, it has been possible to investigate these effects down to the TO of a few Globular Clusters, and the abundance variations appear to persist. Since a TO star is not expected to experience mixing, this calls for another explanation for these variations \\citep{carretta05}. Although it is generally accepted that the stars we presently observe in GCs have been formed out of material which has been inhomogeneously polluted by a previous generation of stars, there is little consensus on which stars were responsible for the pollution. Proposed candidates are IM-AGB stars (D' Antona et al. 2005) and massive stars (see e.g. Decressin et al. 2007a) , including Wolf-Rayet objects (Smith 2006).  It seems therefore timely to compare the N abundance observed in unevolved field and Globular Cluster stars, by making the  effort to place  them on the same abundance scale, and  to use this comparison to shine some light on the formation of nitrogen in the Galaxy and on the formation process of Globular Clusters and on the Globulars - Halo relationship.  We here present abundances for five TO stars (3 \\relax in NGC 6397 and 2 \\relax in NGC 6752) derived from the compact band head of the (0-0) vibrational band of the NH $A^3\\Pi-X^3\\Sigma^-$ system at 3360 \\AA. This feature is stronger than  the UV CN bands used in \\citet{carretta05} and should provide a more accurate measurement. Moreover the N abundance derived from the NH bands does not depend on either C nor O, instead when we use the CN bands the derived N abundance depends on the assumed C and O abundances. The O abundance enters because the majority of C atoms is always locked in the tightly bound CO molecule. For four out of the five stars of our sample preliminary abundances based on the NH bands have been presented in Pasquini et al. (2004, 2007), however we deemed necessary  a reanalysis using the same molecular data and analysis procedure as in Ecuvillon et al. (2004) and Israelian et al. (2004), to allow an unambiguous comparison of the N abundances in the GCs with those in field stars of comparable metallicity. ",
        "conclusions": "\\subsection{ Comparison with  field stars}  Our N abundances for  GC stars can be consistently compared to the results for field stars of Ecuvillon et al. (2004) and Israelian et al. (2004).  Figure \\ref{fignfe} shows the [N/H] abundance  vs. [Fe/H] for field and cluster stars. Most field stars closely follow a 1:1 relationship, with a scatter compatible with the error measurements (quoted as $\\sim \\pm$ 0.15 by Israelian et al. 2004 and given as a reference in the upper left corner of the Figure). The [N/O] ratio is shown in Fig. \\ref{figno}, for the present discussion we ignore the possible complications arising from the fact that the O abundances displayed have been derived in an inhomogeneous way and using different O indicators, and take all measurements at face value. With the {\\em caveat} that any systematic differences between the different analysis will tend to blur any existing pattern. Two points  emerge  from Figure 2 :  \\begin{enumerate}  \\item  The cluster stars do not resemble, in general, the behaviour of the field stars.   \\item  The two clusters do  not show the  same behaviour.  \\end{enumerate}  The first point is somewhat expected, since we have known  for long time of the presence of chemical anomalies in GC stars, which includes the presence of N rich stars. N rich stars, though known in the  field,   are extremely rare, and a typical sample, like the one of Israelian et al. (2004), contains at most one  such object.  The two  NGC 6752 stars show a large spread in N (almost a factor of 20) and the cluster stars, including all stars  observed in  NGC 6397, have a very high N abundance, clearly  departing from the field stars relationship at high significance. This is clearly related to the peculiar formation mechanism of GCs, and we will come back on this point in the next section.  As far as the second point is concerned we believe  that this aspect should be considered with some attention, since it has not been explored too much so far. For this reason we look at the two clusters separately.  NGC 6752 is a cluster showing all possible chemical inhomogeneities and anticorrelations. But, in its anomalies,  the observed trends are  extremely consistent. The difference of a factor $\\sim$20 in N among the two stars is very high, but it is in line with the other large differences recorded for them in Na (factor 10)  and  Li (factor 2.5 ) in these stars and in oxygen (factor 4) in other stars having abundance pattern  similar to star 200613 (in 200613 O has not been measured).  With our new value, the N of star 4428 is  higher, but formally compatible with the field stars having similar metallicity. Even if we do not know how this GC formed, the observations are  compatible with the idea that star 4428 is the prototype of an 'unpolluted' star, perfectly similar in all aspects to its field counterparts of similar [Fe/H], while star 200613 is, at the opposite, the prototype of  highly polluted star, with high N, high Na, low Li and low O, showing therefore the signature of gas which experienced high temperatures where nuclear cycles such as CNO, Ne-Na, and the consequent Li destruction, have occurred.  When plotting N vs. O (Fig. \\ref{figno}),  the situation is similar, unfortunately we have O abundances only for 3 out of five stars. The impression is that the cluster stars (including star 4428)  stand definitely out of the  field relationship. This may be related to the fact that star 4428 has an O/Fe ratio which is lower, by roughly a factor of 2, than field stars of the same [Fe/H], when N is plotted against [Fe/H] star 4428 follows the trend of field stars, but when plotted against oxygen it appears to have too much nitrogen, with respect to field stars with the same [O/H]. Clearly in a cluster like NGC 6752, which has a very strong Na-O anticorrelation, it is difficult to tell which is the ``unpolluted'' oxygen content of the cluster. Among the stars analysed by Gratton et al. (2001) there is one dwarf and two subgiants which have higher [O/H] than star 4428. We have chosen star 4428 as template of the ``unpolluted'' population, because it is the star with the lowest [Na/Fe]. New measurements of N, based on the NH bands for all the stars observed by Gratton et al. (2001), would be desirable, given the weakness of the CN bands and the difficulty to use them to measure N.  Summarising, for NGC 6752, our results confirm that the chemical composition of the stars is consistent with the  assumption that each object is composed by a mixture of  'pristine' material, possibly N rich, but otherwise indistinguishable  from that forming typical halo stars, with highly processed material, N and Na rich, O and  Li  poor.  The NGC 6397 cluster stars are on the other hand more intriguing: the NGC 6397 objects have  very similar N abundance, with a negligible spread,  but  at an absolute level about 10 times higher than field stars of similar metallicity. It is not possible, even taking the highest rate of rapid rotating massive stars, to insert such a high N production into a standard Galactic chemical evolutionary model (see e.g. Chiappini et al. 2006), thus such a high N is the direct proof that the NGC 6397 stars have been affected by  a peculiar  process. In their extensive study of the Globular Cluster system of M 31 \\citet{burstein}, using integrated light spectra, have shown that the older clusters of M 31 have considerably stronger NH bands than Galactic clusters. The reasons for this difference is unclear. NGC 6397 is not among the Galactic Globular Clusters considered by   \\citet{burstein}, but adopting the naive point of view that NH band strength is related to N abundance, one may expect that NGC 6397 \\relax in integrated light should show a stronger NH band than other   Galactic Globular Clusters, although not necessarily as strong as that of the M 31 clusters. It remains to be investigated if the mechanism which causes a strong NH band in the Clusters of M 31 is the same, or akin to, that which causes the high N abundance in NGC 6397.   The oxygen content of the cluster is, on the opposite,  very  low; at least a factor of two lower than the field stars of similar metallicity. Even if  the behaviour of oxygen in metal poor stars is highly debated, oxygen has been found at the level of only [O/Fe] = 0.2-0.3 in NGC 6397 by all authors (Th\\'evenin et al. 2001, Pasquini et al. 2004,  Gratton et al. 2001). Carretta et al. 2005 show that the C in NGC~6397   is at a level of [C/Fe]$\\sim$0,  which is very similar to what is observed in field stars of similar metallicity. This  clearly indicates that for the dwarfs observed in this  cluster the whole CNO balance is different to that of field stars. The Li abundance (Molaro \\& Pasquini, 1994, Pasquini \\& Molaro 1996, Th\\'evenin et al. 2001, Bonifacio et al. 2002, Korn et al. 2006) is, on the other hand, absolutely consistent with the {\\em Spite plateau} \\citep{spite82a,spite82b}, and  Na, with a slightly supersolar value, does not vary at a significant level in the stars so far observed (Gratton et al. 2001).  As far as the chemical composition of the TO stars is concerned, the homogeneity in NGC 6397 stars seems much higher than in NGC 6752, in that  only a large difference of oxygen has been reported by Pasquini et al. (2004) for  two stars of this sample, while no substantial spread is observed in any other element among the main sequence stars. Only three subgiants  show evidence for strong Na and N variations and no real 'anticorrelation' can be deduced by the abundances so far published (Gratton et al. 2001, Carretta et al. 2005, Table 1).  Of course we must be careful in deriving strong conclusions from such a low statistics, and we are aware that some of the previous evidences (or missing evidences) could be simply the result of a combination of a limited sample and large errors, however, we think we may summarise this part of the discussion saying that, in spite of some similarity,  the two clusters show a different behaviour with respect to their N content and, in general, the spread of  light elements amongst their stars: while the NGC 6752 stars abundances show consistent signatures of pollution, the NGC 6397 stars show at the same time the signatures of  CNO (and likely NeNa) processed  material, while showing at the same time no clear anticorrelation, a much larger homogeneity in  abundances, and a Li content similar to the {\\em Spite plateau}.   \\subsection{About GC formation}  One fundamental consideration in the discussion of the formation scenarios is the fact that  anti-correlations persist in evolved stars after the dredge-up (see e.g. Grundhal et al. 2002). This shows that the whole stellar mass, and  not only the external stellar layers, is affected by the chemical anomalies.  As far as this paper is concerned, we may summarise the on-going discussion on the formation of GC in two main points:  \\begin{itemize} \\item The mass of the ``polluting'' stars:  broadly speaking some groups advocate that the CNO cycling has been produced by intermediate mass AGB stars(Ventura et al. 2002, D' Antona et al. 2005), while others favour  massive rotating stars (Maeder and Meynet 2002, Decressin et al. 2007b). Each of these schemes has strong and weak points. In an attempt to summarise them,  we can say that  IM-AGB stars have the advantage of being 'natural' N  producers, able to also produce Li (and explain the behaviour of NGC 6397 through the so called Cameron-Fowler  mechanism; Cameron \\& Fowler, 1971), but having at present two main problems: the difficulty of reproducing the  observed yields for other elements (in particular Na), the fact that for NGC 6397 a natural conspiracy should have produced an ad hoc amount of Li,  exactly the same observed in the Spite plateau (Bonifacio et al. 2002). Finally,  the IM-AGB hypothesis would require a very peculiar IMF, heavily flat-topped (Prantzos and Charbonnel 2005, Chiappini et al. 2006). As far the rotating massive hypothesis is concerned, this scheme mitigates the IMF requirements, and can better explain some of the observed yields, although even the most extreme models cannot account for the observed high range of the Mg-Al correlation, unless some of the employed nuclear cross-sections are affected by large errors (Decressin et al. 2007b).  One important conclusion of both hypotheses is that the ``polluted'' stars should have a much higher  $^4$He abundance. Higher $^4$He abundance is claimed to be observed in one of the Main Sequences of $\\Omega$ Cen \\citep{piotto05} and of NGC 2808 \\citep{piotto2808}; it would naturally explain the blue extension of HB in NGC 6752 and other clusters (D'Antona et al. 2005). On the other hand, a substantially higher $^4$He abundance, should spectroscopically produce  an apparent higher abundance of heavier elements and higher spectroscopic gravity  in $^4$He rich stars (B\\\"ohm-Vitense 1979). A spread in He abundances, should bring about a spread in the abundances of {\\em all} elements, including iron-peak elements. The results of Gratton et al. (2001) in NGC 6397 and NGC 6752 exclude any significant spread in Fe abundances in either cluster.  \\item The second open issue is related to the overall mass of the structures containing the contaminants. Most analysis have so-far analysed self-pollution, that is,  that all the polluted material was   created within the cluster. New schemes consider the possibility of infall from the exterior (e.g. Bekki et al. 2007). In this scheme, for instance, GCs are the remnants of the core of much more massive dwarf  galaxies,  receiving processed gas from the whole galaxy, rather than from its own stars only. This scheme would greatly mitigate the quest for anomalous IMF, but , on the other hand, it opens the parameter space , since a new variable (the infall of material) is now added, and  it has not been  looked in detail as far as chemical evolution is concerned.  We shall finally recall that most attempts  aim to explain all the data observed in the different GC within a unique scheme. This is , of course, a very ambitious and valuable approach, but we would like to point out that the observational data show that different clusters present a different behaviour, and therefore, the attempt to unify all observations in a unique scheme might be premature.  \\end{itemize}    The measured abundances give us the possibility of studying the composition of  the polluting gas in some detail. For NGC 6752 we can safely assume that the two stars observed are at the extremes of the chemical distribution of the stars in the cluster.  In order to analyse pollution, Li is the best element to be considered, since it is destroyed at the temperatures where CNO cycling occurs. The difference of Li between the two stars is of a factor about 2.5.  If no Li production has occurred, this implies a fraction of  60$\\%$ of processed material in the contaminated star. The observed N enhancement would  imply a 25-fold N production in the polluting  gas. This can be assumed as a solid lower limit for the yields, in case no infall is present. In case the polluting gas is somewhat rich in Li the amount of polluting gas  would be larger than  60$\\%$ and, as a consequence, the overproduction of N in the previous generation  lower than what estimated above. We note that a 25-fold  N production is already at the upper limits  for  the current models of   massive stars. Detailed production predictions have been computed e.g. by Decressin et al. (2007b), who find that  massive stars can produce in their winds up to [N/Fe]$\\sim$ 1.5, or a factor 30 at most.  The constraints from oxygen are also relevant: the ratio in O abundance between star 4428 (O-rich) and the most O-poor stars observed in NGC 6752 is  of about 4.5 (Carretta et al. 2005), which would put the lower limit constraint of 22.5$\\%$ to the amount of pristine material, in case the processed material had no O at all.  The NGC 6752 data are consistent with a picture where the stars in this cluster   are formed by a mixture   of ``pristine'' and ``processed'' material.  The maximum observed ``pollution''   amounts to about 60 $\\%$ of the gas forming the stars. This would imply a polluting material  with no  Li, no  O, and N enhanced by 25 times (and Na  by a  similar amount). It may be difficult to produce in a closed system the very   high amount of N required by this scenario. The oxygen difference among the extreme stars suggests an even higher level of polluted material, which would require  some Li production, and  in this case the N (and Na) overabundance should be proportionally lower in the polluting gas.  The abundances of NGC 6397   cannot be accounted by a simple pollution picture, since  the basic variations in oxygen are not accompanied by similar variations in Li  (Bonifacio et al. 2002) and   in N (this paper)."
    },
    "0806/0806.2402_arXiv.txt": {
        "abstract": "In this letter we revisit the idea of reionization feedback on dwarf galaxy formation. We show that primordial minihalos with $v_{cir}<20$~km~s$^{-1}$ stop accreting gas after reionization, as it is usually assumed, but in virtue of their increasing concentration and the decreasing temperature of the intergalactic medium as redshift decreases below $z=3$, they have a late phase of gas accretion and possibly star formation. We expect that pre-reionization fossils that evolved on the outskirts of the Milky Way or in isolation show a bimodal star formation history with 12~Gyr old and $<10$~Gyr old population of stars. Leo~T fits with this scenario. Another prediction of the model is the possible existence of a population of gas rich minihalos that never formed stars. More work is needed to understand whether a subset of compact high-velocity clouds can be identified as such objects or whether an undiscovered population exists in the voids between galaxies. ",
        "introduction": "\\label{sec:introduction}  It is widely assumed that reionization of the intergalactic medium (IGM) suppresses gas accretion and galaxy formation in small mass halos with circular velocity $v_{cir} \\simlt 20$~km~s$^{-1}$, corresponding to a virial temperature $T_{vir}\\simlt 20,000$~K and to dark halo masses $M \\simlt 10^8-10^9$~M$_\\odot$ \\citep[\\eg,][]{Babul:92,Efstathiou:92b}. Hereafter, we will refer to these small halos affected by reionization feedback as ``minihalos''. The reason why gas accretion and star formation is suppressed in minihalos is that, after reionization, the Jeans mass in the IGM exceeds the mass of these minihalos and the IGM gas is unable to condense under the influence of their gravitational potential. Simulations of galaxy formation have confirmed this idea \\citep[\\eg,][]{Bullock:00, Gnedin:00b}. However, a recently discovered dwarf galaxy -- Leo T \\citep{Ryan-Weber:08} -- despite having an estimated mass of about $7 \\times 10^6$~M$_\\odot$ and having properties typical of other ultra-faint dwarf spheroidal (dSph) galaxies, contains gas and is actively forming stars at the present day.  This puzzling observation prompts taking a fresh look at the problem of gas accretion onto minihalos after reionization. However, the model presented in this study stands independently of Leo~T.  We show that the ability of dark minihalos to accrete gas from the IGM depends not only on the ratio of their circular velocity to the IGM sound speed, but also on their concentration, $c$. Typically, the concentration of a halo is $c_{vir} \\sim 5$ at the redshift of virialization but, as the halo evolves in the expanding universe, its concentration increases. The evolution of the halo concentration with redshift can be understood in the context of the theory of cosmological secondary infall of dark matter \\citep{Bertschinger:85} and has been quantified using N-body simulations \\citep{Bullock:01, WechslerB:02}.  In this letter we show that primordial minihalos with $v_{cir}<20$~km~s$^{-1}$ stop accreting gas after reionization, as expected, but in virtue of their increasing concentration and the decreasing temperature of the IGM at $z<3$, they start accreting gas from the IGM at later times. As a result, we expect that pre-reionization fossils \\citep[][hereafter RG05]{RicottiG:05} in the Local Group have a more complex star formation history than previously envisioned. A signature of this model is a bimodal star formation history with an old ($\\sim 12$~Gyr) and a younger ($\\simlt 5-10$~Gyr, depending on the minihalo mass) population of stars. Another prediction of the model is the possible existence of minihalos containing only gas but no stars. This would revive the suggestion that some compact high-velocity clouds (CHVCs) are extragalactic objects \\citep{Blitz:99}. However, the bulk of these objects may be yet undiscovered due to their faint 21~cm and H-$\\alpha$ fluxes. The ongoing survey ALFALFA \\citep{Giovanelli:05, Giovanelli:07} may be able to discover this new population of extragalactic objects.  This letter is structured as follows. In \\S~\\ref{sec:dm} and \\S~\\ref{sec:gas} we describe our model for dark matter and gas accretion and in \\S~\\ref{sec:discussion} we apply the model to the interpretation of Leo~T and observations of CHVCs. Throughout the rest of the paper we use the following cosmological parameters ($h=0.72, \\Omega_{dm}=0.214 , \\Omega_b=0.0438 , \\sigma_8=0.796, n_s=0.963$) from WMAP5 \\citep{Dunkley:08}. ",
        "conclusions": "\\label{sec:discussion}  Although Leo~T contains gas, it properties as similar to those of other ultra-faint dSph galaxies and quite different from those of more massive dwarf irregular galaxies \\citep{BovillR:08}. Its properties are consistent with those of pre-ionization fossils but, in the context of this model, it is difficult to understand how it could hold on to its gas.  We point out that dark minihalos with large concentration are able to accrete gas from the IGM more efficiently than lower concentration halos of the same mass. The concentration of a halo depends on the redshift of virialization as $c(z) \\propto (1+z_{vir})/(1+z)$. Thus, as a result of their growing concentration, early forming minihalos, although affected by reionization feedback, may have a late phase of cold gas accretion from the IGM and form stars (perhaps for the first time) at $z<1$.  This model explains all the observed properties of Leo~T.  Having a distance of $420$~kpc from the Milky Way, Leo~T is just starting to fall into the Milky Way halo, thus it is likely to have evolved in isolation (fulfilling a requirement of the model) and unlikely to have been tidally stripped. Recent observations suggest that Leo~T has a bimodal SFH \\citep{DeJong:08} with a $>12$~Gyr population and a $9$~Gyr population (with SF continuing until few Myrs ago), in agreement with our expectation of a late phase of accretion from the IGM. The bimodality of the SFH should be most pronounced in minihalos with the lowest $v_{cir}$, gradually disappearing as $v_{cir}$ increases toward the $ 20$~km~s$^{-1}$ threshold for pre-reionization fossils (RG05). Leo~T has a gas velocity dispersion $\\sim 7$~km~s$^{-1}$, that, according to our model, implies a dark halo with maximum circular velocity of $11$~km~s$^{-1}$ (Eq.~\\ref{eq:vcir_c}).  Leo~T has a core radius $r_{c} \\sim 100$~pc, $M_{dyn}/M_g \\sim 8-10$ \\citep{SimonG:07, Ryan-Weber:08} and maximum \\HI~ column density $N_H=7\\times 10^{20}$~cm$^{-2}$. These values are all in agreement with our model as shown in \\S~\\ref{sec:gas} and right panel in Fig.~\\ref{fig:core2}. If future observations will show that the dark halo of Leo~T has $v_{cir}^{max} >20$~km~s$^{-1}$ or its SFH does not show any decline after $12$~Gyr ago, our model for Leo~T will be disproved.  Minihalos smaller than the one hosting Leo~T can also accrete gas from the IGM at $z<1$, but may not be able to form stars.  Sufficiently small minihalos are likely to be dark and, if they evolved in isolation, the IGM around them would be nearly metal-free. Cosmological simulations show that the fraction of primordial minihalos that host a luminous galaxy decreases steeply with decreasing $v_{cir}$: $f(lum) \\sim min[1, 50\\%(v_{cir}/10~{\\rm km/s})^{4.6}]$ \\citep{RicottiGS:08}.  A fraction of dark minihalos may be able to accrete gas from the IGM but, if the gas density in their core is below the threshold required for star formation, these objects will resemble CHVCs.  Additional work is needed to determine whether our model is fully consistent with the properties of known CHVCs \\citep{Blitz:99, Braun:99, Robishaw:02}, however the typical column densities we predict in Fig.~\\ref{fig:core2}(left panel) are in agreement with a subset of existing observations.  The typical size and size distribution of CHVCs may agree with the model assuming a mean distance of minihalos of 1~Mpc. We estimate that realistic extension of the diffuse gas in dark minihalos are $\\simlt 5-10$~kpc. This may exclude several known CHVC and/or set an upper limit for their distance from the Milky Way. More extended extragalactic CHVCs are possible for minihalo masses $M>10^8$~M$_\\odot$ that are likely to be luminous. In addition this scenario for more massive minihalos has already been investigated by previous studies and found not fully satisfactory \\citep{Sternberg:02, Maloney:03, Putman:03}.  More work in this area is also motivated by the ongoing survey ALFALFA \\citep{Giovanelli:05, Giovanelli:07} that may be able to discover a even fainter population of extragalactic CHVCs.  Hence, future work shall focus on estimating the number, flux and size distribution of gas rich minihalos in the Local Volume and the local voids, producing synthetic maps for 21~cm and H-$\\alpha$ wavelengths.  \\appendix"
    },
    "0806/0806.0294_arXiv.txt": {
        "abstract": "We present the discovery of fifteen new T2.5-T7.5 dwarfs (with estimated distances between $\\sim$24--93pc), identified in the first three main data releases of the UKIRT Infrared Deep Sky Survey. This brings the total number of T dwarfs discovered in the Large Area Survey (to date) to 28. These discoveries are confirmed by near infrared spectroscopy, from which we derive spectral types on the unified scheme of Burgasser et al. (2006). Seven of the new T dwarfs have spectral types of T2.5-T4.5, five have spectral types of T5-T5.5, one is a T6.5p, and two are T7-7.5. We assess spectral morphology and colours to identify T dwarfs in our sample that may have non-typical physical properties (by comparison to solar neighbourhood populations), and find that one of these new T dwarfs may be metal poor, three may have low surface gravity, and one may have high surface gravity. The colours of the full sample of LAS T dwarfs show a possible trend to bluer $Y-J$ with decreasing effective temperature, and some interesting colour changes in $J-H$ and $z-J$ (deserving further investigation) beyond T8. The LAS T dwarf sample from the first and second main data releases show good evidence for a consistent level of completion to J=19. By accounting for the main sources of incompleteness (selection, follow-up and spatial) as well as the effects of unresolved binarity and Malmquist bias, we estimate that there are 17$\\pm$4 $\\ge$T4 dwarfs in the $J\\le$19 volume of the LAS second data release. Comparing this to theoretical predictions is most consistent with a sub-stellar mass function exponent $\\alpha$ between -1.0 and 0. This is consistent with the latest 2MASS/SDSS constraint (which is based on lower number statistics), and is significantly lower than the $\\alpha\\sim1.0$ suggested by L dwarf field populations, possibly a result of the lower mass range probed by the T dwarf class. ",
        "introduction": "The advent of the large scale Sloan Digital Sky Survey (SDSS; York et al. 2000), the 2-Micron All Sky Survey (2MASS; Skrutskie et al. 2006) and Deep Near Infrared Survey of the Southern Sky (DENIS; Epchtein et al. 1997) were the main factors that led to the identification and study of the two spectral classes beyond M; the L dwarfs (with effective temperature $T_{\\rm eff}\\sim$2300-1450K; Golimowski et al. 2004) have dusty atmospheres and very red near infrared colours, while the even cooler T dwarfs ($T_{\\rm eff}<$1450K) have clear atmospheres (from which the dust has settled), and their blue infrared spectra are dominated by strong CH$_4$ and H$_2$O absorption bands. The spectral typing scheme of Kirkpatrick et al. (1999) established the first properly classified L dwarf population via 2MASS discoveries. The first T dwarf was discovered as a close companion to the early M dwarf GL229 (Nakajima et al. 1995), and the first L dwarf as a companion to the white dwarf GD 165 (Becklin \\& Zuckerman 1988; Kirkpatrick et al. 1993). However, the subsequent growth of the known T (and L) dwarf populations ($>$500 L dwarfs and $>$100 T dwarfs are known at time of writing; Kirkpatrick 2005) has been predominantly achieved via searches of the SDSS (e.g. Leggett et al. 2000; Geballe et al. 2002; Knapp et al. 2004; Chiu et al. 2006) and 2MASS (e.g. Kirkpatrick et al. 2000; Cruz et al. 2007; Burgasser et al. 2002; 2004; Tinney et al. 2005; Looper Kirkpatrick \\& Burgasser 2007) databases.  The coolest $T_{\\rm eff}$ probed by SDSS and 2MASS are currently defined by the eight T7.5-8 dwarfs (as typed by the Burgasser et al. (2006b; B06) scheme) discovered in these surveys. These objects have $T_{\\rm eff}$ in the range 725-950K (Geballe et al. 2001; Saumon et al. 2006; 2007; Leggett et al. 2007; Burgasser, Burrows \\& Kirkpatrick 2006). At lower $T_{\\rm eff}$ it is possible that a new spectral class (pre-emptively called Y, following Kirkpatrick et al. (1999)) may be necessary, if for example, spectral changes occur such as the strengthening of ammonia absorption, and/or the condensation of atmospheric water clouds (Burrows, Sudarsky \\& Lunine (2003)). L and T dwarfs already encompass the temperature scale of transiting ``hot Jupiters'' (e.g. Knutson et al. 2007), and populating the even cooler $T_{\\rm eff}$ regime will allow us to study and understand atmospheres whose $T_{\\rm eff}$ is similar to cooler extra-solar giant planets populations.  The current complement of known L and T dwarfs allows some constraints to be placed on the substellar mass function (e.g. Chabrier 2003 and references therein; Allen et al. 2005; Metchev et al. 2007). However, if the brown dwarf mass function is to be accurately constrained, a significantly larger number of late T dwarfs will be extremely beneficial, since the $T_{\\rm eff}$ distribution of T dwarfs in the $<$950K range is particularly sensitive to mass function variations (e.g. fig 5 of Burgasser 2004). Also, accurate constraints on the brown dwarf birth rate (i.e. their formation history) need the improved statistics that come with larger numbers of both L and T dwarfs.  The UKIRT (UK Infrared Telescope) Infrared Deep Sky Survey (UKIDSS; Lawrence et al. 2007) is a new infrared survey being conducted with the UKIRT Wide Field Camera (WFCAM; Casali et al. 2007). UKIDSS is a set of five sub-surveys, with three wide field surveys -- the Large Area Survey (LAS), the Galactic Cluster Survey (GCS) and the Galactic Plane Survey (GPS) -- and two very deep surveys -- the Deep Extra-galactic Survey (DXS) and the Ultra-Deep Survey (UDS). The LAS is the largest of the wide-field surveys, and will cover 4000 sq degs of sky in four filter bands, going several magnitudes deeper than 2MASS. UKIDSS began in May 2005, and at the time of writing there have been four ESO (European Southern Observatory)-wide releases, including an Early Data Release (EDR) in February 2006 (Dye et al. 2006), and three subsequent main data releases: Data Release 1 (DR1) in 2006 July covering 190 sq degs (Warren et al. 2007b), Data Release 2 (DR2) in 2007 March covering 280 sq degs (including DR1; Warren et al. 2007c), and Data Release 3 (DR3) in 2007 December, covering 900 sq degs (including DR1 and DR2).  These data are providing un-paralleled sensitivity to L and T dwarf populations, and our collaboration has begun a variety of UKIDSS-based searches for these objects. Here we focus on our search for late T dwarfs and potentially new record breaking low-$T_{\\rm eff}$ objects. Previous LAS T dwarfs have been presented by Kendall et al. (2007), Lodieu et al. (2007), Warren et al. (2007a) and Chiu et al. (2008), who have discovered a total of 13 spectroscopically confirmed LAS T dwarfs from the EDR and DR1. This included the first T8.5 dwarf ULAS J0034-0052 (Warren et al. 2007a).  Here we report the discovery of 15 new LAS T dwarfs discovered in DR1, DR2 and some of DR3. Section 2 summarises our selection criteria for identifying candidate T dwarfs as well as possible 400-700K objects. Section 3 describes the follow-up photometry we have obtained to identify spurious objects amongst these candidates, and Section 4 presents our spectroscopic follow-up and confirmation of the 15 new T dwarfs. Section 5 discusses the spectral morphology and colour of the new T dwarfs and how these could relate to their physical properties, and Section 6 presents updated constraints on the size of the LAS T dwarf population, compared to theoretical predictions. Section 7 discusses some future work and gives our conclusions. ",
        "conclusions": "We have discussed our database selection methods and the photometric follow-up that we performed to search for late T dwarfs and even cooler $T_{\\rm eff}$=400--700K objects in the UKIDSS LAS. These techniques have allowed us to make essentially complete follow-up of our candidates from DR1 and DR2 down to $J$=19. We have also followed up some candidates from DR3 and a fraction of our fainter candidates from DR1 and DR2 (to J=19.5). Using a variety of spectroscopic facilities we have measured the spectra of our best candidates, and have spectroscopically confirmed 15 new T dwarfs, bringing the total of number of confirmed LAS T dwarf discoveries to 28.  Compared to typical T dwarf properties, one of our new T dwarfs may be metal-poor, two may have relatively low surface gravity, one may have higher than normal surface gravity, and one may be metal rich or have low surface gravity. These assessments are based on comparisons between the spectral morphology and colours of the T dwarfs with theoretical models trends, and are thus somewhat speculative. However, all indications are that T dwarf spectra are quite sensitive to their physical properties, and it seems clear that a better understanding of these variations would have important implications for our ability to model T dwarf atmospheres, and potentially constrain (via spectroscopy) the mass, age and composition of T dwarf populations in the field (e.g. Pinfield et al. 2006). To this end, it would be very beneficial to identify T dwarfs whose properties could be constrained without the need to study their spectra. Both T dwarfs and L dwarfs of this type may be found in widely separated binary systems (e.g. Gizis et al. 2001; Wilson et al. 2001; Pinfield et al. 2006; Scholz et al. 2003; Burgasser et al. 2000; Luhman et al. 2007) for which the primary object can be used to constrain system age and composition. Gl570D and HD3651B (Burgasser 2000 and Luhman et al. 2007) are good examples, being analysed for age and abundance by Geballe et al. (2001), Burgasser (2007) and Liu Leggett \\& Chiu (2007).  The large volume (and large number of T dwarfs) probed by the LAS should yield many T dwarfs in such binary systems, and in this context we made a basic search for possible companions to all confirmed LAS T dwarfs out to a separation of $\\sim$10,000AU (at the estimated distances in Table 5), by querying the Simbad astronomical database (operated at CDS). We searched for neighboring objects with either a measured spectral type, a parallax, or a high proper motion, for which we could obtain distance constraints from either parallax or from spectral type and photometry. These distance constraints allowed us to rule out (as possible companions) all the neighboring objects, by comparison with the distance constraints of the T dwarfs. However, proper motion measurements for all LAS T dwarfs would facilitate a more general search for common proper motion companions as the number of LAS T dwarfs grows, which could include white dwarf companions that can yield useful age constraints (e.g. Day-Jones et al. 2008) and are readily identified in SDSS (e.g. Kleinman et al. 2004; Gates et al. 2004; Eisenstein et al. 2006).  The number of blue ($J-H<$0.1) LAS T dwarfs discovered with our methods shows statistical evidence for a good level of completeness down to $J$=19. This builds on our previous work (e.g. Lodieu et al. 2007) for which a good level of completion was only achieved to $J\\sim$18.5. The increased size of the LAS T dwarf sample allowed us to place some statistical constraints on the sub-stellar mass function. Indeed, there may be mounting evidence for a steepening (decreasing more rapidly with decreasing mass) sub-stellar mass function in the field, with best-fit $\\alpha$ values of $\\sim 1$ and between -1.0 and 0.0 for L and T dwarf populations respectively. In any event, mass function constraints should be greatly improved by larger samples of LAS T dwarfs as the survey continues to grow. In addition, constraints on the brown dwarf birth-rate should be attainable by building on the LAS L and early T dwarf searches mentioned briefly in Section 2.1 (e.g. fig 3 of Deacon \\& Hambly; fig 10 of Burgasser 2004). For instance, a sample of $\\sim$100 L/T dwarfs per $\\Delta T_{\\rm eff}$=100K range from 1100-1500K (assuming a flat birth-rate) would be capable of ruling out an exponential (e.g. $\\tau_g\\sim$5Gyr) birth-rate at a $\\sim$5$\\sigma$ level of significance. An appropriate sample of $\\sim$400 late L to mid T dwarfs could be identified in the full LAS to $J$=18.5 (for a flat birth-rate).  The discovery of the T8.5 dwarf ULAS J0034-0052 (Warren et al. 2007) in the LAS absolutely demonstrates that UKIDSS is probing unexplored $T_{\\rm eff}$ regimes beyond the previously known late T dwarfs. Having said this, parallax and adaptive optics measurements for this and other very late T dwarfs are important if we are to test for unresolved binarity and better constrain luminosity and $T_{\\rm eff}$. When comparing the colours of ULAS J0034-0052 with all available LAS T dwarfs, it is possible that a blue trend in $Y-J$ with decreasing T$_{\\rm eff}$ may be present. The $J-H$ and $z-J$ colours might also change beyond T8, becoming slightly redder and bluer respectively, although this needs further investigation. These spectral changes would be of great importance not only to our understanding of cool dwarf atmospheres, but also to the potential for finding such objects in the LAS. If it is the case that $T_{\\rm eff}$=400-700K objects become bluer in $Y-J$ and redder in $J-H$ than T8 dwarfs, it could be less problematic to identify such candidates to greater $J-$band depth in multi-band searches of the combined LAS and SDSS databases. Indeed, an evolving set of search criteria takes full advantage of the capabilities offered by large general surveys like UKIDSS and SDSS."
    },
    "0806/0806.2882_arXiv.txt": {
        "abstract": "We conduct a systematic survey of the regions in which distant satellites can orbit stably around the four giant planets in the solar system, using orbital integrations of up to $10^9$ yr. In contrast to previous investigations, we use a grid of initial conditions on a surface of section to explore phase space uniformly inside and outside the planet's Hill sphere (radius $r_{\\rm H}$; satellites outside the Hill sphere sometimes are also known as quasi-satellites). Our confirmations and extensions of old results and new findings include the following: (i) many prograde and retrograde satellites can survive out to radii $\\sim 0.5r_{\\rm H}$ and $\\sim 0.7r_{\\rm H}$, respectively, while some coplanar retrograde satellites of Jupiter and Neptune can survive out to $\\sim r_{\\rm H}$; (ii) stable orbits do not exist within the Hill sphere at high ecliptic inclinations when the semi-major axis is large enough that the solar tide is the dominant non-Keplerian perturbation; (iii) there is a gap between $\\sim r_{\\rm H}$ and $2r_{\\rm H}$ in which no stable orbits exist; (iv) at distances $\\gtrsim 2r_{\\rm H}$ stable satellite orbits exist around Jupiter, Uranus and Neptune (but not Saturn). For Uranus and Neptune, in particular, stable orbits are found at distances as large as $\\sim 10r_{\\rm H}$; (v) the differences in the stable zones beyond the Hill sphere arise mainly from differences in the planet/Sun mass ratio and perturbations from other planets; in particular, the absence of stable satellites around Saturn is mainly due to perturbations from Jupiter.  It is therefore likely that satellites at distances $\\gtrsim 2r_{\\rm H}$ could survive for the lifetime of the solar system around Uranus, Neptune, and perhaps Jupiter. ",
        "introduction": "\\label{sec:intro}  Most of the satellites of the four giant planets in the solar system can be divided into two groups, usually called the regular and irregular satellites. Regular satellites orbit close to the planet (within $\\sim 0.05r_{\\rm H}$, where $r_{\\rm H}$ is the Hill radius\\footnote{The Hill radius is defined as $r_{\\rm H}=a_p(\\mu/3)^{1/3}$, where $a_p$ is the semi-major axis of the planet orbit and $\\mu\\equiv m_p/(m_p+M_{\\odot})$ with $m_p$ the planet mass.}), and move on nearly circular, prograde orbits that lie close to the planetary equator. Irregular satellites are found at distances $\\sim 0.05r_{\\rm H}-0.6r_{\\rm H}$, with large orbital eccentricities and inclinations, on both prograde and retrograde orbits. An alternative division between regular and irregular satellites is given by the critical semi-major axis (e.g., Goldreich 1966; Burns 1986), $a_{\\rm crit}=(2\\mu J_2R^2a_p^3)^{1/5}$; those with $a>a_{\\rm crit}$ are classified as irregular satellites. Here $J_2$ is the planet's second zonal harmonic coefficient (augmented by any contribution from the inner regular satellites) and $R$ is the planet's radius. This critical radius marks the location where the precession of the satellite's orbital plane is dominated by the Sun rather than by the planet's oblateness. The current number ratios of irregular to regular satellites are $55/8$ for Jupiter, $35/21$ for Saturn, $9/18$ for Uranus, and $7/6$ for Neptune (e.g., Jewitt \\& Haghighipour 2007). The regular satellites are likely to have formed within a circumplanetary disk of gas and solid bodies. The kinematic differences between regular and irregular satellites suggest that the latter must have formed through a quite different mechanism, most likely capture from the circumstellar disk (for a recent review, see Jewitt \\& Haghighipour 2007).  The search for irregular satellites of the giant planets has been fruitful in recent years, owing mainly to modern high-sensitivity, large-scale CCDs (e.g., Gladman et al.\\ 1998, 2000, 2001; Holman et al.\\ 2004; Kavelaars et al.\\ 2004; Sheppard \\& Jewitt 2003; Sheppard et al.\\ 2005, 2006). Up to 2007, 106 irregular satellites of the giant planets had been discovered, compared to 53 regular satellites. Two features stand out in the distributions of orbital parameters of these irregular satellites. First, retrograde irregular satellites extend to larger semi-major axes than prograde ones ($\\sim 0.6r_{\\rm H}$ compared to $\\sim 0.4r_{\\rm H}$); second, satellites with orbital inclination in the range $\\sim 60^\\circ-130^\\circ$ relative to the ecliptic are absent.  A number of authors have shown that these features can be explained reasonably well by the requirement that the satellite orbits be stable. H\\'{e}non (1969, 1970) studied the planar circular restricted three-body problem in Hill's (1886) approximation, where the mass ratio $\\mu\\rightarrow 0$ while the radii of interest shrink to zero as $\\mu^{1/3}$. He showed that prograde satellite orbits are stable up to a mean distance from the planet $\\sim 0.4r_{\\rm H}$, while retrograde satellite orbits can be stable at much larger distances from the planet. Thus it is not surprising that retrograde satellites are found at larger distances than prograde ones. Hamilton \\& Krivov (1997) studied the dynamics of distant satellites of asteroids in heliocentric orbits using a ``generalized Tisserand constant'' and, among other conclusions, confirmed that retrograde orbits are more stable than prograde ones. Carruba et al.\\ (2002) used a combination of analytic arguments and numerical integrations to show that high-inclination orbits inside the Hill sphere exhibit large eccentricity oscillations (Kozai oscillations; Kozai 1962) due to secular solar perturbations. They found that orbits with inclinations (relative to the planetary orbital plane) between $55^\\circ$ and $130^\\circ$ are generally unstable, thus explaining the absence of irregular satellites on high-inclination orbits. Nesvorn\\'{y} et al.\\ (2003) performed detailed orbital integrations of the four giant planets plus a grid of test-particle satellites for intervals of $10^6$--$10^8$ yr. They confirmed that retrograde satellites can be stable at larger radii than prograde ones, and that highly inclined orbits are unstable. They argued that the largest semi-major axes at which satellites of the four giant planets could survive for times comparable to the lifetime of the solar system were $\\sim 0.7r_{\\rm H}$ for retrograde satellites and $\\sim 0.4r_{\\rm H}$ for prograde ones, and that these upper limits were achieved only for nearly circular orbits close to the plane of the ecliptic.  Other authors have examined the possibility that stable satellite orbits exist with mean distance from the planet $\\gtrsim r_{\\rm H}$. In the planetocentric frame, the dominant force on such satellites is due to the Sun, rather than the planet\\footnote{Hence these are sometimes called ``quasi-satellites'' (Lidov \\& Vashkov'yak 1994a,b; Mikkola \\& Innanen 1997).}. Nevertheless, the satellite remains close to the planet because it is in a 1:1 resonance in the sense that its heliocentric mean longitude librates around that of the planet; the resulting orbit relative to the planet is a retrograde ellipse with axis ratio 2:1, the short axis pointing towards the Sun, and synodic period equal to the planet's orbital period. The analytical theory of such orbits is described by Jackson (1913), Lidov \\& Vashkov'yak (1994a,b), Mikkola \\& Innanen (1997), Namouni (1999), Mikkola et al.\\ (2006), and others. H\\'{e}non's (1970) numerical analysis of Hill's approximation to the planar circular restricted three-body problem suggests that stable retrograde satellites can exist at arbitrarily large distance from the planet. Benest (1971) confirmed that stable retrograde orbits at large distances persist in the elliptic restricted three-body problem, where the mass ratio and eccentricity were chosen to match those of Jupiter. Wiegert et al.\\ (2000) demonstrated that retrograde satellites of Uranus and Neptune could be stable for up to $10^9$ yr at distances up to $\\sim 10r_{\\rm H}$, suggesting that primordial objects of this type could still exist in the solar system although none are currently known.  Despite the number and quality of these investigations, there are several unanswered questions that lead us to revisit the problem of orbital stability of satellites at large distances from the host planet. (i) Wiegert et al.\\ (2000) found stable satellite orbits beyond the Hill radius only for Uranus and Neptune, not Jupiter or Saturn. What is the reason for this difference? The possibilities include differences in the planetary masses and orbital eccentricities, or different perturbations from neighboring planets. (ii) Wiegert et al.\\ (2000) explored orbits outside the Hill radius, while Nesvorn\\'y et al.\\ (2003) explored orbits inside the Hill radius (indeed, in the former paper the integrations were terminated when the particles entered the Hill sphere of radius $r_{\\rm H}$ around the planet, while in the latter paper the integrations were terminated when the particles exited the Hill sphere). Are there stable satellite orbits that cross the Hill sphere? (iii) As we shall describe further in \\S\\ref{sec:method}, the grids of initial conditions used by Nesvorn\\'y et al.\\ and Wiegert et al.\\ do not provide a complete exploration of the phase space in which stable satellite orbits exist.  The primary goal of this paper is to map out the entire stability region in phase space---both inside and outside the Hill sphere---in which satellite orbits that can survive around the four giant planets for times comparable to the age of the solar system (our main integrations last for up to 100 Myr). We describe our setup in \\S\\ref{sec:method}, and present the results in \\S\\ref{sec:result}. We conclude and discuss our results in \\S\\ref{sec:con}.  Following Fabrycky (2008), we shall define a ``satellite'' of a planet to be a small body whose distance from the planet never exceeds the semi-major axis of the planet, $a_p$. This definition excludes bodies on Trojan orbits around the triangular Lagrange points, bodies on horseshoe orbits, and objects such as asteroid 2003 YN107 (Connors et al.\\ 2004), which oscillates between a horseshoe orbit and an orbit centered on Earth. This definition seems simple and reasonable to us, but other definitions are common in the literature. Many authors define ``satellite'' to be an object that always remains within the Hill sphere of the planet or whose Jacobi constant constrains it to remain within the last closed zero-velocity surface around the planet. Benest (1971) defines a satellite to be a body whose heliocentric orbital frequency is the same as the planet's, but whose synodic frequency around the planet is non-zero. Wiegert et al.\\ use the term ``quasi-satellite'' for an object that remains outside the Hill sphere but whose heliocentric longitude difference from the planet never exceeds $120^\\circ$ and regularly passes through zero. However, the term ``quasi-satellite'' is confusing because it is also used for objects such as 2003 YN107 that spend part of their time on horseshoe orbits and thus are only temporarily satellites in our sense. ",
        "conclusions": "\\label{sec:con}  We have conducted a systematic survey of the stable regions of satellites around giant planets in the solar system, using numerical orbital integrations that include gravitational perturbations from the other planets. We confirm previous results for satellites within the Hill sphere: stable retrograde satellites can exist further out than prograde satellites (e.g., H\\'{e}non 1970; Hamilton \\& Krivov 1997; Nesvorn\\'{y} et al.\\ 2003); and stable orbits cannot exist at high inclinations (e.g., Carruba et al.\\ 2002; Nesvorn\\'{y} et al.\\ 2003).  We also confirm and extend the conclusions of Wiegert et al.\\ (2000) that distant retrograde satellites (``retrograde'' as defined in the rotating frame) can survive well beyond the Hill sphere for at least $10^8$--$10^9$ yr, and probably for the lifetime of the solar system. Uranus and Neptune are the most promising host planets for such distant satellites, since their stability regions have the largest extent (e.g., Fig.\\ \\ref{fig:stability_region_Uranus}-\\ref{fig:stability_region_Neptune}). Jupiter has a smaller stable region (Fig.\\ \\ref{fig:stability_region_Jupiter}), and Saturn appears to have no stable regions beyond the Hill radius (Fig.\\ \\ref{fig:stability_region_Saturn}).  Remarkably, there is a gap between the inner and outer stability zones for retrograde satellites, extending from about $r_{\\rm H}$ to $2r_{\\rm H}$, in which almost no stable orbits exist.  To check whether any of the proposed distant satellites have already been discovered as Centaurs, we take the positions and velocities of the known Centaurs from the IAU Minor Planet Center\\footnote{http://cfa-www.harvard.edu/iau/mpc.html} that have planetocentric distances smaller than the semi-major axes of each of the four giant planets at the last observed epoch. There are 31 Centaurs (Jupiter 1; Uranus 16; Neptune 14) that satisfy the criterion. We numerically integrated these objects along with the Sun and giant planets for $10^8$ yr, but none of them survived as satellites according to the definition in \\S\\ref{sec:intro}.  Searches for satellites far beyond the Hill radius can be carried out either with dedicated deep, wide-angle surveys around the giant planets, or through all-sky surveys such as Pan-STARRS and LSST. The most promising search areas are close to the orbital plane of the planet, since only low-inclination orbits survive (e.g., compare Fig.\\ \\ref{fig:other_henon2d} and Fig.\\ \\ref{fig:uranus_henon_3d}). In Fig.\\ \\ref{fig:obs_ang_vel} (left panel) we show the heliocentric angular velocity of the stable outer retrograde satellites as a function of angular distance from the planet as viewed from the Sun, for Jupiter (black), Uranus (blue) and Neptune (red) respectively, sampled every Myr between $5\\times 10^7$ and $10^8$ yr.   In the middle and right panels of Fig.\\ \\ref{fig:obs_ang_vel} we show histograms of the difference in heliocentric semi-major axis from their host planet and heliocentric eccentricities for the stable outer retrograde satellites. These distributions can be used to cull a large sample for potential satellites. Once a candidate is identified with reliable orbit elements, a long-term orbital integration should be run to confirm its satellite nature.  The discovery and characterization of satellites beyond the Hill sphere would provide rich information about the early formation of the solar system. Fabrycky (2008; also see Kortenkamp 2005) recently performed simulations of capture of neighboring planetesimals from the circumstellar disk during slow planet growth, and found that such distant satellites are a natural outcome for Uranus and Neptune. Thus an inventory of this potential population of bodies would enhance our understanding of the formation of planets and their satellites in the early solar system, and the properties of the primordial planetesimal disk."
    },
    "0806/0806.0625_arXiv.txt": {
        "abstract": "We investigate the properties of quiescent and star--forming galaxy populations to $z\\sim 2$ with purely photometric data, employing a novel rest--frame color selection technique.  From the UKIDSS Ultra--Deep Survey Data Release 1, with matched optical and mid--IR photometry taken from the Subaru XMM Deep Survey and Spitzer Wide--Area Infrared Extragalactic Survey respectively, we construct a $K$--selected galaxy catalog and calculate photometric redshifts.  Excluding stars, objects with uncertain $\\zphot$ solutions, those that fall in bad or incomplete survey regions, and those for which reliable rest--frame colors could not be derived, 30108 galaxies with $K<22.4$ (AB) and $\\zphot\\le 2.5$ remain.  The galaxies in this sample are found to occupy two distinct populations in the rest--frame $U-V$ vs.~$V-J$ color space: a clump of red, quiescent galaxies (analogous to the red sequence) and a track of star--forming galaxies extending from blue to red $U-V$ colors. This bimodal behavior is seen up to $z\\sim 2$.  Due to a combination of measurement errors and passive evolution, the color--color diagram is not suitable to distinguish the galaxy bimodality at $z>2$ for this sample, but we show that MIPS 24$\\mu$m data suggest that a significant population of quiescent galaxies exists even at these higher redshifts.  At $z=1-2$, the most luminous objects in the sample are divided roughly equally between star--forming and quiescent galaxies, while at lower redshifts most of the brightest galaxies are quiescent.  Moreover, quiescent galaxies at these redshifts are clustered more strongly than those actively forming stars, indicating that galaxies with early--quenched star formation may occupy more massive host dark matter halos.  This suggests that the end of star formation is associated with, and perhaps brought about by, a mechanism related to halo mass. ",
        "introduction": "The stellar mass in the local universe lies largely within two distinct classes of galaxies: actively star--forming spirals, and massive ellipticals with evolved stellar populations and little current star formation.  These populations show strongly bimodal behavior in a number of measured and derived quantities, including color \\citep{baldry04}, the 4000\\AA--break strength $D_n(4000)$ \\citep[an indicator of stellar population age;][]{kauffmann03}, and clustering strength \\citep{budavari03}. Relatively low--mass galaxies with high star--formation rates have also been found at high redshifts \\citep[e.g., $z>3$ Lyman--break galaxies or LBGs;][]{steidel96}. Surprisingly, the population of massive, ``dead'' galaxies appears to persist to high redshift as well \\citep[e.g.][]{labbe05,daddi05,kriek06,kriek08b}, even though the universe was only a few Gyr old at that point.  Since the stellar populations of evolved galaxies often have ages that are a significant fraction of the age of the universe at these redshifts, the implication is that they must have formed (and their star formation ceased) far earlier than anticipated by standard hierarchical structure formation models.  Investigating the evolution of star formation activity and the assembly of stellar mass over the past $\\sim 12$\\,Gyr are therefore two key avenues to understanding how the present galaxy population came to be, and large samples of galaxies during this peak formation epoch ($z\\sim 2$) are needed for this. As spectroscopy of faint objects is extremely expensive with regard to telescope time, several highly efficient broadband selection techniques have been designed to weed out low--redshift interlopers.  In addition to the aforementioned LBG selection, which primarily finds unobscured star--forming galaxies at $z>3$, the distant red galaxy (DRG) criterion \\citep[$J-K>2.3$ (Vega) or $>1.34$ (AB);][]{franx03} instead tends to find the most {\\it massive} galaxies at $2<z<3$ \\citep{vdokkum06}, about half of which show signs of heavily--obscured star formation \\citep{papovich06}. Other methods, such as the $BzK$ selection of \\citet{daddi04}, are adept at selecting $z\\sim 2$ galaxies over a wide range of masses and star--formation rates.  Each of these techniques has its advantages and disadvantages, but one point has become clear: to effectively study the mass evolution of high--redshift galaxies, deep near--infrared data are crucial.  The observed $JHK$ bands trace rest--frame optical light at $z\\sim 2-3$ and provide a more reliable indicator of stellar mass than the rest--frame UV (observed optical); indeed, DRGs themselves appear to represent an important fraction, if not the majority, of stellar mass at high redshifts \\citep{rudnick06,grazian07}. Such galaxies are typically faint in the observed optical bands, either from old stellar populations or dust obscuration, and are often missed by purely optical selection techniques \\citep{quadri07b}. Selecting galaxies from deep near--infrared fields is therefore likely to yield a more complete picture of the stellar mass at high redshift than similar studies at optical wavelengths.  Until recently, the small sizes (at most a few square arcminutes) and relatively low efficiencies of the available infrared detectors made large, deep near--IR surveys impractical.  Such projects therefore typically followed the ``pencil--beam'' approach, surveying a single frame to high sensitivity \\citep[e.g., the {\\it Faint Infrared Extragalactic Survey, FIRES};][]{franx03}, or alternatively observing somewhat larger areas at the expense of depth \\citep[e.g., GOODS--S;][]{giavalisco04}.  While such surveys have revealed a great deal about the high--redshift universe, the necessarily small survey volumes proved problematic for statistical studies of galaxy populations (as a result of both cosmic variance and small--number statistics).  Large--format  near--IR detectors on 4m--class telescopes (such as ISPI, WFCAM, and the upcoming VISTA camera) now provide the ability to conduct deep surveys over wide areas, and a number of past, present, and planned projects take advantage of this capability -- examples include the Multiwavelength Survey by Yale--Chile \\citep[MUSYC;][]{gawiser06,quadri07b}, the UKIRT Infrared Deep Sky Survey \\citep[UKIDSS;][]{lawrence07}, and ULTRA--VISTA.  Until this latter survey is complete, the UKIDSS Ultra--Deep Survey (UDS\\footnote{\\url{http://www.nottingham.ac.uk/astronomy/UDS/}}) is the premier near--IR dataset in terms of depth and area.  Overlapping deep optical imaging from the Subaru--XMM Deep Survey \\citep[SXDS;][]{sekiguchi04} and shallow mid--IR data from the {\\it Spitzer} Wide--Area Infrared Extragalactic \\citep[SWIRE;][]{lonsdale03} survey provide complementary data over a broad range of wavelengths.  Here we derive a multiband $K$--selected galaxy catalog from the UDS, SXDS, and SWIRE data, calculating photometric redshifts and rest--frame colors for all objects in the overlapping survey area.  Rather than rely on a specific color selection technique, we instead define samples of galaxies at various redshifts directly from the computed $\\zphot$ values, and further subdivide the samples into star--forming and quiescent galaxies based on their rest--frame colors. The color evolution with redshift, and the clustering of $1<\\zphot<2$ galaxies, are then determined.  Readers who are primarily interested in these science results can find this discussion beginning in \\S\\ref{sec_analysis}.  We present an overview of the survey data in \\S\\ref{sec_surveys}, and the process of preparing and matching the datasets to each other in \\S\\ref{sec_data}.  The extraction of the catalog is then discussed in \\S\\ref{sec_catalog}, and the derivation of photometric redshifts and rest--frame colors in \\S\\ref{sec_photz}.  \\S\\ref{sec_analysis} describes the bimodality of galaxies in rest--frame color space and the separation into quiescent and star--forming samples, and the clustering properties of these two populations are discussed in \\S\\ref{sec_clustering}. Further discussion, caveats, and a summary of the results can then be found in \\S\\ref{sec_discussion}, \\S\\ref{sec_caveats}, and \\S\\ref{sec_conclusions} respectively.  A cosmology with $H_0=70$\\,km\\,s$^{-1}$\\,Mpc$^{-1}$, $\\Omega_m=0.3$, and $\\Omega_\\Lambda=0.7$ is assumed throughout, and magnitudes are quoted in the AB system unless otherwise noted. ",
        "conclusions": "\\label{sec_conclusions} From our matched multiband UDS$+$SXDS$+$SWIRE $K$--selected catalog we have derived reliable photometric redshifts and rest--frame colors.  For a relatively bright ($K<22.4$) subsample of this catalog, we find that: \\begin{enumerate} \\item{Galaxies show strong bimodal behavior in rest--frame $U-V$ vs. $V-J$ color--color space, with one population of ``dead'' (non--star forming) galaxies, and a sequence of dusty, actively star--forming galaxies.  This behavior can be seen out to $z\\sim 2$, so the observed present--day galaxy bimodality was present at least to this redshift; however, at $z\\ga 1.5$ the bimodality is not seen in a single--color--magnitude diagram.} \\item{From $z=0$ to $2.5$, the rest--frame $V-J$ color (at red $U-V$ colors) shows a strong correlation with our estimates of specific star--formation rate, indicating that the $V-J$ color is a reliable tracer of star formation activity even at $z>2$ where the rest--frame color uncertainties are too large for the bimodality to be seen.} \\item{Quiescent and star--forming galaxies at $z>1$ contribute roughly equally to the overall galaxy population at the brightest $V$ luminosities; the less luminous population is dominated by actively star--forming objects.} \\item{The clustering strength of quiescent galaxies from $z=1-2$ appears to be independent of galaxy luminosity, and is consistently stronger than the clustering of the actively star--forming galaxies; thus, ``dead'' galaxies appear to occupy more massive halos than those which are in the process of forming stars.  This suggests a link between halo mass and the early cessation of star formation activity.} \\end{enumerate}  Forthcoming data releases of the UKIDSS UDS, including deep $H$--band data and ultimately reaching a depth of $K=24.9$ (AB), will allow these results to be tested with even fainter, higher--redshift galaxies.  More importantly, the 290--hour ultra--deep {\\it Spitzer} legacy survey of the UDS field ({\\it SpUDS}; PI: J.~Dunlop) currently being undertaken will provide at least an order--of--magnitude increase in exposure time compared to the SWIRE data used in this paper, giving much--improved constraints on photometric redshifts, AGN activity, and rest--frame colors."
    },
    "0806/0806.4230_arXiv.txt": {
        "abstract": " ",
        "introduction": " ",
        "conclusions": "Pressure-driven instabilities occur in static columns when the pressure force pushes the plasma out from the inside of the magnetic field lines curvature, as shown from direct inspection of the ``potential energy\" of the linearized displacement equation (section \\ref{EnPrin}), and from a the dispersion relation of local modes (section \\ref{disprel}). When unstable modes exist, the growth rates are of the order $C_S/R_o$ where $C_S$ is the sound speed and $R_o$ the jet radius. These are very large, comparable to the Kelvin-Helmholtz growth rate (the most studied instability in jets), especially that the ratio of the magnetic energy to the gas internal energy is expected to be of order unity (within an order of magnitude or so). Such instabilities are known to be disruptive in the fusion context when the eigenmodes exhibit a substantial displacement of the plasma outer boundary; such a situation is relevant to magnetically self-confined jets, as the magnetic field in their outer region is predominantly azimuthal, a configuration most favorable to the onset of the instability (sections \\ref{sec:gen} and \\ref{disprel}). However, the presence of a velocity gradient in the outer boundary due to the jet bulk motion is expected to have a substantial stabilizing influence, both in the linear and nonlinear regimes (section \\ref{movcol}).  In its present state, this picture possesses a number of loose ends:  \\begin{itemize} \\item The stabilizing role of an axial velocity gradient needs to be better understood. Not all modes may be stabilized in the linear regime, depending on the details of the equilibrium jet configuration, and the nonlinear mechanism identified in the literature is highly idealized and may not be generic. The one and only simulation of nonlinear stabilization published to date exhibits a very violent relaxation transient, which may still lead to jet disruption. On the other hand, this transient is also an indication that the initial configuration of the simulation is way out of equilibrium, a situation which may not occur in real jets. \\item The role of jet rotation has not yet been correctly investigated. Preliminary results seem to indicate that it is stabilizing \\cite{KLP00}; however, jet rotation may not be an important dynamical factor in the asymptotic jet propagation regime. \\item Most investigations of pressure-driven instabilities rely on a very simple prescription of the equation of state, which raises an issue of principle. Indeed, the very large growth rates usually found for the instability indicate that it develops on time-scales much shorter than the collisional time-scale, and the use of ideal MHD as well as a polytropic equation of state may be questioned in such a context, an issue briefly addressed in Appendix \\ref{app:idealmhd}. A more complex description of the plasma is required to validate the results obtained so far. \\end{itemize}   \\appendix"
    },
    "0806/0806.1623_arXiv.txt": {
        "abstract": "{We use deep multi--colour (UBVRIJK) images mostly taken in the framework of the ESO \\emph{Deep Public Survey} (DPS) to optically identify and derive photometric redshifts for a complete sample of 131 radio sources with $S>0.4$ mJy, observed at both 1.4 and 5 GHz as part of the ATESP radio survey. The availability of multi--wavelength radio and optical information is exploited to infer the physical properties of the faint radio population. In particular we find that, considering both early-type galaxies and quasars as sources with an active nucleus, AGNs largely dominate our sample sub-mJy sample (78\\%). Further radio/optical analysis of such AGN component has revealed a somewhat unexpected class of flat/inverted--spectrum sources with low radio--to--optical ratios ($R<100$), which are preferentially identified with early--type galaxies. Such sources are quite compact ($d<10-30$ kpc), suggesting core-dominated radio emission triggered by low luminosity AGNs. This intriguing class of objects deserves further analysis, and new higher resolution radio observations are currently under way. In parallel we are developing radio source models, for both the AGN and the star-forming components of the sub--mJy radio pupulation. Here we discuss the first results. }  \\FullConference{From Planets to Dark Energy: the Modern Radio Universe\\\\ October 1-5 2007\\\\ The University of Manchester, UK}  \\begin{document} ",
        "introduction": " ",
        "conclusions": "In this paper we have discussed the nature of the sub-mJy radio population, through the radio/optical analysis of the ATESP--DEEP1 sample. The high percentage of identifications ($\\sim 78\\%$) allowed us to make a direct comparison between observations and models, in order to assess the nature and the evolutionary properties of the sub-mJy population, with particular respect to the AGN component. Our radio/optical analysis has revealed a class of flat/inverted spectrum radio sources associated to optically inactive early-type galaxies, which are most plausibly triggered by low-accretion/radiative efficiency AGNs. In order to further investigating this class of objects higher resolution multi-frequency radio observations are on-going.\\\\ We find a general agreement between observations and standard pure luminosity evolution radio source modeling. It is clear however that, to provide more quantative constraints to the models, we need larger deep radio samples together with complete optical follow--up. In particular deeper radio surveys are necessary to probe a possible radio-quiet AGN component, which does not show up at the flux densities probed by our sample ($S>0.4$ mJy). \\\\ A clear understanding of the nature and evolutionary properties of the various components of the sub--mJy population is a very important issue for providing reliable modeling of the nanoJy radio sky will be eventually probed by the SKA, and will be importat matter of study for the new upcoming radio facilities (e.g. LOFAR, EVLA, eMERLIN, etc.)."
    },
    "0806/0806.4006_arXiv.txt": {
        "abstract": "The high value of brane tension has a crucial role in recovering Einstein's general relativity at low energies. In the framework of a recently developed formalism with variable brane tension one can pose the question, whether it was always that high? In analogy with fluid membranes, in this paper we allow for temperature dependent brane tension, according to the corresponding law established by E\\\"{o}tv\\\"{o}s. For cosmological branes this assumption leads to several immediate consequences: (a) The brane Universe was created at a finite temperature $T_{c}$ and scale factor \\thinspace $a_{\\min }$. (b) Both the brane tension and the 4-dimensional gravitational coupling 'constant' increase with the scale factor from zero to asymptotic values. (c) The 4-dimensional cosmological 'constant' evolves with $a$, starting with a huge negative value, passing through zero, finally reaching a small positive value. Such a scale--factor dependent cosmological constant is able to generate a surplus of attraction at small $a$ (as dark matter does) and a late-time repulsion at large $a$ (dark energy). In the particular toy model discussed here the evolution of the brane tension is compensated by energy interchange between the brane and the fifth dimension, such that the continuity equation holds for the cosmological fluid. The resulting cosmology closely mimics the standard model at late times, a decelerated phase being followed by an accelerated expansion. The energy absorption of the brane drives the 5D space-time towards maximal symmetry, becoming Anti de Sitter. ",
        "introduction": "Physics aims for a unified description of nature, tracing back all physical laws to four fundamental interactions. While three of them are quantized, and to certain extent further unified, gravity is still best described classically. In contrast with the rest of the interactions about evolving fields on a flat background, gravity is perceived as the dynamics of geometry. String theory attempts to unify all interactions on different grounds, its basic objects being open or closed strings and higher-dimensional objects, called branes. The co-dimension one brane-world theory (generalizing \\ the early Randall-Sundrum model \\cite{RS2}) carries the original geometric spirit of general relativity, incorporating arbitrary curvature and matter (for a review see \\cite{MaartensLivRev}). The extra dimension is both non-compact and curved (the remaining dimensions required by string and M-theory can still be thought as compactified). Gravity acts in 5-dimensions according to Einstein's equation with a negative cosmological constant $\\widetilde{\\kappa }^{2}\\widetilde{\\Lambda }$, while standard model fields are confined to the brane, a time-evolving 3-dimensional space-like hypersurface.  The projection of the 5-dimensional (5D) Einstein equation onto the brane gives an effective Einstein equation, which in the most generic case reads \\cite{Decomp}% \\begin{equation} G_{ab}=-\\Lambda g_{ab}+\\kappa ^{2}T_{ab}+\\widetilde{\\kappa }^{4}S_{ab}-% \\overline{\\mathcal{E}}_{ab}+\\overline{L}_{ab}^{TF}+\\overline{\\mathcal{P}}% _{ab}\\ ,  \\label{modEgen} \\end{equation}% with $T_{ab}$, the brane energy-momentum tensor; $S_{ab}$, a quadratic expression in $T_{ab}$; and $\\overline{\\mathcal{E}}_{ab}$, the electric part of the 5D Weyl tensor with respect to the brane normal \\cite{SMS}, averaged over the two sides of the brane. The source term $\\overline{L}_{ab}^{TF}$ originates in the asymmetric embedding of the brane and $\\overline{\\mathcal{P% }}_{ab}$ is the pull-back of generic non-standard model fields in 5D \\cite% {Decomp}. The 4-dimensional (4D) and 5D gravitational coupling constants $% \\kappa ^{2}$ and $\\widetilde{\\kappa }^{2}$ are related as $6\\kappa ^{2}=% \\widetilde{\\kappa }^{4}\\lambda $, with $\\lambda $ the brane tension. The 4D cosmological 'constant' $\\Lambda $, apart from contributions of the asymmetric embedding and non-standard model 5D fields, is defined as $% 2\\Lambda _{0}=\\kappa ^{2}\\lambda +\\widetilde{\\kappa }^{2}\\widetilde{\\Lambda } $.  Reference \\cite{Decomp} gives the most generalic form of the gravitational dynamics involving asymmetric embedding \\cite{asymmetry} and non-static 5d space-time due to radiation fields \\cite{ChKN}-\\cite{PalStructure}. Besides cosmological applications in branes embedded into 5d black-hole space-times \\cite{BCG}-\\cite{BraneBlackHole} or into their horizon regions \\cite% {EinBrane}-\\cite{EinBrane3} other aspects of brane-world models have been discussed, including black-hole brane-worlds \\cite{tidalRN}-\\cite% {BraneSwissCheese}, gravitational collapse on the brane \\cite{BGM}-\\cite% {BraneOppSnyder}, stellar models \\cite{GM}-\\cite{Ovalle}, galactic dynamics \\cite{HarkoRC}, the dynamics of clusters of galaxies \\cite{HarkoClusters}, light deflection \\cite{GeDa}, \\cite{GeDefl} and Solar System tests \\cite% {BohmerHarkoLobo}.  A classical fluid membrane needs tension to exist. Similarly its higher dimensional counterpart, the 3-brane, as it evolves, remains a hypersurface due to the brane tension. The strongest bound on the minimal value of $% \\lambda $ was derived by combining the results of table-top experiments on possible deviations from Newton's law, which probe gravity at sub-millimeter scales \\cite{tabletop} with the known value of the 4D Planck constant. In the 2-brane model \\cite{RS1} this gives \\cite{Irradiated} $\\lambda >138.59\\,\\,$TeV$^{4}$. A much milder limit $\\lambda \\gtrsim 1$ MeV$^{4}$ arises from the constraint that the dominance of $S_{ab}$ ends before the Big Bang Nucleosynthesis (BBN) \\cite{nucleosynthesis}. From astrophysical considerations on brane neutron stars an intermediate value $\\lambda >5\\,\\times 10^{8}$ MeV$^{4}$ was derived \\cite{GM}. (All these limiting values are for $c=1=\\hbar $. In units $c=1=G$ the corresponding minimal values of the brane tension are $\\lambda _{tabletop}=4.\\,\\allowbreak 2\\times 10^{-119}$ eV$^{-2}$, $\\lambda _{BBN}=3\\times 10^{-145}$ eV$^{-2}$ and $% \\lambda _{astro}=1.5\\,\\times 10^{-136}$ eV$^{-2}$, respectively \\cite% {BraneOppSnyder}.) For typical stellar densities the condition $\\rho _{star}/\\lambda \\ll 1$ is obeyed with any of these bounds.  In a cosmological context the brane represents our observable Universe. Cosmic expansion is realized through the movement of the brane in the warped extra dimension. During cosmological evolution the temperature of the brane changes drastically. Cosmological branes cool down from a very hot early Universe (whose thermal radiation is able to create a black hole in the fifth dimension \\cite{ChKN}), to the present days low temperature of the Cosmic Microwave Background (CMB).  In Sec. II we explore the possibility of a variable brane tension, discussed in detail in Ref. \\cite{VarBraneTensionPRD}, by introducing a toy model, in which the brane tension literally follows the temperature dependence of the fluid tensions established for membranes. Such a model becomes particularly simple when we assume the continuity equation. We discuss the emerging cosmological model in Sec. III. A numerical solution of this toy-model, which is compatible with cosmological observations is presented in Sec. IV. Finally, Sec. V contains the concluding remarks.  Throughout the paper we follow the notations of Ref. \\cite% {VarBraneTensionPRD}. ",
        "conclusions": "The variation of the brane tension introduces an additional degree of freedom in brane-world models. The particular model discussed here, assuming the E\\\"{o}tv\\\"{o}s law for the temperature dependence of the brane tension, balanced by the energy interchange between the brane and VAdS5 (such that the continuity equation holds), resulted in a monotonic increase with scale factor of the brane tension, gravitational coupling constant and 4D cosmological constant. In the early universe both the brane tension and the 4D gravitational coupling constant are small, enhancing the dominance of brane-world effects. The temperature-dependent 4D cosmological constant, being negative for small values of the scale factor, contributes to mutual attraction, while positive for large $a$, generates dark energy type repulsion.  We established the range of the model parameters allowed by the confrontation with observations, given by Eqs. (\\ref{c1})-(\\ref{c4}). A particular configuration obeying these conditions, with $% R=10^{25},~x_{0}=10^{11},~y_{0}=10^{34}$ was represented on Figs \\ref{aevol}% , \\ref{mevol}. For the allowed range the evolution of the fundamental constants basically occur in the very early universe preceding BBN, after which they asymptote to constant values. Still, feeded by absorbed energy from the VAdS5 regions, they slightly evolve. This process eventually will consume $m\\left( \\tau \\right) $, leaving maximally symmetric AdS5 space-time patches on the two sides of the brane. which further expands in a de Sitter phase."
    },
    "0806/0806.1415_arXiv.txt": {
        "abstract": "At frequencies close to 1 GHz the sky diffuse radiation is a superposition of radiation of Galactic origin, the 3 K Relic or Cosmic Microwave Background Radiation, and the signal produced by unresolved extragalactic sources. Because of their different origin and space distribution the relative importance of the three components varies with frequency and depends on the direction of observation. With the aim of disentangling the components we built TRIS, a system of three radiometers, and studied the temperature of the sky at $\\nu =0.6$, $\\nu = 0.82$ and $\\nu = 2.5$ GHz using geometrically scaled antennas with identical beams (HPBW = $18^{\\circ} \\times 23^{\\circ}$). Observations included drift scans along a circle at constant declination $\\delta=+42^{\\circ}$ which provided the dependence of the sky signal on the Right Ascension, and absolute measurement of the sky temperature at selected points along the same scan circle. TRIS was installed at Campo Imperatore (lat. = $42^{\\circ}~26'$ N, long.= $13^{\\circ}~33'$, elevation = 2000 m a.s.l.) in Central Italy, close to the Gran Sasso Laboratory. ",
        "introduction": "\\label{intro} The diffuse radiation from the sky is a superposition of components. At frequencies between few tens of MHz and few tens of GHz, in terms of brightness temperature we can write \\begin{equation} T_{sky}(\\alpha,\\delta,\\nu)=T_{gal}(\\alpha,\\delta, \\nu)+T_{CMB}(\\nu)+T_{uers}(\\nu) \\label{eq:Tskygen} \\end{equation} \\par\\noindent where $\\alpha$ and $\\delta$ are the Right Ascension and declination of the point at which the telescope axis is aimed. $T_{gal}$ is the Galactic contribution: partially polarized, anisotropically distributed, it tracks the Galactic structure. Its frequency spectrum is a power law \\begin{equation} T_{gal}(\\nu, \\alpha, \\delta)  = K(\\alpha, \\delta) ~\\nu ^{-\\beta({\\alpha,\\delta})} \\label{eq:Tgal} \\end{equation} \\noindent with spectral index $\\beta$ ranging from 2.1 to 3.4 depending on the relative weight of thermal and synchrotron emission, on the energy spectrum of the Cosmic Ray Electrons and on the Galactic Magnetic Field. Compared to the other components of the diffuse radiation it is a foreground and dominates the sky at frequencies below $\\sim 1$ GHz.  \\noindent $T_{CMB}$ is the Cosmic Microwave Background. Relic of the Big Bang, it is substantially unpolarized and isotropically distributed. Its flat frequency spectrum is consistent with the emission of a blackbody with a temperature of $(2.725 \\pm 0.001)$ K \\cite{Fixsen2002}. Compared to the other components it is a true background. At $\\nu \\geq 1$ GHz $T_{CMB}$ definitely overcomes $T_{gal}$ and dominates the sky up to $\\nu \\sim 100$ GHz, above which thermal emissions from an irregular distribution of dust with physical temperature $\\geq 20$ K, gradually overwhelms the other components.  \\noindent $T_{uers}$ is a blend of unresolved extragalactic radio sources isotropically distributed. Its frequency spectrum is a power law \\begin{equation} T_{uers} = K_{uers} \\nu^{-\\gamma_{uers}}  \\label{eq:Tex} \\end{equation} with $\\gamma_{uers} \\sim 2.70$ \\cite{Gervasi2008b}. \\newline \\noindent Each of the above components carries important astrophysical and/or cosmological information, particularly at decimetric wavelengths. At these frequencies for instance we can expect a dip in the flat spectrum of $T_{CMB}$, whose detection can provide direct information on $\\Omega_b$ \\cite{Burigana1991}, the Universe baryon density. In this same region the slope of the power law frequency spectrum of $T_{gal}$ changes, revealing perhaps a knee in the energy spectrum of the cosmic ray electrons responsible for the galactic synchrotron radiation. To disentangle these components, coordinated, multifrequency observations of extended areas of sky are necessary. They will form a data base from which the desired information can be extracted modelling the components and optimizing the model parameters. At frequencies close to 1 GHz it is possible to find in literature maps of $T_{sky}$ which cover all the sky or large parts of it (see Table \\ref{tab1}). Their accuracy is however insufficient to answer many of the questions put forth by present day astrophysical and cosmological models. These maps have been in fact obtained combining data collected at different sites and/or at different times and frequently show artificial structures, e.g. {\\it stripes}. Once {\\it destriped} (see for instance \\cite{Platania2003}), these maps can be used to extract the spectral index of the galactic signal $T_{gal}$. The accuracy of the zero level of the absolute scale of temperature of the same maps is however still insufficient to disentangle the radiation components with the accuracy today required by cosmology. To overcome this difficulty we made absolute and differential measurements of the diffuse radiation at three frequencies, using a set of three radiometers (therefore the project is said TRIS), at 0.6, 0.82 and 2.5 GHz. In this paper we present results of these observations. In \\cite{Gervasi2008a} and \\cite{Tartari2008} (from now on cited as Paper II and III), the astrophysical and cosmological implications of TRIS data are discussed. ",
        "conclusions": "\\label{Discussion} The absolute values of $T_{sky}(\\alpha,\\delta)$ measured by TRIS at 0.6005 GHz, 0.8178 GHz and 2.4278 GHz, at selected points of constant declination $+42^{\\circ}$ are listed in Tables \\ref{tab7}, \\ref{tab8} and \\ref{tab9}. The same data are plotted in Figure \\ref{F5:TRISfinal} together with the drift scans (0.6 GHz and 0.82 GHz are reported in Tables \\ref{tab10} and \\ref{tab11}). Having no reliable drift scan at 2.5 GHz, we built a profile only of the galactic component starting from the Stockert \\cite[]{Reich1986} survey at 1.420 GHz. This map has been convoluted to the beam of the TRIS antennas in order to get the synthetic drift scan at $\\delta =+42^\\circ$. Then we compared this scan with the pure galactic signal we evaluated at 0.6 and 0.82 GHz (see Papers II and III). Using these measurements we evaluated the galactic emission first at 1.420 GHz and then extrapolated it at 2.5 GHz, using the local spectral index evaluated between 0.6 and 0.82 GHz. We used the local spectral index because it is a well defined quantity throughout all the sky provided that absolute galactic emissions measurements are available, which is our case. We used the map at 1.42 GHz because it is the closest one, public domain, at the frequency 2.5 GHz. In Figure \\ref{F5:TRISfinal} at 2.5 GHz, for uniformity with the measured sky profiles plotted at 0.6 and 0.82 GHz, we added over the TRIS absolute sky temperatures a profile (dashed line) obtained summing to the galactic component calculated as described above, the Unresolved Extragalactic Radio Sources contribution \\cite{Gervasi2008b} and Cosmic Microwave Background \\cite{Gervasi2008a}.  The systematic uncertainties are computed propagating through Equation (\\ref{eq:Tant}) and (\\ref{eq:Tsky}) the uncertainties on the measured values of attenuation and temperature of the lossy components between the switch and the antenna mouth. The losses are very small. In some cases they are so small that measurements give only upper limits (see Table \\ref{tab4}) and the range of fluctuation of the true losses is unknown but reasonably smaller and well inside the upper limit, set by the accuracy of the instrumentation used to carry on the measurements. The final uncertainties on the sky temperature obtained propagating these errors are therefore considered as systematic. At 0.6 GHz the great majority of the components is cold ($\\sim 4$ K) and the final systematic uncertainty is small (66 mK). At 2.5 GHz the systematic uncertainty is worse because of the quality of the e.m. matching of the SPTT switch and the presence of an additional section of warm waveguide between the dewar flange and the antenna throat. This correction contributes to the systematics for 284 mK. At 0.82 GHz the systematic uncertainty is very large, 660 mK, essentially because all the components above the switch were warm and at high temperature (see Equation \\ref{eq:DeltaTnoise}). Such a large value, however, does not take into account constraints set by unphysical situations (e.g. negative values of the temperature of the diffuse radiation components and values of the galactic spectral index not supported by models of galactic emission). When astrophysical assumptions are made on these values (for a complete discussion see Paper III) the systematic uncertainty at 0.82 GHz is reduced to $(^{+430}_{-300})$ mK.  The above results hold if the sky radiation is unpolarized. We know however that the galactic component of $T_{sky}$ is partially polarized and that polarization is linear. Since antennas are polarized, our measured values of the sky temperature are better written as \\begin{equation} T_{sky}(\\alpha,\\delta,\\nu) = T_{CMB}(\\nu) + T_{uers}(\\nu) + T_{gal}^{unpol}(\\alpha,\\delta,\\nu) + T_{gal}^{pol}(\\alpha,\\delta,\\nu)~\\cos^2 [\\theta(\\alpha,\\delta,\\nu) - \\theta_0] \\label{eq:Tpol} \\end{equation}  \\par\\noindent where $T_{gal}^{unpol}$ and $T_{gal}^{pol}$ are respectively the unpolarized and polarized components of the galactic signal, and ($\\theta - \\theta_0$) is the angle between the polarization vector of the radiation and the polarization vector of the antenna. \\par\\noindent The true sky temperature is therefore \\begin{equation} T_{sky}^{true}(\\alpha,\\delta,\\nu) = T_{sky}^{unpol}(\\alpha,\\delta,\\nu) + T_{sky}^{pol}(\\alpha,\\delta,\\nu) = T_{sky}(\\alpha,\\delta,\\nu) + \\Delta T^{pol}(\\alpha,\\delta,\\nu) \\label{eq:Ttrue} \\end{equation} \\par\\noindent where \\begin{equation} \\Delta T^{pol}(\\alpha,\\delta,\\nu) = T_{gal}^{pol}(\\alpha,\\delta,\\nu)~\\sin^2 [\\theta(\\alpha,\\delta,\\nu) - \\theta_0] \\label{eq:DeltaT} \\end{equation} \\par\\noindent is the signal lost. Were circular polarization present, an additional term should be added to Equation (\\ref{eq:Ttrue}).  The information today available on the polarization degree and direction of the diffuse radiation is limited. We do not expect circular polarization. Linear polarization on the contrary must be expected but the information available is poor. Below 1 GHz there are only data obtained thirty years ago by \\cite{Brouw1976}. The assessment of the overall sky temperature has been obtained correcting our absolute measurements for the fraction of the linearly polarized signal. In order to do this we used the Brouw and Spoelstra maps of the Stokes Parameters Q$\\&$U in the resampled version prepared by E.Carretti in the framework of the SPOrt program \\cite{Cortiglioni 2004}. Carretti maps having 7$^\\circ$ resolution, were further convolved with the beam of our experiment. Then, considering the relative orientation of the polarization plane (E-plane) of our antennas and the polarization of the celestial signal, we projected the polarized vector into Copolar (collected by our antennas) and Crosspolar components (rejected by our antennas). We conclude that, at TRIS angular resolution, the maximum contribution of the polarized signal to the overall one, around 2$\\%$ at 0.6 GHz and $3\\%$ at 0.82 GHz, comes mostly from regions far from the galactic disk. This result is in agreement with Platania et al. (1998) who gives a maximum contribution of $\\sim 5\\%$ at 0.408 GHz with a $18^\\circ$ beam in a sky region very close to the one scanned by TRIS. A detailed discussion can be found in Paper III. Between 1 and 3 GHz more recent and accurate data are available (see for instance \\cite{Duncan1991}, \\cite{Wolleban2006} and references therein), but with few exception they do not fill all the areas covered by TRIS scans. Full sky models of the galactic synchrotron intensity and linear polarization prepared for feasibility studies of forthcoming space experiments (e.g. \\cite{Giardino2002}, \\cite{Bernardi2003}) are available. However we preferred to evaluate the polarized component starting from real data acquired at frequencies close to the TRIS ones. Considering that at 2.5 GHz with the TRIS angular resolution $(T_{gal}/T_{sky})\\simeq(T_{gal}/T_{CMB})\\leq 0.1$ and $T_{gal}^{pol} \\leq T_{gal}$, we can assume also at 2.5 GHz $(\\Delta T^{pol}/T_{sky}^{true})<0.02$.  In conclusion, neglecting the polarization effects is equivalent to introduce an additional systematic uncertainty on the true values of the sky temperature of less than $3\\%$ (see Table \\ref{tab12}).  Table \\ref{tab13} gives a list of measurements of $T_{Sky}$ present in literature at frequencies comparable to the frequencies used by TRIS, the value of the derived $T_{CMB}$ and the associated uncertainties for $T_{Sky}$ and $T_{CMB}$. All the values of $T_{Sky}$ have been obtained as intermediate steps of experiments aimed at getting the CMB temperature, so papers usually focus the reader attention on the accuracy $\\sigma_{CMB}$ of $T_{CMB}$. The accuracy $\\sigma_{sky}$ of $T_{sky}$ should be better ($\\sigma_{sky}~<~\\sigma_{CMB}$) but extracting it from papers it is not always possible.  When no detailed information is available in literature we assume $\\sigma_{sky} ~\\leq~ \\sigma_{CMB}$. It appears that at 0.6 and 0.82 GHz the TRIS absolute values of the sky temperature are the best today available. At 2.5 GHz TRIS results are less accurate than data obtained in the past by us with the White Mt. and the South Pole collaborations. Nevertheless we report also the TRIS results at 2.5 GHz because they came exactly from the same direction and have been obtained with the same angular resolution of our observation at lower frequencies.  \\noindent TRIS data can be used to  \\noindent i) disentangle the components of the diffuse radiation,  \\noindent ii) extract the cosmological and astrophysical information carried by these components,  \\noindent iii) improve zero level and scale of temperature of the full sky maps of the diffuse radiation at decimetric wavelength in literature (see Tables \\ref{tab1} and \\ref{tab2}).  These types of analysis are intricate and require different approaches, depending on the final aim: astrophysical or cosmological conclusion. They are done in separate Papers II \\cite{Gervasi2008a}, III \\cite{Tartari2008} and \\cite{Gervasi2008b} which accompany present Paper I. Here (Table \\ref{tab14}) we present a summary of the results discussed in the cited papers. We have reduced the error bars on $T_{CMB}$ of a factor 9 at $\\nu=0.6$ GHz and of a factor 7 at $\\nu=0.82$ GHz. These results have been possible because of the improvements in the absolute calibration system and in the foregrounds separation technique. At 2.5 GHz TRIS did not improve the previous measurements, but is in agreement with them."
    },
    "0806/0806.4376_arXiv.txt": {
        "abstract": "This paper summarizes our theoretical understanding of supernova events in a ``back of the envelope'' way. It is intended to aid in the recognition and understanding of those events which are not ``standard'', and which may provide the most insight. ",
        "introduction": "This material is abstracted from previous work by the author; see \\citet{arnett1996,arnett1977} for a more detailed discussion and references to earlier work. It is hoped that by using the template analysis given below, it will be easier to recognize non-standard events. ",
        "conclusions": ""
    },
    "0806/0806.4189_arXiv.txt": {
        "abstract": "We investigate the origins of galaxy morphology (defined by bulge-to-total K-band luminosity) in the $\\Lambda$CDM cosmology using two galaxy formation models based on the Millennium simulation, one by Bower et al. (the Durham model) and the other by De Lucia \\& Blaizot (the MPA model).  Both models have had considerable success in reproducing a number of observed properties of the local and high redshift universe, including star formation rates, the stellar mass function and the luminosity function out to $z\\sim5$.  There are many similarities, but also fundamental disagreements in the predictions of the two models for galaxy morphology.  For example, taking into account uncertainties  in the available observational data, both produce a realistic morphological mix today, but its evolution is very different.  A main cause of this and other differences is the treatment of disk instabilities which play a more prominent role in the Durham model.  Our analysis confirms previous theoretical predictions that elliptical galaxies form most of their stars before the bulk of the galaxy is assembled.  Spirals tend to have later `assembly' times than ellipticals as a consequence of in-situ star formation.  With the exception of the brightest ellipticals (stellar mass $M_{*}\\gsim 2.5\\times 10^{11}h^{-1}M_{\\odot}$), we find that major mergers are {\\it not} the primary mechanism by which most spheroids (ellipticals and spiral bulges) assemble their mass.  In fact, the majority of ellipticals (and the overwhelming majority of spirals) never experience a major merger (above the resolution limit of our simulation.) Most ellipticals and spiral bulges acquire their stellar mass through minor mergers or disk instabilities.  These conclusions are common to both the MPA and Durham models.  The rotation properties of spheroids may help to constrain the importance of disk instabilities in these models. ",
        "introduction": "\\label{sec:intro} The morphological type of a galaxy describes more than merely its physical appearance; it reveals the key processes that have shaped it and continue to affect its evolution.  Hubble's well established classification scheme for galaxies~\\citep{Hubble1926,Hubble1936} has survived with only modest modification~\\citep{DV1959,VDB1960,Sandage1961}, precisely because the disks and spheroids on which it relies are fundamental products of the formation process.  Understanding the origins of such structures is therefore of central importance if we are to build up a coherent picture of galaxy formation and evolution. Modern observational facilities like the Hubble Space Telescope have made it possible to probe progressively higher redshifts and, together with large surveys like the Sloan Digital Sky Survey (SDSS), are beginning to reveal how galaxy morphology changes with time.  A significant trend in the low redshift universe is the `density-morphology' relation~\\citep{Dressler1980,Postman1984,Dressler1997}, which shows that denser regions contain proportionally more early-type galaxies.  This relation holds at least back to $z\\sim1$~\\citep{Postman2005}.  At higher redshifts, between $z\\sim0.5$ and 2, there appears to be a rapid breakdown of the Hubble classification system~\\citep{VanDenBergh2002,Papovich2005} with the proportion of galaxies classed as irregulars or peculiars rising steadily, particularly at faint magnitudes~\\citep[e.g.,][]{Brinchmann2000}.  Nonetheless, a substantial fraction of giant spirals are observed at moderate redshifts~\\citep[e.g.,][]{Ravindranath2004}.  Some authors have claimed that this population shows very little evolution up to the present, implying that at least some Hubble sequence galaxies were in place at $z\\sim1$~\\citep{Marleau1998,Ravindranath2004}, whilst others counter that many of the higher redshift sample would have become bulge dominated in the intervening time, so they are not directly comparable to their $z=0$ counterparts~\\citep{Bell2008}.  Before $z=2$, the majority of observed galaxies are either highly irregular or centrally compact~\\citep{Giavalisco1996,Dickinson2000,Daddi2004}.  The difficulty in linking populations seen at different epochs has ensured that a clear evolutionary picture for morphology remains elusive.  At present, direct simulation of the processes that determine morphology remains too computationally taxing to perform for large samples of galaxies (although significant progress is being made in this respect; see for example \\citealp{Croft2008}).  Semi-analytic techniques constitute a convenient tool to capture the most important characteristics of these simulations, whilst also providing the statistics required to perform detailed analyses.  The galaxy formation models considered in this paper are described in $\\S$~\\ref{sec:MSMODELS}.  Our current understanding of galaxy formation tells us a little about how morphology develops.  It is well known, for instance, that angular momentum acquired by a proto-galactic gas cloud through tidal torques operating during its expansion and collapse will naturally cause it to settle into a disk~\\citep{Hoyle1949,Peebles1969,White1984}.  Beyond this though, the combination of mechanisms required to explain morphological change are poorly understood.  \\citet{Toomre1977} was the first to put forward the idea that mergers could be responsible for the formation of elliptical galaxies.  Indeed, the importance of mergers in the hierarchical cold dark matter (CDM) model makes them a natural candidate for such a role~\\citep{Frenk1985}.  Early CDM galaxy formation models stressed the possibility that, through a combination of mergers and gas accretion, a single galaxy could, at various stages in its lifetime, be identified with several morphological types on the Hubble sequence~\\citep{White1991}, a view supported by gas-dynamical simulations~\\citet{Steinmetz2002}.  Numerous simulations have confirmed at least that major mergers, or very strong gravitational encounters, can lead to the complete destruction of a disk and the formation of a spheroid with properties similar to elliptical galaxies~\\citep[e.g.,][]{White1978,Negroponte83,Barnes1992,Barnes1996,Bekki1997,Naab1999,Springel2005b,Naab2003,Cox2006}. Even in virialised clusters where the relative velocities of galaxies make strong gravitational interactions of any kind infrequent, mergers will still have affected groups collapsing to form the cluster, including those it accretes at later times.  Other phenomena act to change morphology through their effects on gas content and star formation.  For instance,~\\emph{strangulation} (also called~\\emph{suffocation}) occurs when a galaxy with a gaseous atmosphere falls into a larger structure, causing gas cooling to be heavily supressed~\\citep{Larson1980,White1991}.  In the formation models of \\citet{Bower2006} and \\citet{DeLuciaBlaizot2007} considered here, ~\\emph{strangulation} is modelled by stripping all infalling (satellite) galaxies of their hot gas supplies, preventing further cooling.  In combination with efficient feedback from supernovae, this can lead to the reddening of a satellite's stellar population on a relatively short timescale.  Two further effects that are not included in these two models are \\emph{ram pressure stripping}~\\citep{GunnGott1972,Abadi1999} and \\emph{harassment}~\\citep{Moore1996}, both of which are most important for cluster galaxies.  The former involves the stripping of gas as galaxies move through the intracluster medium (ICM), while the latter results from a series of intense, high speed gravitational encounters that can cause bursts of star formation and lead to disk disruption.  Recent additions to the Durham galaxy formation model analysed in this paper \\citep{Bower2006} include a detailed prescription for ram pressure stripping \\citep{Font2008}. In this paper, we use an online database to examine two different galaxy formation models implemented on the $\\Lambda$CDM Millennium Simulation, the first by \\citet{DeLuciaBlaizot2007} (hereafter the MPA model) and the second by \\citet{Bower2006} (hereafter the Durham model).  We look at the morphological content of each model as a function of redshift and track the evolution of the luminosity function of each type.  We go on to determine how a galaxy's formation epoch, and that of its host dark matter halo, affect its final morphology and investigate the role played by major mergers, minor mergers and disk instabilities.  Our main goal is to investigate galaxies of all morphological types and to compare and contrast the predictions of the two models, but we also carry out limited comparisons to observations.  The rest of the paper is laid out as follows.  In $\\S$\\ref{sec:MSMODELS}, we describe the N-body simulation and semi-analytic models that form the basis of this analysis. $\\S$\\ref{sec:DATA} discusses the methods we use to select and retrieve the data and compares some basic properties of morphology in the models with SDSS data.  Finally, our results and conclusions are presented in $\\S$\\ref{sec:RESULTS} and $\\S$\\ref{sec:CONCLUSIONS}. ",
        "conclusions": "\\label{sec:CONCLUSIONS}  We have investigated the origin and evolution of galaxy morphology in the $\\Lambda$CDM cosmology using two different and publicly available semi-analytic galaxy formation models, the Durham and MPA models \\citep{Bower2006,DeLuciaBlaizot2007}, both based on the Millennium dark matter simulation.  Morphology was defined for each galaxy by means of the bulge-to-total ratio in K-band luminosity; we distinguish between the behaviour of bright ($M_{K}-\\rm{5logh}<-22.17$) and faint ($M_{K}-\\rm{5logh}>-22.17$) galaxies.  There are many similarities, but also fundamental disagreements in the predictions of the two models.  One of the largest differences involves the redshift evolution of the morphological fractions: spirals are the most common systems at all times in the MPA model, while ellipticals dominate in the Durham model at redshifts $z>4$. The main cause of this and other differences highlighted below is the different treatment of disk instabilities in the two models.  In the Durham model, instabilities are not only more common at early times, but also have more catastrophic consquences: the sudden conversion of the entire disk, first into a bar and then into a spheroid.  Nevertheless, the two models predict a fairly comparable morphological mix at the present day.  In comparison with data from \\citet{Benson2007}, the substantial systematic uncertainties involved in the process of fitting bulge and disk light profiles mean that only weak constraints can be placed on the fraction of each morphological type as a function of absolute r-band magnitude.  Neither model is obviously inconsistent with the data, but clearly smaller systematic uncertainties would be desirable.  The bright galaxy populations consist of $51\\%$, $15\\%$ and $34\\%$ of spirals, S0s and ellipticals respectively in the Durham model, and $67\\%$, $8\\%$ and $25\\%$ respectively in the MPA model.  For typical Milky-Way type haloes, we find that the formation time of the halo has, at most, only a small effect on the morphology of the central galaxy.  In both models, haloes hosting ellipticals form slightly earlier than haloes hosting spirals or S0s, although this effect is only marginal in the Durham model.  Our analysis confirms previous studies \\citep{Kauffmann1996,DeLucia2006a} that show that ellipticals galaxies form most of their stars before the bulk of the galaxy is assembled.  In fact, the ellipticals form and assemble their stars earlier than galaxies of other present-day morphological types, although the late `assembly' times of spirals are a consequence of in-situ star formation, rather than mergers.  In general, the MPA model predicts slightly earlier star formation for galaxies of all types.  In the Durham model, there is a significant population of S0s and ellipticals that formed at relatively low redshift from the recent collapse of unstable disks, yielding spheroids that never regrow a sizeable stellar disk.  The differences in the stellar formation histories predicted by the two models are, in principle, accessible to observational tests.  Perhaps the most surprising conclusion of this study is that, with the exception of the brightest ellipticals (stellar mass $M_*\\gsim 2.5\\times 10^{11}h^{-1}M_\\odot$), major mergers are {\\it not} the primary mechanism by which most spheroids (ellipticals and spiral bulges) assemble their mass.  Elliptical galaxies are more likely than spirals or S0s to have undergone a recent major merger, but nonetheless more than half of them do not experience one throughout their whole formation history (involving galaxies above the resolution limits of our calculations).  Major mergers are even rarer for spirals, affecting $<4\\%$ of galaxies; their bulges are almost universally formed by minor mergers or disk instabilities.  These conclusions would appear to be quite robust for galaxy formation in a $\\Lambda$CDM universe since they are common to both the MPA and Durham models, in spite of the substantial differences in the detailed modelling of key physical processes.  An important difference between the two models concerns the channel through which less massive ellipticals and bulges acquire their mass.  In the Durham model, disk instabilities are responsible for the bulk of the stellar mass in spheroids at $z=0$, whereas in the MPA model, minor mergers contribute the most.  In the Durham model at least, these results appear to be relatively insensitive to the mass ratio at which a merger is defined as `major'.  Varying the value between 0.25 and 0.35 typically alters the fractions in fig. \\ref{fig:all_frac} by less than a few percent for any given stellar mass; in order for major mergers to become the dominant channel for bulge growth in galaxies larger than $\\sim5\\times10^{10}h^{-1}M_\\odot$, the ratio would need to be lowered to $<0.01$.  Adding a more realistic prescription for gas stripping of satellites \\citep{Font2008} increases the importance of major mergers, but they still contribute $<20\\%$ of spheroid mass in galaxies smaller than $\\sim5\\times10^{11}h^{-1}M_\\odot$.  Mergers as a whole contribute less mass to spheroids in the Durham model because strong supernova feedback suppresses star formation especially strongly in low mass satellite galaxies.  The dominant role of disk instabilities in the production of spheroids, particularly in the Durham model, offers a possible observational test, since spheroids formed this way will tend to rotate faster than spheroids formed through mergers.  However, this test would be complicated by the fact that gas rich major mergers can also produce high $V/\\sigma$ remnants.  What is clear is that a more precise comparison between theory and observations will require more detailed modelling of disk instabilities than we have attempted here.  More generally, semi-analytic techniques offer, at present, only a somewhat crude way to model the processes that establish galaxy morphology.  It is therefore remarkable that such a simplified approach can lead to some important conclusions that appear to be robust to the modelling details.  This approach also serves to identify specific processes that deserve further study, for example, through direct numerical simulation."
    },
    "0806/0806.2146_arXiv.txt": {
        "abstract": "The mass varying neutrino scenario is a model that successfully explains the origin of dark energy while at the same time solves the coincidence problem. The model is, however, heavily constrained by its stability towards the formation of neutrino bound states when the neutrinos become nonrelativistic. We discuss these constraints and find that natural, adiabatic, stable models with the right amount of dark energy today do not exist. Secondly, we explain why using the lightest neutrino, which is still relativistic, as an explanation for dark energy does not work because of a feedback mechanism from the heavier neutrinos. ",
        "introduction": "% Precision observations of the cosmic microwave background \\cite{Bennett:2003bz,Spergel:2003cb,Spergel:2006hy}, the large scale structure of galaxies \\cite{Tegmark:2006az}, and distant type Ia supernovae \\cite{Riess:1998cb,Perlmutter:1998np,Astier:2005qq,Wood-Vasey:2007jb} have led to a new standard model of cosmology in which the energy density is dominated by dark energy with negative pressure, leading to an accelerated expansion of the universe.  A very interesting proposal is the so-called mass varying neutrino (MaVaN) model~\\cite{Hung:2000yg,Gu:2003er,Fardon:2003eh} in which a light scalar field couples to neutrinos, see also \\cite{Peccei:2004sz,Schrempp:2006mk,Ringwald:2006ks,Kaplan:2004dq,Barger:2005mn,Cirelli:2005sg,Barger:2005mh,Gu:2005pq,Li:2004tq,Afshordi:2005ym,Brookfield:2005td,Brookfield:2005bz,Takahashi:2006jt,Spitzer:2006hm,Fardon:2005wc,GonzalezGarcia:2007ib,Bjaelde:2007ki,Ichiki:2007ng,Gu:2007mi,Kaloper:2007gq,Amendola:2007yx,Wetterich:2007kr,Bean:2007nx,Bean:2007ny,Takahashi:2007ru,Gu:2007gy,Bauer:2007gf,Brouzakis:2007aq,Bhatt:2007ah,Anchordoqui:2007iw,Pettorino:2008ez,Mota:2008nj,Bernardini:2008pn,Schrempp:2008zz,Avelino:2008zz,Ichiki:2008rh,Ichiki:2008st,Valiviita:2008iv,Bhatt:2008hr,Li:2008fa} for background. Because of the coupling, the mass of the scalar field does not have to be as small as the Hubble scale but can be larger, while the model still accomplishes late-time acceleration.  In this paper we discuss the different criteria that need to be fulfilled in order to have a stable, adiabatic mass varying neutrino model with the correct cosmology today - i.e. $\\Omega_{\\mathrm{DE}}=0.7$, $\\Omega_{\\mathrm{M}}=0.3$ and $w\\sim-1$. Our aim is to show that it is only possible as long as the scalar field potential resembles a cosmological constant - and hence the model loses its prediction of the current neutrino mass and fails to solve the coincidence problem.  In addition, we discuss the suggestion that the lightest neutrino, which can be relativistic today, may be responsible for dark energy. We find that there is evidence that the relativistic neutrino will feel an instability towards the formation of neutrino nuggets.  In the next section we briefly review the formalism needed to study mass varying neutrinos and in Sec.~\\ref{sec:3} we discuss the different criteria in the MaVaN framework. In Sec.~\\ref{sec:4} we discuss MaVaNs with a relativistic neutrino and in Sec.~\\ref{sec:5} we conclude. ",
        "conclusions": "\\label{sec:5} % Single scalar field models can be used to explain late-time acceleration in the MaVaN scenario. However, in general using these potentials leads to instabilities towards neutrino bound states unless certain criteria are relaxed. In order to obtain a stable or metastable model we direct the reader to Ref.~\\cite{Bjaelde:2007ki} in which one such model is presented. The crucial thing is that the scalar field potential resembles that of a cosmological constant. In this case we simply move the fine-tuning problem to that of explaining the small fraction in the power-law - which unfortunately is also rather unnatural.  Accordingly it has been suggested to include an extra scalar field in the treatment. This has some very nice features and is easily capable of obtaining late-time acceleration as well as $\\Omega_{\\mathrm{DE}}=0.7$ today. However, one drawback is the need for the lightest neutrino to be relativistic today. As was explained above the feedback mechanism will eventually cause the perturbations of the relativistic neutrino to start growing. As a result, the MaVaN fluid will cease to act as dark energy."
    },
    "0806/0806.0375_arXiv.txt": {
        "abstract": "{}% {To study the existence of large-scale convective dynamos under the influence of shear and rotation.}% {Three-dimensional numerical simulations of penetrative compressible convection with uniform horizontal shear are used to study dynamo action and the generation of large-scale magnetic fields. We consider cases where the magnetic Reynolds number is either marginal or moderately supercritical with respect to small-scale dynamo action in the absence of shear and rotation. Our magnetic Reynolds number is based on the wavenumber of the depth of the convectively unstable layer. The effects of magnetic helicity fluxes are studied by comparing results for the magnetic field with open and closed boundaries. }% {Without shear no large-scale dynamos are found even if the ingredients necessary for the $\\alpha$-effect (rotation and stratification) are present in the system. When uniform horizontal shear is added, a large-scale magnetic field develops, provided the boundaries are open. In this case the mean magnetic field contains a significant fraction of the total field. For those runs where the magnetic Reynolds number is between 60 and 250, an additional small-scale dynamo is expected to be excited, but the field distribution is found to be similar to cases with smaller magnetic Reynolds number where the small-scale dynamo is not excited. In the case of closed (perfectly conducting) boundaries, magnetic helicity fluxes are suppressed and no large-scale fields are found. Similarly, poor large-scale field development is seen when vertical shear is used in combination with periodic boundary conditions in the horizontal directions. If, however, open (normal-field) boundary conditions are used in the $x$-direction, a large-scale field develops. These results support the notion that shear not only helps to generate the field, but it also plays a crucial role in driving magnetic helicity fluxes out of the system along the isocontours of shear, thereby allowing efficient dynamo action. } {} ",
        "introduction": "It is evident that the Sun possesses a large-scale magnetic field. The most likely origin of this field is a hydromagnetic dynamo working within or somewhat below the turbulent convection zone (e.g.\\ Ossendrijver 2003). The solar dynamo is generally thought to rely on large-scale shear to produce toroidal magnetic field from poloidal field and a process called the $\\alpha$-effect which produces poloidal field from toroidal field. In the turbulent mean-field dynamo picture an $\\alpha$-effect can arise due to helical motions in a stratified fluid in the presence of either rotation or shear, or both (Rogachevskii \\& Kleeorin \\cite{RogaKle2003}; R\\\"udiger \\& Kitchatinov \\cite{RueKit2006}; R\\\"adler \\& Stepanov \\cite{RaedStep2006}). Other candidates for regenerating large-scale magnetic fields in shearing turbulence include the incoherent $\\alpha\\Omega$-effect (Vishniac \\& Brandenburg \\cite{VishBran1997}; Proctor \\cite{Proctor2007}) and the shear--current effect (Rogachevskii \\& Kleeorin \\cite{RogaKle2003,RogaKle2004}).  Numerical simulations of convection have been used to study dynamo action in local (e.g.\\ Cattaneo \\& Hughes \\cite{CattHugh2006}; Tobias et al.\\ \\cite{Tobiasea2008}) and global (e.g.\\ Brun et al.\\ \\cite{BMT04}, Browning et al.\\ \\cite{Browning_etal06}, Brown et al.\\ \\cite{Brown_etal07}) settings. So far, however, with the exception of the last two papers, these models have had a hard time generating appreciable large-scale magnetic fields, although all the ingredients such as rotation and stratification necessary for an $\\alpha$-effect to occur have been present. Possible reasons for the lack of large-scale dynamo action in the numerical models is arguably the lack of scale separation between the energy carrying scale and the scale of the domain and the absence of large-scale shear in the system. In stratified convection, the lack of sufficient scale separation is difficult to tackle due to the prohibitive computational challenges involved when trying to cover many more vertical scale heights in the domain. Adding shear, however, is easier to accomplish, which is one of the principal aims of the present paper.  An additional point concerning large-scale dynamo action in turbulent convection is the possibility of catastrophic quenching due to small-scale magnetic fields. It has long been argued that small-scale fields cause catastrophic quenching of the large-scale dynamo (e.g.\\ Vainshtein \\& Cattaneo \\cite{VainCatt1992}). This sort of quenching, however, only occurs in specific circumstances, i.e.\\ when there is no magnetic helicity flux out of the domain. This behaviour can be understood in terms of magnetic helicity (Brandenburg \\& Subramanian \\cite{BS05}) which is a conserved quantity in ideal MHD. If magnetic diffusivity is finite, the magnetic helicity can only change diffusively under the aforementioned special circumstances that lead to resistively slow saturation of the large-scale field.  If, however, magnetic helicity is driven out of the domain, catastrophic quenching can be alleviated and the large-scale dynamo is expected to saturate near equipartition field strengths. A promising mechanism capable of driving magnetic helicity fluxes was introduced by Vishniac \\& Cho (\\cite{VishCho2001}), who found that in the presence of shear the flux follows isocontours of constant velocity (Brandenburg \\& Subramanian \\cite{BS05b}; Subramanian \\& Brandenburg \\cite{SB06}). We expect that this shear-driven magnetic helicity flux is of crucial importance for large-scale dynamo action.  The remainder of the paper is organised as follows: the numerical model is described in Sect.~\\ref{sec:model}, and the results and conclusions are presented in Sects.~\\ref{sec:results} and \\ref{sec:conclusions}, respectively. ",
        "conclusions": "\\label{sec:conclusions}  The present simulations provide a clear demonstration that convection can produce large-scale magnetic fields of equipartition strengths in an open Cartesian domain with shear crossing the surface. Our results also demonstrate that this is not possible within a comparable time span using a closed domain with perfectly conducting boundaries. The sensitivity on boundary conditions has also been noted in earlier work with forced turbulence and shear (Brandenburg \\cite{Brandenburg2005}). This can be interpreted as being due to magnetic helicity fluxes allowing the system to dispose of excess small-scale magnetic helicity along contours of constant shear (Vishniac \\& Cho \\cite{VishCho2001}; Brandenburg \\& Subramanian \\cite{BS05b}; Subramanian \\& Brandenburg \\cite{SB06}). In our case this is the vertical direction, but in Tobias et al.\\ (\\cite{Tobiasea2008}), it is the horizontal direction. Thus, in our case shear-driven magnetic helicity flux can escape the domain whereas in theirs it cannot due to the periodicity in the horizontal directions. This could explain why no large-scale dynamo develops in their case. Their model is otherwise similar to ours (same values of $\\Co$ and $\\Pm$, whilst $\\Rm$ is about two times smaller, and $\\Sh$ is twice as large.) Control simulations show that replacing their periodic boundary condition in the $x$ direction by a normal-field boundary condition allows strong large-scale fields also in their setup (Fig.~\\ref{fig:pvs}).  The precise nature of the dynamo found here cannot be pinned down rigorously unless one is able to identify unambiguously the mechanisms that are actively at work. In addition to the $\\alpha$-effect, there are other possible alternatives, in particular the incoherent alpha-shear dynamo (Vishniac \\& Brandenburg \\cite{VishBran1997}; Proctor \\cite{Proctor2007}), which could produce a large-scale magnetic field even without stratification and hence no regular $\\alpha$-effect. This interpretation has been favoured in some recent work where the shear--current effect was found not to be excited (Brandenburg et al.\\ \\cite{Brandea2008}). However, this matter is still under debate and there are arguments in favour of a shear--current effect (e.g.\\ Kleeorin \\& Rogachevskii \\cite{KleRoga2008}).  In order to disentangle the relative roles of different mechanisms in the present work it will be important to establish a closer connection with mean-field theory by determining the components of the $\\alpha$ and turbulent magnetic diffusivity tensors for different boundary conditions and as a function of $\\Rm$. Only when the results are shown to be converged with $\\Rm$ may we expect them to be relevant for the Sun and perhaps other astrophysical bodies.  We emphasize that large-scale dynamo action is normally only seen in the nonlinear stage, whilst the linear stage is dominated by small-scale magnetic fields. It is therefore important to consider mean-field transport coefficients that are affected by the magnetic field. It is in principle even possible that the relevant dynamo mechanism is an intrinsically nonlinear one, as it is in the case of accretion disc turbulence where the turbulence is the result of the dynamo itself (Brandenburg et al.\\ \\cite{Brandea1995}). The present results do not give any indications in this direction, because there seems to be a continuous transition from low-$\\Rm$ to higher-$\\Rm$ large scale dynamo action. For example at $\\Rm=11$, small-scale dynamo action in the usual sense is not expected, and nevertheless, the final saturation strength of the large-scale field is similar to the case of higher-$\\Rm$ dynamos.  Another important aspect to consider in future work is the degree of scale separation. As the Reynolds number increases, the turbulence becomes more vigorous and the convection plumes will no longer penetrate the entire depth of the unstable layer. Therefore, we expect the wavenumber of the energy-carrying eddies to increase, i.e.\\ $\\kef>2\\pi/d$. This means that the degree of scale separation, i.e.\\ the ratio $\\kef/k_1$, will also increase. However, as discussed at the end of Sect.~\\ref{BoxSize}, more scale separation is only expected to enhance large-scale dynamo action."
    },
    "0806/0806.4626_arXiv.txt": {
        "abstract": "We discuss the smooth hybrid inflation scenario in the context of a simple supersymmetric SO(10) GUT in 5D orbifold. In this GUT model, the SO(10) gauge symmetry is broken down to the Pati-Salam (PS) gauge group, SU(4)$_c \\times$ SU(2)$_L \\times$ SU(2)$_R$, by orbifold boundary conditions and all matter and Higgs multiplets are placed only on the brane (PS brane) where only the PS symmetry is manifest. Further breaking of the Pati-Salam group to the Standard Model one is realized by VEVs of the Higgs multiplets $({\\bf 4},{\\bf 1},{\\bf 2}) \\oplus (\\overline{{\\bf 4}},{\\bf 1},{\\bf 2})$. The gauge coupling unification is successfully realized at $M_{\\rm GUT} =4.6 \\times 10^{17}$ GeV after incorporating the threshold corrections of the Kaluza-Klein modes, with the compactification scale (assumed to be the same as the PS symmetry breaking scale) $M_c = v_{\\rm PS}= 1.2 \\times 10^{16}$ GeV. We show that this orbifold GUT model can naturally leads us to the smooth hybrid inflation, which tunes out to be consistent with the WMAP 5-year data with the predicted $M_{\\rm GUT}$ and $v_{PS}$ in the model. ",
        "introduction": "In these decades, terrestrial observational data on neutrino oscillations, B physics as well as the astrophysical ones like Wilkinson Microwave Anisotropy Probe (WMAP) give the very important informations on physics beyond the Standard Model (SM), especially on the grand unified models (GUTs). Among several GUTs, the model based on the gauge group SO(10) is particularly attractive. In fact, SO(10) is the smallest simple gauge group under which the entire SM matter contents of each generation are unified into a single anomaly-free irreducible representation, ${\\bf 16}$. In these SO(10) GUTs, the so-called renormalizable minimal SO(10) model (hereafter the minimal SO(10) GUT) has been paid a particular attention, where two Higgs multiplets $\\{{\\bf 10} \\oplus {\\bf \\overline{126}}\\}$ are utilized for the Yukawa couplings with matters ${\\bf 16}_i$ ($i=1,2,3$ is the generation index) \\cite{Babu} \\cite{Fukuyama1} \\cite{Fukuyama2}. A remarkable feature of the model is its high predictive power for the neutrino oscillation parameters with reproducing charged fermion masses and mixing angles. However, after KamLAND data \\cite{Eguchi:2002dm} was released, it entered to the stage of precision measurements, and many authors performed data-fitting analysis to match up these new data.   Also, Higgs superpotential in the minimal SO(10) GUT has been constructed and the detailed analysis of symmetry breaking patterns have been extensively studied \\cite{Fukuyama:2004xs, Bajc:2004xe}. This construction gives the vacuum expectation values (VEVs) at intermediate energy scales unlike normal SUSY GUTs, which gives rise to a trouble in the gauge coupling unification as well as the necessary scales of the seesaw mechanism and leptogenesis. This mismatch of gauge couplings has been explicitly shown in Ref.~\\cite{Bertolini}, where they are not unified any more and even the SU(2) gauge coupling blows up far below the GUT scale.   In addition to the issue of the gauge coupling unification, the minimal SO(10) model potentially suffers from the problem that the gauge coupling blows up around the GUT scale. This is because the model includes many Higgs multiplets of higher dimensional representations. In field theoretical point of view, this fact implies that the GUT scale is a cutoff scale of the model, and more fundamental description of the minimal SO(10) model would exist above the GUT scale.   In order to solve these problems, the minimal SO(10) GUT has been considered in 5D \\cite{F-K-O} with the warped background geometry \\cite{RS}, where the GUT gauge symmetry is assumed to be broken by VEVs of Higgs multiplets on a brane, as usual in 4D models. Another possibility of constructing GUT models in extra-dimensions is to consider the so-called orbifold GUT \\cite{Kawamura}, where the GUT gauge symmetry is (partly) broken by orbifold boundary conditions. In this paper, we consider a class of SO(10) models in 5D \\cite{Raby}, where SO(10) gauge symmetry is broken into the PS gauge group and further symmetry breaking into the SM gauge group is achieved by VEVs of Higgs multiplets on a brane. In particular, we concentrate on the recently proposed simple SO(10) model \\cite{F-O}. In this model, all matter and Higgs multiplets reside only on a brane (PS brane) where the PS gauge symmetry is manifest, so that low energy effective description of this model is nothing but the PS model in 4D with a special set of matter and Higgs multiplets. At energies higher than the compactification scale, the Kaluza-Klein (KK) modes of the bulk SO(10) gauge multiplet are involved in the particle contents and in fact, the gauge coupling unification was shown to be successfully realized by incorporating the KK mode threshold corrections into the gauge coupling running \\cite{F-O}. The unification scale ($M_{\\rm GUT}$) and the compactification scale ($M_c$) which was set to be the same as the PS symmetry breaking scale ($v_{\\rm PS}$) were found to be $M_{\\rm GUT}=4.6 \\times 10^{17}$ GeV and $M_c=v_{\\rm PS}=1.2 \\times 10^{16}$ GeV.   In this paper, we apply this SO(10) model to the inflationary scenario. The idea of inflation \\cite{InflationRev} has been strongly favored from the view points of not only providing the solutions to the horizon and flatness problems of the standard big bang cosmology but also recent precise cosmological observations on the the cosmic microwave background radiation and the large scale structure in the Universe. Therefore, it is an important task to construct a realistic inflation model based on some well-motivated particle physics model. Among many proposed inflation models, the hybrid inflation \\cite{Linde} \\cite{HIRev} is particularly  attractive because it can be adjusted to a wide class of SUSY models \\cite{Hybrid1}. Variants of hybrid inflation, in particular, applicable to SUSY GUT models have been proposed: The standard \\cite{StdHI}, shifted \\cite{ShiftedHI} and smooth \\cite{SmoothHI} hybrid inflation models. Some of these models are based on the SUSY PS model with one singlet and Higgs multiplets $({\\bf 4},{\\bf 1},{\\bf 2}) \\oplus (\\overline{{\\bf 4}},{\\bf 1},{\\bf 2})$ whose VEVs break the PS symmetry to the SM one. Interestingly, except for the singlet field, the orbifold GUT model of Ref.~\\cite{F-O}, which we are interested in, has the same particle content. Therefore, the GUT model can naturally incorporate the hybrid inflation in it. In the following, we consider the smooth hybrid inflation \\cite{SmoothHI} in this orbifold GUT framework.   This paper is organized as follows: In Section 2 we briefly review the model proposed in Ref.~\\cite{F-O}. Application of this model to inflation is developed in Section 3 and the consistency of the model with current observations is shown. The last section is devoted for conclusions. ",
        "conclusions": "We have discussed the smooth hybrid inflation scenario in the context of the SO(10) GUT model in Ref.~\\cite{F-O}. We have shown that the model can be naturally applicable to the inflation model by introducing only one singlet chiral multiplet, keeping the structure of the gauge coupling unification unchanged. The analysis for the running gauge coupling fixes the model parameters as $v_{\\rm PS}= 1.2 \\times 10^{16}$ GeV and $M_{\\rm GUT} = 4.6 \\times 10^{17}$ GeV. It is very interesting that these parameters are determined independently of the cosmological considerations, nevertheless the inflation scenario with these parameters fits very well with the current cosmological observations, as we have shown."
    },
    "0806/0806.1529_arXiv.txt": {
        "abstract": "We present 1420~MHz polarization images of a $2.5\\arcdeg\\times2.5\\arcdeg$ region around the planetary nebula (PN) \\sh.  The images are taken from the Canadian Galactic Plane Survey (CGPS).  An arc of low polarized intensity (size $0.2\\arcdeg \\times 0.7\\arcdeg$) appears prominently in the north-east portion of the visible disk of \\sh, coincident with the optically identified interaction region between the PN and the interstellar medium (ISM). The arc contains structural variations down to the $\\sim$1\\arcmin\\@ resolution limit in both polarized intensity and polarization angle. Several polarization-angle ``knots'' appear along the arc.  By comparison of the polarization angles at the centers of the knots and the mean polarization angle outside \\sh, we estimate the $RM$ through the knots to be $-43 \\pm 10\\ \\rm{rad}\\,\\rm{m}^{-2}$.  Using this estimate for the $RM$ and an estimate of the electron density in the shell of \\sh, we derive a line-of-sight magnetic field in the interaction region of $5.0 \\pm 2.0$~$\\mu\\rm{G}$.  We believe it more likely the observed magnetic field is interstellar than stellar, though we cannot completely dismiss the latter possibility.  We interpret our observations via a simple model which describes the ISM magnetic field around \\sh, and comment on the potential use of old PNe as probes of the magnetized ISM. ",
        "introduction": "}  The diffuse Galactic synchrotron radiation provides a continuous background of radio emission which is intrinsically highly (up to $\\approx$70\\%) linearly polarized.  This radiation is Faraday-rotated from the point of emission as it propagates through warm ionized gas interwoven with magnetic fields in the disk of the Galaxy; i.e., the angle $\\theta$ of the polarized component of the emission is rotated at wavelength $\\lambda$\\,[m] by \\begin{equation} \\Delta\\theta = RM\\,\\lambda^2\\ [\\rm{rad}], \\end{equation} where $RM$ is the rotation measure [$\\rm{rad}\\,\\rm{m}^{-2}$] and depends on the line-of-sight component of the magnetic field, $B_{\\|}$\\,$[\\mu\\rm{G}]$, the thermal electron density, $n_e$\\,$[\\rm{cm}^{-3}]$, and the path length, $dl$\\,$[\\rm{pc}]$, as \\begin{equation} RM = 0.81\\int{B_{\\|}\\,n_e\\,dl\\ [\\rm{rad}\\,\\rm{m}^{-2}]}. \\end{equation} High-resolution radio polarization images at frequencies $\\lesssim$3~GHz reveal the turbulent imprint of Faraday rotation on the diffuse polarized emission (e.g., \\citealt{Wieringa+1993}; \\citealt{Gray+1999}; \\citealt{Gaensler+2001}; \\citealt{Uyaniker+2003}; \\citealt*{HaverkornKd2003a,HaverkornKd2003c}; \\citealt{Haverkorn+2006b}; \\citealt{Schnitzeler+2007}).  The turbulent nature of the imprint is the product of the random component of the Galactic magnetic field and irregular electron-density distributions in the general interstellar medium (ISM).  Detailed studies and modeling of the diffuse Galactic emission (e.g., \\citealt{Spoelstra1984}; \\citealt*{HaverkornKd2004b}) as well as statistical analyses of the $RM$s of polarized extragalactic sources \\citep{Haverkorn+2006a} indicate that the scale size, or ``cell'' size, for variations in the magnetized ISM range from $\\sim$15~pc to 100~pc.  Depth depolarization then results from the averaging of nonparallel polarization vectors from emission at different cells along the line-of-sight.  Smaller-scale variations are also apparent in polarization images.  Depolarization filaments, or ``canals,'' with a width corresponding to one beam of the observing instrument, indicate the presence of very sharp gradients in $RM$ (see, e.g., \\citealt{Gaensler+2001}; \\citealt{Uyaniker+2003}; \\citealt*{HaverkornKd2004a}).  If the gradient is so steep across the beam as to cause differential rotation of the polarization angle of $\\sim$90\\arcdeg, then complete depolarization occurs.  Beam depolarization in images produced by aperture synthesis telescopes (typically with arcminute resolution) suggests a scale length for $RM$ structures in the magnetized ISM of less than 1~pc.  Interpreting the structure seen in radio polarization images of the Galactic plane is largely left to modeling \\citep[e.g.,][]{HaverkornKd2004b}, as there is generally little correlation between the polarization structures seen in these images and the emission structures seen in total intensity images. Nevertheless, objects of known distance can be used to estimate the line-of-sight distribution of the magnetized ISM, revealing, for example, whether observed polarization structures are generated behind the known object, in the region between the object and the Sun, or perhaps in the immediate vicinity of the object \\citep[see][]{Gray+1999,UyanikerL2002b}.  Moreover, if small-scale depolarization structures can be isolated to a known object, then it may be possible to probe directly the properties of the magnetized ISM within the structures.  \\ion{H}{2} regions and supernova remnants (SNRs) are the two most prevalent (discrete) constituents of the ISM as seen at radio wavelengths, and each class of object is detected in radio polarization images up to a limiting distance determined by the ``polarization horizon'' \\citep*[see][]{KothesL2004}. \\ion{H}{2} regions are detected by way of their depolarizing effects on background diffuse emission, while SNRs are simultaneously a source of polarized synchrotron emission and a Faraday ``screen'' which depolarizes background emission.  However, neither class of object is a particularly good probe of the magnetized ISM.  \\ion{H}{2} regions have very high electron densities and turbulent motions which produce tangled magnetic fields, a combination which leads to virtually complete beam depolarization across the region.  For SNRs, complex models are needed to describe the physical parameters in the shock front at the interface between the rapidly expanding SNR and the ISM.  Another class of object which may potentially serve as a better probe of the magnetized ISM is planetary nebulae (PNe).  Young PNe are relatively strong thermal radio emitters, resulting from high electron densities \\citep[see, e.g.,][]{Bains+2003}, but they are also very compact ($\\ll$1~pc) and not yet interacting with the ISM.  On the other hand, the shells of many old PNe are observed at optical wavelengths to interact with the ISM \\citep[see][]{TweedyK1996}. Moreover, the ISM magnetic field appears to play a significant role in shaping the shells, as evidenced primarily by the visible ``striping'' or filamentary structure of shell gases \\citep*[see][]{TweedyMN1995,SokerZ1997}.  Theoretical treatments show that interactions between PNe and the ISM are an important consideration in the evolution of PN systems moving at even modest speeds ($\\gtrsim$5~$\\rm{km}\\,\\rm{s}^{-1}$) with respect to the ISM \\citep*[see, e.g.,][]{SokerD1997,WareingZO2007}.  If the conditions are right in the interaction region between the PN and the ISM, we may expect to see the Faraday signature of old PNe in radio polarization images.  Such a signature has been identified for the nearby PN Sharpless~2-216 (\\sh) and was first described by one of us in \\citet{Uyaniker2004}.  In this paper, we describe in detail the Faraday-rotation structure in the shell of \\sh.  We present radio polarization images at 1420~MHz of a $2.5\\arcdeg\\times2.5\\arcdeg$ region of the Galactic plane around the position of \\sh.  The images are taken from the Canadian Galactic Plane Survey (CGPS).  In \\S~\\ref{s216}, we summarize the pertinent properties of \\sh.  In \\S~\\ref{obs}, we describe briefly the preparation of the images.  In \\S~\\ref{results}, we describe the structures observed on the visible disk of \\sh\\@ in both polarized intensity and polarization angle, and estimate the $RM$ through the shell of \\sh.  In \\S~\\ref{discuss}, we derive the magnetic field in the shell of \\sh, and interpret the structures and $RM$s in the context of an interaction between the PN and the ISM.  We also discuss the possibility that the observed structures are produced by the stellar field of the host white dwarf or its progenitor.  Finally, in \\S~\\ref{concl}, we summarize our conclusions. ",
        "conclusions": "  1. We presented 1420~MHz polarization images for the $2.5\\arcdeg\\times2.5\\arcdeg$ region in the CGPS around the PN \\sh.  2. A low-polarized-intensity arc, $0.2\\arcdeg \\times 0.7\\arcdeg$ in size, appears in the north-east portion of the visible disk of \\sh. The arc is coincident with the optically-identified interaction region between the PN and the ISM.  3. A second low-polarized-intensity arc appears in the north-west portion of the visible disk of \\sh.  4. The north-east arc contains structural variations down to the $\\sim$1\\arcmin\\@ resolution limit in both polarized intensity and polarization angle.  Several polarization-angle ``knots'' appear along the arc.  5. Via comparison of the polarization angles at the centers of the knots in the north-east arc and the mean polarization angle outside \\sh\\@ (and above $b \\simeq +0.5\\arcdeg$), we estimated the $RM$ through the knots to be $-43 \\pm 10\\ \\rm{rad}\\,\\rm{m}^{-2}$.  5. Using this estimate for the $RM$ and an estimate of the electron density in the shell of \\sh, we derived a line-of-sight magnetic field in the interaction region of $5.0 \\pm 2.0$~$\\mu\\rm{G}$.  6. We believe it more likely the derived magnetic field is interstellar than stellar, though we cannot completely dismiss the latter possibility.  We interpret our observations via a simple model which qualitatively describes the ISM magnetic field around \\sh.  7. It is unclear whether old PNe like \\sh\\@ could be useful probes of the magnetized ISM.  Targeted polarimetric observations at high resolution ($<$1\\arcmin), and possibly at multiple frequencies in the range 1--3~GHz, may help separate the signatures of more PNe from the turbulent ISM."
    },
    "0806/0806.3971_arXiv.txt": {
        "abstract": "As the most common environment in the universe, groups of galaxies are likely to contain a significant fraction of the missing baryons in the form of intergalactic gas.  The density of this gas is an important factor in whether ram pressure stripping and strangulation affect the evolution of galaxies in these systems.  We present a method for measuring the density of intergalactic gas using bent-double radio sources that is independent of temperature, making it complementary to current absorption line measurements. We use this method to probe intergalactic gas in two different environments: inside a small group of galaxies as well as outside of a larger group at a $2 \\mpc$ radius and measure total gas densities of \\nfirst $\\cmc$ and \\nsdss $\\cmc$ (random and systematic errors) respectively.  We use X-ray data to place an upper limit of $2 \\times 10^6$ K on the temperature of the intragroup gas in the small group.  \\end {abstract} ",
        "introduction": "In the local universe, attempts to measure the baryon content in the form of stars, hot X-ray emitting gas, and cold gas account for only one third of the baryon density seen at high redshift \\citep{2004ApJ...616..643F}.  Simulations predict that the remaining ``missing baryons\" exist in a warm-hot intergalactic medium (WHIM) that was shock heated during large-scale structure formation to a temperature range of $10^5 < T < 10^7$ K \\citep{2006ApJ...650..560C,2001ApJ...552..473D}.  Observational confirmation of the WHIM comes primarily from UV absorption lines in the spectra of low redshift quasars \\citep{2000ApJ...534L...1T}.  However, UV absorption line observations are only able to probe a limited temperature range of WHIM gas and there is continued debate as to whether the absorption originates in the extended gaseous halos of individual galaxies, the intergalactic medium (IGM), or from large-scale filaments \\citep{2003ApJS..146..125S}.  Galaxy groups, the most common environment in the local universe, likely contain a significant fraction of the total baryonic mass within their intragroup medium \\citep{1998ApJ...503..518F}.  Until now the only probes of this gas have been UV absorption line and X-ray observations.  In the few groups that have a hot intragroup medium X-ray observations can determine the emission measure, but without some knowledge of the geometry the density of this gas cannot be disentangled.  The X-ray surface brightness typically falls below detectable levels at radii much smaller than the virial radius, while the implied gas mass is still increasing at least linearly with radius \\citep{2000ARA&A..38..289M}.  Groups with a bright quasar behind them have been observed in an attempt to detect absorption from the intragroup medium.  In cases where intervening gas is detected, assumptions about the extent of the absorbing system, the metallicity, ionization fraction, and the intensity of the ionizing radiation field lead to estimates of the intragroup medium density on the order of $n\\sim10^{-4}-10^{-5}~\\mathrm{cm}^{-3}$ \\citep{2004AJ....127..199P,1998ASPC..143..261T}.  Additionally, galaxy groups, according to the hierarchical scenario of the formation of large-scale structure, are the building blocks of rich clusters of galaxies.  Groups are important sites in which to investigate the physical mechanisms responsible for the observed morphology and SFR-density relations \\citep{2003MNRAS.346..601G,2003ApJ...584..210G}.  However, the effectiveness of mechanisms like ram pressure stripping and strangulation on galaxy evolution in the group environment are not well understood due to the uncertainty in the density of the intragroup medium.  Some observational evidence does exist for the stripping of hot gas from galaxies by the intragroup medium, mainly in features like X-ray shocks and tails \\citep{2004ApJ...617..262S,2006MNRAS.370..453R},  as well as HI observations that show compression features opposite extensions in cold disk gas \\citep{2005A&A...435..483K,2007ApJ...671L..33M}.  Another way to measure the density of the IGM is through its effect on the jets of radio galaxies.  As a double lobed radio source travels through intergalactic gas ram pressure causes its jets to be swept back.  Prior to \\citet{1987AJ.....94..587B} it was believed that the combination of the dense intracluster medium and large galaxy velocities needed to form these bent-double radio sources could only be found in rich clusters.  While many bent-doubles do reside in dense clusters, a surprising number are found in lower mass environments like groups of galaxies \\citep{1978A&A....69..253E,1994ApJ...436...67V,1995AJ....110...46D,2001AJ....121.2915B}.  We present measurements of the density of intergalactic gas from radio, optical, and X-ray observations of two bent-double radio sources and their surrounding environments.  Both sources were identified by \\citet{2001AJ....121.2915B} as existing in environments with richness less than Abell class 0. We adopt $H_0=73 \\kmsmpc$. ",
        "conclusions": "We present two measurements of the density of intergalactic gas using radio sources whose jets are bent back by ram pressure.  S1 is probing gas with a density of \\nfirstg (\\nfirst $\\cmc$ assuming a mean molecular weight of $0.6$ and where the first set of errors are random and the second systematic) that is near the center of a small group of galaxies.  In the case of S2 this source is probing gas that is a projected distance of $2 \\mpc$ from the center of a system of galaxies with a density of \\nsdssg (\\nsdss $\\cmc$).  We have not accounted for thermal pressure in the jets, as well as our inability to resolve their true width.  Additional thermal pressure or smaller jets will necessitate a larger IGM density to produce the observed jet geometry.  If we have overestimated the radio galaxy speed or projection effects have led us to a smaller radius of curvature then the necessary density will decrease.  The systematic errors reflect the range of IGM densities possible given the range of jet speeds discussed in Section 3.  These speeds are measured by assuming that flux asymmetries between the jet and counter-jet are the result of relativistic beaming effects.  This beaming depends on the angle of the jets from our line of sight as well as the jet speed.  By looking at large samples of objects and assuming a uniform distribution of angles, a range of jet speeds can be determined.  We lack the inclination information necessary to determine jet speeds for individual sources.  \\citet{2004MNRAS.351..727A} point out that their distribution of jet speeds is asymmetric with the peak toward lower speeds and that selection effects may be pushing the mean toward higher speeds. Slower jet speeds would cause our IGM densities to tip towards the lower end of the range in values.  The inclination of the source also affects the measured radius of curvature. We use a value which is uncorrected for inclination to calculate the IGM densities.  S1 appears relatively unprojected whereas S2 is less easily characterized.  If we assume that a Hubble flow distance is accurate for S2, and that it is traveling along the vector connecting its current position with the cross indicating the center of the nearby system of galaxies, then its inclination from the plane of the sky is $\\sim 75^{o}$.  Deprojecting according to this angle leads to a radius of curvature four times smaller and a corresponding IGM density four times larger.  We can compare the measured IGM densities to those seen in clusters and X-ray bright groups of galaxies where the radial density profile of the hot gas in the IGM is traced using X-ray data.  Gas densities at the center of clusters are on the order of $10^{-25} \\ \\gcmc$ ($0.1 \\cmc$) to $10^{-28} \\gcmc$ ($10^{-4} \\cmc$) at a radius of $1 \\mpc$ \\citep{2006ApJ...640..691V}.  The IGM density we measure near S1 is comparable to the density seen in X-ray observations at a similar radius in the NGC 5044, NGC 533, and ESO 5520200 groups from \\citet{2007ApJ...669..158G}.  Those groups have higher velocity dispersions and may be dynamically older than the group containing S1.  Numerical simulations of the evolution of the IGM predict WHIM gas in the local universe with a density distribution that is broadly peaked at overdensities of $10-20$ and densities in filaments of $10-100$ relative to the critical density ($1.0\\times10^{-29} \\gcmc$ for $H_0=73 \\kmsmpc$) \\citep{2006ApJ...650..560C, 2006MNRAS.370..656D}.  The gas that S2 is tracing has an overdensity of $\\sim 90$.  Considering that this source is $2 \\mpc$ from the center of a large group of galaxies it could be tracing gas in a filament or gas that is only tenuously associated with the nearby group.  Radio galaxies with bent jets in groups of galaxies provide evidence that intragroup gas exists in these systems even when the IGM is not X-ray bright.  The densities derived in this paper indicate that groups of galaxies and their surroundings are likely to contain significant reservoirs of baryons.  Further work is underway with a larger number of sources and will attempt to quantify the contribution of the intragroup medium to the baryon density in the local universe."
    },
    "0806/0806.0118_arXiv.txt": {
        "abstract": "We study effects of winds on the time evolution of isothermal,  self-gravitating accretion discs by adopting a radius dependent mass loss rate because of the existence of the wind. Our similarity and semi-analytical solution describes time evolution of the system in the slow accretion limit. The disc structure is distinct in the inner and outer parts, irrespective of the existence of the wind. We show that existence of wind will lead to a reduction of the surface density in the inner and outer parts of the disc in comparison  to a no-wind solution. Also, the radial velocity significantly increases in the outer part of the disc, however, the accretion rate decreases due to the reduced surface density in comparison to the no-wind solution.  In the inner part of the disc, mass loss due to the wind is negligible according to our solution. But the radial size of this no-wind inner region becomes smaller for stronger winds. ",
        "introduction": "\\label{sec:1} The existence of outflow or wind in many accreting systems  is supported by strong observational evidences. Outflows from Active Galactic Nuclei (AGN) are much more common than previously thought: the overal fraction of AGNs with outflows is fairly constant, approximately $60\\%$, over many order of magnitude in luminosity (Ganguly \\& Brotherton  2008). These mass-loss mechanisms are also observed in microquasars, Young Stellar Object (YSO) and even from brown dwarfs (Ferrari 1998; Bally, Reipurth \\& Davis 2007; Whelan et al 2005). About 30$\\%$ of T Tauri stars present bipolar ejection and this percentage increases to $100\\%$ for class $0$ objects, the earliest stage of star formation.  It is now widely accepted that winds or outflows have their origin in accretion flows (e.g., Blandford \\& Payne 1982). They may provide an additional important sink of mass, angular momentum and energy which their dynamical influence to the accretion flow can not be neglected. However, it remains unclear as how some  part of the accreting gas can be transferred into winds or outflows.    Nevertheless, deviations from Keplerian rotation in some AGNs and the flat infrared spectrum of some T Tauri stars can both be described by self-gravitating discs. Vertical structure of an accretion disc under the influence of self-gravity was studied by Paczy$\\acute{\\rm n}$ski (1978). Some authors studied the effects related to self-gravity of the disc in the radial direction (e.g., Bodo \\& Curir 1992). New class of self-gravitating discs has also been proposed, in which the energy equation is replaced by a self-regulation condition (Bertin \\& Lodato 1999; Lodato \\& Bertin 2001). Mineshige \\& Umemura (1997) extended the previous steady state solutions to the time-dependent case while the effect of the self-gravity of the disc was taken into account. They used an isothermal equation of state, and so their solutions describe a viscous accretion disc in the slow accretion limit. Also, Tsuribe (1999) studied the self-similar collapse of an isothermal viscous accretion disc.   For simplicity, outflow is neglected in most of the steady-state or time-dependent theoretical models of the self-gravitating discs. However, some authors studied the effect of wind or outflow on the radial structure of non-self-gravitating accretion discs (e.g., Knigge 1999). Recently, Combet \\& Ferreira (2008) studied the structure of YSO accretion discs in an approach that takes into account the presence of the protostellar jets. They showed that discs with jet  presents structure different from the standard accretion disc  because of the influences of jets on the radial structure of the disc. Ruden (2004) studied the physics of protoplanetary disc evolution in the presence of a photoevaporative wind.  In this paper, we generalize time-dependent solutions of Mineshige \\& Umemura (1997) for a self-gravitating accretion disc to include winds or outflows. Our parametric model describes mass and angular momentum loss by wind in an isothermal, self-gravitating disc, yet applicable to many types of dynamical disc-wind models. Basic equations are presented in the next section. Properties of our semi-analytical are discussed in section 3. A summary of implications of the results are given in section 4. ",
        "conclusions": "Our similarity solution show that the transition radius increases linearly with time, i.e. $r_{\\rm tr} \\approx \\alpha' c_{\\rm s} t$. Assuming that the radius $r_{\\rm tr}$ is very close to the central part of the disc initially, there is not \"no-wind\" region initially and wind exists at all radii of the disc. At early times of evolution, wind does not modify the surface density significantly. As gas accretes toward the central parts, the surface density profile changes from inside. However,  the accretion rate is reduced  initially because of the wind, in comparison to the no-wind solution. Some fraction of the accreted mass and the angular momentum can be carried away to infinity by the wind, while leaving the remaining flow with smaller density at all parts and enhanced radial velocity at the outer part of the disc.  \\begin{table*}\\label{t} \\centering \\begin{minipage}{150mm} \\caption{Integrals $\\mathcal{I}$ and $\\mathcal{J}$ for the solutions with different input parameters. In the absence of winds/outflows, we have ${\\mathcal I}=1.28$ and ${\\mathcal J}=0$.} \\begin{tabular}{ c c c c   c c c c c c c c} \\hline\\hline &$(\\Gamma_{0},l)=(0.1,1)$&&\\vline&&$(\\Gamma_{0},s)=(0.1,0.7)$&&\\vline&&$(l,s)=(1.0,0.7)$&\\\\ \\hline $s$&${\\mathcal I}$&${100 \\mathcal J}$& \\vline &$l$&${\\mathcal I}$&${100\\mathcal J}$&\\vline&$\\Gamma_{0}$&${\\mathcal I}$&${100\\mathcal J}$&\\\\ $0.1$&$0.83$&$7.19$&\\vline&$0.0$&$1.84$&$9.81$&\\vline&$0.01$&$0.88$&$4.49$\\\\ $0.5$&$0.87$&$5.08$&\\vline&$1.0$&$0.88$&$4.49$&\\vline&$0.05$&$1.04$&$2.66$\\\\ $0.7$&$0.88$&$4.49$&\\vline&$1.5$&$0.58$&$2.86$&\\vline&$0.1$&$1.22$&$0.06$\\\\ \\hline\\hline \\end{tabular} \\end{minipage} \\end{table*}   We can apply our similarity solutions to a self-gravitating disc just before forming a central star, namely, in the runaway collapse phase. This phase of evolution would correspond to the very early (class 0) stage. A molecular cloud core can be approximated by an isothermal gas with sound speed $c_{\\rm s} \\sim 0.2 \\rm Km$ s$^{-1}$ (e.g., Hayashi \\& Nakano 1965). Typical age is considered to be $5\\times 10^{4}$ yr. The radius $r=0.01$ pc at $t_{0}=5\\times 10^{4}$ yr corresponds to $x=1$ in our similarity variable. Then, the radius of the disc $r_{\\rm d}$ becomes \\begin{equation} r_{\\rm d} = 0.01 {\\rm pc} (\\frac{x}{1}) (\\frac{c_{\\rm s}}{0.2 {\\rm Km s^{-1}}})(\\frac{t}{5\\times 10^{4} {\\rm yr}}). \\end{equation}   Now, we can calculate the mass of the disc as follows \\begin{equation} M_{\\rm disc}=\\int_{0}^{r_{\\rm d}} 2\\pi \\Sigma r dr, \\end{equation} or \\begin{equation} M_{\\rm disc} = (9.51\\times 10^{-2} M_{\\odot})(\\frac{t}{5\\times 10^{4} {\\rm yr}}) \\mathcal{I},\\label{eq:Mdisc} \\end{equation} where ${\\mathcal I} = \\int_{0}^{1} x\\sigma dx$. Having our similarity solutions, we can simply calculate this integral for different set of input parameters. When there is not a wind or outflow, the integral becomes ${\\mathcal I} = 1.28$. But existence of a wind or outflow decreases the integral by a factor up to 2 depending on the input parameters (see Table 1). In other words, the mass of a self-gravitating disc with wind is reduced comparing to a similar disc but without wind. So, reduction to the mass of a disc with wind is mainly due to the decreased surface density at all  regions of the disc. This implies that the central mass object is forming up to two  times slower than a system without wind. Of course, this factor may increase if we consider cases, in which more angular momentum is carried away by wind (i.e., larger $l$). We can determine how long is needed to increase the mass of the disc by one solar mass. According to equation (\\ref{eq:Mdisc}), we can write \\begin{equation} \\tau \\simeq \\frac{0.52\\times 10^6}{\\mathcal{ I}} {\\rm yr}. \\end{equation} For a disc without wind, the mass of the disc increases by one solar mass within approximately  $\\tau \\simeq 4.0 \\times 10^{5} {\\rm yr}$. But the integral $\\mathcal{ I}$ decreases by a factor up to $2$ according to Table 1. This allows the formation of a central core with one solar mass in $2$ times slower than a similar system without wind, i.e. $\\tau \\simeq 9 \\times 10^{5} {\\rm yr}$. This result is consistent with numerical simulations of early stages of massive discs with winds/outflows (e.g., Banerjee \\& Pudritz 2007). Numerical study of Banerjee \\& Pudritz (2007) is concentrated towards analyzing massive star formation and inseparable links between gravitational collapse and early wind-driven outflows. They showed that the disc, in the early stages of formation of a high mass star, is more massive than the protostar that is forming within it (see also Banerjee \\& Pudritz 2006). Dominance of  mass of the disc is kept within $7\\times 10^{4} {\\rm yr}$ according to the simulations of Banerjee \\& Pudritz (2006). In our paper, we neglected mass of the central protostar at early stages evolution which is a good approximation at early stages of formation. Also, the other analytical studies show that the central core grows to one solar mass in $1.6\\times 10^{6}$ ${\\rm yr}$ if other effects are ignored (e.g., Tsuribe 1999)   Although we presented the ratio of the total mass loss rate by wind to the accretion rate at each radius of the disc in Figures {\\ref{fig:f1}}, {\\ref{fig:f2}} and {\\ref{fig:f3}}, it is desirable to  calculate how much mass is lost by wind during the early stages. Actually, observations of different systems show that the ratio of mass loss rate by wind to the accretion rate is around $0.1$ (e.g., Konigl \\& Pudritz 2000). Having the self-similar solutions, one can simply show that the total mass-loss rate by the wind is \\begin{equation} {\\dot M}_{\\rm w}=(1.9\\times 10^{-6} M_{\\odot}/{\\rm yr}) {\\mathcal J}, \\end{equation} where ${\\mathcal J}=\\int_{0}^{1} x \\sigma \\Gamma dx$. Obviously, exact value of this integral depends on the profiles of the physical quantities of the disc, i.e. $\\sigma$ and $\\Gamma$. Table 1 shows this integral for the solutions with different sets of the input parameters. We can simply show that during  formation of a core with one solar mass the amount of mass loss by the wind is ${\\mathcal J}/\\mathcal{I}$ times the accreted mass. This ratio varies approximately from $0.001$ to $0.05$ according to Table 1. Thus, our mass loss rate is consistent with observations of protostellar systems.       Our solutions are valid as long as a  central mass has not been formed and the outer disc is not depleted with gas. Thus, there is no surprise  that the similarity solution demands a diverging total mass like a solution without wind, if $x$ tends to {\\it infinity} (Mineshige \\& Umemura 1997). But we restricted our solutions within $x\\leq 1$, and so, the total mass of the disc is finite.  Once the outer part of the disc is depleted with gas and a central mass is formed, we can not apply our similarity solution with wind. However, the effects of wind on the early evolution of self-gravitating discs are remarkably worth to be considered, in particular reduced surface density and the accretion rate at all parts of the disc."
    },
    "0806/0806.0929_arXiv.txt": {
        "abstract": "{} { We provide a quantitative analysis of time-variable phenomena in the photospheric, near-star, and outflow regions of the late-B supergiant (SG) HD~199\\,478. This study aims to provide new perspectives on the nature of outflows in late-B SGs and on the influence of large-scale structures rooted at the stellar surface.} {The analysis is  based primarily on optical spectroscopic datasets secured between 1999 and 2000 from the Bulgarian NAO, Tartu, and Ritter Observatories. The acquired time-series  samples a wide range of weak metal lines, He~{\\sc i} absorption, and both emission and absorption signatures in H$\\alpha$. Non-LTE line synthesis modelling is conducted using FASTWIND for a strategic set of late-B SGs to constrain and compare their fundamental parameters within the context of extreme behaviour in the \\Ha lines.} {The temporal behaviour of HD~199\\,478 is characterised by three key empirical properties: (i) systematic central velocity shifts in the photospheric absorption lines, including C~{\\sc ii} and He~{\\sc i}, over a characteristic time-scale of $\\sim$ 20 days; (ii) extremely strong, variable \\Ha emission with no clear modulation signal, and (iii) the occurrence in 2000 of a (rare) high-velocity absorption (HVA) event in H$\\alpha$, which evolved over $\\sim$ 60 days, showing the clear signature of mass infall and outflows. In these properties HD~199\\,478 resembles few other late-B SGs with peculiar emission and HVAs in H$\\alpha$ (HD~91\\,619, HD~34\\,085, HD~96919). Different possibilities accounting for the phenomenon observed are indicated and briefly discussed.} {At the cooler temperature edge of B SGs, there are objects whose wind properties, as traced by H$\\alpha$, are inconsistent with the predictions of the smooth, spherically symmetric wind approximation. This discordance is still not fully understood and may highlight the role of a non-spherical, disk-like, geometry, which may result from magnetically-driven equatorial compression of the gas. Ordered dipole magnetic fields may also lead to confined plasma held above the stellar surface, which ultimately gives rise to transient HVA events.} ",
        "introduction": "The key limiting assumptions incorporated within current hot star model atmospheres include a globally stationary and spherically symmetric stellar wind with a smooth density stratification. Although these models are generally quite successful in describing the overall wind properties, there are numerous observational and theoretical studies which indicate that hot star winds are certainly not smooth and stationary. Most (if not all) of the time-dependent constraints refer, however, to O-stars and early B supergiants (SGs), while mid- and late-B candidates are currently under-represented in the sample of stars investigated to date.  \\citet{Vink00} have shown that for mid and late-B SGs there is a discrepancy between theoretical predictions and \\Ha mass-loss rates, derived by means of unblanketed model analysis. This finding was confirmed by recent investigations using line-blanketed model atmospheres (see, e.g., \\citealt{crowther06, MP}). The reason for the discrepancies is not clear yet but wind structure and variability might in principle cause them.  Indeed, observations indicate that while winds in late-B SGs are significantly weaker than those in O SGs (e.g. \\citealt{MP}), there is no currently established reason to believe that weaker winds might be less structured/variable than stronger ones (e.g. \\citealt{Markova05, puls06}).  The first extended spectroscopic monitoring campaigns of line-profile variability ($lpv$) in late-B SGs were performed by \\citet{kaufer96a, kaufer96b}, who showed that stellar winds at the cooler temperature edge of the B-stars domain can also be highly variable. Interestingly, in all 3 cases studied by these authors, the variability patterns (as traced by H$\\alpha$) were quite similar consisting of (i) blue- and red-shifted emission with $V/R$ variations  similar to those in Be-stars, and (ii) the sudden appearance of deep and highly blue-shifted absorptions (HVAs). Though the kinematic properties of the HVAs in \\Ha were found to be completely different from those of DACs (Discrete Absorption Components) in the UV spectra of O and early-B stars (e.g., HVAs do not propagate outwards, but instead extend to zero velocity and even indicate mass infall), similar scenarios consisting of large-scale wind structures rooted in the photosphere were suggested to interpret their appearance and development in time.  The present paper is focused on a multi-line investigation of HD~199\\,478 (HR~8020), a B8 Iae star whose stellar and wind properties have been recently determined by means of NLTE model atmosphere analysis of strategic lines in the optical \\citep{MP}. The first extensive monitoring campaign of $lpv$ in the optical spectrum of this star revealed \\Ha variability similar to those described by Kaufer et al. (i.e., peculiar wind emission with $V/R$ variations similar to those in Be stars) with one exception, i.e. no indications for any HVAs were found during that survey \\citep{MV}. \\begin{table*}[t] \\caption{Spectral observations and instruments employed.} \\label{log1} \\begin{tabular}{llllllll} \\hline \\hline \\multicolumn{1}{l}{Observatory} &\\multicolumn{1}{l}{} &\\multicolumn{1}{l}{Telescope} &\\multicolumn{1}{l}{Instrument} &\\multicolumn{1}{l}{{$R$=$\\lambda/\\delta\\lambda$}} &\\multicolumn{1}{l}{window [in \\AA]} &\\multicolumn{1}{l}{region} &\\multicolumn{1}{l}{N$_{\\rm spec}$} \\\\ \\hline \\\\ National Astron. Obs.  &NAO  &2.0m  &coud\\\"e    &15\\,000  &204& \\Ha &45\\\\ National Astron. Obs.  &NAO  &2.0m  &coud\\\"e    &30\\,000  &100& \\Ha &4\\\\ National Astron. Obs.  &NAO  &2.0m  &coud\\\"e    &15\\,000  &204& \\he &46\\\\ Tartu Observatory      &TO   &1.5m  &long-slit  &16\\,000  &127& \\Ha & 9\\\\ Ritter Observatory     &RO   &1.1m  &fiber-fed  &26\\,000  & 70& \\Ha & 7\\\\ \\hline \\end{tabular} \\end{table*} In addition to wind variability significant absorption $lpv$ was also established, which raised suggestions of a link stellar pulsations. To check the pulsational hypothesis however long-term photometric observations were required.  Motivated by the intriguing time-variable properties reported above, we organised and conducted new parallel spectroscopic and photometric monitoring campaigns of HD~199\\,478 during 1999 and 2000.  Details of the spectroscopic analysis are presented here and the results of the photometric survey were recently published by \\citep{PAM}. ",
        "conclusions": "\\label{disk}  Extensive monitoring campaigns of several late-B SGs, namely HD~199\\,478 (present study as well as \\citealt{MV}) and  HD~91\\,619, HD~34\\,085 and HD~96\\,919 \\citep{kaufer96a, kaufer96b, Kaufer97, Israel97}, indicate that their \\Ha profiles exhibit quite similar peculiarities, consisting of a double-peaked emission with V/R variation and occasional episodes of strong absorption indicating simultaneous mass infall and outflows. Such line signatures cannot be reproduced by conventional (i.e. non-rotating, spherically symmetric, smooth) wind models, which instead predict profiles in absorption partly filled in by emission for SGs at this temperature regime.  Discrepancies between observed and predicted \\Ha profiles have been also established for many O and $early$ B SGs, where this finding  was  usually interpreted as an indication for deviations from the adopted spherically symmetric, smooth wind approximations (e.g., \\citealt{Morel, Markova04, Markova05} and references therein).  Following this reasoning, axially symmetric envelopes, modulated, at least in the inner parts, by co-rotating weak magnetic structures have been assumed to explain the appearance and kinematical properties (e.g. double-peaked morphology with $V/R$ variations) of the peculiar \\Ha emission in the spectra of the four late-B SGs noted above \\citep{kaufer96a, MV}.  In addition, to account for the sudden appearance of HVAs in \\Ha and their development in time (e.g. the fast rise over a large velocity range, the lack of unshifted line emission, and the mooted re-appearance over a rotational time-scale), episodic and azimuthally extended, density enhancements in the form of co-rotating spirals rooted in the photosphere \\citep{kaufer96b} or closed magnetic loops similar to those in our Sun \\citep{Israel97} were also suggested.  It is generally mooted that non-radial pulsations (NRPs) and surface magnetic spots may equally be responsible for creating large-scale inhomogeneities in hot star winds \\citep{Full}. However, despite some progress (e.g., \\citealt{Kaufer06}) no convincing evidence of a direct relation between the time-scale of a given cyclical (wind) $lpv$ and the predicted time-scale of recurrent surface features due to a specific pulsation mode, has been derived to date (see \\citealt{townsend07} and references therein).  Note in the particular case of the four stars discussed here, non-radial pulsations due to $g$-modes oscillations were suggested to explain absorption $lpv$ in their spectra \\citep{Kaufer97, MP}. This possibility is partially supported by the present results, which indicate that on the HR diagram, and for parameters derived with FASTWIND, these stars fall exactly in the region occupied by known variable B SGs, for which $g$-modes instability was suggested. Also, the photometric variability of HD~199\\,478 seems to be consistent with a possible origin in terms of $g$-mode oscillations \\citep{PAM}. Thus, it seems very likely that the four late-B SGs in our sample are non-radial pulsators. Although no clear evidence of any causality between photospheric and wind (as traced by H$\\alpha$) variability has been seen so far for these objects (present study as well as \\citealt{Kaufer97}), one might speculate that their winds are perturbed due to pulsational instability, with specific signatures seen in the behaviour of H$\\alpha$.  An alternate possibility is that magnetic fields could be responsible for the appearance of large-scale  structures and wind asymmetries in hot stars. In particular, magneto-hydrodynamical (MHD) simulations for stars with moderately strong rotation, and for stellar and wind parameters typical for O and early B SGs (plus a magnetic dipole aligned to the stellar rotation) showed that depending on the magnetic spin-up, an equatorial compression, dominated by radial $infall$ and/or $outflows$ can be created, with no apparent tendency to form a steady, Keplerian disk \\citep{OD03, DOT08}.  Indeed, due to the lack of strong convection zones associated with hydrogen recombination, normal (i.e. without any chemical peculiarities) hot stars are generally not thought to be magnetically active. However, theoretical considerations (e.g. \\citealt{CM}) supported by more recent observations \\citep{Bychkov, Hurbig07} indicate that this may not necessarily be true and that relatively strong, stable, large-scale dipole magnetic fields are present in different groups of B stars (e.g. SPB, Be,  $\\beta$ Cep  itself etc.)  Thus, it seems likely that in at least some hot stars magnetic fields can be an alternative source of wind perturbations and asymmetries. The potential role of magnetic fields in B SGs remains intriguing, especially because it might provide a clue to understand the puzzling problem of the simultaneous presence of red- and blue-shifted absorption in \\Ha profiles of the four late-B SGs discussed here. \\begin{table} \\caption{Magnetic field strength, $B$ (in G), required to get an equatorial confinement with simultaneous mass infall and outflows around each of our targets. The Keplerian, $R_{\\rm K}$, the Alfven, $R_{\\rm A}$, and the escape, $R_{\\rm E}$, radii (in units of \\Rstar above the photosphere) are calculated following \\citet{OD03}. The HD~34\\,085 estimates correspond to the 2nd entry in Table~\\ref{para}.} \\label{magn_field} \\tabcolsep1.5mm \\begin{tabular}{lrrrr} \\hline \\hline \\multicolumn{1}{l}{Parameters} &\\multicolumn{1}{r}{91\\,619} &\\multicolumn{1}{r}{199\\,478} &\\multicolumn{1}{r}{34\\,085} &\\multicolumn{1}{r}{96\\,919} \\\\ \\hline \\\\ $R_{\\rm K}$                          &0.37        &0.45          &0.38       &0.42 \\\\ $R_{\\rm A}$                          &0.83..3.73  &0.77..4.59    &0.83..3.79 &1.0..4.13 \\\\ $R_{\\rm E}$                          &3.77        &4.58          &3.81       &4.18 \\\\ $B$ ($R_{\\rm K}<R_{\\rm A}<R_{\\rm E}$)&5..100      &5..180        &5..105     &5..85\\\\ \\hline \\end{tabular} \\end{table}   Guided by these perspectives, we employed the scaling relations given in \\citet{OD03} and calculated the Alfven, $R_{\\rm A}$, the Keplerian, $R_{\\rm K}$, and the ``escape\", $R_{\\rm E}$, radii of our targets, using data from Table~\\ref{para} and fixing the magnetic field strength at the values required to create an equatorial confinement. Interestingly, the results listed in Table~\\ref{magn_field}, show that in all four cases a very weak dipole magnetic field can effectively channel the wind outflows, leading to the formation of an equatorial compression with simultaneous radial mass infall and outflow.  With this in mind, new MHD simulations for the case of mid/late-B SGs have been recently initiated. The preliminary results (private communication, Asif ud-Doula) indicate that a pure dipole magnetic field of only a few tens of Gauss is indeed required to obtain a $cool$ equatorial compression (with mass infall and outflow) around a rotating star with stellar and wind parameters as derived with FASTWIND for HD~199\\,478.  (More detailed information about the outcomes of this study will be provided in a forthcoming paper.)  An obvious advantage of the model described above is that it has the potential to at least qualitatively account for some of the puzzling properties of the \\Ha line of our targets. In particular, the sudden appearance of red and blue-shifted absorptions might be explained if one assumes, that due to some reason the plasma in the infalling or outflowing zones of the compression or in both of them (during the HVA episodes) can become optically thick  in the $Lyman$ continuum and L$_{\\alpha}$, thus forcing \\Ha to behave as a resonance line, i.e. to absorb and emit line photons. The kinematic properties of the resulting absorption features \\footnote{(Depending on the size of the \\Ha forming region emission may not appear in the spectrum.)} are difficult to predict from simple qualitative considerations but it is clear that these properties cannot be dominated by stellar rotation, but instead will be controlled by the physical conditions inside the compression.  Concerning the interpretation of the peculiar \\Ha emission, the situation is more complicated since such emission can originate from different parts of the envelope under quite different physical conditions. For example, one can expect that the cool, less dense plasma outside the compression will only emit line photons (via recombinations), producing pure emission feature(s) in H$\\alpha$ .  Also, the cool equatorial compression might contribute to the \\Ha emission, providing the plasma inside the compression can occasionally become optically thin in this line. However, note that even a plasma that is optically thick in L$\\alpha$ and Lyman continuum can, under specific conditions, produce $pure$ emission profiles in \\Ha (e.g. $if$ collisions dominate the \\Ha formation or $if$ due to some reasons the 2nd and 3rd levels of Hydrogen go into LTE (\\citealt{PP96})).  Therefore, very weak dipole magnetic fields might be responsible for creating wind structures in the envelopes of late-B SGs. The models derived via MHD simulations seem to have the potential to account, at least qualitatively, for some of the peculiar characteristics of \\Ha in the spectra of our targets. However, more detailed quantitative analysis is required to investigate this possibility further.  New high signal-to-noise observations to prove/disprove the presence of weak magnetic fields can help to clarify the picture. Of course, due to the low strength of the magnetic fields required one cannot expect to detect these fields directly but indirect evidence such as e.g., the detection of X-ray emission, abundance anomalies, specific periodic variations in UV resonance lines, interferometric observations (for more information see \\citealt{henrichs01}), might also be considered. Note that a weak longitudinal magnetic field of about 130$\\pm$20~G\\footnote{Given that only one spectral line has been used for these measurements an error of 20 Gauss seems somewhat unrealistic and likely represents  only a lower limit.},  might have been detected in HD~34\\,085 ($\\beta$~Ori) \\citep{severny}, but confirmation is lacking.  Finally, at the cooler edge of the B-star temperature regime pure emission profiles in \\Ha can be accounted for if one assumes the winds are clumped. Indeed, a spherically symmetric, clumped wind will mimic wind densities higher than the actual ones, thus giving rise to strong line emission, similar to that in O stars. Such winds may also give rise to wind absorption, providing some of the clumps are optically thick in H$\\alpha$. Detailed numerical simulations and line formation calculations are required to discriminate  between the different possibilities."
    },
    "0806/0806.1434_arXiv.txt": {
        "abstract": "We present the optical to near-infrared (IR) spectrum of the galaxy TSPS J1329$-$0957, a red and bright member of the class of extremely red objects (EROs) at $z$ = 1.26. This galaxy was found in the course of the Tokyo-Stromlo Photometry Survey (TSPS) which we are conducting in the southern sky. The spectroscopic observations were carried out with the Gemini Multi-Object Spectrograph (GMOS) and the Gemini Near Infra-Red Spectrograph (GNIRS) mounted on the Gemini-South telescope. The wide wavelength coverage of 0.6 -- 2.3 $\\mu$m provides useful clues as to the nature of EROs while most published spectra are limited to a narrower spectral range which is dictated by the need for efficient redshift determination in a large survey.  We compare our spectrum with several optical composite spectra obtained in recent large surveys, and with stellar population synthesis models. The effectiveness of using near-IR broad-band data, instead of the spectral data, in deriving the galaxy properties are also investigated. We find that TSPS J1329$-$0957 formed when the universe was 2 -- 3 Gyr old, and subsequently evolved passively to become one of the most massive galaxies found in the $z$ = 1 -- 2 universe. Its early type and estimated stellar mass of $M_{*} = 10^{11.5} M_{\\odot}$ clearly point to this galaxy being a direct ancestor of the brightest elliptical and spheroidal galaxies in the local universe. ",
        "introduction": "}  The formation and evolution of massive (stellar mass $M_* > 10^{11} M_{\\odot}$) galaxies represent one of the key issues in understanding the whole picture of mass assembly in the universe. This class of objects is believed to correspond to the local early-type galaxies, i.e., ellipticals and spheroidals, constituting the brightest end of the luminosity function. They are predicted to form gradually with time through merging of smaller subsets in the hierarchical structure formation scenario based on the $\\Lambda$ cold dark matter (CDM) model \\citep[e.g.,][]{white78, cole94}. While the $\\Lambda$CDM model can successfully reproduce the large scale structure of galaxies, difficulties appear when it deals with  smaller systems such as individual galaxies. Although the hierarchical scenario predicts the time of appearance of the most massive galaxies at $z \\sim 1$, there is growing evidence that fully assembled galaxies exist at $z >1$ \\citep[e.g.,][]{dunlop96, spinrad97, daddi04, fontana04, glazebrook04, cimatti04, saracco05}. They are largely passively evolving, early-type galaxies containing old stellar populations \\citep[e.g.,][]{yan03, yan04, mccarthy04, mignoli05, longhetti05}. The alternative scenario of mass assembly is represented by the monolithic model, in which galaxies are formed in a single collapse and star burst, and subsequently evolve passively \\citep[e.g.,][]{larson74, tinsley77, bruzual80}. In this framework even the most massive galaxies could exist in very early universe. While both the hierarchical and monolithic scenarios are consistent with the observations in the local universe, their differences should become evident at $z > 1$. Recently \\citet{conselice07} showed that the observed mass growth of massive galaxies from $z = 2$ to $z = 0$ is significantly different from the prediction of the Millennium Simulation \\citep{lemson06, delucia06} based on the $\\Lambda$CDM model.   Thus motivated by their unique importance, the massive galaxies beyond redshift $z = 1$ have been a subject of extensive studies over the last decade. Several large survey programs have been completed or are ongoing, such as the Gemini Deep Deep Survey \\citep[GDDS;][]{abraham04}, the K20 Survey \\citep{cimatti02}, and others \\citep[e.g.,][]{doherty05, saracco05, roche06}. The $z > 1$ massive galaxies exhibit very red optical to near-IR colors, either due to evolved stellar populations, or severe reddening in case of dusty starbursts, and they form a class of extremely red objects \\citep[EROs;][]{elston88}. This ERO nature is commonly utilized to pick up the objects from multi-color photometry catalogs for subsequent spectroscopy, usually carried out in the red-visible (0.6 -- 0.9 $\\mu$m) light, to obtain the objects' redshifts \\citep[e.g.,][]{mccarthy04b}.   While the recent studies have been achieving great successes in revealing the mass assembly history out to the redshift $z >$ 2, and placing significant constraints on the galaxy formation scenarios, one should be aware of the fact that most of these works are based on optical spectroscopic observations supplemented with near-IR broad-band photometry, usually fitted with the spectral energy distribution (SED) models constructed from the local galaxy populations. It is based on the implicit assumption that the optical to IR SED, and its evolution with age, of the $z > 1$ galaxies are well represented by those of local counterparts. However, such an assumption has not yet been observationally confirmed due to the lack of the data with sufficiently large (i.e., from optical to near- or mid-IR) spectral coverage. The same shortage of the data has prevented the empirical-template method to be used in the analysis of the observed (spectral or broad-band) SEDs in wide wavelength ranges. Although recent works based on the large surveys have constructed some composite spectra by stacking all the obtained data, such composites have an apparent disadvantage in that different parts of the spectra are represented by systematically different galaxy populations, e.g., blue part by the high-redshift objects and red part by the low-redshift objects. Thus, the accumulation of the observations and analysis of wide spectral regions for individual galaxies are definitely needed at this stage in order to make comprehensive studies of the massive galaxy population existing at $z > 1$.   In this paper we report the first example of that kind of data, a red-visible to near-IR (0.6 -- 2.3 $\\mu$m) spectrum of the massive evolved galaxy TSPS J1329$-$0957 at $z$ = 1.26. It has been discovered in the course of the Tokyo-Stromlo Photometry Survey (TSPS) which we are conducting in the high galactic-latitude sky observed from the southern hemisphere. Introduction of the TSPS, target selection, and the observations and reductions of the optical and near-IR spectra are detailed in \\S \\ref{sec:obs}. We present in \\S \\ref{sec:char} the reduced spectrum which is compared to optical composite spectra generated in the past surveys and to stellar population synthesis models. Our comparisons reveal the basic properties of TSPS J1329$-$0957 as a representative of  passively-evolving galaxies found in the $z$ = 1 -- 2 universe. The effectiveness of using near-IR broad-band data, instead of the spectral data, in deriving the galaxy properties are also investigated. Finally a brief summary appears in \\S \\ref{sec:summary}. Throughout this paper we use the Vega magnitude system and a standard cosmology of $H_{\\rm 0}$ = 70 km s$^{-1}$ Mpc$^{-1}$, $\\Omega_{\\rm M}$ = 0.3, and $\\Omega_{\\Lambda}$ = 0.7. ",
        "conclusions": "}  \\subsection{Reduced Spectrum}  The reduced GMOS and GNIRS spectrum is shown in Figures \\ref{full_spec} and \\ref{composites} ({\\it black}). The \\ion{Ca}{2} HK absorption lines are clearly detected, giving the redshift of this object as $z$ = 1.26. There are the absorption features at the expected positions of \\ion{Mg}{2} $\\lambda$2800 and \\ion{Mg}{1} $\\lambda$2852, though the data quality is rather poor around these features since they are near the detector edge. No other single absorption or emission lines can be identified while some broad features are in common with the composite spectra of past surveys and/or with the model predictions as shown below. The detection of \\ion{Ca}{2} HK lines and absence of any strong emission features such as [\\ion{O}{2}] $\\lambda$3727, as well as the overall SED, clearly indicate that TSPS J1329$-$0957 is an old, passively-evolving galaxy whose light is dominated by evolved stellar population. Considering its redshift and $K_s$-band brightness ($K_s$ = 17.6 mag), this ERO is likely to be among the most massive galaxies with a stellar mass M$_{*}$ exceeding 10$^{11.5}$ M$_{\\odot}$. Hereafter we reveal some basic properties of the galaxy through comparisons with the optical composite spectra obtained in recent large surveys and with the stellar population synthesis (SPS) models.   \\subsection{Comparison with the Optical Composite Spectra\\label{sec:optcomp}}  In Figure \\ref{composites} we compare the optical spectrum of TSPS J1329$-$0957 with the composite spectra generated from the K20 early-type galaxies \\citep[{\\it red};][]{mignoli05}, the Sloan Digital Sky Survey (SDSS) Luminous Red Galaxies (LRGs) at $z$ = 0.30 -- 0.35 \\citep[{\\it green};][]{eisenstein03}, and the local early-type galaxies ({\\it blue}) compiled by \\citet{mannucci01}. Each spectrum has been scaled in accordance with the flux at the wavelength interval 0.34 -- 0.36 $\\mu$m and is expressed in the rest frame. The overall feature of the spectra blueward of $\\sim$4000 \\AA\\ are remarkably similar in spite of the systematic difference in the redshifts of galaxies constituting the composites, i.e., mean redshift $<z>$ = 0.75 for the K20 galaxies, $z$ = 0.30 -- 0.35 for the SDSS LRGs, and $z$ = 0 for the local early types, while $z$ = 1.26 for TSPS J1329$-$0957. It clearly demonstrates the early-type nature of TSPS J1329$-$0957, containing little amount of dust responsible for the reddening of the visible light.  On the other hand, a small but clear difference is seen in the 4000 \\AA\\ break. We show the D4000 index, adopting the definition of \\citet{bruzual83}, of various composite spectra and TSPS J1329$-$0957 in Figure \\ref{d4000}. The plotted composites are those described above plus those generated from the GDDS early-type galaxies at $z$ = 1.3 -- 1.4, the Las Campanas Infrared Survey (LCIRS) pure early-type galaxies \\citep{doherty05}, and the Munich Near-IR Cluster Survey (MUNICS) ``young'' early-type galaxy population defined by \\citet{longhetti05}. Also shown, are the predictions of the SPS models by the \\citet{bc03} code, assuming an  exponentially-declining star formation history with the e-folding time $\\tau$ = 1.0 Gyr, the metallicity $Z$ = $Z_{\\odot}$, 0.4 $Z_{\\odot}$ and 0.2 $Z_{\\odot}$, and the epoch of the onset of star formation at redshift $z_{\\rm f}$ = 7. The figure clearly shows the increase of the observed D4000 index from $z$ = 1.5 toward the local universe, which coincides with the simple galaxy aging represented by the SPS models. The index of TSPS J1329$-$0957 is also on this trend and is consistent with those of the GDDS galaxies and the MUNICS young subsamples within the measurement error. We infer the age of TSPS J1329$-$0957 to be a few Gyr.    \\subsection{Comparison with the SPS Models\\label{sec:sedfit}}  \\subsubsection{Basic Properties of TSPS J1329$-$0957}  Here we compare the overall SED of TSPS J1329$-$0957 with the predictions of  SPS models by the \\citet{bc03} code. In the model calculations, star formation is assumed to start at $t$ Gyr before the galaxy is observed and subsequently reduce its rate exponentially with the e-folding time $\\tau$ Gyr. Dust content is parametrized in terms of the color excess $E_{\\rm B-V}$ with the extinction law of \\citet{calzetti00}. These parameters are varied over the plausible ranges, i.e., $t$ = 0.01 -- 10.0 Gyr, $\\tau$ = 0.01 -- 10.0 Gyr, and $E_{\\rm B-V}$ = 0.0 -- 1.0 mag, and the best-fitting model is searched for by the least ${\\chi}^2$ method. The total stellar mass of the galaxy is determined through normalization of the model SED. We start with the \\citet{salpeter55} initial mass function (IMF) and the solar metallicity ($Z$ = $Z_{\\odot}$). Although other sets of calculations adopting the \\citet{fitzpatrick99} dust extinction raw or the \\citet{chabrier03} IMF were also performed, they provide neither better fits nor significantly different best-fitting results.   We show the best-fitting model with the observed spectrum in Figure \\ref{full_spec} ({\\it green}). The (observed-frame) optical and near-IR parts of the spectrum make similar contributions to the ${\\chi}^2$ where it is minimum. As seen in the figure, the overall agreement between the observation and the model is fairly good over the whole wavelength range. The parameter values comprising the best-fitting model are summarized in the first row of Table \\ref{sps_params}. They are: $t$ = 2.4$^{+0.2}_{-0.5}$ Gyr, $\\tau$ = 0.55$^{+0.10}_{-0.15}$ Gyr, $E_{\\rm B-V}$ = 0.06$^{+0.07}_{-0.03}$ mag, and log $M_{*}$ [$M_{\\odot}$] = $11.49^{+0.01}_{-0.04}$ where the quoted errors denote 90 \\% confidence. The derived age (2 -- 3 Gyr) and dust content ($E_{\\rm B-V} \\la$ 0.1 mag) are quite consistent with the estimates from D4000 index and from the comparison with the composite spectra of the early-type galaxies as described in the previous section. We have also performed the calculations assuming the sub-/super-solar metallicity ($Z$ = 0.4 $Z_{\\odot}$ and 2.5 $Z_{\\odot}$) and found that the best-fit age of the galaxy is $\\sim$1.5-Gyr older/younger than the original ($Z$ = $Z_{\\odot}$) case, which is the consequence of the well-known age-metallicity degeneracy. However, no better fits than the $Z$ = $Z_{\\odot}$ case are achieved in these calculations.   Thus we conclude that TSPS J1329$-$0957 is very likely to be formed 2 -- 3 Gyr before it was observed, and subsequently evolved passively to become one of the most massive objects found in the $z$ = 1 - 2 universe. Its formation epoch corresponds to the redshift $z_{\\rm f}$ = 2.6 when the universe is about 2.5 Gyr old. It is similar to the median formation redshift of the GDDS evolved galaxies studied by \\citet{mccarthy04} ($<z_{\\rm f}>$ = 2.4) and of the MUNICS young subsamples studied by \\citet{longhetti05} ($<z_{\\rm f}>$ = 2.3). The early-type nature and the estimated stellar mass of $M_{*} = 10^{11.5} M_{\\odot}$ clearly indicate that the galaxy is a direct ancestor of the brightest elliptical and spheroidal galaxies present in the local universe.    \\subsubsection{Effectiveness of Near-IR Broad-Band Data}  Most of the previous studies have utilized optical spectra supplemented with near-IR broad-band photometry in deriving the physical properties of EROs by the SED-model fitting. Thus it is useful to compare the fitting results for such data sets with the above results for the whole (optical and near-IR) spectral data. We first consider the optical spectrum with the $K$ ($K_s$)-band photometry, which is the most common data set used in the ERO studies. The $K_s$-band magnitude was calculated from the GNIRS spectrum and the SIRIUS filter transmission curve. We fitted the SPS-model SEDs to this data set in the same way as described above, and deduced the best-fitting parameters. The results are summarized in the second row of Table \\ref{sps_params}; the data set yields smaller $t$,  $\\tau$ and $E_{\\rm B-V}$ than the case of the whole spectral data. The shorter duration of star formation makes the galaxy SED (at a fixed age) redder while the younger age and less dust content make it bluer, thus the overall SED is similar to the case of the whole spectral data. Detailed features in the near-IR spectrum give rise to the difference in the best-fitting parameters for the two data sets. Adding the $J$- and $H$-band photometry to the above data set does not improve the situation (the last row of Table \\ref{sps_params}); the derived $t$, $\\tau$ and $E_{\\rm B-V}$ are tend to be even smaller than the previous case. These three parameters are clearly degenerated and the different data sets find the best SED models on the different locations of this ($t$, $\\tau$, $E_{\\rm B-V}$) degeneracy sequence. On the other hand, the total stellar mass $M_{*}$ is found to be fairly robust. It is consistent with the results of the previous studies \\citep[e.g.,][]{conselice07} reporting the robustness of the stellar-mass estimate as opposed to other parameters (galaxy age, star-formation duration, metallicity, etc.) suffering from various degeneracies.   We conclude that the optical spectrum supplemented with the near-IR broad-band data provide slightly different view of the target galaxy from those deduced from the whole spectral data. However, the basic classification of the galaxy could be achieved without the near-IR spectrum; the three data sets listed in Table \\ref{sps_params} consistently suggest that the galaxy is a few Gyr old system with faded star-formation activity and with little dust. The stellar-mass estimate is fairly robust. We note that the errors of the $JHK_s$ broad-band magnitudes used above have been calculated by appropriately propagating the error of the GNIRS spectrum, which gives 0.05 mag for the $J$ band and 0.03 mag for the $H$ and $K_s$ bands. These are similar to the typical magnitude errors achieved in recent photometric observations. Thus the above data sets of the optical spectrum with the near-IR broad-band photometry well represent those obtained in real observing programs.          }  We present a Gemini/GMOS and GNIRS spectrum of the galaxy TSPS J1329$-$0957, a red and bright member of the ERO population. It has been discovered in the course of the Tokyo-Stromlo Photometry Survey (TSPS) we are developing in the southern sky. The spectrum covers red-visible to near-IR (0.6 -- 2.3 $\\mu$m) wavelength range, providing a useful example to be compared with the same population which has been vigorously discovered in the recent surveys using the 8-m class telescopes and multi-object spectrographs. The \\ion{Ca}{2} HK absorption lines are clearly detected, which gives the redshift of the galaxy as $z$ = 1.26. The detection of \\ion{Ca}{2} HK lines, the absence of any strong emission features, and the overall SED indicate that the galaxy is an old, passively-evolving system whose light is dominated by an evolved stellar population. We compared the observed spectrum with the predictions of the SPS models and found that the model with age 2 --3 Gyr, little amount of dust ($E_{\\rm B-V} \\la$ 0.1 mag), and  stellar mass $M_{*}$ = $10^{11.5} M_{\\odot}$ to best-fit the observations. The inferred properties of the galaxy are supported by the estimate from the  D4000 index and by the resemblance of the observed spectrum to several composite spectra of  early-type galaxies at $z <$ 1. We also investigated the effectiveness of using near-IR broad-band data, instead of the spectral data, in deriving the galaxy properties. It shows that the optical spectrum supplemented with near-IR broad-band data provide slightly different view of the target galaxy from those deduced from the whole spectral data. However, the basic classification of the galaxy could be achieved without the near-IR spectrum and the stellar-mass estimate is fairly robust.   We conclude that TSPS J1329$-$0957 is very likely to have formed 2 -- 3 Gyr before it was observed, and subsequently evolved passively to become one of the most massive objects, of early-type, in the $z$ = 1 - 2 universe. The presence of such a massive galaxy in the early universe might argue against the current representations of the hierarchical formation scenario and might favor the monolithic scenario."
    },
    "0806/0806.3607_arXiv.txt": {
        "abstract": "We present a comparative study of the observed properties of the optical and X-ray afterglows of short- and long-duration $\\gamma$-ray bursts (GRBs).  Using a large sample of 37 short and 421 long GRBs, we find a strong correlation between the afterglow brightness  measured after 11 hours and the observed fluence of the prompt emission.   Both the  optical (R band) and X-ray flux densites ($F_R$ and $F_X$) scale with the  $\\gamma$-ray fluence, $F_{\\gamma}$.    For bursts with a known redshift, a tight correlation exists between the afterglow flux densities at 11 hours (rest-frame) and the total isotropic $\\gamma$-ray energy, $\\Egi$: $F_{R,X} \\propto {\\Egi}^{\\alpha}$, with $\\alpha \\simeq 1$. The constant of proportionality is nearly identical for long and short bursts, when $\\Egi$ is obtained from the Swift data. Additionally, we find that for short busts with $F_{\\gamma} \\ga10^{-7}$ erg cm$^{-2}$, optical afterglows are nearly always detected by reasonably deep early observations.   Finally, we show that the ratio $F_R / F_X$ has very similar values for short and long bursts. These results are difficult to explain in the framework of the standard scenario, since they require that either (1) the number density of the surrounding medium of short bursts is typically comparable to, or even larger than the number density of long bursts;  (2) short bursts explode into a density profile, $n(r) \\propto r^{-2}$ or (3) the prompt $\\gamma$-ray fluence depends on the density of the external medium. We therefore find it likely that either basic assumptions on the properties of the circumburst environment of short GRBs or else the standard models of GRB emission must be re-examined.    We believe that the most likely solution is that the ambient density surrounding typical short bursts is higher than has generally been expected:  a typical value of $\\sim 1$ per cm$^{-3}$ is indicated. We discuss recent modifications to the standard binary merger model for short bursts which may be able to explain the implied density. ",
        "introduction": "Although the past decade has seen great progress in characterizing the afterglows of long-duration GRBs (LGRBs), the afterglows of short-duration GRBs (SGRBs) remain elusive.  The average fluence of an SGRB is an order of magnitude fainter than that of an LGRB, and thus while SGRBs have been localized by {\\it Swift}, about ten long bursts are localized for every short burst.  Optical afterglows of SGRBs have been detected, but are notoriously dim, and large-aperture telescopes are generally required to detect them.  Nearly half of all SGRBs have a fluence less than 10$^{-7}$ ergs cm$^{-2}$ and none of these have a detected optical afterglow.  SGRB X-ray afterglows are commonly detected, but are faint and less likely to be rapidly localized with their positions distributed in real-time, further complicating optical ground-based observations.  As a result, only a few well-sampled, robust optical SGRB light curves exist that span a wide range of time or frequency.  The faintness of the afterglows of short GRBs (compared to the long GRBs) has been attributed to properties of the progenitor and circumburst environment.  While the LGRB-SN connection is strong, the progenitors of SGRBs are less certain.  Observations have shown that long-duration GRBs are associated with the deaths of massive stars (\\citealt{smg+03, hsm+03}; see \\citealt{bw06} for a comprehensive review).  However, a number of models exist for the progenitors of SGRBs.  The leading model is a binary pair of compact objects, dual neutron stars, or a neutron star and black hole, which merge in a dramatic explosion causing a burst of $\\gamma$-rays \\citep{elp+89, npp92, lrg05} powered by accretion onto the newly formed compact object.  It is commonly believed that the difference in the long and short progenitors implies that they reside in different environments (for a recent review, see \\citealt{n07}).  Because LGRBs are associated with massive stars, they will lie in or near dense regions of active star-formation.  SGRBs progenitors can have long lifetimes and do not need to be in galaxies with active star-formation, nor be associated with dense regions.  In fact, the observed diverse host galaxy morphology is roughly consistent with the expected wide range of progenitor lifetimes \\citep{bpb+06}.  Two of the first detected SGRB afterglows, those of GRB 050509B \\citep{gso+05} and GRB 050724 \\citep{bpc+05}, were found near bright, giant elliptical galaxies.  GRB 050709 \\citep{hwf+05} and GRB 051221A \\citep{sbk+06} both occurred in late-type galaxies with ongoing star-formation.  However, within these late-type galaxies, SGRBs do not necessarily trace areas of active star-formation as is the case with LGRBs \\citep{bkd02, fls+06}, rather, they can occur in low luminosity outer regions of the galaxy (e.g. GRB 050709; \\citealt{ffp+05}, GRB 060121; \\citealt{ltf+06}).  Some theories of the formation of compact object binaries include a violent supernova-induced {\\it kick} that can offset them from their host galaxy \\citep{bsr+06} and may place them in these low density environments.  The low environmental densities into which SGRBs occur have an observational implication.  Within the context of the standard afterglow model \\citep[e.g.][]{MR97, spn98} for observing frequencies below the cooling frequency (and above the peak frequency), the afterglow flux is proportional to the density of the external medium to the power $1/2$.  Well-sampled, multiwavelength observations of LGRB afterglows have suggested that at half a day, the cooling frequency generally lies between the optical and X-ray frequencies \\citep[e.g.][]{gwb+98,hbf+99,pk02,sbk+06}.  It is therefore expected that the differences in the density of the circumburst medium between the long and short bursts would affect the late time optical emission, which is below the cooling break.  (Note that according to the standard model, the flux above the cooling break is independent of the external medium density, $n$. Thus, the flux at the X-ray frequency may not be affected by the difference in densities. See discussion in \\S\\ref{sec:analysis} below).  In addition to depending on the environment, we also expect the absolute afterglow brightnesses to be proportional to the total energy released in the burst.  The intrinsic energy release of short burst progenitors is not well-known because the current redshift distribution of SGRBs is fragmentary.  The average redshift of the SGRB distribution is $<z>$ = 0.5 with E$_{ISO} \\sim 10^{48} - 10^{51}$, compared to $<z>$ = 1.9 and   E$_{ISO} \\sim 10^{49} - 10^{53.5}$ for LGRBs.  However, it is unclear to what extend these differences are a result of observational biases and/or true intrinsic differences in luminosity or collimation in comparison with LGRBs.  No SGRB absorption redshifts have been measured, which is likely due to the lower afterglow brightness and possibly a less dense external medium on which to imprint strong absorption lines.  Therefore, because all SGRB redshifts have been measured from optical emission lines in host galaxies, the current sample of redshifts is limited to $z \\la 1$.  \\citet{bfp+07} have noted that the lack of bright galaxies in the error circles of a number of XRT localizations may suggest that a number of short bursts without optical identifications may lie at $z > 1$.  Finally we note that the E$_{ISO}$ values presented here are either based upon the $\\gamma$-ray fluence observed by the Swift Burst Alert Telescope (BAT) or for the smaller non-Swift sample, scaled to the BAT sensitivity. Because the BAT is only sensitive up to 150 keV, these values underestimate the total prompt emission of all bursts.   However, this is most severe on short bursts, which tend to have a harder spectrum than long bursts.  We return to this point in more detail later.  In this paper we compare the optical and X-ray afterglows measured after 11 hours (a standard {\\it Swift} timescale) with the prompt energy emission of both long- and short-duration $\\gamma$-ray bursts, and compare these results with afterglow theory.  We test the predictions of the standard model with a large data set that spans over six decades in energy. We find that SGRBs afterglows are very similar to low-luminosity LGRBs afterglows.  Surprisingly, we find that the afterglow brightnesses depend little upon the classification of long or short; rather, the bursts scale primarily with the high-energy prompt emission.  Our result is difficult to reconcile with standard afterglow theory, and contradicts previous claims that SGRB optical afterglows are dim due to the density of their surrounding medium \\citep{pkn01, n07, b07d}.  This paper is organized as follows: In \\S \\ref{sec:data} we detail the steps taken to gather the data presented herein.  We compare the prompt and afterglow emission of the two populations in \\S 3.  In \\S 4, we explore the assumptions of the standard model, under what conditions the results are valid, and implications related to host galaxies and intrinsic SGRB energies.  We show that the afterglow data implies that either (1) the environment of short GRBs is comparable to that of long GRBs, (2) SGRBs explode into a wind profile, or (3) the prompt emission fluence depends on the environment. In \\S\\ref{sec:statistics} we discuss the statistical significance of our data sample.  We summarize and draw final conclusions in \\S \\ref{sec:summary}. ",
        "conclusions": "\\label{sec:summary}  In this work, we present a comprehensive study of the optical and X-ray afterglows obtained after 11 hours of 37 short and 421 long GRBs, and compared the results to the energy emitted in $\\gamma$-rays during the prompt emission phase. This sample is the largest used so far for this type of study.  We find a strong correlation between the optical and X-ray afterglow brightnesses and prompt $\\gamma$-ray fluence of GRBs, a result similar to one reported contemporaneously by by  \\citet{g+08} \\citet{Kann08}.  However, we also show that the ratio of the optical to X-ray emission in the short burst afterglows is comparable to that of the long-bursts.  This equivalence between burst types, and the absolute ratio we find can only be understood in the framework of the standard theory if the circumburst density of SGRBs is typically comparable to, or even higher than the circumburst density of LGRBs. We point out that in a wind-like scenario for SGRBs can also explain this result if the average mass loss rate of the SGRB progenitor is $\\sim 30$ times lower than the characteristic mass loss rate from a Wolf-Rayet star.  Finally, we have show that in the framework of the standard model, in which the prompt emission fluence is independent of the environmental density, by selecting bursts with $\\gamma$-ray fluence larger than $10^{-7} \\rm{\\, erg \\, cm^{-2}}$ we obtain for short bursts above this fluence an essentially complete sample for analysis.  These results indicate that SGRB afterglows are not necessarily dim due to a lower density in the circumburst environment.  While the average SGRB afterglow is significantly dimmer than the average LGRB afterglow, this appears  to be largely due to the lower total energy release of the burst, as indicated by the fainter prompt $\\gamma$-ray emission.  We can thus conclude that if the prompt emission is a result of internal shocks, then either the external density of long and short bursts is similar, or the afterglow flux is less sensitive to the density of the surrounding medium than predicted.  On the other hand, if the prompt emission is due to an external shock, both quantities the gamma-ray fluence and afterglow intensity will depend on the external density and may scale accordingly.  In this case we could attribute the faintness of {\\it both} the prompt and afterglow emission to the external medium density.   Due to the short lifetimes of the massive star progenitors of long-duration GRBs, they have been traditionally expected to occur in or near giant molecular clouds ($n \\sim 10^{2.5}$ cm$^{-3}$; \\citealt{sal87}).  Some observational evidence supports this hypothesis.  The hosts of LGRBs are typically blue and actively forming massive stars \\citep{ftm+99, sfc+01, ldm+03, chg04, gps+05}. \\citet{fls+06} find that LGRBs are preferentially located on the brightest regions of these galaxies, which they interpret as evidence that GRBs frequently occur on or near massive OB associations.  X-ray observations of GRB afterglows show neutral hydrogen column densities consistent those of Galactic giant molecular clouds \\citep{gw01, jfl+06, smp+07}.  At optical wavelengths, moderate extinction characteristic of dense regions is seen in some bursts \\citep{kkz+07}, while in others the afterglow appears significantly extinguished \\citep{dfk+01, lfr+06, wfl+06, pdl+06, rvw+07}, and at high redshifts where Ly$\\alpha$ absorption can be found in the optical spectrum, large hydrogen column densities (up to $10^{23}\\, {\\rm cm}^{-3}$) are frequently found \\citep[c.f.][]{jfl+07}.  Nonetheless, as noted earlier, the density of the medium into which the relativistic flow of the GRB expands may be dominated by an earlier wind from the progenitor star rather than the ISM.  In contrast, SGRBs are found in both elliptical and star-forming galaxies, a fact consistent with merger scenarios.  The ISM of elliptical galaxies is sparse and has a typical density of only $n \\sim 10^{-2.5}$ cm$^{-3}$ \\citep{fbp+06}.  \\citet{Pan06}, however, estimates a circumburst density of  $10^{-1} < n < 10^{3} \\, {\\rm cm}^{-3}$ for GRB 050724, a burst in an E/S0 galaxy \\citep{bpc+05}. At the same time circumburst densities of $10^{-4} < n < 10^{-1} \\, {\\rm cm}^{-3}$ \\cite{Pan06} and $10^{-3}$ to $10^{-1} \\, {\\rm cm}^{-3}$ \\cite{sbk+06} are found for GRBs 050709 and 051221A, respectively, both of which occurred in apparently star-forming galaxies. As we have shown here, the lower ranges of these two latter densities  are inconsistent with the distribution of ratios of $F_X/F_R$ seen in Figure \\ref{fig:ratio}.  These estimates of the SGRB circumburst density , however, assume a locally uniform medium.  If the SGRBs were to typically form in a windlike structure, this ratio would be acceptable.  Still, the common presence of a wind structure around SGRBs would rule out the standard merger models. The velocity by mass loss during the supernovae, let alone any kicks by the supernovae, would easily move the compact remnants by more than 100 pc from their natal region before a binary merger occurred \\citep{bpb+06}.  A potentially simple explanation of the results presented here is that both the afterglow and the prompt emission are due to an external interaction \\citep[e.g.,][]{Dermer08}. In this case, both the afterglow and prompt emission might be expected to scale similarly. However, it is difficult to explain the strong variability observed in many bursts in such a model \\citep[see, e.g.,][]{Mes06}. Furthermore, observational evidence for a correlation between the density of a circumburst medium and $\\Egi$ is so far either weak or potentially damaging.  \\cite{tko+08} have suggested that SGRBs far from the center of their host may be fainter than those closer in.  However, this work uses a number of associations based on the brightest galaxy in a rather large Swift XRT error circle, and therefore may in some cases assign an incorrect offset, and thus will need to be confirmed on larger, better samples.  In contrast, the only work to examine this issue for long bursts has found a positive correlation between the brightness of the burst and the distance of the burst from the center of the host \\citep{rlb02}.  This is, to first order, the opposite of one would predict if the external medium played a role in the gamma-ray emission.  It has, however, been noted that the bursts GRBG 050509b and 050724, both of which were in early type galaxies were intrinsically weaker than bursts found in star-forming hosts \\cite{Rhoads09}.  Nonetheless, we note again that the work by \\citet{Pan06} suggests that the circumburst medium around at GRB 0507024 was reasonably dense.  Several recent advances in the theory and modelling of the binary merger process and the local insterstellar medium could potentially explain our results.   \\citet{vdh07} has pointed out that the large majority of known double neutron-star binaries have low eccentricities, suggesting that the second born neutron stars in binaries, unlike single neutron stars, typically did not receive substantial ``kick'' velocities at birth. Lower kicks will lead to lower systemic velocities, and a higher probability that the system remains in a moderately dense ISM when the binary merges.   Furthermore, the time to merger may often be relatively short. \\citet{bsf09} estimate that up to 70\\% of NS-NS and NS-BH binaries may merge with 100 Myr after formation.   Indeed, one burst that appears to have likely been a short burst, GRB 060505, lies in an [HII] region on the outskirts of its spiral host \\citep{fwt+06,ocg+07,lk07}.   Some short bursts, however, occur in early-type galaxies.  Here recent result could also provide an explanation for an apparent ISM density of $1\\, {\\rm cm}^{-3}$.   Binary mergers in early-type galaxies are almost certainly dominated by the evolving binaries in globular clusters \\citep{gpm06}, and the most rapid evolution is likely to occur in the densest clusters.   These clusters however, could have ionized gas densities of the required magnitude due to the colliding winds of giants contained in the cluster \\citep{prl09}.  Although we may not fully understand the mechanism, we cannot escape the conclusion that SGRBs are fundamentally similar in emission properties to LGRBs -- particularly the fainter LGRBs that have typically been found at low redshift.  It has been suggested that low-luminosity LGRBs (e.g. GRB 980425, GRB 031203 and GRB 060218) have large, and perhaps nearly quasi-spherical, opening angles when compared to traditional LGRBs \\citep{skb+04, skn+06, lzv+07}. Similarly, of the three SGRBs with estimated opening angles, GRB 050709 and GRB 050724 have large measured angles, $\\Theta \\sim 14^{o}$ \\citep{ffp+05} and $\\sim 25^{o}$ \\citep{gbp+06, mcd+07} respectively, while GRB 051221 has $4^o < \\Theta < 8^o$ \\citep{bgc+06, sbk+06}, which is more typical of traditional high $\\Egi$ LGRB opening angles. Yet both long and short populations appear to largely fall on a simple log-linear relationship between afterglow intensity and gamma-ray luminosity over six orders of magnitude in energy.  Within this relationship there may be strong hints as to the nature of the progenitors of SGRBs, or perhaps a new more fundamental understanding of the GRB emission mechanism."
    },
    "0806/0806.3788_arXiv.txt": {
        "abstract": "To date, there is no core accretion simulation that can successfully account for the formation of Uranus or Neptune within the observed 2-3~Myr lifetimes of protoplanetary disks.  Since solid accretion rate is directly proportional to the available planetesimal surface density, one way to speed up planet formation is to take a full accounting of all the planetesimal-forming solids present in the solar nebula.  By combining a viscously evolving protostellar disk with a kinetic model of ice formation, which includes not just water but methane, ammonia, CO and 54 minor ices, we calculate the solid surface density of a possible giant planet-forming solar nebula as a function of heliocentric distance and time.  Our results can be used to provide the starting planetesimal surface density and evolving solar nebula conditions for core accretion simulations, or to predict the composition of planetesimals as a function of radius.  We find three effects that favor giant planet formation by the core accretion mechanism: (1) a decretion flow that brings mass from the inner solar nebula to the giant planet-forming region, (2) the fact that the ammonia and water ice lines should coincide, according to recent lab results from \\cite{collings04}, and (3) the presence of a substantial amount of methane ice in the trans-Saturnian region.  Our results show higher solid surface densities than assumed in the core accretion models of Pollack et al. (1996) by a factor of 3--4 throughout the trans-Saturnian region.  We also discuss the location of ice lines and their movement through the solar nebula, and provide new constraints on the possible initial disk configurations from gravitational stability arguments. ",
        "introduction": "\\label{introduction}  Previous core accretion simulations of the formation of Uranus and Neptune require the solar nebula lifetime to be at least 10~Myr longer than the observed median protostellar disk lifetime (2--3 Myr; Haisch et al. 2001).  Similarly, the simulation of Saturn's formation by Alibert et al. (2001) allows the planet to form in 2.5~Myr, but only if it migrates from 11.9~AU to 9.5~AU during formation. This scenario is incompatible with the outward migration of Saturn predicted by the Nice model of planetary dynamics (Tsiganis et al. 2005) This research was motivated by the need to create a model of the solar nebula with enough solid mass to form Jupiter, Saturn, Uranus and Neptune via core accretion at heliocentric distances of $5 < R < 20$~AU.  For core accretion simulations, planetesimal mass available for building giant planet cores is usually calculated by invoking the scaled minimum-mass solar nebula with surface density $\\Sigma \\propto R^{-3/2}$, obtained from spreading out the current mass of the planets (Weidenschilling 1977), and assuming a gas/solid ratio of $\\sim 70$. However, a single planet core embedded in a disk of planetesimals grows at the rate of \\begin{equation} {dM_{\\rm core} \\over dt} = \\pi a_{\\rm core}^2 \\Sigma_{\\rm solid} \\Omega F_g , \\label{planetgrowth} \\end{equation} where $F_g$ is a factor taking into account gravitational focusing by the core (Safronov 1969), $\\Sigma_{\\rm solid}$ is the solid surface density of planetesimals, $\\Omega$ is the orbital angular frequency of the core, and $a_{\\rm core}$ is the core's effective capture radius. For Keplerian orbits, we have \\begin{equation} \\Omega = \\left ( {GM_* \\over R^3} \\right )^{1/2} . \\label{kepler} \\end{equation} The $R^{-3/2}$ dependence of $\\Omega$ means that even planets forming in a disk with uniform surface density should, in the end, reach masses that suggest $\\Sigma \\propto R^{-3/2}$ if growth rate is not taken into account.  A near-uniform surface density distribution, in which mass is not concentrated in the inner disk, gives higher surface densities in the outer disk and may speed up the formation of Uranus and Neptune.  Furthermore, the canonical gas/solid ratio of 70, which is based on the composition of Comet Halley (Jessberger et al. 1989), includes only one ice in the solid inventory---H$_2$O.  However, there were certainly other ices present in the solar nebula: the Deep Impact ejecta from Tempel 1 contained $\\sim 1000$~tons of excavated CO, CO$_2$ and CH$_3$OH (A'Hearn 2008).  Taking a full accounting of the solids in the solar nebula may reveal other volatile mass sources besides H$_2$O that condense in the outer solar nebula and improve the formation prospects of Uranus and Neptune.  In this work, we calculate the time-evolving solid surface density available for giant planet formation.  We make two updates to previous solar nebula simulations: (1) we relax the assumption of a steady-state disk and allow for viscous redistribution of mass, which tends to drive the surface density profile toward uniformity, and (2) we add a chemical reaction network to trace the formation and freezeout of volatiles.  This paper is organized as follows.  In \\S \\ref{overview}, we outline our method for calculating solid surface density in the solar nebula: which codes we use, how they are linked, and how they differ from previous solar nebula models.  In \\S \\ref{alphadisk}, we describe our protostellar disk model and use it to calculate the total surface density (gas+solid) and temperature evolution of the disk.  \\S \\ref{chemmodel} details the physics behind our chemical reaction network, the abundances of the major ice species, and the location of the ice lines.  We calculate the solid surface density available for planet formation in \\S \\ref{planetesimal} and present our conclusions in \\S \\ref{conclusions}. ",
        "conclusions": "\\label{conclusions}  By combining a viscously evolving protostellar disk model with a kinetic model of ice formation, we have calculated the time-evolving solid surface density available for giant planet formation.  We find three results that are highly favorable for gas- and ice-giant formation:  \\begin{enumerate}  \\item The total (gas+solid) surface density distribution evolves toward uniformity, with an outward-moving wave depositing mass in the giant planet-forming region between 5 and 20~AU,  \\item The ammonia and water ice lines coincide, providing a mass enhancement of 7\\% at the snow line over pure water condensation, and  \\item Methane and CO condensation beyond 12~AU add mass to Uranus and potentially Neptune's feeding zones.  \\end{enumerate}  The next step is to use a core-accretion model to calculate how fast Saturn, Uranus and Neptune can form with our enhanced surface densities. The same methods used by \\cite{hubickyj05} should work for Saturn. However, an updated method of tracking planetesimal scattering may be required for feeding zones with $R > 15$~AU, where orbital velocities of planetesimals decrease to $\\sim 10$--50 times their escape velocities and self-stirring begins to affect the dynamics of the planetesimal disk.  By calculating solid surface density as a function of radius and time self-consistently with viscous evolution of the protostellar disk, instead of assuming a single gas/solid ratio covers the entire giant planet-forming region of the solar nebula, the research presented here allows core accretion simulations to be fully deterministic.  One area that requires more work is the viscous evolution of an MRI-turbulent disk with a dead zone, where $\\alpha$ varies with radius. Annuli with an inactive midplane would have a lower column-averaged value of $\\alpha$, which could slow or stop the movement of mass from the inner solar nebula to the giant planet-forming region.  However, Turner and Sano (2008) find that a toroidal magnetic field component produced in the midplane by shear drives a laminar accretion flow even through the nonturbulent dead zone.  Since our focus is on increasing the solid mass available for giant planet cores, we have neglected clathrate hydrates in our chemical model, preferring to tap into the large mass reservoir provided by ices. However, noble gas enrichment in the giant planet atmospheres suggests that clathration is very important, if not for building planet cores, then at least for determining atmospheric composition during the late stages of accretion (Hersant et al. 2004).  The presence of PAHs in carbonaceous meteorites, cometary dust and Iapetus implies that these, too, might be an important solid mass reservoir.  It is not yet known whether PAHs survive in the ISM or form primarily in the warm disk. Woods and Willacy (2007) present a mechanism for benzene formation in the inner 3~AU of the solar nebula.  We are investigating the possibility of including PAH chemistry in future work.  Finally, we note that the ALMA Design Reference Science Plan contains a proposal for detecting the snow line in nearby protoplanetary disks.  We submit that the CO ice line would be an equally valuable detection, since this work strongly suggests that CO ice should be part of the solid inventory of the Uranus- and Neptune-building planetesimals.  In a 1 Myr-old solar nebula analog in Taurus (140 pc), the CO ice line would lie 6 resolution units from the star when observing the CO ($J = 3 \\rightarrow 2$) transition.  We thank Julie Moses, Mark Marley and Doug Lin for helpful discussions about the project design.  S.D.R. was supported by grants from the NSF Graduate Research Fellowship program and the Achievement Rewards for College Scientists Foundation.  K.W. and N.T. were supported by the JPL Research and Technology Development Program.  P.B. received funding from NSF Grant AST-0507424 and NASA Origins Grant NNX08AH82G."
    },
    "0806/0806.0953_arXiv.txt": {
        "abstract": "{The Red MSX Source (RMS) survey is an ongoing multi-wavelength observational programme designed to return a large, high-resolution mid-infrared colour-selected sample of massive young stellar objects (MYSOs). We have identified $\\sim$2000 MYSO candidates located within our Galaxy by comparing the colours of MSX and 2MASS point sources to those of known MYSOs. The aim of our follow-up observations is to identify other objects with similar colours such as ultra compact (UC) HII regions, evolved stars and planetary nebulae (PNe) and distinguish between genuine MYSOs and nearby low-mass YSOs.} {A critical part of our follow-up programme is to conduct $^{13}$CO molecular line observations in order to determine kinematic distances to all of our MYSO candidates. These distances will be used in combination with far-IR and (sub)millimetre fluxes to determine bolometric luminosities which will allow us to identify and remove nearby low-mass YSOs. In addition these molecular line observations will help in identifying evolved stars which are weak CO emitters.} {We have used the 15~m James Clerk Maxwell Telescope (JCMT), the 13.7~m telescope of the Purple Mountain Observatory (PMO), the 20~m Onsala telescope and the 22~m Mopra telescope to conduct molecular line observations towards 508 MYSOs candidates located in the 1st and 2nd Quadrants. These observations have been made at the $J$=1--0 (Mopra, Onsala and PMO) and $J$=2--1 (JCMT) rotational transition frequency of $^{13}$CO molecules and have a spatial resolution of $\\sim$20\\arcsec--55\\arcsec, a sensitivity of $T_{\\rm{A}}^*$ $\\simeq$ 0.1~K and a velocity resolution of $\\sim$0.2~\\kms. We complement these targeted observations with $^{13}$CO spectra extracted from the Galactic Ring Survey (GRS), which have a velocity resolution of $\\sim$ 0.21 \\kms\\ and sensitivity $T_{\\rm{A}}^*$ $\\simeq$ 0.13-0.2~K, towards a further 403 RMS sources.} {In this paper we present the results and analysis of the $^{13}$CO spectra obtained towards 911 MYSO candidates. We detect $^{13}$CO emission towards 780 RMS sources which corresponds to approximately 84\\% of those observed. A total of 2595 emission components are detected above 3$\\sigma$ level (typically $T^*_{\\rm{A}} \\ge 0.3$~K), with  multiple components being observed towards the majority of these sources -- 520 sources ($\\sim$56\\%) -- with an average of $\\sim$4 molecular clouds detected along each line of sight. These multiple emission features make it difficult to assign a unique kinematic velocity to many of our sample. We have used archival CS ($J$=2--1) and maser velocities to resolve the component multiplicity towards 175 sources ($\\sim$20\\%) and have derived a criterion which is used to identify the  most likely component for a further 191 multiple component sources. Combined with the single component detections we have obtained unambiguous kinematic velocities for 638  of the 780 MYSOs candidates towards which CO is detected ($\\sim$80\\% of the detections). The 141 sources for which we have not been able to determine the kinematic velocity will require additional line data. Using the rotation curve of Brand and Blitz (1993) and their radial velocities we calculate kinematic distances for all detected components.} {} ",
        "introduction": "\\subsection{Background}  The formation and evolution of massive stars ($>$10$^4$~\\lsun) is still cloaked in mystery despite the fact that they are responsible for some of the most energetic astrophysical phenomena. Massive stars can have an enormous impact on their local environment through their UV radiation, molecular outflows and strong stellar winds as well as depositing huge amounts of processed material back to the ISM. Their impact is not limited to their local environment but can also influence the evolution of their whole galaxy. They are also thought to play an important role in regulating star formation through the propagation of strong shocks into their natal molecular cloud and nearby clouds. These shocks may be responsible for triggering subsequent generations of star formation or disrupting conditions necessary for star formation in nearby clouds (\\citealt{elmegreen1998}).  Our understanding of how massive stars form and evolve lags behind that of low-mass stars, which in broad terms is relatively well understood (\\citealt{shu1987,shu1993}). There are several reasons for this; firstly, massive stars are extremely rare and have very short evolutionary time scales. This results in them generally being located  much farther away than sites of low-mass star formation. Secondly, massive stars are known to form exclusively in clusters making it difficult to identify and attribute derived quantities to individual sources. Finally, they reach the main sequence whilst still accreting material, and are heavily embedded within their natal molecular cloud, hidden away beneath many magnitudes of visual extinction. However, given their profound impact, not only on their local environment, but also on a galactic scale, it is vital that we understand the environmental conditions and processes involved in their birth and the early stages of their evolution.  Due to the observational difficulties, until relatively recently, the only catalogue of well characterised massive young stellar objects (MYSOs) had been limited to 30 or so serendipitously detected sources (\\citealt{Henning1984}) which are mostly nearby. The situation has improved considerably in recent years with a number of studies using a variety of selection criteria (e.g., \\citealt{molinari1996,walsh1997,sridharan2002}). These studies have identified many new MYSOs, however, they are all based on IRAS colours which bias them towards bright, isolated sources and tend to avoid dense clustered environments and the Galactic mid-plane where the majority of MYSOs are expected to be found -- the scale height of massive stars is $\\sim$30\\arcmin\\  (\\citealt{reed2000}). The limited number of MYSOs identified by these studies, and the way they have been selected, means that they are probably not representative of the general MYSO population. This makes statistical studies of many aspects of massive star formation difficult.      \\subsection{The RMS Survey}  There is an obvious need for a large, high-resolution mid-infrared colour-selected sample with sufficient numbers of sources in each luminosity bin to allow statistical studies as a function of luminosity. The Red MSX Source (RMS) survey is an ongoing observational programme designed to return just such a sample with complementary data with which to address these issues.    The formation of massive stars, and the earliest stages of their evolution, take place inside cold ($\\sim$15~K), dense (H$_2$ $>$10$^4$~cm$^{-3}$), compact ($<$0.5~pc) clumps located within giant molecular clouds. Observationally these early stages should be characterised by strong emission at (sub)millimetre wavelengths due to cold dust and gas, with little or no emission at mid-infrared wavelengths as they have not yet begun to warm their surrounding cocoon of dust (\\citealt{beuther2007}).  (Sub)millimetre continuum observations of infrared dark clouds (IRDCs; \\citealt{rathborne2006}), towards IRAS selected sources (e.g., \\citealt{faundez2004,beltran2006}, and towards complete molecular complexes (e.g., \\citealt{motte2007}), have detected a significant number of cold, dense clumps and infrared quite cores.   These cold cores have similar masses ($\\sim$100~\\msun) and densities ($>$10$^5$~cm$^{-3}$) as warmer cores found towards known high-mass star forming regions, which has led to speculation they may represent the some of the earliest stages of massive star formation yet observed.  As the growing protostar begins to warm up its external environment it enters what is known as the hot molecular core stage as exotic molecules are evaporated from the surrounding dust grains (\\citealt{cesaroni2005}). Since the Kelvin-Helmholtz time scale is much shorter than the free-fall time, massive stars can begin fusing hydrogen in their cores while still deeply embedded within their infalling envelope.  As the envelope is dispersed the object emerges as a bright point source in the mid- and then near-infrared. This is the MYSO stage where the infrared luminosity has reached 10$^4$--10$^5$~L$_\\odot$, consistent with young O and early B-type stars (\\citealt{wynn-williams1982}). However, they do not yet ionise their surroundings, probably due to the ongoing accretion (see discussion by \\citealt{hoare2007b,hosokawa2008}). They are usually associated with massive bipolar molecular outflows (e.g., \\citealt{wu2004}) ionised stellar winds that are weak thermal radio sources ($\\sim$1~mJy at 1 kpc; \\citealt{hoare2002}). Finally, when the Lyman continuum radiation does escape the object enters the UCHII region phase, which expands and clears the remaining molecular cloud material in the immediate vicinity (\\citealt{hoare2007a}).   \\citet{lumsden2002} compared the colours of sources from the MSX and 2MASS point source catalogues to those of known MYSOs to develop a colour selection criteria from which we identified approximately 2000 MYSO candidates spread throughout the Galaxy ($|b|<$5\\degr). Sources toward the Galactic centre were excluded (defined as $|l|<10$\\degr) due to confusion and difficulties in calculating kinematic distances.  A problem with a colour selected sample is that the shape of the spectral energy distribution from an optically thick cloud is insensitive to the type of heating source. Therefore there are several other types of embedded, or dust enshrouded object, that have similar colours to MYSOs and contaminate our sample, such as ultra compact (UC) HII regions, evolved stars and a small number of planetary nebulae (PNe).  The RMS survey is a multi-wavelength programme of follow-up observations designed to distinguish between genuine MYSOs and these other embedded or dusty objects (\\citealt{hoare2005,urquhart2007c}) and to compile a database of complementary multi-wavelength data with which to study their properties.\\footnote{http://www.ast.leeds.ac.uk/RMS.} These include high resolution cm continuum observations to identify UCHII regions and PNe (\\citealt{urquhart2007a}), mid-infrared imaging to identify genuine point sources, obtain accurate astrometry and avoid excluding MYSOs located near UCHII regions (e.g., \\citealt{mottram2007}), and near-infrared spectroscopy (e.g., \\citealt{clarke2006}) to distinguish between MYSOs and evolved stars.  Another aspect of our follow-up campaign is to obtain kinematic distances which are crucial for calculating luminosity and hence distinguishing between nearby low- to intermediate-mass YSOs and genuine MYSOs. We have conducted a programme of $^{13}$CO observations towards all RMS sources not previously observed or for which good quality data was not available (see Sect.~\\ref{sect:grs} for details). In addition to being able to determine distances towards all sources where line profiles are detected, these observations allow us to eliminate a significant number of evolved stars ($\\sim$10--15\\%) since they are not generally associated with strong CO emission ($^{12}$CO $J$=1--0 and $J$=2--1 typically less than 1~K; \\citealt{loup1993}). In a recent paper (\\citealt{urquhart2007b}; hereafter Paper I) we presented the results of $^{13}$CO observations made towards 854 RMS source located in the southern Galactic plane (i.e., 180\\degr$<l<$350\\degr). In this paper we present the spectral line data towards 911 RMS sources located in the northern Galactic plane (NGP;10\\degr$<l<$180\\degr). These spectra have been obtained from a combination of 508 targeted observations and 403 spectra extracted from the Galactic Ring Survey (\\citealt{jackson2006}).\\footnote{http://www.bu.edu/galacticring/.} In Sect.~2 we outline our source selection, observational and data reduction procedures. We present our results and statistical analysis in Sect.~\\ref{sect:results}. In Sect.~\\ref{sect:summary} we present a summary of the results and highlight our main findings. ",
        "conclusions": "\\label{sect:summary} We present the results of a programme of molecular line observations of towards 911 candidate massive young stellar objects (MYSOs) located in the 1st and 2nd Quadrants of the Galaxy. The data presented here has been compiled from $^{13}$CO observations made using the JCMT, PMO, Onsala and Mopra telescopes and archival data from the GRS survey.  Emission is detected towards 780 RMS sources with two or more components being seen towards approximately 56\\%. Where available we have used a combination of associated stellar masers (water and methanol) and CS data to identify the cloud associated with an RMS source, and derive a criterion for selecting the most likely cloud associated with each RMS source where no useful tracer is present. In total we have been able to identify the component associated with our MYSO candidates in 379 cases of the 520 multiple detections. Combined with the single component detections we have obtained unambiguous kinematic velocities towards 638 sources ($\\sim$80\\% of the detections). The 141 sources for which we have not been able to determine the kinematic velocity will require additional line data. No $^{13}$CO emission is detected towards 131 RMS sources; these sources are fairly evenly distributed in both longitude and latitude and since they are not associated with any molecular material are very likely to be evolved stars often confirmed by other characteristics such as their IRAS colours.  These data complement the observations of the southern Galactic plane RMS sources presented in Paper~I. The full data set includes $^{13}$CO observations of 1765 MYSO candidates, and together with archival data provides molecular line data for our entire RMS catalogue. In turn this data set forms part of a larger programme of follow-up observations of near and mid-infrared colour-selected sample of MYSO candidates and are significant step in the delivery of the largest sample of MYSOs to date."
    },
    "0806/0806.2217_arXiv.txt": {
        "abstract": "{ We aim at the construction of luminosity and mass functions for Galactic open clusters, based on integrated magnitudes and tidal masses. We also aim at studying the evolution of these functions, with the ultimate purpose of deriving the \\textit{initial} luminosity and mass distributions of star clusters, independent of model assumptions regarding the cluster mass-to-light ratio. Finally we aim at a new determination of the percentage of field stars that have originated in open clusters.  The integrated magnitudes are computed from individual photometry of cluster members selected from the \\ascc catalogue. The cluster masses we assumed to be the estimated tidal mass recently published by us elsewhere. Analysis of the cluster brightness distribution as a function of apparent integrated magnitudes shows that the cluster sample drawn from the \\ascc is complete down to apparent integrated magnitude $I_V = 8$, with 440 clusters and compact associations above this completeness limit. This, on average, corresponds to a completeness area in the solar neighbourhood with an effective radius of about 1 kpc.  The observed luminosity function can be constructed in a range of absolute integrated magnitudes $I_{M_V}= [-10,\\,-0.5]$ mag, i.e. about 5 magnitudes deeper than in the most nearby galaxies. It increases linearly from the brightest limit to a turnover at about $I_{M_V}\\approx-2.5$. The slope of this linear portion is $a=0.41\\pm0.01$, which agrees perfectly with the slope deduced for star cluster observations in nearby galaxies. The masses of the Galactic clusters span a range from a few $M_\\odot$ to $\\log M_c/M_\\odot \\approx 5.5$. The mass function of these clusters can be fit as a linear function with log mass for $\\log M_c/M_\\odot > 2.5$, and shows a broad maximum between $\\log M_c/M_\\odot=1.5$ and 2.5. For $\\log M_c/M_\\odot >2.5$, the linear part of the upper cluster mass function has a slope $\\alpha=2.03\\pm0.05$, again in agreement with data on extragalactic clusters. We regard this agreement as indirect evidence that the tidal masses for Galactic clusters and the luminosity-based masses for extragalactic clusters are on the same scale.  Considering the evolution of the cluster mass function now reveals a slight but significant steepening of the slope with increasing age from $\\alpha=1.66\\pm0.14$ at $\\log t\\leqslant6.9$ to $\\alpha=2.13\\pm0.08$ at $\\log t\\leqslant8.5$. This indicates that open clusters are formed with a flatter (initial) mass distribution than the overall observed cluster mass distribution averaged over all ages. Interestingly, the luminosity function of open clusters does not show the same systematic steepening with age as the mass function does.  We find that the initial mass function of open clusters (CIMF) has a two-segment structure with the slopes $\\alpha=1.66\\pm0.14$ in the range $\\log M_c/M_\\odot=3.37\\dots4.93$ and $\\alpha=0.82\\pm0.14$ in the range $\\log M_c/M_\\odot=1.7\\dots3.37$. The average mass of open clusters at birth is $4.5\\cdot 10^3\\,M_\\odot$, which should be compared to the average observed mass of about $700\\,M_\\odot$. The average cluster formation rate derived from the comparison of initial and observed mass functions is $\\bar{\\upsilon}=0.4\\,\\,\\mathrm{kpc}^{-2}\\mathrm{Myr}^{-1}$. Multiplying by the age of the Galactic disc (T = 13 Gyr) the predicted surface density of Galactic disc field stars originating from dissolved open clusters amounts to $22\\,M_\\odot\\,\\mathrm{pc}^{-2}$ which is about 40\\% of the total surface density of the Galactic disc in the solar neighbourhood. Thus, we conclude that almost half of all field stars were born in open clusters, a much higher fraction than previously thought. ",
        "introduction": "Open clusters constitute an important part of a process transforming gas and dust into stars. They are observed as the most prominent parts in the regions of active star formation, or as tracers of the ceased star formation process in the general Galactic field. However, the role they are playing in this process has still not been fully understood. In spite of their prominence, there are indications that classical open clusters contribute only 10\\% or even less input \\citep{wiel71,misca78,clupop} to the total stellar population of the Galactic disc. This contradiction can be explained either by an early decay of a considerable fraction of newly formed star clusters \\citep[see e.g.][]{tut78,kah01,lamea} or by an insufficient knowledge of cluster formation statistics. In this context, one should note that the most important items of cluster formation like the distribution of cluster masses at birth (i.e., the initial mass function of star clusters) and the cluster formation rate are still poorly known.  Even in the close vicinity of the Sun, the only attempt to construct the luminosity function of open clusters \\citep{vdblaf84} is based on a sample of 142 clusters that, according to the authors, is to 2/3 complete within 400 pc. A mass function of Galactic open clusters is not yet available.  In contrast, extragalactic cluster populations have been actively investigated, especially during the present decade. The luminosity functions of remote clusters are determined both in active and regular galaxies \\citep[see e.g.][ and references therein]{slars,grijs03,giel06}. For some galaxies, luminosity distributions of open clusters are converted to the mass functions via the assumed theoretical mass-luminosity relation \\citep{zhang99}.  Using extragalactic data for the study of cluster population gives certain advantages. When observing external galaxies, one is able to cover large areas of their surfaces and catch the brightest objects, which are intrinsically rare. This is rather difficult in the case of our own Galaxy since they could be hidden by a dusty environment. Also, different selection effects can be taken into account better in an external galaxy because its clusters are located at about the same distance from the Sun. Moreover, studies of extragalactic clusters give a key to our understanding of the cluster formation in different environments.  However, there are a number of problems related to studies based on extragalactic clusters. Samples of extragalactic clusters are restricted to relatively bright objects. There is a danger of a contamination of cluster samples by brightest stars and of losing sparse objects in the strong background. For example, \\citet{slars} finds that his sample in NGC6949 is contaminated to about 20\\% by bright stars. The major problem is the necessity of transforming the observed luminosities into cluster masses. The corresponding transformations are mainly based on the luminosity of a dozen of the most massive stars, whereas cluster masses are defined by numerous stars of lower masses that are below the observing limit.  In our study of the local population of Galactic open clusters, we obtained data that now allow a reliable construction of their luminosity and mass functions. Our cluster sample contains 641 open clusters and 9 associations identified in the all-sky compiled catalogue \\ascc \\citep{kha01}. For each cluster we determined a homogeneous set of cluster parameters including, among others, its distance, age, and tidal mass. The sample is found to be complete up to a distance of about 850~pc from the Sun \\citep[]{clupop}. This completeness limit corresponds to a distance modulus $(V-M_V)$ of about 10--10.5 \\citep[]{clusim} and allow probes of the luminosity function of open clusters that are much deeper than in other galaxies.  In this paper we aim at constructing the luminosity and mass functions of the Galactic open clusters and at comparing them to the data in other galaxies in order to touch the issue of universality of the cluster formation process in different environments. We also consider temporal variations of the two distributions. The ultimate purpose of this study is to construct the initial luminosity and mass distributions of star clusters independent of model assumptions on the M/L-ratio, to derive their parameters (the shape and limits), and to draw conclusions on the impact of these results on our understanding of the role of open clusters in the evolution of the stellar population of the Galactic disc.  The paper has the following structure. In Sect.~\\ref{sec:dadef} we briefly describe the input data and give the definitions used throughout the paper. Sect.~\\ref{sec:lmf} is devoted to the construction of cluster the luminosity and mass functions. In Sect.~\\ref{sec:evol} we examine how cluster mass and luminosity functions evolve with time and build the cluster initial mass function. In Sects.~\\ref{sec:discuss} and \\ref{sec:concl} we discuss and summarise the results. ",
        "conclusions": "\\label{sec:concl}   Using the integrated magnitudes computed from individual photometry of cluster members available in the \\ascc, we found that the sample of 648 clusters drawn from this catalogue is complete for apparent integrated magnitudes brighter than 8 mag, with 440 clusters above this completeness limit. The corresponding completeness distances depending on the brightness of clusters are typically between 1 and 3 kpc, although some prominent clusters can be seen up to distances of 10 kpc and more.  The CPDLF constructed on the basis of the complete sample is observed in the range of $I_{M_V}= [-10,\\,-0.5]$ mag, i.e. about 5 magnitudes deeper than observable in nearby galaxies. It increases from the brightest limit to a turnover at about $I_{M_V}\\approx-2.5$ and slowly falls down towards fainter magnitudes. We consider the turnover to be a real feature reflecting the behaviour of Galactic open clusters. At magnitudes brighter than the turnover position, the CPDLF is linear with a slope $a=0.41\\pm0.01$ (or $\\alpha=2.02\\pm0.02$), which is in perfect agreement with the observed slopes of cluster luminosity functions in nearby spiral galaxies, although these show a higher star formation activity.  The masses of Galactic open clusters cover a range from several tens to hundreds of thousands of solar masses. The CPDMF mimics the general behaviour of the CPDLF: it shows a linear part at high masses ($\\log M_c > 2.5$), a broad maximum between $\\log M_c=1.5$ and 2.5, and a slow decline towards lower masses. In spite of the present-day observations only allowing for a comparison of the CPDMF data of extragalactic open clusters drawn from model mass-luminosity relations, the agreement with published values is impressive. We find that for $\\log M_c >2.5$ the CPDMF can by fitted with high confidence by a power law with a slope $\\chi=1.01\\pm0.04$. We regard the observed coincidence of cluster mass function parameters derived for the Galactic and extragalactic clusters with different methods as indirect evidence that both methods of estimating cluster masses are calibrated consistently and give correct and reliable results.  Inspection of cluster samples with different upper limits for age indicates that cluster mass spectra change with time. At every age the cluster mass function keeps the basic features of the CPDMF, i.e., a quasi-linear high-mass portion, and a non-linear portion at lower masses. With time, the slope of the linear portion increases from $\\chi=0.66\\pm0.14$ at $\\log t\\leqslant6.9$ to $\\chi=1.13\\pm0.08$ at $\\log t\\leqslant8.5$. Then the steepening of the CPDMF slows down and arrives at $\\chi=1.17\\pm0.07$ for $\\log t\\leqslant9.5$. The low-mass portion changes from an approximately flat distribution at $\\log t=6.9$ to a clearly non-linear behaviour displaying a broad maximum with a peak at about 100 $M_\\odot$.  We construct the CIMF from the data of a subsample of the 49 youngest clusters with $\\log t \\leqslant6.9$. The CIMF has a segmented structure with two power-law segments: the power-law slopes are $\\alpha=1.66\\pm0.14$ for $\\log M_c=3.37\\dots4.93$ and $\\alpha=0.82\\pm0.14$ for $\\log M_c=1.7\\dots3.37$.  The luminosity function of open clusters does not show the same systematic steepening with age as does the mass function. It evolves via parallel-shifting along the ordinate without changing the slope. Only the faintest segment of the CLF seems to depend on age and spreads towards the fainter magnitudes with age. The weak dependence on age and the presence of a mixture of clusters with different masses in every magnitude bin makes it difficult to determine the CIMF from the CLF.  When comparing the average mass of the newly formed, youngest clusters $\\olin{M}^{CIMF}_c\\simeq 4.5\\cdot 10^3\\,M_\\odot$ with the average cluster mass from the whole sample ($\\olin{M}^{CPDMF}_c\\simeq 700\\,M_\\odot$), one finds, as expected, that clusters are typically much more massive at birth than later. Since this observation concerns ``normal'' optical clusters, one should acknowledge the importance of a continuous mass loss occurring in open clusters during their evolution with a rate around 3-14 $M_\\odot$ Myr$^{-1}$.  If this effect is not taken into account, one underestimates the role that open clusters play in supplying the Galactic disc with the products of their evolution: stars and stellar remnants, sub-stellar objects, gas and dust. Provided that the cluster formation history has not changed dramatically in the solar neighbourhood during the evolution of the Galactic disc, at least 40\\% of the observed surface density of the disc comes from open clusters. The other half is left for violent events (strong winds, HII zones, or supernova explosions) provoking the fast mass loss during the ``infant'' phase of cluster evolution proposed first by \\citet{tut78} and later by \\citet{kah01}, and recently considered in detail by \\citet{baukr07} and \\citet{weidea07}."
    },
    "0806/0806.1487_arXiv.txt": {
        "abstract": "We describe an approximate statistical model for the sample variance distribution of the non-linear matter power spectrum that can be calibrated from limited numbers of simulations.  Our model retains the common assumption of a multivariate Normal distribution for the power spectrum band powers, but takes full account of the (parameter dependent) power spectrum covariance.  The model is calibrated using an extension of the framework in~\\textcite{habib07} to train Gaussian processes for the power spectrum mean and covariance given a set of simulation runs over a hypercube in parameter space.  We demonstrate the performance of this machinery by estimating the parameters of a power-law model for the power spectrum.  Within this framework, our calibrated sample variance distribution is robust to errors in the estimated covariance and shows rapid convergence of the posterior parameter constraints with the number of training simulations. ",
        "introduction": "The indirect nature of most cosmological observations usually requires numerical simulations of the data in order to infer constraints on cosmological models.  For parameter inference from the cosmic microwave background (CMB), galaxy and weak lensing surveys, and the Lyman~$\\alpha$ forest, the required simulations can be computationally expensive in order to capture the relevant physics, noise sources, and dynamic range.  The computational demands for future observations will only increase as more accurate theoretical predictions are required to match the reduced errors in the data.  In response to this foreseen bottleneck, several tools have recently been under development to reduce computational costs by emulating the output of cosmological simulations for the CMB and galaxy surveys given a training set of simulations~\\cite{pico,pico2,cosmonet,cosmonet2,heitmann06,habib07}.  These tools have been aimed at producing fast estimates of the mean simulation output, but often the error distribution for the data also needs to be inferred from simulations.  Typically, error models are constructed by running many realizations of a forward simulation of the data (or a compressed version of the data) at a fixed point in the model parameter space.  These multiple realizations can be used, for example, to construct covariance matrix estimates for use in inferring cosmological parameter distributions given the data.  If the error distribution is parameter dependent, then many more forward simulations run with varying input parameters could be required~\\cite{seljak03}.  We propose a unified framework for combining estimates of both the data mean and covariance matrix from the same set of simulations for cosmological parameter inference.  Our framework uses an efficient algorithm to interpolate between simulations run at sparse locations in parameter space and allows for propagation of interpolation errors into the inferred cosmological parameter constraints.  This is an extension of the method in Ref.~\\cite{habib07} in that sample covariance estimates at several points in parameter space are interpolated along with the sample mean estimates previously considered.  By requiring sample covariance matrices to be computed at many points in parameter space, our model might appear to require a large increase in the computational resources.  However, we also outline a general method to jointly constrain the covariance matrices for different parameters with the combined simulation realizations covering the whole parameter space.  We focus our validation tests on the prerequisite step of demonstrating the statistical framework when the covariances are already known.  Our model also provides a tool to determine whether the parameter dependence of the errors is important in any given application and a way to incorporate this parameter dependence when it is important (which are issues that can never be addressed by using jacknife covariance estimates from the data).  The parameter dependence of the errors is likely unimportant in any application where the parameters are known {\\it a priori} to be tightly constrained.  However, it may not be clear in any given application what constitutes ``tight'' constraints for the purposes of this approximation.  When the parameters are not tightly constrained, we expect it will probably be important to model the parameter dependence of the errors whenever performing inference from a reduced statistic of the data (because residual parameter dependence of the data can be absorbed into the error distribution for the reduced statistic).  We focus on the non-linear matter power spectrum in this paper as an example of this type of situation.  Because the non-linear matter distribution is non-Gaussian, the power spectrum is not a sufficient statistic and the variance of the power spectrum receives contributions from the (parameter dependent) connected four-point function.  It has already been shown~\\cite{sefusatti06} that the joint covariance of the two-point and three-point functions of the non-linear galaxy distribution has non-trivial and significant parameter dependence.    A simple example where the parameter-dependence of the errors is important is the measurement of the quadrupole of the CMB power spectrum (which has received considerable interest after WMAP reported a value somewhat lower than expectations). The dominant error on the quadrupole actually depends on the value of the quadrupole itself.  So, a naive analysis where one might attempt to construct the error distribution by running Monte Carlo simulations at a fixed point in parameter space would severely bias the inferred value of the quadrupole.  In fact, the properties of the large-scale CMB are simple enough that it is easy to analytically solve for the error distribution of the CMB quadrupole~(e.g. Ref.~\\cite{bond98}) or, using a sampling approach, even calculate the multivariate distribution of a whole set of multipole power amplitudes~\\cite{wandelt04}.  However, in most situations it is likely that the only recourse is to learn about the error distribution from simulations.  For example, the CMB power spectrum error distribution can no longer be calculated analytically once systematic errors and foreground modelling are included, yet the parameter dependence of the error distribution is likely to remain important.  This will not be the case in general, and the importance of modelling the noise variation over the parameter space will have to be decided on a case-by-case basis.   We explain our framework in the context of performing parameter inference from the non-linear matter power spectrum and have therefore limited the model for the (reduced) data error distribution to a multivariate Normal.  This model could be extended, for example, by considering a mixture of multivariate Normal distributions.  We have otherwise kept a general framework that can be applied to a wide array of applications.  This paper is organized as follows.  In Section~\\ref{sc:DMps} we give some background on the statistical properties of the dark matter power spectrum that serve as motivation for our framework.  In Section~\\ref{sc:framework} we describe our model for the power spectrum sample variance distribution and how to calibrate the model using simulations.  We then derive the joint likelihood of the simulation outputs and observed power spectrum for performing parameter estimation.  We test the performance of this framework with a toy model for the power spectrum in Section~\\ref{sc:validation}.  In Section~\\ref{sc:conclusions}, we summarize our results and outline future directions of this work.  A guide to the notation is given in Appendix~\\ref{sc:notation}, a covariance matrix parameterization that fits in our framework is given in Appendix~\\ref{sc:covparam}, and details of the likelihood calculation and evaluation are given in Appendices~\\ref{sc:emlike}, \\ref{sc:proposal}, \\ref{sc:priors}, and \\ref{sc:covmatexpressions}. ",
        "conclusions": "Conclusions}  We have demonstrated an extension to the statistical model of Ref.~\\cite{habib07} to estimate cosmological parameters from the power spectrum using a sample variance distribution calibrated from simulations.  This framework allows modelling of arbitrary, parameter-dependent power spectrum covariance matrices given several realizations of the power spectrum at a fixed number of points in parameter space.  We have focused on modelling the covariance of a multivariate Normal model for the estimated power spectrum in order to capture the correlations induced by filtering a Gaussian CMB or galaxy map or from non-linear graviational evolution in the matter power spectrum.  We tested the calibration of our model from simulations using a toy power-law model for the power spectrum.  In order to focus our tests, we used a simplified emulator that is conditioned on sample means and variances of the simulated power spectra rather than on the individual power spectrum realizations.  For this model, our calibration procedure converges quickly and is quite robust to reducing the number of simulation design points.  We expect that the requirement of computing converged sample covariance estimates at each design point is likely to be a strain on the simulation resources of actual galaxy and weak lensing survey analyses. Therefore, we have described a general formulation of the emulator that allows for constraining parameterized covariance matrices jointly with the other emulator parameters.  Again for our toy model, we have shown that while noisy covariance estimates bias the parameter constraints, the shift is small compared to the width of the parameter posterior distributions.  Our final goal with this work is to develop practical tools to aid in the estimation of cosmological parameters from future measurements of galaxy and cosmic shear power spectra.  As a next step we plan to demonstrate our calibration algorithm using N-body simulations of the dark matter density.  With N-body simulations, our framework provides the means to understand in which regimes modelling of non-linear evolution is important for estimating parameters and, as a related question, how much cosmological information can be extracted from non-linear scales in the dark matter distribution~\\cite{taylor01,rimes05b,neyrinck06b}.  The non-trivial effects of the survey window on the power spectrum covariance discussed in Refs.~\\cite{hamilton05} and~\\cite{sefusatti06} could potentially lead to biases in inferred parameter constraints without the careful modelling allowed by our framework.  In particular, the scaling of the ``beat-coupling'' effect described in Ref.~\\cite{hamilton05} with the fundamental modes in a survey implies extra parameter-dependence in the small-scale power spectrum covariance that could be significant in estimating cosmological parameters.  An emulator for N-body simulations will also provide valuable tests of the full emulator formulation presented here that conditions the GP models on the scatter between power spectrum realizations directly.  In this formulation (and a parameterization as in Appendix~\\ref{sc:covparam}), it may be possible to model the parameter dependent power spectrum covariance without any more simulations than are needed to accurately estimate the mean power.  \\begin{widetext} \\appendix"
    },
    "0806/0806.3482_arXiv.txt": {
        "abstract": "VLBI observations at 5 GHz have revealed that supernova 1979C, in the galaxy M100 in the Virgo cluster, has shell structure. The shell is approximately circular with the 50\\% contour deviating from a circle by an rms of no more than 8\\% of the radius.  The brightness distribution along the ridge may vary. The position of the center of the shell is at R.A.=\\Ra{12}{22}{58}{66758} and decl.=\\dec{15}{47}{51}{9695} (J2000), with a standard error of 0.8 mas in each of the coordinates.  No isolated central component is visible above a flux density limit of $\\sim$150 $\\mu$Jy which corresponds to an upper limit of $\\sim$15 times the corresponding spectral luminosity of the Crab Nebula. The radio lightcurve is clearly falling again and the radio spectrum is now flatter than at earlier times.  SN~1979C is only the fourth, and oldest (optically identified) supernova of which a detailed image could be obtained. ",
        "introduction": "\\objectname{SN~1979C} (type II-L) was discovered in the galaxy \\objectname{M100} \\citep[\\objectname{NGC~4321}, distance: 15.2 Mpc as estimated by][]{Freedman+2001} in the Virgo cluster on 1979 April 19 by G. E. Johnson \\citep{Mattei+1979} and subsequently became one of the most optically and radio luminous Type II supernovae ever found. On the basis of its radio lightcurve, its explosion date is estimated to be 1979 April 4 \\citep[see][]{Weiler+1986}, hereafter $t=0$~yr. VLBI observations started as early as 1982. They resulted in a determination of the angular expansion rate \\citep{Bartel1985, Bartel+1985a} and deceleration. This determination allowed an estimate of the mass-loss to wind-velocity ratio of the progenitor, which was found to be an order of magnitude lower \\citep{SN79C} than that reported from radio lightcurve fitting.  However, despite a quarter century of VLBI observations, the relatively far distance to Virgo has prevented determining the detailed structure of \\objectname{SN~1979C} with VLBI\\@. The determination of the expansion rate and the deceleration had to be based on size estimates made by using a model fit of the supernova brightness distribution.  Clearly, a determination of the supernova's morphology is important for a more accurate determination of the expansion velocity and deceleration. Such more accurate determination would also allow for a direct, geometric, estimate of the distance to \\objectname{SN~1979C} and Virgo using the expanding shock front method \\citep[ESM; see, e.g.,][]{SN93J-4} and resolve the ambiguity in an earlier estimate \\citep{SN79C}. It is also important for revealing the morphological type of the supernova: purely shell-like or center filled, perhaps with a bright compact component. Previously, only three other optically-identified supernovae could be imaged with an angular resolution high enough to unambiguously determine the morphological type. All three, \\objectname{SN 1986J} \\citep{Bartel+1991, SN86J-Sci}, \\objectname{SN 1987A} \\citep{Gaensler+1997, Gaensler+2007}, and \\objectname{SN 1993J} \\citep{Marcaide+1995a, BartelBR1995, SN93J-Sci, SN93J-3} have shell structure, with SN~1986J having a central component in addition to the shell. SN~1979C is the oldest optically-identified supernova that can still be monitored with VLBI\\@.  As such, it provides a bridge between our understanding of supernovae (interacting with the circumstellar medium left over from the progenitor) and supernova remnants (interacting with the interstellar medium).  Only a few supernova remnants which are less than a century old are known, all extragalactic, e.g., in M82 \\citep{McDonald+2001} and in Arp220 \\citep{Parra+2007}.  The youngest Galactic remnants are several centuries old, e.g., the Crab Nebula, which shows a pulsar wind nebula, Cas A, which shows a shell, and G21.5$-$0.9, which shows both a pulsar wind nebula \\citep[e.g.,][]{G21.5} and likely also a shell \\citep{MathesonS2005, Bocchino+2005}.  Since all radio supernovae originate from the explosions of massive stars, they are expected to leave a neutron star or a black hole.  Either a young neutron star or a black hole may be accompanied by a bright, compact radio source near the center of the expanding shell.  Young neutron stars are likely to be associated with radio-bright pulsar wind nebulae, and black holes may have radio bright jets due to an accretion-disk system.  Indeed, pulsar wind nebulae are found around a number of young pulsars, although the youngest known pulsars have ages near 1000~yr. A central radio source possibly associated with a neutron star or black hole was recently found in SN~1986J \\citep{SN86J-Sci}. Such a central radio source might also be seen near the centers of other older supernovae when clearly resolved images are obtained.  Here we report on such an image of SN~1979C, discuss its morphology, radio lightcurve, and spectral evolution, and search for signs of a central compact remnant of the explosion.  The expansion rate, deceleration, and distance determination will be reported elsewhere. ",
        "conclusions": "SN~1979C is the fourth (optically identified) extragalactic supernova for which an image with sufficient angular resolution could be made to determine the nature of the brightness distribution. The other three, as mentioned in the introduction, are SN 1986J, SN 1987A, and SN 1993J\\@.  Our VLBI image shows that SN~1979C has a morphology similar to the other three, namely a fairly to highly round shell, with the brightness possibly being modulated along the ridge. The possible brightness modulation likely arises from density irregularities either in the circumstellar medium or in the ejecta.  The flux density of SN~1979C is now clearly again on the downward trend.  Earlier, it was suggested that the radio lightcurve exhibits a quasi-periodic pattern \\citep{Weiler+1992} which was interpreted as possibly being due to the dynamical modulation of the wind from the red supergiant progenitor by a binary companion \\citep{SchwarzP1996}. With more data up to $t\\sim19$~yr, it was noticed that any periodic variation of the radio lightcurve was not extended beyond $t\\sim10$~yr.  Instead, it was suggested that, with the flattening of the radio lightcurve, the supernova had apparently entered a new stage of evolution \\citep{Montes+2000}.  With our new data up to $t\\sim26$~yr, which show that the lightcurve is again falling after $t \\sim 16$~yr, a new look at the characteristics of the radio lightcurve is warranted. In particular, we think that the evidence for a periodicity in the radio lightcurve or a new stage of evolution now seems weaker.  If at all, then the radio lightcurve appears to be periodic on a logarithmic time scale. However, a closer inspection leads us to think that, rather than being periodic on a logarithmic time scale, the radio lightcurve is modulated on a whole range of time scales. For instance, there appear to be variations on time scales of less than a year. There are also variations on time scales of several years with e.g., a 5-yr bump in the radio lightcurve at $t\\sim7$~yr and a bump lasting at least 15 yr at the end of the observed radio lightcurve.  These variations are likely caused by variations in the density of the ejecta \\citep[for discussions in the case of SN~1993J, see][]{MioduszewskiDB2001} and/or of the circumstellar medium. In case of the variations in the latter, the radio lightcurve can be thought of as a ``time machine'' which records in reverse the mass-loss to wind-velocity ratio of the progenitor for thousands of years before the star died.  The spectrum appears to have flattened somewhat with the spectral index being above the deviation range of $-0.76$ to $-0.63$ given by \\citet{Montes+2000}. A similar behavior is also seen for SN 1993J where the spectrum flattens while the source dimming accelerates\\footnote{Please note that although \\cite{Weiler+2007} claim that the spectral index of SN~1993J remains constant over time, an increase of the spectral index consistent with the measurements of \\citet{SN93J-2} and \\citet{Perez-TorresAM2002} can be clearly seen in their data with the smallest errors and/or scatter.} \\citep{SN93J-2, Perez-TorresAM2002}.  No clearly isolated component was found in the central region of the projected shell of SN~1979C\\@. However, the brightest portion of the shell appears to extend toward the center in projection, and could overlap with any compact component that may be present in the interior.  An isolated central component was in fact found for SN~1986J, with a spectrum peaking at $\\sim$20~GHz while the emission from the remaining parts of the supernova has an optically thin spectrum from at least 1.7 to 22~GHz \\citep{SN86J-CosparII, SN86J-Sci}. This isolated central component may be emission associated with the compact remnant of the explosion.  It is intriguing that the spectrum of SN~1979C shows evidence of flattening with time. Such flattening would be expected for a supernova with emission from a compact remnant that is starting to appear in the center of a shell which is expanding and becoming increasingly transparent at progressively higher frequencies. However, no central component can clearly be seen in SN~1979C, so the observed spectral flattening cannot be ascribed to such a component.  It is probable, therefore, that the spectral flattening is intrinsic to the shell's synchrotron emission.  The flattening must then be due to a change in the population of accelerated particles which produce the synchrotron emission.  Calculations have shown that physical conditions in supernova shocks can effect the spectral index of radio emission \\citep[e.g.,][]{Ellison+2007, EllisonBB2000}.   Our new upper limit on the flux density of any central compact source in SN~1979C of 150~$\\mu$Jy puts strong constraints on the luminosity of such a source. \\citet{BandieraPS1984} argued that young pulsar wind nebulae would have spectral luminosities of 10 to 1000 times that of the Crab Nebula. The spectral luminosity of the central component of SN 1986J is $\\sim$200 times that of the Crab Nebula between 14~GHz and 43~GHz \\citep{SN86J-CosparII}.  In contrast, the upper limit on the 5~GHz spectral luminosity of any central component in SN~1979C is only 15 times the corresponding spectral luminosity of the Crab Nebula, more than an order of magnitude fainter than both that in SN~1986J and our previous limit for SN~1979C \\citep{NB-COSPAR}, and at the low end of \\citet{BandieraPS1984}'s estimate. Any pulsar wind nebula that may exist inside the shell can probably not yet be detected because the shell is likely still opaque at radio frequencies.  Since the supernova shell becomes less dense and more transparent as it expands, any central component will become more easily detectable in the future. Searching for emission from the compact central remnant of the explosion therefore remains an intriguing prospect."
    },
    "0806/0806.3669_arXiv.txt": {
        "abstract": "From the analysis of the color-magnitude diagrams and color functions of four wide LMC fields located from $\\simeq$ 2 to 6 Kpc from the kinematic center of the LMC we present evidence that, while the oldest population is coeval in all fields, the age of the youngest component of the dominant stellar population  gradually increases with galactocentric distance, from currently active star formation in a field at 2\\degr3, to 100 Myr, 0.8 Gyr, and 1.5 Gyr in fields at 4\\degr0, 5\\degr5, and 7\\degr1, respectively. This outside-in quenching of the star formation in the LMC disk is correlated with the decreasing HI column density (which is $\\le 2\\times 10^{20} {\\rm cm}^{-2}$ in the two outermost fields with little or no current star formation). Other work in the literature hints at similar behavior in the stellar populations of irregular galaxies, and in M33. This is observational evidence against the inside-out disk formation scenario in low-mass spirals and irregular galaxies. Alternatively, it could be that the age distribution with radius results from interplay between the evolution with time of the star-forming area of the LMC and the subsequent outward migration of the stars. ",
        "introduction": "\\label{intro}  Modern hydrodynamical simulations suggest that, within the LCDM cosmology,  galaxies form from the inside out, i.e., their inner parts formed first, and they subsequently grew up to their present-day size (e.g. Brook et al.\\ 2006; Ro\\u{s}kar et al.\\ 2008). This can be interpreted either in terms of the age of the oldest population, which would be younger toward the external parts, or in terms of the mixture of populations, which would be younger on average there. However, the formation of a realistic disc galaxy within such simulations remains elusive (e.g. Governato et al. 2007; Abadi et al. 2003) and it is essential to augment such studies with detailed observations.    Observationally, two main avenues have been used to investigate disk formation and evolution. The first relies on the direct comparison of the actual physical sizes (or related parameters) of high-z disk galaxies with those of nearby galaxies (e.g. Trujillo \\& Aguerri 2004; P\\'erez 2004; Reshetnikov et al. 2003). The second focuses  on nearby galaxies, and  typically uses indicators of the ratio of recent to total star formation rate (SFR) as a function of radius (e.g., Mu\\~noz-Mateos et al.\\ 2007; Taylor et al.\\ 2005). Taylor et al.\\ (2005) found, in agreement with previous studies, that late-type spirals and irregular galaxies become on average redder outward. Since no plausible galaxy formation theory predicts positive metallicity gradients, they offer stellar age effects or dust as possible causes of this effect. The use of integrated properties makes it difficult to derive mass fractions as a function of age (and thus to distinguish between these two possibilities), or to determine the age of the oldest population as a function of radius. A direct census of stellar ages in a resolved galaxy would provide the necessary information to trace the characteristics of the stellar population as a function of radius, and thus to provide direct insight on formation and evolution mechanisms.  Nearby galaxies offer this possibility. Indeed, if the oldest main sequence (MS)  turnoffs are reached, the age of the oldest population as a function of radius can be measured and a quantitative age profile can be determined at each radius. In dIrr galaxies, a young  population is routinely found in their central parts, but it disappears  in the outskirts (e.g. Bernard et al.\\ 2007; Vansevicius et al.\\ 2004). In the cases where a CMD reaching the oldest MS turnoffs is available, it has been found that an intermediate age population extends to large distances from the center, and that it very gradually disappears toward the outer part (S. Hidalgo et al.\\ 2008, in preparation; Gallart et al.\\ 2004; No\\\"el \\& Gallart 2007), suggesting a \"shrinking scenario\" of star formation in dwarf galaxies (Hidalgo et al. 2003). A similar result is found in the case of M33 (see Barker et al.\\ 2007).  In this paper we analyze four 35\\arcmin $\\times$ 35\\arcmin \\ fields in the LMC at different galactocentric distances, using CMDs reaching the oldest MS turnoffs, in order to shed light on its stellar population gradients and thus on its formation and evolution. ",
        "conclusions": ""
    },
    "0806/0806.1346_arXiv.txt": {
        "abstract": "{Many experimental results around and above the energies where the solar modulation affects cosmic ion fluxes were quantified, conceptualized  and debated using leaky box models. These models exploit the notion of equilibrium between creation and destruction processes of cosmic ions in an undifferentiated arbitrary volume representing the Galaxy, ignoring the galactic magnetic field, the size of the Galaxy, the position of the solar cavity, the spatial distribution of the sources, the space variation of the interstellar matter and other pertinent observations. Progress in the measurements of the quoted observational parameters rendes obsolete the use of the leaky box models. Specific examples substantiating the inadequacy of the leaky box models are analyzed such as the conversion of the boron-to-carbon flux ratio into grammage and the residence times of cosmic ions in the Galaxy. The unphysical and misleading nature of the leaky box models is ascertained and illustrated at very high energy.}     \\begin{document} ",
        "introduction": "Many experimental results on the primary cosmic radiation at very low energy gathered in the last five decades by balloon-borne instruments and satellite experiments have been expressed using simple theories of galactic cosmic rays referred to as Leaky Box Models. Two basic physical quantities, namely the gas column encountered by Galactic cosmic rays ($grammage$) and the related residence time ($T$), were believed to correctly interpret a number of measurements. Ritual fittings of the grammage above an adjustable energy $E_0$ (or the rigidity $R_0$) based on numerous measurements have the form: \\begin{equation} g = G {r}^{-\\delta} \\end{equation} where $G$ is a constant grammage at the rigidity $R_0$, $\\delta$ a constant and $r$ = $R$/$R_0$, with $R$ the ion rigidity. Numerical values in classical interpolations are, for example: $G$=$10.8$ g/cm$^2$, $\\delta$=$0.6$ and $R_0$=$4$ GV \\cite{ref1}, or $G$=$24.0$ g/cm$^{2}$, $\\delta$=$0.65$ and $R_0$=$5.5$ GV \\cite{ref2}.  Grammage is converted into residence time by: $T$ = $g$/$\\rho$$v$ where $v$ is the ion velocity and  $\\rho$ the matter density. The extrapolation at very high energy of the residence time caused fictitious problems, as pointed out by Hillas \\cite{ref3,ref4}. Figure 1 deliberately shows on a linear scale in energy the B/C flux ratio in the energy band $10$ - $200$ GeV/u indicating no compelling empirical evidence for a decrease of the grammage of the form $r^{-\\delta}$ which, on the contrary, is well established in the region $1 - 4$ GeV/u. The recent measurements of the B/C flux ratio at $700$ GeV/u by the Runjob Collaboration \\cite{ref5}, along with the data shown in figure 1, exclude the functional form \\cite{ref1} above $20$ GeV/u. The data in figure 1 and the B/C flux ratio [5] vividly testify to the inadequacy of the grammage fitting $via$ $r^{-\\delta}$ to physical reality in an area believed in past decades to be the realm of the Leaky Box Models.  \\begin{figure}[t]         % \\begin{center} \\includegraphics [width=0.48\\textwidth]{icrc0705_fig01.eps} \\end{center} \\caption{Measurements of the boron-to-carbon flux ratios above $10$ GeV/u up to the energy of $200$ GeV/u displayed in a linear scale of energy. No decreasing trend of the B/C flux ratio is evident. Data below $10$ GeV/u may be found elsewhere \\cite{ref18}.}\\label{fig1} \\end{figure}  \\begin{figure}[t]         % \\begin{center} \\includegraphics [width=0.48\\textwidth]{icrc0705_fig02.eps} \\end{center} \\caption{Interpolated grammage $via$ $G$ ${r}^{-\\delta}$ (thin lines $a$ and $b$) for two sets of constants mentioned in the Introduction. It is evident that the slope $\\delta$ tuned to the B/C flux ratio in the range $1-5$ GeV/u is not in accord with the experimental data above $10$ GeV/u. The computed grammage (c) is a tiny fraction of that displayed in figure 3.}\\label{fig2} \\end{figure} ",
        "conclusions": ""
    },
    "0806/0806.3496_arXiv.txt": {
        "abstract": "Modified dispersion relations from effective field theory are shown to alter the Chandrasekhar mass limit. At exceptionally high densities, the modifications affect the pressure of a degenerate electron gas and can increase or decrease the mass limit, depending on the sign of the modifications. These changes to the mass limit are unlikely to be relevant for the astrophysics of white dwarf or neutron stars due to well-known dynamical instabilities that occur at lower densities.  Generalizations to frameworks other than effective field theory are discussed. ",
        "introduction": "The principle of Lorentz Invariance is at the heart of the contemporary formulation of physical theory, in particular the Standard Model and general relativity.  Given the centrality of Lorentz Invariance (LI) it is wise to explore as many avenues as possible that test this principle.  One avenue investigated in recent years is the physics of modified dispersion relations (MDR).  A remarkable result of these studies is that astrophysical data significantly limit Planck scale - suppressed modifications (see e.g. \\cite{jlm_rev,d_rev} for reviews). The successful limits on the modifications are due to the high degree of sensitivity of particle process thresholds to Lorentz Violation (LV). Given the delicate interplay between the modifications of the particles involved, it is helpful to look for systems in which the effects of the modification on a single particle type are isolated.  An apparently ideal system is the physics of a degenerate electron gas since the pressure of such a gas supports the gravitational attraction of white dwarfs.  Further, the Chandrasekhar mass limit, about $1.4 M_\\odot$, of white dwarf stars is obtained in the ultra-relativistic limit, precisely where modifications of dispersion relations are expected to be large.  Modified dispersion relations often take the form of an expansion in LV terms \\begin{equation} E^2 =p^2 +m^2 + \\kappa_3 \\frac{p^3}{M_P} + \\kappa_4 \\frac{p^4}{M_P^2} + ... \\end{equation} where the parameters $\\kappa_i$ may differ for different particle species and $M_P$ is the Planck mass, which we take to be $M_P = \\sqrt{ \\hbar / ( 4 \\pi G)} \\approx 3.45 \\times 10^{27}$ eV ($c=1$).  In such MDR models the usual energy-momentum conservation laws hold so there is a preferred frame, which we take to be the one where the cosmic microwave background is isotropic\\footnote{At the conclusion of the paper we comment on these results in MDR frameworks in which LI may not be broken.}.  While these modified dispersion relations could be viewed simply as a phenomenological expansion to test LV, this form has been suggested in a variety of settings including string theory tensor vacuum expectation values, heuristic calculations of the semiclassical limit of loop quantum gravity,  spacetime foam, non-commutative geometry, analogs of emergent gravity, and some braneworld models \\cite{jlm_rev,d_rev}.  Given the energy scale of the modifications, it might seem that testing such modifications might  simply be impossible,   However even the early work \\cite{limits,jlm} demonstrated that particle process thresholds are highly sensitive to these modifications.  Using several particle processes and observed energies, much of the parameter space is ruled out \\cite{limits,jlm}.  For instance the recent work of Macione {\\em et. al.} achieves a limit on the parameter space of electrons  of less than $10^{-5}$ \\cite{new_crab}.  In this paper we report on numerical solutions to the exact equations for the Chandrasekhar mass limit with modified dispersion relations. We extend the analysis of \\cite{camacho}, relaxing the unphysical assumption of constant density.  We find significant differences with the reported results: The mass limit may be raised or lowered depending on the sign of the modifications in the electron dispersion relation; and physical equilibrium radii exist for both signs of the MDR parameter $\\kappa$. Finally we show that, despite the existence of new mass limits, there would be no effect on white dwarf astrophysics.  The effects are only important in the Planck-scale regime far above astrophysically accessible densities.  This paper is organized as follows.  The next two sub-sections are devoted to discussions of MDR and the mass limit.  In section \\ref{mod_mass} we derive the corrections to the mass limit due to MDR.  In the final section \\ref{conclude} we summarize the results, compare with results previously reported in \\cite{camacho}, and comment on the applicability of the calculation to other frameworks with modified dispersion relations.  \\subsection{Modified Dispersion Relations}  To achieve precise limits on the parameters it is necessary to have some additional knowledge of the dynamics of the field theories. This may be achieved in the context of effective field theory, where effective field theory is used to determine the (non-renormalizable) mass dimension five (or higher) LV operators.  Myers and Pospelov found that there are essentially only three operators that simultaneously break local LI and preserve gauge and rotation invariance.  Introducing a preferred frame field $n^a$, the mass dimension 5 operators are \\cite{MP} \\begin{equation} - \\frac{\\zeta}{M_P} n^a F_{ac}\\, n \\cdot \\partial ( n_b \\tilde{F}^{bc} ) + \\frac{1}{2 M_P} n^a \\bar{\\psi} \\gamma_a ( \\xi_1 + \\xi_2 \\gamma_5)(n \\cdot \\partial)^2 \\psi \\end{equation} where $ \\zeta, \\xi_i$ are dimensionless parameters, $\\tilde{F}^{ab}$ is the dual of the usual $F_{ab} = \\partial_a A_b - \\partial_b A_a$ of Maxwell theory.  The dispersion relations for fermions become \\begin{equation} E_\\pm = p^2 +m^2 + \\eta_\\pm \\frac{p^3}{M_P} \\end{equation} where the different signs correspond to the two helicity states. Typically, the modifications become important when they become comparable to the mass terms.  For an electron this occurs at an energy scale of $\\sqrt[3]{m_e^2 M_P} \\sim$ 10 TeV - significantly below the Planck scale.  While these energy scales are beyond terrestrial experiments, astrophysical processes at these energy scales effectively constrain the cubic modifications \\cite{jlm,new_crab,limits}.  The fact that electron cubic MDR have two parameters leads to new effects. At sufficiently high energies electrons are unstable to helicity decay \\cite{jlm_rev}.  If $\\eta_-> \\eta_+$ then the negative helicity electrons will decay producing a photon and positive helicity electron.  While there is no kinematic threshold for this process, an effective threshold may be derived by studying the reaction rate.  The resulting effective threshold is $(m_e^2 M_P/ \\Delta \\eta)^{1/3}$ or about 10 TeV for $\\Delta \\eta = |\\eta_+ -\\eta_-| < 0.4$, as determined by photon stability \\cite{jlm_rev}.  The lifetimes of the negative helicity electron decrease around this effective energy; lifetimes for a 1 TeV electron is about 1 s while for a 50 TeV electron it is reduced to about $10^{-9}$ s \\cite{jlm_rev}.  However the lifetimes increase at higher energies due to the fact that the states become more chiral. The MDR effects on the Chandrasekhar mass only occur at such high energies that the electron population is effectively in one helicity state. Assuming that the astrophysical processes, e.g. accretion, driving the increase in electron energy are sufficiently slow as compared to the lifetimes of the helicity states all the electrons would have positive helicity.  Since the same analysis applies for the other case,  $\\eta_+> \\eta_-$, the resulting degenerate electron gas will all be in the helicity state associated with the smallest parameter.  Hence we will denote the electron parameter with $\\kappa$, where $ \\kappa = \\text{min}(\\eta_+,\\eta_-)$.  For the purposes of studying the effect on the physics of compact stars we assume $\\kappa$ is order one and write the MDR as \\begin{equation} \\label{mdr} E^2 = p^2 \\left[ 1 + \\kappa \\left( \\frac{p}{M_P} \\right)^n \\right] +m^2. \\end{equation} While the case of dimension 5 operators and cubic modifications ($n=1$) is the most interesting, since the analysis easily generalizes we present the mass limit calculation for arbitrary $n$.  The range of validity for this effective description is $p<M_P$.   \\subsection{Chandrasekhar Mass Limit}  White dwarf stars are an end stage of stellar evolution for stellar masses less than $\\sim 8 M_\\odot$.   As the star cools and contracts, the electron wavefunctions begin to overlap appreciably. As the star contracts further, the electrons are forced into higher and higher momentum states to satisfy the Pauli exclusion principle, forming a  degenerate electron gas with relativistic Fermi energy. Because the thermal energy is so much lower than the Fermi energy the temperature is effectively zero. The Chandrasekhar mass limit, about $1.4 M_\\odot$,  of white dwarfs stars arises when the Newtonian gravitational attraction due to the nuclei (often comprised of carbon and helium) is balanced by the outwards pressure of degenerate, cold electrons. The mass limit is obtained for an ultra-relativistic Fermi gas, when $E/m_e \\sim 10^4$ with central densities  in excess of $ \\rho \\sim 10^6$ g cm$^{-3}$. Thus the Chandrasekhar mass limit is derived in the framework of Newtonian gravitation and the statistical mechanics of an ideal, non-interacting, ultra-relativistic gas of electrons at $T=0$.  The Chandrasekhar mass limit is largely of theoretical interest. Other physics comes into play at high densities introducing processes known as dynamical instabilities.  Depending on the details of the star's composition, two processes transform the physics of the star. At densities  $\\rho \\sim 10^9 - 10^{11}$ the electrons acquire sufficient energy to induce inverse $\\beta$-decay and the star turns into a neutron star in a process called neutronization.  At the critical density of  $\\sim 10^{10}$ g cm$^{-3}$ the onset of general relativistic instabilities causes the star to collapse.  These threshold densities are composition dependent so, for instance, the general relativistic instability is irrelevant to iron white dwarfs, which undergo neutronization before gravitational collapse.  For helium and carbon white dwarfs, the situation is reversed, instabilities due to general relativity occur before neutronization. ",
        "conclusions": "\\label{conclude}  Unlike the threshold calculations, which are rich in new features and lead to strong constraints on the parameters (e.g. \\cite{new_crab}), the calculation of the Chandrasekhar mass limit using MDR is straightforward.  The mass limit is raised or lowered according to the sign of the modification. As one can see in the numerical results, the qualitative argument, or from equation (\\ref{approx}) for high $z_c$, the effects become important when $\\delta z_c^n Q^n \\sim 1$ or near the Planck scale, $E/M_P \\sim 1$.  This is well-beyond the scale at which other well-known processes occur so these new mass limits are not directly astrophysically relevant.  This exact analysis corrects the results reported in \\cite{camacho}, where it is suggested that current data rule out the $\\kappa > 0$ case ($\\alpha < 0$ in \\cite{camacho}).  In the analysis of \\cite{camacho} the MDR correction to the radius, $\\delta R$,  is \\begin{equation} \\delta R = \\kappa \\frac{\\hbar}{5 M_P}  \\left( \\frac{9 \\pi M}{m_u} \\right)^{1/3} \\left[1 - \\left( \\frac{M}{\\tilde{M}} \\right)^{2/3} \\right]^{-3/2} \\end{equation} where $\\tilde{M}^{2/3} = (\\hbar/3 \\pi G) ( 9 \\pi /8 m_u)^{4/3}$.  Hence, as $M \\rightarrow \\tilde{M}$ the MDR correction grows arbitrarily large.  This observation led Camacho to conclude that current data seems to rule out positive values of $\\kappa$.  This conclusion rests from the incorrect assumption that the density is uniform.  As the present analysis shows (see figure \\ref{density}) the density is not constant throughout the star.  While the radius does increase with central energies for positive $\\kappa$, the actual radii at these energies are so small that other processes such as neutronization and gravitational collapse occur long before the star evolves to that state.  Current observations seem to indicate white dwarfs with smaller radii than expected, at least for iron cores \\cite{mathews,provencal}.  The conclusion of \\cite{camacho} suggests that solutions with negative $\\kappa$ (or $\\alpha > 0$), when the white dwarf radius in the MDR case is reduced, lead to a possible explanation for the smaller radii. However as seen in the result reported after equation (\\ref{negkapmass}), the model with the realistic density shows significant corrections only at the Planck scale.  The argument rests on fermion statistics and the modified dispersion relation.  Hence, in so far as these relations hold, it would seem that the calculations of the mass limits generalize to frameworks outside the EFT approach. However, as pointed out in \\cite{dsr_chandra}, this neglects the possible changes in the integration measure of the pressure, equation \\ref{Pform}.  For instance, in deformed special relativity (DSR) \\cite{AC,MS}  the modified dispersion relations are regarded as the invariant of a relativity group with two invariant scales - not only the speed of massless modes, $c$, but also an invariant length or energy\\footnote{A particle process threshold analysis of DSR theories does not provide significant constraints on the parameters of the theory.  This is due the fact that in the DSR framework there are no new thresholds and the usual special relativity thresholds are only slightly modified \\cite{dsrthresh}.}. As interpreted by Hossenfelder \\cite{H} the DSR dispersion relations arise from an effective description of gravitational effects at high energy-density.   The calculation of the neutron star mass in DSR was carried out in \\cite{dsr_chandra}.  So barring modification of statistics or of the measure, the present calculations are valid {\\em mutatis mutandis} in a more general context. The new mass limits would appear not to have astrophysical relevance.  Neutronization and the effects of general relativity would still occur at much lower momenta.  Recently there has been work on the impact of MDRs on the structure of astrophysical objets \\cite{bz}.  The MDR effects on the Chandrasekhar mass limit may also be interpreted in the context of generalized uncertainty principle.  The non-linear relation between particle momentum in a high energy-density region and the asymptotic region, induces a modification in the Heisenberg uncertainty principle \\cite{H} \\begin{equation} \\Delta x \\Delta p \\geq \\frac{\\hbar}{2} \\left| \\ev{ \\frac{\\partial p}{\\partial k} } \\right| \\end{equation} Expanding the modification as  $\\partial p/ \\partial k = 1+ \\kappa (p/M_P)^n$ and following the method of estimating the ground state energy of the hydrogen atom from the uncertainty principle, one may show that the pressure is modified as in equation (\\ref{Pform}) (up to a factor of 1/2).  This leads to the same effects as derived in the EFT context."
    },
    "0806/0806.4049_arXiv.txt": {
        "abstract": "{The main sequence binary star 61\\,Cyg (K5V+K7V) is our nearest stellar neighbour in the northern hemisphere. This proximity makes it a particularly well suited system for very high accuracy interferometric radius measurements.} {Our goal is to constrain the poorly known evolutionary status and age of this bright binary star.} {We obtained high accuracy interferometric observations in the infrared K$^\\prime$ band, using the CHARA/FLUOR instrument.  We then computed evolutionary models of 61\\,Cyg~A \\& B with the CESAM2k code. As model constraints, we used a combination of observational parameters from classical observation methods (photometry, spectroscopy) as well as our new interferometric radii.} {The measured limb darkened disk angular diameters are $\\theta_{\\rm LD}(A) = 1.775 \\pm 0.013$\\,mas and $\\theta_{\\rm LD}(B) = 1.581 \\pm 0.022$\\,mas, respectively for 61\\,Cyg~A and B. Considering the high accuracy parallaxes available, these values translate into photospheric radii of ${\\rm R}(A) = 0.665 \\pm 0.005\\ {\\rm R}_{\\odot}$ and ${\\rm R}(B) = 0.595 \\pm 0.008\\ {\\rm R}_{\\odot}$. The new radii constrain efficiently the physical parameters adopted for the modeling of both stars, allowing us to predict asteroseismic frequencies based on our best-fit models.} {The CESAM2k evolutionary models indicate an age around 6\\,Gyrs and are compatible with small values of the mixing length parameter. The measurement of asteroseismic oscillation frequencies in 61\\,Cyg~A \\& B would be of great value to improve the modeling of this important fiducial stellar system, in particular to better constrain the masses.} ",
        "introduction": "Binary stars are numerous in the Galaxy but only few of them can be fully calibrated physically. They are interesting for several reasons, among them the most important for evolutionary modeling is that the two components are coeval, and share the same chemical abundances of the original material. This makes their modeling much easier than for single stars, as several free parameters of the models are shared between the two stars (age and helium content). Possibly the best example of physical calibration of a binary system is provided by our closest neighbour, $\\alpha$\\,Cen (Th\\'evenin et al.~\\cite{thevenin02}). Its two solar-type components were analyzed in detail with the use of photometry, spectrometry, astrometry and interferometry. It resulted in excellent knowledge of the masses of the two components (within less than 1\\%), and, through modeling of the stars, of fundamental parameters not directly accessible to measurements, like the age or the helium  content. Therefore, these two stars are, with Procyon, Sirius and Vega, among the most important benchmark stars for new developments of stellar structure and atmosphere models. The present study focusses on 61\\,Cyg~A (\\object{HD 201091}, \\object{HIP 104214}) and B (\\object{HD 201092}, \\object{HIP 104217}), the nearest stars in the northern hemisphere. The two stars constitute a visual binary pair with a very long orbital period ($\\approx$ 650\\,yrs), also known as \\object{Gl 820}. The parallax of ths system was first measured by Bessel~(\\cite{bessel1838}), and it is now known with extremely high accuracy. The proper motion of more than 5${\\arcsec}$ per year, first determined by Piazzi in the XVIII$^{\\rm th}$ century, makes it one of the fastest moving stars in terms of apparent displacement. Although some of this motion comes from the proximity of 61\\,Cyg to us, the pair has also a high radial velocity of 108 km/s, indicating that 61\\,Cyg is not a member of the thin disk of our Galaxy. The proximity of 61\\,Cyg makes it a northern analog of the numerical modeling benchmark $\\alpha$\\,Cen. Its large parallax also means that this is the easiest low-mass dwarf to resolve interferometrically. The spectral types of its two members (K5V and K7V) ideally complement our previous studies of $\\alpha$\\,Cen A \\& B (G2V+K1V, Kervella et al.~\\cite{kervella03a}, Bigot et al.~\\cite{bigot06}). The masses of 61\\,Cyg A \\& B are controversial, with estimates ranging from approximately 0.74 and 0.46 M$_\\odot$ (Gorshanov et al.~\\cite{gorshanov06}) to 0.67 and 0.59\\,M$_\\odot$ (Walker et al.~\\cite{walker95}). With effective temperatures of about 4\\,400 and 4\\,000\\,K, they shine with luminosities of only 0.15 and 0.08\\,L$_\\odot$. There is no confirmed planet around them, although indications exist that 61\\,Cyg~B could host a giant planetary companion (see. Sect.~\\ref{discussion}). As dimmer versions of solar type stars, they have magnetic cycles similar to that of the Sun, the brighter 8 years, the fainter 11 years (Hempelmann et al.~\\cite{hempelmann06}). Their rotation periods are of the order of 35 days, therefore no rotational distorsion of their photospheres is expected. The abundances of heavy chemical elements have been determined (Luck \\& Heiter~\\cite{luck05}, \\cite{luck06}) in these stars which are found slightly metal poor (-0.2\\,dex), so a priori older than the Sun but belonging to the galactic disk.  In Sect.~\\ref{observations}, we detail the interferometric observations and the corresponding physical parameters we derive (angular diameters and linear radii). Together with the additional observables listed in Sect.~\\ref{constraints}, we propose in Sect.~\\ref{modelfit} a modeling of the two stars using the CESAM2k code. We finally present asteroseismic frequency predictions in Sect.~\\ref{astero}. ",
        "conclusions": "}  We presented high accuracy interferometric measurements of the angular diameters of the two nearby stars 61\\,Cyg~A \\& B: $\\theta_{\\rm LD}(A) = 1.775 \\pm 0.013$\\,mas and $\\theta_{\\rm LD}(B) = 1.581 \\pm 0.022$\\,mas, corresponding to photospheric radii of ${\\rm R}(A) = 0.665 \\pm 0.005\\ {\\rm R}_{\\odot}$ and ${\\rm R}(B) = 0.595 \\pm 0.008\\ {\\rm R}_{\\odot}$. We computed CESAM2k models that reproduce these radii as well as the other observed properties of the two stars. 61\\,Cyg~A \\& B appear as very promising targets for future asteroseismic studies, and we also derived asteroseismic frequencies potentially present in these two stars. The detection of oscillations would bring important constraints to stellar structure models in the cool, low-mass part of the HR diagram, where convection plays a central role. However, in the absence of measured asteroseismic frequencies, it appears difficult to go beyond the present modeling of the binary system. The main reason is that the mass is not constrained sufficiently well by the long period astrometric orbit. We encourage asteroseismic groups to include these two stars in their observing programmes, as the measurement of seismic parameters (in particular the mean large frequency spacing) will bring a decisive constraint to the mass of these stars. We also found that the value of the mixing length parameters for both stars is very well constrained by the radius once the other parameters of the model are fixed, in particular the mass. We reach the same conclusions as Yildiz et al.~(\\cite{yildiz06}) that the $\\alpha$ parameter is mass dependent at least for solar abundance stars. The future {\\it Gaia} mission (Perryman~\\cite{perryman05}) will open the way to a complete calibration of several thousand binaries. Such a large sample will give us a completely renewed view of convection in stars (see e.g. Lebreton~\\cite{lebreton05})."
    },
    "0806/0806.1200_arXiv.txt": {
        "abstract": "We present the discovery of the Vela-like radio pulsar \\palfapsrshort\\ in the Arecibo PALFA survey.  \\palfapsr\\ has a spin period of 81\\,ms, a characteristic age of 21\\,kyr, and a spin-down luminosity $\\dot{E} = 4.6 \\times 10^{36}$\\,ergs s$^{-1}$.  It is positionally coincident with the TeV $\\gamma$-ray source \\hesssrc, which has no other known counterparts. Young, energetic pulsars create wind nebulae, and more than a dozen pulsar wind nebulae have been associated with very-high-energy (100\\,GeV$-$100\\,TeV) $\\gamma$-ray sources discovered with the HESS telescope.  The $\\gamma$-ray emission seen from \\hesssrc\\ is potentially produced by a pulsar wind nebula powered by \\palfapsr; faint X-ray emission detected by \\asca\\ at the pulsar's position supports this hypothesis. The inferred $\\gamma$-ray efficiency is $\\epsilon_{\\gamma} = L_{\\gamma}/\\dot{E} = 3.1$\\,\\% ($1-10$\\,TeV, for a distance of 9\\,kpc), comparable to that observed in similar associations. ",
        "introduction": "In the last few years, approximately 40 new Galactic sources of very-high-energy (VHE) $\\gamma$-ray emission (100\\,GeV$-$100\\,TeV) have been discovered using the High Energy Stereoscopic System (HESS\\footnote{See http://www.mpi-hd.mpg.de/hfm/HESS/public/HESS\\_catalog.htm for a catalog of HESS-detected sources.}) Cherenkov telescope array \\citep[e.g.,][]{aaa+05a}.  These are an exciting new population of sources, which give new insight into non-thermal particle acceleration in Galactic objects such as neutron stars, supernova remnants, and X-ray binaries.  Thus far, close to half of these sources have been established or suggested as being associated with the pulsar wind nebulae (PWNe) of young pulsars, either through direct detection of a PWN or positional coincidence with a young pulsar which is presumed to have a PWN (Table~\\ref{tab:assoc}). Clearly, PWNe are now an important and well-established Galactic source of VHE $\\gamma$-rays.  Since both young pulsars and their PWNe can be very dim, many of the Galactic HESS sources without identified counterparts \\citep[e.g.,][]{aab+08a} may be PWNe, potentially detectable via deep radio or X-ray observations.  In general, radio/X-ray PWNe are associated with extended HESS sources, presumably TeV PWNe, whose peak is offset by several arcminutes from the pulsar\\footnote{Note however that this situation does not appear to hold in the case of the youngest pulsars, e.g. the Crab, where there is no observed offset and the TeV emission is consistent with a point source.  This possibly demonstrates the evolution of TeV PWNe with pulsar age.} (Table~\\ref{tab:assoc}). In some cases the offset, if any, cannot be measured because the position of the pulsar is not known \\citep[see also][]{gal07}.  These offsets have been explained by the hypothesis that the VHE emission is from inverse Compton scattering of ``old'' electrons, which were produced during an earlier epoch in the pulsar's life, off of the ambient photon field (e.g., cosmic microwave background radiation, starlight, and infrared emission from dust; see Aharonian et al. 2005d\\nocite{aab+05d}, de Jager \\& Djannati-Ata\\\"i 2008\\nocite{dd08}). An alternative model has the $\\gamma$-rays produced by the decay of $\\pi^0$ mesons, which are created by the interaction of accelerated hadrons with nuclei in the interstellar medium \\citep{has+06,hahs07}. In both cases, asymmetric crushing of the pulsar wind by the reverse shock of the expanding supernova remnant, in cases where the remnant has expanded into an asymmetric density distribution in the ambient medium, is then responsible for the offset between the pulsar and the peak of the VHE emission \\citep{dd08}.  A high pulsar proper motion may also play a role in the offset. The VHE emission could be a key into the earlier energetic history of these young pulsars as well as completing a broadband (radio to TeV $\\gamma$-rays) picture of the shocked pulsar wind, whose emission would be dominated by synchrotron-emitting electrons below $\\sim 1$\\,GeV and inverse Compton emission from the scattering of these same particles above this energy.  In this Letter, we present the discovery and subsequent timing of the ``Vela-like'' pulsar \\citep[i.e. pulsars having characteristic age $\\tau_c = 10-30$\\,kyr and spin-down luminosity $\\dot{E} \\sim 10^{36}$\\,ergs s$^{-1}$; e.g.,][]{kbm+03} \\palfapsr\\ in the Arecibo PALFA survey of the Galactic Plane.  This young, energetic pulsar is positionally coincident with, and energetically capable of creating, the VHE emission observed from the hitherto unidentified TeV source \\hesssrc\\ \\citep{aab+08a}.  This association suggests that some of the other currently unidentified, extended HESS sources in the Galactic plane may also be related to faint radio pulsars, rather than some new source class.  An archival \\asca\\ image of the area around \\palfapsr\\ and \\hesssrc\\ shows a possible X-ray counterpart, cataloged as \\ascaxray\\ by \\citet{smk+01}, at the pulsar position; this is possibly a synchrotron counterpart to the TeV PWN\\footnote{ Furthermore, in a recent reanalysis of $EGRET$ $\\gamma$-ray data, \\citet{cg08} identified the source EGR J1856+0235, which is analogous to 3EG J1856+0114 in the 3rd $EGRET$ catalog.  \\palfapsr\\ is well within the 95\\% confidence region of EGR J1856+0235, suggesting the pulsar and/or its PWN are also visible in the MeV-GeV range.  We will investigate this further in a follow-up paper.}. ",
        "conclusions": "\\palfapsr\\ is spatially coincident with the unidentified VHE $\\gamma$-ray source \\hesssrc\\ discovered by \\citet[][see Fig.~\\ref{fig:xray}]{aab+08a}.  Here we show that \\palfapsr\\ is energetically capable of powering \\hesssrc\\ and that this association has similar characteristics to the other pulsar/VHE associations in the literature (Table~\\ref{tab:assoc}).  The estimated distance to \\palfapsr, based on its DM and the NE2001 electron density model of the Galaxy \\citep{cl02}, is $\\sim 9$\\,kpc. The uncertainty on this distance is not well-defined, but in some cases the model can be off by as much as a factor of $\\sim 2-3$. Adopting the DM-distance of 9\\,kpc gives a large spin-down flux of $\\dot{E}/d^2 = 5.7/d^2_9 \\times 10^{34}$\\,ergs s$^{-1}$ kpc$^{-2}$, where $d_9$ is the true distance scaled to 9\\,kpc. \\citet{chh+07} find that statistically 70\\,\\% of pulsars with $\\dot{E}/d^2 \\gtrsim 10^{35}$\\,ergs s$^{-1}$ kpc$^{-2}$ are visible as VHE $\\gamma$-ray sources. \\palfapsr\\ is very close to this limit and may exceed it if its distance is overestimated.  Hence, based on its energetics alone, it is likely to be visible as a VHE $\\gamma$-ray source. \\hesssrc\\ has a photon index $\\Gamma = 2.2$ and $1-10$\\,TeV flux $F_{\\rm VHE} = 1.5 \\times 10^{-11}$\\,ergs cm$^{-2}$ s$^{-1}$ (about 15\\,\\% of the Crab's flux in this energy range); these spectral parameters are similar to those measured for the other HESS sources identified with PWNe.  Given the spin-down luminosity of \\palfapsr, this suggests an efficiency $\\epsilon_{\\gamma} = L_{\\gamma}/\\dot{E} = 3.1d^2_9$\\,\\% ($1-10$\\,TeV), comparable to what is seen in other proposed associations (Table~\\ref{tab:assoc}).  \\palfapsr\\ is offset from the centroid of \\hesssrc, $\\alpha = 18^{\\rm{h}}57^{\\rm{m}}11^{\\rm{s}}$, $\\delta = +02^{\\circ}40^{\\prime}00^{\\prime \\prime}$ (J2000, there is a $3^{\\prime}$ statistical uncertainty on this position), by $8^{\\prime}$ (Fig.~\\ref{fig:xray}).  As discussed in \\S 1, this is most likely explained by asymmetric confinement of the pulsar wind. This interpretation is supported by \\palfapsr's offset location on the side of \\hesssrc\\ that appears somewhat compressed (i.e., there is a steep gradient in the $\\gamma$-rays) compared with the rest of the nebula.  If, however, the offset of the VHE emission is due primarily to the proper motion of \\palfapsr, then the direction and rough magnitude of this motion are predictable.  Based on its characteristic age and offset from the centroid of \\hesssrc, \\palfapsr's proper motion should be roughly 23\\,mas yr$^{-1}$ (transverse velocity $v_{\\rm t} = 970d_9$\\,km s$^{-1}$), to the north-west, assuming that the centroid of \\hesssrc\\ marks the birthplace of the pulsar.  The velocity is very large, but not unprecedented for a pulsar \\citep{cvb+05}.  Of course, the velocity will be smaller if the distance to the source is overestimated, or if the offset is at least partially due to an asymmetrically confined pulsar wind. Detecting this proper motion via timing or interferometry would elucidate this further and may be possible in the coming years, though timing noise and the low flux of the pulsar will make this difficult.  \\palfapsr\\ and \\hesssrc\\ are also coincident with the faint \\asca\\ X-ray source \\ascaxray\\ reported by \\citet[][see Fig.~\\ref{fig:xray}]{smk+01}. \\ascaxray\\ was found $3^{\\prime}$ off axis in observations from 1998 April 6 (sequence number 56003000).  It was detected only in the hard band ($2-10$\\,keV) of the Gas Imaging Spectrometers (GISs), with a count rate of 2.6\\,ks$^{-1}$ GIS$^{-1}$ and a significance of 4.3\\,$\\sigma$.  The exposure was $\\sim 13$\\,ks for each of GIS2 and GIS3.  \\palfapsr\\ and \\ascaxray\\ (J2000 position: $\\alpha = 18^{\\rm{h}}56^{\\rm{m}}50^{\\rm{s}}$, $\\delta = +02^{\\circ}46^{\\prime}$) are spatially coincident to within the $1^{\\prime}$ positional uncertainty of sources in the \\citet{smk+01} catalog.  Although the signal-to-noise ratio of the detection of \\ascaxray\\ is modest, its exact spatial coincidence with a young pulsar of relatively high spin-down flux argues that this source is real. Most of the previously established associations of HESS sources with young pulsars also have known X-ray synchrotron PWNe.  \\ascaxray\\ could be an X-ray PWN created by \\palfapsr.  Using CXCPIMMS\\footnote{Available at http://cxc.harvard.edu/toolkit/pimms.jsp.}, with an assumed absorbed power law spectrum with column density $N_{\\rm H} = 1 \\times 10^{22}$\\,cm$^{-2}$ (roughly the total Galactic contribution in this direction) and photon index $\\Gamma = 2$ (typical for X-ray PWNe), we find that \\ascaxray\\ has an unabsorbed flux ($2-10$\\,keV) of $1.6 \\times 10^{-13}$\\,ergs s$^{-1}$ cm$^{-2}$.  This corresponds to an efficiency for the conversion of spin-down energy into X-rays $\\epsilon_{\\rm X} = L_{\\rm X}/\\dot{E} = 0.03d^2_9$\\,\\% ($2-10$\\,keV) that falls into the observed range for Vela-like pulsars \\citep{pccm02}.  There are two additional nearby sources detected by \\citet{smk+01}: AX J185721+0247 and AXJ185750+0240.  It is possible that these are part of some extended X-ray emission related with \\hesssrc, though deeper X-ray observations are needed to investigate this (see below).  There are four short \\swift\\ observations of the region containing \\palfapsr, including two observations specifically targeting \\ascaxray. The deepest of these, from 2007 March 13 (observation ID 36183002), is a 4.1-ks on-axis exposure with the \\swift\\ X-ray Telescope (XRT). This observation does not show any significant emission at the pulsar position; within a circular extraction region of radius $1^{\\prime}$ around the pulsar, the background subtracted number of counts is $1.1^{+4.0}_{-2.8}$.  Using the aforementioned count rate of \\ascaxray, along with the same assumed spectrum, the predicted $0.2-10$\\,keV count rate for XRT from CXCPIMMS is $\\sim 2$\\,counts/ks. Thus, in 4.1\\,ks, there should have been $\\sim 8$\\,counts, consistent with the 2-$\\sigma$ upper limit derived from the data.  Thus, if \\ascaxray\\ is predominantly a point source, it was only at the limit of detectability in this observation.  If it is predominantly an extended nebula, then it would not have been detectable in such a short exposure. Clearly, future, deeper observations, like those that recently discovered a likely X-ray PWN associated with HESS J1718$-$385/PSR J1718$-$3825 \\citep{hfc+07}, will be needed to determine the nature of this candidate X-ray PWN.  We have been granted a $\\sim 60$-ks {\\it XMM-Newton} observation of \\palfapsr\\ and will present an analysis of those data in a follow-up paper.  At least half of the known HESS/pulsar associations are accompanied by radio emission classified as a PWN, a notable exception being HESS J1825$-$137/PSR B1823$-$13. Only a few of the extended HESS sources, e.g. HESS J0835$-$455, HESS J1813$-$178, and HESS J1640$-$465, are known to be accompanied by a supernova remnant (SNR). We have checked available radio imaging data for signs of a PWN or SNR. There is some faint, extended emission in the vicinity of \\palfapsr\\ visible in 1.4-GHz VLA Galactic Plane Survey data \\citep[VGPS,][]{std+06}, though nothing that is clearly indicative of a PWN or SNR.  The surface brightness limit from VGPS is $\\sim 1.8 \\times 10^{5}$\\,Jy sr$^{-1}$ at 1.4\\,GHz.  It is certainly not uncommon for Vela-like pulsars to have faint or no known radio nebula \\citep[e.g. PSR B1823$-$13,][]{bgl89,gsk+03}.  Roughly a third of the cataloged Galactic SNRs \\citep{gre04} are fainter than the surface brightness limit achieved by the VGPS. Deep, dedicated radio imaging observations of \\palfapsr\\ are necessary to investigate this further. Higher resolution 1.4-GHz data from the MAGPIS survey (Helfand et al. 2006\\nocite{hbw+06}; in the vicinity of \\palfapsr, this survey has a sensitivity of 0.2\\,mJy/beam at an angular resolution of $6^{\\prime \\prime}$) reveal no point source which can be definitively associated with \\palfapsr, as expected given the flux density and positional uncertainty of the pulsar."
    },
    "0806/0806.2459_arXiv.txt": {
        "abstract": "Recently, the MILAGRO collaboration reported on the detection of a diffuse multi-TeV emission from a region of the Galactic disk close to the inner Galaxy. The emission is in excess of what is predicted by conventional models for cosmic ray propagation, which are tuned to reproduce the spectrum of cosmic rays observed locally. By assuming that the excess detected by MILAGRO is of hadronic origin and that it is representative for the whole inner Galactic region, we estimate the expected diffuse flux of neutrinos from a region of the Galactic disk with coordinates $-40^{\\circ} < l < 40^{\\circ}$. Our estimate has to be considered as the maximal expected neutrino flux compatible with all the available gamma ray data, since any leptonic contribution to the observed gamma-ray emission would lower the neutrino flux. The diffuse flux of neutrinos, if close to the maximum allowed level, may be detected by a km$^3$--scale detector located in the northern hemisphere. A detection would unambiguously reveal the hadronic origin of the diffuse gamma-ray emission. ",
        "introduction": "\\label{intro}  Recently, the MILAGRO collaboration reported on the detection of multi-TeV diffuse gamma-ray emission from two regions in the Galactic plane: the Cygnus region (located at galactic longitude $65^{\\circ} < l < 85^{\\circ}$) \\cite{MILAGROcygnus} and the portion of the inner Galaxy visible from the location of MILAGRO ($30^{\\circ} < l < 65^{\\circ}$) \\cite{MILAGROinner}. When placed alongside the earlier lower energy measurements by EGRET \\cite{egretdiffuse}, such detections provide an interesting insight into both the spatial distributions of cosmic rays (CR) % in the Galaxy and their energy spectrum.  The diffuse GeV emission detected by EGRET along the galactic plane is believed to be the result of hadronic interactions between CRs and interstellar matter. A strong point in favor of this hypothesis is the fact that the diffuse GeV emission correlates with the spatial distribution of the interstellar gas that constitutes the target for CR interactions \\cite{egretdiffuse}. The same argument might be extended also to the diffuse multi-TeV emission detected by MILAGRO, since both the Cygnus region and the inner Galaxy are characterised by an enhancement in the gas density. In the case of the Cygnus region, the morphology of the multi-TeV emission correlates with the CO emission, which traces the gas density. This suggests again a hadronic origin also for the very high energy diffuse emission \\cite{MILAGROcygnus}. Similarly, for the region close to the inner Galaxy, the rather narrow latitude profile of the multi TeV-emission, concentrated around the dense Galactic disk, might suggest a hadronic interpretation of the gamma-ray data, though also a leptonic one seems feasible \\cite{aa2000,galprophardelectrons,MILAGROinner}.  The expected hadronic gamma-ray emissivity from the Galactic plane can be calculated by solving the transport equation describing the propagation of CRs in the Galaxy and using the measured distributions of the gas density in the Galactic disk (see \\cite{strongreview} for a review). Conventional propagation models are tuned in order to match the CR spectrum observed locally \\cite{galprop,dario,candia}. These models reproduce quite fairly the observed level of the GeV emission (with the notable exception of the so called GeV-excess) \\cite{galpropegret} but fail to reproduce the MILAGRO observations by about one order of magnitude in the case of the Cygnus region \\cite{MILAGROcygnus} and a factor of $\\sim 5$ for the inner Galaxy \\cite{MILAGROinner}. One possible solution of the problem is to assume that the average CR spectrum in the Galaxy is harder than the one we measure at Earth \\cite{wolfendale,aa2000}. % Another possibility is to assume a hard spectrum for the CR electrons \\cite{aa2000,galprophardelectrons} or that the typical CR electron intensity in the Galaxy is significantly higher than that measured locally \\cite{galpropegret}. This would make the inverse Compton contribution to the total gamma-ray emission dominate over the hadronic one at the multi-TeV energies relevant here. Remarkably, both these scenarios (hadronic and leptonic) were proposed to solve the problem of the GeV--excess and can now be generalised to investigate their predictive power at multi-TeV energies.    Thus, the origin of the multi-TeV emission from the central Galactic plane region remains unclear. Either leptonic or hadronic processes are believed to be the dominant contributors to the detected gamma-ray flux, with the possibility of a transition from a dominant hadronic contribution to a leptonic one in the GeV-TeV energy range \\cite{galpropegret}. In the case of the hadronic scenario, one would expect, as it is observed, a narrow extension in latitude of both the EGRET (see Fig.~3 from \\cite{egretdiffuse}) and MILAGRO emission \\cite{MILAGROinner}. In this case the extension of both GeV and multi-TeV emission is determined only by the gas distribution in the Galactic disk, which constitutes the target for proton--proton interactions. In the leptonic scenario, the multi-TeV emission is produced via inverse Compton scattering by $\\approx 100$~TeV electrons which can propagate over a distance of $\\approx 100$~pc before being cooled by synchrotron and inverse Compton losses \\cite{aa2000}. Thus, also in this case a narrow latitude distribution of gamma rays is expected, if the sources of cosmic ray electrons are concentrated around the galactic plane. In this case, however, the extension of the gamma-ray emission would be energy dependent, since higher energy electrons are cooled faster and can propagate shorter distances \\cite{aa2000}. Thus, an accurate comparison of the extension in latitude of the emission detected at GeV and multi-TeV energies is of crucial importance, and future observations with improved angular resolution by GLAST and HAWK might help in discriminating between the two scenarios.  Since the production of gamma rays is accompanied by a corresponding neutrino flux only in the hadronic scenario, neutrino telescopes provide a unique opportunity to disentangle the origin of the multi-TeV radiation. In this paper, we investigate the capability of km$^3$--scale neutrino detectors such as IceCube \\cite{icecube} and KM3NeT \\cite{km3net}, to detect the diffuse flux of neutrinos from the inner Galaxy and possibly reveal the origin of the diffuse emission detected by MILAGRO.   In the past, several estimates of the diffuse neutrino flux from the Galactic plane region have been published. The simplest of these approaches assumes that the CR distribution is constant across the Galactic disk \\cite{stecker,berezinsky,ingelman}, whilst more sophisticated calculations assume a distribution of CR sources in the disk and obtain a steady state CR distribution by solving the diffusion equation \\cite{dario,candia}. All of these approaches, however, result in relatively small neutrino fluxes and are below that required to be detectable by a km$^{3}$-size neutrino detector such as IceCube or KM3NeT, within one year of data acquisition. In these calculations, a CR spectral index of $\\sim 2.7$ is generally assumed, or derived by solving the CR transport equation, since this matches the observed CR spectrum measured at Earth. Such a steep spectrum, normalized to reproduce the diffuse emission observed at GeV energies by EGRET, naturally results in a conservative estimate of the diffuse neutrino flux at multi-TeV energies where km$^{3}$ neutrino telescopes are optimally sensitive. However, no a priori reason exists that such a CR spectrum is applicable throughout the Galaxy. Indications exist for a departure from the observed local CR spectral index of $\\sim E^{-2.7}$ in the diffuse gamma-ray emission observed by H.E.S.S. from the Galactic centre ridge \\cite{H.E.S.S.ridge} and the GeV excess in the diffuse flux from the Galaxy observed in EGRET data, which might be due to a CR spectrum harder than the one observed at Earth \\cite{wolfendale,aa2000}. %  In this work we adopt a phenomenological approach. % We assume that the production of the multi-TeV diffuse flux observed by MILAGRO is entirely due to hadronic interactions. This model requires the assumption of a CR spectrum in the inner part of the Galaxy that is different (harder) than the canonical 2.7 spectral index of the CRs observed at Earth. Consequently we obtain a higher value for the multi-TeV neutrino flux from the same Galactic plane region compared to previous canonical estimates. Any leptonic contribution to the multi-TeV diffuse emission will lower the actual diffuse neutrino flux. Our calculations therefore provide an estimate of the maximal neutrino flux that can be expected, under the constraint of remaining compatible with all the presently available gamma-ray data. Since the MILAGRO detection of the diffuse multi-TeV gamma-rays from the inner Galaxy only covers the angular range $30^{\\circ}<l<65^{\\circ}$ \\cite{MILAGROinner}, to apply this MILAGRO data point to the whole inner Galaxy with coordinates $-40^{\\circ} < l < 40^{\\circ}$ it was assumed that the TeV flux detected by MILAGRO scales in a similar way to the GeV flux detected by EGRET from the two regions.  It is worth noticing that MILAGRO is operating at multi-TeV energies which is also the energy range where km$^3$-scale neutrino telescopes reach their best performances (see \\cite{icecube_perf,km3net_perf} and Sec.~\\ref{sec:telescopes}). This implies that the MILAGRO observations are the most relevant gamma-ray observations to constrain the expected diffuse neutrino flux, since gamma rays and neutrinos are produced roughly at the same energy during CR hadronic interactions with the interstellar medium.  In Sec.~\\ref{sec:gamma} we review the relevant gamma-ray observations that can constrain the expected neutrino flux from the inner Galaxy. We suggest different possibilities to fit the gamma-ray data in scenarios where the emission is dominated by hadronic CR interactions. In Sec.~\\ref{sec:neutrinos} we use such fits to evaluate the expected flux of neutrinos from the inner Galaxy. This has to be considered as the maximal neutrino flux which is compatible to all the available gamma-ray measurements. Any leptonic contribution to the observed gamma-ray emission would lower the expectations. Our conclusions are provided in Sec.~\\ref{sec:concl}. The maximum level of the diffuse neutrino flux allowed by constraints from gamma-ray observations is within the detection capabilities of a km$^3$--scale neutrino telescope located in the northern hemisphere. A detection would unambiguously reveal the hadronic origin of the multi-TeV emission detected by MILAGRO. ",
        "conclusions": "\\label{sec:concl}  In this paper we presented an estimate of the neutrino flux from the inner Galaxy, defined as the region of coordinates $-40^{\\circ} < l < 40^{\\circ}$ and $-2^{\\circ} < b < 2^{\\circ}$. Our calculations rely on the following assumptions: \\begin{enumerate} \\item{the diffuse gamma-ray emission detected at multi-TeV energies by MILAGRO from the region of the Galactic plane of coordinates $30^{\\circ}<l<65^{\\circ}$ is representative for the whole inner Galaxy;} \\item{the diffuse galactic GeV and multi-TeV emission detected by EGRET and MILAGRO respectively follows the same spatial distribution;} \\item{the diffuse gamma-ray emission is entirely of hadronic origin.} \\end{enumerate} These assumptions allow to obtain an almost model independent estimate for the diffuse flux of neutrinos from the inner Galaxy. Such an estimate has to be considered as the maximal expected flux of neutrinos compatible with gamma ray observations, since any leptonic contribution to the observed gamma ray emission would lower our expectations.  The maximal expected neutrino flux is within the detection capabilities of a km$^3$-scale neutrino telescope located in the northern hemisphere. For the most favorable assumptions made in Sec.~\\ref{sec:neutrinos}, we expect from the inner Galactic region 15 neutrinos above 10 TeV per year of exposure against a background of 28 atmospheric neutrinos. At 20 TeV the signal equates the background with about 10 events per year. The detection of these neutrinos would constitute an unambiguous proof of the hadronic origin of the gamma-ray emission detected by MILAGRO. Finally, it is important to stress that MILAGRO is operating in an energy range similar to the one where neutrino telescopes reach their best performances. This makes MILAGRO and future instruments like HAWK, the ideal gamma-ray detectors for multi--messenger studies of sources of astrophysical neutrinos."
    },
    "0806/0806.3755_arXiv.txt": {
        "abstract": "Blazars with strong emission lines were found to be  associated mostly with  broad-line type 1 Active Galactic Nuclei (AGN). Hitherto, evidence for blazars identified with Narrow Line Seyfert\\,1 (NLS1) AGN was limited to very few individual cases. Here we present a comprehensive study of a sample of 23 genuine radio-loud NLS1 galaxies which have the radio-loudness parameters, the ratio of radio (21\\,cm) to optical (4400\\AA) luminosity, greater than 100. The sample, drawn from the SDSS and FIRST, is homogeneous and the largest of this kind. A significant fraction of the sample objects show interesting radio to X-ray properties that are unusual to most of the previously known radio-loud NLS1 AGN, but are reminiscent of blazars. These include flat radio spectra, large amplitude flux and spectral variability, compact VLBI cores, very high brightness temperatures ($10^{11-14}$\\,K) derived from variability, enhanced optical emission in excess of the normal ionising continuum, flat X-ray spectra, and blazar-like SEDs. We interpret them as evidence for the postulated blazar nature of these very radio-loud NLS1 AGN, which might possess at least moderately relativistic jets. We suggest that those steep spectrum radio-loud NLS1 AGN in the sample are of the same population but with their radio jets aligned at large angles to the lines-of-sight. Intrinsically, some of the objects have relatively low radio power and would have been classified as radio-intermediate AGN.  The black hole masses, estimated from the broad Balmer line width and luminosity, are within $10^{6-8}$\\,\\msune, and the inferred Eddington ratios are around unity. Unless the black hole masses are largely under-estimated, our result stretches the low mass end of the black holes of luminous, fast accreting radio-loud AGN to a smaller mass regime (the order of  $10^6$\\,\\msune) in the black hole mass--radio-loudness space where other normal AGN are seldom found. The results imply that radio-loud AGN may be powered by black holes with moderate masses ($\\sim10^{6-7}$\\,\\msune) accreting at high rates (Eddington ratios up to unity or higher). The host galaxies of a few nearby objects appear to be disk-like or merger;  and some of the objects show imprints of young stellar populations in their SDSS spectra. We find that some of the objects, despite having strong emission lines, resemble high-energy peaked BL Lacs in their SED with the synchrotron component peaked at around the UV; such objects constitute an intriguingly high fraction of the sample. The radio sources of the sample are ubiquitously compact. They are smaller than  at most several tens of kilo-parsecs, suggesting a possible link with compact steep-spectrum radio sources. Given the peculiarities of blazar-like NLS1 galaxies, questions arise as to whether they are plain downsizing extensions of normal radio-loud AGN, or whether they form a previously unrecognised population. ",
        "introduction": "As a minority of AGN, radio-loud (RL) AGN differ from their radio-quiet (RQ) counterparts mainly in possessing prominent radio jets and/or lobes that produce strong radio radiation. In observations, a widely used division between RL and RQ AGN is the radio-loudness parameter $R\\approx$10, defined as the ratio between the radio 5\\,GHz to optical $B$-band luminosity \\citep[][]{kel89,sto92}, though it is still a controversy whether $R$ has a bimodal distribution below and above this value. RL AGN are an important laboratory to study the formation of relativistic jets, which is not yet understood so far \\citep[e.g.][]{bla00,cel01,mei03}. Nevertheless, it is widely thought that radio-loudness (formation of radio jets) is possibly depending on the accretion rate/state \\citep[e.g.][]{ho02,mac03,gre06,kor06}, the spin of the black hole \\citep[e.g.][]{bla00,mar03,sik07}, black hole mass \\citep[\\mbhe, e.g.][but see Woo \\& Urry 2002 for a different view]{laor00,lacy01,dun03,mcl04,metca}, and host galaxy morphology \\citep[e.g.][]{balm06,cap06,sik07}, or a combination of some of these factors \\citep[e.g.][]{ball07}.  As an important sub-class of RL AGN, blazar is a collective term for BL Lac objects and flat-spectrum radio quasars (FSRQs). Blazars are characterised by flat radio spectra at above $\\sim$1\\,GHz, fast variability, high and variable polarization, superluminal motion, and high brightness temperatures \\citep[see e.g.][for reviews]{urry95,cel02}. They are now believed to be RL AGN with the orientation of relativistic jets close to the line-of-sight, and hence their  non-thermal jet emission is highly Doppler boosted \\citep{bla78,urry95}. Blazars have distinctive spectral energy distributions (SED) that are characterised by two broad humps in the $\\log\\nu-\\log\\nu f_{\\nu}$ representation. The low-energy hump is commonly interpreted as synchrotron emission, while the high-energy one as inverse Compton emission scattering off the same electron population which produces the synchrotron emission. Depending on the peak frequency of the synchrotron hump, BL Lac objects are divided into Low-energy-peaked BL Lacs (LBL) with the peak around IR--optical wavelengths and High-energy-peaked BL Lacs (HBL) with the peak around UV/soft X-ray energies \\citep{pado95}. Classical FSRQs have generally the peak frequencies of the synchrotron hump similar to or even lower than LBL. Interestingly, it was suggested that the whole blazar family can be arranged in a sequence from HBL to LBL and to FSRQ in the order of decreasing synchrotron peak frequencies and increasing source power---the so called blazar sequence\\footnote{ The blazar sequence can be explained as, for instance, the decrease of the energy of electrons emitting at the SED peaks with increasing energy density of the seed photons for Compton scattering \\citep{ghi98} } \\citep[][]{fos98}. However,  exceptions were also reported in recent years as the presence of powerful FSRQs with HBL-like SEDs (though this remains a matter of controversy), whose synchrotron hump is peaked around the UV/soft X-ray band \\citep[see e.g.][and references therein]{pado03}. Following \\citet{per98}, we term FSRQs with HBL-like SED HFSRQs (High-energy-peaked FSRQs), and those with LBL-like SED---the classical FSRQs---LFSRQs (Low-energy-peaked FSRQs), as parallels to HBLs and LBLs.  Until recently, RL AGN (blazars) with strong broad emission lines were found to ubiquitously have the Balmer line widths greater than 2000\\,\\kmps in full width at half maximum (FWHM). It was found that the line width is correlated with the orientation of the radio axis, with narrower lines as well as stronger optical \\feii emission in objects with smaller viewing angles to the radio axis \\citep[e.g.][]{wills86,jack91}. This was interpreted as arising from a disk-like emission line region. A marked absence of radio-loud quasars (RLQs) was noted below the 2000\\,\\kmps line-width cutoff \\citep{wills86}, that  is the characteristic line-width range for narrow line Seyfert\\,1 galaxies---a sub-class of type\\,1 AGN in contrast to the classical Broad Line AGN (BLAGN). The situation has changed since the detection of radio emission from NLS1 galaxies.  Apart from narrow line-widths of the broad Balmer lines \\citep[originally defined as $<2000$\\,\\kmpse,][]{ost85,goo89}, NLS1 galaxies also show other extreme properties compared to normal broad line AGN \\citep[see][for a recent review]{kom08}, such as strong permitted optical/UV \\feii emission lines \\citep{bg92,gru99,veron01}, steep soft X-ray spectra \\citep{wang96,bol96,gru98}, and rapid X-ray variability \\citep{lei99,kom00}. Observational evidence suggests that NLS1 galaxies tend to have small black hole masses and high Eddington ratios (defined as the bolometric to the Eddington luminosity ratio), \\redd$\\equiv L_{\\rm bol}/L_{\\rm Edd}\\approx 1$ \\citep[e.g.][]{bor02,col04}. In fact, they were found to locate at one extreme end, opposite to BLAGN, of eigenvector\\,1 of the correlation matrix which is believed to be driven primarily by \\redd \\citep{bg92,sul00}. NLS1 galaxies  were once thought to be radio-quiet. This was simply a consequence of the fact that NLS1 galaxies have a low radio-loud fraction \\citep{zhou06,kom06b} and the small number of NLS1 galaxies known till then.  Previous studies of the radio properties of NLS1 galaxies are all based on small samples. These showed that the radio sources, if detected, are compact (less than a few hundred parsec), at modest radio power \\citep{ulv95}, of steep radio spectra \\citep{moran00}, mostly in the RQ regime \\citep{ste03}, and of low RL fraction \\citep{zhou02}. Until recently, there were only several RL NLS1 galaxies identified and studied individually \\citep[see][for a review and references therein]{kom06b}. Studies using (small) samples of RL NLS1 galaxies were carried out by \\citet{kom06b} and \\citet{wha06} for non-radio selected and radio selected objects, respectively. These studies confirmed that, as in `normal' NLS1 galaxies (their RQ counterparts), RL NLS1 AGN are accreting at a high rate close to the Eddington limit. More importantly, as pointed out by  \\citet{kom06b}, with relatively low \\mbh and high $R$, they occupy a previously sparsely populated region in the \\mbhe--$R$ diagram. Most of the sources of that sample are compact steep spectrum sources. It was confirmed that RL NLS1 galaxies are rare ($\\sim7$\\%), and very radio-loud objects with $R>100$ are even more sparse \\citep[see also][]{zhou02,zhou06}, compared to about 10--15\\% for normal BLAGN and quasars \\citep[e.g.][]{ive02}. Several of these studies discussed the starburst versus AGN contribution to the radio emission, and concluded that it is AGN dominated.  Of particular interest, a few outstanding RL NLS1 galaxies came to be known to exhibit blazar properties. RXJ\\,16290+4007 \\citep{sch00,gru04} was known as a blazar and its X-ray emission is dominated by synchrotron emission \\citep{pado02}, or else by soft X-ray emission that is typical for NLS1 galaxies, perhaps from the accretion disk \\citep{kom06b}. The SED of PKS\\,2004-447\\footnote{However, the NLS1 nature of PKS\\,2004-447 is not certain; see  \\citet[][]{zhou03} for a short comment.} \\citep{osh01} was well modeled with the blazar type \\citep{gal06}. Using information on the radio flux variations, very high radio brightness temperatures ($\\gtrsim10^{13}$\\,K) were inferred for J0948+0022 \\citep{zhou03} and 0846+513\\footnote{A formerly known blazar; see Appendix\\,\\ref{sect:indv} for details.} \\citep{zhou05}---both included in our sample, arguing for relativistic beaming \\citep[see also][]{wang06}. They also show some other blazar-like behavior. Recently, \\citet{doi06,doi07} performed high-resolution VLBI observations for several radio-loud NLS1 AGN (three are included in our sample\\footnote{J0948+0022, J1633+4718, and J1644+2619.}) and found that they are unresolved with milli-arcsec resolutions, setting direct lower limits on the brightness temperatures in the range of $10^7-10^9$\\,K. They also found that inverted radio spectra are common in the radio-loudest objects. The authors suggested that Doppler beaming, presumably resulting from highly relativistic jets, can explain naturally the observations. In this respect, the most remarkable object of this kind is perhaps 2MASX\\,J0324+3410, found in our recent work \\citep{zhou07}. It showed rapid variability in the radio, optical, and X-ray bands, and was even claimed to be marginally detected in TeV $\\gamma$-rays \\citep{fal04}. Its non-thermal SED clearly resembles that of HBLs. Moreover, 2MASX\\,J0324+3410 is hosted by a relatively small, apparent disk galaxy with one-armed spiral or ring galaxy morphology. These enigmatic objects present a challenge to current models of both NLS1 galaxies and  blazars, and render a unique opportunity to  study jet formation in black hole systems accreting at high rates. Apparently, they do not seem to belong to any types of AGN currently known. An immediate question is whether they are some exceptional individuals or there exists a population of such objects. What are their general properties in continuum and line emission, central engine and host galaxies? What are their relations with other types of AGN?  Motivated by these issues, we compiled a  RL NLS1 galaxy sample from the Sloan Digital Sky Survey (SDSS) database, aimed at a systematic study with a large and homogeneous sample. Since, as the first step, we focus on genuine RL objects, we consider in this paper objects with the radio-loudness parameter $R>100$ (see \\S\\,\\ref{sect:samplesel} for the definition of $R$) for the following reasons. Firstly, the classification of AGN into RL and RQ is somewhat ambiguous in the range $R=10-50$ due to the presence of (a small number of) so called radio-intermediate AGN spanning in between the bulks of the two classes, which may have a different origin from genuine RL objects \\citep[e.g.][]{fal96a}. The same may also be true for NLS1 galaxies. Secondly, in $R>100$ objects contamination of radio emission from the host galaxy is negligible (see \\S\\,\\ref{sect:disc_sed}). Thirdly, their jet component, if present, can be detected and studied relatively easily in other wavebands than  the radio, such as optical and X-rays. Last, but not the least, they are the least studied objects {\\em as a sample}, due to their extreme rarity, in contrast to less RL NLS1 galaxies as the bulk in the samples of \\citet{kom06b} and \\citet{wha06}. We defer to a later paper a comprehensive and complete treatment of the SDSS--FIRST detected  NLS1 galaxies sample, including radio-intermediate objects, to address questions such as the RL--RQ dichotomy for NLS1 galaxies. Our results show that there exists a population of NLS1 galaxies that are blazar-like, despite their  rarity.  The compilation of the sample and the analyses of multi-waveband data are described in \\S\\,\\ref{sect:sample}. The broad band continuum and emission line properties are presented in \\S\\,\\ref{sect:cont} and \\S\\,\\ref{sect:emline}, respectively. In \\S\\,\\ref{sect:hostgal} the properties of their host galaxies are presented. We discuss other properties such as black hole mass and accretion rate, and isotropic radio emission, as well as the implications of the results in \\S\\,\\ref{sect:discus}. Throughout the paper, we assume a cosmology with $H_{0}$= 70 km\\, s$^{-1}$\\,Mpc$^{-1}$, $\\Omega_{M}=0.3$, and $\\Omega_{\\Lambda}=0.7$. Errors are quoted at the 68\\% confidence level unless mentioned otherwise. We use the following conventions to denote power-law spectral indices: \\av\\ in the frequency domain [$S(\\nu) \\propto \\nu^{\\alpha_{\\nu}}$, specifically \\alpre$\\equiv$\\av\\ in the radio band], and \\alpwvl in the wavelength domain [$S(\\lambda) \\propto \\lambda^{\\alpha_{\\lambda}}$]; and the X-ray photon index $\\Gamma$, defined as $f_{\\rm pho}(E) \\propto E^{-\\Gamma}$, where $f_{\\rm pho}(E)$ is X-ray photon number flux density. ",
        "conclusions": "A homogeneous sample of genuine RL NLS1 galaxies (radio-loudness \\rle$\\ga 100$) is presented for the first time, comprising 23 objects drawn from SDSS spectroscopic objects with FIRST detections. Their NLS1 nature is discovered by analysing systematically the spectroscopic data from the QSO and galaxy database in the SDSS DR5. The sample can be treated as an optically--radio selected one. We carry out  systematic investigations into their multiwaveband properties from radio to X-rays using data available from the archives. The main results are the following.  All radio sources of the sample are compact, less than a few tens of kilo-parsec at least. Several with VLBI observations are unresolved or marginally resolved at scales of about tens of parsec. Among those with available radio spectral indices, the majority are flat-spectrum sources ($\\alpha>-0.5$ around 1.4--5\\,GHz), and a few even have inverted spectra ($\\alpha>0)$. Radio variability is found to be common in most of the flat-spectrum sources. We were able to estimate the brightness temperatures of the radio sources, that are as high as $T_{\\rm B}>10^{11}$\\,K and even exceeding the inverse Compton limit ($10^{12}$\\,K) in four objects. The high brightness temperatures indicate the presence of relativistic beaming. The majority of RL NLS1 galaxies has a low-frequency (isotropic) radio power falling in the range of FR\\,I type radio galaxies (\\llf$\\la 10^{25}$\\,\\whzsre), and the remaining minority in the range of the FR\\,II type. When isotropic, rather than Doppler boosted, emission is considered, a large fraction of the objects can be considered as intrinsically radio-intermediate.  The observed optical luminosities are systematically in excess of the  thermal continua predicted from the \\hb line luminosities by a factor of up to several. The optical slopes of our RL objects, while spanning a wide range from moderately red to blue, are  systematically bluer than that of radio-quiet NLS1 AGN. We interpret this effect as the contribution from a (non-thermal) component, presumably from relativistic jets, to the optical band superimposing the central ionising continuum seen in radio-quiet NLS1 AGN. The optical \\feii emission in the radio-loud NLS1 in our sample is on average stronger than `normal' NLS1 galaxies. Except this, the two subsets of NLS1 galaxies seem to be indistinguishable in other emission line properties.  On the  radio--optical--X-ray effective spectral indices plane (\\alpox vs.\\ \\alproe), our objects have broad-band SEDs well consistent with those of blazars, i.e.\\ both the HFSRQ and LFSRQ types as well as SSRQs. Of particular interest, the fraction of the candidate HFSRQ-type objects seems to be high, though their synchrotron component peaks are only moderately high (at around the UV). This finding extends the recently discovered strongly-lined (FSRQ) HBL-type blazars \\citep{pado03} to NLS1 galaxies.  The continuum shape in the soft X-ray band exhibits a large variety among our RL NLS1 galaxies. There is evidence in one, or possibly two objects for a flat X-ray spectral component ($\\Gamma\\la 1.8$), which we interpret as an inverse-Compton origin from jets as in FSRQs. Most of the objects with estimation of the X-ray spectral slopes show steep spectra. The steep X-ray spectrum is also consistent with what is expected for HFSRQs, that originates from the high energy end of the synchrotron component. Based on this circumstantial evidence, we suggest that such a component may be a significant contributor to the observed X-ray emission, though the X-ray component as seen in radio-quiet NLS1 galaxies may also play a role.  In view of the independent lines of evidence presented here, as well as those found in previous studies for several individuals (either inclusive or exclusive of our sample), we  suggest that there likely exists a population of RL NLS1 galaxies showing the broad-band continuum emission characteristic of blazars, though alternative explanations cannot be ruled out completely. In this scenario,  beamed non-thermal emission from jets, that are at least mildly relativistic, dominates the high-frequency radio band, and contributes significantly to the optical and possibly the X-ray bands. Our finding extends the radio-loud quasar population into a narrow line-width regime of the broad emission lines. We postulate that all genuine radio-loud NLS1 galaxies possess small scale, at least mildly relativistic jets, reminiscent of classical radio-loud AGN. In this scenario, steep-spectrum NLS1 sources (perhaps including some of those with smaller radio-loudness within $R=$10--100) are simply objects of the same population whose radio jets are aligned at a relatively large angle to the line-of-sight.  The host galaxies of two nearby RL NLS1 nuclei, whose images are resolved, have disk morphology, one even in a merger, in line with the finding for the previously published blazar-NLS1 hybrid object 2MASX\\,J0324+3410 \\citep{zhou07}. In several objects imprints of young stellar populations are likely to be present in form of high-order Balmer absorption lines in the SDSS spectra.  Based on the commonly used relation of black hole mass with broad line width and luminosity, our RL NLS1 galaxies have estimated \\mbh in the range of $10^6-10^8$\\,\\msune, provided that this \\mbh scaling relation is applicable to NLS1 and the BLR is not largely flattened. The thus estimated Eddington ratios range from \\redd=0.5 to 3, and peak at \\redde$\\approx 1$ (assuming $L_{\\rm bol}=9 \\lambda L_{\\lambda5100}$). RL NLS1 galaxies deviate from the known trends of radio-loudness versus the Eddington ratio and black hole mass, and are populated in a region in the radio-loudness--\\mbh diagram where normal AGN are seldom found, i.e.\\ moderate \\mbh and large $R$ (\\mbh$\\simeq 10^{6-7.5}$\\,\\msune,  $R\\simeq 10^{1-3}$), confirming and improving upon the previous results suggested by \\citet{kom06b}. As such, RL NLS1 show strong excess radio emission with respect to their \\mbhe, when compared to the previously suggested radio power--\\mbh relation  for normal AGN \\citep[e.g.][]{lacy01}. Our results imply that large radio-loudness ($R=10^{1-3}$) can be produced in AGN with  moderate \\mbh down to $\\sim10^6-10^7$\\,\\msun and accreting at/close to the Eddington rate; and some of these AGN are hosted in late-type galaxies, some even in interacting systems. The presence of relativistic jets in such systems may have interesting implications for understanding  jet formation in AGN in general.  In spite of a likeness  to normal RL AGN in many aspects, RL NLS1 galaxies show some extreme or unusual   properties, though some of the evidence are only circumstantial. These include compact radio morphology, relatively low radio power, a high incidence of HBL/HFSRQ-type SEDs, host galaxies of probably late-type and/or with young stellar populations, and the common properties characteristic of NLS1 galaxies, e.g.\\  smaller black hole masses, higher Eddington ratios, and stronger \\feii emission. All of these lead us to postulate that either they are an extension of normal RL AGN to the lower-\\mbh range, and all those differences can be explained as a consequence of decreasing \\mbhe; or they form a previously unknown population distinct from normal radio-loud AGN. As one interesting possibility, we speculate that low-power CSS  might be the parent population of the flat-spectrum RL NLS1 galaxies with beamed jet emission. For the future, a much larger sample and extensive multi-waveband observations are essential for understanding the nature of radio-loud NLS1 galaxies.  Finally, if the broad-band radiation mechanisms in flat-spectrum radio-loud NLS1 AGN are indeed the same as in blazars, we expect some of these objects to be $\\gamma$-ray emitters in the GeV band. This can be tested by observations of GLAST in the near future."
    },
    "0806/0806.4944_arXiv.txt": {
        "abstract": "Up to now, most stellar-mass black holes were discovered in X-ray emitting binaries, in which the black holes are formed through a common-envelope evolution. Here we give evidence for the presence of a massive black hole candidate as a tertiary companion in the massive eclipsing binary V Puppis. We found that the orbital period of this short-period binary (P=1.45 days) shows a periodic variation while it undergoes a long-term increase. The cyclic period oscillation can be interpreted by the light-travel time effect via the presence of a third body with a mass no less than 10.4 solar mass. However, no spectral lines of the third body were discovered indicating that it is a massive black hole candidate. The black hole candidate may correspond to the weak X-ray source close to V Puppis discovered by Uhuru, Copernicus, and ROSAT satellites produced by accreting materials from the massive binary via stellar wind. The circumstellar matter with many heavy elements around this binary may be formed by the supernova explosion of the progenitor of the massive black hole. All of the observations suggest that a massive black hole is orbiting the massive close binary V Puppis with a period of 5.47 years. Meanwhile, we found the central close binary is undergoing slow mass transfer from the secondary to the primary star on a nuclear time scale of the secondary component, revealing that the system has passed through a rapid mass-transfer stage. ",
        "introduction": "With an orbital period of 1.4545 days and two B-type components stars (spectral types B1 and B3), V Puppis (HR\\,3129 and HD 65818) is one of a few massive semi-detached eclipsing binaries where the less massive components are filling their Roche lobes (e.g., Andersen et al. (1983); Bell et al. (1987a, b); Terrell et al. (2005)). According to the theory of binary evolution (e.g., Sybesma 1986), theses systems were formed through a case A mass-transfer (mass transfer is occurring when hydrogen is burning in the core). Since the discovery of V Puppis more than one century ago, it was extensively investigated photometrically and spectroscopically (e.g. Andersen et al. (1983) and references therein). A weak X-ray source in the field of V Puppis (e.g., Giacconi et al. (1974); Groote et al. 1978); Bahcall et al. (1975)) and a significant amount of circumstellar matter around the binary (e.g., York et al. 1976 and Koch et al. 1981) make it a very interesting binary to study.  Because of its stable light curve and the deep and sharp eclipsing minima (e.g., Andersen et al. (1983)), times of light minimum of V Puppis can be determined precisely. Thus, the variations of the orbital period can in principle be derived with high precision by analyzing the observed-calculated (O-C) diagram. In the present letter, orbital period changes of V Puppis are investigated. Then, based on the period variations, the presence of a massive black hole companion and the evolutionary state of the system are discussed. ",
        "conclusions": "The cyclic change of the orbital period, the presence of a weak X-ray source near V Puppis, and the significant amount of circumstellar matter around the binary star all support the conclusion that a massive black hole is orbiting the massive close binary V Puppis with a period of 5.47 years. Several massive stellar black hole candidates have been discovered recently (e.g., Greiner et al. (2001) and Orosz (2007)), but all candidates are found to be a component in a binary system. We conclude here that we have provided observational evidence for the first massive stellar black hole candidate as a tertiary companion of a massive close binary star. Theoretical investigations have shown that the formation of black hole binaries should experience a common-envelope phase with the spiral-in of the companion in the envelope causing the ejection of the envelope (e.g., Taam \\& Sandquist (2000) and Podsiadlowski et al. (2002)). The mean distance between the black hole candidate and the eclipsing pair is about 5.5 astronomical units (AU). During the supergiant phase, the radius of the precursor of the black hole companion in V Puppis may reach to several AUs. The survival of V Puppis in this evolutionary stage suggests that the tertiary black hole companion must be formed through an evolutionary process that is different from that of the black holes in X-ray binaries (e.g., Woosley \\& Weaver (1995) and Zhang et al. (2007)) and it could give some constraints to the evolution of single massive stars and the formation of isolated black holes.  From the present investigation, we deduce that the progenitor of V Puppis is a triple system that is composed of a short-period detached B-type eclipsing binary and a much more massive tertiary companion. By drawing angular momentum from the inner binary through Kozai oscillation (Kozai 1962) or a combination of Kozai cycle and tidal friction (e.g., Fabrycky \\& Tremaine 2007), the massive tertiary companion may play an important role for the origin the central system, which causes the inner eclipsing binary to have a very short initial orbital period. In the triple system, the massive tertiary companion evolves faster, and finally into a massive black hole via a supernova explosion (e.g., Woosley \\& Weaver (1995) and Zhang et al. (2008)). At the same time, the original more massive component of the eclipsing pair evolves to fill its critical Roche lobe. Then, the binary was undergoing a rapid mass-transfer stage, and after the mass ratio was revered, the central binary finally reached the present semi-detached configuration with a slow secondary-to-primary mass transfer on the nuclear time scale of the present less massive star. Supergiant stars have recently been found to be tertiary companions of OB-type close binaries in the two triple systems HD 167971 and FR Sct (e.g., Leitherer et al. (1987); Davidge \\& Forbes (1988), and Pigulski \\& Michalska (2007)). They resemble the progenitor of V Puppis with the tertiary supergiant companion evolving into a compact object. As for V Puppis, the existence of the black hole in the outer orbit poses strong constrain on the formation of the black hole, the mass loss in the supernova that lead to the formation of the black hole, and the magnitude of the kick received by the black hole, all make the object a much more interesting system for future study."
    },
    "0806/0806.3049_arXiv.txt": {
        "abstract": "We update the SMC, Bridge, and LMC catalogues of extended objects that were constructed by members of our group from 1995 to 2000. In addition to the rich subsequent literature for the previous classes, we now also include HI shells and supershells. A total of 9305 objects were cross-identified, while our previous catalogues amounted to 7900 entries, an increase of $\\approx12\\%$. We present the results in subcatalogues containing 1445 emission nebulae, 3740 star clusters, 3326 associations, and 794 HI shells and supershells. Angular and apparent size distributions of the extended objects are analysed. We conclude that the objects, in general, appear to respond to tidal effects arising from the LMC, SMC, and Bridge. Number-density profiles extracted along directions parallel and perpendicular to the LMC bar, can be described by two exponential-disks. A single exponential-disk fits the equivalent SMC profiles. Interestingly, when angular-averaged number-densities of most of the extended objects are considered, the profiles of both Clouds do not follow an exponential-disk. Rather, they are best described by a tidally-truncated, core/halo profile, despite the fact that the Clouds are clearly disturbed disks. On the other hand, the older star clusters taken isolately, distribute as an exponential disk. The present catalogue is an important tool for the unambiguous identification of previous objects in current CCD surveys and to establish new findings. ",
        "introduction": "\\label{intro}  The Magellanic Clouds are fundamental galaxies for astrophysics owing e.g. to their proximity, chemical compositions, age distributions, star cluster structural properties and related dynamical evolution, and as two close-by interacting galaxies  (e.g. \\citealt{Wes90}, \\citealt{DaCosta91}, \\citealt{Piatti02}, \\citealt{MG03}, \\citealt{DMG02}, \\citealt{BekChi07}). \\citet{Schaefer08} reviews recent estimates of the distance to the Clouds and arrives at the values $d_{\\rm LMC}\\approx50$\\,kpc and $d_{\\rm SMC}\\approx60$\\,kpc.  More than 50 years have elapsed since the first attempts to systematically catalogue extended objects in the Magellanic Clouds (e.g. the nonstellar emission nebulae of \\citealt{Henize56}). We refer as extended objects the emission nebulae, star clusters, associations, and HI shells and supershells. Early catalogues of star clusters included brighter ones in the SMC (\\citealt{Kron56}, \\citealt{Lindsay58}) and LMC (\\citealt{ShaLind63}; \\citealt{LynWes63}). Fainter clusters with deeper photographic material were detected, e.g. by \\citet{HodSex66} and \\citet{Hodge86}. Binary or multiple clusters are another characteristics of many Magellanic Cloud clusters, showing their importance for cluster dynamical evolution (\\citealt{BhaHad88}; \\citealt{Oliveira00}). Examples of catalogues of associations are \\citet{LuHo70} for the LMC, \\citet{Hodge85} for the SMC, and \\citet{BatDem92} in the Bridge. A complementary study to \\citet{Henize56} is the catalogue of nebular complexes by \\citet{DEM76} based on H$\\alpha$ plates.  Hubble Space Telescope (HST) instrumention revealed serendipitously two faint clusters in the LMC bar that were undetected in Sky Survey plates (\\citealt{Santiago98}). These two clusters suggested the existence of an important undetected faint population of clusters. CCD mosaics, e.g. \\citet{Pietrzynski98} in the SMC central parts, started to unveil that elusive cluster population.  \\begin{table*} \\caption[]{New Clusters and Associations} \\label{tab1} \\renewcommand{\\tabcolsep}{8.0mm} \\renewcommand{\\arraystretch}{1.25} \\begin{tabular}{lccl} \\hline\\hline Reference &Analysed Number& New Objects& Acronym \\\\ (1)&(2)&(3)&(4)\\\\ \\hline \\citet{TLD88}   &  1      &        1    &     TLB\\\\ \\citet{WB97}     &  5      &        1    &     WB\\\\ \\citet{Walb99a}  &  2      &        2    &     WBB\\\\ \\citet{Walb99b}   &  4      &        1    &     WDP\\\\ \\citet{WMAB02}   &  5      &        3    &     WMB\\\\ \\citet{HM99} &  3      &        1    &     HCD99-\\\\ \\citet{HM00}    & 2      &        1    &     HRR\\\\ \\citet{HM01}    &  1      &        1    &     HCD01-\\\\ \\citet{HM01}   &  2      &        1    &     HCD02-\\\\ \\citet{HM03}  &  2      &        2    &     HMW\\\\ \\citet{MHMW05}  &  1     &         1    &     MHW\\\\ \\citet{BPF05}     &  1     &         1    &     Bologna\\\\ \\citet{Nota06}; \\citet{Sabbi07}   & 16     &        16    &     NSS\\\\ \\citet{Nakajima05}      &  25     &        18    &     NKD\\\\ \\citet{Testor06}   &   2     &         2     &    TLF\\\\ \\citet{Testor07}     &   2     &         2     &    TLK\\\\ \\citet{Pietrzynski99}     & 615     &       126     &   LOGLE\\\\ \\citet{Gouliermis03}  - Clusters         &  259     &       125     &   GKK-O\\\\ \\citet{Gouliermis03}  - Associations      & 153     &       102    &    GKK-A\\\\ \\citet{Gouliermis07} & 5 & 5 & GKH\\\\ \\citet{HenneK08} & 5 & 0 & HGH\\\\ \\citet{Schm08} &1 &1 & SGDH\\\\ \\hline \\end{tabular} \\begin{list}{Table Notes.} \\item Otherwise stated, objects are essentially all clusters. \\end{list} \\end{table*}   \\begin{figure*} \\resizebox{\\hsize}{!}{\\includegraphics{Fig1.eps}} \\caption{Angular distribution of the 9305 extended objects in the Magellanic System. Clusters older than 4\\,Gyr are shown as large blue circles. The derived LMC and SMC centroids (see Table~\\ref{tabCC}) are indicated by red triangles. The position of 30\\,Dor is shown by the plus sign.} \\label{fig1} \\end{figure*}  Bica \\& Schmitt (1995, updated in \\citealt{BD00}) and \\citet{BSDO99} compiled general catalogues from numerous previous catalogues and many lists sparsely distributed in the literature, and presented new findings based on Sky Survey plates. The total number of extended objects in these catalogues is 7900, including the SMC, intercloud (Bridge) region, and LMC. Cross-identifications that take into account object plate identifications, class, size, positions, and uncertainties were carried out.  As present-day Magellanic Cloud Surveys provide their first results, like the University of Michigan UM/CTIO Magellanic Cloud Emission Line Survey (MCELS)\\footnote{http://www.ctio.noao.edu/mcels/ }, or are scheduled for soon, such as the Visible and Infrared Survey Telescope for Astronomy (VISTA)\\footnote {http://www.eso.org/gen-fac/pubs/messenger/archive/no.127-mar07/arnaboldi.pdf}, we found it timely to update the catalogues of \\citet{BSDO99} and \\citet{BD00}, so that  literature objects  can be easily identified, and new findings in deeper surveys can be confidently asserted.  Besides the update, we will use the present catalogue to investigate structural properties of the Clouds as probed by the large-scale spatial distribution of different classes of objects. We will also examine potential effects of the LMC, SMC, and Bridge tidal fields on the structure of individual objects.  This paper is structured as follows. In Sect.~\\ref{UpCat} we present the updates and additions to the Magellanic System catalogue. In Sect.~\\ref{GlobProp} we discuss statistical properties of the different object classes contained in the catalogue, such as the distributions of apparent size and ellipticity. In Sect.~\\ref{DepGC} we examine dependences of the object parameters with distance to the Clouds centroids. In Sect.~\\ref{CloudStr} we investigate the structure of both Clouds with the spatial distribution of the catalogue objects. Concluding remarks are given in Sect.~\\ref{Conclu}. ",
        "conclusions": "\\label{Conclu}  The primary goal of this paper is to update the catalogue of extended objects in the Magellanic System. With the recent addition of HST, CCD mosaics, and survey data, the number of known objects in the Clouds now reaches 9503, which represents a relatively substantial increase of $\\approx12\\%$ with respect to the previous versions (\\citealt{BH95}; \\citealt{BSDO99}; \\citealt{BD00}). It now includes HI shells and supershells, cross-identifications with the previous literature, and subsequently discovered objects.  Such a number of objects is large enough to allow for a statistically significant analysis of environmental effects on the distribution of structural parameters among the different classes of objects, in the LMC, SMC, and Bridge tidal fields separately. Star clusters present similar distributions of structural parameters in the three MC subsystems. SMC associations (and related objects, emission nebulae, and HI shells and supershells), on the other hand, tend to be larger and more circular than in the LMC. We also detect evidence that the apparent diameter of clusters and associations increase with the distance to each Cloud centroid. The ellipticity presents the opposite trend, especially in the SMC. These relations are consistent with the relative strengths of the LMC, SMC, and Bridge tidal fields. Indeed, the standard model of \\citet{BekChi07} assumes the masses $M_{\\rm LMC}=2\\times10^{10}\\,\\ms$ and $M_{\\rm SMC}=3\\times10^{9}\\,\\ms$ for the LMC and SMC, respectively. Obviously, the Bridge is less massive than the SMC. Such masses can produce significant tidal stress on star clusters and associations, depending on their location.  With respect to the angular distribution of objects, number-density profiles extracted along the LMC bar and in a perpendicular direction can be reasonably well represented by two exponential disks, one for the bar region and the other for the outer parts of the profile (Fig.~\\ref{fig8}). Disk scale-lengths parallel to the bar are $\\approx2.7$\\,kpc and $\\approx0.4$\\,kpc, for the bar region and beyond, respectively. In the perpendicular direction they are $\\approx3.8$\\,kpc and $\\approx1.2$\\,kpc, for the high-surface brightness disk and outer ring, respectively. Similar profiles extracted along equivalent directions in the SMC follow a single exponential disk with $\\approx0.5$\\,kpc of scale-length.  Alternatively, when (angular-averaged) radial number-density profiles are considered, the large-scale structure of both Clouds appears to be best described by a 3-parameter King-like function, characterised by core and halo sub-structures. In this case, the LMC core and tidal radii are $\\rc\\approx2.6\\degr$ and $\\rt\\approx8.2\\degr$, respectively; SMC values are a factor $\\approx0.4$ of the LMC ones. In absolute scale, LMC values are $\\rc\\approx2.3$\\,kpc and $\\rt\\approx7.2$\\,kpc, while SMC ones are about half of these. The tidal/core radii ratio in both Clouds imply low concentration parameters, comparable to those of sparse Galactic GCs.  What emerges from the present work is a scenario where the present-day, (angular-averaged) large-scale structures of both Clouds appear to behave as tidally-truncated systems (which is not unexpected, since they are Milky Way satellites), characterised by well-defined core and halo sub-structures. This picture comes about despite the fact that both Clouds are not spherical systems (Fig.~\\ref{fig1}). Thus, they have undergone severe tidal perturbation when the last dynamical and hydrodynamical interaction between the Clouds took place, about 200\\,Myr ago (\\citealt{BekChi07}). Taken isolately, the older LMC and SMC star clusters, on the other hand, appear to be distributed as an exponential disk. This distribution is possibly reminiscent of the Clouds structure prior to the last interaction.  The LMC/SMC interaction is not unusual in the local Universe. They can be classified into the minor-merger interaction picture, which involves high and low-mass galaxies (e.g. \\citealt{FerPas04}, and references therein). A series of examples of such interactions involving galaxies with prominent bulge and disk has been studied in that paper. Some of the disk galaxies developed the double-disk structure, similarly to the present case of the LMC."
    },
    "0806/0806.3080_arXiv.txt": {
        "abstract": "Clusters of galaxies offer a robust test bed for probing the nature of dark matter that is insensitive to the assumption of the gravity theories. Both Modified Newtonian Dynamics (MOND) and General Relativity (GR) would require similar amounts of non-baryonic matter in clusters as MOND boosts the gravity only mildly on cluster scales. Gravitational lensing allows us to estimate the enclosed mass in clusters on small ($\\sim\\,20\\,-\\,50\\,{\\rm kpc}$) and large ($\\sim$ several 100 kpc) scales independent of the assumptions of equilibrium. Here we show for the first time that a combination of strong and weak gravitational lensing effects can set interesting limits on the phase space density of dark matter in the centers of clusters. The phase space densities derived from lensing observations are inconsistent with neutrino masses ranging from 2 - 7 eV, and hence do not support the 2 eV-range particles required by MOND. To survive, the plausible modifications for MOND may be either an additional degree of dynamical freedom in a co-variant incarnation or mass-varying theories of neutrinos. ",
        "introduction": "The Newtonian Poisson equation, if sourced by purely ordinary baryonic matter, seriously under-predicts the accelerations seen in a wide range of scales. An extra source term, customarily known as either dark matter or a scalar field source in MOND, is needed for consistency (Zhao \\& Famaey 2006; Famaey, Gentile, Bruneton \\& Zhao 2007). The precise nature of dark matter remains unknown within the Newtonian framework as well. The most viable dark matter candidate is a fermionic, neutral particle that condensed from the thermal bath of the early Universe (Kolb \\& Turner 1990). While the detection of neutrinos confirmed the concept of a possible, particulate dark matter candidate, it did very little in closing the wide gap between observed matter density versus that what is required to explain lensing results and motions on galaxy scales. There remains the need for more exotic species of dark matter. In principle, galaxy scale observations ought to be sensitive to the free-streaming scale of these particles.  Another approach that has been followed is to constrain the phase space densities of dark matter particles.  Studying the phase space density of galaxy scale halos derived from observational data, Sellwood (2000) noted that while the peak phase space density of dark matter in galaxies is far from having a universal value, there does appear to be a favoured scale of a few keV. The same keV mass scale reappears in the recent discussions of the cores and dynamical friction in dwarf spheroidals such as Ursa Minor (Kleyna et al. 2005) and Fornax (Goerdt et al. 2006). The dark matter dominated Fornax dwarf spheroidal has globular clusters orbiting at roughly $\\sim$ 1 kpc from its centre.  Goerdt et al. (2006) argue that if the dark matter halo hosting Fornax has a cuspy density profile, the globular clusters would sink to the centre from their current positions within a few Gyr, presenting a puzzle as to why they do indeed survive at the present epoch. They show that this timing problem is alleviated by adopting a cored dark matter halo. In that instance, using numerical simulations and analytic calculations they argue that the sinking time is many Hubble times; and the globulars would effectively halt at the core radius of the dark matter halo. Using the current positions of the globulars Goerdt et al. (2006) therefore conclude that the Fornax dwarf spheroidal has a shallow inner density profile with a finite core radius.  This immediately implies that the dark matter component is warm, with an upper limit to its mass of $\\sim\\,0.5\\,{\\rm keV}$. Such a warm dark matter candidate would suppress structure formation on small scales alleviating another problem the so-called substructure problem that seems to be endemic to Cold Dark Matter (CDM) models.  On the other hand, the flux anomalies of gravitationally lensed quasars argues for the existence of kpc clumps, too dense for keV warm dark matter particles (Metcalf \\& Zhao 2002; Miranda \\& Maccio 2007).  However, in competing gravity theories like MOND, the interpretation of dwarf galaxy scale rotation curves is very different due to the much larger MONDian gravity than predicted by Newtonian theory. With a simple boost of gravity below a scale $a_0\\,\\sim\\,10^{-8}$m/s$^2$, the need for dark matter on dwarf galaxy scales becomes much weaker if we assume MOND (Famaey, Gentile, Bruneton \\& Zhao 2007 and references therein, Sanders \\& McGaugh 2002). Part of the reason that MOND is able to mimic CDM effectively is that there is a common acceleration scale $g\\,\\sim\\,a_0$ in the dark matter cusp of the Navarro, Frenk, White profile (Xu, Wu \\& Zhao 2007), which appears on galaxy to galaxy cluster scales.  This is exactly the scale on which MOND can supplement ordinary gravity, so it is not surprising that MOND and CDM often give comparable fits to data. However, there are still a few tough challenges for MOND (Klypin \\& Prada 2007; Famaey, Bruneton \\& Zhao 2007) even on these scales.  Additionally, Zhao (2005) noted that MOND would require globular clusters and dwarf galaxies to have the same size tidal radius, which appears to contradict current observations.  So while there are ambiguities on galaxy scales, cluster scales are more promising to discriminate between the two theories as the effects of MOND are expected to be mild.  This is due to the fact that MOND boosts the gravitational constant up only by a mild factor $1/\\mu(x)$, where $\\mu(x) = {x \\over 1+x} \\sim 0.5$ in clusters of galaxies (Wu, Zhao, Famaey, Gentile, Tiret, Combes, Angus \\& Robin 2007). The boosting is much larger in dwarf galaxies, by a factor of $\\sim 11$, where $x \\sim 0.1$. In short, the evidence for DM on galaxy scales is weak in the context of alternatives like MOND. However, as we show below, more robust tests definitely derive from lensing in clusters of galaxies, in particular, the potent combination of strong and weak lensing observations.  Clusters of galaxies are the most massive and recently assembled structures in the Universe. In the context of the hierarchical growth of structure in a cold dark matter dominated Universe, clusters are the repository of copious amounts of the dark matter. Gravitational lensing, predicted by Einstein's theory of General Relativity, is the deflection of light rays from distant sources by foreground mass structures is now detected in over a 100 clusters. Dramatic strong lensing occurs when there is a rare alignment of background sources with the dense central region of a foreground cluster. This produces highly distorted, magnified and multiple images of a single background source (Schneider, Ehlers \\& Falco 1992). However, more commonly, the observed shapes of background sources viewed via a foreground cluster lens are systematically elongated, in the so-called weak lensing regime. Coupling strong and weak lensing offers the most reliable probe of the distribution of dark matter on various cosmic scales (Blandford \\& Narayan 1992; Mellier 2002; Schneider, Ehlers \\& Falco 1992).  In particular, the combination of data from these two regimes offers an unprecedented insight into the detailed mass distribution of clusters (Natarajan, Kneib, Smail \\& Ellis 1998; Natarajan, Kneib \\& Smail 2002; Bradac et al. 2006; Jee et al. 2007). The lensing distortion in the shapes of background galaxies viewed through fore-ground mass distributions is independent of the dynamical state of the lens, therefore, unlike other methods for mass estimation there are fewer biases in lensing mass determinations.   Strong lensing studies of the inner regions of several clusters indicate that the dark matter distribution can be represented by the combination of a smoothly distributed, extended component and smaller-scale clumps or subhaloes associated with luminous galaxies (Kneib et al. 1996; Natarajan \\& Kneib 1997; Natarajan, Kneib, Smail \\& Ellis 1998). The smooth component has been detected using weak lensing techniques out to the turn-around radius (typically of the order of several Mpc) in clusters (Kneib et al. 2003; Broadhurst et al. 2005). The lensing derived density profile of the smooth component, and its agreement with profiles computed from high resolution numerical simulations of structure formation in the Universe is currently well studied (Navarro, Frenk \\& White 1997; Navarro et al. 2004; Sand et al. 2004).  In addition, the granularity of the dark matter distribution associated with individual galactic subhaloes holds important clues to the growth and assembly of clusters (Tyson, Kochanski \\& dell'Antonio 1998; Broadhurst, Huang, Frye \\& Ellis 2000; Limousin et al. 2007; Jee et al. 2007).  We exploit the technique of galaxy-galaxy lensing, which was originally proposed as a method to constrain the masses and spatial extents of field galaxies (Brainerd, Blandford \\& Smail 1996). The methodology has since been extended and developed to apply inside clusters (Natarajan \\& Kneib 1996; Natarajan et al. 1998; 2002a; Natarajan, De Lucia \\& Springel 2007). Constraints on the masses of subhaloes associated with galaxies in clusters is also now available for several clusters.  The outline of this paper is as follows: in Section 2 we present the constraints obtained to date on MOND from clusters, in Section 3 we describe our method to derive the central densities of clusters and cluster galaxies, implications thereof for MOND are in Section 4 and we conclude with a discussion in Section 5. ",
        "conclusions": ""
    },
    "0806/0806.3339_arXiv.txt": {
        "abstract": "Both the light curve and spectral evolution of the radiation from a relativistic fireball with extremely short duration are studied, in order to examine the curvature effect for different forms of the radiation spectrum. Assuming a $\\delta$ function emission we get formulas that get rid of the impacts from the intrinsic emission duration, applicable to any forms of spectrum. It shows that the same form of spectrum could be observed at different times, with the peak energy of the spectrum shifting from higher energy bands to lower bands following $E_{peak}\\propto t^{-1}$. When the emission is early enough the $t^{2}f_{\\nu }(t)$ form as a function of time will possess exactly the same form that the intrinsic spectrum as a function of frequency has. Assuming $f_\\nu \\propto \\nu^{-\\beta} t^{-\\alpha}$ one finds $\\alpha=2+\\beta$ which holds for any intrinsic spectral forms. This relation will be broken down and $\\alpha > 2+\\beta $ or $\\alpha \\gg 2+\\beta $ will hold at much later time when the angle between the moving direction of the emission area and the line of sight is large. An intrinsic spectrum in the form of the Band function is employed to display the light curve and spectral evolution. Caused by the shifting of the Band function spectrum, a temporal steep decay phase and a spectral softening appear simultaneously. The softening phenomenon will appear at different frequencies. It occurs earlier for higher frequencies and later for lower frequencies. The terminating softening time $t_{s,max}$ depends on the observation frequency, following $t_{s,max}\\propto \\nu^{-1}$. This model predicts that the softening duration would be linearly correlated with $t_{s,max}$; the observed $\\beta_{min}$ and $\\beta_{max}$ are determined by the low and high energy indexes of the Band function; both $\\beta_{min}$ and $\\beta_{max}$ are independent of the observation frequency. ",
        "introduction": "The successful launch of the Swift satellite (Gehrels et al. 2004) has made great advance for the observations of the X-ray afterglows of gamma-ray bursts (GRBs) in the past few years (for a recent review, see Zhang 2007). Among the many new findings obtained by the Swift instruments, some emerge as a puzzling. One is the spectral evolution which was detected in the tails of some bursts (Campana et al. 2006; Ghisellini et al. 2006; Gehrels et al. 2006; Mangano et al. 2007; Zhang et al. 2007a). The phenomenon has not been predicted by the curvature effect nor by the standard external shock afterglow model. Soon after that, Zhang et al. (2007b) performed a systematic analysis on this issue and found that 33 of the 44 bursts with bright steep decay tails show an obvious spectral evolution --- an observed softening. This suggests that, the detected spectral evolution is not a rare phenomenon, but instead, it is quite common (at least detectable in the majority of GRBs that have bright steep decay tails).  Several attempts have been made to interpret this softening phenomenon. It was proposed that a central engine producing a soft but decaying afterglow emission might be responsible for some of these bursts (Campana et al. 2006; Fan et al. 2006; Zhang et al. 2007b). Some bursts with strong softening might be accounted for by a cooling of the internal-shocked region (Zhang et al. 2007b). A possible thermal component has also been tried. However, Yonetoku et al. (2008) showed that introducing a thermal component is not sufficient to explain all of the spectral softening, and thus additional spectral evolution is required. In investigating the origins for the spectral evolution of GRB 070616, Starling et al. (2008) ruled out the possibility that a superposition of two power-laws causes the evolution. In stead, they considered a possibility of an additional component dominant during the late prompt emission. They proposed that a combination of the spectral evolution and the curvature effect may cause the observed steep decay phase of the light curve. Another scenario is based on the cannonball model of GRBs, which was suggested to be responsible for both the temporal behavior and the spectral softening of the bursts observed by Swift (Dado et al. 2008).  As shown in Zhang et al. (2007b), the spectral softening is accompanied by a very steep decay phase. This phase is seen directly following the prompt emission and is naturally (and generally) regarded as the tail of the prompt emission (Tagliaferri et al. 2005; Barthelmy et al. 2005; Liang et al. 2006). The tail was suspected to arise from the emission of the high latitude of the fireball surface, which is often referred to as the curvature effect (Kumar \\& Panaitescu 2000; Dermer 2004; Dyks et al. 2005; Liang et al. 2006; Panaitescu et al. 2006; Zhang et al. 2006, 2007b; Butler \\& Kocevski 2007; Qin 2008; Starling et al. 2008). A full consideration of the curvature effect includes the delay of time and the shifting of the intrinsic spectrum as well as other relevant factors of an expanding fireball (for detailed explanation and analysis, see Qin 2002; Qin et al. 2004, 2006; Qin 2008). The effect has been intensively studied in the prompt gamma-ray phase, such as the profile of the light curve of pulses, the spectral lags, the power-law relation between the pulse width and energy, the evolution of the hardness ratio and the evolution of the peak energy (Fenimore et al. 1996; Sari \\& Piran 1997; Qin 2002; Ryde \\& Petrosian 2002; Kocevski et al. 2003; Qin \\& Lu 2005; Shen et al. 2005; Lu et al. 2006, 2007; Peng et al. 2006; Qin et al. 2004, 2005, 2006; Jia 2008).  In a recent investigation, Butler \\& Kocevski (2007) concluded that the early emission in $>90\\%$ of early afterglows has a characteristic $\\nu f_{\\nu}$ spectral energy $E_{peak}$, which likely evolves from the $\\gamma$% -rays through the soft X-ray band on timescales of $10^2 - 10^4$ s after the GRB. Many careful analyses revealed that there do exist some bursts with their peak energy $E_{peak}$ decreasing from a higher band to a much lower band. These bursts include: GRB 060124, from 108 keV to 1.3 keV (Butler \\& Kocevski 2007); GRB 060614, from 8.6 keV to 1.1 keV (Butler \\& Kocevski 2007; Mangano et al. 2007); GRB 060904A, from 163 keV to 2.28 keV (Yonetoku et al. 2008); GRB 061121, from 270 keV to 0.95 keV (Butler \\& Kocevski 2007); GRB 070616, from 135 keV to $\\sim$4 keV (Starling et al. 2008). For GRB 070616, the spectral softening evolution was observed even in the prompt emission phase: its duration is $T_{90}=402.4s$, while the softening starts from 285s and extends to 1200s (Starling et al. 2008). Among them, GRB 060614, GRB 060904A and GRB 061121 are members of the Zhang et al. (2007b)'s sample. Although the curvature effect was rejected to interpret the softening by some authors, Starling et al. (2008) insisted that the observed shifting of the peak energy is in agreement with what expected by the curvature effect: the peak energy of the Band function spectrum passes through the $\\gamma$-ray band at a relatively early time, while it passes through the X-ray band at a later time due to the high latitude emission. They proposed that both the curvature effect and a strong spectral evolution cause the steep decline in flux. Based on the explicit illustration of the evolution of the whole spectral form in Butler \\& Kocevski (2007), we suspect that the curvature effect alone might be responsible for both the spectral softening and the accompanied steep decay light curve.  To reveal the pure curvature effect and get rid of the possible impacts from the emission duration, we concern in this paper only extremely short intrinsic emission. Focusing on this emission has two advantages. The first is that the formulas become very simple, and the second is that many key characteristics of the effect can be plainly illustrated. The paper is organized as follows. In Section 2, we present basic formulas of the full curvature effect, applicable to any temporal and spectral forms of emission. We discuss light curves and spectral evolution arising from an extremely short emission in Section 3. In Section 4, we assume an intrinsic spectrum in the form of the Band function and illustrate the corresponding light curve and spectral evolution in detail. Parameters that affect the results are discussed in Section 5. In Section 6, we apply the model to the XRT band and also to a Swift burst. Discussion and conclusions are presented in Section 7. ",
        "conclusions": "We investigate the influence of the curvature effect on both the light curve and the spectrum of late emission of GRBs, attempting to explain the observed softening phenomenon in early X-ray afterglows of Swift. As the first step of investigation, we explore only the case of extremely short intrinsic emission, for which we assume and apply a $\\delta$ function emission. Although how an emission is extremely short is currently unclear (this deserves a detailed investigation in the near future) and the condition that the emission is extremely short might be rare, the investigation is necessary since by considering such emission the possible impacts from the emission duration can be ignored and then the pure curvature effect can plainly be illustrated.  Formulas presented in Qin (2008) are employed to study this issue. Unlike what investigated in Qin (2008) and other relevant theoretical analyses (e.g., Fenimore et al. 1996; Sari et al. 1998; Kumar \\& Panaitescu 2000), we do not limit our study on an intrinsic power law spectrum. Instead, we consider more general spectral form of emission. Assuming a $\\delta$ function emission we get formulas that get rid of the impacts from the intrinsic emission duration, which are applicable to any forms of spectrum. According to these formulas, one would observe the same form of spectrum at different times, with the peak energy of the spectrum shifting from higher energy bands to lower bands. This was detected recently in the early X-ray afterglows of some GRBs (see Butler \\& Kocevski 2007). The peak energy is expected to decline following $E_{peak}\\propto t^{-1}$. In the case when the emission is early enough so that the emitting area on the fireball surface is not far from the line of sight (say, $\\theta \\ll \\pi/2$ or $t \\ll R_{c}/v$), the temporal power law index and the spectral power law index will be well related by $\\alpha =2+\\beta $. As a consequence, the form $t^{2}f_{\\nu }(t)$ will possess exactly the intrinsic spectral form (say, when replacing $t$ with $\\nu$ one will get from $t^{2}f_{\\nu }(t)$ as a function of $\\nu$ that take exactly the same form of the intrinsic spectrum). The $\\alpha =2+\\beta $ relation will be broken down and $\\alpha > 2+\\beta $ or $\\alpha \\gg 2+\\beta $ will hold (see Fig. 4) at much later time when the angle between the moving direction of the emission area and the line of sight is large.  As revealed in Butler \\& Kocevski (2007) and suggested by Starling et al. (2008), we focus our attention to the emission with a Band function spectrum. This spectrum has a power law behavior in both lower and higher energy bands, where the two power laws are smoothly connected (Band et al. 1993). Using this spectral form, we plot the light curve and the spectral evolution with our formulas, expected at 1 keV and other frequencies. The analysis shows that there do exist a temporal steep decay phase and a spectral softening which occur simultaneously. As Fig.2 reveals, both the steep decay light curve and the spectral softening are caused by the shifting of the Band spectrum. As mentioned above, Starling et al. (2008) suggested that the steep decline in the flux of GRB 070616 may be caused by a combination of the strong spectral evolution and the curvature effect. Based on the above argument, we insist that both the spectral evolution and the steep decline observed in GRB 070616 and other Swift bursts are likely to be caused merely by the curvature effect. In addition, we find that, just as what is illustrated in Qin (2008), the $cos\\theta$ term, $R_{c}/v-t$, plays a role in producing the light curve, which ``attaches'' a cutoff tail to the latter (see Fig. 1). The spectral softening terminates when the emission is dominated by that from the high energy portion of the Band function spectrum. Thus, there exists a maximum of the spectral index, $\\beta_{max}$. After the $\\beta_{max}$ appears, it lasts to the end of the emission, including the phase of the cutoff tail. Our analysis shows that the softening due to the curvature effect will appear at different frequencies; it occurs earlier for higher frequencies and later for lower frequencies; a characteristic of the softening, the terminating softening time $t_{s,max}$ (when the $\\beta_{max}$ appears), depends strictly on the observation frequency, which follows $t_{s,max}\\propto \\nu^{-1}$. Although this is concluded based on the assumption of extremely short emission, we tend to believe that its main characters hold in most cases since the light curve from any finite emission is contributed by countless $\\delta$ function emission. The combination of these countless $\\delta$ function emission would change the values of some quantities such as $t_{s,max}$, but would not change the trend. Starling et al. (2008) showed, the spectral evolution of GRB 070616 starts earlier at $ \\gamma$-ray energies (while the X-ray flux is still at an approximately constant level) and begins much later at X-ray energies around the onset of the steep X-ray decay. In terms of the curvature effect, this is due to the following fact: the peak energy of the Band function spectrum passes through the $\\gamma$-ray band at a relatively early time, while it passes through the X-ray band at a later time.  Parameters that might have impacts on the light curve and the spectral evolution are also studied. Whilst a smaller Lorentz factor shifts the spectral softening to larger time scales, smaller values of the fireball radius and the rest frame peak energy make the occurrence of the softening earlier. The analysis shows that the terminating softening time appears much later than the start time of the softening. The former is always larger than the latter by about one order of magnitude. It is therefore predicted that the duration of the softening and the terminating softening time would be linearly correlated. It also shows that the low energy index $\\alpha _{B,0}$ puts a lower limit to the observed spectral index $\\beta$ and the high energy index $\\beta _{B,0}$ confines the upper limit of $\\beta$. The following conclusions are reached: the observed $\\beta_{min}$ and $\\beta_{max}$ are determined by the low and high energy indexes of the observed Band function spectrum; $\\beta_{min}$ and $\\beta_{max}$ for different observation frequencies would remain unchanged as long as the whole softening process appears within the whole available observation time at the concerned frequencies.  As application, we study the light curve and the spectral evolution over the XRT band. That is, the light curve is that has been integrated over the $0.3-10$ keV band and the spectral index is estimated by assuming a power-law over this band. The analysis shows that the light curve slightly deviate from that measured at 1 keV, whilst the spectral evolutionary curve significantly betrays that measured at 1 keV. This suggests that the $\\alpha =2+\\beta $ relation will be violently violated if one measures the light curve and estimates the spectral index over a wide band.  Another application is to check the $\\alpha =2+\\beta $ relation by employing the light curve and spectrum data of GRB 060614. The temporal index $\\alpha$ in the steep decay phase of this burst is evaluated. We compare $\\alpha$ and the spectral index $\\beta$ in the same softening phase (it is also the steep decay phase of the light curve). It shows that the $\\alpha-2$ vs. $t$ curve and the $\\beta$ vs. $t$ curve are roughly in agreement, suggesting that the $\\alpha =2+\\beta $ relation is roughly satisfied in this phase and the softening of this burst might possibly be due to the curvature effect.  What we have investigated are based on the assumption of extremely short emission which ignores the possible impacts from the emission duration. While the contribution of the emission duration might have less affect on the spectrum, it might obviously affect the temporal profile and hence the temporal index. Therefore the $\\alpha =2+\\beta $ relation might not hold in this case. Combining this with the problem arising from the estimation of the spectral index in the softening phase over an energy band rather than at a fixed frequency, it might be able to explain why the $\\alpha=2+\\beta$ relation is not common in Swift bursts (see Liang et al. 2006 and also Zhang 2007 for a detailed discussion). Yonetoku et al. (2008) showed that the two characteristic break energies they considered have a time dependence of $ \\propto t^{-3}-t^{-4}$. This is not in agreement with the prediction of the $E_{peak}\\propto t^{-1}$ law. Perhaps an intrinsic softening might be responsible for this difference when the emission duration is taken into account. Another possibility is that the $E_{peak}$ shifting of this kind of burst is due to structure jets, where in high latitude, $E_{peak}$ would be much smaller than it is in uniform jets. In addition to test the $\\alpha =2+\\beta $ relation and the $E_{peak}\\propto t^{-1}$ law, we suggest to detect the softening at different frequencies as well. The observed relations might not strictly follow what the $\\delta$ function emission model predicts, but as argued above, the trend of the relevant effects would be maintained (which deserves a detailed investigation in the near future).  \\begin{figure}[tbp] \\begin{center} \\includegraphics[width=5in,angle=0]{f10.eps} \\end{center} \\caption{Comparison of the functions of the $t^{3}f_{\\nu }(t)$ vs. $t$ curve (the upper panel) and the $\\nu f_{\\nu}$ vs. $\\nu$ curve (the spectrum; the lower panel) of the emission considered in Fig. 1. The three solid lines in the lower panel are the same lines of 50, 100, and 500 s respectively in Fig. 1.} \\label{Fig. 1_08} \\end{figure}  In addition to these more conventional tests, we propose to try a totally new test which is to check if the $t^{3}f_{\\nu }(t)$ vs. $t$ curve is in agreement with the $\\nu f_{\\nu}$ vs. $\\nu$ curve when replacing $t$ with $\\nu$ in the $t^{3}f_{\\nu }(t)$ and $t$ forms and multiplying a constant to match the corresponding dimensions (see Section 3). Illustrated in Fig. 10 is an example of the comparison (where we compare only the relevant functions and thus do not replace variables or multiply constants to change the dimensions). An advantage of doing so is that we can guess spectrum form merely from the light curve data. For example, when we find that the $t^{3}f_{\\nu }(t)$ form being a perfect power-law function of time, then the spectrum is guessed to be a pure power-law (if the spectral data are found not to obey a power-law, then we will have reasons to doubt if this burst is not affected by the curvature effect). Or, when we find that the $t^{3}f_{\\nu }(t)$ form is a Band function of time, then we will have reasons to guess that the spectrum takes a Band function form. Another usage of plotting the $t^{3}f_{\\nu }(t)$ vs. $t$ curve is to find out the time when the peak energy passes through the observation band (se also Kumar \\& Panaitescu 2000), which can be directly checked by observation and hence becomes a test to the curvature effect as well. According to equation (11), the moment when the peak of $t^{3}f_{\\nu }(t)$ appears is the time when the peak energy passes through the observation frequency $\\nu$, or, it is the time when $E_{peak}=h\\nu$. In plotting the $t^{3}f_{\\nu }(t)$ vs. $t$ curve (see the upper panel of Fig. 10), one can also check the influence of factors other than that of the pure curvature effect by observing and measuring the deviation of a real light curve from that of equation (11). An extra usage of this plot might be a direct comparison of light curves of different bursts, probably enabling us to divide them according to their temporal properties (we strongly suggest a detailed investigation on this issue in the near future).  \\begin{figure}[tbp] \\begin{center} \\includegraphics[width=5in,angle=0]{f11.eps} \\end{center} \\caption{The $t^{3}f_{\\nu }(t)$ (in units of $erg \\cdot cm^{-2} s^{2}$) vs. $t$ curve of GRB 060614 in its steep decay phase. The data are taken from the UNLV GRB Group website.} \\label{Fig. 1_09} \\end{figure}  Displayed in Fig. 11 is the $t^{3}f_{\\nu }(t)$ vs. $t$ curve of GRB 060614. It shows a function of time bearing the Band function form (when replacing $\\nu$ with $t$) attached with a cutoff tail. If we believe that the softening of this burst is due to the curvature effect in the case of extremely short emission, according to Fig. 11 it would be expectable that the peak energy of this burst passes through the corresponding observation energy range (the XRT band) at $\\sim $175 s (or, at $\\sim $175 s, the peak energy of the observed spectrum is just within the energy range of the light curve). Indeed, as presented in Table 4 of Mangano et al. (2007), within the time interval $128-190$ s, the peak energy (when fitting the spectrum with a Band model) is about $2-4$ keV which is well within the $0.3-10$ keV band. A similar result for this burst is also visible in Fig. 1 of Butler \\& Kocevski (2007).  Based on their fits to the composite light curves, Sakamoto et al. (2007) confirmed the existence of an exponential decay component which smoothly connects the BAT prompt data to the XRT steep decay for several GRBs. Yonetoku et al. (2008) also showed that the spectrum of GRB 060904A contains a cutoff tail in its higher energy range, which can be represented by an exponential function. According to the above analysis, the spectral form obviously affects the light curve if the curvature effect is at work and the intrinsic emission is short enough. An intrinsic spectrum with an exponential tail might probably lead to an exponential decay light curve. We notice that the projected factor (the $\\cos \\theta $ term) also produces a very steep decay phase when the angle between the moving direction of the dominant emission area and the line of sight is large enough. However, it is unlikely that many of them (if not only few of them) are due to the $cos\\theta$ term, since this term always appears at a very late time (see Fig. 6 and also Qin 2008). The following factors might also be the cause of this tail: one is the large absorption for higher latitude emission and the other is the limited open angle of jets. Both will lead to a steeper tail. While the former might probably give rise to a smooth decay curve, the latter might probably lead to a sharp feature.  In addition to the shifting of the peak energy, Starling et al. (2008) also observed softening of the low energy power law slope. They measured a softening of the low energy spectral slope from $\\alpha \\sim 0.1 - 1.3$. This implies that the intrinsic spectrum might evolved itself. This cannot be taken into account in the current investigation since a $\\delta$ function emission has no evolution. We suggest to explore in a later investigation the impact from the emission duration and the effect from the intrinsic spectral evolution."
    },
    "0806/0806.0560_arXiv.txt": {
        "abstract": "We present a quantitative comparison between software features of the defacto standard X-ray spectral analysis tool, \\xs\\, and \\isis\\, the Interactive Spectral Interpretation System. Our emphasis is on customized analysis, with \\isis\\ offered as a strong example of configurable software.  While noting that \\xs\\ has been of immense value to astronomers, and that its scientific core is moderately extensible---most commonly via the inclusion of user contributed ``local models\"---we identify a series of limitations with its use beyond conventional spectral modeling.  We argue that from the viewpoint of the astronomical user, the \\xs\\ internal structure presents a \\bbp, with many of its important features \\bfi{hidden} from the top-level interface, thus \\bfi{discouraging} user customization. Drawing from examples in custom modeling, numerical analysis, parallel computation, visualization, data management, and automated code generation, we show how a \\bfi{numerically} scriptable, modular, and extensible analysis platform such as \\isis\\ facilitates many forms of advanced astrophysical inquiry. ",
        "introduction": "\\label{intro} The pursuit of science pushes not only the boundaries of knowledge, but also the limits of technology used to acquire and analyze data from which new knowledge can be distilled.  In this sense software systems for scientific analysis are never truly complete, so it is crucial that they be extensible.  This enables scientists to incorporate custom, experimental codes into the system and drive it in directions not yet supported, or perhaps even envisioned, by its creators.  A corollary is that the mechanisms for and results of the customization should be made as simple, powerful, and general as possible, freeing the practitioner to concentrate more on algorithmic or scientific concerns and less on the plumbing details of a given computational platform.  This paper compares the extent to which two open-source spectral analysis applications actively used in astronomy,  \\xs\\ \\citep{1996ASPC..101...17A} and \\isis\\ \\citep{2002hrxs.confE..17H}, meet these general criteria. The features of both systems are broadly surveyed, with contrasts drawn from a series of examples representative of \\bfi{astronomy in practice}.  A heavy emphasis is placed upon the roles of interpreted arrays, and the capability to quickly generate and numerically manipulate them as powerful enablers of analysis. The emergent theme is that numerical and architectural discontinuities between the internals of \\xs\\ and its controlling \\tcl\\ interpreter, referred to in the large as the \\bbp, limit its full exploitation for scientific tasks that are \\bfi{within the reach of more configurable} interactive analysis systems.  We demonstrate the strong modularity and extensible scriptability of \\isis\\, and suggest that such open and highly configurable systems offer deeper and wider promise for meeting the needs of exploratory scientific analysis.  The latest major versions of \\isis\\ and \\xs, {\\tt 1.4.x} and {\\tt 12.x} respectively, were used in our study.  While \\xs\\ 11 remains in wide use (the first public, non-beta release of \\xs\\ 12 was circa 2005), it is more difficult to extend with a variety of custom models\\footnote{For example, \\xs\\ 11 supports only local models coded in single precision \\fort 77.} and so in terms of extensibility would compare less favorably than does \\xs\\ 12.  Guiding our endeavor will be the sense of what a typically motivated scientist can achieve in each system with reasonable effort. On the premise that a strong programmer can eventually replicate virtually any result using nearly any combination of languages or tools, we are not concerned with what might in principle be achieved {\\em from} either \\xs\\ or \\isis\\, but rather what can be demonstrated {\\em within} each.  By avoiding hypotheticals we aim to transform a potentially open-ended debate into a tractable, concrete discussion. Towards this end we define {\\em working within an application} to mean the execution of commands, or portions thereof, which only access or modify the code or data within the applications addressable memory. In this scheme, for example, issuing an operating system call from \\isis\\ to spawn an external visualization tool would not be ``working in \\isis,\" however enlarging its address space by importing a shared object module, say to smooth an image in memory, would be.  We recognize the utility of integrating complementary but distinct programs into a single logical entity.  Building on trust in known tools, it is a sensible scheme for managing complexity, fostering reuse, and balancing deadlines with goals.  Cooperating applications can, with sufficient effort, be made to appear as seamless as a single program.  \\bfi{One} should not, however, equate the act of \\begin{itemize} \\item dumping internal state to disk files, spawning a series of tools to read this state and create temporary data products along the way, then loading these new data back into the parent application \\end{itemize} with the ability to \\begin{itemize} \\item invoke a function to operate upon a data structure, with both resident in memory. \\end{itemize} \\bfi{In this paper we will argue that the latter offers greater immediacy of use, as well as significantly more algorithmic, data, and visualization flexibility.  It also addresses real scalability concerns \\label{scalability}} as datasets increase in size or tasks increase in time complexity.  A stark example of this is given in \\cite{2005ASPC..347..444D}, which discusses the methods employed to analyze the megasecond observation of Cassiopeia A \\citep{2004ApJ...615L.117H} by the Chandra X-Ray Observatory in 2004, which yielded the first ever map of electron acceleration in a supernova remnant \\citep{2006NatPh...2..614S}.  Over 9 million PHA spectra were extracted from over 300 million events, with hundreds of thousands of spectral fits performed in \\isis\\, to yield arc-second resolution spatial maps of plasma temperatures, line emission, and Doppler velocities.  Custom modules were developed to manage the multi-gigabyte files involved and distribute the spectral extractions and fits across a network of workstations using the Parallel Virtual Machine \\citep{Geist:1994:PPV}.  Based upon the time required to extract a single spectrum from these enormous datasets, with the standard Chandra {\\tt dmextract} tool in CIAO (ca. version 3.2), the authors extrapolated that the spectral extractions alone would have required nearly 10 years had a tool-based approach been taken.  In contrast, each run through the entire PVM-based analysis sequence required less than 12 days to complete. ",
        "conclusions": "Progress in science can be measured by our ability to pose increasingly advanced, open-ended questions and address them with flexible techniques of commensurate sophistication.  Since the age of Newton, one implication of this is that scientists must also practice mathematics, and since the age of Turing it has also meant they must practice programming.  In this paper we have compared two open-source spectral analysis applications, \\xs\\ and \\isis, in an attempt to gauge how their scientific and mathematical reach may be extended with custom programming.  Contrasts between the two have been drawn from the context of current research efforts utilizing high performance computation, numerical modeling, atomic physics, visualization, automated code generation, and data management.  We have demonstrated how these have led to the incorporation of new analytic techniques in \\isis, in ways that are either unexplored or problematic for the current \\xs\\ architecture and its associated \\ftf\\ model of analysis.  \\xs\\ has been of tremendous value to the X-ray community, continues to be actively developed, and has a strong history of community enhancement via local models (and to some extent, \\tcl/TK scripts).  However, we have argued that the top-level simplicity of the \\xs\\ interface, long and rightfully one of the pillars of its appeal, can shroud much of its internal computational and data handling mechanisms. This in turn can render \\xs\\ less adaptable---by the typical user---for evolving research needs.  Commands which operate as black-boxed common denominators of analysis, and permitting only file-based data input, may not be enough to probe some of the more computationally challenging problems facing astronomers.  We conclude that analysis applications such as \\isis, endowed with interpreted whole-array numerical capabilities and an open, modular, and scriptable architecture designed expressly for high configurability, are more favorably equipped to support ad-hoc research needs while not appreciably compromising reproducibility.  Originally envisioned as a tool for Chandra gratings spectroscopy, \\isis\\ has heavily influenced the development of additional Chandra analysis software, while also receiving NASA AISRP funding to continue its evolution within the aforementioned HYDRA project. Although some of the \\isis\\ capabilities we've emphasized do exist in other astronomy software, we are unaware of any publications demonstrating how similar breadths of flexibility and computational power have been collected under the umbrella of a single open-source analysis application and brought to bear on published research in spectral analysis and X-ray astronomy in the manner discussed here."
    },
    "0806/0806.0083_arXiv.txt": {
        "abstract": "Astronomical observations have shown that small carbonaceous molecules can persist in interstellar clouds exposed to intense ultraviolet radiation.  Current astrochemical models lack quantitative information on photodissociation rates in order to interpret these data. We here present {\\it ab initio} multi-reference configuration-interaction calculations of the vertical excitation energies, transition dipole moments and oscillator strengths for a number of astrophysically relevant molecules: C$_3$, C$_4$, C$_2$H, $l-$ and $c-$C$_3$H, $l-$ and $c-$C$_3$H$_2$, HC$_3$H, $l-$C$_4$H and $l-$C$_5$H. Highly excited states up to the 9'th root of each symmetry are computed, and several new states with large oscillator strengths are found below the ionization potentials.  These data are used to calculate upper limits on photodissociation rates in the unattenuated interstellar radiation field by assuming that all absorptions above the dissociation limit lead to dissociation. ",
        "introduction": "\\label{intro} Of the more than 130 different molecules found in interstellar space, an important class is formed by the unsaturated carbonaceous species. Carbohydrides ranging from small molecules such as C$_2$H \\cite{Tucker74} to long chains like C$_8$H \\cite{Bell99} and HC$_{11}$N \\cite{Bell97} have been detected through their millimeter transitions in cold dark clouds like TMC-1 for several decades. Some of these molecules have also been seen in their cyclic form, with c-C$_3$H$_2$ as the best-known example \\cite{Thaddeus85b}. Even negative ions, in particular C$_6$H$^-$, have now been detected \\cite{McCarthy06}. Bare carbon chains are likely present as well but do not have a permanent dipole moment and can therefore not be observed through their pure rotational transitions in emission. Instead, the smallest members of this family, C$_2$ and C$_3$, have been detected in diffuse interstellar clouds, --i.e., clouds which are not completely opaque to visible and ultraviolet radiation-- through their electronic absorptions against bright background stars \\cite{Hobbs79,vanDis84,Maier01}. $c-$C$_3$H$_2$ has been seen in diffuse clouds as well, through absorption at radio wavelengths against background quasars \\cite{Cox88}. The relatively large abundances of these non-saturated molecules, in spite of the fact that there is $10^4$ times more hydrogen than carbon in interstellar clouds, are a vivid demonstration that interstellar chemistry is not in thermodynamic equilibrium. Instead, the kinetics of the reactions that form and destroy these molecules need to be taken into account explicitly in order to explain their abundances. In cold dark clouds, these traditionally involve a series of ion-molecule and neutral-neutral reactions \\cite{Winnewisser93,Herbst03}.    More recently, several of the carbohydrides C$_n$H$_m$ have also been observed in so-called Photon-Dominated Regions (PDRs) \\cite{Hollenbach97}, i.e., clouds which are exposed to intense ultraviolet (UV) radiation. Species like C$_2$H and c-C$_3$H$_2$ have been detected in the Orion Bar PDR \\cite{Hogerheijde95} and even more complex molecules like C$_3$H$_2$ and C$_4$H have been seen in the Horsehead nebula, M38 \\cite{Teyssier04,Pety05}.  Together with the above mentioned observations of diffuse clouds, the data demonstrate that these molecules can exist in UV-exposed environments with abundances comparable to those in cold dark clouds.  Traditional PDR chemical models cannot explain the high abundances of these molecules, leading to speculations that they are perhaps produced by fragmentation of even larger carbonaceous molecules such as polycyclic aromatic hydrocarbons (PAHs) \\cite{Pety05}. However, a major uncertainty in these and other models are the photodissociation rates of the carbohydrides, which are basically unknown.  Other regions in which ultraviolet photons play a role in the chemistry include the envelopes of evolved stars and cometary atmospheres. The same carbohydride molecules have been detected in the envelope of the carbon-rich late-type star IRC+10216, where their similar distributions are a puzzle for photochemical models \\cite{Guelin00}.  In cometary atmospheres, species as complex as C$_4$H have been inferred, whose presence is not readily explained by the standard parent-daughter photochemical models \\cite{Geiss99}.  Many theoretical studies of the electronic structure of small carbohydride molecules exist in the literature \\cite[e.g.,]{Lee85,Botsch89,Jonas92,Cui97,Carter00,Orden98}. However, most of them are limited to the ground and lowest few excited electronic states. Only few studies of higher states exist, most of them performed by Peyerimhoff and her co-workers starting more than 30 years ago, e.g., \\cite{Shih77,Romelt78} As will be shown here, those results are still highly relevant and provide an excellent starting point for further studies.  The aim of this paper is to provide insight into the photodissociation processes of carbohydride molecules through {\\it ab initio} quantum chemical calculations of the vertical excitation energies and oscillator strengths for as large a number of excited electronic states as feasible with current programs.  By assuming that most of the absorptions with energies above the dissociation limit lead to destruction, estimates of the (upper limits of) photodissociation rates under interstellar conditions can be obtained. Even though such calculations are necessarily limited to the smaller members of the carbo-hydride family, they do provide quantitative constraints to test the basic interstellar chemical networks. Specifically, C$_2$H, C$_3$, C$_4$, $l-$ and $c-$C$_3$H, $l-$ and $c-$ C$_3$H$_2$, HC$_3$H, $l-$C$_4$H and $l-$C$_5$H are studied here. Except for HC$_3$H, all these species have been chosen to have been observed in interstellar space. ",
        "conclusions": "We have presented vertical excitation energies, transition dipole moments and oscillator strengths for states up to the 9'th root of each symmetry for several carbonaceous molecules of astrophysical interest.  For lower-lying states, good agreement is generally found with previous studies.  Several new, higher-lying states with large oscillator strengths, often of Rydberg character, are revealed in this work.  The calculated photodissociation rates of the small carbon-bearing molecules studied here are substantial, leading to lifetimes at the edges of interstellar clouds of less than 30 yr. These high rates assume that all absorptions above the dissociation limit indeed lead to dissociation, so that the rates should be viewed as upper limits. Further experimental work is needed to quantify the dissociation efficiencies for the strongest states found in this work."
    },
    "0806/0806.2279_arXiv.txt": {
        "abstract": "{The 11-yr cycle is the best known manifestation of the Sun's activity. While chromospheric cycles have been studied in a number of  solar-like stars, very little is known about how these are reflected in the cyclical behavior of the coronal  X-ray emission in stars other than the Sun.} {Our long-term \\xmm\\ program of long-term monitoring of a solar-like star with a well-studied  chromospheric cycle, HD\\,81809 aims to study whether an X-ray cycle is present, along with studying its characteristics and its relation to the chromospheric cycle.} {Regular observations of \\hd\\ were performed with \\xmm, spaced by 6 months from 2001 to 2007. We studied the variations in the resulting coronal luminosity and temperature, and compared them with the chromospheric Ca\\,{\\sc ii} variations. We also modeled the observations in terms of a mixture of active regions, using a methodology originally developed to study the solar corona. } {Our observations show a well-defined cycle with an  amplitude exceeding 1 dex and an average luminosity approximately one order of magnitude higher than in the Sun. The behavior of the corona of \\hd\\ can be modeled well in terms of varying coverage of solar-like active regions, with a larger coverage than for the Sun, showing it to be compatible with a simple extension of the solar case.} {} ",
        "introduction": "\\label{sec:intro}  Cyclic modulation in the activity level appears to be relatively common among solar-type stars. The best known and studied example is of course the Sun itself, for which the 11-yr cycle was noticed in the form of periodic modulation of the sunspot number already in the mid 19th century. Most activity indicators follow similar cyclical variations (although the study of the Sun's magnetic field shows the true cycle to be 22 yr, with the field polarity reversing each cycle), with the amplitude of the modulation depending strongly on the indicator used. The pioneering and persistent work of O.~Wilson allowed the detection of chromospheric activity cycles on stars other than the Sun, as amply discussed by \\cite{bds+95}.  Chromospheric activity cycles appear to be relatively common in intermediate-activity solar type stars, while less active stars are often found in ``Maunder minimum''-like states, i.e.\\ showing very little long-term variability (\\citealp{bds+95}). The most active stars appear dominated by irregular variability rather than by cyclic behavior. It is, however, not clear if the ``Maunder minimum like'' stars, as well as the Sun itself during the Maunder mininimum, are magnetically less active than the modern solar minimum seems to be, even though they show no HK cycle (\\citealp{js2007}). The present-day Sun has a typical chromospheric modulation amplitude, along the 11-yr cycle, of about a factor of 2 (as measured in the $S$ index).  Evidence of cyclical variations in the coronal emission of stars other than the Sun has so far been elusive. In the Sun the X-ray luminosity follows a cycle with the same periodicity as the chromospheric cycle, but with a much larger amplitude, up to a factor of 100 in the Yohkoh 0.73--2.5 keV band (\\citealp{act96}). Stars with a high activity level, which have been the typical target of X-ray observations, show modest amounts of long-term variability, with no evidence of a cyclic behavior, as for example reviewed by \\cite{ste98b}, and as confirmed from a statistical analysis of the ROSAT All Sky Survey data by \\cite{sl2004}. Evidence of long-term X-ray variability with amplitude of a factor of 2.5, and correlated with the chromospheric $S$ index, was reported by \\cite{hsb+2003} in 61 Cyg A and B (K5V and K7V) spanning 4.5 years of ROSAT HRI data, suggestive of an activity cycle in X-rays. Some weak statistical evidence of solar-like cycles in lower-activity stars was previously reported by \\cite*{hss96} using the RASS observations of the stars with known chromospheric cycles.  To study whether cyclical, high-amplitude variability of the coronal emission is also present on stars other than the Sun, three stellar systems have been monitored since the  beginning of the \\xmm\\ observatory. With a homogeneous temporal basis of almost 7 years, they represent the best available data set for the study of long-term variability and cycle in stars. The three systems monitored are \\hd\\ (discussed in the present paper), \\acen\\ (\\citealp{rsf2005}) and 61 Cyg (\\citealp{hrs+2006}). Of the three, \\hd\\ is the one with the longest available observational baseline and the one which shows the clearest evidence to date of a cyclical behavior.  The initial results of the observations of \\hd\\ were presented by \\cite{fmb+2004} (Paper I). High amplitude variability was evident in the data from 2001 to 2003 discussed in Paper I, with some coherence with the known chromospheric cycle of this star. However, no evidence of a coherent minimum was present in the data up to 2003, which made it impossible to distinguish a cyclic behavior from a simpler long-term trend in the star's activity level; the short data set was therefore completely inadequate for any study of the relation between the chromospheric and coronal activity variations. In this paper we present the results from all available observations, spanning 2001 to 2007, from which the cyclical behavior of the star's corona is evident with little if any remaining ambiguity. ",
        "conclusions": "\\label{sec:concl}  We have presented the results from the first 7 years of monitoring of the solar-like star \\hd\\ (a solar-type star with a well defined cycle in its chromospheric activity) with \\xmm, showing strong evidence of cyclic variability with an amplitude of at least 6 in the ROSAT 0.2-2.5 keV band (and even more considering the outlier observed in June 2002). Our long-term \\xmm\\ program has therefore for the first time resulted in the clear detection of a coronal cycle in a star other than the Sun.  Our observations have however not yet covered a full period of the observed chromospheric cycle. The coverage at the cycle maximum is scarce, and it may be affected by variability of the X-ray emission induced by rotational modulation, so that the single data point does not yet allow us to understand what is the duration of the maximum phase, and to which level is the stochastic variability enhanced with respect to the other phases. The minimum is on the other hand well covered, and it shows a very low level of stochastic variability both in terms of luminosity (a factor $< 2$) and temperature, similar to what is observed in the Sun at the minimum.  \\refcom{At the maximum of activity, our simulations show that the $E\\!M(T)$ distribution is dominated by CO (see Fig.~\\ref{fig:emt}) which determine the values of average temperature and emission measure reported in Fig.~\\ref{fig:fluxtsim}, whereas the contribution from AR to the total coronal emission is negligible. At the maximum, therefore, only the contribution from CO can be univocally determined, whereas for the AR the quoted value of 60\\% is in fact an upper limit to their contribution. When the star is at the minimum of the cycle, the contributions of CO and AR to the $E\\!M(T)$ are comparable; changing their relative weight leads to a rapid change in the average temperature. In this case, therefore, the contributions of CO and AR are reasonably constrained: a doubling of the contribution of CO from 4\\% to 8\\% would lead to an increase in the X-ray luminosity of about 40\\%.}  Therefore, even if the absolute activity level of \\hd\\ at its minimum is much higher than the Sun at minimum, its behavior is very similar. This is somewhat remarkable in that our model shows that a significant fraction of \\hd\\ at minimum is covered by active regions, unlike the Sun which at minimum is dominated by coronal holes. In the future we plan to study in detail the upcoming maximum, with more frequent observations, to decouple the effect of stochastic variability and of rotational modulation of the X-ray emission.  \\hd\\ is the subject of a long-term monitoring program performed with the \\xmm\\ observatory (which also includes $\\alpha$ Cen and 61~Cyg). Two more years of observations (at six months cadence) are already planned on the same target, and we plan to re-propose the target in future \\xmm\\ AOs, to ensure the continuous monitoring. The continuation of the present program will allow us to shed new light on a very basic phenomenon of the Sun, its activity cycles, by studying how similar, somewhat more active stars behave."
    },
    "0806/0806.4252.txt": {
        "abstract": "{}{}{}{}{} % 5 {} token are mandatory  %\\authorrunning{Rodriguez-Fernandez } %\\titlerunning{Dynamics and molecular chemistry in the Galactic center}   \\abstract % context heading (optional) % {} leave it empty if necessary { The gas distribution and dynamics in the inner Galaxy present many unknowns as the origin of the asymmetry of the \\lvd ~ of the Central Molecular Zone  (CMZ). On the other hand, there are recent evidences in the stellar component of the presence of a nuclear bar  that could be slightly lopsided. } % aims heading (mandatory) {Our goal is to characterize the nuclear  bar observed in 2MASS maps and to study the gas dynamics in the inner Milky Way  taking into account this secondary bar.} % methods heading (mandatory) { We have derived a realistic mass distribution by fitting the 2MASS star counts map of Alard (2001) with a model including three  components (disk, bulge and nuclear bar) and we have simulated the gas dynamics,  in the deduced gravitational potential, using a sticky-particles code.} % results heading (mandatory) { Our simulations of the gas dynamics reproduce successfully the main characteristics of the Milky Way for a bulge orientation  of $20\\deg-35\\deg$ with respect to the Sun-Galactic Center (GC) line and a pattern speed of 30-40 \\kmskpc. In our models the Galactic Molecular Ring (GMR) is not an actual ring but the inner parts of the spiral arms, {while the 3-kpc arm and its far side counterpart are lateral arms that contour the bar.} Our simulations reproduce, for the first time, the parallelogram shape of the \\lvd ~ of the CMZ as the gas response to the nuclear bar. This bar  should be oriented by an angle of $\\sim 60\\deg-75\\deg$ with respect to the Sun-GC line and its mass amounts to $(2-5.5)\\,10^9$ \\Msun. We show that the observed asymmetry of the CMZ cannot be due to  lopsidedness of the nuclear bar as suggested by the 2MASS maps. } % conclusions heading (optional), leave it empty if necessary {We do not find clear evidences of lopsidedness in the stellar potential. We propose that the observed asymmetry of the central gas layer can be due to the infalling of gas  into the CMZ in the {\\it l=1.3$\\deg$-complex}.} ",
        "introduction": "Studying the structure and the dynamics of the Milky Way galaxy is necessary to understand its formation and evolution. For example, the large scale star formation rate, the gas transfer to the nucleus and the fueling of the central super-massive black hole are strongly influenced by the large scale dynamics of the Galaxy.  At least three components are needed to explain the rotation curve of the Galaxy: a disc, a bulge and a massive dark halo. In near-infrared images, the bulge is a triaxial structure difficult to distinguish from a bar. The first direct evidence of the barred nature of the Milky Way was found by \\cite{Blitz91}. The analysis of COBE/DIRBE data \\citep{Dwek95, Binney97, Freudenreich98} supported Blitz \\& Spergel results. \\cite{Dwek95} found that the peanut-shaped bulge seen in the COBE image is best fitted by {\\it boxy} Gaussians functions. Since COBE, the direct evidence of the bar has increased based on photometric surveys and star counts in the near-infrared \\citep[see for instance][and references therein]{Lopez-Corredoira00, Lopez-Corredoira05, Babusiaux05}. Although, there is a consensus on the triaxial nature of the bulge, the actual shape and the inclination  is still under debate with currently no clear agreement among different authors. The inclination angles found range from $\\sim 10\\deg-15\\deg$ \\citep{Binney91, Freudenreich98, Lopez-Corredoira97} to $\\sim 30\\deg-45\\deg$ \\citep{Weiner99, Sevenster99, Lopez-Corredoira01} with a number of works giving values in the range $15\\deg-30\\deg$ \\citep{Mulder86, Binney97, Fux99, Englmaier99, Bissantz02, Bissantz03, Babusiaux05, Lopez-Corredoira05}.   In addition to the bar-like {\\it boxy}-bulge, some recent works have also proposed the existence of another bar that would be longer and thinner than the bulge \\citep{Picaud03,  Lopez-Corredoira01, Lopez-Corredoira07,  Benjamin05}. As the bulge, the near side of this long and thin bar will also be located in the first Galactic quadrant. In contrast, the long bar would be oriented at an angle of $\\sim 45\\deg$ with respect to the line of sight towards Galactic center. In the present paper we will not discuss this component as the mass should be dominated by the bulge. Hereafter, we will refer to the triaxial boxy bulge as ``large bar\" or simply ``bulge\".   Due to the difficulty to measure the stellar velocities, most of our knowledge of the Galactic dynamics come from spectroscopic observations of the atomic and molecular interstellar medium \\citep{Rougoor60, Scoville72, Liszt78, Bally87, Dame01}. The gas exhibit large non-circular motions in the inner Galaxy. Among the possible explanations for these kinematics, de Vaucouleurs (1964) proposed the gas response to a bar potential. This explanation has also been discussed in other observational works as those of \\cite{Kerr67}, \\cite{Bania77} or \\cite{Liszt78}. Nevertheless, it is the comparison of the observations with numerical simulations of the gas dynamics what have allowed to better understand the structure of the Galaxy and to derive important dynamical information as the pattern speed of the bar. The first attempt was that of \\cite{Mulder86} solving the gas-dynamical equations to get the gas flow in a weak bar for the whole Galaxy.   { A major step in our understanding of the gas kinematics in the Galactic Center was the paper by \\cite{Binney91}. These authors made a comparison of the CO(1-0) and HI longitude-velocity  diagrams (hereafter \\lvd) of \\cite{Bally87} and \\cite{Burton78}, with the $lv$ locus of closed stellar orbits in a bar potential (x1 and x2 orbits). The \\lvd \\  of the molecular gas in the inner $\\sim 4 \\deg$ of the Galaxy shows a parallelogram shape that resembles the $lv$ locus of the \\textit{cusped} x1 orbit for a bar with inclination of $\\sim 16 \\deg$. In spite of the globally successful interpretation of the CO and HI \\lvd s, the \\cite{Binney91} model cannot account for the observed asymmetry of the CO  parallelogram, which is lopsided towards positive longitudes while the modeled parallelograms are centered at longitude $0\\deg$. In addition the expected upper right vertex of the locus of the \\textit{cusped} x1 orbit has a velocity close to 0, while the upper right vertex of the observed CO parallelogram shown in Fig. 2 of \\cite{Binney91} has velocity of $\\sim 150$ \\kms. }  The qualitative insight on the gas dynamics that one could obtain by comparing the \\lvd~ of the gas with that of a few stellar orbits, should be tested performing simulations of the gas dynamics that take into account the dissipative nature of the gas clouds. This was done by \\cite{Jenkins94}, who tried to reproduce the \\lvd ~ of the CMZ using a sticky-particles code but with moderate success. { Using a combination of SPH and collisionless particles, \\cite{Lee99} were able to reproduce an inner ring with the right inclination in the \\lvd \\ and also some of the gas clouds with non-circular velocities in the range $l<|6\\deg|$.}  To our knowledge, the works by \\cite{Binney91}, \\cite{Jenkins94} and \\cite{Lee99} are the only dynamical models that have tried to interpret the dynamics of the complex Galactic center region (the central 500 pc). Other models of the Galactic dynamics have been mainly devoted to larger scales, from the works of \\cite{Mulder86} and Wada et al (1994) to the more recent works by \\cite{Fux99}, \\cite{Englmaier99}, \\cite{Weiner99} and \\cite{Bissantz03}. The bar pattern rotation speed for most of those models are in the range from 40 to 60 \\kmskpc.  Regarding the corotation radius, it is estimated to be in the range 3-4 kpc by some works \\citep{Englmaier99, Bissantz03, Habing06} and in the range 4-5 kpc by other works \\citep{Combes96, Sevenster99, Fux99}. In contrast, \\cite{Binney91} placed the corotation at 2.4 kpc and found a pattern speed of 80 \\kmskpc. As pointed out by \\cite{Combes96}, the discrepancy between Binney et al. estimation and other  estimations can be due to the fact that the molecular gas distribution in the Galactic center region used by Binney et al. could be the response to a small nuclear bar instead of the large  bar of the Galaxy.   Indeed, there are recent evidences of the presence of a distinct structure in the nuclear region of the Galaxy. The star counts map made by \\cite{Alard01} with 2MASS data shows an excess of counts in the inner four degrees of the Galaxy that resembles a small nuclear bar. This bar could be lopsided with the barycenter displaced to negative longitudes. \\cite{Nishiyama05} have also found evidence of a structure different to the large bar or bulge in the inner degrees of the Galaxy.  In the present paper, we want to investigate the gas response to the mass distribution inferred from the 2MASS data. In particular, we model the central mass distribution as a small nuclear bar. We are interested in the role of this nuclear bar to explain the differences among the results of previous works \\citep{Binney91, Fux99, Combes96}. We also want to study whether a lopsided nuclear bar could explain the apparent lopsidedness of the molecular gas distribution in the nuclear region, whose complex dynamics are far from being understood. Therefore, we have modeled the 2MASS star counts maps of Alard fitting different parametric functions using three components: disc, bulge/bar and nuclear bar. We use that mass distribution to determine the galactic stellar potential and to  model the gas dynamics using a N-body code that takes into account cloud collisions.  { Nested bars are common in external galaxies. As much as one third of barred galaxies have a secondary nuclear bar \\citep{Laine02}. In some cases the secondary bar is coupled with the primary bar \\citep{Shaw93} while in other cases the two bars rotate with different speeds \\citep[][]{Wozniak95}. The present paper is the first one dealing with the gas flow in the Milky Way nested bars and we will assume that both bars are dynamically coupled (they rotate with the same speed). As we discuss in Sect. 8, this seems a good assumption since our simulations explain for the first time some characteristics of the inner Galaxy. }  The paper is organized as follows: in the next section we summarize our understanding of the distribution and kinematics of the interstellar gas. In Sects. 3 and 4 we present the 2MASS star counts model and the fits to the data. The gas dynamics simulations are described in Sect. 5, and the results are presented in Sects. 6 an 7. We discuss the nature of the nuclear bar in Sect. 8. Our conclusions are summarized in Sect. 9.    \\begin{figure*}[ht!] \\centerline{\\includegraphics[angle=-90,width=16cm]{f1_dame.png}} \\caption{Longitude-velocity ($lv$) diagram of the CO(1-0) emission \\citep{Dame01}. The solid lines trace the position of some remarkable features as the locus of the spiral arms, the 3-kpc arm and the Connecting Arm. {  The black dashed lines indicate the contour of the Galactic Molecular Ring. The solid circles are the terminal velocities measurements of \\cite{Fich89} using CO while the triangles are the terminal velocities determined from the HI data of \\cite{Burton93}. The boxes mark the position of the Sagittarius, Scutum, Norma and Centaurus tangent points, located respectively at $l \\sim 50 \\deg ,\\ 30 \\deg, \\, -30\\deg$ and $-50\\deg$. The lines concerning the Nuclear Disk or Central Molecular Zone are blue but they are better shown in the next figure. } All these lines can be used to compare with the figures showing the simulations results. } \\label{fig:dame} \\end{figure*}   %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%5 ",
        "conclusions": "We have presented gas flow models in the mass distribution derived from 2MASS star counts using a model with three components (disk, bulge and nuclear bar). Our dynamical models are the first ones that include a central mass cusp (the nuclear bar) constrained by observations. For the first time, we have obtained good models of the Milky Way from the scales of the spiral arms to the Central Molecular Zone (CMZ) in the Galactic center (GC). The best models are found for  a bulge orientation  of $20-35\\deg$ with respect to the Sun-GC line and a pattern speed of 30-40 \\kmskpc. Which places the corotation radius at 5-7 kpc. The simulations reproduce: \\begin{itemize} \\item the spiral arms, giving in particular the good tangent point for the Carina arm. \\item the Galactic Molecular Ring (GMR), which is not an actual ring but the inner parts of the spiral arms and it is found approximately at the  position of the ultra-harmonic resonance \\item { the 3-kpc arm and its far-side counterpart, which are the lateral arms that surrounds the bar} \\item the HI ring, with a very similar size to that derived from the observations \\item the Central Molecular Zone (CMZ) \\end{itemize} In addition,  we postulate the existence of small spiral arms arising from the extremities of the CMZ and linking the HI ring.   In our simulations, the CMZ is the gas response to the nuclear bar. These simulations reproduce, for the first time,  the parallelogram shape of the \\lvd ~ of the CMZ. Using this shape we have been able to constrain the nuclear bar orientation, which is of $\\sim 60\\deg-75\\deg$ with respect to the Sun-GC line, in excellent agreement with the 2MASS fits and the results of \\cite{Sawada04}. We have also studied the observed asymmetry of the CMZ and we have concluded that it  cannot be due to the lopsidedness of the nuclear bar suggested by the 2MASS maps. Gas dynamical simulations in a lopsided potential cannot reproduce the observations, and in any case, we do not see a strong direct evidence in 2MASS data of an intrinsic lopsidedness of the stellar potential.  We propose an alternative scenario to explain the observed asymmetry of the CMZ. The asymmetry can be due to gas falling into the CMZ trough the inner spiral arm  in the {\\it l=1.3$\\deg$-complex} region. We have also discussed that the {\\it l=1.3$\\deg$-complex} shows  significant differences with the, otherwise symmetric, CMZ that extends from Sgr C to Sgr B. In particular, the  {\\it l=1.3$\\deg$-complex} shows all the signatures of the shocks expected if it is an interaction region where new material is falling into the CMZ.  { In our models both bars rotate with the same pattern speed. The success of these models in explaining many characteristics of the Milky Way as the Central Molecular Zone (CMZ) imply that the observations are compatible with that assumption. Furthermore, in our models the  nuclear bar is leading the large bar by $\\sim 55 \\deg$ as expected for coupled nested bars \\citep{Combes94}. Therefore, in spite of the high mass of the nuclear bar  (10-25 $\\%$ of the mass of the bulge) the dynamical decoupling my still not be effective. However, this question requires further studies and gas flow models with two pattern speeds or self-consistent gas and stars dynamics simulations.}"
    },
    "0806/0806.3703_arXiv.txt": {
        "abstract": "{ We discuss a new technique of studying magnetic fields in diffuse astrophysical media, e.g. interstellar and intergalactic gas/plasma. This technique is based on the angular momentum alignment of atoms and ions in their ground or metastable states. As the life-time of atoms in such states is long, the alignment induced by anisotropic radiation is susceptible to weak magnetic fields ($1{\\rm G}\\gtrsim B\\gtrsim0.1\\mu$G). The alignment reveals itself in terms of the polarization of the absorbed and emitted light. A variety of atoms with fine or hyperfine splitting of the ground or metastable states exhibit the alignment and the resulting polarization degree in some cases exceeds 20\\%. We show that in the case of absorption the polarization direction is either parallel or perpendicular to magnetic field, while more complex dependencies emerge for the case of emission of aligned atoms. We show that the corresponding studies of magnetic fields can be performed with optical and UV polarimetry. A unique feature of these studies is that they can reveal the 3D orientation of magnetic field. In addition, we point out that the polarization of the radiation arising from the transitions between fine and hyperfine states of the ground level can provide yet another promising diagnostics of magnetic fields, including the magnetic fields in the Early Universe. We mention several cases of interplanetary, circumstellar and interstellar magnetic fields for which the studies of magnetic fields using ground state atomic alignment effect are promising.}  \\addkeyword{Atomic processes} \\addkeyword{ISM: magnetic fields} \\addkeyword{polarization}  \\begin{document} ",
        "introduction": " ",
        "conclusions": "\\subsection{Range of applicability}  An incomplete list of objects where effects of alignment should be accounted for arises from our studies, which include this paper, as well as YLa,b. These include diffuse medium in the early Universe, quasars, AGNs, reflection nebulae, high and low density ISM, circumstellar regions, accretion disks and comets. One can easily add more astrophysical objects to this list. For instance, Io sodium tail can be studied the same way as sodium tail of comets.  In general, in all environment when optical pumping is fast compared with the collisional processes we expect to see effects of atomic alignment and magnetic realignment of atoms. The wide variety of atoms with fine and hyperfine structure of levels ensures multiple ways that the information can be obtained. Comparing information obtained through different species one can get deep insights into the physics of different astrophysical objects. If the implications of atomic alignment influenced the understanding of particular features of the Solar spectrum, then the studies of atomic alignment in a diffuse astrophysical media can provide much deeper and yet unforeseen changes in our understanding of a wide variety of physical processes.  As the resolution and sensitivity of telescopes increases, atomic alignment will be capable to probe the finer structure of astrophysical magnetic fields including those in the halo of accretion disks, stellar winds etc. Space-based polarimetry should provide a wide variety of species to study magnetic fields with.  \\subsection{Polarization of absorption lines} The direction of polarization has the same pattern, namely, either $\\parallel$ or $\\perp$ to the magnetic field on the plane of sky (YLa). In fact, this should be applicable to all absorption lines (including molecular lines) regardless of their different structures.  This fact is very useful in practice. It means that even we do not have an exact prediction and precise measurement of the degree of polarization of the absorption lines. We can have a 2D mapping of magnetic field on the plane of sky within an accuracy of $90^o$ once we observe their direction of polarizations. In this sense, it has some similarity with the Goldreich-Kylafis effect although it deals with radio emission lines.  For absorption lines, there is inevitably dilution along line of sight, which adds another dependence on the ratio of alignable column density and total column density $N_a/N_{tot}$. For different species, this ratio is different. The ratio should be close to 1 for highly ionized species which only exist near radiation sources. The same is true for the absorption from metastable state (see Paper I). Combining different species (with different $N_a/N_{tot}$), it is possible to acquire a tomography of the magnetic field {\\it in situ}.  \\subsection{Polarization of emission lines}  Compared to absorption lines, emission lines are more localized and therefore the dilution along the line of sight can be neglected. The disadvantage of emission lines compared to absorption lines is that the direction of the polarization of emission lines has a complex dependence on the direction of the magnetic field and the illusion light. Therefore, the use of emission lines is more advantageous when combined with other measurements.  We considered the situation that atoms are subjected to the flow of photon that excite transitions at the rate $\\tau^{-1}_R=B_{lu}{\\bar J}^0_0$ which is smaller than the Larmor precession rate $\\nu_L$, but larger than the rate of disalignment due to collisions $\\tau^{-1}_c$. For the cold gas with 30 hydrogen atoms per cubic cm, the characteristic range over which the atoms can be aligned by an O star is $\\lesssim 15$pc for HI due to spin exchange collisions (with rate $C_{10}/n_H=3.3\\times 10^{-10}{\\rm cm}^3{\\rm s}^{-1}$).  \\subsection{Radio Emission Lines Influenced by Optical Pumping}  {\\it Anisotropic} optical pumping also influences the magnetic dipole transitions between the sublevels of the ground state. We briefly discussed this in \\S\\ref{mag_dipole} and YLb,c. We showed that the emission arising from such atoms is polarized, which provides a new way of studying magnetic fields. Apart from apparent galactic and extragalactic applications, this may be an interesting process to study magnetic fields at the epoch of reionization, which hopefully will be available with the instruments are currently under construction. Atomic alignment has some similarity to Goldreich-Kylafis effect, which also measures magnetic field through magnetic mixing (Goldreich \\& Kylafis 1982; Girart, Crutcher, \\& Rao 1999). However, Goldreich-Kylafis effect is due to the radio pumping. Atomic alignment, on the other hand, happens with ground states as a result of optical and UV pumpings.  \\subsection{Comparison to using aligned grains}  In our studies we were focused on new ways to study magnetic field that atomic alignment of the atoms/ions with fine and hyperfine structure provides. Being alignable in their ground state, these species can be realigned in weak magnetic fields, which is extremely good news for the studies of weak magnetic fields in astrophysical diffuse gas. The latter studies are currently very limited, with polarimetry based on grain alignment  being the most widely used technique (see Whittet 2005 and ref. therein). However, in spite of the progress of grain alignment theory (see Lazarian 2007 for a review), the quantitative studies of magnetic fields with aligned atoms are not always possible. Atoms, unlike dust grains, have much better defined properties, which allows, for instance, tomography of magnetic fields by using different species\\footnote{As we discussed in YLa long-lived alignable metastable states that are present for some atomic species between upper and ground states may act as proxies of ground states. The life time of the metastable level may determine the distance from the source over which the atoms are aligned being on metastable level. Absorption from such metastable levels can be used as diagnostics for magnetic field in the star vicinity.}, which would be differentially aligned at different distances from the source. In addition, ions can trace magnetic fields in the environments in which grains cannot survive.  Incidentally, for the grain alignment, anisotropic radiation has also been identified as a major driving agent (see Dolginov  \\& Mytrophanov 1976, Draine \\& Weingartner 1996, Lazarian \\& Hoang 2007a). Another class of grain alignment mechanisms includes the mechanical alignment (see Gold 1952, Roberge et al. 1995, Lazarian \\& Hoang 2007b), which arises from streaming of grains in respect to gas. Interestingly enough, atomic alignment due to collisions with atoms and ions is feasible as well (see Fineschi \\& Landi Degl'Innocenti 1992). For instance, one may expect that atoms produced by charge exchange of directed flow of ions and atoms to be aligned.  \\subsection{Studying time variations of magnetic fields}  In fact, rather than compare advantages and disadvantages of different ways of studying magnetic fields, we would  stress the complementary nature of different ways of magnetic field studies. The subject is starved for both data and new approaches to getting the data. Note, that atomic realignment happens on the Larmor precession time, which potentially allows to study the dynamics of fast variations of magnetic fields, e.g., related to MHD turbulence.  Note, that for interplanetary studies, one can investigate not only spatial, but also temporal variations of magnetic fields. This can allow cost effective way of studying interplanetary magnetic turbulence at different scales."
    },
    "0806/0806.1530_arXiv.txt": {
        "abstract": "Red quasars are a population, characterized by significant extinction in UV, which could be explained by absorption of dusty gas on a scale of a few kpc. We show that the enhanced radiation-pressure drives the dusty gas to supersonically expand and produces shocks. The shocks energize electrons to be relativistic via the first Fermi acceleration. As a balance of shock acceleration and synchrotron emission and inverse Compton scattering, the maximum Lorentz factor of the electrons reaches as $\\sim 10^6$. The shocked interstellar medium appears as extended multiwavelength fuzz, in which synchrotron emission from the electrons peaks at near infrared or UV bands and inverse Compton scattering around 1.0GeV$-$0.1TeV. Future multiwavelength images of the fuzz would provide new clues to study the details of radiative feedback if red quasars could be a certain phase in evolutionary chains of galaxies. ",
        "introduction": "Radiative feedback from quasars to surroundings of their hosts has been realized for many years. Several effects have been suggested, including the Compton heating (Begelman 1985; Chang et al. 1987; Wang et al. 2005), acoustic wave heating (Chelouche 2007), proximity effects (Bajtlik et al. 1988; Goncalves et al. 2008; Wild et al. 2008; Prochaska \\& Hennawi 2008). The strong feedback due to radiation may quench the inflows fueling the central black hole (e.g. Binney \\& Tabor 1995) and star formation so as to arise the relation between black hole masses and dispersion velocities (Silk \\& Rees 1998; King 2003; Ciotti \\& Ostriker 2007; Fabian et al. 2006). However, the direct large scale evidence for feedback is scarce, and even absent in most cases.  Traditionally quasars are selected based on their bluer color, however, the existence of significant population of red quasars has been the subject of debate for a number of years. The percentage of red quasars is only $\\sim 6\\%$ or so from the large database of Sloan Digital Sky Survey (SDSS) (Richards et al. 2003), $\\sim 20-30\\%$ in the FIRST-2MASS red quasar survey (Glikman et al. 2007), and $\\sim 20-30\\%$ in {\\em Spitzer} survey (Brown et al. 2006; Lacy et al. 2007). The fraction of red quasars remains a big debate, however this might be due to heterogeneous definition or selection effects of surveys. There is increasing evidence for dust extinction as one of the major explanations for red quasars from SDSS (Richards et al. 2003), {\\em GALEX} data (Trammell et al. 2007), {\\em Spitzer} survey (Brown et al. 2006; Lacy et al. 2007) and {\\em HST} images (Urrutia et al. 2008) as well as from numerical simulations (Di Matteo et al. 2005; Hopkins et al. 2007). The debris, ejected from mergers or starburst triggering the central black holes, surrounds the nucleus prior to the galactic winds (Sanders et al. 1989; Yun et al. 2004; Fabian et al. 2006 and Hopkins et al. 2007), reddening quasars. The well-known roles of dust particles in outflows, as results of strong coupling between the charged grains with gas (e.g. Draine \\& Salpeter 1979), have been realized in the dynamics of dusty gas (Chang et al. 1987; Murray et al. 2005; Fabian et al. 2006). It is then expected that the swept medium by quasar radiation results in extinction of quasars.  The present Letter makes an attempt to connect the action of the radiative feedback to the phenomena of red quasars and predicts some observational features. We argue that radiation from quasar is partly converted to kinetic luminosity of supersonic outflowing gas and then partly tansformed to the kinetic energy of accelerated electrons via Fermi mechanism, producing extended fuzz with non-thermal emission from radio to $\\gamma-$rays around red quasars. ",
        "conclusions": "In this Letter, we show that the enhanced radiation pressure inevitably leads to strong shocks in the dusty gas of both homogeneous and clumpy ISM. The first Fermi shock acceleration results in production of relativistic electrons in the shocked ISM, from which produces non-thermal emission via synchrotron radiation and inverse Compton scattering, displaying multiwavelength fuzz from radio, infrared and $\\gamma-$ray bands. The current radio telescopes, MIPS, HESS and VERITAS are capable of detecting the fuzz. The multiwavelength image of the fuzz provides an interesting and new clue to understand the interaction between quasars and their environments.  The predicted $\\gamma-$ray emission from fuzz could have significant contribution to cosmic $\\gamma-$ray background if {\\em all} quasars have to undergo such a red phase during their cosmic evolution. We will issue this problem in the future."
    },
    "0806/0806.4409_arXiv.txt": {
        "abstract": "The extrasolar system OGLE-2006-BLG-109L is the first multiple-planet system to be discovered by gravitational microlensing~\\citep{Gaudi:2008Science}; the two large planets that have been detected have mass ratios, semimajor axis ratios, and equilibrium temperatures that are similar to those of Jupiter and Saturn; the mass of the host star is only $0.5 M_\\odot$, and the system is more compact than our own Solar system.  We find that in the habitable zone of the host star, the two detected planets resonantly excite large orbital eccentricities on a putative earth-mass planet, driving such a planet out of the habitable zone. We show that an additional inner planet of $\\gtrsim0.3\\Mearth$ at $\\lesssim0.1$ AU would suppress the eccentricity perturbation and greatly improve the prospects for habitability of the system.  Thus, the planetary architecture of a potentially habitable OGLE-2006-BLG-109L planetary system -- with two ``terrestrial'' planets and two jovian planets -- could bear very close resemblance to our own Solar system. ",
        "introduction": "\\label{s:Introduction} It is possible that the OGLE-2006-BLG-109L system harbors other planets, including earth-like planets, that are beyond the sensitivity of the microlensing observations.  Here we use dynamics to constrain the prospects for an earth-like habitable planet in this system.  We adopt \\cite{Kasting:1993Icarus}'s definition of the habitable zone (HZ) as the region in which an Earth-like planet with an N$_2$--CO$_2$--H$_2$O atmosphere would be able to maintain liquid oceans on its surface.  For an Earth-like planet, the inner boundary of the HZ is the distance from the star below which the planet would experience a runaway greenhouse and lose its oceans, and the outer boundary is the distance from the star beyond which the planet would condense CO$_2$ clouds.  For OGLE-2006-BLG-109L, the HZ is $\\sim0.25$--$0.36\\AU$ from the star.  For comparison, and using the same definition, the HZ in our solar system is in the 0.95--1.37 AU heliocentric distance range.  Both these estimates are reported in \\cite{Kasting:1993Icarus}, in their table III. ",
        "conclusions": ""
    },
    "0806/0806.4253_arXiv.txt": {
        "abstract": "We present a unified model of infrared (IR), optical, ultraviolet (UV), and X-ray light curves for the 1983 outburst of GQ Muscae (Nova Muscae 1983) and estimate its white dwarf (WD) mass.  Based on an optically thick wind model of nova outbursts, we model the optical and IR light curves with free-free emission, and the UV 1455 \\AA\\ and supersoft X-ray light curves with blackbody emission.  The best fit model that reproduces simultaneously the IR, optical, UV 1455 \\AA, and supersoft X-ray observations is a $0.7 \\pm 0.05~M_\\sun$ WD for an assumed chemical composition of the envelope, $X=0.35-0.55$, $X_{\\rm CNO} =0.2-0.35$, and $Z = 0.02$, by mass weight.   The mass lost by the wind is estimated to be $\\Delta M_{\\rm wind} \\sim 2 \\times 10^{-5} M_\\odot$. We provide a new determination of the reddening, $E(B-V) = 0.55 \\pm 0.05$, and of the distance, $\\sim 5$~kpc.  Finally, we discuss the strong UV flash that took place on JD 2,445,499 (151 days after the outburst). ",
        "introduction": "Despite their overall similarity, the optical light curves of classical novae show a wide variety of timescales and shapes \\citep[e.g.,][]{pay57}. Various empirical time-scaling laws have been proposed in the attempt to recognize common patterns and to unify the nova light curves.  For example, \\citet{mcl42} proposed a compression of the time scale obtained by normalizing the time to $t_m$, the time for the brightness to decrease by $m$ magnitudes (usually $m=2$ or 3).  The results were however unsatisfactory because slow novae reached their late stages relatively earlier than faster novae.  A different approach was adopted by \\citet{vor48} who, based on a large collection of nova light curves, found that $m(t)$, the visual magnitude at time $t$, is conveniently represented, on average, by $m(t)= m_0 + b_1 \\log (t-t_0)$ at early stages and $m(t)= m_1 + b_2 \\log (t-t_0)$ at later stages with $b_1 = 2.5$ and $b_2 > b_1$.  The underlying problem of both the above approaches is that, during the decline phase from maximum, the flux in the visual band, initially dominated by free--free emission, becomes more and more affected by the increasing contribution from the emission lines, which finally dominates.  The presence of these two competing and heterogeneous emission mechanisms is most likely the main cause for the difficulty to find a suitable normalization parameter for the nova light curves.  A more appropriate way to monitor the evolution of the visual continuum is to use the Str\\\"omgren $y$ filter, which is designed to avoid strong emission lines and, in particular, the [\\ion{O}{3}] line \\citep[e.g.,][]{kal86, loc76}.  In the case of V1668 Cyg, \\citet{kal86} has indeed shown that the continuum $y$ flux declines much faster than the H$\\beta+$[\\ion{O}{3}] fluxes.  Consequently, the decline rate is much faster in the $y$ band than in the visual band. This was demonstrated by \\citet{kal86}, who showed that the $y$ magnitude of V1668 Cyg declines with a slope of $b_1= 3.5$ at early stages and it then changes to $b_2= 6.5$ at later stages, values that are sensibly larger than those reported by \\citet{vor48}.  The observations of V1500 Cyg by \\citet{gal76} suggest that the visual and infrared continua are well represented by free-free emission during the early decay phase.  The same applies to GQ Mus: \\citet{kra84} reported that, during the early 40 days (diffuse enhanced and Orion-phases), the energy spectrum of $0.3 - 10\\mu$ was well represented by optically thin free-free emission \\citep[see also][]{din86}.  The flux in the optical and infrared appears to decay with time as $F_\\lambda \\propto t^{-\\alpha}$ in various speed classes of novae \\citep[e.g.,][for V1500 Cyg]{enn77, woo97}.  An interesting attempt to unify the light curves of novae was proposed by \\citet{ros99a, ros99b}, who suggested that the main parameter determining the shape of the light curve is the radius of the ejected shell, $\\log R_{\\rm shell} = \\log(t) + \\log (v)$, where $t$ is the time from optical maximum, and $v$ is the velocity of the ejecta.  He plotted various nova light curves in the $m-\\log R_{\\rm shell}$ diagram and found that light curves overlapped each other.  The rather large scatter of the data in that diagram is probably due to the fact that the expansion velocity of novae varies with time \\citep[see][]{cas05}.  \\citet{cas02} studied the time evolution of the UV continuum flux in twelve CO and ONeMg novae and found a strong correlation between the $t_3$ time and the time of maximum flux in the 1455 \\AA\\ continuum (see Fig. 4 in their paper).  In a subsequent paper, \\citet{cas05} studied the time evolution of the UV emission lines in seven CO novae and found a strong correlation between the line ionization potential and the time of maximum emission normalized to the $t_3$ (Fig. 5 in their paper).  These results strongly suggest that novae do indeed evolve following a common pattern being time-normalized by the $t_3$ time, which mainly depends on the white dwarf mass \\citep[e.g.,][]{liv92}.  Recently, \\citet[][hereafter referred as Paper I]{hac06kb} found that the visual and IR light curves of several novae follow a universal decline law, and interpreted that in terms of free-free emission \\citep[Paper I;][hereafter referred as Paper II]{hac07k}; in particular, the time-normalized light curves were found to be independent of the white dwarf mass, the chemical composition of ejecta, and wavelength.  They also showed that the UV 1455 \\AA\\ light curve, interpreted as blackbody continuum, can also be time-normalized by the same factor as in the optical and IR.  The authors, in the above quoted papers, determined the white dwarf mass and other parameters for a number of relatively well-observed classical novae.  In the present paper, we apply the above universal decline law to the classical nova GQ Muscae 1983. GQ Mus is the first classical nova in which the supersoft X-ray turnoff was detected with the X-ray satellite {\\it ROSAT} \\citep{oge93, bal98, ori01}.  This object was also well observed by the UV satellite {\\it IUE} \\citep[e.g.,][]{kra84} and in near-IR bands of {\\it JHKL} \\citep[e.g.,][]{whi84}.  The next section summarizes the basic observational characteristics of GQ Mus.  In \\S 3, we revisit the UV light curves of GQ Mus obtained with the {\\it IUE} satellite.  In \\S 4, we briefly introduce our method based on our optically thick wind model. Light curve fittings of GQ Mus in the X-ray, UV, optical, and near IR bands are shown in \\S 5.  Discussion and conclusions follow in \\S\\S 6 and 7, respectively.  Finally, in the Appendix, we assess the problem of the contribution of the emission lines to the visual and infrared photometric bands for GQ Mus and for other well known novae: V1500 Cyg, V1668 Cyg, and V1974 Cyg. ",
        "conclusions": "We have applied the ``universal decline law'' of classical novae described in \\S\\ref{model_nova_outburst} to GQ Muscae 1983 and derived various parameters of the nova. Our main results may be summarized as follows:  1. We  show that the ``universal decline law'' reproduces well the observed light curves of GQ Mus in the optical and in the near infrared $I$, $J$, $H$, and $K$ bands.  2. Our blackbody light curve model for the UV 1455 \\AA\\ band can reproduce the observed UV 1455 \\AA\\ fluxes except two UV flashes on day 37 and 151.  3. The UV flash on day 151 described in \\S\\ref{uv_evolution_cont} was accompanied by a mass ejection episode seen as a fast wind with a terminal velocity of about 3200 km~s$^{-1}$.  4. An analysis of the {\\it IUE} reprocessed data of GQ Mus indicates $E(B-V)= 0.55 \\pm 0.05$, a value that is larger than previously reported.  5. We find that the mass of the WD component of GQ Mus is $0.7 \\pm 0.05 ~M_\\sun$ for the adopted envelope chemical composition of $X=0.35-0.55$, $X_{\\rm CNO}= 0.2-0.35$.  This value has been derived by comparing predicted fluxes from our models with the observations in the supersoft X-ray, in the UV, and in the optical and near infrared, obtained at different times.  6. We have estimated a mass of $\\Delta M_{\\rm wind} \\sim 2 \\times 10^{-5} M_\\odot$ lost by the wind.  7. We have estimated a distance of $d \\sim 5$ kpc."
    },
    "0806/0806.1689_arXiv.txt": {
        "abstract": "Remnant planetesimals might have played an important role in reducing the orbital eccentricities of the terrestrial planets after their formation via giant impacts. However, the population and the size distribution of remnant planetesimals during and after the giant impact stage are unknown, because simulations of planetary accretion in the runaway growth and giant impact stages have been conducted independently. Here we report results of direct N-body simulations of the formation of terrestrial planets beginning with a compact planetesimal disk. The initial planetesimal disk has a total mass and angular momentum as observed for the terrestrial planets, and we vary the width (0.3 and 0.5AU) and the number of planetesimals (1000-5000). This initial configuration generally gives rise to three final planets of similar size, and sometimes a fourth small planet forms near the location of Mars. Since a sufficient number of planetesimals remains, even after the giant impact phase, the final orbital eccentricities are as small as those of the Earth and Venus. ",
        "introduction": "The rocky planets are believed to have formed via the accretion of small planetesimals. The formation mechanism of planetesimals and their initial mass and spatial distribution are however still controversial. The standard picture of accretion of terrestrial planets from planetesimals is as follows. During the early stages of planetesimal accretion, larger planetesimals grow faster than smaller ones owing to their stronger gravity \\citep{Gre78, WS89}. Through this runaway growth stage, a few tens of Mars-size protoplanets form with mutual separation of 10 Hill radii \\citep{KI98, KI02}. The growth time scale of protoplanets is estimated to be  $\\sim$ 0.1--1Myr around 1AU, and is longer with larger distance from the Sun \\citep{WS93, KI02}. As long as some amount of nebular gas and/or planetesimals remain, its damping effect stabilizes the orbits of protoplanets, preventing mutual collisions \\citep{Iwa02, KomI02}. As the amount of remnant gas and/or planetesimals decreases, the orbital eccentricities of protoplanets increases due to their mutual interactions. Eventually their orbits become chaotic and late time giant impacts occur. During this giant impact stage, whose time scale is considered to be $\\sim$ 100Myr, the current terrestrial planets form \\citep[e.g.][]{CW98, Agn99, Kok06}. The orbital eccentricities of planets immediately after giant impacts are likely to be much larger than those of the current terrestrial planets. Therefore, interactions with remnant gas and/or planetesimals is expected to reduce their eccentricities. Reviews for the processes described above are given by \\citet{Cha04,Nag07}.  Whether remnant gas or remnant planetesimals is more important for reducing eccentricities primarily depends on the time scale of gas dissipation. If the time scale is long, the velocity dispersion of planetesimals is suppressed by the gas drag. Hence, the gravitational focusing effect of protoplanets is enhanced, resulting in a fast clean up of remnant planetesimals and a lower eccentricity distribution of the final planets \\citep{AW02, KomI02, KomI04, Nag05, Ogi07}. On the other hand, if the time scale of gas dissipation is short, planetesimals remain unaccreted by protoplanets for a longer period of time. In this case, remnant planetesimals may be required to reduce the eccentricities \\citep{Cha01, Obr06, Ray06}. In this paper, we examine the latter scenario. In other words, we ignore the effects of the gas drag and the tidal interaction between a gas disk and protoplanets. The effects of gas will be investigated in a future study.  There have been several attempts to examine the effect of remnant planetesimals based on direct $N$-body simulations \\citep{Cha01, Obr06, Ray06}, and with simulations using a hybrid-code \\citep{KB06}. Direct $N$-body simulations usually adopt lunar to Mars size protoplanets surrounded by smaller planetesimals as initial conditions. Direct $N$-body simulations suggest that the eccentricities of final planets are further reduced as the total mass of planetesimals increases. Even for the same total mass of planetesimals, the damping effect is stronger with a larger number of smaller planetesimals \\citep{Obr06, Ray06}. However, the total mass and mass distribution of the remaining planetesimals are unknown, unless mass evolution in the runaway stage is followed.  The hybrid code of \\citet{KB06, BK06}, is able to follow planetary accretion through both the runaway and giant impact stages. In their code,  the mass and velocity distributions of planetesimals contained in multi annuli are solved by a statistical approach whereas orbits of protoplanets are solved using direct $N$-body calculations, that include the effect of interaction between planetesimals and protoplanets. Statistical approaches based on the local approximation produce consistent results with those obtained from direct $N$-body simulations in the runaway stage \\citep[e.g.][]{Ina01}. However, it is questionable if statistical approaches can accurately follow the late stage of planetary accretion because orbital eccentricities of remnant planetesimals are usually very large.  Here we report results of direct $N$-body simulations beginning with a planetesimal disk until the end of planetary accretion in the terrestrial region. We consider compact planetesimal disks (initial disk widths of $\\le 0.5$AU), whose total masses and total angular momenta are the same as those of the present terrestrial planets. These initial conditions are used since accretion simulations beginning from compact disks are usually computationally less expensive than those from wider disks.  This is the case even with the same initial number of particles, because of the rapid decrease in the number of particles through accretion. Another reason to adopt compact disks is that total angular momenta of final planetary systems obtained from most of previous simulations are much larger than for the terrestrial planets, as these simulations usually have a super-massive Mars \\citep[e.g.][]{Cha01, Ray06}. This excess angular momentum is likely due to initially extended disks. Though Jupiter removes angular momentum, mostly from the asteroid region, its effect does not seem to be sufficiently strong in the terrestrial region. As one possibility for this issue, we consider initially compact disks, supposing that they result from, for example, dust migration due to the gas drag prior to formation of planetesimals \\citep[e.g.][]{YS02}.  In \\S~2, we explain the numerical methods used in this study. We show results of simulations in \\S~3. We compare our results with previous simulations in \\S~4. In \\S~5, we give some physical interpretations for our simulation results using analytic estimations. We summarize our results in \\S~6. ",
        "conclusions": "We have conducted direct $N$-body simulations of the formation of terrestrial planets beginning with a compact planetesimal disk, with the total mass and angular momentum being those for the current terrestrial planets. In most of the cases, a planetesimal disk results in three planets of similar size, and sometimes a fourth small planet forms around the location of Mars. Since a sufficient number of planetesimals remain even after giant impacts of protoplanets, orbital eccentricities of the final planets are as small as those of the current Earth and Venus. This is a very nice success of our simulations which test the standard model for planet formation.  The final eccentricities of planets are nearly in the equilibrium state for the energy partitioning with remnant planetesimals, meaning that the mutual interactions of planets in the final state is unimportant in most cases. The final eccentricities of planets depends on the initial mass of planetesimals only very weakly, and on the surface density relatively strongly. These dependences are interpreted in terms of the effective mass of remnant planetesimals $(m_{\\rm eff} = \\langle m^2 \\rangle/\\langle m \\rangle)$, which determines the strength of the gravitational scattering effects of planetesimals. The mass distribution of remnant planetesimals is approximately represented by a power law distribution, $dn \\propto m^{q}dm$, with $q \\sim -2$, with the upper cut off mass (we call it the minimum oligarchic mass) which increases very weakly with the initial mass of planetesimals. Therefore, the dependence of the effective mass on the initial mass of planetesimals is very weak as well.  In a few of our simulations, planet-planet interactions are important even at the end of simulations and the orbits of final planets can be much more eccentric than for our terrestrial planets. This situation seems to appear when the gradient of population of planetesimals in the radial direction is large during the giant impact stage. However, the number of our simulations is still too small to statistically discuss the conditions that are responsible for final planetary orbits."
    },
    "0806/0806.4065_arXiv.txt": {
        "abstract": "\\hspace*{\\parindent} We consider cosmological perturbations caused by modulated inflaton velocity. During inflation, the inflaton motion is damped and the velocity is determined by the parameters such as couplings or masses that may depend on light fields(moduli). The number of e-foldings is different in different patches if there are spatial fluctuations of such parameters. Based on this simple idea, we consider ``modulated inflation'' in which the curvature perturbation is generated by the fluctuation of the inflaton velocity. This talk is based on our recent papers\\cite{Matsuda:Modulatedinflation3}. ",
        "introduction": "Let us first discuss common ideas for cosmological perturbations in terms of the $\\delta N$ formalism. Assuming that $H$ is a constant during inflation, the formula for $\\delta N$ is given by $\\delta N \\simeq H \\delta t$, where $\\delta t$ is the fluctuation related to the time passed after horizon crossing. According to the traditional inflationary scenario, the curvature perturbation is generated at the horizon crossing\\footnote{See the first picture in Fig.\\ref{fig:normal}.}. Another idea for the cosmological perturbation is ``at the end'' scenario\\cite{Attheend}\\footnote{See the second picture in Fig.\\ref{fig:normal}.}. In the `` at the end'' scenario, the curvature perturbation is generated at the end, where the fluctuation of the goal-line is induced by the light field (${\\cal M}$); $\\delta \\phi_e \\simeq \\phi_e' \\delta {\\cal M}$. Of course, considering the fluctuation of the distance ($\\delta\\phi$) to obtain $\\delta t$ is very natural, as it obviously causes the fluctuation of the number of e-foldings $\\delta N$. However, we know that the time elapsed after horizon crossing $\\sim |t_N-t_e|$ depends not only on the distance $\\sim|\\phi_N-\\phi_e|$ but also on the inflaton velocity $\\dot{\\phi}$. Therefore, if the inflaton velocity depends on a light field, the fluctuation of the light field may lead to the fluctuation of the inflaton velocity; $\\delta\\dot{\\phi}\\sim (\\dot{\\phi})'\\delta {\\cal M}$, which eventually causes $\\delta t$ and $\\delta N$. The effect would be obvious when (1) many massless degrees of freedom appear in the inflaton trajectory, or (2) the inflaton mass is slightly larger than the Hubble parameter and there is no significant perturbation from $\\delta \\phi_N$. Based on the above simple speculation, we will focus on modulated\\cite{Kofman:modulated} inflaton velocity, where spatial fluctuation of the time elapsed during inflation $\\delta t$ is caused by the fluctuation of the inflaton velocity.\\footnote{See the red-dotted line in the last picture in Fig.\\ref{fig:normal}.} Moreover, if the fundamental parameters are determined by the moduli in the underlying (string) theory, it would be natural to think that the Planck scale may also depend on moduli. What happens if the Planck scale is determined by moduli fields that may have fluctuations during inflation? We will see that the perturbation related to $\\delta M_p(\\cal M)$ is explained in terms of $\\delta \\dot{\\phi}$. Let us see more details of this ``modulated velocity scenario'', considering simple examples. \\begin{figure}[ht] \\begin{center} \\begin{picture}(400,250)(0,0) \\resizebox{15cm}{!}{\\includegraphics{susy08-1.eps}} \\end{picture} \\caption{In the first(Traditional) and the second(At the end) scenarios, $\\delta t$ is caused by the fluctuation of the distance between the start-line at $\\phi_N$ and the goal-line at $\\phi_e$, while in the last scenario $\\delta t$ is caused by $\\delta \\dot{\\phi}$.} \\label{fig:normal} \\end{center} \\end{figure}  \\subsection{Modulated velocity from the inflaton potential}  For the first example, we consider standard kinetic terms and conventional interaction term in the potential, $V(\\phi, {\\cal M})$. In this case, the definition of the number of e-foldings for constant $H$ is \\begin{equation} N = \\int H dt =\\int H\\frac{\\dot{\\phi}d\\phi +\\dot{\\cal M} d{\\cal M}}{\\dot{\\phi}^2 + \\dot{\\cal M}^2}. \\end{equation} If there is no bend in the trajectory\\cite{multi-inflation}, the perturbation related to the inflaton velocity is expanded as \\begin{equation} \\label{delta-vel} \\delta N \\simeq -\\int^{\\phi_N}_{\\phi_e}\\frac{H}{\\dot{\\phi}^2}\\left( \\delta\\dot{\\phi}-\\dot{\\phi}A\\right)d\\phi, \\end{equation} where we consider linear scalar perturbations of the metric, $ds^2 = -(1+2 A)dt^2 + 2aB_idx^idt + a^2\\left[(1-2\\psi)\\delta_{ij}+2E_{ij} \\right]dx^i dx^j$. From the conventional energy and momentum constraints we can find \\begin{equation} \\dot{\\phi}\\left( \\dot{\\delta \\phi}-\\dot{\\phi}A\\right) -\\ddot{\\phi}\\delta\\phi \\simeq \\dot{\\phi}\\left(\\dot{\\delta \\phi}-\\dot{\\phi}A\\right) \\propto \\frac{k^2}{a^2}, \\end{equation} which suggests that the factor $\\dot{\\delta \\phi}-\\dot{\\phi}A$ appearing in the delta-N formula decays after horizon crossing. Therefore, the factor $e^{-2Ht}$ must be included in the calculation, even if the perturbation $\\dot{\\phi}\\delta\\dot{ \\phi}$ itself is supposed to be a constant. Of course, $e^{-2Ht}$ in the integral does not lead to the exponential suppression after the integration. It is easy to see that there is no exponential suppression for large $N$($\\sim Ht$), \\begin{equation} \\int^{t_e}_0 \\delta C  H e^{-2Ht} dt \\simeq \\frac{1}{2} \\delta C, \\end{equation} where the actual suppression factor is not significant. In fact, the corrections from terms proportional to $k^2/a^2$ have been disregarded in previous studies, since $k^2/a^2$ is obviously small at a distance. However, if they appear in the equation of $\\dot{\\cal R}$, these terms may yield significant correction to ${\\cal R}$ after integration, as we can see easily from the above equation.  \\subsection{Modulated velocity from the inflaton kinetic term}  Next, we consider moduli-dependent kinetic term $\\sim \\frac{1}{2} \\omega({\\cal M}) g^{ab}\\phi_{,a}\\phi_{,b}$. In this case, the number of e-foldings is \\begin{equation} N = \\int H dt \\simeq \\int \\frac{H}{\\dot{\\phi}}d\\phi\\simeq \\int \\frac{3H^2}{V_\\phi} \\omega d\\phi. \\end{equation} Again, we assume no bend. $\\delta N$ that is caused by the moduli perturbation is \\begin{equation} \\delta N \\simeq \\int \\frac{3H^2}{V_\\phi} \\omega' \\delta {\\cal M} d\\phi\\simeq \\frac{\\omega' N}{\\omega}  \\delta {\\cal M}. \\end{equation} Note that unlike the perturbation caused by the potential, the constraints from the energy and momentum does not yield $\\frac{k^2}{a^2}$ factor for the perturbation related to the kinetic term\\cite{kinetic-term}. ",
        "conclusions": "We considered cosmological perturbations caused by modulated inflaton velocity. The velocity perturbation has been disregarded in previous studies for multi-field inflation. However, the perturbation may lead to significant results, as we have shown in this presentation. The sources of such perturbations are clear in the $\\delta N$ formalism. Important results are: \\begin{itemize} \\item Small deviation from the standard perturbation may be explained by the modulated velocity. \\item The curvature perturbation can be generated after horizon crossing even if the both boundaries are completely flat. \\item It is possible to add significant non-gaussianity to the conventional perturbations after horizon crossing. \\end{itemize}"
    },
    "0806/0806.3473_arXiv.txt": {
        "abstract": "We present a radial velocity study of the rapidly rotating B-star Regulus that indicates the star is a single-lined spectroscopic binary.  The orbital period (40.11~d) and probable semimajor axis (0.35 AU) are large enough that the system is not interacting at present.  However, the mass function suggests that the secondary has a low mass ($M_2 > 0.30 M_\\odot$), and we argue that the companion may be a white dwarf.  Such a star would be the remnant of a former mass donor that was the source of the large spin angular momentum of Regulus itself. ",
        "introduction": "%  Regulus ($\\alpha$ Leo; HD~87901; HR~3982; HIP~49669) is a nearby ($d=24.3\\pm0.2$ pc; \\citealt{van07}) intermediate mass star of spectral type B7~V \\citep{joh53} or B8~IVn \\citep{gra03}. It is one of a number of nearby B- and A-type stars exhibiting extremely fast rotation.  The very broad shape of its photospheric absorption lines indicates a projected rotational velocity of $V\\sin i = 317 \\pm 3$ km~s$^{-1}$ \\citep{mca05}.  The full picture of its fast spin came with the first interferometric observations of Regulus with the CHARA Array optical long baseline interferometer \\citep{mca05}.  These observations showed that the star is rotationally flattened and gravity darkened at its equator.  Models of the spectrum and interferometry demonstrate that the star has a rotation period of 15.9~hr, with an equatorial velocity equal to $86\\%$ of the critical velocity, where centripetal acceleration balances gravity.  The fast spin of Regulus is puzzling given its probable age ($\\approx 150$ Myr; \\citealt*{ger01}).  Stars born as fast rotators are expected to slow down relatively quickly after birth, and only again achieve rapid rotation at the conclusion of core H-burning through a redistribution of angular momentum \\citep{eks08}.  Thus, it is surprising to find rapid rotation in Regulus, a star which is still in the middle of its core H-burning stage.  On the other hand, stars that are members of interacting binaries can experience large changes in spin due to tidal interactions and mass exchange.  \\citet{lan08} discuss how mass transfer may lead to the spin up of the mass gainer in a large fraction of these binaries.  Depending on the initial separation and mass ratio, the system may merge or it may widen following mass ratio inversion, leaving the donor remnant in a large orbit that shuts down mass transfer once the donor's envelope is lost.  We know of several examples of such post-mass transfer binaries with rapid rotators, including the Be X-ray binaries with neutron star companions \\citep{coe00} and Be binaries with He-star companions \\citep{gie98,mai05,pet08}.  Regulus does have a known wide companion ($\\alpha$~Leo~B at a separation of $\\approx 175\\arcsec$, which is itself a binary consisting of K2~V and M4~V pair; \\citealt{mca05}), but this companion has far too great a separation to have ever interacted directly with Regulus.  There are no known closer companions, but the last significant radial velocity investigation was made in 1912 -- 1913 by \\citet{mel23}.  The scatter in the results introduced by the broad and shallow appearance of the spectral lines may have discouraged other investigators, but this early work and others \\citep*{mau92,fro26,cam28,pal68} suggest that any velocity variations present are relatively small.  However, a low mass donor remnant would probably create only a modest reflex motion in Regulus, so the lack of demonstrated variability is not unexpected.  We have made spectroscopic observations of Regulus on many occasions over the last few years, and here we present a summary of the velocities measured in these spectra.  We find that Regulus is in fact a low amplitude, single-lined, spectroscopic binary, and we discuss the possible nature of the companion. ",
        "conclusions": "%  The orbital variation has a small semiamplitude that eluded detection in earlier studies.  Consequently, the derived mass function is also small (Table~3), and we show in Figure~2 the constraints on the possible masses from the mass function. \\citet{mca05} used model fits to derive a probable mass of the primary star of $M_1 = 3.4 \\pm 0.2 M_\\odot$, and the boundaries of this range are indicated by the vertical dotted lines.  Larger orbital inclinations are favored for random orientations, and it is possible that the orbital inclination is comparable to the spin inclination of Regulus, $i\\approx 90^\\circ$ \\citep{mca05}.  Thus, the mass of the companion may be close to the minimum mass shown (for $i= 90^\\circ$) of $M_2 > 0.30\\pm 0.01 M_\\odot$.  \\placefigure{fig2}     %  A companion this small may be a low mass white dwarf or main sequence star.  If Regulus was spun up by mass transfer in an interacting binary, then the remnant of the donor star is probably a low mass white dwarf \\citep{rag01,wil04}. Indeed, the lowest mass white dwarfs are usually found in binary systems \\citep*{mar95} where they lost a significant fraction of their mass, and some reach masses as low as $0.17 M_\\odot$ \\citep{kil07}.  Models by \\citet{wil04} for mass transfer during H-shell burning (their ``evolutionary channel 1'') often lead to remnant and gainer masses and orbital periods similar to the case of Regulus.  Since Regulus and its wider companion $\\alpha$~Leo~B are not too old ($<150$ Myr; \\citealt{ger01}), a white dwarf companion would not have progressed too far along its cooling track and would probably have an effective temperature $>16000$~K \\citep*{alt01}, much higher than that of the primary B7~V star.  Consequently, we might expect to observe a modest FUV flux excess if the companion is a white dwarf, and, in fact, \\citet{mor01} find that the spectral energy distribution of Regulus is about a factor of two brighter in the 1000 -- 1200~\\AA ~range than predicted by model atmospheres for a single B7~V star.  We note for completeness that a neutron star companion is probably ruled out because it would require a very small inclination (see Fig.~2) and an unlikely evolutionary scenario in which only a modest amount of mass transfer occurred before the supernova explosion (no more than the present mass of the primary).  On the other hand, the companion could be a low mass, main sequence star (making Regulus one of the most extreme mass ratio binaries among the massive stars after exclusion of the massive X-ray binaries).  Adopting the minimum mass, the companion would be an M4~V star that would be too faint to alter significantly the spectral energy distribution from that for the primary alone (see Fig.~4 in \\citealt{mca05}). It is very unlikely that the companion is the outer component of a once compact triple system. If the fast spin of Regulus is the result of prior mass transfer, then an M4~V star in the orbit we find would have been too close to the central binary for orbital stability (and would have likely been ejected or become a merger product).  Thus, in binary models for rapid rotation, the companion cannot be a low mass main sequence star but must be the remnant of the donor.  The angular separation of the binary is probably small. If we assume masses of $M_1=3.4 M_\\odot$ and $M_2=0.3 M_\\odot$, then according to Kepler's third law, the semimajor axis is 0.35~AU.  Thus, for a distance of 24.3~pc, the maximum angular separation will be approximately 15 mas, too small for detection by speckle interferometric or adaptive optics techniques. The binary could be resolved in principle by lunar occultation methods or optical long baseline interferometry, but the fact that there is no evidence of a binary from these methods is consistent with the expected faintness of the companion \\citep*{han74,rad81,rid82,mca05}. For example, the magnitude difference in the $K$-band is probably close to $\\triangle m \\approx 10$ and 6 mag for the cases of a white dwarf and an M4~V star companion, respectively.   Thus, the flux of the companion has no influence on the analysis of the interferometry presented by \\citet{mca05}.  Our study has led to the discovery of a binary companion to the twenty second brightest star in the sky, and it may, like Sirius, offer another example of a bright star that is orbited by a faint white dwarf.   If the companion is a white dwarf, then it may be the closest case of a star stripped to its core by mass transfer in a close binary. The best opportunity to test the white dwarf hypothesis will come from very short wavelength observations where a hot companion may outshine the B7~V primary."
    },
    "0806/0806.4854_arXiv.txt": {
        "abstract": "The acceleration of charged particles (electrons and protons) in flaring solar active regions is analyzed by numerical experiments. The acceleration is modelled as a stochastic process taking place by the interaction of the particles with local magnetic reconnection sites via multiple steps. Two types of local reconnecting topologies are studied: the Harris-type and the X-point. A formula for the maximum kinetic energy gain in a Harris-type current sheet, found in a previous work of ours, fits well the numerical data for a single step of the process. A generalization is then given approximating the kinetic energy gain through an X-point. In the case of the multiple step process, in both topologies the particles' kinetic energy distribution is found to acquire a practically invariant form after a small number of steps. This tendency is interpreted theoretically. Other characteristics of the acceleration process are given, such as the mean acceleration time and the pitch angle distributions of the particles. ",
        "introduction": "\\label{S-Introduction} The study of the energetic particle acceleration process during solar flares still remains an open and challenging problem for solar physics. Solar flares are the manifestation of the energy release process in the solar corona and atmosphere. It is well established that during this energy release process the free magnetic energy is converted, through magnetic reconnection, into heating, bulk motion of the flaring plasma and particle acceleration.  In the framework of the theory of magnetic reconnection process during solar flares, a number of works in the literature have addressed the problem of particle acceleration. Several authors studied a magnetic field topology with an induced homogeneous electric field, with both fields constant in time, using analytical methods or numerical integration of test particles.  The magnetic field topology is either an X-point with or without the presence of a guide field (Bulanov 1980; Deeg, Borovsky and Duric 1991; Bruhwiler and Zweibel 1992; Moses, Finn and Ling 1993; Mori, Sakai and Zhao 1998; Browning and Vekstein 2001; Zharkova and Gordovskyy 2005; Hannah and Fletcher 2006) or a Reconnecting Current Sheet (RCS) (Martens 1988; Martens and Young 1990; Litvinenko 1996; Zharkova and Gordovskyy 2004). The purpose is to study the final kinetic energy distribution of accelerated particles, the condition of adiabaticity of their orbits and the level of charge separation at the ejection points. More realistic steady X-point topologies were derived using the MHD equations (Craig and Litvinenko 2002; Heerikhuisen, Litvinenko and Craig 2002).  In order to model the burst effects many authors incorporate a time dependent electric field in X-points (Hamilton \\etal\\ 2003; 2005; Petkaki and MacKinnon 1997; 2007) or adopted a numerical code as Tajima \\etal\\ (1987). Wood and Neukirch (2005) adopted an X-point with a nonhomogeneous electric field that was stronger at the center of the X-point. Finally, Kliem (1995) used O-points and X-points combinations to study the effect of the fragmentation of RCS on particles acceleration. On the other hand, several studies (e.g. Benz \\etal\\ 1994; Saint-Hilaire and Benz 2002) show that more than $40 \\%$ of the released energy in solar flares is deposited to high energetic particles, indicating the close relation between the particle acceleration process and the way (and amount) of the energy released in these highly energetic events. Despite this, a very common approach to building acceleration models (see Miller \\etal\\ 1997; Anastasiadis 2002 for reviews) is to consider that the different processes (i.e. energy release, acceleration, transport and radiation) are decoupled. The main reason for such a consideration is the very different temporal and spatial scales on which the different processes evolve. This difference notwithstanding, it should not be neglected that all the processes are interwoven in a way rendering necessary to develop {\\it global models} for solar flares, i.e., models taking into account the interplay between the individual processes.  Parker (1998) first proposed that the free magnetic energy could be released in the solar corona through {\\it multiple dissipation sites}. This assumption implies essentially the fragmentation of the energy release in both space and time during solar flares. As several observations, either related to solar flare parameters (e.g. Dennis 1985; Benz 1985; Crosby, Aschwanden and Dennis 1993; Crosby \\etal\\ 1998; Aschwanden \\etal\\ 1995; 2000) or to emission processes in the solar corona (e.g. Mercier and Trottet 1997; Krucker and Benz 1998), can be interpreted on the basis of such a scenario, several models of particle acceleration were developed that incorporate concepts from the general theory of dynamical systems, such as the {\\it complexity} of the energy release process. For example, Anastasiadis and Vlahos (1991,1994) assumed the presence of multiple shock waves acting on the particles. Such shocks have the form of small-scale, short-lived magnetic discontinuities inside a flaring active region. Another possibility is the acceleration of particles by multiple DC electric fields (Anastasiadis, Vlahos and Georgoulis 1997; Anastasiadis \\etal\\ 2004; Vlahos, Isliker and Lepreti 2004; Dauphin, Vilmer and Anastasiadis 2007). Such fields are associated with the variety of dissipation sites inside an active region. A simulation of the energy release process can then be made by a cellular automaton (CA), the dynamical substrate of which incorporates the assumption that the active region evolves according to the dynamical laws of a system being on the status of `self-organizing criticality' (see Isliker, Anastasiadis and Vlahos 2001 for details on CA models for solar flares).  In the solar corona, large-scale and long lasting current sheets, in which a large number of particles could be accelerated, are unlikely to be present. Recent MHD simulations show that large current sheets are not stable for a long time and are quickly fragmented (Onofri, Isliker and Vlahos 2006). Furthermore, Hughes \\etal\\ (2003) have shown that solar flares can be considered as cascades of reconnecting small-scale magnetic loops inside an overall simple large magnetic field topology. In other words, the large-scale magnetic field topology probably determines the location of the magnetic energy release regions only on the scale of the active region. In addition, the complex magnetic environment of an active region together with the turbulent photospheric motions should be taken into account, since these external factors drive the system continuously by adding new stresses to the existing large-scale topologies. In particular, such a type of driving leads necessarily to the formation of multiple reconnecting current sheets (RCS) configurations.  In the present work, we consider the above description as the basis for constructing numerical simulations of the acceleration of particles (electrons and protons) within multiple 3D current sheets developing in an active region. Our basic model is a sequence of encounter events of a population of particles with local current sheets with physical parameters changing stochastically. In order to understand the outcome of such a process, it is important to study first just one step of the process, namely the interaction of the particle with just one RCS. This was the basic motive behind two recent papers of ours (Efthymiopoulos, Gontikakis and Anastasiadis 2005; Gontikakis, Efthymiopoulos and Anastasiadis 2006) in which we explored the particle dynamics and acceleration in a single Harris-type RCS configuration by numerical and analytical means. A basic outcome of our study is an analytical formula yielding the maximum possible kinetic energy gain of the particles passing through the RCS as a function of the particles' initial kinetic energy and of the physical parameters of the current sheet. Our formula generalizes previous formulae given by Speiser (1965) and Litvinenko (1996), containing the  later as asymptotic limits. In the present paper we further generalize this formula to an approximate formula valid for single X-point topologies. Following the assessment of this `one-step' process, we then pass to modelling the acceleration of particles inside a complex flaring  active region as a stochastic process. In this we repeat the methodology used in Gontikakis, Anastasiadis and Efthymiopoulos (2007) for Harris-type RCSs, and we provide a further theoretical analysis of the results obtained there in.  In summary, the simulations are done as follows: particles are initially interacting with a single RCS. This changes their kinetic energy distribution. After this interaction, the particles are considered to perform a free flight following the magnetic field lines until they reach another RCS present in the active region. This process is in principle repeated ad infinitum. Nevertheless, one finds that the kinetic energy distribution tends to acquire a limiting form following a few such steps, after which no appreciable change is seen in the distribution. This fact was observed in our previous study and we here provide a theoretical explanation for it. Furthermore, we expand our study by considering also multiple encounters of the particles with local X-point reconnecting magnetic field configurations. In this we find again the tendency of the kinetic energy distribution to reach a limiting form. We thus conclude that this phenomenon is probably generic, i.e., independent of the details of the reconnection topology.  In the next section (Section~\\ref{S-Set}) we outline the basic characteristics of our numerical set up. In Section~\\ref{S-Single} the interaction of particles with a single RCS is studied, followed by our results for the multiple encounters (Section~\\ref{S-Multiple}). Section~\\ref{S-Summary} summarizes our results. ",
        "conclusions": "\\label{S-Summary}  In the present work we studied the acceleration of electrons and protons in a multiple particle-RCS interaction scenario, assuming two different reconnection topologies, the Harris-type current sheet and the X-point. Our main conclusions are the following:  1) The particles accelerated close to a X-point, with specific field parameters ($\\xi_\\perp, \\xi_\\parallel, \\epsilon$) gain a smaller amount of kinetic energy than particles accelerated through a Harris-type current sheet with the same model parameters. This is because particles accelerated through an X-point encounter a stronger average perpendicular magnetic field component.  2) The particles escape from a Harris-type current sheet in longer times than from an X-point.  3) A previously derived analytical formula (Equation~\\ref{emax}), yielding the maximum kinetic energy gain for particles accelerated through a Harris type topology (Gontikakis, Efthymiopoulos and Anastasiadis 2006), provides a satisfactory approximation to the kinetic energy gains computed numerically. (Equation~\\ref{emax}) can also be modified to yield an estimate for the kinetic energy gain of particles accelerated by an X-point if the perpendicular component $\\xi_\\perp$ is replaced by an effective field $<\\xi_\\perp>$ which is the mean perpendicular magnetic field along an orbit.  4) A number of similarities are also found for the acceleration of particles through Harris-type RCSs and X-points. An important fraction of protons with initially thermal distributions are not accelerated at all. The pitch angle distributions of the accelerated electrons and protons, present sharp peaks for angles $<10$\\degr.  5) When the particles encounter multiple RCSs with parameters obtain randomly from a uniform sample, the particles kinetic energy distribution tends to acquire a limiting form after some iterations. This tendency was previously observed (Gontikakis, Anastasiadis and Efthymiopoulos 2007) in the case of multiple encounters with Harris-type RCSs, and we here confirm it also in the case of X-points. A theoretical explanation is given by considering the probability of an RCS to act as efficient accelerator (parameterized by the percentage $w_0$ of kinetic energy gain with respect to some initial energy $E_0$). In particular we demonstrate analytically that this probability decreases as $E_0$ increases.   \\begin{acks} We wish to thank Dr. I. Contopoulos for stimulating discussions on the problem of magnetic reconnection as well as the anonymous referee and Dr. K-L. Klein for their comments and suggestions that improved our manuscript. The work of CG and CE was supported in part by the Research Committee of the Academy of Athens. AA would like to thank the members of the Particle Acceleration Working Group of the CESRA-2007 workshop for many stimulating discussions. \\end{acks}"
    },
    "0806/0806.3968_arXiv.txt": {
        "abstract": "We present point-source catalogs for the $\\approx2$~Ms exposure of the \\chandra\\ Deep Field-South (\\hbox{CDF-S}); this is one of the two most-sensitive \\hbox{X-ray} surveys ever performed. The survey covers an area of $\\approx436$ arcmin$^{2}$ and reaches on-axis sensitivity limits of $\\approx1.9\\times10^{-17}$ and $\\approx1.3\\times10^{-16}$ \\flux\\ for the \\hbox{0.5--2.0} and 2--8~keV bands, respectively. Four hundred and sixty-two \\hbox{X-ray} point sources are detected in at least one of three \\hbox{X-ray} bands that were searched; 135 of these sources are new compared to the previous $\\approx1$~Ms \\hbox{CDF-S} detections. Source positions are determined using centroid and matched-filter techniques; the median positional uncertainty is $\\approx0\\farcs36$. The X-ray--to--optical flux ratios of the newly detected sources indicate a variety of source types; $\\approx$55\\% of them appear to be active galactic nuclei while $\\approx$45\\% appear to be starburst and normal galaxies. In addition to the main \\chandra\\ catalog, we provide a supplementary catalog of 86 \\hbox{X-ray} sources in the $\\approx2$~Ms \\hbox{CDF-S} footprint that was created by merging the $\\approx250$~ks Extended \\chandra\\ Deep Field-South with the \\hbox{CDF-S}; this approach provides additional sensitivity in the outer portions of the \\hbox{CDF-S}. A second supplementary catalog that contains 30 \\hbox{X-ray} sources was constructed by matching lower significance \\hbox{X-ray} sources to bright optical counterparts ($R<23.8$); the majority of these sources appear to be starburst and normal galaxies. The total number of sources in the main and supplementary catalogs is 578. $R$-band optical counterparts and basic optical and infrared photometry are provided for the \\hbox{X-ray} sources in the main and supplementary catalogs. We also include existing spectroscopic redshifts for 224 of the X-ray sources. The average backgrounds in the 0.5--2.0 and 2--8~keV bands are 0.066 and 0.167~counts~Ms$^{-1}$~pixel$^{-1}$, respectively, and the background counts follow Poisson distributions. The effective exposure times and sensitivity limits of the \\hbox{CDF-S} are now comparable to those of the $\\approx2$~Ms \\chandra\\ Deep Field-North (\\hbox{CDF-N}). We also present cumulative number counts for the main catalog and compare the results to those for the \\hbox{CDF-N}. The soft-band number counts for these two fields agree well with each other at fluxes higher than $\\approx2\\times10^{-16}$~\\flux, while the \\hbox{CDF-S} number counts are up to $\\approx25\\%$ smaller than those for the \\hbox{CDF-N} at fluxes below $\\approx2\\times10^{-16}$~\\flux\\ in the soft band and $\\approx2\\times10^{-15}$~\\flux\\ in the hard band, suggesting small field-to-field variations. ",
        "introduction": "One of the greatest successes of the {\\it Chandra X-Ray Observatory} (\\chandra) has been the characterization of the sources creating the 0.5--8 keV cosmic \\hbox{X-ray} background (CXRB), and the deepest \\chandra\\ surveys form a central part of this effort. The two deepest \\chandra\\ surveys, the \\chandra\\ Deep Field-North and \\chandra\\ Deep Field-South (\\hbox{CDF-N} and \\hbox{CDF-S}, jointly CDFs; see \\citealt{Brandt2005} for a review), have each detected hundreds of \\hbox{X-ray} sources over $\\approx 450$~arcmin$^2$ areas with enormous multiwavelength observational investments. They have measured the highest sky density of accreting supermassive black holes (SMBHs) to date and have also enabled novel X-ray studies of starburst and normal galaxies, groups and clusters of galaxies, large-scale structures in the distant universe, and Galactic stars.  As part of an effort to create still deeper \\hbox{X-ray} surveys, we proposed for substantial additional exposure on the \\hbox{CDF-S} during \\chandra\\ Cycle~9. The \\hbox{CDF-S} has superb and improving coverage at optical, infrared, and radio wavelengths; it will continue to be a premiere multiwavelength deep-survey field for the coming decades as additional large facilities are deployed in the southern hemisphere. Furthermore, owing to the 1~Ms of \\chandra\\ exposure already available \\citep[][hereafter G02]{Giacconi2002}, the \\hbox{CDF-S} is a natural field to observe more sensitively. Although our proposal was not approved in the peer review, subsequently 1~Ms of Director's Discretionary Time was allocated for deeper \\hbox{CDF-S} observations. The allocated observations were successfully executed in 2007 September, October and November, raising the \\hbox{CDF-S} exposure to $\\approx 2$~Ms and improving its sensitivity to be comparable to that of the \\hbox{CDF-N} \\citep[e.g.,][hereafter A03]{Alexander2003}. Additional sky coverage at such flux levels is critically important as it substantially improves the statistical sample sizes of the faintest X-ray sources and also allows a basic assessment of the effects of cosmic variance. Furthermore, approximately doubling the exposure on previously detected sources substantially improves the constraints on their positions, spectral properties, and variability properties.  In this paper, we present up-to-date \\chandra\\ source catalogs and data products derived from the full $\\approx 2$~Ms \\hbox{CDF-S} data set along with details of the observations, data processing, and technical analysis. Detailed subsequent investigations and scientific interpretation of the new \\hbox{CDF-S} sources will be presented in future papers, e.g., studies of heavily obscured and Compton-thick active galactic nuclei (AGNs), high-redshift AGNs, AGN spectra and variability, starburst and normal galaxies, and clusters and groups of galaxies. In \\S2 we describe the observations and data reduction, and in \\S3 we present the main and supplementary point source catalogs and describe the methods used to create these catalogs. In \\S4 we estimate the background and sensitivity across the survey region. We also present basic number-count results for point sources in \\S5. We summarize in \\S6.  The Galactic column density along the line of sight to the \\hbox{CDF-S} is remarkably low: \\hbox{$N_{\\rm H}=8.8\\times 10^{19}$~cm$^{-2}$} \\citep[e.g.,][]{Stark1992}. The coordinates throughout this paper are J2000. A $H_0=70$~km~s$^{-1}$~Mpc$^{-1}$, $\\Omega_{\\rm M}=0.3$, and $\\Omega_{\\Lambda}=0.7$ cosmology is adopted. ",
        "conclusions": "We have presented catalogs and basic analyses of \\hbox{X-ray} point sources detected in the $\\approx2$~Ms \\hbox{CDF-S}, which is one of the two deepest \\chandra\\ surveys. The key points from this work are the following:  \\begin{enumerate} \\item The entire \\hbox{CDF-S} consists of 23 separate observations with 1.911~Ms of combined exposure. The survey covers an area of 435.6 arcmin$^{2}$.  \\item The main \\chandra\\ source catalog consists of 462 sources that were detected using {\\sc wavdetect} with a false-positive probability threshold of $1\\times10^{-6}$. These sources were detected in up to three X-ray bands: 0.5--8.0~keV, 0.5--2.0~keV, and 2--8~keV; 135 of these sources are new.  \\item The first supplementary \\chandra\\ source catalog contains 86 sources that were generated by merging the $\\approx250$~ks \\hbox{E-CDF-S} with the \\hbox{CDF-S}, which provides additional sensitivity in the outer regions of the \\hbox{CDF-S}.  \\item The second supplementary \\chandra\\ source catalog contains 30 sources that were detected at a lower X-ray significance threshold of $1\\times10^{-5}$ and that have bright optical counterparts ($R<23.8$).  \\item Source positions for the main and supplementary \\hbox{CDF-S} plus \\hbox{E-CDF-S} \\chandra\\ catalogs have been determined using centroid and matched-filter techniques; the median positional uncertainty is $\\approx0\\farcs36$.  \\item The basic X-ray and optical properties of the point sources indicate a variety of source types. More than half of the sources in the main \\chandra\\ catalogs appear to be AGNs. Of the 135 newly detected sources, $\\approx$55\\% appear to be AGNs while $\\approx$45\\% appear to be starburst and normal galaxies. The majority of the sources in the supplementary optically bright catalog are expected to be normal and starburst galaxies.  \\item The average backgrounds in the 0.5--2.0 and 2--8~keV bands are 0.066 and 0.167~counts Ms$^{-1}$~pixel$^{-1}$, respectively. Thus these observations are nearly photon limited near the aim point and could be extended to substantially greater depths with further exposure. The background count distributions are very close to Poisson distributions. The on-axis flux limits in the 0.5--2.0~keV and 2--8~keV bands are $\\approx1.9\\times10^{-17}$ \\flux\\ and $\\approx1.3\\times10^{-16}$ \\flux, respectively.  \\item Compared to the other deepest \\chandra\\ survey, the $\\approx2$~Ms \\hbox{CDF-N}, the \\hbox{CDF-S} has similar effective exposure coverage and sensitivity limits. The cumulative number counts of these two fields are consistent with each other to within $\\approx$1~$\\sigma$ at fluxes above $\\approx2\\times10^{-16}$~\\flux\\ in the soft band. The \\hbox{CDF-N} number counts are up to $\\approx25\\%$ higher than the \\hbox{CDF-S} number counts at the faintest fluxes in the soft and hard bands, indicating small field-to-field variations.  \\end{enumerate}"
    },
    "0806/0806.2063_arXiv.txt": {
        "abstract": "Although $^{79}$Se represents an important branching in the weak $s$ process, the stellar neutron capture cross sections to this isotope have not yet been measured experimentally. In this case, experimental data is essential for evaluating the important branching in the $s$-process reaction path at $^{79}$Se. The total cross section of $^{78}$Se at a stellar energy of $kT$ = 25 keV has been investigated with a combination of the activation technique and accelerator mass spectrometry (AMS), since off\\-line decay counting is prohibitive due to the long terrestrial half life of $^{79}$Se (2.80$\\pm$0.36 $\\times$10$^5$~y \\cite{HJJ02}) as well as the absence of suitable $\\gamma$-ray transitions. The preliminary result for the total Maxwellian averaged cross section is $<\\sigma>_{30~keV}$= 60.1$\\pm$9.6~mbarn, signif\\-icantly lower than the previous \\mbox{recommended} value. In a second measurement, also the partial cross section to the 3.92~min-isomer was determined via $\\gamma$-spectroscopy and yielded $<\\sigma>_{30~keV}$(part.)= 42.0$\\pm$2.0~mbarn. ",
        "introduction": "The origin of the elements heavier than iron can be almost completely ascribed to neutron capture processes characterized by much longer and much shorter time scales compared to average $\\beta$-decay half-lives. The respective nucleosynthesis processes, known as the slow ($s$) and the rapid ($r$) neutron capture process, contribute in approximately equal parts to the total elemental abundances in the mass range above iron. The $s$ process can be subdivided into two fractions, corresponding to different mass regions, temperature ranges and neutron exposures. The \"weak\" component, responsible for the elements with A$<$90, is driven by the neutron production via the $^{22}$Ne$(\\alpha,n)$$^{25}$Mg reaction at temperatures of T= 2-3 $\\times$10$^8$ K. The \"main\" component is mainly based on the $^{13}$C$(\\alpha,n)$$^{16}$O reaction, which operates at temperatures of 10$^8$ K. In this case the $^{22}$Ne$(\\alpha,n)$$^{25}$Mg contributes only a small part of the total exposure, but affects the f\\-inal abundance pattern. This main component is producing the $s$ abundances in the region A$\\leq$90$\\leq$A$\\leq$208. The astrophysical environment for the weak $s$ process are massive stars with core He and shell C burning, whereas the main component needs He shell f\\-lashes in low mass TP-AGB stars.  Our measurements refer to the region of the \"weak\" $s$ process. Among the nuclei involved, the long-lived radioactive isotopes $^{63}$Ni, $^{79}$Se, and $^{85}$Kr assume key positions, because the $\\beta$-decay rate becomes comparable to the neutron capture rate ($\\lambda_\\beta \\approx \\lambda_n$). This competition leads to branchings of the synthesis path, as sketched in Fig.~\\ref{branch}. The branching point isotopes are important, since they can be used for a diagnosis of the temperature and neutron density by comparing the ratios of isotopes at masses above and below the respective branching with observations.  The split of the reaction path at $^{79}$Se causes part of the f\\-low to bypass $^{80}$Kr, whereas at $^{82}$Kr both f\\-lows merge back. The strength of the resulting branching is ref\\-lected in the abundance of the $s$-only isotopes $^{80,82}$Kr, which are shielded from the decay chains of the $r$ process by $^{80,82}$Se. Analysis of the local abundance pattern in the mass region 80$\\leq$A$\\leq$82 yields the effective half-life of $^{79}$Se at the $s$-process site.  \\begin{figure*}[!hb] \\includegraphics{branching.eps}% \\caption{\\label{branch}Sketch of the $s$-process f\\-low in the Se-Br-Kr region. The half-lives given here are terrestrial half-lives.} \\end{figure*}  It was noticed that the effective branching ratios for $^{79}$Se and $^{85}$Kr are almost equal. This leads to a conf\\-lict, since the terrestrial half-life of both isotopes differ by roughly four orders of magnitude. The solution of this puzzle is due to $\\beta$ decays from thermally populated excited states. The ground state $\\beta$-decay of $^{79}$Se is f\\-irst-forbidden unique, but allowed decays are possible from the isomer at 95.7~keV, and from states at 128~keV and 137~keV, which are populated to 1.0\\%, 0.4\\% and 1.2\\%, respectively, at $kT$=30~keV \\cite{WaB82,KBW89}. Apart from this, stellar decay rates are also inf\\-luenced by the high degree of ionization in the hot stellar plasma. This allows \"bound-state\" $\\beta$-decays (Q$_\\beta$$<$300~keV), which contribute $\\approx$25\\% to the stellar rate \\cite{TaY83}. It was shown for $^{79}$Se (see Fig.~7 in \\cite{KlK88}), that the $\\beta$-decay rate is almost constant for T$<$10$^8$K, but drops quickly to the order of a few years when low-lying f\\-irst excited states become thermally populated. Since the temperature dependence of the half-life is well known from calculations (see \\cite{KlK88} and \\cite{TaY83}), the branching at $^{79}$Se can be interpreted as an $s$-process thermometer \\cite{WBK86}. Nevertheless, this leads to the uncommon fact that the temperature dependence is better known than the terrestrial half-life.  For $^{78}$Se, no experimental information on its stellar (n,$\\gamma$) cross section to the ground and isomeric state in $^{79}$Se existed so far. Both cross sections were determined by means of the activation technique. In combination with AMS this technique can be extended to hitherto inaccessible cases, e.g. to reactions producing very long-lived nuclei with very weak or completely missing $\\gamma$ transitions. The application of AMS counting in stellar neutron reactions, which has been demonstrated recently for $^{62}$Ni \\cite{Nas04} and also $^{58}$Ni \\cite{RDF06}, has the further advantage of being independent of uncertain $\\gamma$-ray intensities or half-lives. ",
        "conclusions": ""
    },
    "0806/0806.2125_arXiv.txt": {
        "abstract": "We propose two distinct atom interferometer gravitational wave detectors, one terrestrial and another satellite-based, utilizing the core technology of the Stanford $10\\,\\text{m}$ atom interferometer presently under construction.  Each configuration compares two widely separated atom interferometers run using common lasers.  The signal scales with the distance between the interferometers, which can be large since only the light travels over this distance, not the atoms. The terrestrial experiment with two $\\sim 10 \\text{ m}$ atom interferometers separated by a $\\sim 1 \\text{ km}$ baseline can operate with strain sensitivity $ \\sim \\frac{10^{-19}}{\\sqrt{\\text{Hz}}}$ in the 1 Hz - 10 Hz band, inaccessible to LIGO,  and can  detect gravitational waves from solar mass binaries out to megaparsec distances. The satellite experiment with two atom interferometers separated by a  $\\sim 1000 \\text{ km}$ baseline can probe the same frequency spectrum as LISA with comparable strain sensitivity $ \\sim \\frac{10^{-20}}{\\sqrt{\\text{Hz}}}$.  The use of ballistic atoms  (instead of mirrors) as inertial test masses improves systematics coming from vibrations, acceleration noise, and significantly reduces spacecraft control requirements.  We analyze the backgrounds in this configuration and discuss methods for controlling them to the required levels. ",
        "introduction": "Gravitational waves offer a rich, unexplored source of information about the universe \\cite{Schutz:1999xj, Cutler:2002me}.  Many phenomena can only be explored with gravitational, not electromagnetic, radiation.  These include accepted sources such as white dwarf, neutron star, or black hole binaries whose observation could provide useful data on astrophysics and general relativity.  It has even been proposed that these compact binaries could be used as standard sirens to determine astronomical distances and possibly the expansion rate of the universe more precisely \\cite{Schutz:1986gp}.  Gravitational waves could also be one of the only ways to learn about the early universe before the surface of last scattering.  There are many speculative cosmological sources including inflation and reheating, early universe phase transitions, and cosmic strings.  For all these applications, it is important to be able to observe gravitational waves as broadly and over as large a range of frequencies and amplitudes as possible.   In this article we expand on a previous article \\cite{Dimopoulos:2007cj}, giving the details of our proposal for an Atomic Gravitational wave Interferometric Sensor (AGIS).  We develop proposals for two experiments, one terrestrial, the other satellite-based.  We will see that, at least in the configurations proposed here, it is primarily useful for observing gravitational waves with frequencies between about $10^{-3} \\text{~Hz}$ and $10 \\text{~Hz}$.  In particular, the terrestrial experiment is sensitive to gravitational waves with frequencies $\\sim 1 - 10 \\text{~Hz}$, below the range of any other terrestrial gravitational wave detector.  This arises in part from the vast reduction in systematics available with atom interferometers but impossible with laser interferometers.  The satellite-based experiment will have peak sensitivity to gravitational waves in the $\\sim 10^{-3} - 1 \\text{~Hz}$ band.  The use of atomic interferometry also leads to a natural reduction in many systematic backgrounds, allowing such an experiment to reach sensitivities comparable to and perhaps better than LISA's with reduced engineering requirements.  The ability to detect gravitational waves in such a low frequency band greatly affects the potential sources.  Binary stars live much longer and are more numerous at these frequencies than in the higher band around 100 Hz where they are about to merge.  Stochastic gravitational waves from cosmological sources can also be easier to detect at these low frequencies.  These sources are usually best described in terms of the fractional energy density, $\\Omega_\\text{GW}$, that they produce in gravitational waves.  The energy density of a gravitational wave scales quadratically with frequency (as $\\rho_\\text{GW} \\sim h^2 f^2 M^2_\\text{pl}$).  Thus, a type of source that produces a given energy density is easier to detect at lower frequencies because the amplitude of the gravitational wave is higher.  This makes low frequency experiments particularly useful for observing cosmological sources of gravitational waves.  There are several exciting proposed and existing experiments to search for gravitational waves including broad-band laser interferometers such as LIGO and LISA, resonant bar detectors \\cite{Astone:2006ut, Zendri:2002em}, and microwave cavity detectors \\cite{Ballantini:2005am}.  Searching for gravitational waves with atomic interferometry is motivated by the rapid advance of this technology in recent years.  Atom interferometers have been used for many high precision applications including atomic clocks \\cite{Oskay}, metrology, gyroscopes \\cite{PhysRevLett.78.2046}, gradiometers \\cite{PhysRevLett.81.971}, and gravimeters \\cite{0026-1394-38-1-4}.  We consider the use of such previously demonstrated technology to achieve the sensitivity needed to observe gravitational waves.  Further, we consider technological advances in atom interferometry that are currently being explored, including in the Stanford 10 m interferometer, and the possible impact these will have on the search for gravitational waves.   The different sections of this paper are as independent as possible.  In section \\ref{Sec:AI} we give a brief description of atom interferometry.  In Section \\ref{Sec: GW Signal} we calculate the gravitational wave signal in an atom interferometer.  In section \\ref{Sec: Earth} we discuss an experimental configuration for observing this signal on the earth and the relevant backgrounds.  In section \\ref{Sec: Space} we give the setup and backgrounds for a satellite-based experiment.  In section \\ref{Sec: Sources} we give a brief summary of the astrophysical and cosmological sources of gravitational waves that are relevant for such experiments.  In section \\ref{Sec: Sensitivities} we give a description of the projected sensitivities of the earth- and space-based experiments.  In Section \\ref{Sec: Conclusion} we compare this work with previous ideas on atom interferometry and gravitational waves and summarize our findings. ",
        "conclusions": "\\label{Sec: Conclusion}  \\subsection{Comparison with Previous Work}  Previous studies on the role of atom interferometers in gravitational wave detection concluded that they would be of limited use in probing the gravitational wave spectrum. Our proposal differs significantly from these efforts owing to the central role played by light pulse interferometry in our setup.  The work of \\cite{Chiao+Speliotopoulos}, \\cite{Foffa:2004up} and \\cite{Roura:2004se} used material mirrors like diffraction gratings to execute the interferometer. The gravitational wave signal in the configurations considered in these papers is  $\\sim k h d$ where $k$ is the momentum of the atom, $h$ the amplitude of the gravitational wave and $d$ the distance between the mirrors. It is experimentally difficult to make the distance between these mirrors bigger than $\\sim 1$ m. Even if the distance between the mirrors were to be increased, the experiment would still be difficult since the separation between the two arms of the atom's wave function must also be equally scaled. These considerations forced the authors to conclude that  an unrealistic atom flux would be needed to see a gravitational wave. The use of material mirrors suffers from the additional drawback that the mirrors would be subject to vibration noise. The mirrors would have to be placed on vibration isolation stacks so this interferometer would be subject to the same limitations as light based interferometers like LIGO.  The work of \\cite{Delva:2006qa} and \\cite{Tino:2007hs} described atom interferometers which used light pulse interferometry. However, these authors did not consider the effect of the gravitational wave on the light pulses used to execute the interferometry. Without this effect, the phase shift in the interferometer is $\\sim k h d$ where $d$ is the separation between the two arms of the interferometer (see discussion in Section \\ref{Sec: Signal Results}). Since the separation between the two arms of the interferometer cannot be easily scaled, these authors were also forced to consider unrealistic atom fluxes. Moreover, these papers did not discuss strategies to handle crucial backgrounds to gravitational wave detection like vibration and laser phase noise.  In this paper, we point out that the effect of the gravitational wave on the light pulses used to execute the interferometer is crucial and can be easily scaled to increase the signal. When the interferometer is operated by a laser at a distance $L$, a gravitational wave of amplitude $h$ causes a phase shift $\\sim k h L$. This signal increases as long as $L$ is smaller than the wavelength of the gravitational wave.  Unlike the separation between the two arms of the atom's wave function, the distance between the atom and the laser can be easily scaled.   With $L \\sim 10$ km, the signal in this interferometer is $10^4$ larger than the signal in the configurations previously considered. In addition to boosting the signal, the configuration considered in this paper offers an effective way to deal with vibration and laser phase noise. By using the same laser to run two widely separated interferometers and measuring the differential phase shift between the two interferometers, this setup drastically suppresses the effects of vibrations and laser phase noise. Our setup thus achieves a large, scaleable enhancement in the signal while simultaneously suppressing backgrounds thereby making it possible to search for gravitational waves with current technology.  The SAGAS \\cite{SAGAS} project that uses atom  interferometry and ion clock techniques to explore gravity in the outer solar system was proposed. SAGAS will improve current bounds on stochastic gravitational waves in the frequency band $10^{-5} \\text{ Hz} - 10^{-3} \\text{ Hz}$ but is not expected to be sensitive to known sources of gravitational radiation. In contrast, our proposal  will search for gravitational waves in the $10^{-3} \\text{ Hz} - 10 \\text{ Hz}$ band at sensitivities that can detect gravitational waves from expected sources.  \\subsection{Summary}  We have proposed two experiments, terrestrial and satellite-based, to observe gravitational waves using atom interferometry.  Both experiments rely on similar underlying ideas to achieve a large, scaleable enhancement to the gravitational wave signal while naturally suppressing many backgrounds.  A differential measurement is performed between two atom interferometers run simultaneously using the same laser pulses.  The lasers provide a common `ruler' for comparison of the two interferometers.  The distance between the interferometers can be large because only the light travels over this distance, not the atoms.  The signal still scales with this distance and so can be competitive with light interferometers.  In a sense, the atom interferometers are the analogue of the mirrors in a light interferometer and it is the distance between them that determines the size of the signal.  Further, many backgrounds are naturally suppressed by this method.  Laser phase noise, which must be cancelled between the two arms of a light interferometer, is here cancelled by the differential measurement between the two atom interferometers.  Since this subtraction is between two interferometers along the same laser axis with only vacuum in between, vibrations of the lasers (and any optics) are cancelled as well.  The atoms themselves, the analogues of the mirrors in a light interferometer, are in free fall and are unaffected by vibrations.  This removes one of the major backgrounds which prohibits terrestrial light interferometers from achieving sensitivity to lower frequencies.  For example, Advanced LIGO will lose sensitivity below $\\sim 10$ Hz due to direct (non-gravitational) coupling to vibrations (see e.g. \\cite{Hughes:1998pe}).  Similarly, in the satellite-based experiment the atoms can be far from the satellite, greatly reducing the engineering requirements on the control of the satellites.  Satellite position control and laser noise are two of the major hurdles for an experiment such as LISA.  For similar gravitational wave sensitivity, these requirements are significantly reduced for our atom interferometer proposal.  Of course, new backgrounds may enter in an atomic experiment.  We have attempted to consider all the relevant backgrounds and show that they are controllable with practical technology in a realistic setup.  Many backgrounds will require careful engineering, just as for any gravitational wave detector.  We are certainly not experts in every relevant area of expertise necessary for such experiments, but this experiment seems possible and exciting enough to merit more serious consideration.  An interesting consequence of having a differential measurement between two interferometers along the same baseline is that this setup would have sensitivity to scalar-type perturbations, that would, for example, change the length of the perimeter of the LISA triangle \\footnote{Thanks to Bob Wagoner for pointing this out.}.  LIGO lacks sensitivity to these signals since two perpendicular laser arms are used to remove backgrounds including laser phase noise.  Our setup confers sensitivity to overall changes in the length of a single arm since each arm is a laser phase noise free combination.  In LISA these Sagnac channel events would be vetoed.  One interesting signal of this type would arise from large mass dark matter particles passing near the detector \\cite{Adams:2004pk, Seto:2004zu}.  There are many avenues for improvement of these proposals in the future.  The atom statistics may be improved with improved cooling techniques, ultimately limited only by the limit on the density of the cloud from atom-atom interactions and on the total number of atoms from opacity of the cloud.  The use of squeezed atom states may also allow significant improvements in atom statistics beyond the standard quantum limit.  Improved sensitivities could also come from better classical and atom optics including multi-photon LMT beamsplitters, higher laser powers and larger laser waist sizes.  There may also be clever ideas for improved atomic systems, for example which suppress the spontaneous 2-photon transition rate without suppressing the stimulated rate.  It is difficult to predict what advances will be made in the future.  Nevertheless, the rapid advance of atom interferometry motivates us to consider a range of sensitivity curves that illustrate the possibilities not just for current but also near future technology.  The proposed gravitational wave detectors may allow the observation of low frequency sources in the band $10^{-3} - 10 \\text{~Hz}$.  This is a very exciting range for astrophysical and cosmological sources.  Compact binaries including black holes, neutron stars, and white dwarfs live for a long period in this band.  Such low frequencies also allow enhanced sensitivity to a stochastic background of gravitational waves, assuming at least two such AGIS detectors are built.  Many cosmological sources arising from physics beyond the Standard Model could be present in this range including inflation and reheating, early universe phase transitions, or cosmic strings.  The observation of gravitational waves has the potential to reveal significant information about new physics at both the shortest and longest length scales."
    },
    "0806/0806.2639_arXiv.txt": {
        "abstract": "We present new {\\em Spitzer Space Telescope} observations of stars in the young ($\\sim$5 Myr) $\\gamma$ Velorum stellar cluster.  Combining optical and 2MASS photometry, we have selected 579 stars  as candidate members of the cluster.  With the addition of the Spitzer mid-infrared data, we have identified 5 debris disks around A-type stars, and 5-6 debris disks around solar-type stars, indicating that the strong radiation field in the cluster does not completely suppress the production  of planetesimals in the disks of cluster members. However, we find some evidence that the frequency of circumstellar primordial disks is lower, and the IR flux excesses are smaller than for disks around stellar populations with similar ages.   This could be evidence for a relatively fast dissipation of circumstellar dust by the strong radiation field from the highest mass star(s) in the cluster. Another possibility is that  $\\gamma$ Velorum stellar cluster is slightly older than reported ages and the the low frequency of primordial disks reflects the fast disk dissipation observed at $\\sim$5 Myr. ",
        "introduction": "\\label{sec:int}  In recent years the sensitivity of the {\\em Spitzer Space Telescope} has made it possible to expand our knowledge of dusty disk emission dramatically, enabling statistically significant studies of disk emission as a function of mass and age over a much wider wavelength range than previously possible \\citep[e.g.,][]{lada06,aurora06,hernandez07a,hernandez07b}. In addition to confirming the rapid decrease of optically-thick disk emission previously inferred from near-infrared observations \\citep{strom89, haisch01}, {\\em Spitzer} observations have provided strong evidence that disk dissipation timescales depend significantly on host star mass \\citep{lada06,carpenter06,hernandez07a}, with a timescale for inner disk dissipation of 5-7 Myr \\citep{haisch01,hartmann05b,hernandez07a} for low mass stars ($\\lesssim$1M$\\sun$), and $<$3 Myr \\citep{hernandez05} for intermediate mass stars (2-10{$\\msun$}). {\\em Spitzer} observations have also provided new clues to disk structure through Spectral Energy Distribution (SED) analysis, from slight indications of faster dust settling in inner disks \\citep{aurora06} to strong evacuation of the inner disks in some stars that retain optically thick outer disks \\citep{calvet02,calvet05,dalessio05a,espaillat07a}. Spitzer data have also been used to show that debris disks around intermediate mass stars are more frequent and have larger 24 $\\mu$m excesses around 10-13 Myr, when models predict that dust is produced by collisional cascades due to the formation of large solid bodies ($\\sim$1000 km) in the disk \\citep{kenyon05,hernandez06,currie08}.  The possible effects of environment on disk frequency and structure are less clear.  In clusters with high mass stars, disk photoevaporation due to the strong ultraviolet radiation field can be important, and is directly observed in the Orion Nebula Cluster \\citep{richling00,adams04,hollenbach04,throop05,clarke07,megeath07}. \\citet{balog07} found a reduction in disk frequency of about a factor of two in the relatively young cluster NGC 2244, but only in the innermost regions (d$\\lesssim$0.5 pc).  NGC 2244 is very rich and contains seven O type stars, so that the ultraviolet radiation field should be very strong.  The discovery of a young cluster of low-mass stars around the $\\gamma$ Vel system \\citep{pozzo00}, with the central binary consisting of an O7.5 star and a Wolf Rayet star \\citep{demarco99,demarco00}, provides an opportunity to study the possible effects of central star ultraviolet radiation fields on disk evolution. The Wolf Rayet star is classified as WC8 (carbon-rich), implying a progenitor mass of $>$40$\\msun$ \\citep{demarco99,crowther07}. In this paper we report {\\em Spitzer} observations of infrared emission in this cluster. In \\S\\ref{sec:obs} we present the observational data used in this contribution, and discuss membership in \\S\\ref{s:mem}. We analyze the infrared data \\S\\ref{sec:disks}, and compare with disk frequencies in other young stellar populations in \\S\\ref{sec:comp}, with our main results summarized in \\S\\ref{sec:conc}. ",
        "conclusions": "\\label{sec:conc} We have used the IRAC and MIPS instruments on board the Spitzer Space Telescope to conduct a study of disks around the $\\gamma$ Velorum cluster. Since the central object is a binary system consisting of the closest known Wolf Rayet star and a high mass O star, a strong UV radiation field and stellar winds are present in the cluster. Using optical photometry of  X ray sources (2XMM) and members confirmed by spectroscopy (Jeffries R.D. et al, in preparation), 579 photometric candidates were selected as possible members of the cluster. The level of contamination by non-members depends on the V-J color range, showing the highest level of contamination ($\\sim$68\\%) at V-J = 1.5 - 3.5, where the field giant branch crosses the young stellar population. Combining optical, 2MASS and Spitzer data we have detected infrared excess in 29 stars. One of the infrared excess stars is a Be star. We report 5 debris disks around A type stars, 5 debris disks around solar type stars (spectral type range F - early-K) and one solar type stars with infrared excess produced by a very massive debris disks or by a primordial disk with a high degree of dust settling. Seventeen disk bearing low mass stars (K5 or later) were found in the cluster with a range of disk properties.  We classified these objects in three classes using the infrared excess at 5.8{\\micron} and the SED slope [8.0]-[24]:  9 class II stars, 7 evolved disks star and one pretransitional candidate. We found that 76\\% of the stars bearing primordial disks have color V-J$>$4 ($\\sim$later than M3.5), indicating a mass dependent timescale for disk dissipation in the $\\gamma$ Velorum cluster, similar to results in other young star populations \\citep{carpenter06,hernandez07a,hernandez07b,lada06}.  We found that the disk frequency observed in the $\\gamma$ Velorum cluster is lower than in other young stellar groups with similar ages ($\\sim$5 Myr), and comparable to values found in older populations ($\\sim$10Myr). Empirical isochrones support the age of $\\sim$5 Myr for the overall population of the $\\gamma$ Velorum cluster. The relative low frequency observed in the cluster is consistent with: (1) rapid disk dissipation due to the influence of the strong radiation fields and winds from the WR and O star cluster members or (2) the age of the cluster is slightly older and the low frequency reflects the fast dropping in disk frequency observed at $\\sim$5 Myr.  More than 75\\% of the primordial disk stars have inner disk emission lower than the median observed in the disk population of other 5 Myr clusters, indicating that  the disks of the $\\gamma$ Velorum cluster have a larger degree of dust settling.  We find no primordial disks closer than  $\\sim$0.75 pc from the high mass central binary.   At projected distance $>$1 pc, the disk fraction seems to be constant around 5-7\\%. This result is similar to that found in the NGC 2244 cluster \\citep{balog07} in which primordial disks are affected only in the immediate vicinity of the O stars ($<$0.5pc). Our results reinforce an emerging trend in studies of disk frequency, specifically that optically thick disk dissipation is rapid for $\\sim 90$\\% of systems, lasting less than 5 Myr, while $\\sim 10$\\% of disks can last for 8 Myr or more."
    },
    "0806/0806.3854_arXiv.txt": {
        "abstract": "We present synthetic OH Zeeman splitting measurements of a super-Alfv\\'{e}nic molecular cloud model. We select dense cores from synthetic $^{13}$CO maps computed from the largest simulation to date of supersonic and super-Alfv\\'{e}nic turbulence. The synthetic Zeeman splitting measurements in the cores yield a relation between the magnetic field strength, $B$, and the column density, $N$, in good agreement with the observations. The large scatter in $B$ at a fixed value of $N$ is partly due to intrinsic variations in the magnetic field strength from core to core. We also compute the relative mass-to-flux ratio between the center of the cores and their envelopes, ${\\cal R}_{\\mu}$, and show that super-Alfv\\'{e}nic turbulence produces a significant scatter also in ${\\cal R}_{\\mu}$, including negative values (field reversal between core center and envelope). We find ${\\cal R}_{\\mu} < 1$ for 70\\% of the cores, and ${\\cal R}_{\\mu} < 0$ for 12\\%. Of the cores with $|B_{\\rm LOS}|>10$~$\\mu$G, 81\\% have ${\\cal R}_{\\mu}<1$. These predictions of the super-Alfv\\'{e}nic model are in stark contrast to the ambipolar drift model of core formation, where only ${\\cal R}_{\\mu}>1$ is allowed. ",
        "introduction": "The process of star formation in molecular clouds involves a complex interaction of turbulent velocity fields, gravitational forces, and magnetic fields. While the role of turbulence has been investigated in depth only in recent years, the importance of the magnetic field in star formation was recognized long ago. On small scales, the magnetic field may support sub-critical prestellar cores against gravitational collapse for an ambipolar drift time \\citep{Lizano+Shu89}, and provide a way to shed angular momentum by magnetic breaking \\citep{Mouschovias77,Mouschovias79} and by helping to power winds and jets \\citep{Pudritz+Norman83}. On larger scales, a strong enough magnetic field may even prevent the collapse of a giant molecular cloud (GMC), potentially explaining the low star-formation rate in the Galaxy \\citep{Shu+87}.  As an alternative to this scenario of large scale magnetic support, \\citet{Padoan+Nordlund99MHD} proposed a super-Alfv\\'{e}nic model of molecular clouds. In this model, the turbulent flows can compress the gas in all directions through shocks. Compressions with a component perpendicular to the magnetic field can locally increase the field strength, and the net result is a correlation of magnetic field and gas density, though with a large scatter. While in this model the mean magnetic field of a GMC is very low, its value inside dense cores and filaments (where it is usually observed) can be very large. Based on a number of tests comparing simulations with observations \\citep{Padoan+Nordlund99MHD,Padoan+04power}, the super-Alfv\\'{e}nic model seems to explain the observations better than models assuming a strong magnetic field.  The most direct way to infer the magnetic field strength in GMCs is by Zeeman splitting measurements. These measurements are very difficult, and are available in GMCs only for relatively dense cores, primarily from OH emission lines \\citep[e.g.][]{Crutcher+93,Troland+Crutcher08}. Only a handful of these measurements have provided detections of the field strength. Zeeman splitting is only sensitive to the line-of-sight component of the magnetic field averaged along the line of sight, and weighted by the OH emission (roughly speaking gas density and OH abundance). Thus, Zeeman splitting does not directly provide the large-scale volume-averaged magnetic field strength, needed to understand the nature of the fragmentation process that leads to star formation.  In this Letter, we relate the large-scale mean magnetic field to local Zeeman splitting measurements by a {\\it deductive method}, predicting the observational properties of the theoretical model, rather than inducing the theory from the observations. The super-Alfv\\'{e}nic molecular cloud model is obtained from a numerical experiment of three-dimensional, supersonic, magneto-hydrodynamic (MHD) turbulence. Synthetic Zeeman splitting measurements are then simulated through radiative transfer calculations. This work is partly motivated by i) an improved interpretation of existing Zeeman splitting measurements by Crutcher (2008, pers. comm.), and ii) a new observational test of ambipolar drift by \\citet{Crutcher+08b}.  \\begin{figure*}[t] \\epsscale{1.15} \\plottwo{f1a.eps}{f1b.eps} \\caption{\\emph{Left:} Simulated $^{13}$CO (1-0) map of the model in the $z$-axis direction. The locations of the cloud cores are shown with squares. The circles indicate the locations of telescope beams used in the synthetic observations of three cores. \\emph{Right:} Line-of-sight magnetic field strength as calculated from Zeeman splitting.} \\end{figure*}   A conclusion of the new analysis of Zeeman splitting measurements by Crutcher (2008, pers. comm.) is that the observed scatter in the relation between the absolute value of the line-of-sight magnetic field strength, $| B_{\\rm LOS}|$, and the column density, $N$, is partly due to a large scatter in the intrinsic magnetic field strength, $B$, at any given $N$. As shown before \\citep{Padoan+Nordlund99MHD}, and as illustrated in the present work, this is a fundamental prediction of the super-Alfv\\'{e}nic model. If the mean field were very strong, its local value could not have a large scatter, and the observed scatter in $| B_{\\rm LOS}|$ could only arise as a result of random orientations, which seems to be ruled out by the new analysis of Crutcher (2008, pers. comm.).  The new observational test of ambipolar drift by \\citet{Crutcher+08b}, is based on the mass to flux ratio of the envelope relative to that of the core -- henceforth referred to as the \"relative mass-to-flux ratio\", ${\\cal R}_{\\mu}$, as the ambipolar drift model of core formation predicts ${\\cal R}_{\\mu}>1$ (e.g. \\citet{Ciolek+Mouschovias94} -Fig.~3a). In this Letter we show that the super-Alfv\\'{e}nic model predicts a significant scatter in ${\\cal R}_{\\mu}$ (including negative values), with ${\\cal R}_{\\mu}<1$ in 81\\% of the cores with $| B_{\\rm LOS}|>10$~$\\mu$G. Values of ${\\cal R}_{\\mu}<1$ in cores formed by turbulent flows were first found by \\cite{Vazquez-Semadeni+05}, where a plausible mechanism for explaining that result was proposed. ",
        "conclusions": "We have found that the super-Alfv\\'{e}nic model reproduces well the observed $|B_{\\rm LOS}|-N$ relation in dense cores, despite its very low mean magnetic field strength. Because the large scatter of $| B_{\\rm LOS}|$ in the simulation is partly due to intrinsic variations of $B$ from core to core, and not only to the effect of the random orientation, we interpret the observations to also be consistent with intrinsic variations in $B$ from core to core. This interpretation of the observations, confirmed by the new analysis of Crutcher (2008, pers. comm.), suggests that the mean magnetic field on a larger scale cannot be as large as in the ambipolar drift model.  The super-Alfv\\'{e}nic model also predicts that OH Zeeman splitting measurements of cores with central field values of $| B_{\\rm LOS}|>10$~$\\mu$G should yield ${\\cal R}_{\\mu}<1$ in more than $4/5$ of the cores, and one or more field reversals in the envelope beams, relative to the central one, in 28\\% of the cores. In contrast, the ambipolar drift model of core formation predicts ${\\cal R}_{\\mu} >1$ and no field reversal. If future observational measurements of ${\\cal R}_{\\mu}$ \\citep{Crutcher+08b} were to yield ${\\cal R}_{\\mu}<1$ in most cores, as in the super-Alfv\\'{e}nic model, the ambipolar drift model of core formation would seem to be ruled out."
    },
    "0806/0806.4968_arXiv.txt": {
        "abstract": "{About 2500 planetary nebulae are known in our Galaxy but only 224 have central stars with reported spectral types in the Strasbourg-ESO Catalogue of Galactic Planetary Nebulae (Acker et al. 1992; Acker et al. 1996).} {We have started an observational program aiming to increase the number of PN central stars with spectral classification.} {By means of spectroscopy and high resolution imaging, we identify the position and true nature of the central star. We carried out low resolution spectroscopic observations at CASLEO telescope, complemented with medium resolution spectroscopy performed at Gemini South and Magellan telescopes.} {As a first outcome of this survey, we present for the first time the spectra of the central star of the PN  \\object{Th 2-A}. These spectra show emission lines of ionized C and O, typical in Wolf-Rayet stars.} {We identify the position of that central star, which is not the brightest one of the visual central pair. We classify it as of type [WO 3]pec, which is consistent with the high excitation and dynamical age of the nebula.} {} {} ",
        "introduction": "The central stars of planetary nebulae (CSPN) play an important role in the understanding of post-AGB stellar evolution. Planetary Nebulae (PNs) are the relics of intermediate-mass stars, and they are important probes of stellar evolution, stellar Populations, and cosmic recycling. Understanding the evolutionary status of particular classes of CSs is relevant in a broader astrophysical sense.   About 60\\% of the PNe listed in the Strasbourg - ESO Catalogue of Galactic Planetary Nebulae (Acker et al. 1992; Acker et al. 1996) include data concerning the CSPN, but only 30\\% of these actually have spectral classification. Among the 224 spectral types reported in that catalog, there are 75 stars presenting emission lines typical of Wolf-Rayet (WR) stars (Acker \\&  Neiner 2003, hereafter AN03).  The [WO] type (Oxygen sequence) was introduced by Smith \\& Aller (1969) to classify a group of 7 CSPN showing an optical WR-like spectrum with strong \\ion{O}{vi} emission lines. Subsequently, Barlow \\& Hummer (1982) defined WO stars as a new WR subtype. The most prominent feature of these objects is the strong emission of the \\ion{O}{vi} $\\lambda\\lambda$3811-34 \\AA\\, doublet, other strong emission lines in the spectra of [WO] stars are the \\ion{C}{iv} + \\ion{He}{ii} $\\lambda\\lambda4658-86$ \\AA \\ blend and \\ion{C}{iv} $\\lambda\\lambda5801-12$ \\AA \\ doublet, features which are also exhibited strongly by early WC stars. At the moment there are only four WO stars known in our Galaxy (van der Hucht, 2001; Drew et al. 2004) and about twenty [WO] (AN03; Gesicki et al. 2006).   We have started an observational program at CASLEO with the goal of studying the spectra of hitherto unknown CSPNe. In this paper we present one of our first results, i.e. the discovery of a new [WO] star. ",
        "conclusions": "\\label{resultados}  The central star of Th\\,2-A (object A in Fig.~\\ref{foto}) was described by Kerber et al. (2003) as being the ``well defined photometric center'' of the nebula (PHOT in their identification class). Ciardullo et al.\\ (1999) determined its Johnson $V$-band magnitude $V$ = $17.08$ using HST images. They included this PN in their survey for resolved companions, because they noted a nearby companion (star B), which resulted to have a $\\sim 50\\%$ probability of being a chance superposition. On the other hand, we did not find any spectroscopic study reported in the bibliography, related with the CSPN of Th\\,2-A.  \\begin{table} \\caption[]{Emission lines identified in the LCO spectrum of the central star of the PN Th 2-A over the 3520-9430\\AA\\ wavelength range. First column shows the central wavelength measured in our spectrum, when a deblend was not possible the centroid of the blend is indicated. The marks \":\" and \"n.m.\" refers to uncertain values and values no measured respectively.} \\label{star} \\centering \\setlength{\\tabcolsep}{0.5mm} \\begin{tabular}{l c c c } \\hline\\hline\\noalign{\\smallskip} Line [\\AA]   & Ion & FWHM [\\AA] & $W_\\lambda $ [\\AA] \\\\ \\noalign{\\smallskip}\\hline\\noalign{\\smallskip} 3822  & \\ion{O}{vi}               &   80 &  580  \\\\ 4118  & \\ion{O}{v}                &   31 &  67   \\\\ 4222  & \\ion{C}{iv} ?             &  58  & 22   \\\\ 4342  & \\ion{He}{ii}, \\ion{C}{iv} &  40  & 20 \\\\ 4515  & \\ion{O}{vi},\\ion{O}{v}    &  101 & 44   \\\\ 4674  & \\ion{C}{iv},\\ion{He}{ii}  &  66  &  385    \\\\ 4856  & \\ion{He}{ii}, \\ion{C}{iv} &  50  & 9 \\\\ 5101  & \\ion{O}{v} ?              &  9   & 71 \\\\ 5287  & \\ion{O}{vi}               &  54  & 33   \\\\ 5413  & \\ion{C}{iv}, \\ion{He}{ii} &  88  & 59 \\\\ 5468  & \\ion{C}{iv}               &  66  & 19  \\\\ 5590  & \\ion{O}{v}                &  92  &  82   \\\\ 5803  & \\ion{C}{iv}               &  86  &  190  \\\\ 6067  & \\ion{O}{viii}, \\ion{O}{vii} & 70 & 14  \\\\ 6197  & \\ion{O}{vi}               & 71   & 12 \\\\ 6383: & \\ion{O}{iii} ?            & n.m.     & n.m. \\\\ 6436: & \\ion{He}{ii}, \\ion{O}{v}  & n.m.     & n.m.  \\\\ 6555  & \\ion{He}{ii}, \\ion{C}{iv} &   79 & 63 \\\\ 7060  & \\ion{C}{iv}               &  88  & 44 \\\\ 7567: & ?                         & n.m. & n.m.   \\\\ 7725  & \\ion{C}{iv}, \\ion{O}{vi}  & 95   & 158 \\\\ 8844: & ?                         & n.m. & n.m. \\\\ 9111: & \\ion{O}{iii} ?            & n.m. & n.m. \\\\ \\hline \\end{tabular} \\end{table}   \\begin{figure*} \\centering \\includegraphics[width=14.0cm]{fig-2-b.eps} \\includegraphics[width=13.5cm]{fig-2-r.eps} \\caption[]{Blue and red spectra of Th 2-A obtained in December 2007 using LCO. The most intense nebular lines have been truncated (``X''-labeled) in order to show more details of the spectra. Note the broad lines from the central star at 3822 \\ion{O}{vi}, 4670 \\ion{C}{iv} and 5805 \\ion{C}{iv}.} \\label{esp} \\end{figure*}   The inspection of our (A+B) CASLEO spectrum revealed that the CSPN Th2-A was one of the rare [WO] stars, thus claiming our attention. Searching for published information about this star, we found the above cited Ciardullo et al.'s  work and retrieved higher spatial resolution HST-WFPC2 images from the MAST archive. We noted that certainly the B star seems more centered in the nebula than A. Moreover, we performed a preliminary aperture-photometry, % and estimated that the B star is bluer than A. But, in order to get confidence about the true identification of the central star of the PN, we obtained a high spatial resolution spectrum along the A and B stars (see fig~\\ref{foto}). Then, the GMOS-S spectrum (see observation description) confirmed that the B star is in fact the CSPN and A star is a late type one. These high spatial resolution images and spectra will be presented and discussed elsewhere in a thorough analysis of the whole planetary nebula system.    In our optical spectra of star B, we identified broad and intense \\ion{C}{iv} and \\ion{O}{vi} emission lines, that are typical of WR stars, thus confirming that this spectrum was originated in the nucleus of the PN. We identified other faint emission lines of \\ion{He}{ii}, \\ion{O}{v}, and \\ion{O}{vii} (see Table~\\ref{star} for a list of all the identified features), but some of the faintest ones remained unidentified. We specially noted three emission lines which we identify as \\ion{O}{v} $\\lambda5590$, \\ion{O}{vi} $\\lambda5290$, and \\ion{C}{iv} $\\lambda5440$, which are usually observed in WO-type objects. In Fig. \\ref{esp}, we show the LCO spectrum of Th 2-A and identify the emission lines detected in it. We show this spectrum because it covers the largest spectral range.   We performed a quantitative classification of the spectrum of the CSPN of Th 2-A using the criteria described by AN03, based on dereddened line intensities ratios and FWHM of emission lines, we obtained that most of the emission lines strongly indicate a [WO3]-type classification, but the FWHM of the \\ion{C}{iv} $\\lambda$ 5806 is similar to the one measured in [WO4]pec-type stars. Then, we suggest a [WO3]pec spectral classification for the CSPN of Th 2-A. A note of caution has to  be added, the line ratios indicate a type between [WO2] and [WO4] and the intrinsic scatter in the classification scheme is also non-negligible as pointed out by AN03.   Th 2-A is an old PN with a dynamical age of 7~kyr estimated in a very simplified way, i.e. ignoring the effects of the velocity gradient and the acceleration over time, and assuming a distance of $2.07$ Kpc (Phillips 2004), an angular size of $27\\farcs3$ (Tylenda et al. \\cite{ty}) and an expansion velocity \\textit{V}$_{\\mathrm{exp}}=18 $\\kms\\ (Meatheringham et al. \\cite{me}). The estimations of kinematical ages in PN are still matter of debate (Schonberner et al. 2005), but the classification of the CSPN fits well in the evolution sequence i.e. post-AGB $\\rightarrow$ [WC11] to [WC4] $\\rightarrow$ [WO4] to [WO1]. Moreover, the hot spectral-type estimated for the CSPN Th 2-A is consistent with the high excitation state of the nebula (Kingsburgh \\& Barlow, 1994).         Finally, we want to remark that it is important to perform multi-epoch observations that will help to determine its binary nature (if any). Certainly, there is a surprising lack of surveys that can detect binaries, letting the actual PN binary fraction as unknown, thus some theories that propose that binary interactions have an important role can not be tested (e.g. \\cite{moe06}). Moreover, a full quantitative analysis using stellar atmospheric modeling will be worth while."
    },
    "0806/0806.4159_arXiv.txt": {
        "abstract": " ",
        "introduction": "During six Earth gravity assist flybys, significant anomalous velocity increases of a few mm/s were observed (Antreasian and Guinn,~1998 and Anderson~$et~al.$,~2008) using both doppler frequency data and ranging methods. These are known as the flyby anomalies, and, so far, no explanations have been found that can account for them. Explanations considered and rejected to date have included: anomalous thruster activity, computer software glitches, troposphere and ionosphere effects and others (see Antreasian and Guinn,~1998 and Lammerz\\\"ahl~$et~al.$~(2006).  Anderson~$et~al.$~(2008) analysed the available data, some of which are summarised in table 1 (columns 1-4), and intriguingly managed to show that the six velocity anomalies observed so far (dv) fitted a formula, given, with slightly modified notation, by \\begin{equation} dv = 3.099 \\times 10^{-6} \\times v_{\\infty} \\times (cos {\\phi}_1 - cos {\\phi}_2) \\end{equation} where $v_{\\infty}$ is the hyperbolic excess velocity, ${\\phi}_1$ is the incident angle of the trajectory and ${\\phi}_2$ is the exit angle. Equation (1) shows that the anomalous velocity gain $dv$ depends on the difference between the incident latitude and the latitude of the exit trajectory. For example, the NEAR (Near Earth Asteroid Rendezvous) probe approached at low latitude and left on a polar trajectory and its velocity jump was large, whereas Messenger approached and left on an equatorial trajectory and only a very small jump was seen. This led Anderson~$et~al.$~(2008) to make the interesting suggestion that the cause may be somehow related to the Earth's rotation, although they did not suggest a cause. A possible cause is suggested in this paper.  McCulloch~(2007) proposed a model in which the inertial mass reduces slightly as the acceleration decreases: a modification of inertia due to a Hubble-scale Casimir effect (hereafter: MiHsC). This is interesting, because if we take the unusual step of summing all the accelerations seen by NEAR on its flyby, then on its equatorial approach it would see high accelerations as the masses comprising the planet rotate towards and away from it, and many of the acceleration vectors would point at the craft, but on its polar exit trajectory, NEAR would see much less acceleration since the Earth's acceleration vectors, pointing at the spin axis, would not point at the craft. Therefore, MiHsC predicts a lower post-flyby inertial mass for NEAR, which, through conservation of momentum, implies an increase in its speed. In this paper it is shown that the increase in speed predicted by MiHsC agrees quite closely with the observed flyby anomalies. ",
        "conclusions": "Six well-observed Earth flybys with unexplained velocity anomalies were modelled using a theory that assumes that inertia is due to a form of Unruh radiation, and varies with acceleration due to a Hubble-scale Casimir effect.  The theory reproduces three of the observed flyby anomalies within error bars, without the need for adjustable parameters, and explains the recently-observed dependence of the anomalies on the latitude difference between the incident and exit trajectories. The errors for the other flybys were 1, 2 and 3 mm/s.  It should be stressed that the suggested model only seems to apply to unbound trajectories and the reason for this is unknown. Also, the definition of acceleration used here is different to that used in Milgrom~(1994,~1999) and Ignatiev~(2007). The method uses the same physics as McCulloch~(2007) (eg: Eq. 2), but differs in including more detailed information about accelerations. The results of McCulloch~(2007) would be largely unchanged using this enhanced scheme since the Pioneer craft maintained a low latitude with respect to the Sun (Parthasarathy and King, 1991)."
    },
    "0806/0806.3071_arXiv.txt": {
        "abstract": "We have analysed a large data set of \\ion{O}{6}{} absorber candidates found in the spectra of 3702 SDSS quasars, focusing on a subsample of 387 AGN sightlines with an average S/N$\\geq 5.0$, allowing for the detection of absorbers above a rest-frame equivalent width limit of W$_{r} \\geq 0.19 \\AA${} for the \\ion{O}{6}{} 1032 \\AA{} component.\\\\ Accounting for random interlopers mimicking an \\ion{O}{6}{} doublet, we derive for the first time a secure lower limit for the redshift number density $\\Delta N / \\Delta z$ for redshifts $z_{abs} \\geq 2.8$. With extensive Monte Carlo simulations we quantify the losses of absorbers due to blending with the ubiquitous Ly$\\alpha${} forest lines, and estimate the success rate of retrieving each individual candidate as a function of its redshift, the emission redshift of the quasar, the strength of the absorber and the measured S/N of the spectrum by modelling typical Ly forest spectra. These correction factors allow us to derive the 'incompleteness and S/N corrected' redshift number densities of \\ion{O}{6} absorbers :$\\Delta N _{OVI, c} / \\Delta z_{c} (2.8 < z < 3.2) = 4.6 \\pm 0.3$, $\\Delta N _{OVI, c} / \\Delta z_{c} (3.2 < z < 3.6) = 6.7 \\pm 0.8$,and $\\Delta N _{OVI, c} / \\Delta z_{c} (3.6 < z < 4.0) = 8.4 \\pm 2.9$.\\\\ We can place a secure lower limit for the contribution of \\ion{O}{6}{} to the closure mass density at the redshifts probed here: $\\Omega _{OVI} (2.8 < z < 3.2) \\geq 1.9 \\times 10^{-8} h^{-1}$. We show that the strong lines we probe account for over 65\\% of the mass in the \\ion{O}{6}{} absorbers; the weak absorbers, while dominant in line number density, do not contribute significantly to the mass density.\\\\ Making a conservative assumption about the ionisation fraction, $\\frac{OVI}{O}${}, and adopting the \\citet[]{anders1989}{} solar abundance values, we derive the mean metallicty of the gas probed in our search : $\\zeta (2.8 < z < 3.2) \\geq 3.6 \\times 10^{-4} h$, in good agreement with other studies. These results demonstrate that large spectroscopic datasets such as SDSS can play an important role in QSO absorption line studies, in spite of the relatively low resolution. ",
        "introduction": "The cosmic evolution of the majority of the baryons, as probed by absorption features of neutral hydrogen towards the sightlines of QSOs, has come to be well understood over the past years of extensive research : the availability of powerful echelle spectrographs on 8m-class telescopes like HIRES/Keck and UVES/VLT has opened the window for high resolution studies of the Lyman $\\alpha${} forest down to very low column densities \\citep[]{hu1995, lu1996, kim1997, kirkman1997}. On the theoretical front, models incorporating gas dynamics, radiative cooling, and photoionisation have been spectacularly successful in reproducing the observed properties of the parts of the gas traced by neutral HI throughout much of the history of the universe \\citep[]{cen1994, miraldaescude1996, hernquist1996, petitjean1995, dave1997}. The emerging picture paints baryonic gas being encountered in a wide variety of physical conditions. Most of the high redshift Lyman $\\alpha${} forest, which is the main repository for baryons at that time of the universe, exhibits relatively cool temperatures (T$\\sim 10^{4}${} K) in low-to-medium overdensity structures, governed by photoionisation, but there are clear signs of a additional highly ionised component, traced by the prominent absorption lines of CIV and OVI.\\\\ The existence of metal absorption lines in the spectra of quasars, already well established early on in the study of QSO absorption lines \\citep[]{bahcall1968}, and then later linked to the forest (\\ion{C}{4}{} absorbers seen at the same redshift as Lyman $\\alpha${} absorbers by \\citet[]{meyer1987}, poses interesting problems, and in fact allows to address several astrophysical phenomena in a unique setting. First of all, metals allow insight into the structure of the intergalactic medium (IGM) itself : its temperature and small-scale velocity structures can be probed via the width of such lines. Secondly, the existence of metal absorbers reveals the impact of star formation on the IGM - either the metals had to be produced $\\it in{} situ${} by Population III stars in small-scale regions, or a transport mechanism from galaxies like galactic winds into the IGM has to be invoked. Furthermore, studying the ratios of metal column densities of different species and ionic stages can be used to constrain the spectral shape of the metagalactic UV background.\\\\ The hot phase of the IGM (WHIM) with temperatures above T $\\sim 10^{5}${} K is still comparatively poorly understood. Because collisional ionisation of hydrogen becomes significant at these temperatures, Lyman $\\alpha${} becomes a worse tracer of this phase, and thus budgets of the IGM content based solely on hydrogen will underestimate the contribution of the hot gas \\citep[]{simcoe2004}. The high ionisation potential and high abundance renders oxygen a prime candidate for studying the IGM phase under these conditions, possibly tracing shock-heated material. Furthermore, OVI has been long recognised as ideal candidate to trace the low density component (log N$_{HI} < 14.0$) of the photoionised Lyman forest at z$>2${} with optical ground-based spectroscopy \\citep[]{hellsten1998}, with \\ion{C}{4}{} being another possible tracer. The low density in this regime allows for the production of OVI via photoionisation by the intergalactic UV and soft X-ray radiation field \\citep[]{haardt1996}.\\\\ Searches for the signatures of intervening OVI absorbers in QSO spectra have focused on two techniques. A direct line detection, based upon the identification of the characteristic OVI doublet at rest-frame wavelengths of 1032 and 1038 \\AA, at low redshifts must be performed via space-based observations, and ground-based detections hence become possible only for absorber redshifts beyond $z_{abs} \\sim 2.0$, when the lines are shifted redward of the atmospheric cutoff at 3000 \\AA.  \\citet[]{burles1996} conducted the first systematic survey for OVI absorbers at $z_{abs} \\sim$1.0 with FOS aboard HST, studying a sample of 11 QSOs and identifying 9 good candidates, thus establishing that the number density per unit redshift of these absorbers is similar to, if not greater than, that for CIV and MgII absorbers at the same redshift. More recently, \\citet[]{thom2008_I, thom2008_II}{} have used HST/STIS to search for intervening OVI absorbers towards 16 QSOs at $z_{em} = 0.17 - 0.57$, retrieving a sample of 27 candidates with rest-frame equivalent widths $W_{r} \\geq$5m\\AA. The biggest sample of such absorbers at low $z_{abs}${} has been collected by \\citet[]{danforth2007}, who obtain 83 OVI systems by analysing 650 Ly$\\alpha${} absorbers with HST and FUSE spectroscopy, enabling sound statistical analyses of the population.\\\\ Towards higher redshift, despite the advantage of having the features shifted into the wavebands for ground-based optical spectroscopy, this search method is hampered by the ever increasing density of the Ly$\\alpha${} forest, rendering an unambiguous detection of the absorber doublet in its midst more difficult. Hence, most searches have employed high-resolution and high signal-to-noise (S/N) spectroscopy of a handful of QSO sightlines, to build up samples of such OVI absorbers. In order to ensure the highest likelihood of detecting the doublet securely, these studies have focused on absorber and emitter redshifts below $z \\sim$3.0. To date, the number of securely identified OVI absorber components with this method in the redshift range $2.0 \\leq z_{abs} \\leq 3.0${} is of order 350 \\citep[]{simcoe2004, simcoe2006, bergeron2005, carswell2002, fox2007, lopez2007, fox2008}.\\\\ An alternative approach to characterising the metal content of the Lyman forest is the pixel-to-pixel-correlation method, pioneered by \\citet[]{cowie1998}{} for CIV. With this method, statistical correlations between the mean optical depths of various transitions enable an estimate of the metallicity of low density gas. Over the same redshift range as mentioned for the direct line identification searches, \\citet[]{songaila1998, schaye2000} and most recently \\citet[]{aracil2005}{} assert the existence of OVI in gas of optical depth as low as $\\tau _{HI} = 0.1$. For a summary of the results and a discussion of the ambiguity in their interpretations see \\citet[]{aguirre2002} and \\citet[]{pieri2004}.\\\\ In the first paper of our series \\citep[]{frank2007}, we have described how we utilise the large number of high-redshift QSOs in the Sloan Digital Sky Survey (SDSS) spectroscopic sample to explore a different locus of the parameter space for direct searches : low resolution ($R \\sim 1800$) and low to medium S/N at emission redshifts $z_{em} \\geq 2.8$. Frank et al. (2010) describes the sample selection, the search algorithm and presents the candidates found in over 3700 sightlines. While it is clear that the lower quality of the spectra as compared to the high-resolution studies does not allow us to probe such weak absorbers as in the above mentioned searches, the large combined pathlength covered by the SDSS QSOs' sightlines enables us to systematically build up a large sample of strong \\ion{O}{6}{} absorbers at unprecedented high redshifts.\\\\ In this paper, we explore the statistics of the subsample of the candidates towards sightlines with the best S/N, and derive, for the first time at redshift beyond $z_{abs} \\geq 2.8$, secure lower limits for the redshift path density, dn/dz, of strong \\ion{O}{6}{} absorbers. Furthermore, by a Monte-Carlo analysis we can determine the fraction of absorbers lost due to noise and the blending with ubiquitious Lyman forest lines, and hence derive a good estimate for the path density probed with our analysis (section 2). Enforcing a strict lower limit to the {\\it total}{} column density found for our sample, we derive in section 3 a limit for the total mass density of \\ion{O}{6}{} at these high redshifts, and then in section 4 we estimate the minimum metallicity of the gas traced by our absorber sample by limiting the ionisation of oxygen to the maximum value achieved in both the collisionally and photoionised cases. We compare our results with other studies, and conclude in section 5, where we also propose different scenarios for the origin of these absorbers and point out the need for higher resolution and/or signal-to-noise spectra in order to discriminate between them.\\\\ Unless specifically noted otherwise, we use a cosmology with H$_{0}$=71 km s$^{-1}$Mpc$^{-1}$, $\\Omega _{M}$ = 0.27 and $\\Omega _{\\lambda}$=0.73 in this study. Abundances are given by number relative to hydrogen, and we utilise the \\citet[]{anders1989}{} solar abundance values. ",
        "conclusions": "We have analysed the \\citet[]{frank2007}{} data set of \\ion{O}{6}{} absorber candidates found in the spectra of 3800 SDSS quasars, focusing on a subsample of 387 AGN with an average S/N$\\geq 5.0${} per pixel inside the Lyman forest region, allowing for the detection of absorbers above a rest-frame equivalent width limit of W$_{r} \\geq 0.19 \\AA${} for the \\ion{O}{6}{} 1032 \\AA{} component.\\\\ Accounting for random interlopers mimicking an \\ion{O}{6}{} doublet, we derive for the first time for redshifts $z_{abs} \\geq 2.8${} a secure lower limit for the redshift number density $\\Delta n / \\Delta z$ :\\\\ $\\Delta n _{OVI} / \\Delta z (2.8 < z < 3.2) \\geq 1.31 $\\\\ $\\Delta n _{OVI} / \\Delta z (3.2 < z < 3.6) \\geq 0.71 $\\\\ $\\Delta n _{OVI} / \\Delta z (3.6 < z < 4.0) \\geq 0.45 $\\\\ In order to quantify the losses of absorbers due to blending with the ubiquitous Ly$\\alpha${} forest lines, we have modelled typical Lyman forests following canonical line density and column distribution functions, and estimate the success rate of retrieving each individual candidate as a function of its redshift, the emission redshift of the quasar, the strength of the absorber and the measured S/N for the detection spectrum. Furthermore, we have excluded the parts of the redshift pathlength in each spectrum that fall below the S/N which allows for the detection of lines above the rest-frame W$_{r} = 0.19 \\AA$. These correction factors allow us to derive the 'incompleteness and S/N corrected' redshift number densities of \\ion{O}{6} absorbers :\\\\ $\\Delta n _{OVI, c} / \\Delta z_{c} (2.8 < z < 3.2) = 4.6 \\pm 0.3$\\\\ $\\Delta n _{OVI, c} / \\Delta z_{c} (3.2 < z < 3.6) = 6.7 \\pm 0.8$\\\\ $\\Delta n _{OVI, c} / \\Delta z_{c} (3.6 < z < 4.0) = 8.4 \\pm 2.9$\\\\ These values are lower than the ones found in other studies at redshifts $2 \\leq z_{abs} \\leq 2.5$, owing to the fact that the low SDSS resolution does not allow us to probe systems as weak as in those high-resolution and high S/N studies. Simply comparing these studies with our approach by dropping the weak lines out of the statistics for the high-resolution is not possible : multiple weak and thus for SDSS unresolvable components of a complex absorber can form features that will be detected by our algorithm, yet the number of lines identified in such cases strongly depends on the details of the velocity and column density structures.\\\\ Transforming the measured equivalent width of the \\ion{O}{6}{} candidates into a column density is not possible with the limited information obtainable with the SDSS resolution. We can, however, place a secure lower limit for the detectable \\ion{O}{6} column density : log $N_{OVI} \\geq 14.0$. With this limit, we can put a constraint on the contribution of \\ion{O}{6}{} to the closure mass density at the redshifts probed here:\\\\ $\\Omega _{OVI} (2.8 < z < 3.2) \\geq 1.9 \\times 10^{-8} h^{-1}$\\\\ Interestingly, this limit is only a factor of 2-3 below the density estimates of the high-resolution studies mentioned above. Noting that the correction factors due to blending with Ly$\\alpha${} and due to the decrease of the available path because of noise are roughly of the same order, this indicates that the weak \\ion{O}{6}{} absorbers do not contribute significantly to the mass density. Indeed, we have shown that the strong \\ion{O}{6}{} systems that we probe account for $\\geq 65$\\%{} of $\\Omega_{b}$(\\ion{O}{6}). The large number of sightlines allows us for detection of strong \\ion{O}{6}{} absorbers despite the low number-density of the lines, while also overcoming the problem of cosmic variance. While high-resolution, high S/N studies of QSO absorption lines are invaluable for IGM studies, we have shown here that low resolution studies also provide important and complementary results, as long as they cover a large number of sightlines.\\\\ Making the conservative assumption that the ionisation fraction, $\\frac{OVI}{O}${}, can never be above 0.22 (for both the photoionised and the collisionally ionised cases), and adopting the Anders \\& Grevesse solar abundance values, we derive for the mean metallicity of the gas probed in our search :\\\\ $\\zeta (2.8 < z < 3.2) \\geq 3.6 \\times 10^{-4} h$\\\\ in good agreement with other studies. We take this fact, i.e. that the limits on number density, mass density {\\bf and, to a lesser degree, the mean gas metallicity} we derive agree well with these studies, as further {\\bf indication that our search strategy yields sensible results.}\\\\ Having obtained these observational data will allow us to confront them with simulations that model the intergalactic medium's metal enrichment, and serve as important constraint for some of the model parameters. We point out that a similar strategy can easily be applied to similar, even larger data-sets (as e.g. the BOSS survey project for SDSSIII will provide in the near future), which will enable without doubt much smaller (statistical) errors.\\\\  The SDSS is managed by the Astrophysical Research Consortium for the Participating Institutions. The Participating Institutions are the American Museum of Natural History, Astrophysical Institute Potsdam, University of Basel, University of Cambridge, Case Western Reserve University, University of Chicago, Drexel University, Fermilab, the Institute for Advanced Study, the Japan Participation Group, Johns Hopkins University, the Joint Institute for Nuclear Astrophysics, the Kavli Institute for Particle Astrophysics and Cosmology, the Korean Scientist Group, the Chinese Academy of Sciences (LAMOST), Los Alamos National Laboratory, the Max-Planck-Institute for Astronomy (MPIA), the Max-Planck-Institute for Astrophysics (MPA), New Mexico State University, Ohio State University, University of Pittsburgh, University of Portsmouth, Princeton University, the United States Naval Observatory, and the University of Washington."
    },
    "0806/0806.0628_arXiv.txt": {
        "abstract": "We investigate an important question of Hawking-like radiation as seen by an infalling observer during gravitational collapse. Using the functional Schrodinger formalism we are able to probe the time dependent regime which is out of the reach of the standard approximations like the Bogolyubov method. We calculate the occupation number of particles whose frequencies are measured in the proper time of an infalling observer  in two crucially different space-time foliations: Schwarzschild and Eddington-Finkelstein. We demonstrate that the distribution in Schwarzschild reference frame is not quite thermal, though it becomes thermal once the horizon is crossed. We approximately fit the temperature and find that the local temperature increases as the horizon is approached, and diverges exactly at the horizon. In Eddington-Finkelstein reference frame the temperature at the horizon is finite, since the observer in that frame is not accelerated. These results are in agreement with what is generically expected in the absence of backreaction. We also discuss some subtleties related to the physical interpretation of the infinite local temperature in Schwarzschild reference frame. ",
        "introduction": "Understanding quantum radiation from collapsing objects, despite tremendous progress made in the last thirty years, still represents a major challenge in theoretical physics. An asymptotic observer, watching gravitational collapse of a massive object, will start registering radiation of quanta coming from the fields excited by the non-trivial metric of the background space-time. As the radiation progresses, it acquires more and more thermal features. Finally, when the horizon is formed, radiation becomes completely thermal, in agreement with the fact that radiation from a pre-existing horizon must be thermal, as shown for the first time by Hawking \\cite{Hawking:1974sw}. In this picture, radiation is treated in the context of quantum mechanics, while the background is described by classical general relativity. Moreover, the background is held fixed, i.e. the form of the metric and even the mass of the collapsing object (or a black hole) are held fixed. It is extremely difficult to include any form of backreaction in the realistic models (see \\cite{Candelas:1980zt,Russo:1992ax,Vachaspati:2007ur,Mottola:2006ew, Mikovic:1997ty,Mikovic:1996bh} for some attempts).   In the absence of any backreaction, the radiation flux as seen by an asymptotic observer is constant at late times. If integrated over the infinite time (an asymptotic observer sees formation of a horizon only after infinite amount of his time) this flux diverges. However, in classical general relativity, an infalling observer (say an observer which is falling together with a collapsing object) will cross the Schwarzschild radius in the finite amount of his proper time. This time interval corresponds to an infinite interval of the asymptotic observer. The question then is what such an observer would see. The answer will likely strongly depend on the space-time foliation one chooses. If an observer is infalling in Schwarzschild reference frame, his trajectory in any given moment corresponds to an accelerated observer. Though such an observer reaches the Schwarzschild radius in finite time, his acceleration diverges there and the temperature he measures will diverge. Such an observer would need to encounter all of the radiation that an asymptotic observer would see, only in finite time. Will such an observer be burned by an intensive radiation as he is sailing toward the Schwarzschild radius? Not necessarily. A \"particle\" has no invariant meaning. The distribution of produced particles (occupation number as a function of the frequency) strongly depends on the frequency of produced \"particles\". Since proper time coordinates for an asymptotic and infalling observers are different, it is likely that these two different observers will register different distributions of particles.  A very important complementary picture is given by Eddington-Finkelstein coordinates where the singularity at the horizon is absent. Such an observer would still register radiation due to the time dependent metric, however we do not expect that such radiation will be thermal in the whole frequency range. In addition, since the Schwarzschild radius is not a singular point, we do not expect the temperature of such radiation to diverge there.  Unfortunately, arguments of this type have remained only qualitative so far. The reason is that the notion of particles is well defined only in the asymptotically flat regions (e.g Minkowski or Schwarzschild). In the Bogolyubov method, which is widely used in this context, one matches the coefficients between the two asymptotically flat spaces, i.e. Minkowski at the beginning and Schwarzschild at the end of the gravitational collapse. The mismatch of these two vacua gives the number of produced particles. What happens in between is beyond the scope of the Bogolyubov method. Thus, the question what an infalling observer (in different space-time foliations) would register during the collapse can not be answered in the context of the Bogolyubov method. However, recently developed functional Schrodinger formalism, goes beyond the approximations of the Bogolyubov method. In this context, one can successfully attack the question of radiation as seen by an infalling observer, which is the goal of this paper. ",
        "conclusions": "We studied the question of quantum radiation from a collapsing object from a point of view of an infalling observer. Precisely, we calculated the occupation number of particles whose frequencies are measured in the proper time of an infalling observer. As the collapsing object approaches its own Schwarzschild radius it excites fields that are propagating in this dynamical background. In the context of the functional Schrodinger formalism, the information about the excited radiation is contained in the time dependent wave function. The distribution function of the radiation depends on the wave function overlap between the initial state (vacuum) and the state at some later time. This formalism allows us to track the time evolution of the radiation distribution function. This was not possible in approximations usually used in similar setups, e.g. in Bogolyubov method where one matches the coefficients between the two static asymptotically flat spaces, i.e. Minkowski at the beginning and Schwarzschild at the end of the gravitational collapse.  We demonstrated several interesting results. Quantum radiation accompanies gravitational collapse since the metric describing the collapse is time dependent. However, the characteristics of the emitted radiation strongly depend on the space-time foliation of an observer. For the infalling observer in Schwarzschild coordinates the radiation distribution function is not quite thermal, though it becomes thermal when the collapsing object reaches its own Schwarzschild radius. We call such radiation Hawking-like or pre-Hawking radiation (as opposed to thermal Hawking radiation from a pre-existing horizon). In the absence of any backreaction, an asymptotic observer will observe a divergent flux of particles at infinity. However, the number of particles and their energies have different meaning for different observers. The radiation distribution function (i.e. the particle occupation number as a function of frequency) depends on the time coordinate that the observer is using. Naively, an infalling observer would need to encounter an infinite number of particles, but in finite amount of his proper time. However, it is only the occupation number of particles that is divergent, not the actual number of detected particles. An infalling observer can not detect particles whose wavelength is larger than his detector and therefore would miss most of them.  By fitting the approximate temperature corresponding to the radiation distribution function, we showed that the local temperature near the Schwarzschild radius, $R_s$, is large and diverges exactly at $R_s$. This is expected since this local temperature has to be infinitely resifted in order to get a finite temperature of radiation as seen by the asymptotic observer. After all, an infalling observer in Schwarzschild coordinates is accelerated at any given moment, with acceleration which diverge exactly at $R_s$. Such an accelerated observer should ``see\" divergent temperature at the horizon in agreement with the Unruh effect \\cite{Unruh:1976db}. This still does not mean that the horizon is a violent place for the infalling observer since local vacuum polarization energy density is usually thought to cancel out the divergence of the temperature energy density, which yields a finite total energy density..  The situation in Eddington-Finkelstein coordinates is quite different. This coordinate system is not singular at the horizon, and an infalling observer in these coordinates is not accelerated. We find that the distribution of the particle occupation number is not thermal in the whole frequency range. By fitting the temperature only in the low frequency range ($\\bar{\\omega}<R_s^{-1}$) we find the temperature $T \\approx 0.7/R_s$. This temperature is somewhat greater than the Hawking temperature as measured at infinity by an asymptotic observer $T_\\infty = 1/(4\\pi R_s)$,  which is in agreement with results obtained in different approaches (see e.g. \\cite{Brynjolfsson:2008uc}). However, the spectrum never becomes thermal in the whole frequency range, so it is difficult to talk about the temperature as seen by such an observer.    \\appendix"
    },
    "0806/0806.2288_arXiv.txt": {
        "abstract": "{  A thorough critical literature survey has been carried out for reliable measurements of oxygen and neon abundances of planetary nebulae (PNe) and \\hiir. By contrasting the results of PNe and of \\hiir, we aim to address the issues of the evolution of oxygen and neon in the interstellar medium (ISM) and in the late evolutionary phases of low- and intermediate-mass stars (LIMS), as well as the currently hotly disputed solar Ne/O abundance ratio. Through the comparisons, we find that neon abundance and Ne/O ratio increase with increasing oxygen abundance in both types of nebulae, with positive correlation coefficients larger than 0.75. The correlations suggest different enrichment mechanisms for oxygen and neon in the ISM, in the sense that the growth of neon is delayed compared to oxygen. The differences of abundances between PNe and \\hiir\\, are mainly attributed to the results of nucleosynthesis and dredge-up processes that occurred in the progenitor stars of PNe. We find that both these $\\alpha$-elements are significantly enriched at low metallicity (initial oxygen abundance $\\la$ 8.0) but not at metallicity higher than the SMC. The fact that Ne/O ratios measured in PNe are almost the same as those in \\hiir, regardless of the metallicity, suggests a very similar production mechanism of neon and oxygen in intermediate mass stars (IMS) of low initial metallicities and in more massive stars, a conjecture that requires verification by further theoretical studies. This result also strongly suggests that both the solar neon abundance and the Ne/O ratio should be revised upwards by $\\sim$0.22~dex from the Asplund, Grevesse \\& Sauval values or by $\\sim$0.14~dex from the Grevesse \\& Sauval values.} ",
        "introduction": "It is generally assumed (e.g. \\citealt{henry1989}) that $\\alpha$-elements, such as oxygen and neon, are not altered significantly by the various nucleosynthesis and dredge-up processes that occurred during the late evolutionary stages of LIMS, which range in mass from about 0.8 to 8\\,$M_{\\sun}$. Therefore, the abundances of $\\alpha$-elements measured in the descendants of LIMS, such as Asymptotic Giant Branch (AGB) stars and PNe, should reflect the metallicity, $Z$, of the ISM, from which the parent stars were formed. Under the assumption, abundances determined for PNe have been widely used to constrain the chemical evolution of the Milky Way and of nearby galaxies.  Recently, there is however some observational evidence that the above assumption is questionable, especially in low $Z$ galaxies, where it has been found that oxygen abundances of PNe are higher than the average value of \\hiir\\, (e.g. SMC and LMC; Leisy \\& Dennefeld 1996, 2006). The enrichment of oxygen in PNe in lower metallicity environments becomes more apparent by comparing the Magellanic Clouds and the Galaxy. On the other hand, oxygen destruction has been found in PNe evolved from massive progenitors (e.g. NGC 6302; \\citealt{pottasch1999}; we define the progenitor stars which were initially more massive than about 5\\,$M_{\\sun}$ as massive progenitors, where the ON cycle begins to occur). There is also evidence that asymmetric PNe may be enriched in neon, as suggested by the observed relationships between the abundances and morphological types of PNe (\\citealt{corradi1995}; \\citealt{stanghellini2000}).  The various nucleosynthesis [e.g. the Hot Bottom Burning (HBB)] and dredge-up processes occurred in the late evolutionary epochs of LIMS, which eventually determined the chemical composition of the ejected nebulae. These depend strongly on the initial metallicity as well as mass. Recent theoretical calculations have indeed revealed several channels of altering oxygen and/or neon abundances by some processes that occurred in the PN progenitors. Oxygen destruction, for example, can occur through the CNO cycle specifically via the ON cycle in massive progenitors, while the 3rd dredge-up can efficiently modify the nebular oxygen abundance by transporting the freshly manufactured material to the stellar surface. \\cite{charbonnel2005} emphasized that the surface abundance of oxygen in massive progenitors is a result of competition between the efficiencies of the 3rd dredge-up (for production) and the HBB (for destruction). We note that neon production may also become significant for a narrow range of stellar mass (\\citealt{KL2003}).  One of the purposes of the current study is to clarify whether oxygen and neon abundances are significantly modified in LIMS and if so, the modifications that affect our current understanding of the chemical evolution of galaxies. To address this issue, we have collected all recent reliable determinations of oxygen and neon abundances of PNe and \\hiir\\, and present a thorough analysis of available data in this paper.  As a consequence of the application of a time-dependent 3D hydrodynamical model of the solar atmosphere in place of the earlier 1D model, the metal content in the solar convection zone has been reduced by almost a factor of two. Amongst them, abundances of key elements C, N and O have been lowered by approximately 0.2--0.3~dex from the earlier widely adopted values (see a review by \\citealt{asplund2005}; henceforth AGS05). However, those large downward revisions have caused serious conflicts with other solar physics studies. To fit the helioseismological measurements, for example, \\citet{bahcall2005} suggested that the neon abundance needs be raised from 7.84 to 8.29 or the Ne/O ratio from 0.15 to 0.42. Recent Chandra measurements of 21 nearby solar-type stars by \\cite{drake05} concluded that the solar Ne/O ratio should be around 0.41. Finally, observations of Galactic PNe yield an average Ne/O ratio of about 0.25 from collisionally excited lines (\\citealt{WL07}). As we shall show in Section~3, the depletion or enrichment of oxygen in PNe is insignificant at solar metallicity. Thus accurate determinations of oxygen and neon abundances of PNe and \\hiir\\, can be used to constrain and calibrate solar oxygen and neon abundances and their ratio. This serves as the goal of the current paper.  The paper is organized as follows. In Section 2, we present our critical literature survey. In Section 3, we compare the oxygen and neon abundances in PNe and in H~{\\sc ii} regions and discuss their possible evolution mechanisms and histories. In Section 4, we discuss the neon abundances and the Ne/O ratio in the Sun. In the last Section, we conclude by summarizing the main results. ",
        "conclusions": "In this letter, we present a critical literature survey of recent measurements of oxygen and neon abundances of PNe and of \\hiir\\, in the Milky Way and other nearby galaxies. We find that there is significant oxygen and neon production in LIMS at metallicities lower than the SMC, but not at higher metallicities. We find that oxygen destruction is probably insignificant. We show that neon and oxygen are probably enhanced by the same amounts in PNe at low metallicities, a result not predicted by the current theory of nucleosynthesis for LIMS.  We find that the Ne/O ratio increases with increasing oxygen abundance in PNe and in \\hiir, suggesting a different enrichment history of neon and oxygen in the ISM and thus probably different production mechanisms of these two $\\alpha$-elements in massive stars, as predicted by current theoretical calculations.  Both PNe and \\hiir\\, in the Galactic disk yield a consistent Ne/O ratio of 0.25, higher than the solar value of 0.18 (GS98) or 0.15 (AGS05). We suggest that the solar Ne/O ratio and the absolute neon abundance need to be revised upwards by about 0.22~dex from the values of AGS05 or by 0.14~dex from those of GS98. A recent discussion based on the Ne/S ratios measured in \\hiir\\, in nearby galaxies by \\citet{rubin2008} also points to an upwards revision to the solar neon abundance.  $$\\\\$$ The work is partially supported by Grant \\#10325312 of the National Natural Science Foundation of China. We would also like to thank the referee for useful comments."
    },
    "0806/0806.2846_arXiv.txt": {
        "abstract": "We present the X-ray source catalog in the \\subaru /\\xmm\\ deep survey. A continuous area of 1.14 \\de\\ centered at R.A.\\ = 02h18m and Dec.\\ = --05d is mapped by seven pointings with \\xmm\\ covering the 0.2--10 keV band. From the combined images of the EPIC pn and MOS cameras, we detect 866, 1114, 645, and 136 sources with sensitivity limits of $6\\times10^{-16}$, $8\\times10^{-16}$, $3\\times10^{-15}$, and $5\\times10^{-15}$ \\ergs\\ in the 0.5--2, 0.5--4.5, 2--10, and 4.5--10 keV bands, respectively, with detection likelihood $\\geq$7 (corresponding to a confidence level of 99.91\\%). The catalog consists of 1245 sources in total including 32 extended-source candidates. The averaged \\logn s are in good agreement with previous results, bridging the flux range between \\chandra\\ deep surveys and brighter surveys. The \\logn s show significant spatial variation among pointings on a scale of 0.2 \\de . Analyzing the auto correlation function, we detect significant clustering signals from the 0.5--2 keV band sample, which can be fit with a power law form $(\\theta/\\theta_c)^{-0.8}$ with a correlation length of $\\theta_c=5.9$\\byo $^{+1.0^{{\\prime}{\\prime}}}_{-0.9^{{\\prime}{\\prime}}}$ when the integral constraint term is included. In the 2--10 keV band, however, the clustering is not significant with a 90\\% upper limit of $\\theta_c < 1.5$\\byo . ",
        "introduction": "The \\subaru /\\xmm\\ Deep Survey (SXDS; \\citealt{sek07}) is, along with COSMOS \\citep{sco07,has07}, one of the largest multi-wavelength survey projects with an unprecedented combination of depth and sky area over a contiguous region of $>$ 1 \\de . The main aims of the SXDS are to make an accurate measurement of the global properties of the universe without being affected by cosmic variance and to reveal the evolution of the large scale structure. The SXDS consists of a wealth of multi-wavelength data taken by the most modern observing facilities; X-ray imaging/spectroscopic data in the 0.2--10 keV band taken by the European Photon Imaging Camera (EPIC; \\citealt{str01,tur01}) onboard \\xmm\\ \\citep{jan01}, multi-color (B, V, R, i$^\\prime$, z$^\\prime$) deep optical images by \\scam\\ on the \\subaru\\ telescope \\citep{fur07}, deep near infrared maps (J, H, K) observed as the Ultra Deep Survey (UDS) in the United Kingdom Infrared Deep Sky Survey (UKIDSS; \\citealt{law07}), mid- and far-infrared data (3.6--160 $\\mu$m) taken with the \\spitzer\\ Space Observatory\\footnote{the \\spitzer\\ Legacy Survey and part of the \\spitzer\\ Wide-area InfraRed Extragalactic (SWIRE) survey}, the submillimeter (850 $\\mu$m) map by the SCUBA Half Degree Extragalactic Survey (SHADES; \\citealt{mor05})\\footnote{The SXDS field is also the field of SCUBA-2 Cosmology Legacy Survey.}, and the deep radio image (1.4 GHz) by the Very Large Array \\citep{sim06}. The SXDS field is centered at R.A.\\ = 02h18m and Dec.\\ = --05d and the total field of the Subaru \\scam\\ and \\xmm\\ EPIC images covers an area of 1.3 and 1.14 \\de, respectively. The overall survey design and details of each survey are summarized in \\citet{sek07}.  The data of \\xmm\\ constitute a major component of the SXDS project. X-ray surveys are a powerful tool to trace the cosmological evolution of active phenomena in the universe, including Active Galactic Nuclei (AGNs) and clusters/groups of galaxies. The main constituents of X-ray sources that make up the X-Ray Background (XRB) are AGNs (for a recent review, see \\citealt{bra05}). Their dominant populations are ``obscured'' AGNs \\citep{set89}, where the central engine is blocked by dust and/or gas in the line of sight. In particular, surveys by hard X-rays above 2 keV are the most efficient approach to detect these obscured AGN populations of various luminosity classes with least bias, thanks to their strong penetrating power against the photo-electric absorption of matter and small contamination from stars. In fact, the surface density of AGNs detected in the most sensitive X-ray surveys \\citep[][]{ale03} far exceeds that detected in the optical bands \\citep[e.g.,][]{wol03}, thus providing us with the most complete and clean samples for AGN studies including heavily obscured, low luminosity ones that dominate the whole AGN populations by number.  The tight correlation between the mass of a Super Massive Black Hole (SMBH) in the galactic center and that of the galactic bulge \\citep{fer00,geb00} indicates strong links between the growth of SMBHs and star formation in the past. Since AGNs are phenomena that mark the process of growth of SMBHs in galactic centers, to elucidate the evolution of AGNs is a fundamental issue for understanding the history of the universe. Combination of ultra-deep pencil beam surveys and large area surveys have revealed the cosmological evolution of the X-ray luminosity function of AGNs \\citep[][]{ued03,laf05,bar05,has05}. The global accretion history of the universe is similar to that of star formation \\citep[see e.g.,][]{fra99,mar04}. An important finding is that the number density of more luminous AGNs have a peak at higher redshifts compared with less luminous ones. This behavior is called ``down-sizing'' or ``anti-hierarchical'' evolution, which is opposite to a naive expectation from the standard structure formation theory of the universe. Similar trends were also obtained in the star formation history \\citep[e.g.,][]{cow96,kod04}. These facts imply two modes of ``co-evolution'' of galaxies and SMBHs for different masses. It could be explained by the feedback from supernova and AGNs \\citep[e.g.,][]{gra04}. Deep multi-wavelength surveys enable us to observe various populations in different evolutionary stage, giving the best opportunities to investigate how galaxies and SMBHs co-evolved with cosmic time.  Despite this major progress on AGN evolution in recent years, there are several critical issues to be resolved by current and future X-ray surveys, even if we limit the scope only to X-ray detected AGNs. Immediate objectives include (1) whether or how the fraction of obscured AGNs evolves with redshift, (2) the number density of Compton thick AGNs (those whose line-of-sight absorption exceeds $N_{\\rm H} > 10^{24}$ cm$^{-2}$ ), (3) the evolution of the AGN luminosity function at high redshifts ($z \\simgt 4$) to be compared with the results from optical QSO surveys, and (4) the evolution of clustering properties of AGNs as a function of luminosity and type. The SXDS X-ray survey, one of the few ``wide'' and ``deep'' surveys, will give important steps to answer these questions and to establish the average properties of the AGN populations. Population synthesis models of the XRB utilize the results from deep surveys in a limited area at the faintest flux levels \\citep{ued03,gil07}, which would be subject to cosmic variance when we discuss the XRB spectrum with $\\approx$10\\% accuracy. This is crucial when we discuss a relatively minor contribution of some X-ray sources, such as Compton thick AGNs. To constrain the statistical properties of rare objects, such as high redshift QSOs, much larger cosmic volume is necessary than in the currently available deep surveys. By surveying a continuous area with a large depth, we can also study the evolution of the large scale structure probed by AGNs. It helps us to understand the physical conditions that triggers AGN activity for different luminosity class, type (1 or 2), and redshift. At the flux limits of the SXDS, we can detect the most dominant populations of X-ray sources that contribute to the XRB. In all of these studies, high completeness of identification (redshift determination) is crucial. The multi-wavelength data are particularly useful in identifying X-ray sources even if optical spectroscopy is difficult for a population of ``optically faint'' AGNs \\citep{ale01}.  In this paper, we describe the X-ray data of the SXDS based on the \\xmm\\ survey performed in 2000 and 2003. We present the whole X-ray catalog, and basic statistical properties of the detected X-ray sources. The X-ray data have been used in a number of studies in the SXDS projects, detection of ultra luminous X-ray sources in nearby galaxies \\citep{wat05}, mid-IR and radio selected Compton-thick AGNs \\citep{mar07}, radio sources \\citep{sim06}, Lyman $\\alpha$ sources \\citep{sai07,ouc07}, and optical variability selected AGNs \\citep{mor07}. A series of papers on optical identification of the X-ray sources are forthcoming (Akiyama \\etal , in preparation). The paper is organized as follows. \\S~2 summarizes the observation of the \\xmm . In \\S~3 we describe the detailed procedure of data analysis, and present the source list and its statistics. \\S~4 and \\S~5 shows the results of \\logn s and auto correlation functions, respectively. In \\S~6 long-term time variability of X-ray sources is studied in a selected field. \\S~7 summarizes the conclusions. Throughout the paper, the cosmological parameters \\cosmo\\ = (100$h$ km s$^{-1}$ Mpc$^{-1}$, 0.3, 0.7) are adopted. ",
        "conclusions": "We have presented the first X-ray source catalog and their basic X-ray properties in the SXDS field, based on the seven \\xmm\\ pointings performed between 2000 and 2003 that cover a continuous area of 1.14 \\de\\ in the 0.2--10 keV band. The catalog consists of 1245 sources in total, consisting of those detected either in the 0.3--0.5 keV, 0.5--2 keV, 2--4.5 keV, 4.5--10 keV, 0.5--4.5 keV, and 2--10 keV with detection likelihood larger than 7. From empirical sensitivity maps derived by simplified simulation, we obtain \\logn s in various energy bands bands, which bridge the flux range between \\chandra\\ deep surveys and brighter surveys. Clustering properties of X-ray sources are also studied by means of auto correlation function. We detect significant signals that can be fit by $(\\theta/\\theta_c)^{-0.8}$ with $\\theta_c=5.9$\\byo $^{+1.0^{{\\prime}{\\prime}}}_{-0.9^{{\\prime}{\\prime}}}$ for the 0.5--2 keV selected sample at an estimated median redshift of $z_{\\rm eff}\\approx 1.3$, and an 90\\% upper limit of $\\theta_c < 1.5$\\byo\\ for the 2--10 keV selected sample at $z_{\\rm eff}\\approx 1.1$. Two pointing data separated by 2.5 years indicate the importance of flux variability on a time scale of years. Our results establish the average properties of X-ray sources at fluxes where a majority of the XRB is produced, being least affected by cosmic variance. The data presented here constitute a major component of the SXDS project and shall be used for a number of research programs in combination with other wavelengths data."
    },
    "0806/0806.1462_arXiv.txt": {
        "abstract": "The \\ion{Ar}{17} X-ray line group principally due to transitions $1s^2 - 1s2l$ ($l=s$, $p$) near 4~\\AA\\ was observed in numerous flares by the RESIK bent crystal spectrometer aboard {\\it CORONAS-F} between 2001 and 2003. The three line features include the \\ion{Ar}{17}  $w$ (resonance line), a blend of $x$ and $y$ (intercombination lines), and $z$ (forbidden line), all of which are blended with \\ion{Ar}{16} dielectronic satellites. The ratio $G$, equal to $[I(x)+I(y) + I(z)]/I(w)$, varies with electron temperature $T_e$ mostly because of unresolved dielectronic satellites. With temperatures estimated from {\\it GOES} X-ray emission, the observed $G$ ratios agree fairly well with those calculated from {\\sc chianti} and other data. With a two-component emission measure, better agreement is achieved. Some \\ion{S}{15} and \\ion{S}{16} lines blend with the Ar lines, the effect of which occurs at temperatures $\\gtrsim 8$~MK, allowing the S/Ar abundance ratio to be determined. This is found to agree with coronal values. A nonthermal contribution is indicated for some spectra in the repeating-pulse flare of 2003 February~6. ",
        "introduction": "High-resolution soft X-ray spectral observations have provided important information about the physical properties of solar flares and active regions, including electron temperatures $T_e$, densities $N_e$, and ionization state. Those from the bent crystal spectrometer RESIK (REntgenovsky Spektrometr s Izognutymi Kristalami) on the Russian {\\it CORONAS-F} spacecraft \\citep{syl05a} cover the range 3.3--6.1~\\AA, which includes strong emission lines of Si, S, and Ar ions and continuum, and weaker lines of Cl and K ions. The wavelength resolution $\\Delta \\lambda$ is between 8~m\\AA\\ (3.4~\\AA) and 17~m\\AA\\ (6.1~\\AA). The instrument operated successfully between 2001 (shortly after the launch of {\\it CORONAS-F} on July~31) and May 2003, obtaining numerous flare and active region spectra. RESIK spectra have already been used for K, Ar, Cl, S, and Si abundance determinations in flares \\citep{syl06}. Lines in the 3.3--6.1~\\AA\\ range have not been well observed by previous spectrometers with the exception of the He-like S (\\ion{S}{15}) lines at $\\sim 5$~\\AA. Of particular interest here are the He-like Ar (\\ion{Ar}{17}) lines between 3.94~\\AA\\ and 4.00~\\AA\\ and associated dielectronic satellites emitted by Li-like Ar (\\ion{Ar}{16}). There are four prominent \\ion{Ar}{17} lines, due to transitions from the $1s2p\\,^1P_1$, $1s2p\\,^3P_2$, $1s2p\\,^3P_1$, and $1s2s\\,^3S_1$ levels to the ground level $1s^2\\,^1S_0$ (designated lines $w$, $x$, $y$, and $z$ respectively), with wavelengths 3.949~\\AA, 3.966~\\AA, 3.969~\\AA, and 3.994~\\AA. Dielectronic \\ion{Ar}{16} satellites, particularly $q$ (3.981~\\AA), $k$ (3.990~\\AA), and $j$ (3.994~\\AA) feature prominently.  The \\ion{Ar}{17} lines have been observed during a few solar flares with the Flat Crystal Spectrometer (FCS) on {\\it Solar Maximum Mission}, with the \\ion{Ar}{16} $k$ satellite being resolved from \\ion{Ar}{17} line $z$ but $j$ indistinguishable from line $z$. RESIK flare spectra do not resolve lines $k$ and $z$ or the \\ion{Ar}{17} $x$ and $y$ lines, so there are three main features, the blend of line $w$ with unresolved weak \\ion{Ar}{16} satellites on the long-wavelength side of $w$ ($w'$), the blend of lines $x$, $y$, and other weak \\ion{Ar}{16} satellites [$(x+y)'$], and the blend of line $z$ with the two prominent satellites $j$ and $k$ ($z'$). The relative fluxes of these three line features can thus be examined over a variety of flare conditions, RESIK's sensitivity being much improved (by a factor $\\sim 60$) over the FCS. In particular, the temperature dependence of the ratio $G$, which in terms of the observed features is $[I(x+y)' + I(z')]/I(w')$, can be examined. At temperatures $\\gtrsim 8$~MK, the effect of sulphur lines blending with the three \\ion{Ar}{17} line features becomes apparent on the $G$ ratio. These are the $1s-3p$ (Ly-$\\beta$) line of \\ion{S}{16} (3.991~\\AA) and $1s^2-1s5p$ line of \\ion{S}{15} (3.998~\\AA), blending with the \\ion{Ar}{17} $z$ line, and the $1s^2-1s6p$ line of \\ion{S}{15} (3.949~\\AA) blending with the \\ion{Ar}{17} $w$ line. These lines were not considered in an earlier analysis of FCS spectra \\citep{phi93}. ",
        "conclusions": "We have discussed RESIK observations of the $G$ ratio for the \\ion{Ar}{17} lines at $\\sim 4$~\\AA\\ over a wide range of activity. For 98 flare spectra as well as four {\\it SMM}\\/ FCS spectra, there is a clear trend of $G$ that is close to the theoretical ratio. The agreement is improved for a two-component DEM. From this it is deduced that the Ar/S abundance ratio is close to 4.9, its coronal value, so indicating a coronal origin for the flare plasma. Two spectra with anomalous $G$ during the 2003 February~6 flare are noted, one of which is exactly at the time of an impulsive, nonthermal burst in {\\it RHESSI} 12--25~keV emission. There is thus some suggestion that part of the \\ion{Ar}{17} emission is nonthermal in origin."
    },
    "0806/0806.3521_arXiv.txt": {
        "abstract": "{% This article quickly summarizes the performances and results of the GOLF/SoHO resonant spectrometer, thus justifying to go a step further. We then recall the characteristics of the multichannel resonant GOLF-NG spectrometer and present the first successful performances of the laboratory tests on the prototype and also the limitations of this first technological instrument. Scientific questions and an observation strategy are discussed.} ",
        "introduction": "The SoHO satellite first milestone of the ESA Horizon 2000 programme, is a real success due to the simultaneous presence of internal and external solar probes. It enriches the present transition period between a classical view of stars and a dynamical view that we are beginning to constrain. The success is partly due to helioseismology which reveals the properties of millions of modes not yet within reach for any other type of stars. This fact enables a strong scientific return with some problems being solved and totally new questions emerging.  Seismology of the Sun and stars will be pursued during the next decade with a strong impact in Astronomy.  The successors of SoHO are SDO in the United States (launch in 2009) and Solar Orbiter in Europe (launch in 2014-2015). But these missions have not been designed to probe the deep solar interior.  After more than a solar cycle of continuous observation with SoHO and three decades with the BiSON ground network, it is time to critically review the results obtained by the European community to highlight the advantages and the limitations of the resonant spectrometer instrument and to deduce which directions will lead to improvements.  In section 2, we recall the characteristics of the GOLF instrument and its results. The third section deals with the objectives of the GOLF-NG instrument, successor of GOLF,  the performances obtained in the laboratory and the necessary improvements to get a scientifically relevant instrument.  The last section will be devoted to (1) the questions that we need to solve in the future,   (2) the solar proposal at the ESA Cosmic Vision level and (3) the way we might observe during the next decade. ",
        "conclusions": ""
    },
    "0806/0806.1524_arXiv.txt": {
        "abstract": "We use a set of semi-numerical simulations based on Zel'dovich approximation, friends-of-friends algorithm and excursion set formalism to generate reionization maps of high dynamic range with a range of assumptions regarding the distribution and luminosity of ionizing sources and the spatial distribution of sinks for the ionizing radiation. We find that ignoring the inhomogeneous spatial distribution  of regions of high gas density where recombinations are important -- as is often done in studies of this kind -- can lead to misleading conclusions regarding the topology of reionization, especially if reionization occurs in the photon-starved regime suggested by \\lya\\ forest data. The inhomogeneous spatial distribution of recombinations significantly reduces the mean free path of ionizing  photons and the typical size of coherently ionized regions. Reionization proceeds then  much more as an outside-in process. Low-density regions  far from ionizing sources become ionized before regions of  high gas density not hosting sources of ionizing radiation.  The spatial  distribution of sinks of ionization radiation also significantly affects shape and amplitude the power spectrum of fluctuations of 21cm emission. The slope of the 21cm power spectrum as measured by upcoming 21cm experiments should be able to  distinguish to what extent the topology of reionization proceeds outside-in or inside-out while the evolution of the amplitude of the   power spectrum with increasing ionized mass fraction should be sensitive to the spatial distribution  and the luminosity of ionizing sources. ",
        "introduction": "The reionization of neutral hydrogen is an important milestone in the evolution of the Universe. The epoch of reionization has received  a major boost  of attention  recently due  to a series of observational advances which suggest that the process is complex and that the reionization of hydrogen extends  over wide redshift range from $6 \\lesssim z \\lesssim 15$ (for reviews see \\citeNP{cf06a,fob06}).  We are about  to enter  an  exciting phase as  planned 21cm observations are expected to settle the  questions when and how the Universe was reionized. It  thus timely  to develop more accurate and detailed analytical and numerical models  in order to extract the maximum information about the physical processes relevant  for reionization from the expected large and complex  future data sets.  Currently operating and upcoming low-frequency radio observations (e.g., GMRT\\footnote{http://www.gmrt.ncra.tifr.res.in/}, 21CMA\\footnote{http://web.phys.cmu.edu/~past/}, MWA\\footnote{http://www.haystack.mit.edu/ast/arrays/mwa/}, LOFAR\\footnote{http://www.lofar.org/}, SKA\\footnote{http://www.skatelescope.org/})  of redshifted 21cm emission of neutral hydrogen should also probe the  topology of the neutral (or ionized) regions at high redshifts.   Unfortunately, modelling the expected data sets is not straightforward  because of the dauntingly wide range of physical scales involved and our  lack of knowledge of many details of the relevant physical processes.   Full numerical simulations including radiative transfer effects  are still computationally extremely challenging.  Modelling the  smallest mass haloes contributing  to the  ionizing emissivity at early epochs (with a total mass  $M \\sim 10^8 M_{\\odot}$) requires linear scales $\\lesssim 0.1$Mpc  while at the same time,  the size of the simulated regions need to extend over  $100$ Mpc or more in order to probe the largest coherently ionized regions in the final stages  of reionization. Despite such challenging requirements, considerable  progress has been made in performing radiative transfer simulations of  ionization maps of representative regions of the Universe (see e.g. \\citeNP{gnedin00,cfw03,pn05,impmsa06,ica++06,imsp07,mlz+07}). Most radiative transfer simulations are, however, still  rather limited in dynamic range and equally important also limited in their ability to explore the large parameter space of plausible assumptions regarding the spatial distribution and time evolution of the ionizing emissivity.  This is one of the reasons why modelling  the evolution of ionized regions  analytically using excursion-set-like formalisms has become a widely used and useful tool ({\\it e.g.} \\citeNP{fzh04b}). Such methods are well adapted to  obtain estimates of  the size distribution of ionized regions for arbitrary  models of the luminosity function and time evolution of the ionizing emissivity. One has, however, to keep in mind that these models make  quite drastic simplifying assumptions. The shapes of ionized bubbles are e.g. assumed to be spherical and the (relative) spatial distribution of sources and sinks of ionizing radiation are not properly taken into account.  Ideally,  one would like to compare realistic models of  the ionization state  of the IGM  with a large dynamic range for a wide range of assumptions  with future observations. For this purpose a variety of  semi-numeric formalisms  have recently been proposed which are based on performing an excursion-set formalism on the initial Gaussian random field. The models  predict the spatial distribution of the (integrated) ionizing emissivity as well as the spatial distribution of  ionized regions \\cite{zlm+07,mf07,gw08}. They incorporate many of the relevant physical  processes and allow the modeller to produce  21cm maps for representative volumes of the Universe with a modest computational effort.  These studies suggest that  reionization proceeds strictly inside-out with dense regions ionized first and reionization slowly progressing into the large underdense region as time goes on \\cite{fzh04b,wm07,mf07,mlz+07}.  This  appears, however,  in conflict with what is expected and observed for the post-overlap phase where the low-density regions are found to be highly ionized while high-density  regions remain neutral because of their high recombination rate \\cite{mhr00,wl03,cf05,cf06b}.  These neutral regions determine the photon mean free path and manifest themselves as Lyman-limit systems in QSO absorption spectra. Based on the this low-redshift intuition derived from studying  the intergalactic medium at $z\\sim 2-4$ with Ly$\\alpha$ forest data,  one is thus  drawn to the conclusion that reionization  must have proceeded -- at least to some extent -- outside-in rather than inside-out in the final stages. The obvious suspect for resolving this apparent  contradiction is the role recombinations play  in these simulations. Most of these models assume a spatially uniform distribution  of recombinations and hence do not take into account the self-shielding and shadowing of high-density regions. \\citeN{fo05}  have attempted to model this   by introducing the concept of recombination-limited bubbles in analytic studies of  the size distribution of ionized bubbles, which has been implemented in simulations by \\citeN{lzfmhz08}. As we will show in this paper it is   important to realistically  model the spatial distribution of the sinks of ionizing radiation due to recombinations when modelling the topology of reionization.  Many of the models also assume that reionization occurs rather fast diminishing the relative importance of recombinations. This appears,  however, to be in conflict with the ionizing emissivity   inferred  from the opacity of the \\lya\\ forest in QSO absorption spectra which suggests that reionization occurs slowly in a photon-starved regime \\cite{bh07}.   We will study here the effects of the inhomogeneous spatial distribution of recombinations on ionization maps and present a consistent picture of reionization combining the concepts of growing bubbles in the pre-overlap phase with the expected  presence of neutral clumps  in the post-overlap phase. Our modelling  is similar in spirit to  other semi-numerical  models of this kind  and in many aspects we (need to) make similar  approximations  and simplifications.  In order to determine whether a high-density clump can remain neutral or self-shielded against ionizing radiation, it is  necessary to determine its position with respect to the nearest sources of ionizing photons. An important requirement for a realistic model  of the spatial distribution of density-dependent recombinations is thus  (i) a realistic representation of the baryon distribution  and more importantly, (ii) the location of the sources of ionizing sources with respect to the density field. Note that we will here  concentrate on a qualitative  understanding of the physical effects of a spatially inhomogeneous distribution of recombinations on the topology of reionization.   The paper is organized as follows: We describe our method for generating the ionization maps in Section 2. Section 3 discusses our  main results  for the modelling of a single source and representative volumes  of the Universe. In Section 4 we check the consistency of our modelling with \\lya\\ forest data. In Section 5 we present predictions for the evolution of the  power spectrum and probability distribution of 21cm emission and discuss prospects for the first generation low-frequency instruments LOFAR and MWA. Section 6 contains our conclusions. Throughout the paper,  we assume a flat Universe with cosmological parameters  $\\Omega_m = 0.26$, $\\Omega_{\\Lambda} = 0.74$, $\\Omega_b h^2 = 0.022$, and $h=0.73$.  The  parameters defining the linear dark  matter power spectrum we use are $\\sigma_8=0.9$, $ n_s=1$, $\\de n_s/\\de \\ln k =0$ \\cite{vhl06}.  \\begin{figure*} \\rotatebox{270}{\\resizebox{0.4\\textwidth}{!}{\\includegraphics{denplushaloplot.ps}}} \\caption{The density field (left panel) and the location of collapsed haloes  (right panel) at $z=6$ for our fiducial simulation (see text for details).  The thickness of the slice shown is 1$h^{-1}$Mpc.} \\label{fig:denplushalo} \\end{figure*} ",
        "conclusions": "We have used here semi-numerical simulations to investigate  the role of the  spatial distribution of sinks and sources of ionizing radiation on the topology of hydrogen  reionization.  Our main results are the following.   \\begin{itemize}   \\item The combination of Zel'dovich approximation, halo-finder and excursion set formalism  is a powerful tool to calculate realistic ionization maps with high dynamic range  at a  very  moderate computational   cost.  \\item  Enforcing  consistency with the \\lya\\ forest data helps to significantly shrink  the otherwise rather unconstrained parameter space of models of reionization.  In the  photon-starved regime of reionization suggested by the \\lya\\ forest data recombinations  are much more important than in models with high ionizing emissivity where reionization occurs quickly.  Taking into account a realistic spatially inhomogeneous distribution of sinks  of ionizing radiation has a large effect on the topology of  reionization in the photon-starved regime.  \\item  Initially reionization proceeds inside-out with the high-density regions hosting  the sources of ionizing sources becoming ionized first. In the later stages of  photon-starved reionization the sinks of ionizing region in our models remain neutral and  reionization proceeds deep into the underdense regions before slowly  evaporating  denser regions not hosting ionizing sources where recombinations are important.  This  reversal to a more outside-in progression in the late stages of reionization  is more pronounced if the  emission of ionizing radiation is  restricted to massive  highly-clustered and rare sources.  \\item  If the emission of ionizing radiation is restricted to rare sources   reionization proceeds more quickly and  the sizes  of coherently ionized regions are  significantly larger.  The latter results in an about factor two or more larger  mean free path for ionizing photons.  \\item Like other studies we find that the amplitude of the 21cm power spectrum and its  evolution in the later stages of reionization is mainly sensitive to the space density of ionizing  sources.   The sensitivity to the space density of ionizing sources is, however,  significantly increased if a realistic spatially  inhomogeneous distribution of sinks of ionizing radiation is taken  into           account. The slope of the power  spectrum is very sensitive to the spatial distribution  of sinks of ionizing radiation.  \\item  Measurements of the amplitude and slope of the 21cm power spectrum at scales  corresponding  to    $k \\sim  0.3 - 0.5 h{\\rm Mpc}^{-1}$  with the upcoming low-frequency   instruments LOFAR and MWA have excellent prospects to reveal important information on the spatial distribution of sinks and sources  of ionizing radiation and the speed of  reionization if the daunting tasks of accurate calibration  and foreground removal are mastered  successfully. The PDF of the  21cm brightness distribution contains important complimentary information. Measuring the PDF  will, however, unfortunately most likely require higher sensitivity than can be achieved with first generation 21cm experiments.  \\end{itemize}  Our modelling here has involved  a number of significant simplifications. The spatial  distribution of dark matter   modelled in the Zel'dovich approximation was used as an proxy for the spatial distribution  of the IGM. The ionizing emissivity  of sources and recombination in dense region was modelled  only in an approximate integrated fashion and  the dynamical effects of the ionization radiation  on the gas were neglected. Despite the large particle number used in the simulations resulting in a substantial dynamic range there were still clear deficiencies in modelling high-density regions and low-mass collapsed objects/mini-haloes. We nevertheless think that our simulations  have caught the essential properties of the topology of  the epoch of reionization. Our simulations suggest that the idea that reionization proceeds strictly inside-out from beginning to nearly to the end may need revision if reionization indeed occurs in a photon-starved regime as suggested by  the \\lya\\ forest data."
    },
    "0806/0806.3717_arXiv.txt": {
        "abstract": "We have performed nonlinear ideal magnetohydrodynamic simulations of the long term evolution of a magnetized low-density ``bubble\" plasma formed by a radio galaxy in a stratified cluster medium.  It is found that about 3.5\\% of the initial magnetic energy remains in the bubble after $\\sim 8 \\times 10^{9}$~years, and the initial magnetic bubble expansion is adiabatic.  The bubble can survive for at least $8 \\times 10^9$~years due to the stabilizing effect of the bubble magnetic field on Rayleigh-Taylor and Kelvin-Holmholtz instabilities, possibly accounting for ``ghost cavities\" as observed in Perseus-A\\@.  A filament structure spanning about 500~kpc is formed along the path of bubble motion. The mean value of the magnetic field inside this structure is $\\sim 0.88$~$\\mu$G at $\\sim8\\times10^9$~years. Finally, the initial bubble momentum and rotation have limited influence on the long term evolution of the bubble. ",
        "introduction": "An unsolved problem in active galactic nuclei (AGN) feedback on clusters is how to account for the the morphology and stability of buoyant bubbles and their interactions with the ambient intracluster medium (ICM) \\citep{mn07}. \\citet{fstacji06} showed that in the Perseus cluster, such bubbles can stay intact far from cluster centers where they were inflated by AGN jets. However, studies of kinetic-energy dominated jets in the purely hydrodynamic limit did not explain the observed long term persistence of the buoyant bubbles, which are prone to Rayleigh-Taylor (RTI) and Kelvin-Holmhotz (KHI) instabilities and fragment entirely within $100\\;\\unit{Myr}$, contrary to observations [but see \\citet{rmfs05,ps06,ga07}].  Appreciable magnetic energy has been observed in both cluster and radio lobe plasmas \\citep{oek00,kdl01,chh05}. The magnetic fields could play a vital role in the dynamics of the rising bubble, as shown by a series of 2D magnetohydrodynamic (MHD) studies for bubbles of several $10^8$~years \\citep{bk01,rdrr04,jd05}. \\citet{sg08a} showed that uniform strong magnetic fields do not suppress RTI completely, but sheared fields at the bubble interface quench the instability. This implies that not only the field strength but also the configuration matters for bubble stability. \\citet{rebhp07} did a comprehensive study of the influence of different magnetic field configurations upon the stability of a rising bubble. They found that the internal bubble helical magnetic field moderately stabilizes the bubble, however the bubble has an initial plasma parameter $\\beta\\equiv\\left<2nT/B^2\\right>\\gg1$,  still in thermal energy-dominated regime. Recently, \\citet{llf06} proposed a different bubble magnetic field configuration for jet/lobes. Here we adopt the field configuration of \\citet{llf06} and focus on the late stage of the evolution of magnetically dominated bubbles in 3D\\@. We show that this spheromak-like magnetic field configuration strongly stabilizes the instabilities and prevents the bubble from breaking up, possibly forming the intact but detached ``ghost cavity\" observed in systems such as Perseus-A\\@. This {\\it Letter} is organized as follows. In Sec.~\\ref{setup}, we describe the problem setup. The simulation results and discussions are given in Sec.~\\ref{results}. ",
        "conclusions": "\\label{results}   Energy and density evolution reveal the different stages of the bubble. Figure~\\ref{energy}(\\emph{top}) presents the time evolution of various  energies at the early stage, in which the gravitational energy $E_g=\\int \\rho\\psi dV$ and the internal energy $E_T=\\int p/(\\gamma-1)dV$, where $dV$ is the infinitesimal volume and the integral is over the entire computation domain, have been reduced by their initial values $E_{g,0}$ and $E_{T,0}$, respectively. The total energy is defined as $E_{\\rm total}=E_m+E_k+E_T+E_g$, where the  kinetic energy is defined as $E_k=\\int 1/2\\rho v^2 dV$ and the  magnetic energy is defined as $E_m=\\int B^2/2 dV$. At the early stage ($t\\lesssim5$), the kinetic energy increases since the bubble accelerates upward due to buoyancy, and the magnetic energy decreases due to work done in expansion from the weak initial Lorentz force and conversion to shock and wave energy \\citep{nll06}. The gravitational energy increases because there is a net outward mass flow in the axial direction. The passage of the shock wave heats and compresses the ICM and alters its pressure gradient.  The total energy $E_{\\rm total}$ is almost constant before $t\\sim10$, at which time the shock and wavefront reach the boundary. Conservation of the total energy for $t\\lesssim10$ is not strictly satisfied due to numerical diffusion.  The initial bubble expansion ($t\\le20$) is approximately adiabatic, which gives $E=Vp\\propto V^{1-\\gamma}$, where $V$ is the bubble volume. Therefore $E_2/E_1=(V_1/V_2)^{2/3}$, where the subscript $1$ and $2$ indicate the initial and final state, respectively. At $t=20$, the radius of the bubble has increased from $2$ to $\\sim5$, which gives $E_2/E_1\\sim16\\%$, roughly matching the amount of magnetic energy in the bubble [$20\\%$ from Fig.~\\ref{energy}(\\emph{top})].  For later times [Fig.~\\ref{energy}(\\emph{bottom})], fitting with $E_m(t)/E_{m,0}=\\exp(-t/\\tau_{\\rm dis})$, we have: (1)~a fast dissipation stage ($t\\lesssim10$) when magnetic energy dissipation time $\\tau_{\\rm dis}\\sim11$; (2)~a slow dissipation stage ($t\\gtrsim10$) when $\\tau_{\\rm dis}\\sim114$. Both times are much smaller than the numerical dissipation time $\\tau_{\\rm res}$ at the corresponding stages (see discussions at the end of this section). The kinetic energy is oscillating while it is decaying slowly, which can be understood as follows. Gravity pulls the bubble down to the denser ICM, compressing the bubble and causing an increase in magnetic energy (since the magnetic flux is nearly conserved).  This results in the magnetic energy oscillating in phase with the kinetic energy with period $\\sim90$ ($3.5\\times10^9\\;\\unit{yr}$). At $t=200$ after several periods of decaying oscillations, about 3.5\\% of the initial magnetic energy remains.  The magnetic field suppresses instabilities and therefore the bubble remains intact longer. Figure~\\ref{overview} presents the typical density distribution (logarithmic scale) in 2D $x$--$z$ slices at $y=0$ at different times ($t=0,7.5,20,50,100,125$). The white solid contour lines indicate contours of constant magnetic field strength $|B|$. Consistent with Fig.~\\ref{energy}(\\emph{bottom}), after $t\\gtrsim50$, the bubble undergoes a slowly decaying oscillation between $z\\sim19$ and $z\\sim24$. We have also performed simulations of an unmagnetized bubble and found that it disintegrates after $t=20$, whereas the magnetic bubble still clearly differentiates itself from the ambient medium at $t=200$ (Fig.~\\ref{overview}). The formation of an ``umbrella\" or a thin protective magnetic layer on the bubble working surface suppresses instabilities \\citep{rebhp07}.  The stronger magnetic field is also found to move the bubble faster and push the bubble farther away from its initial position.  The position of the bubble top as a function of time is shown in Fig.~\\ref{pos}. This position is calculated by plotting the axial profile of the density along the line of $(x,y)=(0,0)$ and finding the location of the first density jump from the top of the domain. Comparing the unmagnetized run (dash line) with the magnetized run (both without initial momentum/rotation), we can see that the rising speed increases from $~0.36$ to $0.6$.  Note that the rough estimate of the terminal speed $v_t$ based on the initial gravitational acceleration $g$ and size $r_d$, both position-dependent, is $v_t\\approx4/3\\sqrt{2gr_d}\\sim2$, assuming force balance between buoyancy and viscous drag force at the final stage \\citep{mn07}.  We investigate the effect of the magnetic field on bubble stability in more detail.  Following \\citet{nll07}, we first study KHI of the bubble. At both sides of the bubble, the magnetic field is almost uniform, not twisted like at the top of the bubble. From linear analysis, the instability criterion for nonaxisymmetric KHI surface modes is \\citep{hr02} : $\\Delta V > V_{As}=[(\\rho_b+\\rho_e)/(4\\pi \\rho_b\\rho_e)(B_b^2+B_e^2)]^{1/2}$, where $\\Delta V\\equiv|V_b-V_e|$ is the velocity shear and $V_{As}$ is the surface Alfv\\'en speed. The subscripts $b$ and $e$ indicate the bubble and external medium, respectively. The nonaxisymmetric body modes would be important if \\citep{hr99}: (1) $V_b>V_f$, in which $V_f$ is the fast magnetosonic speed; or (2) $C_sV_A/(C_s^2+V_A^2)^{1/2}<V_b<V_s$, in which $V_A$ is the Alfv\\'en speed and $V_s$ is the slow magnetosonic speed. The definitions of $V_{f,s}$ are: $V_{f,s}=\\big\\{1/2\\{C_s^2+V_A^2\\pm[(C_s^2+V_A^2)^2-4C_s^2V_A^2\\cos^2\\theta]^{1/2}\\}\\big\\}^{1/2}$. Since we focus on the $x$-direction only, $V_A^2\\cos^2\\theta$ are taken to be $B_x^2/\\rho$. Figure~\\ref{khi} (\\emph{top}) displays the transverse distribution of the bulk flow speed $V=(v_x^2+v_y^2+v_z^2)^{1/2}$, the density $\\rho$, and the magnetic field strength $B$ at $t=7.5$ at $z=8.0$. A distinct velocity shear is identified at $x\\sim3.5$. Across this shear, $\\Delta V\\sim0.75$ and $V_{As}\\sim0.85$. The inequality $\\Delta V<V_{As}$ holds. This means that the KHI surface modes are completely suppressed. Figure~\\ref{khi} (\\emph{bottom}) displays the transverse distribution of $V$, $V_f$, and $V_s$ at $t=7.5$ at $z=8.0$ . We see that the bulk flow $V$ lies between $V_f$ and $V_s$. The inequality $V_s<V<V_f$ is satisfied in the body of the buoyant bubble. This rules out the KHI body modes as well.  RTI could also be suppressed by the magnetic field. For the idealized case of two conducting fluids separated by a contact discontinuity with a uniform magnetic field parallel to the interface undergoing constant acceleration $g$, \\citet{chan61} demonstrated that RTI on a scale $L$ parallel to the field requires $B<B_c\\equiv[Lg(\\rho_h-\\rho_l)]^{1/2}$, in which $\\rho_h$ and $\\rho_l$ are the densities in the heavy and light fluids respectively. Modes perpendicular to the field are unaffected.  At the top of the bubble ($z\\sim11.6$) (Fig.~\\ref{rti}), $\\rho_h$ and $\\rho_l$ are found to be $0.15$ and $0.06$, respectively, and the gravitational acceleration $g$ is calculated to be $0.95$. If $L$ is chosen to be the computation domain size, the maximum possible mode wavelength in the simulations, then the critical magnetic field strength is $B_c\\sim1.6$, which is larger than the magnetic field  ($\\sim0.8$) at that location. This means that only part of the parallel modes are suppressed regardless of the perpendicular modes. This is not consistent with the simulation results. However, as pointed out in \\citet{sg08a}, the twisting nature of the bubble field at the top of the bubble introduces a current sheet at the surface, and the changes in the direction of the field at the interface must be on very small scales to inhibit the interchange modes. A more detailed study of this effect is beyond the scope of this paper and will be the subject of future study.  Our simulation also provides one possible explanation of the morphology and origin of the large scale magnetic field and the generation of ``ghost cavities\" observed in many clusters. Like \\citet{rdrr04}, a magnetized high-density tail remains as the bubble rises (Fig.~\\ref{overview}). The tails have an elongated morphology (Fig.~\\ref{overview}), resembling H$\\alpha$ filaments, which is found to indicate the history of the rising bubble.  Interestingly, the magnetic field also helps stabilize this ``filament\" as it is still visible at $t=200$.  \\citet{nb04} has argued that thermal conduction has to be strongly suppressed in the ICM; otherwise, such cold filaments would be rapidly evaporated. As in \\citet{rebhp07}, our results provide a possibility that thermal conduction may be locally weaker in the bubble wake, thus preventing or slowing down filament evaporation. The magnetic field is distributed between $z\\sim2.5$ and $z\\sim25$ at $t=200$ with peak value around $0.156$ ($\\sim1.8\\;\\unit{\\mu G}$) or mean value around $0.078$ ($\\sim0.88\\;\\unit{\\mu G}$), close to the estimates of wider cluster fields \\citep{ct02,tfa02}. This large-scale ($\\gtrsim500\\;\\unit{kpc}$) magnetic field structure also mimics the morphology of the second class (``Phoenix\") of ``radio relics\" elongated from the cluster center to the periphery observed in A 115 \\citep{gfg01}.   The simulation results mean that the magnetic bubble rising from $\\sim125\\;\\unit{kpc}$ with an initial magnetic energy of $2\\times10^{59}\\;\\unit{ergs}$, will spread magnetic fields between $62.5\\;\\unit{kpc}\\lesssim z\\lesssim600\\;\\unit{kpc}$ and keep a magnetized bubble between $475\\;\\unit{kpc}$ and $600\\;\\unit{kpc}$ for at least $\\sim8\\times10^{9}\\;\\unit{yr}$. Therefore one could reproduce the intact but detached ``ghost cavity\" observed in some systems such as Perseus-A.  Real radio lobes possess complex, jet-driven flows, different from the static bubble that we use for our initial conditions.  We study the effects of additional variations in the initial bubble by having: (1)~an initial uniform injection speed $v_{\\rm inj}$ or (2)~an initial uniform rotation $\\omega$. \\citet{ntll08} reported that the radio lobe gains a momentum of $v_{\\rm inj}\\sim0.91C_{s0}=1.174$, and the rotation speed $\\omega$ at the edge of the bubble reaches $\\sim1.5C_{s0}/r_d=0.9675$ after jet-driven inflation and interaction with the ambient ICM, although neither are not uniform inside the bubble. Here we take these two values as our initial values for $v_{\\rm inj}$ and $\\omega$ and idealize both to be uniform throughout the bubble.  Fig.~\\ref{pos} presents the time evolution of the position of the bubble top with initial uniform injection velocity $v_{\\rm inj}=1.174$ (long dash) and initial uniform rotation speed $\\omega=0.9675$ (dash dot). It is seen that these two variables only have some influence on the evolution of the bubble during the early stage. The bubble with an initial rotation has a reduced buoyancy resulting from rotation-driven instabilities that in turn lead to a reduced pressure differential between the bubble and the ICM. The evolution, and structure of the bubble in the later times are, however, essentially the same.  The effects of numerical diffusion are estimated as follows. If we fit the net toroidal magnetic flux $\\psi_t=\\int B_y dS$ (only positive $B_y$ is selected) as $\\psi_t(t)/\\psi_t(t=0)\\equiv\\exp(-t/\\tau_{\\rm res})$, the ``resistive\" dissipation time due to numerical diffusion is $\\tau_{\\rm res}\\sim 38$ for $t\\le10$, $\\tau_{\\rm res}=138$ for $10\\le t\\le50$, and $\\tau_{\\rm res}=507$ for $50\\le t \\le 100$. Therefore the numerical diffusion is not important on the time scales of interest $(t\\lesssim200)$. Increasing numerical resolution can further reduce the numerical diffusion.  The oldest bubbles known so far are a few hundred $\\unit{Myr}$ old \\citep{mn07}, which are much younger than the long-sustained ($\\sim$ a Hubble time) magnetized bubbles reported in this \\emph{Letter}. The simulations show that at later times $t\\gtrsim50$, the density and X-ray luminosity contrast between the bubble and ICM plasma at the periphery of the cluster is small, which makes the direct X-ray observation of the bubble faraway from the parent galaxy very difficult. But the radio bubble at the borders of the clusters may still be detectable through: (1) Faraday rotation measurement of the outskirts of the cluster if an outside synchrotron source exists; (2) or radio observation from re-energization of fossil radio bubble plasma by shocks or mergers. It is highly possible for the bubbles to pile up in the outer atmospheres of the clusters in deep Chandra observation, if those observations are technically viable.  It is conceivable that part of the initial magnetic energy of the bubble accelerates cosmic rays. Throughout the lifetime of a cluster, there could be multiple AGNs injecting jets/lobes into the ICM\\@.  Both the cosmic rays and the remaining magnetic fields (distributed over large scales) could provide the energy sources for phenomena such as radio relics and radio haloes that have been observed for a number of clusters \\citep{fgsbr08}."
    },
    "0806/0806.1238_arXiv.txt": {
        "abstract": "General Relativity (GR) is a phenomenologically successful theory that rests on firm foundations, but  has not been tested on cosmological scales. The advent of dark energy (and possibly even the requirement of cold dark matter), has increased the need for testing modifications to GR, as the inference of such otherwise undetected fluids, depends crucially on the theory of gravity. In this work I outline a general scheme for constructing consistent and covariant modifications to the Einstein equations. This framework is such that there is a clear connection between the modification and the underlying field content that produces it. I conclude by a simple metric based modification of the fluctuation equations for which the background is exact $\\Lambda$CDM  and present its impact on observations of the cosmic microwave background radiation. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0806/0806.1712_arXiv.txt": {
        "abstract": "Extragalactic HI astronomy is half century old. Its maturity dramatically increased in the 1970s, with the commissioning of new powerful facilities. Its contributions to Cosmology are important, from the observation of galaxy rotation curves that showed the presence of dark matter in galaxies, to the measurement of cosmological parameters and the mapping of the large scale structure of the Universe. The Arecibo telecope has played a key role in these developments. It is also currently engaged in a number of experiments that utilize its L--band feed array to map thousands of square degrees of the sky and obtain the most sensitive large--scale view of the low $z$ HI Universe. ",
        "introduction": "I have been asked by the organizers of this workshop to devote opening remarks to a brief review of achievements in extragalactic HI astronomy and how these set the stage for current efforts in the field, particularly the ALFALFA survey. Because of the judgement passed not long ago by a national review committee, a dark shadow has been cast on the future --- and very existence in the next decade -- of the most beautiful of radio observatories, which hosts this workshop today. Finding myself in disagreement with some of the conclusions of said review, I make no apology for the fact that my report will lean towards the body of discovery Arecibo has contributed in extragalactic spectroscopy, while in the last session of this workshop, I chance some predictions on the possible large--scale projects Arecibo may be able to attack in the next decade.  Baryons make up about 4.5\\% of the mass/energy budget of the Universe, and only 1/6 of its matter density. At $z=0$ the vast majority of baryons are thought to exist in the form of coronal and intergalactic gas, at temperatures $>10^{5}$ K; $\\Omega_{stars}$ is a tiny 0.0027 and $\\Omega_{cold~gas}$ an even smaller 0.0008, of which a bit over half is neutral Hydrogen~\\cite{omega}, the target of this workshop's observational concerns. This unimpressive budgetary datum could well prompt the question: why do we care about extragalactic radio spectroscopy? First, HI is easy to detect at 21 cm wavelength, most of the emission originates in optically thin regions and cold gas masses are reliably measured; the abundance of cold gas is a reliable indicator of star forming potential for an extragalactic system. Second, the distribution of HI, which can be found at larger galactocentric distances than other easily detectable components in a galaxy, makes it an excellent tracer of the large--scale dynamics of its host. Third, scaling relations of disks, such as that between luminosity and rotational width, make HI measurements good cosmological tools: for example in the measurement of $H_\\circ$, peculiar velocities, the convergence depth of the Universe and the local matter density field. Fourth, because of its presence at relatively large galactocentric distances, HI is more vulnerable than the stellar component to external influences and thus constitutes a good tracer of tidal interactions, mergers and other environmental effects. Fifth, it can be the dominant baryonic component in low mass galaxies and thus provide a reliable census of low mass systems in the galactic hierarchy.  The 21 cm line of neutral Hydrogen was first detected in 1951 from Milky Way clouds and in 1953 from the Magellanic Clouds. It took nearly two decades of pioneering work, largely by M.S. Roberts, to bring the number of galaxies detected in HI to exceed one hundred \\cite{roberts75}. During this period, the near totality of extragalactic HI observations were made with single dishes, at Green Bank, Nan\\c cay, Effelsberg and Owens Valley. In the 1970s, mature, new aperture synthesis telescopes at Westerbork and the VLA were commissioned, followed two decades later by the GMRT, the largest synthesis telescope currently operating at 21cm. In the early 1970s, the primary reflecting surface of the Arecibo telescope was upgraded, making operation at L--band possible. At the same time, lower $T_{sys}$ receivers and broad band spectrometers also became available. The sensitivity of the 305m telescope provided the complementary counterpart to the high angular resolution capabilities of the synthesis intruments. Sensitivity was unfortunately not matched initially at Arecibo by breadth of field of view. With a field of view of a single--beam of $\\sim 3.3'$, unbiased, sensitive, blind, large--scale surveys remained impractical. This situation changed with the implementation of focal plane arrays in large single dishes. The first effective use of that technology for HI spectroscopy took place in the 1990s at Parkes and it yielded the important HIPASS catalog \\cite{hipass} ~of about 5000 HI sources. ALFA, a seven--feed focal plane array, was commissioned at Arecibo in 2004, making possible full--fledged, wide solid angle survey operations, which started in early 2005. ",
        "conclusions": ""
    },
    "0806/0806.0378_arXiv.txt": {
        "abstract": "We perform cosmological hydrodynamic simulations to determine to what extent galaxies lose their gas due to photoheating from an ionizing background. We find that the characteristic mass at which haloes on average have lost half of their baryons is $M_{\\rm c}\\sim 6.5 \\times 10^9 \\,h^{-1} \\ {\\rm M_\\odot}$ at $z = 0$, which corresponds to a circular velocity of $25 \\ {\\rm km} \\ {\\rm s}^{-1}$. This is significantly lower than the filtering mass obtained by the linear theory, which is often used in semianalytical models of galaxy formation. We demonstrate it is the gas temperature at the virial radius which determines whether a halo can accrete gas. A simple model that follows the merger history of the dark matter progenitors, and where gas accretion is not allowed when this temperature is higher than the virial temperature of the halo, reproduces the results from the simulation remarkably well. This model can be applied to any reionization history, and is easy to incorporate in semianalytical models. ",
        "introduction": "A UV-background produced by quasars and stars acts to quench star formation in small galaxies by photoheating their gas, which gets too hot to be confined in their shallow potential wells. This was originally pointed out by \\citet{dzn67}, and investigated in the context of the cold dark matter (CDM) model by \\citet{cr86}. Depending on the spectral shape of the ionizing background, the typical virial temperature below which haloes lose their gas is around $10^4$~K. Photoheating associated with reionization therefore inhibits star formation in dwarf galaxies and affects the faint end of the galaxy luminosity function \\citep[e.g.][]{efs92}.\\\\  Since \\cite{rees86} argued that in dark matter haloes with circular velocities around 30 km s$^{-1}$ gas can be confined in a stable fashion, with radiative cooling balancing photoheating, many attempts to quantify the effects of an ionizing background on galaxy formation have been made using semianalytical calculations \\citep[e.g.][]{br92, efs92, sgb94, ngs99, bul00, ben02a, ben02b, som02}, spherically symmetric simulations \\citep{tw96, kit00}, and three-dimensional cosmological hydrodynamic simulations \\citep{qke96, ns97, whk97, gne00}. The effects of self-shielding have been investigated using three-dimensional radiative hydrodynamic simulations by \\citet{su04a, su04b}.\\\\  However there is still debate about the value of the characteristic mass, $M_{\\rm c}$, below which galaxies are strongly affected by photoionization. \\citet{gne00} argued that $M_{\\rm c}=M_{\\rm F}$, the {\\it filtering mass} that corresponds to the scale over which baryonic perturbations are smoothed in linear perturbation theory (see Appendix~\\ref{FILTERINGMASS}). \\citet{hoe06} used simulations to argue that $M_{\\rm c}\\ll M_{\\rm F}$, in particular at low redshift. They argued that $M_{\\rm c}$ follows from considering the equilibrium temperature, $T_{\\rm eq}$, between photoheating and radiative cooling at a characteristic overdensity of $\\Delta \\simeq 1000$; $M_{\\rm c}$ haloes have a virial temperature $T_{\\rm vir} \\simeq T_{\\rm eq}(\\Delta = 1000)$.\\\\  The relation $M_{\\rm c}=M_{\\rm F}$ proposed by \\citeauthor{gne00} is often used in semianalytic modelling for dwarf galaxy formation in order to describe the quenching of star formation \\citep{bul00, ben02a, ben02b, som02}. An important application of this relation is to estimate whether reionization can explain the apparent dearth of satellite galaxies in the Milky Way, as compared with the high abundance of dark matter subhaloes \\citep{moo99}. Interestingly, \\citet{som02} and \\citet{no06} argued that models underestimate the number of dwarf satellite galaxies in the Local Group if the effect is as strong as inferred from the filtering mass. \\\\  Given the uncertainty in $M_{\\rm c}$, we were motivated to perform high-resolution hydrodynamic simulations of galaxy formation in a $\\Lambda$CDM universe with a time-evolving UV-background, and to measure the evolution of the baryon fraction as function of halo mass. We compare our results to those of \\citeauthor{gne00}  and \\citeauthor{hoe06}, and formulate an intuitive and simple model that reproduces our simulations very well. The model can be incorporated easily into any semianalytic model of galaxy formation.\\\\  This paper is organised as follows. Details of our simulations are described in Section~2. Section~3 contains our numerical results which focus on the baryon fraction of the simulated haloes, and the characteristic mass at which the baryon fraction is on average half the cosmic mean. We present a simple physical model that describes the evolution of the baryon fraction of individual haloes and compare it to the simulations in Section~4. Discussion and a summary follow in Sections~5 and 6. ",
        "conclusions": "We have performed high-resolution cosmological hydrodynamical simulations with an imposed, spatially uniform but time dependent UV-background to investigate to what extent haloes lose baryons due to photoevaporation. Our numerical results are consistent with the simulations of void regions by \\citet{hoe06}, but inconsistent with those found earlier by \\cite{gne00}.  The effect of a photoionizing background on galaxy formation is much weaker than previously assumed in semianalytic models \\citep[e.g.][]{bul00, ben02a, ben02b, som02}  based on the filtering mass approach proposed by \\citet{gne00} (see Appendix~\\ref{FILTERINGMASS}).\\\\  We have presented a novel model which uses the merger tree of each halo, and very simple physical arguments about gas retention and accretion, to predict how much the baryonic mass is in haloes of given total mass. The model is in excellent agreement with the simulations. Gas accretion is modelled very accurately if one simply assumes that gas at the outskirts of the halo, at an overdensity of a third of the virial gas density, cannot accrete if it is hotter than the halo's virial temperature. A slightly more involved model does even better in predicting baryon fractions as function of halo mass and merger history, as compared to simulations. These models allow us to predict the baryon fraction of individual haloes and hence the characteristic mass $M_{\\rm c}(z)$, at which haloes lose half of their baryons, for any assumed reionization history, by computing the predicted baryon fraction as a function of mass and redshift. The predicted $M_{\\rm c}(z)$ evolution is in very good agreement with the full simulations.\\\\  Our simulations and models both neglect self-shielding \\citep{su04a}, and radiative transfer and non-equilibrium effects \\citep{ah99}. We have performed a numerical experiment in which we artificially increase the photoheating rate at higher density, to estimate how our results might be changed by these effects. The modified heating rate does not change the mean IGM temperature and hence also the filtering mass remains the same, yet $M_{\\rm c}$ determined from the simulations increases significantly. Our models correctly predict this increase, demonstrating that the key physical quantity regulating the baryon fraction in haloes is neither the mean temperature of the IGM favoured by \\cite{gne00}, nor temperature of gas of $\\Delta \\simeq 1000$ favoured by \\citeauthor{hoe06}, but the temperature of the gas hanging around haloes ($\\Delta \\simeq 1/3 \\Delta_{\\rm vir}$). \\\\  Our models are easy to implement in any semianalytic model of galaxy formation, provide one can compute the equilibrium temperature at $\\Delta = 1/3 \\Delta_{\\rm vir}$ as a function of redshift for the assumed UV-background."
    },
    "0806/0806.0187_arXiv.txt": {
        "abstract": "We present N abundances for 21 bright giants in the globular cluster NGC 6752 based on high-resolution UVES spectra of the 3360\\AA~NH lines. We confirm that the Str{\\\" o}mgren $c_1$ index traces the N abundance and find that the star-to-star N abundance variation is 1.95 dex, at the sample's luminosity. We find statistically significant correlations, but small amplitude variations, between the abundances of N and $\\alpha$-, Fe-peak, and $s$-process elements. Analyses using model atmospheres with appropriate N, O, Na, and Al abundances would strengthen, rather than mute, these correlations. If the small variations of heavy elements are real, then the synthesis of the N anomalies must take place in stars which also synthesize $\\alpha$-, Fe-peak, and $s$-process elements. These correlations offer support for contributions from both AGB and massive stars to the globular cluster abundance anomalies. ",
        "introduction": "\\label{sec:intro}  Several decades have passed since the first identifications of the star-to-star abundance variations of light elements in globular clusters \\citep{popper47,osborn71,nf79}. Hydrogen burning via the CNO, Ne-Na, and Mg-Al chains at high temperatures qualitatively accounts for the observed trends: the abundances of C and O are low when N is high, O and Na are anticorrelated as are Mg and Al \\citep{langer95,denissenkov98,karakas03}. While the amplitude of the variation may differ from cluster to cluster, the abundance patterns have been found in every well studied Galactic globular cluster \\citep{smith87,kraft94,gratton04}. Such abundance patterns have also been identified in stars in extragalactic globular clusters \\citep{letarte06,johnson06}. However, the origin of these abundance anomalies remains elusive.  Internal nucleosynthesis and mixing can account for C, N, and Li in giant stars whose abundances exhibit a dependence upon evolutionary status \\citep{ss91,grundahl02}. These abundance patterns, C and Li destruction along with N production, are also found in field stars as the evolve up the RGB \\citep{gratton00b}. However, the identification of C, N, O, Na, Mg, and Al abundance variations in unevolved cluster stars \\citep{briley96,gratton01,grundahl02,cohen05} demonstrates that an external pollution mechanism must be the dominant source since unevolved stars have insufficient internal temperatures to run the necessary nuclear reactions and lack a mechanism to mix the products to surface layers. Unlike C, N, and Li, abundance variations for the elements O to Al have rarely, if ever, been identified in field stars, indicating that the abundance anomalies are likely associated with some (presently unknown) property of the cluster environment. \\citet{fulbright07} found low O and high Na and Al abundances in two bulge giants. However, they suggest that these two stars may be members of the bulge globular cluster NGC 6522. Intermediate-mass asymptotic giant branch (AGB) stars \\citep{cottrell81,ventura05c} and massive stars \\citep{prantzos06,smith06,decressin06} are candidates for the external pollution scenario. However, neither scenario currently provides a satisfactory explanation.  Nitrogen abundance variations in globular cluster giants were inferred from early observations of CH and CN variations \\citep{zinn77,dacosta80}. The NH molecular lines arguably offer the best probe of N abundances, in contrast to analyses of the CN lines which require knowledge of the C and O abundances. However, few studies have measured N from NH as well as the abundances for additional elements in a large homogeneous sample. In particular, few studies have measured N from NH using high-resolution spectra for a large sample. In this paper, we perform such an analysis of bright giants in the globular cluster NGC 6752 to measure the amplitude of the N abundance variation and to identify correlations between N abundances and light element abundances as well as possible correlations between N abundances and the abundances of heavy elements. ",
        "conclusions": "\\label{sec:summary}  In this paper we present measurements of N abundances in 21 bright giant stars near the RGB bump of the globular cluster NGC 6752. The sample was chosen to span the full range of the Str{\\\" o}mgren $c_1$ index at this luminosity. The amplitude of the N abundance variation at the sample's luminosity is 1.95 dex. We confirm that the N abundances are correlated with the $c_1$ index. We find a linear relation between the N abundances and a new index $cy$ = $c_1$ $- (b-y)$. We apply this new relation to large numbers of RGB stars and find that the [N/Fe] distribution increases towards higher values. Broader implications are that all globular clusters show large star-to-star variations in their $cy$ indices (i.e., N abundances) at all evolutionary stages.  We find that the N abundances are correlated with the light elements O, Na, Mg, and Al, a feature that is seen in every well studied globular cluster. However, we find for the first time that the N abundances are also correlated with the abundances of Si and heavier elements. While such correlations are statistically significant, the amplitude of the variation is small ($<$ 0.2 dex variation in [X/Fe] as [N/Fe] varies by 2 dex). We attribute the detection to the large sample size, high quality data, and most importantly to the homogeneous and precise stellar parameters obtained from the \\citet{grundahl99} Str{\\\" o}mgren photometry.  Analysis using model atmospheres with the appropriate N, O, Na, and Al abundances gives very similar abundances to those derived using model atmospheres with scaled-solar compositions. Therefore, the correlations between N and heavier elements are unlikely to be due to using models with scaled-solar compositions. In fact, for the subset of elements investigated (Si, Y, and Zr), the correlations with N would be even more significant had we used models with appropriate compositions.  Of great interest would be to search for such correlations in an additional sample of stars at a different luminosity within this cluster and in other clusters. If correlations between light and heavy elements are again identified, they would offer support to both the AGB and massive star pollution scenarios for explaining the star-to-star abundance variation of light elements."
    },
    "0806/0806.3266_arXiv.txt": {
        "abstract": "Considering gravitino dark matter scenarios with a long-lived charged slepton, we show that collider measurements of the slepton mass and its lifetime can probe not only the gravitino mass but also the post-inflationary reheating temperature $\\TR$. In a model independent way, we derive upper limits on $\\TR$ and discuss them in light of the constraints from the primordial catalysis of $^6$Li through bound-state effects. In the collider-friendly region of slepton masses below 1~TeV, the obtained conservative estimate of the maximum reheating temperature is about $\\TR=3\\times 10^9\\,\\GeV$ for the limiting case of a small gluino--slepton mass splitting and about $\\TR=10^8\\,\\GeV$ for the case that is typical for universal soft supersymmetry breaking parameters at the scale of grand unification. We find that a determination of the gluino--slepton mass ratio at the Large Hadron Collider will test the possibility of $\\TR> 10^9\\,\\GeV$ and thereby the viability of thermal leptogenesis with hierarchical heavy right--handed Majorana neutrinos. ",
        "introduction": "Big bang nucleosynthesis (BBN) is a powerful tool to test physics beyond the Standard Model; cf.~\\cite{Dimopoulos:1988ue,Jedamzik:2004er,Kawasaki:2004yh,Kawasaki:2004qu,Kohri:2005wn,Jedamzik:2006xz,Kawasaki:2008qe} and references therein. Indeed, in supersymmetric (SUSY) theories, severe constraints appear owing to the existence of the gravitino $\\gravitino$ which is the gauge field of local SUSY transformations and whose mass is governed by the SUSY breaking scale. As the spin-3/2 superpartner of the graviton, the gravitino is an extremely weakly interacting particle with couplings suppressed by inverse powers of the (reduced) Planck scale $\\MPl=2.4\\times 10^{18}\\,\\GeV$~\\cite{Wess:1992cp}.  Accordingly, once produced in thermal scattering of particles in the hot primordial plasma~\\cite{Ellis:1984eq,Moroi:1993mb,Ellis:1995mr,Bolz:1998ek,Bolz:2000fu,Pradler:2006qh,Pradler:2006hh,Rychkov:2007uq}, unstable gravitinos with a mass $\\mgr \\lsim 5~\\TeV$ have long lifetimes, $\\tau_{\\gravitino} \\gsim 100~\\seconds$, and decay during or after BBN. Since the decay products affect the abundances of the primordial light elements, successful BBN predictions imply a bound on the reheating temperature after inflation $T_{\\Reheating}$ which governs the abundance of gravitinos before their decay~\\cite{Ellis:1984eq,Moroi:1993mb,Ellis:1995mr,Bolz:1998ek,Bolz:2000fu,Pradler:2006qh,Pradler:2006hh,Pradler:2007ne,Rychkov:2007uq}: $T_{\\Reheating}\\lsim 10^8\\,\\GeV$ for $\\mgr \\lsim 5~\\TeV$~\\cite{Kohri:2005wn,Kawasaki:2008qe}.  We consider SUSY extensions of the Standard Model in which the gravitino $\\gravitino$ is the lightest supersymmetric particle (LSP) and a charged slepton $\\slepton$---such as the lighter stau $\\stau$---the next-to-lightest supersymmetric particle (NLSP). Assuming R-parity conservation, the gravitino LSP is stable and a promising candidate for dark matter; cf.~\\cite{Moroi:1993mb,Bolz:1998ek,Asaka:2000zh,Bolz:2000fu,Ellis:2003dn,Feng:2004mt,Cerdeno:2005eu,Steffen:2006hw,Pradler:2006qh,Pradler:2006hh,Pradler:2007ne,Rychkov:2007uq,Steffen:2007sp} and references therein.  Because of the extremely weak interactions of the gravitino, the NLSP typically has a long lifetime before it decays into the gravitino. If these decays occur during or after BBN, the Standard Model particles emitted in addition to the gravitino can affect the abundances of the primordial light elements.  For the charged slepton NLSP case, the BBN constraints associated with hadronic/electromagnetic energy injection have been estimated~\\cite{Ellis:2003dn,Feng:2004mt,Cerdeno:2005eu,Steffen:2006hw}. Taking into account additional constraints from large-scale-structure formation, apparently viable gravitino dark matter scenarios had been identified (see, e.g., benchmark scenarios $A_{1,2}$--$C_{1,2}$ in Ref.~\\cite{Steffen:2006hw}) which are attractive for two reasons: (i)~$T_{\\Reheating} > 10^9\\,\\GeV$ is possible so that thermal leptogenesis with hierarchical right-handed neutrinos~\\cite{Fukugita:1986hr,Davidson:2002qv,Buchmuller:2004nz,Blanchet:2006be,Antusch:2006gy}% \\footnote{Note that flavor effects~\\cite{Nardi:2006fx,Abada:2006fw,Blanchet:2006be,Antusch:2006gy} do not change the lower bound $T_{\\Reheating} > 10^9\\,\\GeV$ required by successful thermal leptogenesis with hierarchical right-handed neutrinos~\\cite{Blanchet:2006be,Antusch:2006gy}. However, in the case of (nearly) mass-degenerate heavy right-handed Majorana neutrinos, resonant leptogenesis can explain the baryon asymmetry at smaller values of $T_{\\Reheating}$~\\cite{Flanz:1994yx,Covi:1996fm,Pilaftsis:1997jf,Anisimov:2005hr}. Another example for a framework in which the limit $T_{\\Reheating} > 10^9\\,\\GeV$ is relaxed is non-thermal leptogenesis; see, e.g., \\cite{HahnWoernle:2008pq} and references therein.} remains a viable explanation of the baryon asymmetry~\\cite{Bolz:1998ek,Fujii:2003nr,Cerdeno:2005eu,Steffen:2006hw,Pradler:2006qh} and (ii)~the gravitino mass can be close to the NLSP mass $\\mslepton$, i.e., $0.1\\,\\mslepton\\lesssim\\mgr<\\mslepton$, so that a kinematical $\\mgr$ determination appears viable~\\cite{Buchmuller:2004rq,Martyn:2006as,Hamaguchi:2006vu}. With a kinematically determined $\\mgr$, one would be able to measure the Planck scale $\\MPl$ at colliders~\\cite{Buchmuller:2004rq,Martyn:2006as,Hamaguchi:2006vu} and to test the viability of thermal leptogenesis in the laboratory~\\cite{Pradler:2006qh}. Indeed, an agreement of the $\\MPl$ value determined in collider experiments with the one inferred from Newton's constant~\\cite{Yao:2006px} $G_{\\rm N} = 6.709\\times 10^{-39}\\,\\GeV^{-2}$ would provide evidence for the existence of supergravity (SUGRA) in nature~\\cite{Buchmuller:2004rq}.  Already in the early paper~\\cite{Dimopoulos:1989hk} it had been realized that long-lived negatively charged massive particles $\\champ$ (such as long-lived $\\slepton^-$'s) can form primordial bound states, which can affect the abundances of the primordial light elements. It was however only recently~\\cite{Pospelov:2006sc} when it was realized that bound-state formation of $\\champ$ with $\\Hefour$ can lead to a substantial overproduction of primordial $\\Lisix$ via the catalyzed BBN (CBBN) reaction \\begin{equation} (\\Hefour\\champ)+\\deuterium \\rightarrow \\Lisix + \\champ \\, . \\label{Eq:CBBNof6Li} \\end{equation} With an $\\champ$ abundance that is typical for an electrically charged massive thermal relic~\\cite{Asaka:2000zh}, this reaction becomes so efficient that an $\\champ$ lifetime of $\\tau_{\\champ} \\gtrsim 5\\times 10^3~\\seconds$ is excluded by observationally inferred values of the primordial $\\Lisix$ abundance~\\cite{Pospelov:2006sc}.  In the considered gravitino LSP scenarios, this bound applies directly to the lifetime of the $\\slepton$ NLSP~\\cite{Pospelov:2006sc,Cyburt:2006uv,Steffen:2006wx,Pradler:2006hh,Hamaguchi:2007mp,Kawasaki:2007xb,Pradler:2007is,Pradler:2007ar,Steffen:2007sp,Kawasaki:2008qe}: $\\tau_{\\slepton}\\lesssim 5 \\times 10^3~\\seconds$.% \\footnote{While numerous other bound-state effects can affect the abundances of $^6$Li and the other primordial light elements~\\cite{Kohri:2006cn,Kaplinghat:2006qr,Cyburt:2006uv,Bird:2007ge,Jedamzik:2007cp,Jedamzik:2007qk}, the approximate $\\tau_{\\slepton}$ bound is found to be quite robust; cf.~\\cite{Pospelov:MPP-2008-63}.} This implies $\\mgr\\lesssim 0.1\\,\\mslepton$ in the collider-friendly mass range of $\\mslepton\\lesssim 1~\\TeV$. Accordingly, the region $0.1\\,\\mslepton\\lesssim\\mgr<\\mslepton$, in which the kinematical $\\mgr$ determination appears feasible, seems to be excluded by BBN constraints~\\cite{Steffen:2006wx}.  Moreover, within the framework of the constrained minimal supersymmetric Standard Model (CMSSM), we have found that the $\\Lisix$ constraint implies the upper limit $\\TR \\lesssim 10^7~\\GeV$ in gravitino dark matter scenarios with unbroken R-parity and typical thermal $\\stau$ NLSP relic abundances~\\cite{Pradler:2006hh,Pradler:2007is,Pradler:2007ar}. For a standard cosmological history, this finding clearly disfavors successful thermal leptogenesis within the CMSSM in the case of hierarchical heavy right-handed Majorana neutrinos.  For a gravitino LSP mass range that is natural for gravity-mediated SUSY breaking, the $\\Lisix$ constraint can even point to a CMSSM mass spectrum which will be difficult to probe at the Large Hadron Collider (LHC)~\\cite{Cyburt:2006uv,Pradler:2006hh,Pradler:2007is,Pradler:2007ar}.  This letter provides a model independent study of gravitino LSP scenarios with a charged slepton NLSP that has a $^6$Li-friendly lifetime of $\\tau_{\\slepton}\\lesssim 10^4~\\seconds$ and a collider-friendly mass of $\\mslepton\\lesssim 1~\\TeV$. In contrast to~\\cite{Pradler:2006hh,Pradler:2007is,Pradler:2007ar,Kersten:2007ab}, the investigation presented in this work is not restricted to a constrained framework such as the CMSSM or to the gravitino mass range suggested by gravity-mediated SUSY breaking. Thereby, generic results are obtained with a range of validity that includes models with gauge-mediated SUSY breaking and/or non-standard mass spectra.  The remainder of this letter is organized as follows. In the next section we review that $\\mgr$ could be determined by measuring $\\mslepton$ and $\\tau_{\\slepton}$ at colliders. We give also an upper limit on $\\mgr$ that depends on $\\mslepton$ and $\\tau_{\\slepton}$. This limit allows us to derive a lower limit for the gravitino density $\\OmegaGr$. Since $\\OmegaGr$ cannot exceed the dark matter density $\\OmegaDM$, this leads to con\\-servative $\\TR$ limits which can be probed in measurements of $\\mslepton$ and $\\tau_{\\slepton}$ (and the gluino mass $m_{\\gluino}$). We use the derived expressions to translate the $\\tau_{\\slepton}$ constraint from CBBN of $^6$Li into robust upper limits on $\\TR$ that depend on the mass ratio $m_{\\gluino}/\\mslepton$. Finally, we show that the requirement $\\TR> 10^9\\,\\GeV$ needed for successful standard thermal leptogenesis provides an upper limit on this ratio and thereby a testable prediction for LHC phenomenology. ",
        "conclusions": "The observation of a (quasi-) stable heavy charged slepton $\\slepton$ as the lightest Standard Model superpartner at future colliders could provide a first hint towards the gravitino LSP being the fundamental constituent of dark matter. Under the assumption that the gravitino is the LSP and a long-lived $\\slepton$ the NLSP, we have shown that measurements of the slepton mass $m_{\\slepton}$ and its lifetime $\\tau_{\\slepton}$ can provide the gravitino mass $\\mgr$ and upper limits on the reheating temperature $\\TR$ after inflation, which can be tightened once the gluino mass $m_{\\gluino}$ is measured.  The conceivable insights into $\\mgr$ are a direct consequence of the SUGRA Lagrangian and of R-parity conservation. Thus, they will be valid for any cosmological scenario provided the $\\gravitino$ is the LSP and a $\\slepton$ the NLSP. The possibility to probe $\\TR$ at colliders results from the $\\TR$ dependence of the thermally produced gravitino density and the fact that this density cannot exceed the dark matter density. Thereby, the given upper limits on $\\TR$ rely not only on the SUGRA Lagrangian and R-parity conservation but also on the assumptions of a standard cosmological history and on (gauge) couplings that evolve at high temperatures as expected from their standard MSSM renormalization group running. Nevertheless, under these assumptions, the presented $\\TR$ limits are very robust since they have been derived in a model independent way from a conservative lower limit on the (thermally produced) gravitino density.  Prior to collider measurements of $m_{\\slepton}$, $\\tau_{\\slepton}$, and $m_{\\gluino}$, primordial nucleosynthesis can be used to constrain $\\mgr$ and $\\TR$. In particular, we have found that the constraint $\\tau_{\\slepton}\\lesssim 5\\times 10^3\\,\\seconds$ from CBBN of $^6$Li implies a truly conservative limit of $\\TR \\lesssim 3\\times 10^9\\,\\GeV$ in the collider-friendly region of $m_{\\slepton}\\lesssim 1\\,\\TeV$ and in the limiting case of a small gluino--slepton mass splitting at the weak scale: $c=m_{\\gluino}/\\mslepton\\simeq 1$. Indeed, with $c=7$---as obtained in scenarios with universal soft SUSY breaking parameters at the scale of grand unification---our robust conservative upper limit is $\\TR \\lesssim 10^8\\,\\GeV$ and thus as severe as in scenarios with an unstable gravitino of mass $\\mgr \\lesssim 5~\\TeV$.  The $\\tau_{\\slepton}$ constraint from CBBN of $^6$Li has allowed us to derive upper limits on $\\TR$ that depend on $\\mslepton$ and $c$ only. In particular, we find that the condition for successful thermal leptogenesis with hierarchical right-handed heavy Majorana neutrinos, $\\TR> 10^9\\,\\GeV$ ($3\\times 10^9\\,\\GeV$), implies an upper limit of $c<3$ ($1.5$) for $\\mslepton<1~\\TeV$. This is a prediction of thermal leptogenesis that will be testable at the upcoming LHC experiments.  \\medskip"
    },
    "0806/0806.4393_arXiv.txt": {
        "abstract": "{} {I investigate the shape of the electron cosmic ray spectrum in the range up to $\\sim 1000$~keV, assuming that the acceleration process at the shock results in a power law in momentum, and that downstream of the shock the spectrum is affected by Coulomb interactions with background electrons only. }{ In the non-relativistic regime one can analytically determine the energy of an electron starting with a certain energy, and use this result to produce an electron cosmic ray spectrum, modified by Coulomb losses. }{ An analytic expression for the electron spectrum is obtained that depends on the parameter $n_{\\rm e}t$, which can be estimated from a similar parameter used to characterize the line spectra of supernova remnants. }{ For the brightest supernova remnants $n_{\\rm e}t > 10^{11}$~cm$^{-3}$s, and most of the electrons accelerated to $< 100$~keV have lost their energy. Because of its high radio flux, Cas A is the most likely candidate for non-thermal bremsstrahlung. Although it has $n_{\\rm e}t \\sim 2\\times 10^{11}$~cm$^{-3}$s, one may expect to pick up non-thermal bremsstrahlung above 100~keV with current hard X-ray detectors. } ",
        "introduction": "Over the last five years there has been considerable progress in our understanding of high energy cosmic ray acceleration by supernova remnants (SNRs). The spatial distribution and spectral shape of X-ray synchrotron emission from all historical SNRs and some other SNRs has made it clear that electron can be accelerated to energies up to $\\sim 100$~TeV \\citep{koyama95}, that diffusive shock acceleration works very efficiently, with diffusion close to the Bohm limit \\citep{vink05b,vink06d,parizot06,stage06}, and that magnetic fields are amplified by cosmic ray streaming \\citep{vink03a,berezhko03c,bamba05,voelk05}. Moreover, the morphology of Tycho's SNR shows that at least in this remnant, but also probably in Kepler's SNR, cosmic ray acceleration is so efficient that it gives rise to enhanced compression ratios \\citep{warren05,ellison04}. In addition, considerable progress has been made in the field of TeV astronomy, with many SNRs being established as sources of TeV \\grays\\ \\citep[e.g.][]{aharonian04,aharonian05,albert07}. However, it has not yet been firmly established whether the \\gray\\ emission is dominated by inverse Compton scattering from TeV electrons, or from neutral pion decay, caused by TeV ion cosmic rays.  In contrast,% little has happened concerning observations of low energy cosmic rays. This is unfortunate, since from a theoretical perspective, the initial stages of the acceleration process, from thermal energies up to energies where Fermi acceleration efficiently operates, is complex and not well understood \\citep[see e.g.][for a review]{malkov01}. In fact, the processes by which particles are injected into the Fermi acceleration process, are likely to be different for electrons and ions, unlike Fermi acceleration itself. Our current understanding of the initial stages of acceleration is largely based on either in situ observations of inter-planetary shocks, and on computer modeling using hybrid or particle in cell codes \\citep[e.g.][]{bykov99,schmitz02,lee04}. There was some hope that hard X-ray observations of SNRs might provide  observational information on at least the electron component of low energy cosmic rays \\citep{asvarov90,favata97,vink97,vink03a}. However, it is generally agreed that the non-thermal hard X-ray emission from SNRs such as Cas A \\citep{the96,allen97,favata97,vink03a,renaud06}, SN1006 \\citep{allen99,kalemci06b}, Tycho \\citep{allen99}, Kepler and RCW 86 \\citep{allen99} is due to X-ray synchrotron from the highest energy electron cosmic rays, rather than from bremsstrahlung from low energy electrons. Nevertheless, the issues remains of interest as current and future hard X-ray detectors push the detection of hard X-ray emission to higher photon energies. At high enough energies, the steepening X-ray synchrotron spectrum is likely to be overtaken by non-thermal bremsstrahlung.  In this paper I discuss the shape of the low energy cosmic ray electron spectra,  given the importance of Coulomb-losses. These Coulomb losses in general alter the shape of the low electron cosmic ray spectrum as generated near the shock front. Note that a similar situation was investigated by \\citet{sarazin99} for higher energy electrons in clusters of galaxies.  \\begin{figure} \\centerline{ \\includegraphics[angle=-90,width=0.45\\textwidth]{electron_trajectories.ps} } \\caption{ The electron energy as a function of $n_{\\rm e} t$ ($\\lambda_{\\rm ee} = 30.9$) for different values of energy at $t=0$~s (50~keV, 100~keV, 200~keV...6400~keV). The solid lines show the exact, but implicit solutions to Eq.~\\ref{eq-losses}, i.e. Eq.~\\ref{eq-implicit}. The dotted lines are the non-relativistic solutions to Eq.~\\ref{eq-losses}, i.e. Eq.~\\ref{eq-e_t}. \\label{fig-eltrajectories} } \\end{figure} ",
        "conclusions": "I have derived an analytic expression for the low energy electron cosmic ray spectrum, assuming that the cosmic ray spectrum, as produced by the diffusive acceleration process, is a power law in momentum, and is only affected by Coulomb losses downstream of the shock. As may be expected, the parameter governing the Coulomb losses depends on $n_{\\rm e}t$, the product of (background) electron density and the age of the accelerated electron population. This parameter is similar to the ionization time scale of the hot, shock heated plasma, and can be easily estimated from the line spectra of SNRs.  For the SNRs with high densities, and large $n_{\\rm e}t$, one does not expect a large population of accelerated electrons below 100~keV. This is in particular the case for the very bright radio source Cas A. Nevertheless, Cas A remains a good candidate for searching for non-thermal bremsstrahlung, because of its high ion density and because its radio brightness indicates a large amount of accelerated electrons. Based on reasonable, but somewhat uncertain, assumptions, the results presented here imply that only above $\\sim 100$~keV one expects to pick up the non-thermal bremsstrahlung component from Cas A. This may be accomplished with future deep Suzaku HXD observations, or additional deep INTEGRAL observations."
    },
    "0806/0806.1055_arXiv.txt": {
        "abstract": "The cross-correlation between high redshift galaxies and 21 cm emission from the high redshift intergalactic medium (IGM) promises to be an excellent probe of the Epoch of Reionization (EoR). On large scales, the 21 cm and galaxy fields are anti-correlated during most of the reionization epoch.  However, on scales smaller than the size of the H II regions around detectable galaxies, the two fields become roughly uncorrelated. Consequently, the 21 cm-galaxy cross power spectrum provides a tracer of bubble growth during reionization, with the signal turning over on progressively larger scales as reionization proceeds. The precise turnover scale depends on the minimum host mass of the detectable galaxies, and the galaxy selection technique. Measuring the turnover scale as a function of galaxy luminosity constrains the characteristic bubble size around galaxies of different luminosities. The cross spectrum becomes positive on small scales if ionizing photons fail to escape from low mass galaxies, and these galaxies are detectable longward of the hydrogen ionization edge, because in this case some identifiable galaxies lie outside of ionized regions. LOFAR can potentially measure the 21 cm-galaxy cross spectrum in conjunction with mild extensions to the existing Subaru survey for $z=6.6$ Lyman-alpha emitters, while the MWA is slightly less sensitive for detecting the cross spectrum. A futuristic galaxy survey covering a sizable fraction of the MWA field of view ($\\sim 800$ deg$^2$) can probe the scale dependence of the cross spectrum, constraining the filling factor of H II regions at different redshifts during reionization, and providing other valuable constraints on reionization models. ",
        "introduction": "\\label{sec:intro}  Detecting 21 cm emission from the high redshift intergalactic medium (IGM) will potentially revolutionize our understanding of cosmic reionization. The 21 cm signal promises direct, three-dimensional information regarding the state of the IGM during reionization (e.g. Scott \\& Rees 1990, Madau et al. 1997, Furlanetto et al. 2006a). Unfortunately, experimental challenges are substantial. In particular, astrophysical foregrounds are expected to be four to five orders of magnitude larger than the signal from the high redshift IGM. However, {\\it known} foreground contaminants are spectrally smooth and should be distinguishable from the high redshift 21 cm signal itself (Zaldarriaga et al. 2004).  Given that we anticipate observational complications, it is important to develop diagnostics to confirm that the detected 21 cm signal indeed originates from the high redshift IGM. One such approach is to measure the cross correlation between 21 cm emission and a high redshift galaxy survey (Furlanetto \\& Lidz 2007, Wyithe \\& Loeb 2007). Since most of the anticipated foregrounds come from low redshift -- primarily galactic synchrotron -- and not from high redshift galaxies, the mean 21 cm-galaxy cross power spectrum signal is largely immune to foreground contamination (Furlanetto \\& Lidz 2007). Detecting a 21 cm-galaxy cross correlation should hence confirm that the detected 21 cm signal comes from the high redshift IGM. Moreover, continuing efforts are pushing galaxy surveys towards higher redshifts, and it is natural to consider the information that may be gleaned from combining galaxy and 21 cm surveys. Detecting galaxies at very high redshift is extremely challenging (e.g. Stark et al. 2007, Bouwens et al. 2008), but we will show here that a cross spectrum detection may already be possible with modest extensions to the Subaru survey (Kashikawa et al. 2006; see also Wyithe \\& Loeb 2007, Furlanetto \\& Lidz 2007).  In addition to these important practical advantages, the 21 cm-galaxy cross correlation is potentially sensitive to the size and filling factor of H II regions, the clumpiness of the IGM, and the nature of the ionizing sources. The 21 cm-galaxy cross correlation will also provide a more direct tracer of the interplay between the reionizing sources and the surrounding IGM, than the 21 cm auto power spectrum. The 21 cm-galaxy cross correlation should hence provide a unique and powerful probe of the Epoch of Reionization (EoR) and early structure formation. Here we follow up on earlier work by Wyithe \\& Loeb (2007) and Furlanetto \\& Lidz (2007), and focus on modeling the 21 cm-galaxy cross power spectrum, and exploring the insights that future surveys will provide regarding the EoR.  The outline of this paper is as follows. In \\S \\ref{sec:cross} we establish notation, describe the reionization simulations used in our analysis, and examine the basic simulated signal. We then characterize (\\S \\ref{sec:evol_sig}) the dependence of the cross spectrum on redshift and ionization fraction. In \\S \\ref{sec:sources} we illustrate how the signal is sensitive to the properties of the ionizing sources. In \\S \\ref{sec:recombs} we describe its variation with the abundance of Lyman-limit systems. We then examine the signal's dependence on the way in which high redshift galaxies are selected (\\S \\ref{sec:selection}), contrasting the results for Ly-$\\alpha$ selected galaxies with galaxies selected through other techniques. In \\S \\ref{sec:detectability} we briefly discuss the statistical power of future surveys to constrain reionization through measurements of the 21 cm-galaxy cross power spectrum. Finally, we summarize our main results and conclude in \\S \\ref{sec:conclusions}.  Throughout we consider a $\\Lambda$CDM cosmology parameterized by: $n_s = 1$, $\\sigma_8 = 0.8$, $\\Omega_m = 0.27$, $\\Omega_\\Lambda = 0.73$, $\\Omega_b = 0.046$, and $h=0.7$, (all symbols have their usual meanings), consistent with the WMAP constraints from Spergel et al. (2007) and Komatsu et al. (2008). ",
        "conclusions": "\\label{sec:conclusions}  In this paper, we considered the scientific return of future 21 cm-galaxy cross power spectrum measurements. A strong cross spectrum measurement ultimately requires detecting a sizable number of high redshift galaxies over a large field of view, which presents a significant observational challenge. Nonetheless, we showed that a detection of the cross spectrum may be achieved in the near future by combining LOFAR and the Subaru survey for LAEs at $z \\sim 6.6$. We estimate that a $\\sim 3-\\sigma$ detection is feasible, provided the IGM is $\\gtrsim 20\\%$ neutral at this redshift, and that the Subaru survey's sky coverage is extended from $0.25$ deg$^2$ to $\\sim 2$ deg$^2$. This detection would already be quite valuable, as it would help confirm that the detected 21 cm signal comes from the high redshift IGM, and not from foreground contamination, which should mostly be uncorrelated with high redshift galaxies (Furlanetto \\& Lidz 2007).  Futuristic galaxy surveys covering $100$s of square degrees on the sky, can be combined with the MWA, LOFAR, and other 21 cm surveys, to move beyond a mere detection of the cross spectrum signal and map out its detailed scale dependence. The galaxy surveys required for these measurements are clearly very challenging, but rapid progress is being made in this direction as deep, widefield surveys are being designed to study baryonic acoustic oscillations and/or weak-lensing at high redshift (e.g., ADEPT, HETDEX\\footnote{http://www.as.utexas.edu/hetdex/}, CIP, and others). Another option is to sparsely sample the MWA or LOFAR fields, in order to capture the large-scale modes (Furlanetto \\& Lidz 2007). We have shown that the 21 cm-galaxy cross spectrum is a relatively direct tracer of bubble growth during reionization. Measuring the turnover scale as a function of galaxy luminosity constrains the luminosity dependence of the characteristic bubble size. This information is difficult, or impossible, to obtain with the 21 cm auto spectrum alone. In order to extract the most information out of the cross spectrum, it should be combined with measurements of the galaxy auto spectrum and luminosity function, which will help to constrain the galaxy luminosity-halo mass correlation.  A further interesting feature of the simulated signal is that the cross-correlation changes sign on large scales near the beginning of reionization (Figure \\ref{fig:cross_v_z}). At this early phase of reionization, our results may, however, be modified by spin temperature fluctuations, which we presently neglect. Future work should incorporate these fluctuations. If our signature holds up, the change in sign of the cross correlation would provide a very interesting observational indicator of the earliest phases of reionization. Finally, we found that the 21 cm-galaxy cross power spectrum might provide an interesting observational signature of scenarios where ionizing photons fail to escape from low mass halos. Provided galaxies in these low mass halos are detectable longward of the ionization edge, we expect the cross spectrum to change sign and turn positive on small scales. Generally speaking, the 21 cm-galaxy cross spectrum is a more direct tracer of the impact of galaxies on the surrounding IGM than the 21 cm auto spectrum. As such, it can potentially provide a wealth of information about the EoR and early structure formation."
    },
    "0806/0806.1113_arXiv.txt": {
        "abstract": "We perform a linear stability analysis for corrugations of a Newtonian shock, with particle pressure included, for an arbitrary diffusion coefficient. We study first the dispersion relation for homogeneous media, showing that, besides the conventional pressure waves and entropy/vorticity disturbances, two new perturbation modes exist, dominated by the particles' pressure and damped by diffusion. We show that, due to particle diffusion into the upstream region, the fluid will be perturbed also upstream: we treat these perturbation in the short wavelength (WKBJ) regime. We then show how to construct a corrugational mode for the shock itself, one, that is, where the shock executes free oscillations (possibly damped or growing) and sheds perturbations away from itself: this global mode requires the new modes. Then, using the perturbed Rankine--Hugoniot conditions, we show that this leads to the determination of the corrugational eigenfrequency. We solve numerically the equations for the eigenfrequency in the WKBJ regime for the models of \\citeauthor{amatoblasi2005}, showing that they are stable. We then discuss the differences between our treatment and previous work. ",
        "introduction": "\\label{sec:intro}  One of the major uncertainties still surrounding particle acceleration around shocks is the level at which this saturates: \\ie, the total energy content in non-thermal particles, per unit volume, $\\cal{R}$, in units of the incoming fluid kinetic energy density. This saturation level plays a very important role, for instance, in discussions of the origin of cosmic rays as observed at Earth: it has to be rather large ($\\ga 0.1$) to allow SNe to provide the observed flux. Also, discussions on the origin of UHECRs are influenced by similar considerations: the idea that UHECRs are accelerated by GRBs \\citep{vietri1995,waxman1995} draws some measure of support by the coincidence between the energy release rate of GRBs in $\\gamma$-ray photons and that required to account for observations at Earth \\citep{waxman2002,vietrial2003}.  It is certainly possible that this level of saturation is somehow connected with the very uncertain particle injection mechanism, and that saturation occurs, in many astrophysically important situations, at very low levels. Yet, given the arguments concerning SNe and GRBs mentioned above, it appears that at least occasionally saturation must occur at rather large levels. This is especially so in the case of Galactic cosmic rays: if in fact we were to discard SNe as sources of cosmic rays, all other possible classes of sources, being both less numerous and less energetic, would force us to require saturation at even larger levels\\footnote{The situation is much less clear for UHECRs, because many of the proposed classes of sources are purely hypothetical, so that there are no constraints on their properties.}.  A distinct possibility is that the saturation level is determined not by the injection mechanisms, but by modifications which the particles' pressure induces in the shock properties. After all, in the test particle limit the particles' spectrum is ultraviolet divergent (though marginally so, of course), and one may hope that, since the particles' back-reaction on the fluid makes acceleration less efficient (by reducing the velocity jump around the gas sub-shock), a convergent spectrum will be obtained, with a definite value for the parameter ${\\cal R}$. After \\citeauthor{malkov1997}'s seminal papers \\citep{malkov1997,malkoval2000,malkovdrury2001}, fully self-consistent solutions with particles' pressure properly included have been obtained by \\citet{amatoblasi2005}. In these models, it appears that particles can carry away an arbitrarily large fraction of the incoming energy flux, for sufficiently large Mach number at upstream infinity. Even more intriguing is the fact that these solutions have particles' spectra which are more, not less, ultraviolet divergent than those in the test particle limit. In fact, this divergence is arrested only by arbitrarily limiting the largest individual energy to some fiducial value \\citep{amatoblasi2005}.   These solutions then seem to beckon for a stability analysis, in the hope that they are found to be unstable once ${\\cal R}$ exceeds some critical value. The instability we envision is of course the corrugational instability for shocks, whereby ripples on the shock surface become of larger and larger amplitude\\footnote{In all of this paper, the term corrugation instability is taken to include also the spontaneous emission of sound waves by the shock.}. If this instability were to exist, we might easily imagine that the shock is substituted by a region of strong fluid turbulence, where particles moving a few Larmor radii perceive only a small velocity difference at the two ends of their free wandering, and thus acceleration to high energies is made somewhat less efficient.  However, it is well-known \\citep{landaulifshitz1987} that shocks in polytropic fluids are stable against this kind of perturbations, and against the spontaneous emission of sound waves as well. The hope for the existence of an instability is based on a rather more subtle argument. When the shock surface flaps, it sheds in the downstream region pressure waves as well as entropy and vorticity perturbations. We will show later that the last two do not couple to particle perturbations, but pressure waves do, thus generating small perturbations in the particles' distribution function, $\\delta\\!f$. These particles will however return upstream by means of diffusion, and here generate, by their sheer perturbed pressure, some more fluid perturbations. Thus, not all energy shed by the shock is lost forever toward upstream infinity, as is the case in the purely hydrodynamical case, but some fraction of it returns to the shock to generate more flapping. In fact, since pressure waves are strongly damped by diffusion, and nearly all particles return to the upstream section since the shock is Newtonian, we may conjecture that most energy shed by the shock flapping makes it back to feed more flapping, even though account must be taken of diffusion and phase mixing. Is it possible that this sets up some strong reinforcement, making the whole system unstable? This is the question we address in this paper.  This question has of course been studied before, under somewhat different conditions.  The diffusion coefficient of non-thermal particles by the fluid ($\\equiv D$) has been taken (by other authors) as independent of the particle impulse: this is conventionally referred to as the two-fluid approximation, because the Boltzmann equation for the non-thermal particles can then easily be recast into an equation for their pressure, thusly erasing all references to the underlying distribution function. Besides making the two-fluid approximation, \\citet{monddrury1998} also neglected diffusion altogether. A more complex analysis has been presented by \\citet{toptygin1999}, who included energy transport and particle injection at the shock, but still in the two-fluid approximation. In the following, we shall make no such approximations: we will consider a finite diffusion coefficient $D$, and it will be allowed to be an arbitrary function of both $p$, the particle momentum, and of $\\rho$, the local fluid density, so as to at least mimic the increase in magnetic field strength due to flux freezing. The arbitrary nature of the dependence of $D$ on $p$ automatically prevents the use of the two-fluid approximation, and forces us to use the full Boltzmann-Skilling equation.  A word of caution is in order about some assumptions. We will neglect all energy in the form of magnetic field and Alfv\\'en waves; of course this is necessary because the zero-th order solutions of Blasi and collaborators do not include these effects, but in our case this neglect requires one extra assumption, {\\it i.e.} that the time scales for energy to accumulate into any of these energy sinks, $T_{in}$, and to flow out of each of them, $T_{out}$, be ordered like this: $T_{out} \\lesssim T_{in}$. If this inequality were severely violated, the magnetic field or Alfv\\'en waves might acquire a significant fraction of the total energy, despite their negligibility in the zero-th order solution, and make our treatment completely irrelevant.  {}  In the following, we shall follow closely the treatment of the shock corrugational instability given by \\citet{landaulifshitz1987}. In particular, we shall consider a steady-state shock in its own frame, located at $z=0$, with fluid coming from the left, and exiting from the right, so that all speeds are $>0$. Exactly like \\citeauthor{landaulifshitz1987}, we shall consider perturbations generated by the shock flapping, so that there can be no incoming waves, from either upstream or downstream infinity. In Section \\ref{sec:homog}, we shall consider perturbations in a homogeneous medium; this is perhaps a tad boring, but it contains a number of new results which are of the utmost importance later on. In Section \\ref{sec:pertup} we briefly discuss perturbations upstream, {\\it i.e.} where the flow is inhomogeneous; we show here that we can easily obtain the perturbations in the WKBJ limit $k_y\\rightarrow \\infty$, which restricts our analysis to the regime $\\lambda \\lesssim L$, where $\\lambda$ is the perturbation wavelength perpendicular to the shock, and $L$ is the typical size of the region of inhomogeneity upstream. In Section \\ref{sec:fitboundary} we discuss the boundary conditions on the perturbed particle distribution function, and derive how it relates to the amplitude of the modes to which particle perturbations couple. We present in Section \\ref{sec:fluifcond} the perturbed Rankine--Hugoniot conditions and in Section \\ref{sec:eigenf} what fixes the global corrugation mode eigenfrequency. In Section \\ref{sec:results} we present our numerical computations for the stability of the exact solutions of the zero-th order problem by \\cite{amatoblasi2005}. In Section \\ref{sec:discussion}, we will compare our results with other works in the literature, and briefly summarize our work. ",
        "conclusions": "\\label{sec:discussion}  In essence, our method is exact, except for the short-wavelength (WKBJ) approximation necessary to treat analytically perturbations in the inhomogeneous upstream precursor. One may however wonder where our method differs from previous work \\citep{monddrury1998,toptygin1999}, which has reported the existence of corrugational instabilities.  \\citet{monddrury1998} have reported the existence of both genuine corrugational instabilities, and of spontaneous emission of sound waves, for some (not all) of their models. There is of course a number of differences between this paper and theirs: we do not use the two-fluid approximations, and are interested in small-wavelength perturbations, contrary to them; we also notice the existence of perturbations in the upstream fluid, which they do not discuss.  \\citet{toptygin1999} included a number of novelties in his treatment, but he too did not notice that particles would diffuse upstream, so that he neglects the third and the d-modes altogether. Despite this, he does have perturbations upstream (pressure waves!), because he remarks that, for sufficiently long-wavelength perturbations, these become supersonic with respect to the fluid alone. This occurs because particles and fluid are tightly coupled  in long wavelength perturbations by diffusion, which traps non-thermal particles, so that the restoring pressure is the sum of particles' and fluid's pressures, which is larger than the pure fluid speed. This is of course correct, and exists in our computations as well.  Contrary to these authors, we have found that shocks with non--vanishing particle pressure are stable to corrugational instabilities, even in the region of parameter space where multiple solutions are possible. We believe that the reason for this discrepancy lies in our inclusion of diffusion, and the abandonement of the two-fluid approach, as we now argue. We stated in the Introduction that the only possibility for the shock destabilization (polytropic shocks without particles are well--known to be corrugationally stable) lies in setting up a loop whereby perturbations shed by the shock in the downstream region return to the upstream region via particle diffusion and excite more perturbations; thus perturbation energy is not lost to downstream infinity, but returns to the shock to create more havoc. However, the very same mechanism that brings particles back upstream, also causes perturbations' dampening: we have seen that diffusion leads to damping both pressure waves and the so--called third modes. Also, we have established that a part of the perturbation occurs in the $d$--mode (where $d$ stands for damping), where particles are perturbed in phase--space, but the total perturbation to the pressure vanishes. Thus, some fraction of the perturbation goes into a totally useless form (the $d$--mode), the new mode (the third mode) is always damped, and even pressure waves acquire a damping which is altogether neglected in the two--fluid approximation. In the end, while diffusion brings particles (and perturbations) back to the shock, the simultaneous damping is so strong to make the excitation of an instability ineffective.  We should remark that damping of pressure waves is weakest for the largest wavelengths, a limit which is unaccessible to us because of our WKBJ approximation, but is exactly the limit investigated by \\citet{monddrury1998}. While we have not found a large wavelength beyond which the shock becomes unstable, we cannot exclude that a proper treatment of the perturbations in the space--dependent precursor may yield a transition to the unstable regime.    A further comment is in order: of all possible geometries, the planar one is probably the least likely to display instability. Consider in fact a spherically symmetric explosion like a SuperNova. In this case, the perturbations (except of course for entropy and vorticity) generated downstream do not escape to infinity, as they do in the planar case, but return to the shock because the downstream region is finite and because they are deflected by a spatially-dependent refraction index; once they reach the shock,  they may generate further perturbations. The situation is even more promising when the shock is due to an accretion flow, or is stalling: in fact, except for the presence of particles, this is exactly the scenario proposed (see \\citealt{laming2007} for an analytic approach and discussion) for the generation of asymmetries in proto-neutron stars: in this case, the mechanism for the instability of a stalled accretion shock is the reflection by the hard star surface into outgoing pressure waves of advected entropy  perturbations, which return to the shock to generate more mischief. In the problem with particles, there is the extra complication due to diffusion, to be overcome to generate instability. Given the relative complexity of even a pure fluid analysis \\citep{laming2007}, it is likely that this problem will require a numerical approach."
    },
    "0806/0806.0924_arXiv.txt": {
        "abstract": "In the context of six-dimensional conical braneworlds we consider a simple and explicit model that incorporates long distance modification of gravity and regularization of codimension-2 singularities. To resolve the conical singularities we replace the codimension-2 branes with ring-like codimension-1 branes, filling in the interiors with regular caps. The six-dimensional Planck scale in the cap is assumed to be much greater than the bulk Planck scale, which gives rise to the effect analogous to brane-induced gravity. Weak gravity on the regularized brane is studied in the case of a sharp conical bulk. We show by a linear analysis that gravity at short distances is effectively described by the four-dimensional Brans-Dicke theory, while the higher dimensional nature of gravity emerges at long distances. The linear analysis breaks down at some intermediate scale, below which four-dimensional Einstein gravity is shown to be recovered thanks to the second-order effects of the brane bending. ",
        "introduction": "The most intriguing possibility for the origin of the current acceleration of the Universe~\\cite{Acceleration} may be modifying gravity at very long distances. Within the framework of braneworlds, Dvali, Gabadadze, and Porrati (DGP) proposed a model in which gravity looks five-dimensional (5D) at long distances and 4D at short distances~\\cite{DGP}. In their model, gravity becomes weaker on large scales due to its leakage from our 4D brane to the higher dimensional space, while the induced-gravity term on the brane maintains the 4D nature of short-range gravity. Although the original motivation of the idea of~\\cite{DGP} was to realize 4D gravity even in the presence of infinitely large extra dimensions, it was soon noticed that the model gives rise to the self-acceleration, i.e., the accelerating expansion of the Universe without a cosmological constant~\\cite{SA, Lue_Rev, Koyama_CC}. The modified Friedmann equation in the DGP braneworld is given by~\\cite{SA} \\begin{eqnarray*} H^2-\\frac{H}{r_c}=\\frac{8\\pi G}{3}\\rho, \\end{eqnarray*} where $r_c$ is the crossover scale above which gravity is modified. This allows for the late-time accelerating solution, $H\\to H_{\\infty}=r_c^{-1}$. Despite multiple problematic features of the DGP model~\\cite{Rubakov:2003zb, Ghost0, Ghost1, Ghost2,Kaloper:2005az,Gregory:2007xy}, it is still worth exploring the arena of IR modified gravity in the braneworld context because more elaborated brane models may provide consistent IR modification which has possible relevance to the cosmological constant problem, in terms of a fully covariant and nonlinear theory of gravity.   So far the consequences of brane-induced gravity have been often discussed in codimension-1 models, but it seems quite interesting to consider more general, codimension-$N\\geq 2$ braneworlds with induced gravity. Infinitely thin branes with $N\\geq 2$, however, suffer from certain singularities which must be dealt with some care~\\cite{DG}. The pathology is clearly illustrated by the shock wave solution on thin codimension-2 branes~\\cite{Kaloper}: the shock wave profile is perfectly 4D over all distance scales on the brane, but it vanishes (even infinitesimally) away from the brane. Namely, gravity confinement is too effective. This is due to the fact that the Green function in the space transverse to the brane diverges at the origin. (The divergence is the generic feature for $N\\geq2$.) If one is to resolve this short distance singularity somehow, the regularization at UV will affect the behavior of gravity at IR~\\cite{Kiri, Kaloper, Seesaw, Reg-Rubakov, Reg-DGP, Softly}. (See also different approaches to construct higher codimension braneworlds with induced gravity~\\cite{Int, Cascading}.)   In this paper we consider a 6D brane model that incorporates long distance modification of gravity and regularization of codimension-2 singularities. The codimension-2 branes in our model are regularized in the manner of~\\cite{Kaloper}, replacing them with ring-like codimension-1 branes and filling in the interiors with regular caps. The scheme here is similar to what is widely used in flux-stabilized models of 6D braneworlds~\\cite{Peloso, Papa, Reg-TK, UVcaps}. Then, instead of introducing brane-induced gravity term directly, we use the possible mechanism for realizing {\\em effectively} DGP-type gravity, that is, asymmetry between the two sides of the brane~\\cite{Padilla, KoyamaKoyama}. While the standard way to study codimension-1 brane models in five dimensions is assuming $\\mathbb{Z}_2$-symmetry across the brane~\\cite{RS, Z2}, the bulk-cap system in regularized braneworlds has an asymmetric configuration in general~\\cite{nonZ2}. In addition to this natural geometrical asymmetry, we allow the cap Planck scale, $M_\\rml$, to differ from that in the bulk, $M_\\rmr$. If $M_\\rml\\gg M_\\rmr$, this leads to an explicit realization of the regularization of brane-induced gravity~\\cite{Reg-DGP}. The purpose of the present paper is to investigate in detail the behavior of weak gravity in such a braneworld. Even if gravity is weak, linearization is not justified below a certain distance scale, as in the DGP model~\\cite{Deffayet,Lue:2001gc,Gabadadze:2004iy,Gruzinov:2001hp,Tanaka, KScp}. We identify the scale and then carefully take into account the second order effects of the brane bending to explore the seminonlinear regime.         This paper is organized as follows. In the next section we present our background model of regularized braneworlds. We then study the behavior of linear perturbations in Sec.~III, showing that brane gravity at short distances is described by the 4D Brans-Dicke theory. As an example, in Sec.~\\ref{sec:pointsource} we compute the gravitational field of a static point source modeled by a loop of a string along the ring-like brane. Using this example we confirm that gravity looks five- or six-dimensional at long distances. In Sec.~IV we perform a nonlinear analysis to see how 4D Einstein gravity emerges at even shorter distances. Sec.~V is devoted to conclusions. ",
        "conclusions": "Summarizing, we have studied weak gravity in the 6D regularized braneworld having the mechanism of long distance modification of gravity. We resolved the codimension-2 singularity by capping the apex of the cone-shaped bulk and replacing the conical brane with a ring-like codimension-1 brane. The brane contains a scalar field which cancels the pressure along the compact direction. In regularizing the codimension-2 brane, we have assumed the Planck scale in the cap to be much greater than the bulk Planck scale. This assumption is along the line of the regularization of brane-induced gravity of~\\cite{Reg-DGP}. We have shown that the model {\\em effectively} gives rise to DGP-type gravity without introducing induced-gravity terms on the brane. In the case of the sharp conical bulk $(\\rho_0\\ll r_c\\lesssim r_0)$, we have obtained the following result in the regime where linearization is justified: \\begin{itemize} \\item the 4D Brans-Dicke theory with $\\omega_{{\\rm BD}}=-r_c/r_0$ for $\\mathtt{r}\\ll r_c, r_0$; \\item 5D gravity for $r_c\\ll\\mathtt{r}\\ll r_0$; \\item 6D gravity for $\\mathtt{r}\\gg r_c, r_0$. \\end{itemize} (Of course, the second regime appears if $r_c\\ll r_0$.) The linear analysis is not valid below $R_*=(r_c^2R_g)^{1/3}$, where $R_g$ is the gravitational radius of the source. We have performed a seminonlinear analysis for $\\mathtt{r}\\ll R_*$ by carefully taking into account quadratic terms in the brane bending. Our finding is: \\begin{itemize} \\item 4D Einstein gravity is reproduced for $\\mathtt{r}\\ll R_*$. \\end{itemize}   The present model is very simple in that the background is given by a locally flat spacetime. In spite of this simplicity, the system is shown to have a rich and interesting structure. It is easy to extend the model to more general background configurations which are curved due to the cosmological constant and other fields living in the bulk. The similar regularization scheme can be applied to such cases as well. Bulk-cap asymmetry will then help to reproduce 4D gravity on the brane~\\cite{Padilla, KoyamaKoyama}.     As mentioned in the Introduction, long distance modification of gravity is often motivated by explaining the current acceleration of the Universe. It would be interesting to construct cosmological solutions and explore whether the accelerating expansion is possible\\footnote{ It has been known that the Lovelock term in the bulk leads to induced gravity on the brane~\\cite{Ig, Lovelock, Lovelock2}. Self-accelerating solutions were found in the framework of 6D Gauss-Bonnet braneworlds in~\\cite{Charmousis}. } due to gravity leakage in the present model. We hope to address this issue in the near future."
    },
    "0806/0806.0323_arXiv.txt": {
        "abstract": "Supercriticality of the same kind as that in a nuclear pile can take place in high-energy astrophysical objects producing a number of impressive effects. For example, it could cause an explosive release of the energy of a cloud of ultrarelativistic protons into radiation. More certainly, supercriticality should be responsible for energy dissipation of very energetic relativistic fluids such as ultrarelativistic shocks in gamma-ray bursts and jets in active galactic nuclei (AGNs). In this case, the photon breeding process operates. It is a kind of the converter mechanism  with the high-energy  photons and $e^+e^-$ pairs converting into each other via pair production and inverse Compton scattering. Under certain conditions, which should be satisfied in powerful AGNs, the photon breeding mechanism becomes supercritical: the high-energy photons breed exponentially until their feedback on the fluid changes its velocity pattern. Then the system comes to a self-adjusting near-critical steady state. Monte-Carlo simulations with the detailed treatment of particle propagation and interactions demonstrate that a jet with the Lorentz factor $\\Gamma\\approx 20$ can radiate away up to a half of its total energy and for $\\Gamma=40$ the radiation efficiency can be up to 80 per cent. Outer layers of the jet decelerate down to a moderate Lorentz factor 2--4, while the spine of the jet has the final Lorentz factor in the range 10--20 independently on the initial  $\\Gamma$. Such sharp deceleration under the impact of radiation must cause a number of interesting phenomena such as formation of internal shocks and an early generation of turbulence. ",
        "introduction": "The term ``supercriticality'' has different meanings. Here we use the same meaning as in nuclear physics: the condition of exponential runaway behaviour of the system (nuclear explosion or nuclear pile). So far  two kinds of supercritical phenomena   were proposed in high-energy astrophysics: an explosive energy release of a cloud of ultrarelativistic protons into photons (``proton bomb'') and dissipation of the bulk energy of relativistic fluids into radiation.  Let us consider a magnetized cloud of relativistic protons with typical Lorentz factor $\\gamma_p$. These protons do not radiate much, because the synchrotron radiation and the inverse Compton scattering rates for protons are suppressed by factor $(m_e/m_p)^2$ with respect to electrons. Suppose $N$ seed photons with energy $\\varepsilon_t \\sim 200{\\rm MeV}/\\gamma_p$ interact with protons producing pions decaying into high-energy electrons or photons. They, in turn, produce a wide spectrum of softer photons, with the number $\\xi N$ of them on average interacting with the high-energy protons again. If  $\\xi > 1$  then the process has a supercritical character: the number of photons rises exponentially until the protons cool down, so that the supercriticality condition breaks down and the energy release terminates. If the energy supply to protons continues then the system reaches supercriticality and explodes again. Such scenario was proposed and studied with numerical simulations \\cite{SS91} and analytically \\cite{KM92}.  The supercriticality condition depends on the radiation spectrum emitted by the particles produced in photo-meson interaction. The hardest possible spectrum is the fast-cooling power-law $dN/d\\varepsilon \\propto \\varepsilon^{-3/2}$. The softest spectrum is that of the saturated pair cascade  $dN/d\\varepsilon \\propto \\varepsilon^{-2}$ \\cite{sve87}, which takes place in a dense soft photon field.  The cross-section of photo-meson production multiplied by the fraction of energy transfered to pions, $x$, is (see \\cite{ber01}) $\\sigma (p\\gamma) x \\sim 0.7\\times 10^{-28} {\\rm cm}^2$ in the photon energy range $\\sim 250$ MeV--2 GeV in the proton rest frame. The number of produced photons which can interact again is $N \\sim 1/\\gamma_p$ in the first case and $N \\sim 1/\\gamma_p^2$ in the second case. Then condition $\\xi > 1$ translates to $\\rho_E > 10^{25}$ erg cm$^{-2}$ in the fast cooling case and  $\\rho_E > 10^{25}/\\gamma_p$ erg cm$^{-2}$ in the cascade case, where $\\rho_E$ is the column energy density of protons across the cloud.   Where one can find such a cloud of ultrarelativistic protons? A possible site for this phenomenon, gamma-ray bursts (GRBs), was proposed by  \\cite{KGM02}. Indeed, all scenarios of GRBs involve ultrarelativistic shocks, where the space between the shock front and the contact discontinuity is filled with ultrarelativistic protons.  Another possible class of supercritical phenomena can take place in relativistic flows. A first example  is the  supercritical pair loading in AGN jets near  the accretion disc's hot corona which emits hard X-rays \\cite{Lev96}. If the jet Lorentz factor at such distance is $\\Gamma \\sim 10$, then cool pairs moving with the jet can up-scatter X-rays to MeV energies and these photons in their turn, interacting with the same X-rays, can produce new pairs.  A more powerful and universal way of dissipation of relativistic fluids into radiation, is associated with the converter mechanism suggested in general form by Derishev et al. \\cite{Der2003}  and  in a more specific form of a runaway (supercritical) electromagnetic cascade in ultrarelativistic shocks by Stern \\cite{St2003}. We later performed a numerical study of the converter mechanism, in its electromagnetic version, operating in relativistic jets of blazars  \\cite{SP06,SP08} and found that it has a supercritical character: the high-energy photons and relativistic $e^+e^-$ pairs breed exponentially. ",
        "conclusions": ""
    },
    "0806/0806.0862_arXiv.txt": {
        "abstract": "Galaxies hosting $z\\sim 2$ quasars are the high-$z$ progenitors of today's massive `red-and-dead' galaxies. With close pairs of quasars at different redshifts, a background quasar can be used to study a foreground quasar's halo gas in \\emph{absorption}, providing a wealth of information about feedback, quenching, and the physics of massive galaxy formation.  We present a Keck/HIRES spectrum of the bright background quasar in a projected pair with angular separation $\\theta = 13.3''$ corresponding to $\\rperp =108$\\,kpc at the redshift of the foreground quasar $z_{\\rm fg}= 2.4360 \\pm 0.0005$, precisely determined from Gemini/GNIRS near-IR spectroscopy. Our echelle spectrum reveals optically thick gas ($\\mnhi \\approx 10^{19.7} \\cm{-2}$) coincident with the foreground quasar redshift. The ionic transitions of associated metal-lines reveal the following properties of the foreground quasar's halo: (1) the kinematics are extreme with absorption extending to $+780 \\mkms$ relative to $z_{\\rm fg}$; (2) the metallicity is nearly solar; (3) the temperature of the predominantly ionized gas is $T \\lesssim 20,000$K; (4) the electron density is $n_{\\rm e} \\sim 1 \\cm{-3}$ indicating a characteristic size $\\sim 10 - 100$\\,pc for the absorbing `clouds'; (7) there is little (if any) warm gas $10^{5}\\,{\\rm K} \\lesssim T \\lesssim 10^{6}\\,{\\rm K}$; (8) the gas is unlikely illuminated by the foreground quasar, implying anisotropic or intermittent emission. The mass of cold $T\\sim 10^{4}{\\rm K}$ gas implied by our observations is significant, amounting to a few percent of the total expected baryonic mass density of the foreground quasar's dark halo at $r\\sim 100\\,{\\rm kpc}$. The origin of this material is still unclear, and we discuss several possibilities in the context of current models of feedback and massive galaxy formation. ",
        "introduction": "\\label{sec:intro}  Over the course of a quasar's lifetime, the accretion of material onto its $\\sim 10^9\\,\\msol$ supermassive black hole will liberate an enormous energy $E=\\epsilon M_{\\rm BH} c^2 \\simeq 2\\sci{62}\\,{\\rm ergs}$\\footnote{This energy is appropriate for a typical quasar near the peak of activity at $z\\sim 2.5$, a radiative efficiency $\\epsilon\\sim 0.1$, and an Eddington ratio of $0.1$.} affecting its environment from pc to Gpc scales. On cosmological scales, the collective emission from quasars dominates the ultraviolet background radiation field \\citep{hm96} that maintains roughly $90\\%$ of the universe as a highly ionized plasma \\citep{gp65}, while the harder photons produce the cosmic X-ray background \\citep{fi99,erz02,Ueda03,Cao05}. On Mpc scales, the ultraviolet flux photoionizes the nearby IGM resulting in the proximity effect, or the lower optical depth for \\lya\\ absorption at the quasar redshift \\citep{bdo88,sbd+00}. On the smallest scales of 1-100\\,pc, the interplay between black hole accretion, radiation, and the surrounding gas, produce the characteristic emission from the well-studied broad and narrow emission line regions \\citep[see][]{ferost06}.  But the degree to which a quasar can influence its host galaxy on scales $\\sim\\,{\\rm kpc-Mpc}$ is much less clear. Indeed, just a few percent of the energy $\\sim 10^{62}\\,{\\rm ergs}$ emitted by a quasar is comparable to the binding energy of a massive galaxy $\\sim 10^{60}$\\,erg, and thus capable of ejecting its interstellar medium (ISM) or shock-heating it to temperatures $T\\sim 10^{7}$\\,K.  The bulges of all local galaxies harbor supermassive black holes \\citep{KormRich95}, the masses of which are tightly correlated with the properties of their host spheroids, measured on scales $\\sim$ kpc \\citep{Magorrian98,FerrMerr00,Gebhardt00}.  This has led many to speculate that some ``feedback'' mechanism couples the quasar phase of rapid supermassive black hole growth with the evolution of its host galaxy \\citep[e.g.][]{SilkRees98,kh00,WyitheLoeb03,gds+04,ScannOh04,Springel05,Scann05,Kawata05,Menci06,Cox06,LuMo07,Sijacki07,Hopkins08}.  Besides explaining the correlation between black holes and bulges, feedback from an AGN has also been invoked on even larger scales, as the energy source which ``quenches'' star-formation in massive galaxies, leaving them ``red and dead''.  The observed bimodality in the galaxy color-magnitude diagram \\citep{Strateva01,baldry04,bell04,blanton05,faber07} and the sharp cutoff at the bright end of the galaxy luminosity function \\citep{bbf+03} both point to some physical mechanism which shuts off star-formation in massive galaxies. Otherwise, large quantities of gas should have accreted onto the progenitors of ellipticals resulting in an unseen population of very massive blue galaxies. It is generally believed that the coupling between the quasar and the galaxy is ``kinetic'' in nature, involving galactic-scale outflows which suppress gas accretion and inject heat, ultimately shutting off star-formation \\citep[e.g.][]{kh00,ScannOh04,Springel05,Cox06,hhc+07}.   An important laboratory for studying the astrophysics of quasar feedback and quenching are luminous high redshift radio galaxies \\citep[HzRGs; for a review see][]{Mccarthy93,mileyd08}.  In these sources outflowing collimated relativistic plasma, or ``radio jets'', can extend to several tens of kpc from the optical host galaxies, providing the most compelling evidence that an AGN can impact its galactic scale surroundings. It is unclear, however, whether this jet energy $\\gtrsim 10^{60}$\\,erg \\citep[e.g.][]{Miley80} couples strongly to the ambient ISM \\citep{bc89}.  In addition to the tremendous radio power, $z\\gtrsim 2$ HzrGs are often associated with giant Ly$\\alpha$ recombination nebulae, with sizes of up to $\\sim 200$\\,kpc, sometimes extending well beyond the radio emission \\citep[see][for a review]{villar07}.  Ionizing radiation from the central AGN, thought to be obscured from our vantage point, is likely to play a significant role in exciting this emission. But the extreme kinematics of these nebulae \\citep[${\\rm FWHM}\\gtrsim 1000\\mkms$][]{nle+06,reuland07,villar07}, their complex irregular morphology, the tendency for the line emission to be aligned with the axis of the radio jets, and their near solar metallicities \\citep{vernet01,humphrey08}, strongly suggest that jet-gas interactions or feedback from the AGN are giving rise to a large-scale outflow, and shocks from these motions could also be triggering the emission. But the comoving number density of HzRGs and giant \\lya\\ nebulae, $n \\sim 10^{-8}\\ , {\\rm Mpc}^{-3}$ \\citep{mileyd08}, is about three orders of magnitude smaller than the number density of luminous quasars. Thus only about one in a thousand of the supermassive black holes in the Universe exhibit such dramatic evidence for feedback during their active phase\\footnote{Because the implied lifetime of the jets and nebulae are $\\sim 10^8$ years, they cannot represent a short lived phase of evolution in every quasar.}, which would do little to quench star-formation in the entire population of local massive galaxies. Is feedback occurring in the typical quasar in the Universe?              While it has been argued that heating from a quasar is responsible for quenching massive galaxies \\citep{sdh05,hopkins07}, an alternative scenario is related to the physics of structure formation in a hierarchical Universe.  This idea goes back to a classic argument first made by \\citet{ro77} that the ability to cool on a dynamical timescale is what sets the upper bound for the mass of luminous galaxies. At low halo masses the cooling time of gas is shorter than the dynamical time resulting in a ``cold accretion'' mode, limited only by the supply of infalling material.  For halos with $M \\gtrsim 10^{12}\\msol$, the cooling time exceeds the dynamical time, which is the criterion for the formation of a stable accretion shock \\citep{db06}.  In these more massive halos, subsequent infalling gas is shock heated and pressure supported by a hot halo in quasi-hydrostatic equilibrium. However, cosmological hydrodynamical simulations have shown that although the majority of gas is shock-heated to the virial temperature, some level of ``cold mode'' accretion still takes place in even the most massive halos \\citep{kkw+05,ocvirk08}. It is unknown whether the formation of a hot-halo alone can quench future star-formation and explain the red and dead galaxy population. \\citet{db08} recently proposed a quenching scenario whereby the gravitational energy of cosmological accretion can serve as a heat source to quenching star formation, provided that a cold ($T\\sim 10^4$\\,K) phase of pressure-confined (and thus long lived) gas clumps ($M\\sim 10^5-10^8\\msol$) form, which can penetrate the halo and deliver the necessary heat.  Understanding the role that quasar feedback, hot halos, cold accretion, and pressure confined clumps play in quenching star-formation, requires studying the physical state of the gas on scales 10\\,kpc-1\\,Mpc in the high redshift progenitors of local massive red galaxies.  \\citet[][see also Wechsler et al.\\ 1998]{cst+07} model the clustering and number density of $z\\sim 2-3$ star-forming galaxies and deduced host dark halo masses $M \\lesssim 10^{12}~\\msol$. These low masses allowed them to convincingly argue that $z\\sim 2-3$ star-forming galaxies \\emph{do not} evolve into the quenched red and dead galaxies that we see today, but rather evolve into blue $\\sim L_{\\ast}$ galaxies similar to the Milky Way. But the strong clustering of luminous quasars at $z\\sim 2$ implies larger dark halo masses $\\gtrsim 10^{13}~\\msol$ \\citep{croom01,pmn04,croom05}, indicating that they \\emph{are indeed the progenitors of local red and dead galaxies}. Of course, we are only observing a small fraction $\\sim 10^{-2}-10^{-3}$ of these progenitors as quasars at any given time because the quasar lifetime is much shorter than the age of the Universe.  However quasar spectra have thus far provided little insight into the state of gas in their galactic environments for two reasons. First, the large ionizing flux from the quasar typically photoionizes the hydrogen in and around the galactic host \\citep{hp07,phh08,cmb+08}, which would otherwise be detected as strong \\ion{H}{1} absorption at the quasar redshift. Second, the interpretation of material that is detected in absorption along the line-of-sight to a quasar is limited by the unknown distance between the material and the quasar. A systematic search for the signatures of photoionized gas near quasars (e.g.\\ via \\ion{N}{5}, \\ion{O}{6} absorption) is only now being carried out on statistical datasets \\citep{tripp08,fbp08}. Observers do frequently identify narrow associated absorption lines (NAALs) or so-called ``associated absorbers'' which are believed to arise in the quasar environment. Some of these are attributed to gas on sub-pc scales \\citep[e.g.][]{elvis00}, but other examples exhibit ionization states that suggest the gas lies at several tens of kpc \\citep{dcr+04,rps+07}. Note that these associated absorbers are distinct from the broad absorption line (BAL) features which arise from dense clumps of material that have been ionized and accelerated by the quasar, and are generally believed to be confined to sub-kpc scales \\citep[but see][]{dab+01}.  In a series of papers, we have introduced a novel technique to study the physical state of gas in the ISM and halo of luminous quasars, which has the potential to provide powerful constraints on feedback, quenching and the physics of massive galaxy formation.  Namely, we use a background quasar (b/g quasar) sightline to probe the state of gas in absorption in the vicinity of a foreground quasar \\citep[f/g quasar; see also][]{bhm+06}.  Although such projected quasar pair sightlines are extremely rare, \\citet[][see also Hennawi 2004]{hso+06} showed that it is straightforward to select $z\\gtrsim 2$ projected quasar pairs from the imaging and spectroscopy provided by the Sloan Digital Sky Survey \\citep[SDSS;][]{yaa+00}.  To date about 90 pairs of quasars have been uncovered with impact parameter $R < 300~{\\rm kpc}$ and $z_{\\rm fg} > 1.6$\\footnote{The lower limit on redshift is motivated by the ability to detect redshifted \\lya\\ absorption above the atmospheric cutoff $\\lambda > 3200$~\\AA}. Spectroscopic observations of the b/g quasar in each pair reveals the nature of the IGM transverse to the f/g quasar on scales of a few 10\\,kpc to several Mpc. This approach has the advantage of tracing diffuse gas over a wide range of density and temperature, ranging from cold neutral material $T\\sim 100\\,{\\rm K}$ to collisionally ionized plasma $T \\approx 10^6$\\,K, and with column densities in the range $N\\sim 10^{12}-10^{22}\\,{\\rm cm^{-2}}$.  In Paper~I \\citep{hpb+06}, we searched 149 b/g quasar spectra for optically thick absorption in the vicinity of $1.8 < z_{\\rm fg} < 4.0$ luminous f/g quasars, and uncovered a sample of 27 new quasar-absorber pairs with impact parameters ranging from $30\\,{\\rm kpc}-2.5\\,{\\rm Mpc}$.  In Paper~II \\citep{hp07}, we analyzed the clustering of these transverse absorbers with the foreground quasar and measured a large clustering signal on galactic scales.  We also refer the reader to the manuscript by \\cite{tytler08} who consider several applications of quasar pair spectroscopy.   In this paper, we present the first high resolution spectrum of the b/g quasar in the close projected quasar pair \\sdssj.  By mining the sky for very rare close associations of quasars \\citep{thesis,hso+06,hpb+06}, we previously discovered this rare system with angular separation $\\theta = 13.3''$ corresponding to impact parameter $\\rperp =108$\\,kpc at the redshift of the foreground quasar $z_{\\rm fg}= 2.4360 \\pm 0.0005$, precisely determined from Gemini/GNIRS near-IR spectroscopy.  The spectral and photometric properties of \\fgqso\\ make it an unremarkable quasar at $z \\sim 2.4$. We estimate a bolometric luminosity of $L_{\\rm QSO} \\simeq 1.4\\sci{46}{\\rm erg\\, s^{-1}}$ placing it near the `knee' of the $z\\sim 2.5$ quasar luminosity function \\citep{croom04,richards06}, and corresponding to a supermassive black hole $M_{\\rm BH} \\simeq 1.1\\sci{9}\\left(\\frac{f_{\\rm Edd}}{0.1}\\right)^{-1}\\,\\msol$, if it accretes at one tenth of the Eddington limit.  The impact parameter of our background sightline $\\rperp =108$\\,kpc, easily resolves the expected virial radius $r_{vir} = 250\\,{\\rm kpc} (M/10^{13.3} \\msol)^{1/3}$ of the f/g quasar host, and pierces its halo at about the `cooling radius', where gas shock-heated to the virial temperature should take about a Hubble time to cool.  The only remarkable thing about \\sdssj\\ is that it has a $z_{\\rm bg} = 2.53$ b/g quasar in close projection which is bright enough ($r=19.0$) for high resolution spectroscopy. Our Keck HIRES Echelle spectrum of \\bgqso, the first ever to probe the halo gas of a f/g quasar, resolves the velocity fields of the absorbing gas and allows us to measure precise column densities for \\ion{H}{1} and the ionic transitions of metals like Si, C, N, O, and Fe. These measurements allow us to place constraints on the physical state of the gas near the f/g quasar, such as its kinematics, temperature, ionization structure, chemical enrichment patterns, volume density, the size of the absorbers, the intensity of the impingent radiation field, as well as test for the presence of hot collisionally ionized gas.  In $\\S$~\\ref{sec:obs} we present the observations and provide column density measurements.  We constrain the ionization state, estimate relative chemical abundances, and constrain the electron density of the gas in $\\S$~\\ref{sec:ion}.  In $\\S$~\\ref{sec:discuss} we further discuss our results and how they relate to other observations of quasars, absorption line systems, and high redshift galaxies. The implications of our results for models of feedback, quenching, and massive galaxy formation are presented in $\\S$~\\ref{sec:model}, and we conclude with a summary in $\\S$~\\ref{sec:summ}.  The reader who is not concerned with the details of the observations and absorption line modeling can read Table~\\ref{tab:summ} summarizing those results and then skip to $\\S$~\\ref{sec:discuss}.  Throughout the manuscript, we use the cosmological parameters $\\Omega_m = 0.30$, $\\Omega_\\Lambda =0.70$, $h=0.70$, consistent to within $1-\\sigma$ with the parameters measured by the WMAP experiment \\citep{wmap05} and we adopt the solar chemical composition compiled by \\cite{gas07}.   \\begin{figure} \\begin{center} \\includegraphics[width=6.8in]{f1.eps} \\end{center} \\caption{GNIRS spectrum of \\fgqso\\ centered on the H$\\beta$-[\\ion{O}{3}] emission line complex (H$\\beta$~$\\lambda 4861$, [\\ion{O}{3}]~$\\lambda 4959$, and [\\ion{O}{3}]~$\\lambda 5007$) which is redshifted into the $H$-band. This spectrum was used to determine the systemic redshift of the quasar, $z_{\\rm fg} = 2.4360 \\pm 0.0005$, from the strong [\\ion{O}{3}]~$\\lambda 5007$ emission line. % The lower red histogram is the 1-$\\sigma$ noise and the smooth blue curve overplotted on the data is the composite quasar template of \\citet{vanden01} redshifted to $z_{\\rm fg}$. The upper panel shows the atmospheric transmission for the corresponding wavelengths.} \\label{fig:GNIRS} \\end{figure} ",
        "conclusions": "\\label{sec:discuss}   In this section we discuss the observed properties of \\slls\\ in terms of the galactic environment around high $z$ quasars.  We examine the velocity field of this gas, its chemical abundances, and then test whether \\fgqso\\ may be shining on the absorbing material.  As our discussion focuses on a single system, the degree to which we can generalize these results is questionable. It is nevertheless fruitful to explore various scenarios for interpreting the observations. Table~\\ref{tab:summ} summarizes the physical characteristics of \\slls\\ derived in the previous section.   \\subsection{Kinematic Constraints} \\label{sec:kinematic}  The most robust measurements for \\slls\\ from our high resolution spectrum of \\bgqso\\ are on the velocity field of the gas.  Both the absolute (i.e.\\ redshift) and relative velocities constrain the origin of the gas and its relation to the foreground quasar \\fgqso.  The redshift of \\fgqso\\ was measured from our observations of the rest-frame optical [\\ion{O}{3}]~$\\lambda 5007$ emission line using the GNIRS spectrometer on the Gemini-S telescope ($\\S$~\\ref{sec:obs}, Figure~\\ref{fig:GNIRS}). We achieved a velocity precision of $\\approx 40$\\kms\\ giving $z_{\\rm fg} = 2.4360 \\pm 0.0005$.  Figures~\\ref{fig:HI} and \\ref{fig:metals}, which present the velocity profiles of \\slls\\ relative to $z_{\\rm fg}$ show strong \\ion{H}{1} and metal-line absorption from $\\delta v \\approx +70\\mkms$ to +700\\,\\kms.  These absorption profiles have two notable characteristics.  First, the gas along the \\bgqso\\ sightline is asymmetrically distributed relative to $z_{\\rm fg}$, i.e.\\ within a 2000\\kms\\ window centered on $z = z_{\\rm fg}$, there is significant absorption only at $z > z_{\\rm fg}$.  The nearest `cloud' with $\\tau_{Ly\\alpha} > 1$ and $\\delta v < 0 \\mkms$ lies off of Figure~\\ref{fig:HI} at $\\delta v \\approx -1500\\mkms$, which corresponds to $\\approx 6 h^{-1}$ Mpc (proper) under the assumption of Hubble expansion.  Second, the total velocity width is very large, $\\delv \\approx 650 \\mkms$.  One may compare this value against the velocity widths measured for systems with comparable \\ion{H}{1} surface density, i.e.\\ the damped \\lya\\ (DLA) and super Lyman limit systems.  Regarding the former, the velocity width of \\slls\\ exceeds greater than $99\\%$ of the $\\delv$ values for randomly selected DLAs at $z>2$ \\citep{pw97,pw01}.  Regarding the SLLS population \\citep{peroux_slls03}, none have as large of a velocity width.  The equivalent width of the metal-line transitions (a proxy for velocity width) is also large compared to systems selected on the basis of metal-line absorption \\citep[e.g.\\ the \\ion{Mg}{2} systems;][]{ppb06}. In these respects, {\\it the velocity field of the gas near \\fgqso\\ is extreme}.  Several factors lead us to conclude that the distance between the gas in \\newline \\slls\\ and \\fgqso\\ is comparable to the impact parameter.  Recall that \\slls\\ exhibits a multi-component kinematic structure with unusually large velocity width between the components. If the relative motions of the clouds and the quasar are to be interpreted as Hubble flow, their implied line-of-sight separation would be $\\approx 2.8$ \\,Mpc (for $\\Delta v \\approx 700\\mkms$). The chance alignment of several galaxies across this large distance is unlikely however, which we can quantify with clustering arguments.  Quasars arise in highly biased regions of the Universe. They are observed to cluster strongly on $\\sim 10$\\,\\hMpc\\ scales \\citep{croom01,pmn04,croom05,myers07a,shen07} and the correlation function steepens considerably on smaller scales $\\sim 100$\\,\\hkpc\\ \\citep{thesis,hso+06,myers07b,myers08}.  High redshift quasars also exhibit a large cross-correlation amplitude with star-forming galaxies \\citep{as05,coil06}, and it is possible that galaxies clustered around the quasar are giving rise to the \\lya\\ and metal-line absorption seen in \\slls.  Another possibility is that the absorption in \\slls\\ is not due to a nearby galaxy, but is rather halo gas distributed around the quasar with some density profile. Either scenario results in a clustering pattern around the quasar, which we measured in \\cite{hp07}. The probability $P(<r|\\rperp, v_{\\rm max})$ that the total distance between the quasar and the absorber is less than $r$, given a fixed impact parameter $\\rperp$ and the knowledge that absorption occurs within a velocity interval $\\pm v_{\\rm max}$, is simply an integral over this quasar-absorber correlation function.   \\begin{figure} \\begin{center} \\includegraphics[width=6.8in]{f6.eps} \\end{center} \\caption{Probability distribution of the distance from the foreground quasar. The function $P(<r|R_{\\perp}, v_{\\rm max})$ represents the probability that the distance from the foreground quasar is less than $r$, given a fixed impact parameter $R_{\\perp}$ and knowledge that the absorber lies in the velocity interval $\\pm v_{\\rm max}$. The impact parameter for \\sdssj\\ is used and a velocity interval of $v_{\\rm max} = 700\\mkms$ was assumed.  This probability is computed from the quasar absorber correlation function $\\xi_{\\rm QA}=(r\\slash r_0)^{-\\gamma}$ measured by \\citet{hp07}, with $\\gamma=2$ and $r_0 = 5.8$\\,\\hMpc. The solid (black) curve illustrates the probability distribution accounting for clustering with $\\xi_{\\rm QA}$. The dashed (red) curve shows the probability if there is no clustering. The upper x-axis shows the ionizing photon flux $\\Phi_{\\rm QSO}$ implied by the SDSS magnitude of \\fgqso\\ at the given distance.} \\label{fig:cluster} \\end{figure}   Figure~\\ref{fig:cluster} shows this cumulative probability for $v_{\\rm max} = 700\\mkms$ and the impact parameter of $\\rperp = 108$\\,kpc, from which we deduce that the probability for the gas in \\slls\\ lying within 500\\,kpc of \\fgqso\\ is $\\approx 70\\%$. Perhaps an even stronger argument that the gas in \\slls\\ lies at a distance comparable to the impact parameter is that the physical characteristics (ionization state, \\ion{H}{1} surface density, chemical abundances) of the subsystems comprising \\slls\\ show only modest variations. Whereas, several of these same properties --- specifically the metallicity, relative abundance pattern, and the velocity field --- are highly anomalous compared to the population of intervening optically thick absorbers. In particular, the only place that such large N/O ratios are observed is in quasar environments. Finally, it is worth emphasizing that the very small sizes deduced for the subsystems of \\slls\\ ($d \\lesssim 100$\\,pc) imply that their relative separations are also comparable to the impact parameter. The alternative, which would be for the subsystems to be part of a single structure with size $d \\lesssim 100$\\,pc, would imply an unphysically large velocity shear across this small distance.    \\begin{figure} \\begin{center} \\includegraphics[width=6.0in]{f7.eps} \\end{center} \\caption{ Observed gas-phase (a) $\\alpha$/Fe and (b) N/$\\alpha$ abundances as a function of alpha abundance [$\\alpha$/H] for galaxies intervening quasar sightlines \\citep[DLAs;][]{pw01,pwh+07}, the three subsystems associated with \\fgqso\\ (blue dotted), and the \\nh-weighted average values (blue solid). We note the gas associated with \\fgqso\\ has systematically higher metallicity and higher $\\alpha$/Fe, N/O ratios than the majority (or all) DLAs.  This indicates gas that was enriched in a `starburst' fashion, i.e.\\ a system with a very high star formation efficiency \\citep{cmv03}. } \\label{fig:relabnd} \\end{figure}   \\subsection{Clues from Chemical Enrichment} \\label{sec:chem}  Clues to the origin of \\slls\\ are encoded in the enrichment pattern of its gas.  In particular, the relative abundances constrain the timescales and intensity of star-formation of the galaxy (or galaxies) that produced the observed metals.  Because the gas is partially ionized and we do not observe transitions from all ionization states, we proceed cautiously to interpret the observed ionic ratios.  One robust conclusion from our abundance analysis ($\\S$~\\ref{sec:ion}; Table~\\ref{tab:summ}) is that the gas is highly enriched.  The observed O$^0$/H$^0$ ratio indicates [O/H]~$>-1$ across the $\\approx 650 \\mkms$ velocity interval spanning the entire system and we derive an integrated average abundance: [O/H]$_{\\rm TOT} > -0.5$\\,dex.  This metallicity is expressed as a lower limit for several reasons. First, ionization corrections to the observed O$^0$/H$^0$ ratios will give larger [O/H] values (Figure~\\ref{fig:ioncorr}).  Second, the lower limit assumes the nearly maximal \\nhi\\ values permitted by the data for the three subsystems. If we adopt a modest ionization correction and lower \\nhi\\ values, the oxygen metallicity for \\slls\\ approaches the solar abundance. Finally, a proper derivation of the oxygen metallicity for the SLLS is to take an \\nh-weighted average of each subsystem. This is especially important for predominantly ionized gas. Using our fidcuial values, we find $\\rm <[O/H]> = -0.23$\\,dex assuming no ionization corrections. We emphasize that even a 1/3 solar metallicity ([O/H]~$= -0.5$\\,dex) represents a very high enrichment level. It exceeds the metallicity of nearly all optically thick absorbers at these redshifts: 98\\% of sightlines randomly intersecting galaxies \\citep[DLAs;][]{pgw+03} and 95\\% of the super-LLS population \\citep{pdd+07}.  Such high enrichment levels are generally observed only in gas within and around star-burst galaxies \\citep{prs+02,simcoe06} and in quasar environments \\citep{dhs03,dcr+04,agk+07}.  The observations also constrain the relative abundances of O, Fe, and N, e.g.\\ the O/Fe and N/O ratios.  These diagnostics are valuable for constraining the star formation history for \\slls\\ because each element is associated with a unique nucleosynthetic site and has a unique production timescale. Oxygen is made in massive stars and one predicts minimal delay ($t < 100$\\,Myr) for its nucleosynthesis and release to the ISM following the onset of star formation. The production of nitrogen and iron, however, is expected to be delayed by several 100\\,Myr to $\\approx 1$\\,Gyr because these elements are made in intermediate mass stars and Type\\,Ia SN respectively. By examining the relative abundances of O, N, and Fe, therefore, one may constrain the duration of the star-formation episode that produced the metals.  Unfortunately, the relative abundance ratios of O, N, and Fe are more sensitive to the ionization state of the gas.  The most conservative treatment is to report lower limits to O/Fe and N/O from the observed ionic column density ratios of O$^0$/Fe$^+$ and N$^0$/O$^0$.  As with O$^0$/H$^0$, corrections to these ionic ratios are positive definite (Figure~\\ref{fig:ioncorr}). We measure lower limits to [O/Fe] of $+0.1, +0.5, +0.8$\\,dex for subsystems~A,B,C respectively and conclude the gas has a super-solar $\\alpha$/Fe ratio.  For even a modest ionization correction, the integrated average $\\alpha$/Fe ratio approaches ten times the solar ratio and exceeds the value measured for nearly all intervening absorption line systems. In Figure~\\ref{fig:relabnd}, we present an estimate for these values for the various subsystems and also for the \\nh-weighted averages against measurements for a set of high-$z$ galactic observations.  In the local universe, large $\\alpha$/Fe ratios are observed in metal-poor stars \\citep[e.g.][]{mcw97}, massive early-type galaxies \\citep{tfw+00}, and highly depleted ISM gas \\citep[e.g.][]{ss96}. The first two examples reflect nucleosynthetic enhancements and are attributed to the Type~II supernovae of massive stars.  The latter example, however, results from the highly refractory nature of Fe relative to O, i.e.\\ one observes enhanced O/Fe gas-phase abundances in the ISM because Fe is preferentially depleted into dust grains. The observed $\\alpha$/Fe enhancement in \\slls\\ could be due to one or both of these effects, but our expectation is that differential depletion is not dominant. A high depletion level would be unusual for gas with modest surface and volume densities that is predominantly ionized with a temperature of $T \\approx 20,000$K. It is more likely that the super-solar $\\alpha$/Fe ratios reflect an intrinsic, nucleosynthetic enhancement.  Granted the high metallicity of \\slls, the $\\alpha$/Fe enhancement suggests a star formation history representive of bright spheroids \\cite[e.g.\\ bulges, elliptical galaxies][]{gfs+07}. Chemical evolution modeling suggests these stellar systems underwent a short period of intense star formation where the gas was rapidly enriched by supernovae from massive stars before Type~Ia SN could contribute \\citep[e.g.][]{matteucci94}.   We conclude that the metals for \\slls\\ were produced in a starburst lasting less than $\\sim 1$\\,Gyr.  The N/O ratio of the gas offers additional constraints on the star-formation history.  Similar to O/Fe, ionization corrections may be important and the conservative approach is to adopt the N$^0$/O$^0$ ratio as a lower limit to N/O.  We find [N/O]~$> -0.7$\\,dex for subsystems~A and B and [N/O]~$> -0.2$\\,dex for subsystem~C (Table~\\ref{tab:summ}). The observations, specifically the large N$^+$/N$^0$ ratio, indicate non-zero ionization corrections. We estimate the correction from photoionization modeling to be $+0.25$ to $+0.5$\\,dex for $\\log U = -4$ to $-3$\\,dex and conclude that N/O is at least 1/3 the solar ratio in each subsystem and may be super-solar in subsystem~C.  These values are larger than N/O values observed for quasar absorption line systems \\citep[Figure~\\ref{fig:relabnd}, e.g.][]{henrypro07}. Indeed, solar N/O ratios are commonly observed only in quasar environments \\citep{hamannf99,dhs03,dcr+04,agk+07}. This further reflects the high metallicity of \\slls\\ because N behaves as a `secondary element' with its yield scaling as the square of the enrichment level (M/H)$^2$. Therefore, one expects near-solar or even super-solar N/O abundances in high metallicity environments. The exception to this expectation is if the system is too young for many intermediate mass stars to have cycled through the AGB phase. The observation of nearly solar N/O in \\slls, therefore, implies the gas was enriched over an episode of at least a few 100\\,Myr \\citep{mp93,rsm+02}.   As noted above, the high metallicity and solar or super-solar N/O and $\\alpha$/Fe ratios for \\slls\\ describe a chemical abundance pattern that matches the abundances commonly measured for gas near quasars. This includes the high metallicities inferred from emission lines from the quasar broad line region and the relative abundances derived for narrow associated absorption lines \\citep{dhs03,dcr+04}.  We are inclined, therefore, to causally connect the enrichment pattern of \\slls\\ with \\fgqso.  A key question that follows is whether the metals were produced by the quasar's host galaxy or whether they originated in a neighboring galaxy or galaxies.  The latter hypothesis would imply that galaxies near quasars have enrichment histories similar to that inferred for the quasar environment, which could be tenable considering that early-type galaxies near the centers of modern clusters have very uniform colors and spectra indicating a common age of formation and star-formation history \\citep[e.g.][]{ble92}. Alternatively, as we discuss further in the next section, the metals may have been generated during an early episode of star formation in the host galaxy of \\fgqso, and then transported to large radii $\\rperp \\approx 100$\\,kpc via a large scale outflow in $\\approx 100$\\,Myr time.  To achieve a high N/O ratio while maintaining a super-solar $\\alpha$/Fe ratio, the data suggest a starburst lasting several 100\\,Myr but less than $\\approx 1$\\,Gyr.  It is intriguing that the implied sequence of early, intense star-formation followed by an optically bright quasar phase and possibly a large scale outflow, follows the general prescription for quasar activity proposed by Hopkins and collaborators \\citep{Springel05,Cox06,hhc+07}.   \\begin{figure} \\begin{center} \\includegraphics[width=6.8in]{f8.eps} \\end{center} \\caption{ Apparent column density $N_a$ profiles for the NV~1242 transition, and the \\ion{C}{4} and \\ion{O}{6} doublets for velocities relative to $z_{\\rm fg} = 2.4360$. The profiles are significantly contaminated by coincident line-blending which is best revealed by comparing the $N_a$ profiles for the pair of transitions in each doublet.  After accounting for this line-blending, we conclude that there is weak (or negligible) \\ion{N}{5} and \\ion{O}{6} absorption in subsystem~C and at $\\delta v < 0 \\mkms$. Values and upper limits to the integrated ionic column densities are given in Table~\\ref{tab:high}. } \\label{fig:highion} \\end{figure}   \\begin{deluxetable}{lcccccccc} \\tablewidth{0pc} \\tablecaption{HIGH-ION COLUMN DENSITIES \\label{tab:high}} \\tabletypesize{\\footnotesize} \\tablehead{\\colhead{Ion} & \\colhead{\\fgqso$^a$} & \\colhead{Subsystem A$^b$} & \\colhead{Subsystem B$^c$} & \\colhead{Subsystem C$^d$}} \\startdata $\\log \\N{C^+}$ &$12.7\\pm 0.11$&$13.4\\pm 0.05$&$14.0\\pm 0.05$&$13.3\\pm 0.06$\\\\ $\\log \\N{N^{+4}}$ &$<13.9$&$<13.3$&$<13.8$&$<11.9$\\\\ $\\log \\N{O^{+5}}$ &$<14.9$&$<14.4$&$<14.8$&$<14.0$\\\\ \\enddata \\tablenotetext{a}{Integrated AODM column densities for the interval containing \\fgqso: $-200 \\mkms < \\delta v < +50 \\mkms$ relative to $z_{fg}=2.4360$.} \\tablenotetext{b}{Integrated AODM column densities for Subsystem~A: $+50 \\mkms < \\delta v < +200 \\mkms$ relative to $z_{fg}=2.4360$.} \\tablenotetext{b}{Integrated AODM column densities for Subsystem~B: $+200 \\mkms < \\delta v < +600 \\mkms$ relative to $z_{fg}=2.4360$.} \\tablenotetext{d}{Integrated AODM column densities for Subsystem~C: $+600 \\mkms < \\delta v < +780 \\mkms$ relative to $z_{fg}=2.4360$.} \\end{deluxetable}    \\subsection{Constraints from High-Ion Observations} \\label{sec:high}  The current paradigm of baryons within galactic halos envisions a hot, diffuse medium that has been heated by shocks during virialization and/or feedback from supernovae within the galaxy.  Embedded within this diffuse medium may be cooler, dense clouds that may be photoionized by the local or background UV radiation field.  In the Galaxy, this simple picture is supported by the observations of the high velocity clouds \\citep[cool, dense clumps, e.g.][]{wv97,pds+02} and the widespread detection of \\ion{O}{6} absorption which traces a hotter, diffuse medium \\citep{sws+03}.  This model is also supported by quasar absorption line surveys, i.e.\\ the association of \\ion{Mg}{2}, \\ion{C}{4}, and \\ion{O}{6} gas with galactic halos \\citep[e.g.][]{s93,clw01,fpl+07}.  On theoretical grounds, it has been argued that a hot diffuse component pressure confines a population of cold clouds, which are unlikely to be massive enough to be self-gravitating \\citep{mm96,mb04}.   With these scenarios in mind, we are motivated to search for absorption from a diffuse, hot component that could be associated with the halo of \\fgqso.  The sightline to \\bgqso\\ intersects the presumed galactic halo of \\fgqso\\ at an impact parameter $\\rperp = 108$\\,kpc. As we will see in \\S~\\ref{sec:NFW}, the virial radius of the dark matter halo expected to host \\fgqso\\ is $r_{\\rm vir} \\simeq 250\\,{\\rm kpc}$, so that our background sightline easily resolves the virial radius.  In $\\S$~\\ref{sec:ion}, we discussed the detection and analysis of low and intermediate-ion transitions observed in \\slls.  We also commented on the general absence of strong absorption by high-ions like O$^{+5}$, C$^{+3}$, and N$^{+4}$.  These ions trace gas either with a large ionization parameter ($\\log U > -2$) or with a high temperature $(T > 10^5$\\,K).  Their observed column densities, therefore, constrain the nature of a diffuse and/or hotter phase within the halo of \\fgqso.  In Figure~\\ref{fig:highion}, we present the apparent column densities, $N_a^i \\equiv 10^{14.5761} \\ln(1/I^i) / (f\\lambda)$, of the \\ion{C}{4}, \\ion{N}{5} and \\ion{O}{6} profiles (smoothed by 5 pixels) for the velocity interval spanning the quasar redshift and the three subsystems of \\slls.  The C$^{+3}$ and O$^{+5}$ ions exhibit multiple transitions which allow one to identify line-blending by visual inspection, i.e.\\ regions where the two profiles diverge significantly.  The \\ion{C}{4}~1548 and 1550 profiles track each other closely for $\\delta v > +200 \\mkms$ indicating a positive detection without blending at these velocities. At $\\delta v < +200 \\mkms$, the profiles diverge because the \\ion{C}{4}~1548 profile of component C blends into the \\ion{C}{4}~1550 profile of component A (see also Fig~\\ref{fig:metals}).  Therefore, the \\ion{C}{4}~1548 profile sets an upper limit to the column density of C$^{+3}$ at these velocities.  At $\\delta v \\le 0\\mkms$, there is no statistically significant detection of \\ion{C}{4}.  The \\ion{N}{5} and \\ion{O}{6} doublets lie within the \\lya\\ forest and therefore are more likely to suffer from line blending.  Examining the \\ion{O}{6} transitions, the two profiles diverge at nearly all velocities.  This indicates that the line-profiles are severely affected by blends, most likely \\lya\\ and \\lyb\\ absorption from gas at lower redshift.  The only interval where the \\ion{O}{6} profiles roughly coincide is at $400 \\mkms < \\delta v < 550 \\mkms$.  One may optimistically interpret the absorption as a positive detection of O$^{+5}$ ions, but the more conservative approach is to report even this absorption as an upper limit.  Turning to \\ion{N}{5}, we present only the \\ion{N}{5}~1242 transition because the \\ion{N}{5}~1238 profile is heavily blended with a strong \\lya\\ system at $z=2.502$. We also expect that the broad absorption at $\\delta v < 200 \\mkms$ is associated with coincident \\lya\\ absorption; meanwhile, the optical depth in subsystem~C ($\\delta v < 600 \\mkms$) is not statistically significant.  The only plausible \\ion{N}{5}~absorption, therefore, is within subsystem~B where one notes the optical depth profile of \\ion{N}{5}~1242 resembles that of \\ion{C}{4} in the interval $300 \\mkms < \\delta v < 500 \\mkms$. We caution, however, that the \\ion{O}{6} profiles do not track \\ion{N}{5} in this velocity interval.   This draws into question a positive detection of \\ion{N}{5} absorption because the O$^{+5}$ and N$^{+4}$ ions trace gas with similar physical properties and one generally expects coincident \\ion{N}{5} and \\ion{O}{6} absorption. It is possible, therefore, that the apparent \\ion{N}{5} profile instead tracks an unidentified blend and we report the measured optical depth as an upper limit to $\\N{N^{+4}}$ in subsystem~B.  Table~\\ref{tab:high} summarizes the column densities of the high-ions in subsystems A, B, and C and in the velocity interval containing $z = z_{\\rm fg}$.  For C$^{+3}$ and N$^{+4}$, the column densities are lower than 10$^{14} \\cm{-2}$ and fall below $10^{13} \\cm{-2}$ in several regions. For O$^{+5}$, the upper limits are a factor of ten higher owing to noisier data and a higher frequency of line-blending. All of these values lie below the $\\gtrsim 10^{14} \\cm{-2}$ column densities measured for the low-ions of CNO in \\slls\\ (Table~\\ref{tab:colm}).  This has significant implications for the physical conditions of diffuse gas associated with the galactic halo of \\fgqso.   \\begin{figure} \\begin{center} \\includegraphics[width=6.0in]{f9.eps} \\end{center} \\caption{ Upper panel: Predicted \\ion{H}{1} column densities for a plane-parallel slab lying at a distance of 100\\,kpc from \\fgqso\\ for a range of total H column ($\\log N_H = 19, 19.5, ..., 21$) and as a function of volume density $n_{\\rm H}$. As $n_{\\rm H}$ decreases, the effective ionization parameter increases and the gas eventually becomes optically thin (i.e.\\ $\\mnhi < 10^{17.2} \\cm{-2}$).  Lower panel: Predicted column densities of several metal ions for a plane parallel slab with $N_{\\rm H} = 10^{20.5} \\cm{-2}$ illuminated by \\fgqso\\ at a distance of 100\\,kpc.  As the assumed volume density of the gas is lowered, the gas is more photoionized and CNO are pushed into the high-ion states of C$^{+3}$, N$^{+4}$, and O$^{+5}$.  The predicted column densities for a metallicity of 1/3 solar are high and easily detectable even in a low-resolution spectrum (e.g.\\ SDSS).  These predicted column densities are to be compared against the observed values (Figure~\\ref{fig:highion}, Table~\\ref{tab:high}).  } \\label{fig:evap} \\end{figure}     \\begin{figure} \\begin{center} \\includegraphics[height=6.8in,angle=90]{f10.eps} \\end{center} \\caption{ Predicted column densities for high-ion states of CNO and Si assuming collisional ionization equilibrium (CIE) for a gas with total H column density $N_{\\rm H} = 10^{21.8} \\cm{-2}$ and 1/3 solar metallicity.  This hydrogen column density corresponds to the estimated baryonic surface density at a 100\\,kpc impact parameter from a dark matter halo with $M=10^{13} \\msol$ and an NFW density profile. The arrows on the curves indicate constraints to the temperature based on upper limits to the observed ionic column densities. One notes that the data rule out $3 \\sci{4} \\, \\rm K < T_{CIE} < 2 \\sci{6} \\, \\rm K$. } \\label{fig:cie} \\end{figure}   A direct conclusion from these observations is that the sightline does not intersect a large column density of optically thin, photoionized material.  In the lower panel of Figure~\\ref{fig:evap}, we present the predicted column densities of C$^{+3}$, N$^{+4}$, and O$^{+5}$ ions for a wide range of physical density assuming an absorber with total hydrogen column density $N_{\\rm H} = 10^{20.5} \\cm{-2}$ and metallicity [M/H]~$=-0.5$\\,dex.  For these calculations, we assume the gas is illuminated by the \\fgqso\\ at a distance of $r = \\rperp$. Focusing on the parameter space where the gas is optically thin ($n_{\\rm H} < 0.3 \\cm{-3}$), one predicts column densities exceeding $10^{14.5} \\cm{-2}$ for all of the ions. Even if we assume a 10$\\times$ lower metallicity, we can rule out such a phase of gas along the sightline.   The absence of photoionized high-ions like O$^{+5}$ is not particularly surprising, however, because the ambient medium in the outer galactic halo surrounding a high-$z$ quasar almost certainly is at a temperature so high that collisional ionization dominates over photoionization.  This material may be observed independently of the gas contributing to \\slls.  A galaxy with mass $M > 10^{12} \\msol$, for example, has an implied virial temperature in excess of $10^{6}$\\,K and the dark matter halo expected to host \\fgqso would have a virial temperature $T_{\\rm vir}\\sim 10^{7}\\,{\\rm K}$ (see $\\S$~\\ref{sec:NFW}).  In Figure~\\ref{fig:cie} we present predicted column densities for C$^{+3}$, N$^{+4}$, and O$^{+5}$ ions assuming collisional ionization equilibrium \\citep[CIE;][]{sd93}, $N_H = 10^{21.8} \\cm{-2}$ and [M/H]~$= -0.5$\\,dex.  This H column density is the predicted value at $\\rperp = 108\\,$kpc in a $M = 10^{13.3} \\msol$ halo with a Navorro-Frank-White (NFW) profile (concentration $c = 4$) and assuming a gas fraction equal $f_g = 0.17$ equal to the cosmic baryon fraction \\citep[][see $\\S$~\\ref {sec:NFW}]{wmap05}.  The figure demonstrates that the gas must have $T> 10^{6}$\\,K to be consistent with the observed upper limits to $\\N{O^{+5}}$.  The dependence of $\\N{O^{+5}}$ on $T$ is so strong that even if we assume considerably lower $N_H$ and [M/H] values for the diffuse component, the gas must have $T > 10^6$\\,K.  A cooler phase is only allowed if $\\log N_H + {\\rm [M/H]} < 17.7$.  The principal conclusion of this section is that a diffuse medium in the halo of \\newline \\fgqso, must have a temperature exceeding 10$^6$\\,K.  This is consistent with the virial temperature $T_{\\rm vir}\\sim 10^{7}\\,{\\rm K}$ for the gas in the $M \\gtrsim 10^{13}\\msol$ dark matter halo that is expected to host the f/g quasar (see $\\S$~\\ref {sec:NFW}).  Note however that the absence of strong high-ion transitions associated with \\fgqso\\ contrasts with similar observations (using a b/g quasar sightline) of $z\\sim 2.3$ star-forming galaxies where strong \\ion{O}{6} and \\ion{N}{5} absorption is detected at similar impact parameters \\citep{simcoe06}.  In addition, the damped \\lya\\ systems (DLAs) also frequently exhibit \\ion{O}{6} absorption which reflects a highly ionized, diffuse medium in the halos of these galaxies \\citep{fpl+07}. The average $\\N{O^{+5}}$ for the DLAs exceeds $10^{14} \\cm{-2}$.  The DLA sightlines presumably intersect galactic halos at much smaller impact parameters than $\\rperp = 108$\\,kpc.  Nevertheless, the high-ion material in DLAs is undoubtedly associated with the galactic halo \\citep{wp00a,mps+03,pcw+08} and may extend to many tens kpc.  The fundamental difference is that DLA sightlines will rarely penetrate the galaxies hosting high $z$ quasars and likely trace lower mass halos with significantly lower virial temperatures.    \\subsection{Is \\fgqso\\ Shining on \\slls?} \\label{sec:shining}  Of considerable interest to studies of quasars and the buildup of supermassive black holes in the Universe \\citep[e.g.][]{yu02,hoprich07,shankar07} is whether their optical-UV emission is isotropic and or intermittent.  Evidence that quasars emit anisotropically or intermittently comes from the null detections of the transverse proximity effect in the Ly$\\alpha$ forests of projected quasar pairs \\citep[][but see Jakobsen et al. 2003]{Crotts89,DB91,FS95,LS01,Schirber04,Croft04}, the anisotropic clustering of optically thick absorbers near quasars \\citep{hp07,phh08}, and the absence of fluorescence detections in optically thick gas proximate to quasars \\citep[Hennawi \\& Prochaska 2008, in prep.; but see][]{ask+06}. Together, these studies provide a statistical case that quasars emit anisotropically or intermittently.  For an individual quasar, one can use a background sightline to observe nearby gas in absorption, and test whether it is illuminated by studying its ionization state. The general approach is to (1) derive the ionization parameter $U$ of the gas near the quasar using ionic metal line transitions (2) estimate the volume density of the gas, and (3) compare the inferred intensity of the impingent radiation field with that estimated for the nearby quasar.  \\cite{wfw+07} and \\cite{gsp08} have conducted this test for highly ionized gas near several $z \\sim 2$ quasars.  The \\cite{wfw+07} study was deemed inconclusive.  \\cite{gsp08}, meanwhile, identified several absorbers proximate to quasars with high ionization parameters and estimated an ionizing radiation field that exceeds the EUVB.  Instead, the estimated fluxes roughly match the values calculated for the nearby quasars under the assumption that they emit isotropically.  These authors, however, did not place empirical constraints on the gas density, did not consider the effects of collisional ionization, and did not explore the effects of non-solar relative abundances.  In these respects, we consider their results to be suggestive but inconclusive.   We can similarly test for illumination of \\slls\\ by \\newline \\fgqso. In $\\S$~\\ref{sec:ion}, we compared the observed ratios of N, Si, Al, Fe, and C atoms and ions against photoionization models assuming a single phase, plane-parallel slab of gas illuminated by a quasar spectrum ($f_\\nu \\propto \\nu^{-1.57}$). For subsystems A+C of \\newline \\slls, the majority of the observational constraints imply $\\log U < -3.25$.  The only exception is the large N$^+$/N$^0$ ratio which indicates $\\log U \\approx -3$ for subsystem~A,  and $\\log U > -4$ for subsystem~C. Under the constraints of a single-phase model and a plane-parallel geometry, one cannot identify an equilibrium photoionization model that reproduces all of the observed ionic ratios. By the same token, the observations are inconsistent with any collisional ionization equilibrium model.   We do not know if these inconsistencies highlight errors in the atomic data (e.g.\\ dielectronic recombination rates), the oversimplification of a single-phase, plane-parallel geometry, an erroneous shape for the radiation field, and/or the presumption of equilibrium models.  We will proceed under the relatively conservative assumption that the ionization parameter $\\log U < -3$ and we will comment on implications for significantly smaller values. We also remind the reader that the observed ionic ratios are roughly consistent with a $T \\approx 20,000$K collisional ionization equilibrium model.  Any contribution of collisional ionization processes to the observations would lower the estimated flux of ionizing radiation.   The other key physical constraint is the volume density of the gas $n_{\\rm H}$.  For a predominantly ionized gas, this can be estimated from the electron density $n_{\\rm e}$.  Our analysis of the fine-structure levels of the C$^+$ and Si$^+$ ions indicate $n_{\\rm e}$ less than approximately $5 \\cm{-3}$ for each of the subsystems. These upper limits could be improved by obtaining higher S/N observations of the undetected \\ion{Si}{2}$^*$~1264 transition at $\\lambda \\approx 4350$\\AA.  In any case, we will adopt this conservative limit to the electron density and assume $n_{\\rm H} \\approx n_{\\rm e}$ which is appropriate for a predominantly ionized medium.  Adopting $\\log U < -3$ and $n_{\\rm H} < 5 \\cm{-3}$, we infer an upper limit to the ionizing photon flux,  \\begin{equation} \\Phi_{\\rm obs} = U n_{\\rm H} c \\; < \\; 1.5\\sci{8} \\;\\; {\\rm photons \\; s^{-1} \\; cm^{-2}} \\;\\;\\; . \\label{eqn:phiobs} \\end{equation} We may compare this value with the ionizing flux of \\fgqso\\ under the assumptions that the gas lies at a distance equal to the impact parameter and that the quasar emits isotropically. Combining the SDSS optical photometry of \\fgqso\\ with an assumed power-law spectral shape ($f_\\nu \\propto \\nu^{-1.57}$), we derive an AB magnitude $m_{912} = 21.42$ at 1\\,Ryd.   At $z_{\\rm fg} = 2.436$, this gives a specific luminosity $L_{912} = 1.16\\sci{30} {\\rm erg \\; s^{-1} \\; cm^{-2} \\; Hz^{-1}}$ and we estimate an ionizing photon flux at $r = \\rperp = 108$\\,kpc of  \\begin{equation} \\Phi_{\\rm QSO} = 9\\sci{7} \\; {\\rm photons \\; s^{-1} \\; cm^{-2}} \\;\\;\\; , \\label{eqn:phiqso} \\end{equation} assuming the quasar emits isotropically. Therefore, the conservative upper limit to $\\Phi_{\\rm obs}$ roughly matches the predicted photon flux from \\fgqso\\ at the observed impact parameter.  In this conservative observational limit, we do not have strong evidence for anisotropic emission.  If we adopt less conservative but more realistic constraints on the ionization parameter and volume density ($\\log U < -4; n_{\\rm H} < 1\\cm{-3}$), we set an upper limit to the observed photon flux $\\Phi'_{\\rm obs} < 3\\sci{6} \\;\\; {\\rm photons \\; s^{-1} \\; cm^{-2}}$. To maintain the null hypothesis that \\slls\\ is illuminated, we would require the gas be located at $r > 500$\\,kpc.  This large distance is disfavored by other arguments.  The extreme velocity field, high metallicity, and quasar-like relative abundances suggest the gas is local to the quasar host galaxy.  At $r=500$\\,kpc one could associate \\slls\\ with a protocluster containing \\fgqso, but not its galactic halo.  Furthermore, the strong clustering of optically thick absorbers around quasars at $z\\sim 2.5$ implies a high probability that $r \\approx \\rperp$ \\citep{hp07}.  In Figure~\\ref{fig:cluster}, we show the probability of the absorber lying at a distance less than $r$ from the f/g quasar, given the knowledge that it is within a velocity interval $v_{\\rm vmax} = \\pm 700 \\mkms$ and has an impact parameter $\\rperp = 108$\\,kpc.  The upper x-axis of the figure shows the corresponding cumulative probability of the ionizing flux $\\phi_{\\rm QSO}$, assuming that \\fgqso emits isotropically.  We calculate a 70\\% probability that \\slls\\ is located within 500\\,kpc of the quasar.  We independently test whether the transverse direction is illuminated by searching for \\lya\\ fluorescence from the optically thick gas near the b/g quasar sightline. By coadding a series of longslit exposures of \\sdssj\\ using Gemini/GMOS, amounting to a 7200s total integration, we have not detected any extended emission \\citep[Hennawi \\& Prochaska 2008, in prep.; but see][]{ask+06} at the f/g quasar redshift near the impact parameter of the b/g quasar sightline. The gas clouds comprising \\slls\\ have estimated sizes of $\\sim 10-100\\,{\\rm pc}$, making them so small that it would be impossible (with current sensitivity) to detect \\lya\\ emission from the \\emph{actual} clouds we see in absorption in \\slls.  However, the high covering factor of optically thick gas near foreground quasars with $\\rperp \\lesssim 100\\,{\\rm kpc}$ characterized statistically by \\citet{hpb+06}, together with the detection of three such optically thick subsystems in \\slls, argue that the covering factor of optically thick gas is indeed roughly unity at $\\rperp$ near \\fgqso. If this gas is illuminated it should be emitting detectable \\lya\\ photons. The fluorescent surface brightness of clouds illuminated by \\fgqso\\ at a distance of $\\rperp = 108\\,{\\rm kpc}$ would be \\newline ${\\rm SB}_{\\rm Ly\\alpha} < 4.8 \\sci{-17} {\\rm ergs \\, s^{-1} \\, \\cm{-2} \\, arcsec^{-2}}$. From a preliminary analysis, we conservatively estimate that we could have detected extended emission a factor of ten fainter than this expectation, implying that the distance to the clouds would have to be $r > \\sqrt{10}\\,\\rperp$ to render them undetectable if the foreground quasar emits isotropically. Thus, isotropic emission from \\fgqso\\ is ruled out unless all of the optically thick gas lies beyond $r > 340\\,{\\rm kpc}$.   \\subsection{Why Aren't Similar Absorbers Observed along the Line of Sight to Quasars?} \\label{sec:nal}  \\cite{hp07} have demonstrated that high-$z$ quasars exhibit a high covering fraction to optically thick absorbers at impact parameters $\\rperp \\lesssim 200$\\,kpc, and argued that the incidence of such transverse absorption exceeds the incidence along the line-of-sight to bright high-$z$ quasars.  \\cite{phh08} similarly concluded that there are fewer damped \\lya\\ systems near $z>2$ quasars then expected from clustering arguments and \\cite{wkw+08} have recently reported a similar deficit of strong \\ion{Mg}{2} absorbers along the line-of-sight compared to the transverse direction.  Together these observations argue that quasars emit ionizing radiation anisotropically or intermittently. Hence a paucity of optically thick absorbers are detected along the line-of-sight because this gas, which is by construction illuminated, is photoionized by the quasar.  One test of this hypothesis is to search for an enhancement of highly ionized material along the sightline, e.g.\\ \\ion{Si}{4}, \\ion{C}{4}, \\ion{N}{5}, \\ion{O}{6}. Indeed, \\cite{wkw+08} report a higher incidence of \\ion{C}{4} absorption along the same sightlines where the lower ionization potential transition, \\ion{Mg}{2}, is deficient. There is also a high incidence of highly ionized absorbers within several thousand \\kms\\ of the quasar redshift commonly termed narrow associated absorption line systems \\citep[NAAL;][]{dcr+04}.  We now explore the physical conditions required for \\fgqso\\ to photoionize/photoevaporate \\slls\\ and also the expected properties of \\slls\\ if it were illuminated by the quasar, and thus exposed to a more intense radiation field.  These calculations are similar to those performed by \\cite{cmb+08} and we refer the reader to their work for additional analysis.  In Figure~\\ref{fig:evap} we present a series of calculations from the Cloudy software package for constant density gas slabs placed 100\\,kpc from \\fgqso\\ along the sightline to Earth.  The quasar's specific luminosity was assumed to be a power-law ($L_\\nu \\propto \\nu^{-1.57}$) and was normalized to SDSS photometry: $L_\\nu = L_{912} (\\nu/\\nu_{912})^{-1.57}$ with $L_{912} = 1.16 \\sci{30} {\\rm erg \\, s^{-1} \\, Hz^{-1}}$.  In all of the calculations we assume a gas metallicity [M/H]~$=-0.5$\\,dex with solar relative abundances.  In the upper panel of the Figure, we show the \\ion{H}{1} column densities for slabs with a series of total H column densities\\footnote{These calculations are nearly self-similar with respect to the ionization parameter $U$, i.e.\\ we derive similar results if we modify the distance to the slab, the quasar luminosity and $n_{\\rm H}$ provided that $U$ remains constant.}  $\\log N_H = 19, 19.5, ..., 21$ as a function of the gas density $n_{\\rm H}$.  For slabs with $N_H = 10^{19} \\cm{-2}$, the gas has $\\mnhi < 10^{17.2}$ and is hence optically thin to ionizing photons, provided $n_{\\rm H} < 10 \\cm{-3}$; while slabs with $N_H = 10^{21} \\cm{-2}$ are optically thin only if $n_{\\rm H} < 0.1 \\cm{-3}$.  In $\\S$~\\ref{sec:ion}, we estimated $N_H \\approx 10^{20} \\cm{-2}$ and $n_{\\rm H} \\approx 1 \\cm{-3}$ for subsystems A and C based on the relative column densities of several ion pairs, our estimate of the \\ion{H}{1} column density (uncertain by $\\pm 0.4$\\,dex), and an analysis of \\ion{C}{2} and \\ion{Si}{2} fine-structure transitions.  According to the calculations presented in Figure~\\ref{fig:evap}, this gas would have $\\mnhi \\approx 10^{17} \\cm{-2}$ if it were at 100\\,kpc from \\fgqso\\ along our sightline.  If the volume density is just a little lower (permitted by the fine-structure analysis), the gas would be optically thin.  We conclude that subsystems~A and C of \\slls\\ would be photoionized by \\fgqso\\ if it were placed along our sightline.  Subsystem B, however, exhibits a larger \\nhi\\ value and has a larger implied $N_H \\approx 10^{20.3} \\cm{-2}$ value.  Assuming $n_{\\rm H} = 1 \\cm{-3}$, this system would have $\\mnhi \\approx 10^{17.3} \\cm{-2}$ if illuminated by the \\fgqso\\ at a distance of 100\\,kpc. To be optically thin and illuminated at this distance requires $n_H < 1 \\cm{-3}$. It is important to note that \\fgqso\\ is relatively faint compared to the brighter quasars studied to establish the low incidence of optically thick absorbers near quasars \\citep[e.g.][]{reb06,phh08,wkw+08}. These quasars are a factor of $\\sim 5-10$ brighter and the resulting ionization parameters are higher by the same factor.  Thus clouds with properties similar to \\slls\\ could still appear optically thin if illuminated, resulting in no conflict with previous observations. Similarly, it would be very fruitful to characterize the incidence of proximate optically thick absorbers as a function of quasar luminosity.   If an absorber like \\slls\\ is photoionized by a nearby quasar, the system may still give rise to several discernible features.  If the \\ion{H}{1} column density exceeds $10^{14} \\cm{-2}$, it would exhibit a strong \\lya\\ absorption feature.  This would be difficult to distinguish from the ambient \\lya\\ forest, however.  If this gas has even a modest metallicity, it may show unique metal-line signatures. The lower panel of Figure~\\ref{fig:evap} presents the predicted column densities of several high-ions assuming $N_H = 10^{20.5} \\cm{-2}$, [M/H]~$=-0.5$\\,dex, and that the cloud is at a distance of 100\\,kpc from \\fgqso\\ along our sightline.  Once the cloud becomes optically thin (at $n_{\\rm H} \\lesssim 0.5 \\cm{-3}$), large column densities of C$^{+3}$, N$^{+4}$, and O$^{+5}$ ions are expected.  The alkali doublets associated with these ions saturate at column densities of $\\approx 10^{14} \\cm{-2}$ and, therefore, would yield absorption lines that could be detected at even the modest resolution of the SDSS spectroscopic survey.  These column densities scale directly with the assumed metallicity and are also sensitive to the $N_{\\rm H}$ value. Combining the high incidence of optically thick gas observed in the transverse direction with the high ionization parameter along the sightline, and using our knowledge of the physical properties of the (unilluminated) absorbing clouds from \\slls, we conclude that that a large fraction of bright quasars should exhibit strong \\ion{C}{4}, \\ion{N}{5}, and \\ion{O}{6} absorption at $z \\approx z_{\\rm em}$ along their sightlines.  Absorption systems with strong high-ion absorption near quasars with $z_{\\rm abs} \\approx z_{\\rm em}$ or small velocity intervals ($\\Delta v < 1000 \\mkms$) have been detected in quasar spectra, and are referred to as narrow associated absorption line (NAAL) systems \\citep[e.g.][]{wbp93,sp00}.  The most comprehensive survey of these NAALs using high resolution spectra was performed by \\cite{dcr+04} who analyzed the UVES key project of $z \\approx 2$ quasars.  Of the 22 quasars in the sample, \\cite{dcr+04} report that 16 exhibit significant \\ion{C}{4} absorption within $\\pm 5000 \\mkms$ of $z_{\\rm em}$ for a total of 34 \\ion{C}{4} systems, 15 of which also exhibit \\ion{N}{5} absorption.  \\cite{dcr+04} studied the properties of a subset which allowed for detailed photoionization modeling.  The characteristics of this subset are similar in metallicity, density, and relative abundances to \\slls\\ (see also $\\S$~\\ref{sec:chem}). Furthermore, the authors used observed limits on the volume density of the gas (from fine-structure analysis) and estimates of partial coverage to infer distances of 10 to 200\\,kpc from the quasar.  We conclude, therefore, that if \\slls\\ experienced an $\\approx 10$ times higher ionization parameter, it would show characteristics common to NAAL systems.  Indeed, \\fgqso\\ is approximately 2.5\\,mag (10$\\times$) fainter than the quasars of the \\cite{dcr+04} sample.  We speculate that an absorption-line survey of intrinsically faint quasars will show a high incidence of absorbers with properties similar to \\slls\\ and a correspondingly lower incidence of NAAL systems.  One could perform such a survey with current and future echellette spectrometers on 10m-class telescopes.   \\label{sec:summ}  In this paper, we introduced a novel technique to study the physical state of gas in the halos of luminous quasars, which has the potential to provide powerful constraints on the physical processes governing the formation of massive galaxies.  By mining the sky for very rare close associations of quasars \\citep{thesis,hso+06,hpb+06}, we previously discovered \\sdssj, a close projected quasar pair with angular separation $\\theta = 13.3''$ corresponding to $\\rperp =108$\\,kpc at the redshift of the foreground quasar $z_{\\rm fg}= 2.4360 \\pm 0.0005$, precisely determined from Gemini/GNIRS near-IR spectroscopy. The spectral and photometric properties of \\fgqso\\ make it an unremarkable quasar at $z \\sim 2.4$. It has an SDSS $i$-band magnitude of $i=20.5$, from which we estimate a bolometric luminosity of $L_{\\rm QSO} \\simeq 1.4\\sci{46}{\\rm erg\\, s^{-1}}$, corresponding to a supermassive black hole mass $M_{\\rm BH} \\simeq 1.1\\sci{9}\\left(\\frac{f_{\\rm Edd}}{0.1}\\right)^{-1}\\,\\msol$, if the black hole accretes at one tenth of the Eddington limit.  The luminosity of \\fgqso\\ places it near the `knee' of the $z\\sim 2.5$ quasar luminosity function \\citep{croom04,richards06}, and the clustering of such quasars indicates they inhabit massive dark halos $M \\sim 10^{13.3}\\msol$ \\citep{croom01,pmn04,croom05}, making them the progenitors of massive red-and-dead galaxies observed today. Our impact parameter $\\rperp$ easily resolves the expected virial radius $r_{vir} = 250\\,{\\rm kpc} (M/10^{13.3} \\msol)^{1/3}$, and pierces the halo of \\fgqso\\ at about the cooling radius, where gas shock heated to the virial temperature should take about a Hubble time to cool.  Rather, the only remarkable thing about \\fgqso\\ is that it forms a close projection with a b/g quasar ($z_{\\rm bg} = 2.53$) which is bright enough ($r=19.0$) for high resolution spectroscopy. Our Keck HIRES Echelle spectrum of \\bgqso, the first ever to probe the halo gas of a f/g quasar, resolves the velocity fields of the absorbing gas and allows us to measure precise column densities for \\ion{H}{1} and the ionic transitions of metals like Si, C, N, O, and Fe. These measurements allow us to place constraints on the physical state of the gas near the f/g quasar, such as its kinematics, temperature, ionization structure, chemical enrichment patterns, volume density, the size of the absorbers, the intensity of the impingent radiation field, as well as test for the presence of hot collisionally ionized gas. We first summarize the results for our single sightline, and then discuss their implications for massive galaxy formation and the quasar phenomenon.   \\subsection{Model Independent Constraints}  In our Keck/HIRES spectrum of \\bgqso, we identify a super Lyman limit system (SLLS) with a redshift $z=2.44$.  A Voigt profile fit to the \\lya\\ and \\lyb\\ profiles imply a total \\ion{H}{1} column density $\\mnhi = 10^{19.65 \\pm 0.15} \\cm{-2}$.  The \\ion{H}{1} absorption occurs redward of $z_{\\rm fg}$, spanning from $\\delta v = +50 \\mkms$ to $+780 \\mkms$.  There is no \\ion{H}{1} absorption detected ($\\mnhi < 10^{13.5} \\cm{-2}$) in the velocity interval $\\delta v = -1500 \\mkms$ to 0\\kms.  We identify a series of metal lines coincident in velocity with the \\ion{H}{1} absorption (distributed in three primary components) that includes transitions from a range of ionization states of C, O, Fe, Si, Al, and N.  We observe an integrated O$^0$/H$^0$ ratio of $\\log[\\N{O^0}/\\N{H^0}] = -3.9$\\,dex which indicates a highly enriched gas.  Ignoring ionization corrections to this ratio, which should be negligible, the average gas metallicity is [O/H]~$= -0.5$\\,dex.  Both the extreme kinematics and high metallicity of this system are highly uncommon for intervening SLLSs.  The Doppler parameters of the low and intermediate ions are small $(b \\lesssim 5\\mkms)$ indicating a gas temperature $T \\simeq 20,000$K. At all velocities we observe an ionic ratio N$^+$/N$^0 > 1$ which implies the gas is predominantly ionized.  We observe weak, (if any) \\ion{C}{4}, \\ion{N}{5}, and \\ion{O}{6} absorption indicating a modest (or negligible) quantity of metal-enriched gas with $T \\approx 10^5$ to 10$^6$K in the halo surrounding \\fgqso.    \\subsection{Model Dependent Constraints}  Under the assumption that the gas is photoionized by a hard radiation field (e.g.\\ the nearby quasar or the EUVB), we estimate an ionization parameter $\\log U = -3.0 \\pm 0.3$ for the gas.  This implies an ionization fraction $x = 0.96$ for the two lower column density components, and $x=0.2$ for the component with the largest column. The total implied hydrogen column density of the system is $N_H = 10^{20.6} \\cm{-2}$.  Applying ionization corrections to the observed N$^0$/O$^0$ ratio, we infer solar to super-solar N/O abundances.  At high $z$, such large N/O values are only found in gas associated with quasar environments.  At $\\delta v \\approx +700 \\mkms$, we detect absorption from the excited fine-structure level of C$^+$, \\ion{C}{2}*~1335.  Under the assumptions that electron collisions dominate the C$^+$ excitation and that $\\N{C^+} = \\N{Si^+} + \\log(\\rm C/Si)_\\odot$, we estimate an electron density $n_{\\rm e} = 0.2$ to $10 \\cm{-3}$.  By comparing this volume density to the ionization corrected column density, we determine characteristic sizes $\\ell \\sim 10$ to 100\\,pc for the absorbing `clouds'.  The properties of the SLLS toward \\bgqso\\ -- $z \\approx z_{\\rm fg}$, extreme kinematics, high metallicity, and a solar N/O ratio -- argue that this gas is located within the halo of \\fgqso.  There is an additional statistical argument: the strong clustering of optically thick absorbers around $z \\sim 2$ quasars \\citep{hp07} implies the SLLS is likely to lie at a distance near the observed transverse distance (see Figure~\\ref{fig:cluster}).  We summarize in the next subsection several interpretations drawn from associating \\slls\\ with the galactic halo of \\fgqso.  We tested the hypothesis that the \\slls\\ is illuminated by ionizing radiation from f/g quasar by comparing the radiation field intensity inferred from analysis of the SLLS against the value expected from \\fgqso.  Adopting the most conservative parameters for the SLLS ($\\log U = -3$ and $n_{\\rm H} = 5 \\cm{-3}$), illumination of the SLLS cannot be ruled out if it is located at $r = \\rperp = 108$\\,kpc from \\fgqso. Adopting more realistic values ($\\log U = -4$ and $n_{\\rm H} < 1 \\cm{-3}$), we conclude that \\slls\\ is not illuminated by \\fgqso, unless the SLLS is located at an unlikely distance of greater than 500\\,kpc.  We further demonstrated that if a system like \\slls\\ were located along the sightline to a bright quasar at $r < 100$\\,kpc, and hence illuminated, then it would exhibit properties characteristic of the narrow associated absorption line (NAAL) systems.  This would include strong absorption by high-ion transitions of \\ion{O}{6}, \\ion{N}{5}, and \\ion{C}{4}, optically thin \\ion{H}{1} absorption, and solar chemical abundances.  The absence of strong high-ion absorption in this SLLS further suggests it is not illuminated by \\fgqso\\ and it also indicates that there is not a large reservoir of warm ($T\\approx 10^5$ to 10$^6$K) gas in the halo surrounding the quasar.  We suggest that there may still be a diffuse shock heated medium but that it has a high temperature ($T>10^6$K), characteristic of the virial temperature of the massive dark matter halo $M \\gtrsim 10^{13}\\msol$ expected to be hosting the f/g quasar. If it exists, this material undetectable in absorption with the transitions accessible with our b/g quasar spectrum.  \\subsection{Interpretation and Outlook}  The covering factor of cold $T\\sim 10^4{\\rm K}$, neutral, optically thick gas is nearly unity for transverse sightlines to $z \\sim 2$ quasars with $\\rperp \\lesssim 100\\,{\\rm kpc}$ \\citep{hpb+06}, although this high column density absorbing gas is rarely observed along the line-of-sight to individual quasars \\citep{reb06,phh08}. The explanation for this anisotropic absorption is that the transverse direction is much less likely to be illuminated by the f/g quasar ionizing flux, either because of obscuration or intermittent emission \\citep{hpb+06,hp07,phh08}. Besides this anisotropic absorption pattern, several independent lines of evidence corroborate this picture for the case of \\sdssj. First, we modeled the ionization state of \\slls\\, which indicates that the gas is unlikely illuminated by \\fgqso\\ (see \\S~\\ref{sec:shining}). Second, we failed to detect high-ion transitions like \\ion{N}{5} or \\ion{O}{6} that should have been seen if the gas were highly ionized by a hard radiation field (see \\S~\\ref{sec:high}). Third, we showed that if gas clouds similar to those in \\slls\\ were at a similar distance along the line-of-sight (and hence illuminated), they would explain the population of much more highly ionized associated NAALs, commonly detected in quasar spectra (see \\S~\\ref{sec:nal}).  Fourth, fluorescent \\lya\\ emission is not observed near $\\rperp$ in \\sdssj\\ \\citep{qpq4}; whereas the clouds responsible for the high covering factor of optically thick gas \\citep{hpb+06} near quasars (three such clouds were detected in \\slls), would be emitting detectable \\lya\\ photons if illuminated \\citep{qpq4}. The upshot of this anisotropic illumination picture is that the b/g quasar sightline probes the physical state of gas near \\fgqso\\ which is \\emph{shadowed and hence unaltered by the intense ionizing radiation emitted by the f/g quasar}.  By combining the statistical covering factor measured by \\citet{hp07} with the total $N_{\\rm H}$ column determined for \\slls, we argued in \\S~\\ref{sec:cold} that the amount of cold gas at $r \\sim 100$\\,kpc is significant, amounting to $\\sim 3\\% (\\frac{f_{\\rm g}}{\\Omega_{\\rm b}\\slash \\Omega_{\\rm m}})^{-1}$ of the total expected gas density of the f/g quasar's dark matter halo, if the gas fraction $f_{\\rm g}$ is equal to the cosmic baryon fraction $\\Omega_{\\rm b}\\slash \\Omega_{\\rm m}$.  Similarly, if one assumes the material is distributed in a thin shell of radius $\\rperp$ which subtends $\\Omega = 2\\pi$, the implied gas mass is $M \\sim 3\\sci{11}\\msol$. Although we have deduced much about the physical state of this cold gas (see Table~\\ref{tab:summ}), its origin is still unclear.  The biggest clue could lie in its extreme enrichment patterns. The nearly solar metallicity of \\slls\\ and its roughly solar N/O relative abundance make it anomalous relative to the population of intervening SLLS and DLAs (see Figure~\\ref{fig:relabnd}).  Indeed at $z \\sim 2.5$ such high metallicities have only been observed on kpc scales in starburst galaxies \\citep{pss+01,sas+03} and solar N/O has only been observed in the broad-line regions of quasars \\citep{dhs03,agk+07} or in the NAALs \\citep{dcr+04}.  How do we explain the presence of such a large cold gas mass of highly enriched material with a broad-line-region-like enrichment pattern at such a large distance from the quasar $r\\sim 100\\,{\\rm kpc}$?  It is tempting (and fashionable) to associate the absorbing material with a large scale outflow or galactic wind.  If we are observing dense material swept up by an outflow then the energetics are extreme ${\\dot E_{\\rm w}} \\sim 10^{45}\\,{\\rm erg\\,s^{-1}}$. A starburst cannot drive this wind unless we are willing to consider unprecedented star-formation rates ${\\dot M_{\\ast}}\\gtrsim 10^4\\,\\msol\\,{\\rm yr^{-1}}$. Radiation pressure feedback from the quasar cannot drive an outflow with this power even if its bolometric luminosity were completely absorbed by dust grains (along a different direction than our line-of-sight).  However, the power of the flow is only a few percent of \\fgqso's estimated bolometric luminosity, or similarly, its energy is of order a percent of the radiated rest-mass energy required to grow a $\\sim 10^9\\msol$ black hole. So the feedback hypothesis suggests a coupling between black hole accretion and outflow which is comparable to the value used by simulators to reproduce the supermassive black hole $M_{\\rm BH}-\\sigma$ correlation \\citep[e.g.][]{Springel05}. Furthermore, radiative losses from such an energetic flow could be significant, so our estimate of the coupling factor must be considered a lower limit.  If we are in fact observing feedback in \\slls, then the kinematics, radial extent, and high metallicity of the emergent outflow bear intriguing similarities to the giant \\lya\\ nebulae observed in high-z radio galaxies.  However, it is difficult to explain why such an energetic outflow would result in so little material at intermediate temperature $T \\sim 10^{4.5} - 10^{6}{\\rm K}$, which we would have easily detected (but did not, see \\S~\\ref{sec:high}).  But we may not be observing an outflow at all. The observed kinematics in \\newline \\slls, although extreme relative to the intervening SLLS population, are consistent with the expected gravitational motions if the f/g quasar is indeed hosted by a massive dark matter halo $M\\gtrsim 10^{13}\\,\\msol$, as indicated by the strong clustering of $z\\sim 2.5$ quasars \\citep{croom01,pmn04,croom05}. In this alternative scenario, the absorption is being caused by cold clouds, with sizes $r_{\\rm abs} = 10-100\\,{\\rm pc}$, unit covering factor and implied volume filling factor $C_{\\rm v} \\sim 10^{-5}-10^{-4}$, which are either infalling or virialized in the deep potential well of the massive dark matter halo.  We estimated the pressure of these clouds to be $P \\sim 10^{5}\\,{\\rm K\\,cm^{-3}}$, which intriguingly matches the pressure of the $T_{\\rm vir} \\sim 10^{7}\\,{\\rm K}$ shock-heated gas expected to permeate the massive dark matter halo hosting the f/g quasar. This pressure equilibrium is reminiscent of a large class of galaxy formation models that postulate cold $T \\sim 10^{4}{\\rm K}$ clouds pressure confined by a hot shock heated virialized medium \\citep{mm96,mb04,db06}, and these scenarios might explain the lack of significant intermediate temperature $T\\sim 10^{5-6}{\\rm K}$ gas near the f/g quasar.  However, if we are indeed detecting pressure confined cold clouds undergoing gravitational motions, why should these clouds have such a high metallicity? This question is all the more puzzling considering that the expected cooling time of the tenuous virialized hot gas would be comparable to the Hubble time at $r\\sim 100\\,{\\rm kpc}$. While it may be more plausible to associate the cold gas and metals with star-forming galaxies clustered around the quasar, they would need to be extremely abundant $n \\sim 0.1\\,{\\rm Mpc}^{-3}$ to produce the near unit covering factor of cold gas and metals.  These faint $L \\sim 10^{-3} L_\\ast$ dwarf galaxies would need to enrich spheres of $r_{\\rm abs}\\sim 20\\,$kpc to solar metallicity, which seems implausible in light of the low metallicities observed in most DLAs (see Figure~\\ref{fig:relabnd}) and the similarly low values seen locally in dwarf galaxies.  \\begin{figure} \\begin{center} \\includegraphics[width=6.0in]{f11.eps} \\end{center} \\caption{Comparison of absorption line properties of $z\\sim 2.5$ f/g quasars to $z\\sim 2.5$ star forming galaxies. The lower panel plots the \\ion{H}{1} column density versus impact parameter for absorbers, detected in a b/g quasar spectrum, coincident with the f/g quasar (or galaxy) redshift. The (red) circles are thirteen projected quasar pairs uncovered in \\citet{hpb+06} with $\\rperp < 350\\,{\\rm kpc}$, and (blue) squares are seven f/g $z\\sim 2.5$ star-forming galaxies studied in absorption against a bright b/g quasar by \\citet{simcoe06}. Downward pointing arrows are upper limits on \\nhi\\ for the f/g quasars, which reflects our inability to measure reliable column densities below $\\mnhi \\lesssim 18.5$ from the moderate resolution (FWHM $\\sim 125 \\mkms$) b/g spectra used in \\citet{hpb+06}. Because the \\citet{simcoe06} study used a high resolution b/g spectrum they do not suffer from this limitation. The gray shaded region indicates the range of abundances observed in the low density IGM.  } \\label{fig:simcoe} \\end{figure}    With just a single sightline, we cannot distinguish between a quasar powered outflow or cold clouds undergoing gravitational motions. However, similar observations of a statistical sample of projected quasar pairs stand to teach us a tremendous amount about the physics of massive galaxy formation.  This is nicely illustrated by the preliminary comparison of the absorption line properties of $z \\sim 2.2$ quasars to $z\\sim 2.3$ star-forming galaxies presented in Fig~\\ref{fig:simcoe}. The primary motivation for this comparison is that, across the two populations, we are afforded a mass baseline of more than an order of magnitude.  The clustering of star-forming galaxies at $z\\sim 2$ indicates that they inhabit dark matter halos with $M \\lesssim 10^{12}~\\msol$, making them the likely progenitors of $\\sim L_{\\ast}$ galaxies that inhabit the ``blue-cloud'' of the color magnitude diagram today \\citep{cst+07}; whereas, the stronger clustering of quasars at $z\\sim 2$ implies larger halo masses $M \\gtrsim 10^{13}~\\msol$ \\citep{croom01,pmn04,croom05}, making them progenitors of massive red-and-dead galaxies on the red sequence. Hence in comparing quasars to star-forming galaxies at $z\\sim 2$, \\emph{we are effectively comparing the progenitors of quenched galaxies to those of un-quenched galaxies.}   The lower panel of Figure~\\ref{fig:simcoe} plots the \\ion{H}{1} column density of absorbers at the f/g quasar (or galaxy) redshift versus the impact parameter to the b/g quasar sightline. The (red) circles are the thirteen f/g quasars ($\\langle z_{\\rm fg}\\rangle = 2.2$) with a background quasar within $\\rperp < 350\\,{\\rm kpc}$ uncovered by \\citet{hpb+06}, and (blue) squares are seven f/g star-forming galaxies ($\\langle z_{\\rm fg}\\rangle = 2.3$) studied in absorption against a bright b/g quasar by \\citet{simcoe06}. The vertical rectangles in the upper panel illustrate the range of metallicities encountered in the individual components of each system at each impact parameter. The only f/g quasar metallicity measured thus far is for \\sdssj\\ (this work), whereas \\citet{simcoe06} measured metallicities near six f/g galaxies.  Two notable features of Fig~\\ref{fig:simcoe} warrant further discussion. First, at $\\rperp \\simeq 100\\,{\\rm kpc}$, one galaxy and one quasar metallicity have been measured, and both indicate abundances near solar. \\citet{simcoe06} attributed the high metallicity of their smallest impact parameter system (MD103) to galaxy formation feedback, and we similarly cited the high metallicity of \\sdssj\\ as the most compelling argument for an outflow powered by the quasar. By mapping out the run of abundance with impact parameter in this plot, one could characterize the size of the enriched regions around protogalaxies with a significant mass lever-arm. If feedback is responsible for the metals at large impact parameters, measuring the abundance-impact parameter relation would constrain the energetics and transport processes characterizing the relevant feedback mechanism. Furthermore, this measurement would be fundamental to any discussion of the enrichment history of the IGM and the ICM.  The second noteworthy feature of Figure~\\ref{fig:simcoe} is that on small scales the f/g quasars appear to have \\nhi\\ column densities significantly larger than the f/g galaxies.  This comparison is complicated by the fact that the dark matter halos that host the quasars are expected to be larger and more massive than those hosting the galaxies. For reference, \\citet{cst+07} estimated that the average halo mass of $z\\sim 2$ star-forming galaxies to be $M = 10^{12}\\msol$, which implies a virial radius $r_{\\rm vir} = 89\\,{\\rm kpc}$; whereas, at $z\\sim 2.4$ we are using $M = 10^{13.3}\\msol$ implying $r_{\\rm vir} = 250\\,{\\rm kpc}$ \\citep{pmn04}.  At an impact parameter of $\\rperp = 100\\,{\\rm kpc}$, the ratio of the total hydrogen column densities of these dark matter halos is $N^{\\rm QSO}_{\\rm H}\\slash N^{\\rm gal}_{\\rm H} = 5$, whereas the quasars in Fig~\\ref{fig:simcoe} have \\ion{H}{1} column densities a factor $\\gtrsim 100$ times larger than the galaxies for $\\rperp \\lesssim 200\\,{\\rm kpc}$.  Although we caution that the statistics in the lower panel of Figure~\\ref{fig:simcoe} are still very poor, it is nevertheless interesting to speculate about the implications of this tentative result.  Why would the quasars have a much larger reservoir of cold gas at $\\rperp \\sim 100\\,{\\rm kpc}$, of order a few percent of the total expected gas density (see \\S~\\ref{sec:cold}), than star-forming galaxies? Can this difference be attributed to a distinct feedback mechanism operating in quasars that does not occur in star-forming galaxies?  If instead the absorbers at $\\rperp \\sim 100\\,{\\rm kpc}$ arise from cold gas accretion rather than outflows, then the larger relative density of cold gas around quasars is particularly puzzling. Cosmological hydrodynamical simulations predict the opposite trend with dark halo mass: cold accretion accounts for a smaller fraction of the total accreted gas in higher mass halos where shocks become stable throughout the halo and the cold filamentary mode of accretion disappears \\citep{kkw+05,db06,ocvirk08}. Whether the cold gas around quasars is ejected by feedback or accreted in gravitational collapse, what is the ultimate fate of this material? Does it eventually fall back onto the quasar host galaxy or do the gas and metals blend into the IGM. Do we expect to observe a similar cold halo of gas around nearby massive elliptical galaxies, since their progenitors are $z\\sim 2$ quasars like \\fgqso?  Although preliminary, the putative trends in Figure~\\ref{fig:simcoe} provoke fundamental questions about feedback, quenching and the physics of massive galaxy formation.  Yet this discussion represents just two measurements (metallicity and \\nhi) gleaned from our analysis of the spectrum of \\bgqso. Other properties of the gas such as its kinematics, temperature, ionization structure, relative abundance patterns, the presence or absence of a hot collisionally ionized phase, and the volume density and size of the clouds, have a similar potential to constrain the physics of massive galaxy formation if they can be mapped out for statistical samples. Such a statistical analysis is well within reach. To date, our pair confirmation program \\citep{thesis,hso+06,hpb+06} has confirmed about 90 pairs of quasars with impact parameter $R < 300~{\\rm kpc}$ and $z_{\\rm fg} > 1.6$, and about a comparable number still remain to be discovered in the existing SDSS photometric data.  Of the known pairs, only about five are bright enough $r\\lesssim 19$ for high resolution (FHWM$\\sim 10\\mkms$) echelle spectroscopy like the HIRES spectrum used to study \\slls. However, higher throughput but lower resolution echellette spectrographs, such as the Echellette Spectrograph and Imager \\citep[ESI,][]{ESI} on Keck II, the Magellan Echellette (MagE) on the Magellan Clay Telescope, or the planned X-Shooter spectrograph \\citep{xshooter} for the VLT, can deliver spectra with FHWM$\\sim 30-50\\mkms$ and the required signal to noise ratio ($\\sim 10$ per resolution element), in 1-2 hours for $r \\lesssim 21.5$, which would make \\emph{all} of the quasar pairs known observable.  Although these instruments have lower spectral resolution, they roughly resolve the \\ion{H}{1} Lyman series affording precise estimates of \\nhi. Furthermore, such data would yield metallicity and relative abundance estimates to a precision of 0.3\\,dex, and would allow for the construction of photoionization models with an accuracy comparable to that achieved in this work.  The most important aspect of the approach taken to understanding galaxy formation in this study and \\citet{simcoe06} (see also Adelberger et al.\\ 2003), is that they provide the first observational constraints on the physical state of the \\emph{gas} on scales $\\gtrsim$ kpc in high-$z$ protogalaxies. Although the ideas behind quenching and feedback were introduced to explain the observed properties of local galaxy stellar populations, such as the bimodality in the color magnitude diagram \\citep{Strateva01,baldry04,bell04,blanton05,faber07} and the sharp cutoff in the luminosity function \\citep{bbf+03}, these are fundamentally gas dynamics problems.  Many of the relevant hydrodynamical and physical processes can be (or are nearly) resolved by simulation grids with current technology. Conversely, predicting how this gas physics manifests itself in the stellar populations of local or high-$z$ galaxies requires that the uncertain ``sub-grid'' physics of star formation be inserted by hand, or with the aid of semi-analytical recipes.  Thus, observational constraints on the gas provide especially fruitful comparisons to theory. If the techniques presented here for a single object can be expanded to statistical samples, that will constitute a significant step on the road toward understanding massive galaxy formation."
    },
    "0806/0806.0265_arXiv.txt": {
        "abstract": "Damped Doppler shift oscillations have been observed in emission lines from ions formed at flare temperatures with the Solar Ultraviolet Measurements of Emitted Radiation spectrometer on the \\textit{Solar and Heliospheric Observatory} and with the Bragg Crystal Spectrometer on \\textit{Yohkoh}. This Letter reports the detection of low-amplitude damped oscillations in coronal emission lines formed at much lower temperatures observed with the EUV Imaging Spectrometer on the \\textit{Hinode} satellite. The oscillations have an amplitude of about 2 km s$^{-1}$, and a period of around 35~min. The decay times show some evidence for a temperature dependence with the lowest temperature of formation emission line (\\ion{Fe}{12} 195.12~\\AA) exhibiting a decay time of about 43~min, while the highest temperature of formation emission line (\\ion{Fe}{15} 284.16~\\AA) shows no evidence for decay over more than two periods of the oscillation. The data appear to be consistent with slow magnetoacoustic standing waves, but may be inconsistent with conductive damping. ",
        "introduction": "In recent years, oscillatory phenomena have been detected in the solar corona using spacecraft imaging instruments \\citep[e.g.,][]{Aschwanden1999,Schrijver2002} and spectroscopic instruments \\citep[e.g.,][]{Wang2002,Kliem2002,Mariska2006}. Detection and characterization of coronal oscillatory phenomena with spectroscopic instruments that simultaneously capture many emission lines is particularly important, since it offers the possibility of refining our understanding of the temperature-dependent behavior of the oscillations. It may also be possible to make simultaneous observations in density-sensitive pairs of emission lines, further constraining the characteristics of the oscillations.  The EUV Imaging Spectrometer (EIS) on \\textit{Hinode} produces stigmatic spectra in two 40~\\AA\\ wavelength bands centered at 195 and 270~\\AA. EIS can image the Sun using 1\\arcsec\\ and 2\\arcsec\\ slits to produce line profiles and 40\\arcsec\\ and 266\\arcsec\\ slots to produce monochromatic images. When the slits are used, the EIS mirror can be moved to construct spectroheliograms by rastering an area of interest. This can provide a snapshot of the thermal structure and dynamics of a region of interest, but with typical exposure times ranging from a few seconds to a minute or more, it is difficult to obtain high time cadence dynamical data in this manner. When a higher cadence is important, EIS can operate in a sit-and-stare mode in which the slit covers a fixed location on the Sun and takes repeated exposures. A more global context can then be provided by preceding and/or following the sit-and-stare data with EIS spectroheliograms. Further context is often available from observations taken simultaneously with the \\textit{Hinode} X-Ray Telescope (XRT). An overall description of EIS is available in \\citet{Culhane2007}, XRT is described in \\citet{Golub2007}, and the \\textit{Hinode} mission is described by \\citet{Kosugi2007}.  In this Letter we show that damped oscillations are present in some spectral lines observed in sit-and-stare observations taken with EIS. While the earlier spectroscopic observations detected oscillations in plasma at flare temperatures \\citep[e.g.,][]{Wang2003,Mariska2006}, the data presented here show the presence of oscillations at typical active region temperatures. The oscillations appear to be standing slow mode magnetoacoustic waves, but the temperature dependence of the decay times may be inconsistent with conductive damping. ",
        "conclusions": "Defining the maximum amplitude displacement for the oscillations \\citep{Wang2003} as \\begin{equation} A = A_0/(\\omega^2 + \\lambda^2)^{1/2} , \\end{equation} we calculate a value for $A$ in the \\ion{Fe}{14} 274.20~\\AA\\ line of about 860~km. Thus, if the oscillations are transverse to the line of sight, the range of motion would only be on the order of 1\\arcsec, making them very difficult to detect in images.  The Doppler shift data shown in Figure~\\ref{fig:eis-row-osc} looks similar to the Doppler shift oscillations observed with the Solar Ultraviolet Measurements of Emitted Radiation instrument (SUMER) on the \\textit{Solar and Heliospheric Observatory} (\\textit{SOHO}) \\citep[e.g.,][]{Wang2003}. Those oscillations were primarily seen in emission lines from \\ion{Fe}{19} and \\ion{Fe}{21} and appeared to be the response to flare-like heating in active region loops. The amplitudes were much larger than those reported here---averaging about 98 km~s$^{-1}$. The average period reported was 17.6~min, but ranged up to 31.1~min.  As we discussed in \\S2, we believe that the observed oscillation is the result of a heating event in a nearby structure that appears to be magnetically connected to the area where the oscillations are observed. Assuming we are observing a loop, that the oscillations are near one footpoint and that the other footpoint is at an ($x$, $y$) position of approximately ($860$, $-130$), we compute a distance between the two footpoints of roughly 218~Mm. A semicircular loop with this diameter would then have a length of approximately 342~Mm. For a sound speed of 200 km s$^{-1}$, the sound travel time between the footpoints is about 29~min. This is within a factor of two of the value expected for a standing magnetoacoustic wave oscillating in the fundamental mode. Since the locations of the footpoints are highly uncertain, especially since they are very close to the limb, and the loop geometry is difficult to determine, we believe this this level of agreement is good and suggests that we are observing a damped standing mode slow magnetoacoustic wave.  A strong indicator that the oscillations are standing mode would be the presence of a $1/4$ period phase shift between the intensity and the Doppler shift \\citep{Sakurai2002}. Both the \\ion{Fe}{12} and \\ion{Fe}{13} line intensities show local peaks at about 10 and 20 min into the interval used for fitting the Doppler shift data, suggesting a period for any intensity oscillation of about 10 min. The \\ion{Fe}{14} and \\ion{Fe}{15} emission lines show local peaks at about 30 and 50 min into the fitting interval, suggesting a period for any intensity oscillation of of about 20 min. In both cases, these periods are shorter than those listed in Table~\\ref{table:results}. For a sound wave, $v = c_s \\delta \\rho/\\rho$, where $c_s$ is the sound speed and $\\rho$ is the density. Thus, for the formation temperatures of the lines showing Doppler shift oscillations, $\\delta \\rho/\\rho \\approx 0.01$, a very small density fluctuation, which would not be detectable at the signal-to-noise levels in the intensity data. Thus, we believe that the most that can be said about the observed intensities is that they are not inconsistent with standing mode slow magnetoacoustic waves.  The ratio of the decay time to the wave period can provide some additional information on the damping mechanism. For the SUMER oscillations the average value for that ratio is 0.85 \\citep{Wang2003}. For the three cooler lines in Table~\\ref{table:results} the values are 1.2, 1.3, and 2.6, respectively, with the value for the \\ion{Fe}{15} line being greater than 2.8. The result that the ratio for the \\ion{Fe}{15} line is potentially much larger than that for the other three lines would be inconsistent with the conclusion from SUMER observations that the observed slow magnetoacoustic waves are damped by thermal conduction \\citep{Ofman2002}. Because of the strong temperature dependence of the thermal conductivity, we would expect the damping rate to increase with increasing temperature \\citep{Porter1994}.  Because the corona is threaded with a magnetic field, a diverse array of oscillation modes is possible \\citep[e.g.,][]{Roberts2000,Roberts2003}. The sensitivity and temperature coverage of EIS, especially when combined with the high time cadence imaging that XRT can provide, opens up the possibility of detecting new wave modes and new excitation mechanisms. Similarly, this increased sensitivity may make it possible to detect additional dissipation mechanisms. For example, recent work by \\citet{Ofman2007} and \\citet{Selwa2007} show that wave leakage may be important in the complex magnetic geometries of active regions.  While EIS is capable of detecting Doppler shifts of of less than 0.5 km s$^{-1}$, the result is dependent on an accurate removal of the orbital Doppler shift fluctuation, which has an amplitude of about 20 km s$^{-1}$. Since the observation took place over more than one orbit of the spacecraft, the orbital variation can be accurately modelled and we believe it has been correctly removed. In addition, the \\textit{Hinode} pointing exhibits a jitter in both the $x$ and $y$ directions that can have a range of up to 4\\arcsec\\ over a 120~min observation. We have looked at the jitter over the period of the sit-and-stare observation and see no evidence for any correlation between the jitter and the detected oscillations. Note, however, that this jitter means that the EIS 1\\arcsec\\ slit is averaging over more than 1\\arcsec\\ in both $x$ and $y$ during the observation. The fact that we see the oscillation taking place over at least seven pixels in the $y$-direction suggests that the oscillating structures are sizable, and the jitter should not cause a significant problem.  To further verify this observation, we have performed the same analysis on the oscillatory Doppler-shift signal marked in the lower portions of Figure~\\ref{fig:eis-sns-v}. The oscillations observed over that position range (nine pixels), have similar amplitudes, a shorter period (about 25~min), and similar damping times. There is, however, less evidence for a temperatures dependence to the ratio of the damping time to the period. We believe that this second event in this EIS sit-and-stare data set confirms the presence of these oscillations at active region temperatures, but shows the need for additional studies on the damping mechanism."
    },
    "0806/0806.4266_arXiv.txt": {
        "abstract": "An infinite family of axially symmetric thin disks of finite radius is presented. The family of disks is obtained by means of a method developed by Hunter and contains, as its first member, the Kalnajs disk. The surface densities of the disks present a maximum at the center of the disk and then decrease smoothly to zero at the edge, in such a way that the mass distribution of the higher members of the family is more concentrated at the center. The first member of the family have a circular velocity proportional to the radius, representing thus a uniformly rotating disk. On the other hand, the circular velocities of the other members of the family increases from a value of zero at the center of the disks until a maximum and then decreases smoothly until a finite value at the edge of the disks, in such a way that for the higher members of the family the maximum value of the circular velocity is attained nearest the center of the disks. ",
        "introduction": "A fact usually assumed in astrophysics is that the main part of the mass of a typical spiral galaxy is concentrated in a thin disk (\\cite{BT}). Accordingly, the obtention of the gravitational potential generated by an idealized thin disk is a problem of great astrophysical relevance and so, through the years, different approaches has been used to obtain such kind of thin disk models.  Wyse and Mayall (\\citeyear{WM}) studied thin disks by superposing an infinite family of elementary disks of different radii. Brandt (\\citeyear{BR}) and Brandt and Belton (\\citeyear{BB})  constructed flat galaxy disks by the flattening of a distribution of matter whose surface of equal density were similar spheroids. A simple potential-density pair for a thin disk model was introduced by Kuzmin (\\citeyear{KUZ}) and then rederived by Toomre (\\citeyear{T1,T2}) as the first member of a generalized family of models.  The Toomre models are obtained by solving the Laplace equation in cylindrical coordinates subject to appropriated boundary conditions on the disk and at infinity. The Kuzmin and Toomre models of thin disks, although they have surface densities and rotation curves with remarkable properties, represent disks of infinite extension and thus they are rather poor flat galaxy models. Accordingly, in order to obtain more realistic  models of flat galaxies, is better to consider methods that permit the obtention of finite thin disk models.  A simple method to obtain the surface density, the gravitational potential and the rotation curve of thin disks of finite radius was developed by \\cite{HUN1}. The Hunter method is based in the obtention of solutions of Laplace equation in terms of oblate spheroidal coordinates, which are ideally suited to the study of flat disks of finite extension. By superposition of solutions of Laplace equation, expressions for the surface density of the disks, the gravitational potential and its rotational velocity can be obtained as series of elementary functions.  The simplest example of a thin disk obtained by means of the Hunter method is the well known Kalnajs disk (\\cite{KAL}), which can also be obtained by flattening a uniformly rotating spheroid (\\cite{WM,BR,BB}). The Kalnajs disk have a well behaved surface density and represents a uniformly rotating disk, so that its circular velocity is proportional to the radius, and its stability properties  have been extensively studied (see, for instance, Hunter (\\citeyear{HUN1,HUN2}), \\cite{KAL} and \\cite{K-A}).  In this paper we use the Hunter method in order to obtain an infinite family of thin disks of finite radius. We particularize the Hunter general model by considering a family of thin disks with a well behaved surface mass density. We will require that the surface density be a monotonically decreasing function of the radius, with a maximum at the center of the disk and vanishing at the edge, in such a way that the mass distribution of the  higher members of the family be more concentrated  at the center.  The paper is organized as follows. In Sec. 2 we present a summary of the Hunter method used to obtain the thin disk models of finite radius and also we obtain the general expressions for the gravitational potential, the surface density and the circular velocity. In the next section, Sec. 3, we present the particular family of models obtained by imposing the required behavior of the surface densities and then, in Sec. 4, we analyze its physical behavior. Finally, in Sec. 5, we summarize our main results. ",
        "conclusions": "We presented an infinite family of axially symmetric thin disks of finite radius obtained by means of a particularization of the Hunter method. The disk models so obtained are generalizations of the well known Kalnajs disk, which corresponds to the first member of the family. The particularization of the Hunter model was obtained by requiring that the surface density was a monotonically decreasing function of the radius, with a maximum at the center of the disk and vanishing at the edge, in such a way that the mass distribution of the  higher members of the family were more concentrated  at the center.  We also analyzed the rotation curves of the models and we find for the first member of the family, the Kalnajs disk, a circular velocity proportional to the radius, representing thus a uniformly rotating disk, whereas for the other members of the family the circular velocity increases from a value of zero at the center of the disks until reach a maximum at a critical radius and then decreases to a finite value at  the edge of the disk. Also we find that the value of the critical radius decreases as the value of $m$ increases.  We believe that the obtained thin disk models have some remarkable properties and so they can be considered as appropriated realistic flat galaxy models, in particular if the  superposition of these thin disks with appropriated  halo distributions (\\cite{BT}) is considered. We are now considering some research in this direction. We are now also working in the non axially symmetric generalization of the here presented disks models and also in the obtention of the relativistic generalization of them for the axially symmetric case."
    },
    "0806/0806.3670_arXiv.txt": {
        "abstract": "PSR~B0540-69 is a young pulsar in the Large Magellanic Cloud that has similar properties with respect to the Crab Pulsar, and is embedded in a Pulsar Wind Nebula. We have analyzed the complete archival RXTE dataset of observations of this source, together with new Swift-XRT and INTEGRAL-IBIS data. Accurate lightcurves are produced in various energy bands between 2 and 60 keV, showing no significant energy variations of the pulse shape. The spectral analysis shows that the pulsed spectrum is curved, and is best fitted up to 100 keV by a log-parabolic model: this strengthens the similarities with the Crab pulsar, and is discussed in the light of a phenomenologic multicomponent model. The total emission from this source is studied, the relative contributions of the pulsar and the PWN emission are derived, and discussed in the context of other INTEGRAL detected pulsar/PWN systems. ",
        "introduction": "PSR~B0540-69 is situated in the Large Magellanic Cloud, and is one of the few known extragalactic pulsars. Its pulsed emission was discovered by Seward et al. (1984) in the soft X-ray band with the \\emph{Einstein} observatory, and thereafter also in the optical (Middleditch \\& Pennypacker, 1985) and radio bands (Manchester et al., 1993).  This pulsar is usually referred as the ``Crab Twin'', because of the similar period and period derivative with respect to the well known Crab Pulsar (PSR~B0531+21, see Table \\ref{t:crabvs0540} for a comparison). However, their pulse shapes are strongly different: for the Crab it consists of two peaks, with a phase separation of about 0.4, that are visible at all wavelengths, while PSR~B0540-69 has a single, broadly symmetric pulse about 0.5 wide in phase, with some structure at the maximum.  \\begin{table} \\caption{Comparison between various parameters for Crab Pulsar (PSR~B0531+21) and PSR~B0540-69. Data are from the ATNF Pulsar Database, available at \\texttt{http://www.atnf.csiro.au/research/pulsar/psrcat/}. }\\label{t:crabvs0540} \\centering \\begin{tabular}{lcc} & Crab \t& PSR~B0540-69 \\\\ \\hline Period $P$ (ms)\t\t\t\t\t\t\t& 33.5\t& 50.5 \\\\ Period derivative $\\dot{P}$ ($\\times 10^{13}$ s/s)\t& 4.2\t\t& 4.8 \\\\ Braking index $n$ \t\t\t\t\t\t\t& 2.5\t\t& 2.1 \\\\ Distance $d$ (kpc)\t\t\t\t\t\t\t& 2\t\t& 50 \\\\ Characteristic age $\\tau$ (yr)\t\t\t\t\t& 1240\t& 1670 \\\\ Magnetic field $B_{s}$ ($\\times 10^{12}$ G)\t\t& 3.78 \t& 4.96 \\\\ Spin-down luminosity ($\\times 10^{38}$ erg/s)\t\t& 4.6\t\t& 1.5 \\\\ \\hline \\end{tabular} \\end{table}   This pulsar, like the Crab, is also embedded in a Pulsar Wind Nebula (PWN) that shines in the optical to X-ray band (Serafimovich et al., 2004). This pulsar/PWN  system has been observed many times in the X-rays, by ROSAT (Finley et al., 1993), GINGA (Deeter et al., 1999), BeppoSAX  (Mineo et al., 1999), ASCA (Hirayama et al., 2002), RXTE (de Plaa et al., 2003),  and Chandra (Gotthelf \\& Wang, 2000; Hwang et al., 2001; Kaaret et al., 2001; Petre et al., 2007). This source is not detected in $\\gamma$-rays ($>50$ MeV), with only upper limits reported by EGRET (Thompson et al., 1994).  In the optical/near UV band PSR~B0540-69 is a faint $\\sim$22 magnitude source with a power law spectrum having a steep spectral index $\\alpha_{\\nu} \\sim 1$ (Serafimovich et al., 2004) while the Crab has a flat spectrum, $\\alpha_{\\nu} \\sim 0$, in the same band.  In the X-rays, this pulsar has a flux of a few milliCrab, and a spectrum fitted by various authors as a power law with photon indices in the range from 1.8 to 2.0. De Plaa et al. (2003), using high resolution RXTE spectra, showed evidence for a curved, \\emph{log-parabolic} spectrum, i.e. with the form $F(E) = KE^{-a-b\\,\\mathrm{Log}E}$. A similar spectrum has been observed also for the Crab pulsar (Massaro et al., 2000, 2006) and few other sources (Cusumano et al., 2001; Mineo et al., 2004).   PSR~B0540-69's  PWN (also called N158A) has a torus-like structure, with a major radius of $\\sim$4\\arcsec\\ and with the presence of jets along the direction of proper motion. At variance with the Crab, the PWN of PSR~B0540-69 is further surrounded by a Supernova Remnant (SNR) thermal shell, with a radius of about 30\\arcsec, that has radio and soft X-ray emission. Optical (HST; Serafimovich et al., 2004) and X-ray (Chandra; Hwang et al., 2001; Kaaret et al., 2001; Petre et al., 2007) observations showed that the various regions of the PWN/SNR system have different brightnesses and spectra, not described by simple power law models.  A first analysis of PSR~B0540-69 using INTEGRAL data was performed by Slowikowska et al. (2007), who reported pulse profiles in different energy bands and a total, phase-averaged, spatially-extended spectrum.  In this paper we present a comprehensive X/$\\gamma$-ray study of this pulsar-PWN system, analyzing in detail the INTEGRAL observations together with new SWIFT/XRT data. We also exploit the very large RXTE observation dataset of this pulsar, which spans a period of more than ten years.  This paper is structured as follows. In Sect. \\ref{obs-datared} we describe the X-ray observations and data reduction and we present our results on the pulsar and PWN spectra. In Sect. \\ref{spectra} and \\ref{conclusions} we discuss these results and their interpretation also in the context of pulsar-PWN high-energy emission models. ",
        "conclusions": "PSR~B0540-69 is an interesting young pulsar, showing many similarities but also some differences with respect to the Crab. The study of this pulsar, although made difficult by its high distance and apparent faintness, could help the theoretical understanding of the characteristics of the high-energy emission from the pulsars and their PWNs.  With the analysis of the complete RXTE dataset of observations of PSR~B0540-69, exploiting for the spectral analysis the longest gain epoch that spans more than half of these data, together with that of new SWIFT/XRT and INTEGRAL/IBIS data, we have compiled an accurate and up-to-date picture of the emission from this source in the X and soft $\\gamma$-rays, from 0.7 to 200 keV. Thanks to the full HEXTE dataset we have also determined the contribution of pulsed hard X-ray emission to the total one.  Regarding the pulsed emission, we confirm the presence of a curved spectrum, peaking at about 12 keV, well described by a log-parabolic law up to $\\ga$100 keV. We have shown also that it presents some similarities with the Crab pulsed emission.  We have studied in detail the total (pulsar+PWN) emission from this source, in the energy range 0.7--100 keV, and we have shown that the major (about 75--80\\%) contribution comes from the synchrotron nebula. The comparison of PSR~B0540-69 with other IBIS pulsar/PWN systems shows that it has the highest conversion efficiency in the hard X/soft $\\gamma$-ray band, with similar spectral properties with respect to the other sources. Most of these INTEGRAL PSR/PWN systems have been associated with HESS sources. PSR~B0540-69 looks different, in this respect, not  showing evidence of emission in TeV range. We have discussed the possible instrumental limitations, due to the low flux and declination, but there is also the possibility that this source genuinely lacks TeV emission.  Future observations, especially in the hard X-ray band, will further refine this picture. High angular resolution long Chandra X-ray observations will permit an accurate study of the synchrotron nebula to derive its physical parameters. Additional INTEGRAL pointings will increase the statistics of the total emission, while the advent of hard X-ray focusing telescopes, like Simbol-X (Ferrando et al., 2005), will help to disentangle the pulsed emission from the nebular one also in the HEXTE and IBIS energy range. Moreover, deep GLAST observations could be useful to estabilish more strict upper limits to the $\\gamma$-ray emission from this source."
    },
    "0806/0806.4099_arXiv.txt": {
        "abstract": "The most recent results from direct searches for dark matter particles in the galactic halo are examined in terms of an effective Minimal Supersymmetric extension of the Standard Model at the electroweak scale without gaugino masses unification. We show that the annual modulation effect at  8.2 $\\sigma$ C.L. recently presented by the DAMA Collaboration, as the result of a combined analysis of the DAMA/NaI and the DAMA/LIBRA experiments for a total exposure of 0.82 ton yr, fits remarkably well  with what expected for relic neutralinos for a wide variety of WIMP distribution functions. Bounds derivable from other measurements of direct searches for dark matter particles are analyzed. We stress the role played by the uncertainties affecting the neutralino--quark couplings arising from the involved hadronic quantities. We also examine how present data on cosmic antiprotons can help in constraining the neutralino configurations selected by the DAMA effect, in connection with the values of the astrophysical parameters. Perspectives for measurement of antideuterons possibly produced  in the galactic halo by  self--annihilation of neutralinos belonging to the DAMA configurations are examined. Finally, we discuss how findings at LHC would impact on these issues. ",
        "introduction": "\\label{sec:intro}   We are at present witnessing a great activity in running and preparing experimental projects of direct searches for dark matter (DM) particles in the galactic halo \\cite{taup2007}.  A number of new results came out recently. Most notably, the DAMA Collaboration, by analyzing the data obtained with the DAMA/LIBRA experiment \\cite{dama/libra}, has confirmed the evidence for the annual modulation effect already measured with the previous DAMA/NaI experiment \\cite{dama/nai}. As a result of the combined analysis of the DAMA/NaI and the DAMA/LIBRA experiments, for a total exposure of 0.82 ton yr,  the effect of the annual modulation is now at 8.2 $\\sigma$ C.L. \\cite{dama/libra}. Other Collaborations have reported upper bounds on the WIMP--nucleon scalar cross section \\cite{taup2007,xe,cdms,kims} (as usual here  WIMP  stands for a generic Weakly Interacting Massive Particle), using other approaches and scaling from different nuclei.  In this paper we discuss the relevance of these experimental results for the most widely discussed candidate for cold dark matter, the neutralino. Neutralino configurations (concerning here mainly light neutralinos) which fit the annual modulation data are extracted from the DAMA results and confronted with the current measurements on galactic antiprotons. Perspectives of seeing signals due to these configurations in forthcoming measurements of galactic antideuterons are examined. The current experimental upper bounds on the WIMP-nucleon scalar cross section are also scrutinized.  In all our considerations we explicitly keep into account the uncertainties affecting the neutralino--quark couplings because of the involved hadronic quantities, typically the pion--nucleon sigma term $\\sigma_{\\pi N}$, the quantity $\\sigma_0$ which is related to the SU(3) symmetry breaking and a quark mass ratio $r$ (see definitions later on).  In the analysis of the DAMA data, we discuss both the case in which the channeling effect \\cite{channeling} is included in their analysis as the one where this effect is not taken into account; these two cases are treated separately. We recall that the channeling effect occurs when an ion, traversing a crystalline structure along a path (quasi-) parallel to crystallographic axes or planes, has a deeper penetration than in an amorphous material of the same composition \\cite{drob}. In Ref. \\cite{channeling} it is shown that the occurrence of the channeling effect makes the response of the DAMA NaI(Tl) detector more sensitive to WIMP-nucleus interactions than previously estimated.  Finally, in the present paper it is stressed that CERN LHC (Large Hadron Collider), which hopefully will shed light on the Higgs sector and supersymmetric extensions of the Standard Model (SM), has a remarkable discovery potential in terms of the population of  light neutralinos which fit the annual modulation data.  The plan of the paper is the following. In Sect. II the main features of the supersymmetric model employed in the present paper are delineated. Sect. III is devoted to a recollection of the main phenomenological properties used in our analysis, and Sects. IV--V to a detailed comparison of our theoretical evaluations to the current results of WIMP direct detection. In Sect. VI we discuss the constraints implied by measurements of galactic antiprotons and in Sect. VII we examine the perspectives for forthcoming measurements on cosmic antideuterons. Impact of LHC on these issues is discussed in Sect. VIII, and finally conclusions are drawn in Sect. IX. ",
        "conclusions": "\\label{sec:conclusions}  In this paper we have analyzed the most recent experimental data on direct searches for dark matter particles in the galactic halo. We have discussed the various features  involved in an interpretation of these data in terms of relic neutralinos; the supersymmetric scheme employed is an effective MSSM scheme at the electroweak scale, without  gaugino--mass unification at a Grand Unified scale. The role of  the uncertainties affecting the neutralino--quark couplings, because of the involved hadronic quantities, are critically discussed and included in our considerations.  First we have considered in detail the results of the DAMA Collaboration which, by a combined analysis of the DAMA/NaI and the DAMA/LIBRA experiments, for a total exposure of 0.82 ton yr,  provide evidence of an annual modulation effect at 8.2 $\\sigma$ C.L. Comparison of our theoretical evaluations with the DAMA data has been carried out both in the case in which channeling effect is included in the DAMA analysis and in the one where it is excluded. The DAMA results presented here refer uniquely to the case in which the putative WIMP has a coherent interaction with the nuclei in the detector and are therefore translated in terms of regions  in the $m_{\\chi}$--$\\xi \\sigma_{\\rm scalar}^{(\\rm nucleon)}$ plane, for different selections of the galactic distribution functions \\cite{dama3}. We wish also to recall  that DM--detector interaction mechanisms or DM candidates different from the one considered in the present paper could be the origin of the annual modulation effect \\cite{dama/libra}.   The roles the hadronic uncertainties and of the effect of channeling are kept separated at any stage of our discussion on DAMA data; the analysis has been carried out in such a way that, once hopefully some of the afore mentioned uncertainties are resolved in the future, one can easily employ our present results to narrow down the physical relevant regions.   By considering first the DAMA data  when a convolution of galactic distribution functions is considered, we have shown that the annual modulation region is largely compatible with the theoretical predictions for relic neutralinos with masses $m_{\\chi} \\lsim$ 100 GeV, and in particular for neutralinos within the low--energy funnel for $m_{\\chi} \\lsim$ 50 GeV. This occurs, whether or not the channeling effect is included.   We have then pursued our analysis by employing the cored isothermal sphere DF and  shown that relic neutralinos fit quite well the annual modulation data. For the case where the channeling is included according to the modeling of Ref. \\cite{channeling}, experimental data favor values of  $v_0$ and $\\rho_0$ which are in the low-medium side of their own physical ranges, {\\it i.e.} $v_0 \\simeq$ (170 -- 220) km sec$^{-1}$ and $\\rho_0 \\simeq (0.3 - 0.4)$ GeV cm$^{-3}$ and neutralino masses in the mass range $m_{\\chi} \\simeq (7 - 30)$ GeV. In case of no-channeling,  low values of $v_0$ and $\\rho_0$ ($v_0 \\simeq$ 170 km sec$^{-1}$ and $\\rho_0 \\simeq 0.2$ GeV cm$^{-3}$) appear to be somewhat disfavored, though in this case neutralinos with $m_{\\chi} \\simeq$ 60--100 GeV could be involved.  The agreement between experimental data and theoretical evaluations looks very good for the other combinations of $v_0$ and $\\rho_0$ values, with an overall preference for neutralinos of low mass. We have also commented about the results one would obtain when some other DF is selected.  The neutralino populations selected on the basis of the annual modulation data have been analyzed in terms of the antiproton fluxes which they would produce in our halo. We have shown that many of them are fully compatible with the current stringent bounds on cosmic antiprotons, especially for values of local dark matter density $\\rho_0$ and local rotational velocity $v_0$ in the low side of their physical ranges, and for values of the diffusion parameters not too close to the values of their maximal set.  It was also derived that forthcoming measurements of galactic antideuterons will have good chances to explore fluxes generated by the self--annihilation of the neutralino configurations compatible with the annual modulation data.  The upper bounds on   $\\xi \\sigma_{\\rm scalar}^{(\\rm nucleon)}$,  derived by other experiments of WIMP direct searches \\cite{taup2007,xe,cdms,kims}, have  been analyzed by considering various WIMP galactic DFs and by including the uncertainties induced by the previously mentioned hadronic quantities. It is concluded that no conservative upper bound can now exclude the light  ($m_{\\chi} \\lsim$ 50 GeV) neutralino population within the model discussed in Sect. \\ref{sec:susy}.  We have pointed out that LHC is expected to  allow an efficient exploration of the supersymmetric parameter regions compatible with light neutralinos and the annual modulation data. Due to the characteristic features of the physical regions involved, the measurements of LHC  could be able to prove or disprove the supersymmetric model considered in the present paper and then validate or not our interpretation of the annual modulation effect in terms of relic neutralinos.   It is finally worth noting that the relic neutralinos of light masses, considered in the present paper and previously in Refs. \\cite{lowneu,ind}, can be of relevance also in many interesting astrophysical contexts different from the ones discussed before. For instance, by their self--annihilations, these neutralinos may be the origin of electron--positron pairs which could interact with photons of the cosmic microwave background producing a peculiar Sunyaev--Zeldovic effect \\cite{cola,cola2}. Also, positrons and electrons created by annihilation of light neutralinos in the galactic center might generate a synchrotron radiation responsible for an excess of microwave emission from a region within $\\sim 20^o$ of the galactic center \\cite{fink} (see also Ref. \\cite{haze}). Light neutralino dark matter could also imply a sizable primordial $^6$Li abundance \\cite{kj}, a feature particularly interesting in view of recent new determinations of this abundance in metal--poor halo stars \\cite{asplund}."
    },
    "0806/0806.0384_arXiv.txt": {
        "abstract": "Current models of galaxy evolution suggest that feedback from active galactic nuclei is needed to explain the high-luminosity cutoff in the galaxy luminosity function.  Exactly how an AGN outflow couples with the ambient medium and suppresses star formation remains poorly understood.  However, we have recently uncovered an important clue to how that coupling might work.  Observations of H$\\alpha$ emission and blue light from the universe's most luminous galaxies, which occupy the centers of galaxy clusters, show that star formation happens only if the minimum specific entropy of the intracluster gas is $\\lesssim 30 \\, {\\rm keV \\, cm^2}$.  Here we suggest that this threshold for star formation is set by the physics of electron thermal conduction, implying that conduction is critical for channeling AGN energy input toward incipient star-forming regions and limiting the progress of star formation. ",
        "introduction": "\\setcounter{footnote}{0}  What determines the upper limit to the luminosity of a galaxy?  This question has received a great deal of attention lately because the answer seems to involve feedback from a galaxy's nucleus.  Semi-analytic models of galaxy formation show that feedback of some kind is necessary to account for the observed characteristics of the galaxy luminosity function.  Models without feedback overpredict the number of both low-luminosity and high-luminosity galaxies.  Plausible amounts of feedback from supernovae bring predictions for low-luminosity galaxies into agreement with observations, but getting things right at the high-luminosity end is more difficult---the amount of supernova feedback needed to quench star formation in these systems is implausibly large \\citep{Benson03}.  Numerical simulations of galaxy formation treat the physics of merging and cooling more realistically but still do not solve the problem of the high-luminosity cutoff.  In current simulations, late-time cooling of hot gas in the central galaxies of clusters produces so many stars at $z \\sim 0$ that the galaxies have luminosities several times larger than the observed cutoff in the galaxy luminosity function and colors that are much too blue \\citep[e.g.,][]{Saro06}.  A closely related cooling problem has long bedeviled our understanding of galaxies and galaxy clusters, and X-ray observations strongly suggest that the solution lies with active galactic nuclei.  In roughly half of nearby clusters, the cooling time at $r \\lesssim 100$~kpc is less than a Hubble time, implying that gas at the center of the cluster must cool and condense, if there is no mechanism to offset cooling.  However, the star-formation rates in these systems are at most 1\\% to 10\\% of the expected gas condensation rate, and X-ray spectroscopy shows that the the bulk of the central gas remains suspended at temperatures $\\gtrsim 1/3$ of the virial temperature \\citep[see][for a review]{pf06}.  High-resolution imaging with {\\em Chandra} has revealed that many clusters and giant elliptical galaxies with short central cooling times also have cavities in their ICM that coincide with radio-emitting plasma outflows from the active nucleus in the central galaxy.  In many cases the AGN energy output inferred from these cavities is similar to the X-ray cooling rate of the central gas---a strong piece of circumstantial evidence in favor of AGN feedback as the mechanism that limits cooling and star formation in high-mass halos \\citep{Birzan04,mn07}.  Not all clusters with short central cooling times have obvious cavities, but that finding can be understood if AGN heating is episodic.  The most recent versions of semi-analytic models have therefore incorporated schematic implementations of AGN feedback and find that ``radio-mode\" feedback is a plausible explanation for both the observed cutoff in the galaxy luminosity function and the fact that the largest galaxies in the universe are red, not blue \\citep{Croton06,Bower06}.  To match the observations, radio-mode feedback must occur preferentially in high-temperature systems and must provide enough feedback energy to suppress star formation without making many stars.  In this context, AGN feedback is most effective in systems experiencing hot-mode accretion because an extended hot galactic atmosphere is necessary to halt and thermalize the energy of the AGN outflow.  Here we consider the implications of two recent studies showing that star formation in brightest cluster galaxies (BCGs) is closely linked to the entropy\\footnote{As is customary in the X-ray cluster field, we quantify entropy using the adiabatic constant $K = kT n_e^{-2/3}$ \\citep[see][for more background]{Voit05}.} structure of the intracluster medium (ICM).  Section~2 presents evidence that star formation occurs only in BCGs whose central entropy is $\\lesssim 30 \\, {\\rm keV \\, cm^2}$, equivalent to a cooling time $\\lesssim 5 \\times 10^8$~years.  The star formation rates in these systems are generally much smaller than the gas cooling rates in the absence of feedback, and understanding why star formation still proceeds at a reduced rate is likely to tell us something about the coupling between AGN feedback and the ICM.  In section~3 we suggest that the critical central entropy threshold for star formation is determined by electron thermal conduction, because multiphase structure cannot persist in the ICM if the central entropy is much larger than the observed threshold.  If conduction is indeed responsible for governing the rate of star formation, then it is also likely to be a major conduit of thermal energy in the cluster core and a crucial part of the AGN feedback loop.  Section~4 summarizes the paper.   \\begin{figure}[t] \\includegraphics[width=3.5in, trim = 1.2in 0.2in 1.6in 0.7in]{f1.eps} \\caption{ \\footnotesize Dependence of H$\\alpha$ emission in BCGs on central entropy from Cavagnolo et al. (2008a).  Central entropy $K_0$ has been measured by fitting {\\em Chandra} entropy profiles with a radial power law plus a constant, and $L_{{\\rm H}\\alpha}$ values are a heterogeneous set taken from the literature.  Orange filled circles show clusters with detected H$\\alpha$, and black filled circles give upper limits for clusters without detections (see Cavagnolo et al. 2008b for more details).  Notice that strong H$\\alpha$ emission appears only when $K_0 \\lesssim 30 \\, {\\rm keV \\, cm^2}$. \\label{Ha-K0}} \\end{figure} ",
        "conclusions": "Two recent studies have shown that the star-forming properties of BCGs depend on the entropy structure of the host cluster's ICM.  Clusters with central entropy $\\lesssim 30 \\, \\Kunits$, corresponding to a cooling time $\\lesssim 5 \\times 10^8$~years often have BCGs with H$\\alpha$ emission (Cavagnolo et al. 2008b) and blue cores (Rafferty et al. 2008) indicative of ongoing star formation, while none of the clusters with high central entropy show the same features.  We suggest that this entropy threshold for star formation is set by the physics of thermal conduction.  Clusters typically have entropy profiles that approximately follow the law $K(r) = K_0 + (150 \\, \\Kunits)(r/ 100 \\, {\\rm kpc})^{1.2}$, with differing values of $K_0$  (Cavagnolo et al. 2008a).  If $K_0  > 30 \\, \\Kunits$, then conduction can carry heat energy toward the center of the cluster faster than the ICM can radiate it away, as long as the suppression factor $f_c \\gtrsim 0.2$.  Therefore, we should not expect clusters with higher central entropy levels to have a persistent multiphase medium in their cores.  Conversely, conduction in lower-entropy clusters cannot fully compensate for radiative cooling, allowing star formation to proceed.  \\citet{gor08} have recently presented a model explaining how AGN feedback can suppress star formation in these thermally unstable cluster cores.  The correspondence we find between the observed star-formation threshold and the predictions of conductive models indicates that conduction indeed is important in cluster cores and may be crucial to the AGN feedback mechanism thought to regulate star formation in these systems.  However, detailed hydrodynamical simulations including conduction will be necessary to determine more precisely how the phase structure of the ICM depends on a cluster's entropy profile.  \\vspace*{1.0em}  This work was partially supported at MSU by NASA grants NNG04GI89G and NNG05GD82G.  PEJN was supported by NASA grant NAS8-01130"
    },
    "0806/0806.2662_arXiv.txt": {
        "abstract": "Recent theoretical studies have revealed the possibly important role of the  capture and annihilation process of weakly interacting massive particles (WIMPs) for the first stars.  Using new evolutionary models of metal-free massive stars, we investigate the impact of such ``dark matter burning'' for the first stars in different environments of dark matter (DM) halos, in terms of the ambient WIMP density ($\\rho_\\chi$).  We find that, in agreement with existing literature, stellar life times can be significantly prolonged for a certain range of $\\rho_\\chi$ (i.e., $10^{10} \\la \\rho_\\chi [\\dmrho] \\la 10^{11}$ with the current upper limit for the spin-dependent elastic scattering cross section $\\sigma_0^\\mathrm{SD} = 5 \\times 10^{-39} \\mathrm{cm^2}$).  This greatly enhances the role of rotationally induced chemical mixing in rotating stars, in favor of abundant production of primary nitrogen, massive helium stars and long gamma-ray bursts, from the first stars.  We also find that stars with $\\rho_\\chi > 2 \\times 10^{11}~\\dmrho$ may not undergo nuclear burning stages, confirming the previous work, and that ionizing photon fluxes from such DM supported stars are very weak.  Delayed metal enrichment and slow reionization in the early universe would have resulted if most of the first stars had been born in DM halos with  such high $\\rho_\\chi$, unless it had been lowered significantly below the threshold for efficient DM burning on a short time scale. ",
        "introduction": "In the current framework of the $\\Lambda$CDM cosmological model, the possibility that dark matter (DM) annihilation effects may impact stellar evolution has recently received renewed attention. After the  original works of Bouquet, Dearborn, Freese, Gould, Griest, Krauss, Olive,Press, Raffelt, Renzini, Salati, Silk, Spergel, Srednicki and Wilczek in the`80s and early `90s, several authors have recently re--examined the effects that DM, if made of weakly interacting massive particles (WIMPs), would have on compact objects  \\citep{Moskalenko:2007ak, Bertone:2007ae} and on the zero-age main sequence of low--mass stars \\citep{Fairbairn:2007bn}.  This exciting activity has been motivated by  a double scope: any peculiar and distinguishable feature of WIMP  annihilation on observable stellar quantities is extremely precious in the ``quest'' for dark matter evidence; on the other hand, all possible  effects impacting the life of celestial objects must be taken into account by astrophysicists in the current precision era.  In particular, the first stellar episode at high redshift occurs under very  different conditions from those in the present universe.  The higher concentration of dark matter,  the short Hubble time which prevents DM self--annihilation  from severely affecting the central density, and the characteristic formation of a single PopIII star in the center of  the halo are the most favorable conditions for DM annihilation effects to be very efficient in the first stars.  In their pioneering work,  \\citet{Spolyar:2008qv} first found the possibility of very high DM density in primordial halos at high redshift to such an extent that energy released from DM annihilation at the center may halt the gravitational collapse of the baryonic cloud, calling such a DM powered object a {\\it dark} star.  \\citet[hereafter I08]{Iocco:2008xb} and \\citet{Freese:2008ur} also noticed that WIMP capture is most efficient in Population III stars.  More recently, \\citet{freese08a, freese08b} and \\citet{Iocco:2008rb} further investigated the role of annihilation of adiabatically contracted DM in the formation of the first stars.  \\citet{Iocco:2008rb} have also followed the evolution from the pre-main sequence phase to helium exhaustion in the presence of WIMP capture and annihilation (hereafter, DM burning), showing that this can severely delay the evolution of pre--MS objects in the early universe, as well as extend their MS lifetimes.  All of these studies motivate us to explore possible consequences of DM burning for the final fate of the first stars and their feedback effects on the evolution of the early universe, even if the role of DM annihilation in the formation of the first stars still remains subject to many uncertainties \\citep[see][for a recent review]{FS3proc}.  In this Letter, we address the issue by discussing the evolution of the first stars of $20\\le M/\\mathrm{M_\\odot} \\le 300$ up to the carbon burning stage (Sect.~\\ref{sect:results}).  We also investigate the interplay of rotation with DM burning in the evolution of the first stars of $100~\\mathrm{M_\\odot}$, given the particular importance of rotation for the evolution and deaths of metal poor massive stars \\citep[e.g,][]{Meynet08, Yoon08}. Implications of our results for the history of reionization in the early universe are briefly discussed (Sect.~\\ref{sect:discussion}). ",
        "conclusions": "\\label{sect:discussion}  Our results indicate that DM burning should not significantly alter our view  on the final fate of the non--rotating first stars: pair--instability supernovae for $ 140 \\la  M/\\mathrm{M_\\odot} \\la 260$, and core--collapse events for other masses \\citep{Heger02}, although their life times may be significantly prolonged.  However, the impact of DM burning appears more important for rotating stars.  The quasi-chemically homogeneous evolution can be rather easily realized even with moderate rotation velocities with $10^{10} \\la \\rho_\\chi~[\\dmrho] \\la 10^{11}$.  Such evolution can lead to production of massive helium stars that emit large amounts of helium ionizing photons, as well as abundant production of primary nitrogen, as discussed in \\citet{Yoon08}.  Note also that the quasi-chemically homogeneous evolution scenario (CHES) is one of the favored ones for the production of long GRBs from metal poor stars~\\citep{Yoon05, Woosley06}.  Our result therefore indicates that DM burning might promote the production of long gamma--ray bursts from the first stars of $12 \\la M/\\mathrm{M_\\odot} \\la 60$ via the CHES channel \\citep[see][]{Yoon06}.  DM burning in the first stars must have consequences in the history of reionization in the early universe.  Table~1 lists the number of hydrogen and helium ionizing photons emitted from $100~\\mathrm{M_\\odot}$  models.   If $\\rho_\\chi \\la 2\\times10^{11} \\dmrho$, the total number of ionizing photons increases proportionally to the DM density, as a direct consequence of the life--prolonging effect of DM burning. For a given $\\rho_\\chi$ rotation does not significantly alter hydrogen ionizing photon counts nor the lifetime.  However, for models that take path of the CHES, the helium ionizing photon counts is increased by more than a factor of 2 compared to the non--rotating case.  If $\\rho_\\chi$ is very large, on the other hand,  the surface temperature of the star drops significantly enough (Fig. 1) that the  total number of ionizing photons is reduced even with much longer  lifetimes due to the DM burning.  For example,  with $\\rho_{\\chi}= 2\\times10^{12}~ \\dmrho$, it would have to take $\\sim$ 10 Gyr to emit as many hydrogen ionizing photons as a non--DM burning counterpart.  Therefore, if most of the first stars  had been born with such high $\\rho_\\chi$ their contribution to reionization would have been dramatically reduced.  The effect of the temperature drop on the number of helium ionizing photons is even more prominent: with $\\rho_\\chi = 2\\times10^{12}~\\dmrho$, it decreases by more than 19 orders of magnitude  compared to the other cases with lower $\\rho_\\chi$.  Future study on the history of helium ionization at high redshift might therefore be a strong probe of DM burning in the first stars.  In this study we assume that the background DM halo density stays constant throughout the stellar evolution.  This assumption may be valid if the stellar life times are shorter than about 100 Myr -- which is the expected merger timescale of DM halos at  at z $\\sim$ 20~\\citep[e.g.,][]{Lacey93}  -- as in our model sequences with $\\rho_\\chi \\la 4\\times10^{10}~\\dmrho$.  The evolution of DM burners in halos with higher $\\rho_\\chi$, however, should be critically determined by the change of DM halo environments.  If $\\rho_\\chi$ is sufficiently reduced due to merger events and/or to displacement of the star from the densest region of the DM halo, the DM burners will  become ``normal'' stars, dominated by nuclear burning.  The star may then die quickly as we expect for normal stars, or become a ``born--again'' DM burner if $\\rho_\\chi$ increases again in later stages for some reason. The detailed history of the feedback from the first stars (e.g. metal enrichment and reionization) on the evolution of the early universe may depend on the nature and evolution of DM halos where the first stars are formed.  This issue should be addressed in future work.  \\begin{table} \\begin{center} \\caption{Number of hydrogen and helium ionizing photons from $100~\\mathrm{M_\\odot}$ star models}~\\label{tab1} \\begin{tabular}{crccc} \\tableline \\tableline $v_\\mathrm{rot}/v_\\mathrm{K}$ & $\\rho_\\chi~[\\dmrho]$  &  $N_\\mathrm{H}$  & $N_\\mathrm{He}$  & Duration \\\\ \\tableline 0.0               &      0.00     &  $1.2 \\times 10^{64}$  & $2.2 \\times 10^{62}$    & 3.2 Myr \\\\ 0.0               &      $2\\times10^{10}$     &  $2.1 \\times 10^{64}$  & $3.4 \\times 10^{62}$    & 5.5 Myr \\\\ 0.0               &      $4\\times10^{10}$     &  $8.5 \\times 10^{64}$  & $1.5 \\times 10^{63}$    & 22.3 Myr \\\\ 0.0               &      $10^{11}$     &  $2.9 \\times 10^{65}$  & $6.3 \\times 10^{62}$    & 100.0 Myr\\tablenotemark{*} \\\\ 0.0               &      $2\\times10^{11}$     &  $2.0 \\times 10^{65}$  & $1.7 \\times 10^{61}$    & 100.0 Myr\\tablenotemark{*} \\\\ 0.0               &      $2\\times10^{12}$     &  $1.3 \\times 10^{62}$  & $9.4 \\times 10^{43}$    & 100.0 Myr\\tablenotemark{*} \\\\ \\tableline 0.1               &      0.00    &  $1.5 \\times 10^{64}$  & $2.5 \\times 10^{62}$    & 3.4 Myr \\\\ 0.1               &      $2\\times10^{10}$ &  $2.7 \\times 10^{64}$  & $5.2 \\times 10^{62}$    & 6.0 Myr \\\\ 0.1               &      $4\\times10^{10}$ &  $8.7 \\times 10^{64}$  & $3.8 \\times 10^{63}$    & 19.6 Myr \\\\ \\tableline \\end{tabular} \\vspace{-0.5mm} \\tablenotetext{*}{The numbers are calculated  only for the first 100 Myr.} \\end{center} \\end{table}"
    },
    "0806/0806.2348_arXiv.txt": {
        "abstract": "{X-ray spectroscopy is a sensitive probe of stellar winds.  X-rays originate from optically thin shock-heated plasma deep inside the wind and propagate outwards throughout absorbing cool material.  Recent analyses of the line ratios from He-like ions in the X-ray spectra of O-stars highlighted problems with this general paradigm: the measured line ratios of highest ions are consistent with the location of the hottest X-ray emitting plasma very close to the base of the wind, perhaps indicating the presence of a corona, while measurements from lower ions conform with the wind-embedded shock model.  Generally, to correctly model the emerging X-ray spectra, a detailed knowledge of the cool wind opacities based on stellar atmosphere models is prerequisite. A nearly grey stellar wind opacity for the X-rays is deduced from the analyses of high-resolution X-ray spectra. This indicates that the stellar winds are strongly clumped. Furthermore, the nearly symmetric shape of X-ray emission line profiles can be explained if the wind clumps are radially compressed.  In massive binaries the orbital variations of X-ray emission allow to probe the opacity of the stellar wind; results support the picture of strong wind clumping. In high-mass X-ray binaries, the stochastic X-ray variability and the extend of the stellar-wind part photoionized by X-rays provide further strong evidence that stellar winds consist of dense clumps.} ",
        "introduction": "All OB-stars with spectral types earlier than B1.5 have X-ray luminosities $L_{\\rm X}$ exceeding $10^{30.5}$\\,erg\\,s$^{-1}$. Remarkably, their X-ray luminosity roughly correlates with stellar bolometric luminosities as $L_{\\rm X}$$\\approx$$10^{-7}L_{\\rm bol}$ (Seward et al. \\cite{sew79}, Bergh\\\"ofer et al. \\cite{berg97}).  The high-resolution X-ray spectra of two O-stars, $\\zeta$\\,Pup and $\\zeta$\\,Ori, are shown in Figure\\,1. The spectra are dominated by strong emission lines that reflect stellar wind abundances, e.g. nitrogen lines are more prominent in the spectrum of the more evolved star $\\zeta$\\,Pup. X-ray spectra of single non-magnetic massive stars can be explained as emitted by optically thin plasma heated to  temperatures $kT_{\\rm X}$$\\lsim$$0.6$\\,keV (e.g. Wojdowski \\& Schulz \\cite{ws05}).  There is a number of hypotheses to explain how X-rays are generated in stellar winds. Lucy \\& White (\\cite{lw80}) put forward a theory where X-rays  originate from bow shocks around a population of blobs (see also Cassinelli et al., these proceedings). Cassinelli \\& Swank (\\cite{csw83}) proposed an analogy with coronal stars where X-rays originate from magnetically confined regions at the wind base. Owocki et al.\\ (\\cite{ocr88}) performed numerical simulations of the evolution of instabilities in radiatively driven stellar winds and predicted a development of reverse shocks with postshock temperatures in the range $10^{6}-10^{7}$\\,K (i.e. as inferred from the observed X-ray spectra).  Feldmeier et al.\\ (\\cite{feld97}) developed a model explaining the X-rays as originating from cooling zones behind shock fronts and were able to reproduce the low-resolution {\\em Rosat} spectra of $\\zeta$\\,Pup and $\\zeta$\\,Ori.  In massive binaries a copious amount of X-ray emission is produced by the collision of two powerful stellar winds. X-ray spectra of colliding wind binaries reveal that the X-ray emitting plasma departs from collisional equilibrium (e.g. Pollock et al. \\cite{amtp05}). Colliding wind systems tend to have somewhat higher X-ray luminosities and their X-ray spectra can be ``harder'' compared to the single stars. \\begin{figure}[H] \\begin{center} \\includegraphics[width=0.7\\columnwidth]{oskinova-fig1.ps} \\caption{High-resolution X-ray spectra of $\\zeta$\\,Pup {\\it (top panel)} and $\\zeta$\\,Ori {\\it (bottom panel)} obtained with {\\em XMM-Newton}. Prominent emission lines are identified.} \\label{fig:der} \\end{center} \\end{figure} Another type of X-ray emitting plasma is found in high-mass X-ray binary systems (HMXBs) that consist of an OB supergiant and a degenerate companion, usually a neutron star (NS) deeply embedded in the stellar wind of the primary.  Stellar wind accretion onto a NS powers strong X-ray emission that has a characteristic power-law spectrum, while the surrounding stellar wind is photoionized by this X-ray emission.  Independent on the formation mechanism, the X-rays registered by the observer are attenuated in the stellar wind and interstellar medium. The transport of X-rays in stellar winds can often be treated as pure absorption (Baum et al.\\ \\cite{baum92}, Hillier et al.\\ \\cite{hil93}). This leads to a great simplification of the radiative transfer problem and allows modeling even in the case when the absorbing medium is inhomogeneous. ",
        "conclusions": "The spectral and temporal properties of the X-ray emission from single stars, colliding wind binaries and high-mass X-ray binaries provide consistent evidence of wind clumping. The multitude of data requires strongly clumped winds where clumps are separated by a few tenths of the stellar radius in the wind acceleration zone. The approximation that clumps are optically thin at {\\em all} wavelengths in the X-ray band is not justified and cannot be applied universally. Clumped stellar winds have a reduced opacity for X-rays, compared to a homogeneous wind, due to a reduction of empirically determined mass-loss rates and to the effect of wind porosity."
    },
    "0806/0806.0517_arXiv.txt": {
        "abstract": "{In the recent literature there has been some doubt as to the reliability of CO multi-transitional line observations as a mass-loss-rate estimator for AGB stars.}{Using new well calibrated CO radio line observations, the main aim of the work presented here is to carefully evaluate the reliability of CO mass-loss-rate estimates for intermediate- to high-mass-loss-rate AGB stars with different photospheric chemistries.}{Mass-loss rates for 10 intermediate- to high-mass-loss-rate AGB stars are derived using a detailed non-LTE, non-local radiative transfer code based on the Monte-Carlo method to model the CO radio line intensities. The circumstellar envelopes are assumed to be spherically symmetric and formed by constant mass-loss rates. The energy balance is solved self-consistently and the effects of dust on the radiation field and thermal balance are included. An independent estimate of the mass-loss rate is also obtained from the combination of dust radiative transfer modelling with a dynamical model of the gas and dust particles.}{We find that the CO radio line intensities and shapes are successfully reproduced for the majority of our objects assuming a constant mass-loss rate. Moreover, the CO line intensities are only weakly dependent on the adopted micro-turbulent velocity, in contrast to recent claims in the literature. The two methods used in the present work to derive mass-loss-rates are consistent within a factor of $\\sim$\\,3 for intermediate- to high-mass-loss-rate objects, indicating that this is a lower limit to the uncertainty in present mass-loss-rate estimates. We find a tentative trend with chemistry. Mass-loss rates from the dust/dynamical model are systematically higher than those from the CO model for the carbon stars and vice versa for the M-type stars. This could be ascribed to a discrepancy in the adopted CO/H$_{2}$-abundance ratio, but we caution that the sample is small and systematic errors cannot be excluded.}{} ",
        "introduction": "\\label{s:intro} Mass loss is the single most important process during the final evolution of low- and intermediate-mass stars on the asymptotic giant branch (AGB) \\citep{bloe95}. The mass loss determines the lifetime on the AGB, and thus the maximum luminosity and radius reached \\citep[e.g., ][]{marigira07}. Furthermore, since AGB stars are major contributors of mass to the interstellar medium \\citep{sedl94,dorshenn95}, and to the integral luminosity of intermediate-age stellar systems, an understanding of the mass-loss phenomenon is also important for studying extra-galactic populations \\citep{battdeme04} and for the chemical evolution of galaxies [\\cite{schrsedl01}; see also \\cite{karalatt07} for a recent paper on yields from AGB stars]. The existence of mass loss is well established \\citep{will00}, but its details are less well known. This is unfortunate since even a modest mass-loss-rate change, e.g. by a factor of two, will have a profound effect on the evolution of the star, its nucleosynthesis, and its mass return to the interstellar medium \\citep{forechar97}.  Although mass loss on the AGB can reach rates as high as 10$^{-4}$\\,M$_{\\sun}$\\,yr$^{-1}$, the wind speed is quite low ($\\leq$\\,30\\,km\\,s$^{-1}$). The temperature of the outer parts of the atmosphere is low enough for dust to condense. These facts, combined with a high stellar luminosity ($\\sim$\\,10$^{4}$\\,L$_{\\odot}$), have lead to the conclusion that the winds of AGB stars are, at least partly, dust-driven, i.e., the mass-loss from the star is driven by radiation pressure on grains and the gas is dragged away from the star by momentum-coupling between the grains and the gas \\citep[e.g.,][]{lamecass99}. The effects of the gas-dust interaction on the mass loss and wind formation have been studied by several authors \\citep[e.g.,][]{habi96,simietal01}. Time-dependent dynamical models studying the effects of stellar pulsations \\citep{bowe88}, and further including detailed time-dependent dust formation \\citep{hofnetal95}, have progressed to dynamical models in which the motion of the gas and dust particles is calculated and the radiative transfer is solved \\citep{sandhofn03}. These models work well for carbon stars, and support the dust-driven wind scenario, but the formation of silicate grains close enough to the star to drive mass-loss rates of the observed magnitude still poses a problem for M-type stars \\citep{woit06,hofnande07}.  Most likely, the average mass-loss rate increases as the star evolves along the AGB \\citep{habi96,schoolof01}. In addition to this steady increase, modulations of the mass-loss rate on shorter time scales are present \\citep{habi96}. Indications of this can be seen in observations of circumstellar CO line emission around some AGB stars with double-component line profiles \\citep{knapetal98,wintetal03}, and in observations of stellar and ambient Galactic light scattered by dust in circumstellar envelopes (CSEs) \\citep{maurhugg00,maurhugg01,maurhugg06}. A process which can have a strong effect on the mass-loss properties of an AGB-star is the He-shell flash (or thermal pulse) \\citep{olofetal93,olofetal00,stefscho00,schoetal05b}. Other processes may be important during the very last stages when the star sheds the remaining few tenths of a solar mass of its hydrogen envelope (during a few hundred years) before it leaves the AGB \\citep{skinetal97}. Finally, the presence of a binary companion may influence the magnitude as well as the characteristics (e.g., the geometry) of the mass loss.  To make progress in the study of AGB stars, it is important to establish reliable mass-loss-rate estimators, which preferably are useable over a wide range of mass-loss rates and time scales, and for different chemical types (M-, S-, and C-stars). Estimating AGB mass-loss rates through observations of circumstellar CO radio line emission, in combination with modelling of the radiative transfer, is considered to be one of the most reliable methods. CO line emission has also been used in a number of studies aimed at tracing the mass-loss history \\citep{schoetal02,kempetal03,teysetal06,decietal07}, and at identifying drastic mass-loss-rate changes at the end of the AGB \\citep{hesketal90,justetal96}. Detailed studies of the mass-loss properties of carbon stars \\citep{schoolof01,schoetal02}, M-type AGB stars \\citep{olofetal02,delgetal03}, and S-type AGB stars \\citep{ramsetal06} have been performed on samples dominated by low-mass-loss-rate objects. It was concluded that for this type of objects the CO radio line emission is a good mass-loss-rate estimator, although not entirely free of problems.  It has been known for some time that for the higher mass-loss-rate objects the situation is less clear \\citep{saha90,kast92,schoolof01}. A higher mass-loss rate leads to a decrease in the drift velocity between the dust and the gas and therefore a less efficient heating of the gas through collisions with dust grains. Since CO line emission is also a dominant coolant, the effect of an increased mass-loss rate is increased cooling of the circumstellar gas. This, in combination with saturation of the spectral lines, leads to a situation where the integrated CO line intensities are less dependent on the mass-loss rate when it exceeds about 10$^{-5}$\\,M$_{\\odot}$\\,yr$^{-1}$; the higher the frequency of the lines the lower the mass-loss rate when the effect sets in. This makes CO line emission a questionable mass-loss-rate estimator at mass-loss rates in excess of 10$^{-5}$\\,M$_{\\odot}$\\,yr$^{-1}$, an unfortunate effect since we know that mass-loss rates up to 10$^{-4}$\\,M$_{\\odot}$\\,yr$^{-1}$ are reached during the very final stage, and this is also when the stars make their largest contribution to the return of nuclear-processed material to the interstellar medium. In addition, \\citet{kempetal03} concluded that standard circumstellar radiative transfer models could not explain some CO line intensity ratios, and that the mass-loss-rate estimate is sensitively dependent on the (unknown) local turbulent line width.  An alternative way to estimate the mass-loss rate, not subject to the same problems, is to combine modelling of the dust radiative transfer with solving the coupled equations of motion for the gas and the dust particles \\citep{habietal94}. The gas expansion velocity is a strong function of the dust-to-gas mass-loss-rate ratio and therefore a solution of the kinematical equations that fits the gas expansion velocity and the dust optical depth, obtained through dust radiative transfer modelling, results in an independent estimate of the mass-loss rate.  Presented here are a detailed radiative transfer modelling of primarily new, well-calibrated, and high signal-to-noise ratio observations of CO  $J$\\,=\\,1\\,$\\rightarrow$\\,0, 2\\,$\\rightarrow$\\,1, 3\\,$\\rightarrow$\\,2, 4\\,$\\rightarrow$\\,3, and 6\\,$\\rightarrow$\\,5 emission from a sample of 10 well-studied high-mass-loss-rate AGB stars, and the resulting mass-loss-rate estimates. In addition, dust emission modelling combined with a dynamical model give independent estimates of the mass-loss rates. The dust emission modelling also provides the dust information necessary for the CO modelling, e.g., the dust radiation field. By obtaining the mass-loss rates for the stars using two independent methods, we can evaluate the reliability of the mass-loss-rate estimators, and also look for trends with e.g. mass-loss rate and chemistry. A similar investigation comparing mass-loss-rates derived from ISO-SWS observations of the silicate dust features with mass-loss rates from CO modelling found in the literature on a sample of low-mass-loss-rate M-type AGB stars has been performed by \\citet{herahony05}.  \\begin{table} \\caption{The sample with spectral types, variable types, periods ($P$), luminosities ($L_{\\star}$), and distances ($D$).} \\label{sample} $$ \\begin{array}{p{0.3\\linewidth}ccccccc} \\hline \\noalign{\\smallskip} \\multicolumn{1}{l}{{\\mathrm{Source}}} & \\multicolumn{1}{c}{{\\mathrm{Spec.}}} & \\multicolumn{1}{c}{{\\mathrm{Var.}}} && \\multicolumn{1}{c}{{\\mathrm{P}}} & \\multicolumn{1}{c}{L_{\\star}} && \\multicolumn{1}{c}{D}\\\\ & \\multicolumn{1}{c}{\\mathrm{Type}} & \\multicolumn{1}{c}{\\mathrm{Type}} && \\multicolumn{1}{c}{[\\mathrm{days}]} & \\multicolumn{1}{c}{[\\mathrm{L}_\\odot]} && \\multicolumn{1}{c}{[\\mathrm{pc}]} \\\\ \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} \\object{LP And} & \\mathrm{C} & \\mathrm{Mira} &&  620 & \\phantom{0}8200 && \\phantom{1}620 \\\\ \\object{CW Leo} & \\mathrm{C} & \\mathrm{Mira} && 630 & \\phantom{0}8300 && \\phantom{1}120 \\\\ \\object{RW LMi} & \\mathrm{C} & \\mathrm{SRa} && 640 & \\phantom{0}8500 && \\phantom{1}390 \\\\ \\object{V384 Per} &  \\mathrm{C} & \\mathrm{Mira} && 535 & \\phantom{0}7000 && \\phantom{1}590 \\\\ \\object{AFGL 3068} & \\mathrm{C} & \\mathrm{Mira} && 696 & \\phantom{0}9300 && 1090 \\\\ & & & && && \\\\ \\object{TX Cam} & \\mathrm{M} & \\mathrm{Mira} && 557 & 11900 && \\phantom{1}440 \\\\ \\object{GX Mon} & \\mathrm{M} & \\mathrm{Mira} && 527 & 11100 && \\phantom{1}630 \\\\ \\object{WX Psc} & \\mathrm{M} & \\mathrm{Mira} && 660 & 14600 && \\phantom{1}720 \\\\ \\object{IK Tau}\t& \\mathrm{M} & \\mathrm{Mira} && 500 & 10400 && \\phantom{1}300 \\\\ \\object{IRC--10529}\t& \\mathrm{M} & \\mathrm{Mira} && 680 & 15100 && \\phantom{1}890 \\\\ \\noalign{\\smallskip} \\hline \\end{array} $$ \\end{table}  The sample and the observational data are presented in Sects~\\ref{s:sample} and \\ref{s:obsdata}. In Sect.~\\ref{s:circum} the adopted circumstellar model is discussed. In Sect.~\\ref{s:COmod} the CO radiative transfer model is described and CO as a mass-loss rate estimator is discussed. In Sect.~\\ref{s:dynmod}, we describe how dust radiative transfer and the dynamical model is used to estimate mass-loss rates. The results are presented in Sect.~\\ref{s:res}, and outlined for each individual object in Sect.~\\ref{s:ind}. In Sect.~\\ref{s:discuss} we discuss and compare the two methods, and in Sect.~\\ref{s:conc} we give our conclusions. ",
        "conclusions": "  \\begin{itemize} \\item{The mass-loss rate estimates obtained using the two different methods agree within a factor of $\\sim$\\,3, suggesting that this is (at least) the uncertainty in present mass-loss-rate estimates, e.g., using CO radio line emission. The average of the ratio between the two mass-loss-rate estimates is $\\langle\\dot{M}_{\\rm{dyn}}/\\dot{M}_{\\rm{CO}}\\rangle=2.3\\pm1.0$ for the carbon stars, and $\\langle\\dot{M}_{\\rm{dyn}}/\\dot{M}_{\\rm{CO}}\\rangle=0.8\\pm0.5$ for the M-type stars (excluding GX Mon).} \\item{In most cases, it is not possible to set an upper limit to the mass-loss rate for the high-mass-loss-rate objects by modelling only the integrated CO line intensities. The analysis can be improved by taking the shape of the line profile into account.} \\item{We are able to successfully model the CO($J$\\,=\\,1\\,$\\rightarrow$\\,0 to 4\\,$\\rightarrow$\\,3) data of all stars using a constant mass-loss rate. However, in most cases there are various indications that our 'simple' circumstellar model does not fully reproduce the data. Effects of time-variable mass loss, and non-isotropic mass loss may be present.} \\item{We do not find any trends with mass-loss rate when comparing the two models, but we find a possible dependence on chemical type. Generally, the CO model estimates a lower mass-loss rate for the carbon stars than the dynamical model, while the opposite behaviour can be observed for the M-type stars. We conclude that a CO/H$_{2}$-ratio increased and decreased by about a factor of 1.5 for the M- and C-stars, respectively, would give a very good agreement between the methods, but caution that this is a tentative conclusion considering that the sample is small, and there may be other systematic errors.} \\end{itemize}"
    },
    "0806/0806.3208_arXiv.txt": {
        "abstract": "{Optical observations from 1989 of the Small Magellanic Cloud (SMC) B[e] supergiant star LHA\\,115-S\\,23 (in short: S\\,23) revealed the presence of photospheric He{\\sc i} absorption lines, classifying S\\,23 as a B8 supergiant. In our high-resolution optical spectra from 2000, however, we could not identify any He{\\sc i} line. Instead, the spectral appearance of S\\,23 is more consistent with the classification as an A1 supergiant, maintaining the so-called B[e] phenomenon.} {The observed changes in spectral behaviour of S\\,23 lead to different spectral classifications at different observing epochs. The aim of this research is, therefore, to find and discuss possible scenarios that might cause a disappearance of the photospheric He{\\sc i} absorption lines within a period of only 11 years.} {From our high-resolution optical spectra, we perform a detailed investigation of the different spectral appearances of S\\,23 based on modern and revised classification schemes. In particular, we derive the contributions caused by the interstellar as well as the circumstellar extinction self-consistently. The latter is due to a partly optically thick wind. We further determine the projected rotational velocities of S\\,23 in the two epochs of spectroscopic observations. } {Based on its spectral appearance in 2000, we classify S\\,23 as A1\\,Ib star with an effective temperature of about 9000\\,K. This classification is supported by the additional analysis of the photometric $UBV$ data. An interstellar extinction value of $E(B-V) \\simeq 0.03$ is derived. This is considerably lower than the previously published value, which means that, if the circumstellar extinction due to the stellar wind is neglected, the interstellar extinction, and hence the luminosity of the star, are overestimated. We further derive a rotation velocity of $\\varv\\sin i \\simeq 150$\\,km\\,s$^{-1}$, which means that S\\,23 is rotating with about 75\\% of its critical speed. The object S\\,23 is thus the fourth B[e] supergiant with confirmed high projected rotational velocity. The most striking result is the apparent cooling of S\\,23 by more than 1500\\,K with a simultaneous increase of its rotation speed by about 35\\% within only 11\\,years. Since such a behaviour is excluded by stellar evolution theories, we discuss possible scenarios for the observed peculiar variations in S\\,23. } {} ",
        "introduction": "B[e] stars belong to one of the puzzling phenomena in modern astrophysics. While many of their members have been known for decades, the number of new detections of B[e] stars is growing tremendously (see, e.g., Kraus \\& Miroshnichenko \\cite{KM}). The classification by Lamers et al. (\\cite{Lamers98}) of B[e] stars according to their evolutionary phase made clear that B[e] stars occur in the pre- as well as in the post-main sequence phase, and the most popular class is formed by the B[e] supergiants with confirmed members predominantly in the Magellanic Clouds (see the recent review by Zickgraf \\cite{Zickgraf2006}).  The derivation of stellar parameters of B[e] stars is complicated because in most cases the stellar photospheres are hidden inside the dense and optically thick winds. Thus, the spectra are contaminated with wind continuum emission in the form of free-free and free-bound emission as well as with an enormous amount of emission lines that are generated in the circumstellar wind material, which is usually non-spherically symmetric. Evidence for non-sphericity thereby comes from polarimetric observations, e.g., by Magalh\\~aes (\\cite{Magalhaes}), Magalh\\~aes et al. (\\cite{Magalhaesetal}), and Melgarejo et al. (\\cite{Melgarejo}), in rare cases even from optical imaging with the Hubble Space Telescope, as was the case for the galactic B[e] star Hen\\,2-90 (Sahai et al. \\cite{Sahai02}; Kraus et al. \\cite{Kraus05}), and from interferometric observations with, e.g, the VLTI/AMBER and VLTI/MIDI instruments (e.g., Domiciano de Souza et al. \\cite{Armando}). An observed strong infrared (IR) excess emission due to warm and hot circumstellar dust (e.g., Zickgraf et al. \\cite{Zick85}, \\cite{Zickgraf86}, \\cite{Zickgraf89}, \\cite{Zick92}; Gummersbach et al. \\cite{Gummersbach}; Borges Fernandes et al. \\cite{Marcelo}) indicates the presence of a high-density disk. This disk must further be the location of the neutral material like O{\\sc i} (Kraus et al. \\cite{Kraus07}) and the molecular gas, expressed by CO first-overtone bands (McGregor et al. \\cite{McGregorA}; \\cite{McGregorB}; \\cite{McGregor}; Morris et al. \\cite{Morris}; Kraus et al. \\cite{CO}) and TiO band emission (Zickgraf et al. \\cite{Zickgraf89}). The proposed scenario for B[e] supergiants thus suggests that their circumstellar material consists of a disk-like structure formed by a high density equatorial wind with low outflow velocity, and a normal B-type line-driven wind of high velocity and much lower density in polar directions (Zickgraf et al. \\cite{Zick85}; Zickgraf \\cite{Zickgraf2006}). A possible mechanism for the formation of outflowing high density disks might be rapid rotation of the central star (see, e.g., Bjorkman \\& Cassinelli \\cite{Bjorkman}; Cur\\'e \\cite{Cure04}; Cur\\'e et al. \\cite{Cure05}). In addition, if the star is rotating rapidly, the wind material in the equatorial plane recombines in hydrogen in the vicinity of the stellar surface, as has recently been shown by Kraus (\\cite{Kraus06}). And for at least three B[e] supergiants projected rotational velocities at a significant fraction of their critical speed have indeed been found (Gummersbach et al. \\cite{Gummersbach}; Zickgraf \\cite{Zickgraf2000}; \\cite{Zickgraf2006}), hinting to the possible connection between the presence of a circumstellar disk and stellar rotation.  For the Small Magellanic Cloud (SMC) star \\object{LHA 115-S 23} (AzV 172, 2MASS J00555380-7208596, in the following refered to as S\\,23) that is investigated in this paper, no rotational velocity has been derived yet. In addition, not much is known about the orientation of the system, although it has been suggested to be seen more or less edge-on, based on the low expansion velocity derived from the Balmer lines (Zickgraf et al. \\cite{Zick92}, referred to as ZSW in the following). This object was first classified by Azzopardi \\& Vigneau (\\cite{AV}) as a B8 supergiant, and the detection of several photospheric He{\\sc i} absorption lines in the optical mid-resolution spectrograms of ZSW is in agreement with the classification of S\\,23 as a late B-type supergiant. These authors further fixed the interstellar extinction at $E(B-V) = 0.10$, derived an effective temperature of $T_{\\rm eff} = 11\\,000$\\,K, and a surface gravity of $\\log g = 2.0$. This led to $M_{\\rm bol} = -6.4, \\log L_*/L_{\\odot} = 4.46, R_* = 45\\,R_{\\odot}$, and the assignment of the luminosity class Ib. The spectra of ZSW showed further strong Balmer line emission, [O{\\sc i}] emission, as well as many Fe{\\sc ii} and [Fe{\\sc ii}] emission lines. In combination with the pronounced infrared excess in the photometric data, which can be ascribed to warm or hot circumstellar dust, ZSW assigned S\\,23 to the group of B[e] supergiants.  About eleven years later, during our observing mission of Magellanic Cloud B[e] supergiants, we observed S\\,23 using a high-resolution optical spectrograph. Our spectra confirm the presence of the strong Balmer emission lines as well as of the [O{\\sc i}], [Fe{\\sc ii}], and Fe{\\sc ii} emission lines. But interestingly, even with a more than 10 times higher resolution, our spectra do not display any He{\\sc i} line either in absorption nor in emission. The absence of the  He{\\sc i} lines speaks more in favour of a cooler central object than a late B supergiant. This paper is therefore aimed to investigate in detail the spectral appearances of S\\,23 at the two epochs of observations and to discuss possible scenarios that might cause the observed peculiar variations of S\\,23.  The paper is structured as follows: our observations are described in Sect.\\,\\ref{obs}. Then, we perform the spectral classification of S\\,23 for the two different observing epochs, using modern, improved classification schemes (Sect.\\,\\ref{specclas}) and determine the contributions of both, interstellar and circumstellar extinctions (Sect.\\,\\ref{extinc}). From the detected photospheric Mg{\\sc ii} absorption lines, we next derive the projected rotation velocity of S\\,23 at the two epochs of observations (Sect.\\,\\ref{rotation}). In Sect.\\,\\ref{discus}, we discuss possible scenarios that might explain the observed peculiarities, before we summarise our major results in Sect.\\,\\ref{concl}. ",
        "conclusions": "We investigated the SMC star S\\,23, which was formerly classified by ZSW as a B8\\,Ib star, belonging to the class of B[e] supergiants. Our optical spectra show clear differences to the previously published one. These differences especially concern the absence of all photospheric He{\\sc i} absorption lines. Our spectra are therefore consistent more with a spectral classification of S\\,23 as A1\\,Ib supergiant. This classification is further supported by our analysis of the photometric $UBV$ data. During this analysis, we could derive self-consistently the contributions of the interstellar extinction, and of the circumstellar extinction caused by the stellar wind. A circumstellar extinction caused predominantly by dust could be excluded.  We further determined the projected rotation velocity of S\\,23 from an analysis of its photospheric Mg{\\sc ii} line. We found a value of $\\varv\\sin i \\simeq 150$\\,km\\,s$^{-1}$, which claims that S\\,23 is rotating with at least 75\\% of its critical velocity. Consequently, S\\,23 is the fourth member of the B[e] supergiant star group with confirmed high projected rotational velocity.  The most puzzling result is the fact that within a period of only 11 years, between the two sets of spectroscopic observations, the star seems to have cooled by more than 1500\\,K, and at the same time increased its rotation speed by about 35\\%. Such behaviour is definitely ruled out by any stellar evolution scenario. And also the scenario of a shell phase, as is often reported for classical Be stars, has severe difficulties in explaining all the observed peculiarities in S\\,23. Therefore, we presented and discussed several other alternatives: a rapidly rotating star undergoing a high precession, the scenario of an eclipsing binary, and the idea of He abundance and/or temperature inhomogeneities on the stellar surface, as they have been reported from CP stars. All these suggestions can explain most of the curiosities of S\\,23. However, the question, which of these scenarios is the correct one, or whether even all scenarios have to be discarded, could not be solved. Nevertheless, we suggested observational tests that can be performed to distinguish or discard the different scenarios. Definitely more observations with a much better time resolution are needed to disentangle the mystery of the peculiar variations in S\\,23."
    },
    "0806/0806.1129.txt": {
        "abstract": "We present high-speed, three-colour photometry of seven short period (P$_{\\rm orb} \\le 95$ mins) eclipsing CVs from the Sloan Digital Sky Survey. We determine the system parameters via a parametrized model of the eclipse fitted to the observed lightcurve by $\\chi^2$ minimization.  Three out of seven of the systems possess brown dwarf donor stars and are believed to have evolved past the orbital period minimum. This is in line with predictions that 40--70 per cent of CVs should have evolved past the orbital period minimum. Therefore, the main result of our study is that the missing population of post-period minimum CVs has finally been identified. The donor star masses and radii are, however, inconsistent with model predictions; the donor stars are approximately 10 per cent larger than expected across the mass range studied here. One explanation for the discrepancy is enhanced angular momentum loss (e.g. from circumbinary discs), however the mass-transfer rates, as deduced from white dwarf effective temperatures, are not consistent with enhanced angular momentum loss. We show it is possible to explain the large donor radii without invoking enhanced angular momentum loss by a combination of geometrical deformation and the effects of starspots due to strong rotation and expected magnetic activity.  Choosing unambiguously between these different solutions will require independent estimates of the mass-transfer rates in short period CVs.  The white dwarfs in our sample show a strong tendency towards high masses. We show that this is unlikely to be due to selection effects. The dominance of high-mass white dwarfs in our sample implies that erosion of the white dwarf during nova outbursts must be negligible, or even that white dwarfs grow in mass through the nova cycle. Amongst our sample there are no Helium core white dwarfs, despite predictions that 30--80 per cent of short period CVs should contain Helium core white dwarfs. We are unable to rule out selection effects as the cause of this discrepancy. ",
        "introduction": "\\label{sec:introduction} Cataclysmic variable stars (CVs) are a class of interacting binary system undergoing mass-transfer, usually via a gas stream and accretion disc, from a Roche-lobe filling secondary to a white dwarf primary. A bright spot is formed at the intersection of the disc and gas stream, giving rise to an `orbital hump' in the lightcurve at phases $0.6-1.0$ due to foreshortening of the bright-spot. For a basic review of CVs, see \\cite{warner95a}. The secular evolution of CVs represents a long-standing problem in astrophysics, with wide ranging implications for all close binary systems \\cite[see][for example]{kolb93}. Close binary star evolution is driven by angular momentum loss, sustaining mass-transfer which in turn results in changes in the binary's orbital period. Since the orbital period is easy to determine,  it is accurately known for many CVs, and the aim of CV evolutionary theory has been to reproduce the observed period distribution. Salient features which must be explained  are the period gap -- a dearth of systems with periods between 2 and 3 hours, and the  period minimum -- a sharp cut-off in the period distribution at 76.2 mins \\citep{knigge06}. The period gap is thought to be caused by a sudden reduction in angular momentum loss when the donor star becomes fully convective. The cause of this reduction is the cessation of magnetic braking \\citep[e.g][]{robinson81}. Whilst observational evidence for a dramatic change in magnetic braking rate at the convective boundary exists \\citep{delfosse98,reiners07}, both the theoretical justi\u00decation \\citep{tout92} and the observational basis \\citep{andronov03} for the so-called disrupted magnetic braking model have been challenged. Despite this, disrupted magnetic braking remains the best explanation to date for the period gap\u00d5s existence.  The period minimum is caused by the response of the donor star to ongoing mass loss \\citep[e.g][]{paczynski81}. Typically, the donor shrinks in response to mass loss, causing the orbital period to decrease. When the donor nears the substellar limit, however, two effects come into play. Firstly, the thermal timescale  of the donor exceeds the mass-transfer timescale, and the donor cannot shrink rapidly  enough in response to mass loss. Secondly, for substellar donors, changes in the internal structure mean that the donor expands in response to mass loss. Both these effects mean that continuing mass loss results in an increasing orbital period. Thus there exists a period minimum, the exact value of which depends upon the interplay between the thermal timescale of the donor, and the mass-transfer rate. Furthermore, since the number density of CVs, at a given period, scales inversely with the rate at which its period evolves, a significant ``spike'' of systems is expected to accumulate at the minimum period.  Significant problems exist with the explanation of the period minimum. Firstly, the observed minimum period is around 10 minutes longer than the predicted value \\cite[see][for a review]{kolb02}. Also, the predicted period spike is not seen in the orbital period distribution. Moreover, whilst significant numbers of CVs should have passed the period minimum, with the exact figure ranging from around 70 per cent \\citep{kolb93,howell97} to nearer 40 per cent \\citep{willems05}. Despite this, it has proved very hard to identify these so-called \u00d2post-bounce\u00d3 systems \\citep{littlefair03,patterson05}. Until recently, it was not known whether this reflected a genuine absence of post-bounce systems, or a selection effect, caused by the insensitivity of CV surveys to low mass-transfer rate systems \\citep{kolb99}, and the   difficulty of observing very low mass  donor stars against the backdrop of the accretion disc and white dwarf \\citep{littlefair03}.  Recent developments have allowed both these shortcomings to be overcome. The  Sloan Digital Sky Survey (SDSS) \\citep{york00} goes much fainter than previous surveys, and as objects are selected on the basis of their spectra, CVs need not show outbursts  to be included. The SDSS sample could therefore contain a large number of post-period minimum systems. Furthermore, the availability of {\\sc ultracam} \\citep{dhillon07} means that it is possible to obtain high quality lightcurves of the eclipses of faint CVs within the SDSS. This is significant because the lightcurves of eclipsing CVs allow the system parameters to be determined to a high degree of precision \\citep{wood86a}, even when the donor star is not directly visible. Taking advantage of this technique, \\cite{littlefair06} recently presented the first robust identi\u00decation of a post-bounce CV. Here we expand upon those results, by determining the system parameters of a small sample of SDSS CVs with orbital periods below 95 minutes.  Our sample consists of seven systems discovered in the SDSS.  The systems selected are SDSS J090350.73+330036.1, SDSS J103533.03+055158.4, SDSS J122740.83+513925.0, SDSS J143317.78+101123.3, SDSS J150137.22+550123.3, SDSS J150240.98+333423.9 and SDSS J150722.30+523039.8 (hereafter SDSS 0903, 1035, 127, 1433, 1501, 1502 and 1507, respectively). The systems were flagged as high inclination by the presence of broad, double peaked, emission lines and were subsequently found to be eclipsing by follow-up observations. SDSS 1433, 1501, 1502 and 1507 were reported as eclipsing in the discovery papers \\citep{szkody04,szkody05,szkody06,szkody07}. SDSS 1035 was found to be eclipsing from time-resolved spectroscopy \\citep{southworth06}. Amateur observations in outburst revealed SDSS 1227 to be eclipsing \\citep{shears07}, whilst \\cite{dillon08} discovered eclipses in SDSS 0903. Our sample consists of all the eclipsing dwarf novae within the SDSS, with orbital periods below 95 minutes, which were known of and visible at the time of observations. As such, selection effects within our sample should be minimal. Analysis of, and results from, SDSS 1035 and SDSS 1507 are presented in \\cite{littlefair06} and \\cite{littlefair07}, respectively. In this paper, we present the analysis of the remaining systems, and discuss the implications of our results for the evolution of short period cataclysmic variables. ",
        "conclusions": "We present high-speed, three-colour photometry of a small sample of short period, eclipsing CVs taken from the Sloan Digital Sky Survey. We determine the system parameters via a parametrized model of the eclipse fitted to the observed lightcurve by $\\chi^2$ minimization.  Three out of the seven systems possess brown dwarf donor stars and have thus evolved past the orbital period minimum. To the extent that our poor statistics and ignorance of selection effects will allow, this number confirms predictions that $\\sim$40 per cent of CVs should have evolved past the orbital period minimum.  The donor star masses and radii are inconsistent with model predictions, with the majority of donor stars being $\\sim$10 per cent larger than expected across the mass range studied here. One explanation for the discrepancy is enhanced angular momentum loss (e.g. from circumbinary discs), however the mass-transfer rates as  deduced from white dwarf effective temperatures are not consistent with enhanced mass-transfer rates. Alternatively, we find the larger radii can be explained with a combination of the effects of geometrical distortion and starspots. Choosing between these explanations will require better estimates of the mass-transfer rates in these systems. This in turn will require refined estimates of the white dwarf temperature, together with additional work modeling the effect of accretion on the effective temperature of white dwarfs.  The white dwarfs in our sample show a strong tendency towards high masses. We show that this is unlikely to be due to selection effects, and instead is probably a real property of the short period SDSS CVs. The dominance of high mass white dwarfs within our sample implies that the white dwarfs in CVs are not significantly eroded by nova outbursts, or may in fact increase in mass over many nova cycles. Amongst our sample there are no He-core white dwarfs, despite predictions that 30--80 per cent of short period CVs should contain He-core white dwarfs. We are unable to rule out selection effects as the cause of this discrepancy. One white dwarf in our sample may be smaller than expected for a CO-core white dwarf, and possibly has an Fe-core, adding to the small number of Fe-core white dwarf candidates in the literature."
    },
    "0806/0806.4334_arXiv.txt": {
        "abstract": "s{The footprint of inflation in the polarization pattern of the CMB is expected at more than an order of magnitude below the limits of current polarization measurements. Large receiver arrays are mandatory for achieving the required sensitivity. We describe here the approach of the Q/U Imaging ExperimenT (QUIET) which will use coherent detector arrays. Pseudo-correlation receivers have been produced at 40 (90) GHz in small massproducable chip packages. Deployment of the first arrays with 19 (91) receivers is taking place in 2008 in the Atacama Desert in Chile and an expansion to $\\sim$1000 receivers is foreseen for the future. The two frequencies and the selection of observing regions with minimal foreground and good overlap with other upcoming experiments will enable the detection of tensor to scalar ratios $r\\sim10^{-2}$.} ",
        "introduction": "Measurements of the Cosmic Microwave Background Radiation helped significantly in establishing the current cosmological model \\cite{wmap}. The CMB temperature anisotropies have already been measured to a high precision, while only in recent years experiments achieved sensitivities to detect and measure the much smaller polarization anisotropies. The polarization pattern is divided into the E- and B-modes where the former derive from the primordial density fluctuations and the latter from primoridal gravity waves and from the lensing of the E-modes by matter in the line-of-sight. The yet undetected B-mode signal from primordial gravity waves can offer unique insights into the inflationary period \\cite{infl}. Its size is model-dependent, parametrized by the tensor-to-scalar ratio $r=T/S$, but it is expected to be at least one order of magnitude smaller than the E-modes so that a significant increase in sensitivity is required of future experiments. Since current receivers are nearing fundamental limits this leap in sensitivity is only feasible with large receiver arrays. ",
        "conclusions": "The first coherent detector array at 40~GHz will start data taking in mid 2008. With a 90 GHz array following in the winter the sensitivity will allow measurements of the E-mode spectrum at an unprecedented precision. Expansions of the prototype arrays to up to 1000 detectors are foreseen and will provide the sensitivity to access the B-modes from primordial gravity waves down to levels of $r \\sim 10^{-2}$."
    },
    "0806/0806.1985_arXiv.txt": {
        "abstract": "We study the dynamical structure of a cooling dominated rotating accretion flow around a spinning black hole. We show that non-linear phenomena such as shock waves can be studied in terms of only three flow parameters, namely, the specific energy (${\\cal E}$), the specific angular momentum ($\\lambda$) and the accretion rate (${\\dot m}$) of the flow. We present all possible accretion solutions. We find that a significant region of the parameter space in the ${\\cal E}-\\lambda$ plane allows global accretion shock solutions. The effective area of the parameter space for which the Rankine-Hugoniot shocks are possible is maximum when the flow is dissipation free. It decreases with the increase of cooling effects and finally disappears when the cooling is high enough. We show that shock forms further away when the black hole is rotating compared to the solution around a Schwarzschild black hole with identical flow parameters at a large distance. However, in a normalized sense, the flow parameters for which the shocks form around the rotating black holes are produced shocks closer to the hole. The location of the shock is also dictated by the cooling efficiency in that higher the accretion rate (${\\dot m}$), the closer is the shock location. We believe that some of the high frequency quasi-periodic oscillations may be due to the flows with higher accretion rate around the rotating black holes. ",
        "introduction": "The accretion onto black holes is thought to be an essential process  in black hole astrophysics since the gravitational potential energy released by the infalling matter is directly related to the emitted power from the vicinity of a black hole \\citep{fkr02}. In the last three decades, a number of works have reported extensive studies of accretion disk models around gravitating objects by the several groups both theoretically as well as through numerical simulations \\citep{ss73,lb78,eck88,c89,nh94,yk95,c96a,ly97,ft04,nt73,pt74,hetal84a,hetal84b,g99}. In 1980s it was realized that an accreting matter becomes transonic around a galactic and extragalactic black hole (BH) in order to satisfy the inner boundary conditions at the event horizon. A fully general relativistic theoretical work of \\citet{c96b} showed that multi-dimensional sub-Keplerian flows exhibit discontinuous shock transitions for a significant region of the flow parameter space spanned by the specific energy and the specific angular momentum in a non-dissipative flow. Recently, it has been pointed out that the shock waves formed around a BH could be a very useful tool in explaining the observational results. The hot, dense post-shock flow could be the natural site for generating high energy photons \\citep{cm06}. In particular, \\citet{cm06} reported that non-thermal electrons may be processed around the shock fronts by particle acceleration mechanism that could produce emitted radiation up to few MeVs by inverse Comptonization of non-thermal synchrotron radiations. In addition, the observed quasi-periodic oscillations (QPOs) in wide frequency range may be originated due to some type of resonance oscillation of the post-shock flow \\citep{msc96,cam04,otm07}, or due to dynamical oscillations of the sub-Keplerian flows \\citep{rcm97}. This oscillation is governed by the infall time scale of matter in the post-shock region. Since a shock may form very close to a spinning BH, it could possibly explain the origin of high frequency QPOs also. Thus, it would be worthwhile to study the behavior of the shock waves around a rotating BH for more realistic cases by addition of cooling processes inside the disk. This is all more important since the post-shock pressure which holds the shocks in place, depends on the cooling processes.  So far, major progresses on the theoretical front have taken place only when the accretion flow around the rotating BH is thin and when the energy dissipation processes are inefficient \\citep{c96b}. In reality, the accreting matter should be affected by the various cooling processes, such as bremsstrahlung and synchrotron cooling, most significantly the latter as the magnetic field is expected to be present in the disk. Some theoretical work along this line has recently been completed \\citep{d07}. However, the effects of synchrotron cooling on the dynamical structure of the accreting flow around a rotating black hole has not yet been studied well. Indeed, the study of accretion flows with various energy dissipation processes around a compact object in the general relativistic limit is very complex. Recently, \\citet{skcsm06} proposed a pseudo-Kerr potential which describes the space-time geometry around a rotating BH quite satisfactorily for the black hole spin parameter, $a_k$, $-1 \\leq a_k \\leq 0.8$. This pseudo-potential, which is an extension of the concept laid down while writing the Paczynski-Wiita potential \\citep{pw80} in Schwarzschild geometry allows us to study very complex bahaviour of dissipative accreting matter around a rotating black hole in a very simple and elegant way.  In an optically thick and geometrically thin accretion flow, cooling processes reduce the bulk energy of the flow during accretion while viscosity not only heats the flow but also redistributes the angular momentum in the disk. The cooling is governed by the density and temperature of the flow while viscosity is assumed mainly due to ion collisions or its coupling with the magnetic field. A realistic disk can have some viscosity, but as \\citet{cd04,c96a} already pointed out, standing shock exits even in presence of significant viscosity. Thus, in this paper we shall explore the properties of shock waves in presence of cooling alone. In other words, our focus here is to deal with the accretion flows having low angular momentum upto the marginally bound value. However, we argue that basic conclusions regarding our idealized inviscid solution would be qualitatively similar to the viscous solution. We wish to report viscous solutions elsewhere in the future.  The suitability of the pseudo-Kerr potential that we employ has been well tested for both particle dynamics and fluid dynamics \\citep{skcsm06,smskc06} which allows us to use this for our present study with sufficient confidence. We identify all possible accretion solution topologies including the ones which may allow standing shock transitions. In this work, we only focus on thin, non-dissipative shock waves. We separate various regions in the parameter space according to the solution topologies. We identify a special region where standing shock conditions (Rankine-Hugoniot shock conditions [RHCs], \\citealt{ll59}) are not satisfied but accretion solution continues to have three critical points. Solutions of this kind generally exhibit oscillating shock as in \\citet{rcm97} when Schwarzschild geometry was used. In the present paper, we ignore bremsstrahlung cooling process since it is very inefficient compared to the synchrotron loss \\citep{cc00}, especially in flows around stellar mass black holes. We also do not consider Compton cooling effect. These will be carried out in future.  In the next Section, the governing equations are presented. In \\S 3 and \\S 4, the accretion solutions are shown. We discuss the shock properties in \\S 5. In the final Section, we make concluding remarks. ",
        "conclusions": "In this paper, we have studied the properties of cooling dominated accretion flow around a rotating black hole by solving a set of equations that regulate the dynamical structure of the flow. A special consideration is given to synchrotron cooling that strongly affects the disk properties as well as the emitted spectrum and luminosity that are observed.  We obtain the global accretion solutions with and without the shocks in terms of a few flow parameters, namely, energy, angular momentum, black hole rotation parameter and the accretion rate, which effectively acts as the cooling parameter. We find that the accreting matter experiences a centrifugal force which acts as a barrier, inducing a shock formation. We show that the global shocked accretion solution can be obtained for a significant region of the parameter space even when the cooling is significant. Using a conventional accretion disk model we expect the accretion to take place when the angular momentum is close to the marginally stable value. Our calculation shows that the region is actually broader, in terms of both the angular momentum and energy.  The discussion regarding the nature of the sonic point has been reported in many occasions \\citep{cd04,dc04}. However, a detailed analysis was not presented before for a cooling dominated flow around a rotating black hole. Our present work suggests that a large region of the parameter space provides a stable saddle type sonic point. In Fig. 2, we demonstrated that the parameter space for the stable saddle type sonic point is gradually reduced with the increase of cooling efficiency.  We show that the standing shocks form closer to a spinning black hole as the accretion rate is enhanced. At the post shock region, the density and the temperature is relatively high compared to the pre-shock flow and thus cooling is more efficient there. For a higher cooling, the post-shock matter cools faster reducing the thermal pressure drastically. This forces the shocks to move inward to maintain pressure balance across them.  One of the aims of the present work was to study the effect of black hole rotation parameter on the dynamical structure of cooling dominated global solutions. We find that for flows with identical outer boundary condition ({\\it e. g.,} same energy and angular momentum at the outer edge) shock recedes from the black hole horizon with the increase of black hole rotation parameter ($a_k$). However if we choose the relevant angular momentum for each case, such as the marginally stable angular momentum the shock location moves in with the increase of $a_k$. The range of $a_k$ for which the stationary shocks are formed is restricted for a flow of given angular momentum. Shocks are possible around a rapidly rotating black hole when the flow angular momentum is relatively low. Since that produces a very low centrifugal pressure, the shock can form very close to the black hole for a rapidly spinning black hole.  We identify the region of parameter space for the formation of a standing shock. We find that the effective region of the parameter space for the stationary shock shrinks when the accretion rate is enhanced. This suggests that the possibility of shock formation decreases for higher accretion rate. In addition, we also separate a region where the Rankine-Hugoniot relation is not satisfied. In the context of invidcid flows, it has been observed that the flow parameters from such a region give rise to oscillating shocks \\citep{rcm97}. The reason is that the higher entropy at the inner sonic point forces the flow to pass through it by generating extra entropy at the shock. But since RHCs are not satisfied the shock can not settle itself at a given location. Thus the cause of oscillation is sufficiently generic and we suspect that exactly the same thing will happen in the present case. Most importantly, since rotating black holes may have shocks very close to the horizon, the frequencies of such oscillations are expected to be higher.  Our present findings suggest that shocks, standing or oscillating, do form around the spinning black holes and it may be an essential ingredient since shocks could successfully explain the observed stationary \\citep{cm06} as well as time dependent behaviour of the radiations from the black hole candidates \\citep{cam04,otm07}. We demonstrated that the shocks form closer to the black hole as cooling is increased. This will enhance the QPO frequency as it is proportional to the infall time scale \\citep{cm00,msc96} and thus the QPO frequency may vary in a wide range starting from mHz to KHz depending on the accretion rate. This understanding also generally agrees with the observational results. Recent reporting of the outbursts of GRO 1655-40 showed a clear evidence of the QPO frequencies increasing monotonically from about $90$mHz to $17$Hz \\citep{cd05} in a matter of $~ 15 $ days and this could be easily fitted using the shock propagating at a constant velocity (Chakrabarti et al. 2005).  The formalism presented here does not include outflows/jets which may be generated from the inner part of the disk as a result of deflection of inflowing matter due to excess thermal pressure at the shock front \\citep{dcnc01,cd07,dc08}. Since the outflows/jets are ejected evacuating the inner part of the disk, it will necessarily reduce the post-shock pressure and therefore, shock front has to move in to retain the pressure balance. This suggests that the result should be affected if the accretion-ejection mechanism is considered together. We plan to consider this study in a future work and it will be reported elsewhere.  We have approximated the effect of general relativity using the pseudo-Kerr gravitational potential. This pseudo-Kerr potential has been successfully tested to retain most, if not all, of the salient features of the flows in a Kerr metric. The use of this approach allows us to find out the non-linear shock solutions in a curved space-time geometry in a simpler way. We believe that our basic results would be qualitatively the same with fully general relativistic calculations, especially for $a_k <0.8$ for which the pseudo-Kerr potential was found to be satisfactory."
    },
    "0806/0806.4272_arXiv.txt": {
        "abstract": "The motion of test particles in the gravitational fields generated by the first four members of the infinite family of generalized Kalnajs discs (\\cite{GR}), is studied. In first instance, we analyze the stability of circular orbits under radial and vertical perturbations and describe the behavior of general equatorial orbits and so we find that radial stability and vertical instability dominate such disc models. Then we study bounded axially symmetric orbits by using the Poincar\\'e surfaces of section and Lyapunov characteristic numbers and find chaos in the case of disc-crossing orbits and completely regular motion in other cases. ",
        "introduction": "A fact usually assumed in astrophysics is that the main part of the mass of a typical spiral galaxy is concentrated in a thin disk (\\cite{BT}). Accordingly, the study of the gravitational potential generated by an idealized thin disk is a problem of great astrophysical relevance and so, through the years, different approaches has been used to obtain such kind of thin disk models(see Kuzmin (\\citeyear{KUZ}) and Toomre (\\citeyear{T1,T2}), as examples). A simple method to obtain the surface density, the gravitational potential and the rotation curve of thin disks of finite radius was developed by \\cite{HUN1}, the simplest example of disk obtained by this method being the Kalnajs disk (\\cite{KAL}), which can also be obtained by flattening a uniformly rotating spheroid (\\cite{WM,BR,BB}). In a previous paper (\\cite{GR}) we use the Hunter method in order to obtain an infinite family of thin disks of finite radius, an infinite family of generalized Kalnajs discs with a well behaved surface mass density.  Closely related with the above study is the analysis of the motion of test particles in the gravitational field generated by such disklike distributions of matter. In particular, the study of orbits in the equatorial plane is of clear astrophysical relevance due to its relation with the dynamics of intergalactic stellar motion or the flow of particles in accretion disks around black holes. Also, a knowledge of the disc's internal kinematics is relevant for its subsequent statistical analysis, and the study of external particles motion help us to understand the behavior of stars belonging to the galaxy's remaining component (for example, the halo). Many of them cross back and forth through the disc, experiencing a fairly abrupt change in the gravitational force field. This fact gives rise to a large variety of chaotic and regular orbits, as it was pointed out by Hunter (\\citeyear{Hunter}), in the case of Kuzmin-like discs, and Martinet et al (\\citeyear{Martinet1}; \\citeyear{Martinet2}; \\citeyear{Martinet3}; \\citeyear{Martinet4}), in the case of Schmidt's models (Schmidt \\citeyear{Schmidt}).  In agreement with the above considerations, in this paper we shall focus on the kinematics around the generalized Kalnajs discs, introduced by Gonz\\'alez \\& Reina (\\citeyear{GR}). They form an infinite family of axially symmetric finite thin discs, whose first member is precisely the well-known Kalnajs disc (\\citeyear{KAL}).  The paper is organized as follows. First, in section \\ref{sec:kal}, we present a summary of the main aspects of the generalized Kalnajs discs, the surface densities, the gravitational potential and the motion equations. Then, in section \\ref{sec:equ} we shall focus on equatorial orbits, i.e. $z=0$ trajectories ($\\xi=0$, $\\eta=\\sqrt{1-R^{2}/a^{2}}$ inside the disc and $\\xi=\\sqrt{R^{2}/a^{2}-1}$, $\\eta=0$ outside). At first we shall study the stability under radial and vertical perturbations  of circular orbits, then examining the principal features and conditions of general equatorial orbits. In the next section, section \\ref{sec:cros}, we present numerical solutions of (\\ref{em}) describing some representative Poincar\\'e surfaces of section. As it is expected, we find chaotic sections for disk-crossing orbits and regular sections for other cases. Some meridional plane orbits are plotted and the Lyapunov characteristic numbers (LCN) are calculated. Finally, in section \\ref{sec.con}, we summarize our main results. ",
        "conclusions": "\\label{sec.con}  One remarkable fact, suggested by the analysis \\textbf{performed} in Sect. 2, is that the stability of circular orbits under radial and vertical perturbations increases with the parameter $m$ (for $m\\geq 2$). More larger $m$ corresponds to more stable models, when we deal with circular orbits. Another important fact is that the range of axial angular momentum for which we can find bounded motion inside the disc, decreases with $m$. Thus, loosely speaking, more and more stable models have less and less possibilities to maintain their particles inside the disc. These considerations have special relevance in the search of equilibrium distribution functions (they are $E,\\ell$-dependent) characterizing such galaxy models.  The numerical calculations showed in Sect. 3 confirm the analysis made by Hunter about disc-crossing orbits. There exist a chaotic motion induced by the presence of the disk and, despite the current versions of KAM theorem do not apply (due to the discontinuous force field), there is also a significant range of regular orbits. Moreover, since in our case we deal with finite disc models, a distinction between disc-crossing and non-disc-crossing orbits is sometimes necessary. In the latter case we did not find chaotic motion, even in extreme situations where there are saddle points outside the field source, and the disc and disc-free region are connected (Figures \\ref{Akalnajs1} and \\ref{Akalnajs2}).  Although one would be tented to think that such hyperbolic exterior points can induce chaos at disc-free regions, what really happens is that the stochastic motion tends toward disc regions and a completely regular motion is developed outside there. These considerations have special relevance in galaxy models with thin disc plus halo components. Particles belonging to the halo component will follow a motion with the features mentioned above and, as it was showed by some authors, this fact determines decisively the internal structure of such stellar systems (see Ostriker, Spitzer and Chevalier (\\citeyear{Ostriker}))."
    },
    "0806/0806.3664_arXiv.txt": {
        "abstract": "The standard explanation for large pulsar glitches involves transfer of angular momentum from an internal superfluid component to the star's crust. This model requires an instability to trigger the sudden unpinning of the vortices by means of which the superfluid rotates. This Letter describes a new instability that may play this role. The instability, which is associated with the inertial r-modes of a superfluid neutron star, sets in once the rotational lag in the system reaches a critical level. We demonstrate that our model is in good agreement with observed glitch data, suggesting that the superfluid r-mode instability may indeed be the mechanism that triggers large pulsar glitches. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0806/0806.2042_arXiv.txt": {
        "abstract": "We present measurements of the transmission of radio/microwave pulses through salt in the Cote Blanche salt mine operated by the North American Salt Company in St.~Mary Parish, Louisiana.  These results are from data taken in the southwestern region of the 1500~ft. (457~m) deep level of the mine on our third and most recent visit to the mine. We transmitted and received a fast, high-power, broadband pulse from within three vertical boreholes that were drilled to depths of 100~ft. (30~m) and 200~ft. below the 1500~ft. level using three different pairs of dipole antennas whose bandwidths span 125~to~900~MHz.  By measuring the relative strength of the received pulses between boreholes with separations of 50~m and 169~m, we deduce the attenuation of the signal attributed to the salt medium. We fit the frequency dependence of the attenuation to a power law and find the best fit field attenuation lengths to be $93 \\pm 7$~m at 150~MHz, $63 \\pm 3$~m at 300~MHz, and $36 \\pm 2$~m at 800~MHz. This is the most precise measurement of radio attenuation in a natural salt formation to date. We assess the implications of this measurement for a future neutrino detector in salt. ",
        "introduction": "The observation of cosmic rays with energy higher than the Greisen-Zatsepin-Kuzmin~(GZK) cutoff at $\\sim~10^{19.5}$~eV~\\cite{gzk} implies a corresponding flux of ultra high-energy (UHE) neutrinos, with energy in the $10^{17}-10^{19}$~eV range~\\cite{ess}.  These secondary neutrinos are created via photomeson production of cosmic rays on the 2.7~K cosmic microwave background.  Detection of these neutrinos would provide unique information about the origin of primary cosmic rays and the nature of their sources.  Although predictions for the flux of UHE neutrinos differ by orders of magnitude, a reasonable set of parameters puts the rate of UHE neutrinos at the level of 10/km$^2$/century~\\cite{ess}. A detector volume of hundreds of cubic kilometers water equivalent is required to detect a significant number of neutrinos in this energy regime.  Optical techniques are widely used in neutrino detectors, but the volume that can be instrumented is constrained by attenuation lengths of tens of meters at optical frequencies in detector media (such as ice), so optical detectors have limited sensitivity to rare high-energy events.  G.~Askaryan first predicted coherent radio emission from ultra high-energy showers~\\cite{askaryan}, and the effect has been confirmed in accelerator beam tests~\\cite{sandsaltice}. The energy of the coherent radio emission that results from the development of a negative charge excess in the shower depends quadratically on shower energy, and for showers with energy greater than $10^{16}$~eV, dominates the emitted Cherenkov power spectrum.  Askaryan also proposed a few materials as detection media that occur naturally in very large volumes and are expected to have long attenuation lengths in the radio regime, including ice and salt. Radio attenuation lengths longer than 1~km have been measured in ice near the South Pole~\\cite{southpoleice}.    Formations of salt rock could therefore be a viable detector for high-energy neutrinos if attenuation in the radio regime is indeed low. Domes of rock salt occur naturally in many parts of the world, including the Gulf Coast region of the United States.  In these formations, the salt originates from dried ocean beds which have been buoyed upward due to geological forces through a process called diapirism.  Through this process, the salt becomes very pure as impurities are extruded. Such salt domes, with typical dimensions of several square kilometers by several kilometers deep, are thought to be good candidates for a neutrino detector. For a more detailed discussion of the dielectric properties of salt and the application to neutrino detection, see Reference~\\cite{hockley}. ",
        "conclusions": "We have measured field attenuation lengths in the frequency range 125-900~MHz using broadband pulses transmitted and received by dipole antennas in boreholes drilled 100~ft. into the floor of the 1500~ft. deep level of the the Cote Blanche salt mine.   The best fit power law that describes the measured frequency-dependent attenuation lengths is given by \\begin{equation} L_{\\alpha}=\\left(32\\pm2\\right)\\cdot \\left( \\nu/1~\\mathrm{GHz} \\right)^{-0.57\\pm0.06} \\end{equation} with statistically dominated uncertainties.  This is the most precise measurement of radio attenuation in a natural salt formation to date.  The main result of this paper is derived from data that we took on our third trip to the Cote Blanche Salt Mine.  The data taken from the three trips indicates that the room and pillar system of mining leads to regions of salt near the walls and ceilings of the corridors with short attenuation lengths ($\\sim10$~m) which could be due to fracturing brought about by the explosives used in that method.  We find that the salt that is just under the floor of a corridor does not suffer the same degradation, which may be due to the cut along the floor that is made prior to the detonation of the explosives to blast away the salt.  Although we expect there to be variation in salt properties between salt domes as well as regions of salt within the same dome, we have compared our measurements with previous measurements made at Cote Blanche and elsewhere. Our result is not inconsistent with attenuation lengths measured at the Hockley Salt Mine.  We have attempted to compare our measurements with the longest transmission distance reported by GPR experts in the same mine. Our data generally showed greater attenuation than can be reconciled with that result.  However, due to the variation observed in our data, some of our measurements do point to salt that could have been clear enough to be consistent.  That we strain to reconcile our numbers with the longest transmission distance reported there may indicate that there are regions of the Cote Blanche mine with salt that has less attenuation than in the regions where our measurements were made.  We have modeled a SalSA array of $10 \\times 10 \\times 10$ dipole antennas in boreholes that reach 3~km depth to assess the sensitivity of a neutrino detector in salt with the attenuation lengths that we report here to compare with the simulation study described in the Hockley attenuation length paper. We find the lower best fit attenuation lengths measured at Cote Blanche reduce the sensitivity of such a detector measured in terms of event rate by a factor of 9, but still maintain an expectation of order one event per year using a standard model for the expected ultra-high energy neutrino flux.  In order for a next-generation neutrino detector in salt that measures $\\sim10$~GZK events/year to be feasible, attenuation lengths longer than the $\\sim100$~m lengths reported here will likely be necessary.  However, since this result is based on a limited region of the dome, and past GPR results in the same mine may have pointed clearer salt, there are reasons to be optimistic that there is clearer salt in other locations whose radio loss has not yet been precisely measured.  For this to be pursued further, we believe that developing a portable radar system is a compelling next step so that the radio loss of the salt can be sampled in many locations on a single visit to a mine."
    },
    "0806/0806.3933_arXiv.txt": {
        "abstract": "The available information on isotopic abundances in the atmospheres of low-mass Red Giant Branch (RGB) and Asymptotic Giant Branch (AGB) stars requires that episodes of extensive mixing occur below the convective envelope, reaching down to layers close to the hydrogen burning shell (Cool Bottom Processing).  Recently \\cite{Busso:2007jw} suggested that dynamo-produced buoyant magnetic flux tubes could provide the necessary physical mechanisms and also supply sufficient transport rates.  Here, we present an $\\alpha-\\Omega$ dynamo in the envelope of an RGB/AGB star in which shear and rotation drain via turbulent dissipation and Poynting flux.  In this context, if the dynamo is to sustain throughout either phase, convection must resupply shear.  Under this condition, volume-averaged, peak toroidal field strengths of $\\left<B_\\phi\\right>\\simeq3\\times10^3$ G (RGB) and $\\left<B_\\phi\\right>\\simeq5\\times10^3$ G (AGB) are possible at the base of the convection zone.   If the magnetic fields are concentrated in flux tubes, the corresponding field strengths are comparable to those required by Cool Bottom Processing. ",
        "introduction": "In low-mass Red Giant Branch (RGB) and Asymptotic Giant Branch (AGB) stars, matter must circulate in episodes of extended mixing from the convective zone to the radiative region where nuclear processing can occur (the so-called Cool Bottom Processing; CBP; \\citealt{Wasserburg:1995cj}; \\citealt{Charbonnel:1998hi}; \\citealt{Nollett:2003il}; \\citealt{Herwig:2005eb}).  While the mechanism driving mixing remains under investigation, CBP predicts the chemical and isotopic evolution for a given bulk transport rate by utilizing the path integral of the nuclear reactions over the mass and independent of the mixing mechanism (\\citealt{Boothroyd:1994to}; \\citealt{Wasserburg:1995cj}).  The necessary transport rates and processing temperatures are constrained by RGB/AGB isotopic measurements and presolar meteoritic dust grain abundance ratios (\\citealt{Choi:1998zt}; \\citealt{Amari:2001xw}; \\citealt{Nollett:2003il}; \\citealt{Nittler:2005hq}).  Weak circulation ($\\dot{M}\\sim10^{-6}-10^{-8}$ $M_\\odot$/yr) in which material is transported from the convective envelope to just above the H-burning shell, processed and then returned has had success in matching: (i.) $^{12}$C/$^{13}$C, C and N abundances in RGB stars, (ii.) $^{12}$C/$^{13}$C and N/O in AGB stars and (iii.) $^{26}$Al/$^{27}$Al, $^{18}$O/$^{16}$O, $^{17}$O/$^{16}$O and $^{12}$C/$^{13}$C in circumstellar dust grains found in meteorites that must come from AGB sources (\\citealt{Wasserburg:1995cj}; \\citealt{Nollett:2003il}; \\citealt{Herwig:2005eb}).  Several physical origins have been proposed for the assumed mixing including: shear and thermohaline instabilities, meridional circulation and internal gravity waves (\\citealt{Zahn:1992kb}; \\citealt{Denissenkov:1996ye}; \\citealt{Denissenkov:2003qm}; \\citealt{Chaname:2005hs}; \\citealt{Eggleton:2006fb}; \\citealt{Charbonnel:2007wu}).  Purely rotation-induced mixing appears insufficient as isotopic changes are minimal during the RGB and lead to a quenching of s-process nucleosynthesis during the AGB (\\citealt{Siess:2004oz}; \\citealt{Palacios:2006qa}).  Recently, it was proposed that buoyant magnetic flux tubes can induce the necessary transport rates required for Cool Bottom Processing (\\citealt{Busso:2007jw}). Constraints on the field strengths for the required mixing rates were calculated while the fields were assumed to originate from a dynamo operating in the RGB and AGB interiors.  In particular,the proposed mechanism would provide a means of rapid transport of material to the convective mantle and may provide a solution to the problem of Li rich stars.  Magnetically mediated outflows resulting from dynamo amplification have been proposed as the origin of bipolarity in post-Asympotitc Giant Branch stars (post-AGB) and planetary nebula (PN) (\\citealt{Pascoli:1997kc}; \\citealt{Blackman:2001sy}; \\citealt{Nordhaus:2006oq}; \\citealt{Nordhaus:2007il}).  However, dynamo-induced shaping requires differential rotation.  This can be supplied by binary companions or sustained throughout the AGB phase by convective redistribution of angular momentum (\\citealt{Nordhaus:2007il}).  Convection may resupply differential rotation (analogous to the $\\lambda$-effect in the Sun) and sustain the magnetic fields throughout the RGB/AGB stages (\\citealt{Rudiger:2004zv}; \\citealt{Nordhaus:2007il}).  In particular, it may be that a weaker dynamo (supplying sufficient mixing) is sustained throughout the RGB/AGB phases, but that a binary companion is required to power bipolar jets in post-AGB stars \\citep{Nordhaus:2008os}.  We investigate the origin of large-scale magnetic fields via a dynamical, $\\alpha-\\Omega$ dynamo operating at the base of the convection zone in an initially 1.5 $M_\\odot$ main sequence star during the RGB and AGB phases.  Our stellar models are identical to those used in \\cite{Busso:2007jw}, and allow for direct comparison.  In both phases, subadiabatic zones similar to the solar tachocline exist.  We allow a fraction of the turbulent energy cascade to resupply shear, thus sustaining the dynamo.  The field penetrates the shear zone and may transport material close to the H-burning shell.  We compare our results to the field strengths and penetration depths needed for adequate magnetic mixing in low-mass RGB and AGB stars. ",
        "conclusions": "\\cite{Busso:2007jw} recently suggested that buoyant toroidal fields, may supply the necessary chemical transport rates to circulate material from near the hydrogen burning shell to the convective envelope required by Cool Bottom Processing in low-mass RGB and AGB stars.  In this paper, we have shown how an $\\alpha-\\Omega$ dynamo operating in the interior of an RGB/AGB star could supply the fields necessary to accomplish the needed transport.  The fields diffuse in a magnetic cycle period to the depths required by Cool Bottom Processing.  If the dynamo is to sustain throughout either phase, convection must resupply differential rotation.  This can be accomplished if a fraction of the turbulent energy cascade reinforces shear ($\\leq15\\%$ for the RGB star; $\\leq7\\%$ for the AGB star).  For the model RGB star, mean toroidal field strengths of $\\left<B_\\phi\\right>=2.6\\times10^3$ G are sustainable at the base of the convection zone.  If the magnetic field is concentrated in flux tubes, the volume averaged fields, $\\left<B^2\\right>$, are related to the flux tube fields, $\\left<B_t^2\\right>$, by $\\zeta<1$, such that $\\left<B^2\\right> \\sim \\frac{\\zeta}{1-\\zeta}\\left< B_t^2\\right>$.  For the model RGB star, $\\zeta=0.36$ will satisfy CBP requirements.  For the model AGB star, stronger fields are possible with sustained toroidal field strengths of $\\left<B_\\phi\\right>=4.8\\times10^3$ G possible at the base of the convective zone.  This requires $\\zeta=0.23$ to match CBP flux tube requirements.  We have found that dynamo-produced fields provide a plausible mechanism for the physical origin of extra mixing in low-mass RGB and AGB stars.  However future work, in particular, realistic rotation profiles, the complex relationship between convection and shear, intermittent magnetic mixing, the dynamics of mass transport via flux tubes all need to be studied further.  Equally important is to independently constrain the value of the diffusion coefficients in the shear and convection zones to determine if they are consistent with what our model implies.  Finally, $\\geq$ 2-D studies of this problem are needed."
    },
    "0806/0806.3791_arXiv.txt": {
        "abstract": "We have conducted a deep and uniform 1.1\\,mm survey of the GOODS-N field with AzTEC on the James Clerk Maxwell Telescope (JCMT).  Here we present the first results from this survey including maps, the source catalogue, and 1.1\\,mm number-counts.  The results presented here were obtained from a 245\\,arcmin$^2$ region with near uniform coverage to a depth of 0.96--1.16\\,mJy\\,beam$^{-1}$.  Our robust catalogue contains 28 source candidates detected with S/N $\\ge3.75$, only $\\sim$1--2 of which are expected to be spurious detections.  Of these source candidates, 8 are also detected by SCUBA at 850\\,\\um\\ in regions where there is good overlap between the two surveys.  The major advantage of our survey over that with SCUBA is the uniformity of coverage.  We calculate number counts using two different techniques: the first using a frequentist parameter estimation, and the second using a Bayesian method.  The two sets of results are in good agreement.  We find that the 1.1\\,mm differential number counts are well described in the 2--6\\,mJy range by the functional form $dN/dS = N' (S'/S)\\mathrm{exp}(-S/S')$ with fitted parameters $S' = 1.25\\pm0.38$\\,mJy and $dN/dS = 300\\pm90$\\,mJy$^{-1}$\\,deg$^{-2}$ at 3~mJy. ",
        "introduction": "\\label{intro}  Identifying and studying the galaxies at high redshift that will evolve into today's normal and massive galaxies remains a major goal of observational astrophysics.  Galaxies discovered in deep sub-millimetre and mm-wavelength surveys \\citep[e.g.][]{Smail1997, Hughes1998, Barger1998, Blain1999A, Barger1999A, Eales2000, Cowie2002, sescott2002, Webb2003, Borys2003, Greve2004, Laurent2005} are generally thought to be dominated by dusty, possibly merger-induced starburst systems and active galactic nuclei (AGN) at redshifts $z>2$ with star formation rates as high as SFR $\\sim 1000 ~ \\mathrm{M}_{\\sun} \\mathrm{yr}^{-1}$~\\citep{Blain2002}.  The high areal number density of these sub-mm and mm-detected galaxies (SMGs), combined with their implied high star formation rates and measured FIR luminosities \\citep[$L_{\\mathrm{FIR}}\\sim10^{12} \\mathrm{L}_{\\sun}$,][]{Kovacs2006,Coppin2008}, makes their estimated contribution to both the global star formation density and the sub-mm background radiation as high as 50\\% at $z\\sim2$ \\citep[e.g.,][]{Borys2003,Wall2008}.  Their observed number counts imply strong evolution between $z=2$ and today \\citep[e.g.][]{sescott2002,Greve2004,Coppin2006}.  The high star formation rates at early epochs of SMGs generally match the expectation for rapidly forming elliptical galaxies, a view supported by the high rate of mergers seen locally in samples of ultra-luminous infrared galaxies \\citep[ULIRGs;][]{Borne2000}, which are plausible local counterparts of distant SMGs.  Together, these characteristics have led many observers to surmise that SMGs are likely to evolve into the massive galaxies observed locally \\citep[e.g.,][]{Dunlop1994,Smail1997,Bertoldi2007} and may hold important clues to the processes of galaxy and structure formation in general at high redshift.  GOODS-N is one of the most intensively studied extragalactic fields, with deep multi-wavelength photometric coverage from numerous ground-based and space-based facilities. These include Chandra in the X-ray \\citep{Alexander2003A}, HST in the optical and NIR \\citep{Giavalisco2004}, Spitzer in the NIR--MIR (Chary et al. in prep., Dickinson et al. in prep.), and the Very Large Array in the radio \\citep[][Morrison et al. in prep.]{Richards2000}, as well as highly complete spectroscopic surveys from ground-based observatories \\citep[e.g.][]{Wirth2004,Cowie2004}.  This field is therefore ideally suited for deep mm-wavelength studies of SMGs: the extensive coverage in GOODS-N allows the identification of SMG counterparts in X-ray, UV, optical, IR, and radio bands, as well as constraints on photometric redshifts and investigation of SMG power sources and evolution.  Deep mm surveys of blank fields are needed in order to constrain the faint end of the SMG number counts, while large areal coverage is required to constrain the bright end.  Together they provide strong constraints on evolutionary scenarios.  Previous sub-mm surveys of GOODS-N have been carried out with SCUBA on the JCMT \\citep{Hughes1998,Barger2000,Borys2003,Wang2004,Pope2005}.  The `Super-map' of the GOODS-N field, which was assembled from all available JCMT shifts covering the field, contains 40 robust sources at 850\\,\\um\\ down to an average sensitivity of 3.4\\,mJy ($1 \\sigma$) and covers 200\\,arcmin$^2$ \\citep{Borys2003,Pope2005}.  However, the r.m.s. is highly non-uniform ranging from 0.4 mJy to 6 mJy (see Fig.~\\ref{fig_goodsn_cover}).  That non-uniformity presents serious complications for comparisons with multi-wavelength data. \\begin{figure} \\centering \\includegraphics[width=\\hsize]{cov_regions.eps} \\caption{AzTEC and SCUBA coverage contours for the GOODS-N region demonstrates our uniform coverage.  The dark rectangular contour corresponds to the AzTEC region with a map r.m.s. $\\le$1.16\\,mJy at 1.1\\,mm, the coverage region presented here.  The grey contours, according to increasing line thickness, are the 850\\,\\um\\ SCUBA contours for r.m.s. values of 4\\,mJy, 2.5\\,mJy, and 0.5\\,mJy respectively.  The underlying map is the IRAC 3.6\\,\\um\\ image from the Spitzer legacy program (Dickinson et al. in prep).  The AzTEC map represents a significant improvement in the uniformity of coverage at faint flux levels.} \\label{fig_goodsn_cover} \\end{figure}  In this paper we present a new 1.1 mm survey of the GOODS-N field made with AzTEC \\citep{Wilson08} at the 15-m James Clerk Maxwell Telescope (JCMT) on Mauna Kea, Hawaii.  This map is the deepest blank-field survey undertaken during the AzTEC/JCMT observing campaign, and is one of the largest, deepest, and most uniform mm-wavelength maps of any region of the sky.  Our map covers $245$\\,arcmin$^2$ and completely encompasses the $16.5^{\\prime}\\times10^{\\prime}$ {\\em Spitzer} GOODS-N field and all of the previous GOODS-N sub-mm and mm-wavelength fields, including the original HDF map of \\citet{Hughes1998} and the SCUBA GOODS-N `Super-map' \\citep[indicated in Fig.~\\ref{fig_goodsn_cover} here and presented in][]{Borys2003,Pope2005}.  The large number and high stability of the AzTEC bolometers has enabled us to produce a map with small variations in r.m.s., from 0.96--1.16\\,mJy, across the 245 min$^2$ field.  This uniformity is a drastic improvement over the SCUBA GOODS-N `Super-map.' The sensitivity variations of the AzTEC and SCUBA maps are compared in Fig.~\\ref{fig_goodsn_cover}.  In this work, we extract a catalogue of mm sources from the map and calculate number counts towards the faint end of the 1.1-mm galaxy population.  The main results we discuss here were obtained from the AzTEC data alone; data from other surveys have been used only as tools to check the quality of our map.  A second paper will address counterpart identification of our AzTEC sources at other wavelengths (Chapin et al. in prep.).  We present the JCMT/AzTEC observations of GOODS-N in \\S~\\ref{obs}, data reduction and analysis leading to source identification in \\S~\\ref{ana}, properties of our source catalogue in \\S~\\ref{sources}, the number counts analysis in \\S~\\ref{nc}, the discussion of results in \\S~\\ref{results}, and the conclusion in \\S~\\ref{conclusion}. ",
        "conclusions": "\\label{conclusion}  We have used the AzTEC instrument on the JCMT to image the GOODS-N field at 1.1~mm.  The map has nearly uniform noise of 0.96--1.16$\\,{\\rm mJy}\\,{\\rm beam}^{-1}$ across a field of $245\\,{\\rm arcmin}^2$.  A stacking analysis of the map flux at known radio source locations shows that any systematic pointing error for the map is smaller than 1\\,arcsec in both RA and Dec.  Thus, the dominant astrometric errors for the 36 source-candidates with S/N$\\geq$3.5 are due to noise in the centroid determination for each source.  Using a S/N$\\geq$3.75 threshold for source robustness, we identify a subset of 28 source candidates among which we only expect 1--2 noise-induced spurious detections.  Furthermore, of the 11 AzTEC sources that fall within the considered region of the SCUBA/850-\\um, 8 are detected unambiguously.  This AzTEC map of GOODS-N represents one of the largest, deepest mm-wavelength surveys taken to date and provides new constraints on the number counts at the faint end (down to ${\\sim}\\,2\\,{\\rm mJy}$) of the 1.1~mm galaxy population. We compare two very different techniques to estimate the number density of sources as a function of their intrinsic flux--a frequentist technique based on the flux histogram of detected sources in the map similar in spirit to that of \\citet{Laurent2005}, and a Bayesian approach similar to that of \\citet{Coppin2006}.  Reassuringly, the two techniques give similar estimates for the number counts.  Those results are in good agreement with the number counts estimates of \\citet{Greve2004} but differ significantly from those of \\citet{Maloney2005}.  The 1.1~mm number counts from this field are consistent with a direct flux scaling of the 850~\\micron\\ SCUBA/SHADES number counts \\citep{Coppin2006} within the uncertainty of the two measurements, with a flux density scaling factor of $2.08\\pm0.18$.  If we assume that the two instruments are detecting the same population of sources, we obtain a grey body emissivity index of $2.84\\pm0.32$ for the dust in the sources.  While there is no evidence based on the number counts that 1.1~mm surveys select a significantly different population than 850~\\micron\\ surveys, we caution that the number counts alone cannot really test this hypothesis.  A more thorough study of whether AzTEC is selecting a systematically different population than SCUBA can come only from comparison of the redshifts and multi-wavelength SEDs of the identified galaxies, which we will describe in Chapin et al. (in prep.), the second paper in this series.  There is also a survey of GOODS-N with MAMBO at 1.25\\,mm performed by Greve et al. (in prep.).  A comparison between these two millimetre maps and, possibly, the SCUBA `Super-map' is reserved for a future paper (Pope et al. in prep.).  This AzTEC/GOODS-N map is one of the large blank-field SMG surveys at 1.1~mm taken at the JCMT. Combined with the AzTEC surveys in the COSMOS \\citep{Scott08} and SHADES (Austermann et al. in prep.) fields, these GOODS-N data will allow a study of clustering and cosmic variance on larger spatial scales than any existing (sub)mm extragalactic surveys."
    },
    "0806/0806.3272_arXiv.txt": {
        "abstract": "Recent improvements in the capabilities of low frequency radio telescopes provide a unique opportunity to study thermal and non-thermal properties of the cosmic web. We argue that the diffuse, polarized emission from giant radio relics traces structure formation shock waves and illuminates the large-scale magnetic field. To show this, we model the population of shock-accelerated relativistic electrons in high-resolution cosmological simulations of galaxy clusters and calculate the resulting radio synchrotron emission. We find that individual shock waves correspond to localized peaks in the radio surface brightness map which enables us to measure Mach numbers for these shocks. We show that the luminosities and number counts of the relics strongly depend on the magnetic field properties, the cluster mass and dynamical state. By suitably combining different cluster data, including Faraday rotation measures, we are able to constrain some macroscopic parameters of the plasma at the structure formation shocks, such as models of turbulence.  We also predict upper limits for the properties of the warm-hot intergalactic medium, such as its temperature and density. We predict that the current generation of radio telescopes (LOFAR, GMRT, MWA, LWA) have the potential to discover a substantially larger sample of radio relics, with multiple relics expected for each violently merging cluster.  Future experiments (SKA) should enable us to further probe the macroscopic parameters of plasma physics in clusters. ",
        "introduction": "The plasma within and between galaxies is magnetized.  Despite many observational efforts to measure galactic and intergalactic magnetic fields, their properties and origins are not currently well understood. The magnetic fields influence the physics of the plasma in several important ways.  They couple the collisionless charged particles to a single but complex fluid through the Lorentz force, and trace dynamical processes in the Universe. Magnetic pressure and tension mediate forces and provide the plasma with additional macroscopic degrees of freedom in terms of Alfv{\\'e}nic and magnetosonic waves. They cause the turbulent cascade to become anisotropic towards smaller scales and suppress transport processes such as heat conduction and cosmic ray diffusion across the mean magnetic field. They are essential for accelerating cosmic rays by providing macroscopic scattering agents which enables diffusive shock acceleration (first order Fermi process) and through magneto-hydrodynamic turbulent interactions with cosmic rays leading to second order Fermi acceleration. They illuminate distant cosmic ray electron populations by enabling synchrotron emission and tell us indirectly about violent high-energy astrophysical processes such as formation shock waves or $\\gamma$-ray bursts. The magnetic fields in spiral galaxies are highly regular, showing alignment with the spiral arms. They are believed to arise from weak seed fields amplified by dynamo processes, driven by differential rotation in galactic disks. The seed fields could have been produced by many sources, ranging from stellar winds and jets of active galactic nuclei, to plasma instabilities and battery effects in shock waves, in ionization fronts, and in neutral gas-plasma interactions. More hypothetical ideas for the seed field origins invoke primordial generation in early Universe processes, such as phase transitions during the epoch of inflation.  In order to understand more about magneto-genesis, we need to study the least processed plasma possible that still shows some primordial memory. This points us to the magnetized plasma in intergalactic space, in particular to the plasma in galaxy clusters. There, magnetic fields show a smaller degree of ordering compared to spiral galaxies.  However, their primordial properties may be masked in clusters because of processing by turbulent gas flows, driven by galaxy cluster mergers, and the orbits of the member galaxies.  For an overview on the present observational and theoretical knowledge the reader is pointed to the review articles by \\citet{1987QJRAS..28..197R,1993A&ARv...4..449W, 1994RPPh...57..325K, 1996ARA&A..34..155B, 1999ARA&A..37...37K, 2001SSRv...99..243B, 2001PhR...348..163G, 2002ARA&A..40..319C, 2002RvMP...74..775W}. This work aims at closing a gap between theoretically motivated phenomenological models of large scale magnetic fields and actual observational non-thermal phenomena associated with them.  Diffuse radio synchrotron emission has already been observed in more than 50 galaxy clusters \\citep{2008SSRv..134...93F}.  The emission is associated with the entire intra-cluster medium (ICM).  The synchrotron emission process demonstrates the presence of highly relativistic electrons (cosmic ray electrons, CRe) with a Lorentz factor typically up to $\\gamma \\sim 10^4$ and magnetic fields within the ICM.  The diffuse radio emission can be classified into two categories: radio halos and radio relics. Giant radio halos are centrally located, trace the thermal emission and show no sign of polarization, while radio relics are located at the periphery of clusters, are polarized and are elongated in appearance. There exist a number of classes of radio objects that have been referred to over the years as ``radio relics'' \\citep[][ and references therein]{2004rcfg.proc..335K}. Two of these are associated with extinct or dying active galactic nuclei (AGN). These either host a synchrotron cooling radio plasma from a past AGN outburst that created the radio lobes or are revived ``radio ghosts'' where an aged radio relic has been re-energized by a merger or an accretion shock \\citep{2001A&A...366...26E}.  The focus of this paper is on a third type of radio relic emission, sometimes referred to as radio ``gischt''\\footnote{The name ``gischt'' derives from a German word for the crest on top of waves that are breaking at the shore thus resembling the radio emission of freshly injected electrons by formation shocks.} \\citep{2004rcfg.proc..335K}, that shows diffuse emission on scales up to 1~Mpc.  Diffusive shock acceleration at structure formation shocks can energize a primary population of relativistic electrons that emit synchrotron radiation \\citep{1998A&A...332..395E,2001ApJ...562..233M} in a magnetic field that can be amplified by the post-shock turbulence.  Prominent examples for this class of radio relics can be seen in Abell 3667 \\citep{1997MNRAS.290..577R}, Abell 2256 \\citep{1976A&A....52..107B,1978MNRAS.185..607M, 1979A&A....80..201B, 1994ApJ...436..654R, 2006AN....327..553C}, Abell 3376 \\citep{2006Sci...314..791B}, and, more recently, Abell 2255 \\citep{2008A&A...481L..91P} and Abell 521 \\citep{2008A&A...486..347G}. All galaxy clusters with observed radio relic emission are merging or show signs of ongoing dynamical activity, but not all dynamically active galaxy clusters are observed to have relics. This raises the question of whether diffuse radio emission is a property of a special subset of clusters, or a universal property, with many relics too faint to be seen by current telescopes.  From CMB measurements we know that the Universe is composed of 4.6\\% baryonic matter, {\\it e.g.} \\citet{2008arXiv0803.0547K}. However, when observing the local Universe ($z<1$), we can account for fewer than half of these baryons \\citep{2004IAUS..220..227F,2005ApJ...624..555D}. This is known as the missing baryon problem. The current cosmological paradigm of large scale structure formation provides a solution to the missing baryon problem. As the Universe evolves, large scale structure grows from small density perturbations imprinted during an earlier epoch. In the hierarchical scenario of structure formation, structure grows from small to large scales, with baryons flowing in a filamentary web with clusters at the interstices. The temperature of baryons deviates from adiabatic cooling associated with the Hubble expansion by increasing multiple times in discrete steps -- always corresponding to a passage through a structure formation shock. Before they are shock-heated to the virial temperatures $kT \\sim 1 - 10$~keV of galaxy groups and clusters, where they can be observed through their thermal bremsstrahlung emission, they are predicted to reside in the warm-hot intergalactic medium (WHIM).  Temperatures in the WHIM are in the range of $10^5\\,\\mbox{K}<T<10^7\\,\\mbox{K}$ \\citep{1998ApJ...509...56H, 1999ApJ...514....1C, 2001ApJ...552..473D, 2004ApJ...611..642F, 2005ApJ...620...21K}.  We will investigate whether it is possible for diffuse radio emission associated with these formation shocks to be used as a tracer of the WHIM boundaries and so indirectly observe the WHIM.  \\begin{figure*} \\begin{center} \\begin{minipage}[t]{0.55\\textwidth} \\centering{\\it \\large Mach numbers} \\end{minipage} \\hfill \\begin{minipage}[t]{0.43\\textwidth} \\centering{\\it \\large Radio Gischt at shocks} \\end{minipage} \\resizebox{0.53\\hsize}{!}{\\includegraphics{figures/f1a.eps}}% \\resizebox{0.46\\hsize}{!}{\\includegraphics{figures/f1b.eps}}\\\\ \\end{center} \\caption{Structure formation shocks triggered by a recent merger of a large galaxy cluster ($M\\simeq 10^{15} h^{-1}\\,\\rmn{M}_\\odot$) dissipate the associated gravitational energy. Left: the Mach number of shocks weighted by the energy dissipation rate is shown by the colour (while the brightness displays the logarithm of the dissipation rate). Right: three-color image of energy dissipation rate at shocks (shown with a color scale ranging from black over red to yellow) and radio synchrotron emission at 150 MHz from shock-accelerated relativistic electrons (blue and contours with levels starting at $7\\times 10^{-4}$ mJy arcmin$^{-2}$ and increasing with a factor of 15, respectively). This ``radio gischt'' emission traces structure formation shock waves, highlights the intermittent nature of mass accretion in galaxy clusters, and illuminates magnetic fields that are amplified by turbulence that can be excited by these shock waves. } \\label{fig:machnumbers} \\end{figure*}  A larger sample size of these diffuse radio sources is required to delve deep into the details of the non-thermal processes working within galaxy clusters. However, the combination of low surface brightness, diffuseness and small dynamical range in the sensitivity of current radio telescopes makes the detection of this particular radio emission difficult. With the current capabilities of the Giant Meter Radio Telescope (GMRT, \\citealt{1995JApAS..16..427A}) and the imminent arrivals of the Low Frequency ARray (LOFAR, \\citealt{2003NewAR..47..405R}), the Murchison Wide-field Array (MWA, \\citealt{2006NewAR..50..173M}) and the Long Wavelength Array (LWA, \\citealt{2005ASPC..345..392K}), and eventually the construction of the Square Kilometre Array (SKA, \\citealt{2004NewAR..48.1119K}), powerful low frequency radio telescopes are positioned to further increase our understanding of diffuse radio emission and give us insight into the following important topics: \\begin{itemize} \\item  the strength and coherence scale of magnetic fields on scales of galaxy clusters, \\item the process of diffusive shock acceleration of electrons, \\item the existence and properties of the WHIM, \\item the exploration of observables beyond the thermal cluster emission which are sensitive to the dynamical state of the cluster. \\end{itemize}  \\indent In the course of this work we will consider how radio relic emission can shed light on each of these topics.  To do this, we adopt a simplified model for the shock-accelerated population of electrons. The key figure illustrating these considerations is shown in Fig.~\\ref{fig:machnumbers}.  In our simulations, we can visualize properties of structure formation shocks that are triggered by a recent merger of a large galaxy cluster ($M\\simeq 10^{15} h^{-1}\\,\\rmn{M}_\\odot$) and dissipate the associated gravitational energy. In the left panel, the shock Mach numbers, weighted by the energy dissipation rate, are encoded by the colour, while the brightness displays the logarithm of the dissipation rate. This shows that most of the energy is dissipated in weak flow shocks internal to the cluster, while the shock waves become strongest and steepen as they break at the shallower peripheral potentials of the clusters and within filaments. The right panel shows a three-color image of of energy dissipation rate at shocks (shown in red and yellow) and radio synchrotron emission at 150 MHz from shock-accelerated relativistic electrons (blue and contours, modeled according to \\citealt{2008MNRAS.tmp..277P}).  This ``radio gischt'' emission traces structure formation shock waves and highlights the intermittent nature of mass accretion, in particular along the filament extending from the cluster center to the upper left of the image and for the giant radio relic to the lower-right of the cluster.  This radio emission illuminates magnetic fields that are amplified by magneto-hydrodynamic instabilities that are associated with these shock waves.  The paper has been structured accordingly: in Sect.~\\ref{sec:methodology} we outline our methodology; we describe our results and discuss them in Sects.~\\ref{sec:results} and \\ref{sec:discussion}; and present our conclusions in Sect.~\\ref{sec:conclusions}. ",
        "conclusions": "\\label{sec:conclusions}  The intermittency and inhomogeneous nature of structure formation shocks are characterized by a highly non-Gaussian distribution function.  This requires numerical simulations to study the implied non-universality of the induced radio relic (or gischt) emission. It is hard to conceive of an adequate analytical approach to this problem. Observing the polarized emission of a sample of relics at different frequencies enables us to gain insight into the non-equilibrium processes at work -- in particular into the interplay of large scale magnetic fields and structure formation shocks.  The relevant observables of the relics include morphology, spectral shape, relic luminosity function and Faraday rotation measure. The theoretical implications of radio relic observables are as follows:  We model the shock acceleration of electrons at formation shocks and find that the {\\em morphology} of radio relics unambiguously characterises the underlying structure of dissipating shock waves (Fig.~\\ref{fig:machnumbers}). The resulting simulated relics are very similar to the observed relics and thus support our hypothesis.  Their positions identify regions that are not in equilibrium and where the electron and ion temperatures are expected to show strong deviations due to the comparatively long Coulomb mean free path that governs their equilibration process.  The {\\em relic luminosity function} is sensitive to the combination of normalization and scaling properties of the magnetic field with thermal energy density as well as the electron acceleration efficiency, the mass and dynamical state of a cluster.  Thus, it can provide hints about the processes that generate these large scale fields and can help to disentangle the dominant transport processes which include effects from magnetic flux freezing and growth by turbulent dynamos.   The {\\em rotation measure (RM) map} is sensitive to the line-of-sight integrated magnetic field. From the mean and variance of RM maps, we can infer the location of the formation shock with respect to the cluster center as the variance increases as a function of integration length (Eqn.~\\ref{eq:RMsquared3}). This helps in constraining the geometry of the merger. Deprojecting the 2D RM power spectrum enables one to measure the 3D magnetic power spectrum, under the assumption that the behaviour of the electron density along the line of sight can be obtained from X-ray measurements. The peak of the 3D power spectrum yields the total magnetic energy and the magnetic coherence length $\\lambda_B$. Performing this procedure for different relics or for different regions of one large relic allows us to estimate the variance of $\\lambda_B$ across the cluster and might possibly tell us about the nature of MHD turbulence.  We found that the correlation between $n_\\e$ and $B$ biases the {\\em rms} magnetic field strength derived from RM maps high if this is not taken into account; we note that most works have done so. This correlation should be a natural consequence of MHD effects such as flux freezing. If systematic errors associated with RM studies are smaller than statistical ones, we find that measurements of the small scale slope of the RM power spectrum are accurate enough to differentiate between Kolmogorov ($k^{-5/3}$) and Burgers ($k^{-2}$) turbulence spectra. The interpretation of these slopes is however not straight forward and needs to account for an additional flattening due to small scale fluctuations in $n_{\\mathrm{e}}$.  The {\\em spectral index} of a power law spectrum of a radio relic is a measure of the shock strength of that relic. The median spectral index of a distribution of relics probes a distribution of virializing shocks and can give an indication of CR proton injection. This is of particular relevance for questions concerning the pressure contribution of non-thermal components and enables comparison with predictions of hydrodynamical simulations.  The shape of the spectrum is sensitive to the acceleration mechanism of the relativistic electrons and to their cooling processes. The variation of the spectrum over the relic allows one to infer in situ magnetic field strengths by comparing the synchrotron and IC cooling times to the advection time downstream provided that the magnetic energy density is not much smaller than the CMB energy density. This might constrain models for the magnetic amplification at shocks in high beta plasmas. A radio relic's luminosity is roughly correlated with the shock strength. Thus, it is favourable to look for radio relics in large, dynamically disturbed clusters, or use relic detections as a proxy for dynamical activity of clusters \\citep{2001A&A...378..408S}.  We demonstrated that the combination of {\\em the relic spectral index with deprojected X-ray and SZ profiles} allows one to indirectly infer upper limits on the density and temperature of the warm hot intergalactic medium. Simulations show that the WHIM is not uniform, rather it is characterized by highly inhomogeneous structure that shows intermittent accretion events that are channeled mostly through filaments.  We predict that there will be a large sample of polarized radio relics from a considerable number clusters in the near future. This sample should allow one to constrain macroscopic model parameters, which are expected to be higly non-Gaussian, using a joint analysis method on the radio observables. For example, the combination of the relic luminosity function, RM power spectra, X-ray and SZ measurements should constrain the models of diffusive shock acceleration and large scale magnetic fields. Future work will address the details of this procedure."
    },
    "0806/0806.4387_arXiv.txt": {
        "abstract": "We present multi-wavelength observations of the brightest galaxies in four X-ray luminous groups at $z\\sim0.37$ that will merge to form a cluster comparable in mass to Coma.  Ordered by increasing stellar mass, the four brightest group galaxies (BGGs) present a time sequence where BGG-1, 2, \\& 3 are in merging systems and BGG-4 is a massive remnant ($M_{\\ast}=6.7\\times10^{11}M_{\\odot}$).  BGG-1 \\& 2 have bright, gravitationally bound companions and BGG-3 has two nuclei separated by only 2.5~kpc, thus merging at $z<0.5$ increases the BGG mass by $\\gtrsim40$\\% ($t_{MGR}<2$ Gyr) and $V$-band luminosity by $\\sim0.4$ mag.  The BGGs' rest-frame $(B-V)$ colors correspond to stellar ages of $>3$ Gyr, and their tight scatter in $(B-V)$ color ($\\sigma_{BV}=0.032$) confirms they formed the bulk of their stars at $z>0.9$.  Optical spectroscopy shows no signs of recent ($<1.5$ Gyr) or ongoing star formation.  Only two BGGs are weakly detected at \\mipsmu, and X-ray and optical data indicate the emission in BGG-2 is due to an AGN.  All four BGGs and their companions are early-type (bulge-dominated) galaxies, and they are embedded in diffuse stellar envelopes up to $\\sim140$~kpc across.  The four BGG systems must evolve into the massive, red, early-type galaxies dominating local clusters.  Our results show that: 1) massive galaxies in groups and clusters form via dissipationless merging; and 2) the group environment is critical for this process. ",
        "introduction": "Whether massive cluster galaxies at $z\\sim0$ build up a significant fraction of their stellar masses at $z<1$ via merging continues to be an intensely debated question.  In the paradigm of hierarchical formation, galaxies continue to grow via accretion of smaller satellites \\citep{peebles:70}, and recent models predict that brightest cluster galaxies assemble half their stellar masses at $z<0.5$ \\citep{delucia:07b}.  However, direct merging is not favored in virialized clusters due to their high velocity dispersions ($\\sigma_{1D}\\sim1000$~\\kms).  Brightest cluster galaxies (BCG) are particularly compelling tracers of mass growth because they are the most massive galaxies in the universe, can be observed to high redshifts \\citep{gunn:75}, and should have a rich merger history.  A number of BCGs in the local universe have multiple nuclei \\citep{hoessel:85,lauer:86,laine:03} which suggests recent merging.  However, studies of BCGs in the more distant universe disagree as to whether BCGs have grown significantly in stellar mass since redshift $z\\sim1$ \\citep{brough:02,whiley:08}.  While merging of red, luminous members is observed in groups and dynamically young clusters \\citep{vandokkum:99,tran:05b,mcintosh:07}, the merging frequency and thus its importance remains uncertain.  The key may be environment: merging is rare in established clusters but more frequent in groups that are accreted onto clusters. Because groups have a higher fraction of passive early-type galaxies relative to the field \\citep{wilman:05,jeltema:07}, merging between such systems will form massive, passive galaxies.  Our study focuses on the brightest galaxies in four X-ray luminous galaxy groups at $z\\sim0.37$ discovered in the {\\it Las Campanas Distant Cluster Survey}\\footnote{The LCDCS was designed to find galaxy clusters at $z\\sim0.35-0.9$ and was sensitive to groups only at the lowest redshifts.} \\citep{gonzalez:01}.  These four galaxy groups are gravitationally bound to each other and form an extended structure that will assemble into a galaxy cluster comparable in mass to Coma by the present epoch \\citep[hereafter G05]{gonzalez:05}.  By studying galaxies in the group environment prior to cluster assembly, we directly test whether galaxy-galaxy merging on group scales can drive stellar growth in the massive galaxies observed in local clusters, and whether a large fraction of new stars is formed in the process. Throughout the paper, we use $H_0=70$~km~s$^{-1}$~Mpc$^{-1}$, $\\Omega_M=0.3$, and $\\Omega_{\\Lambda}=0.7$; at $z=0.37$, this corresponds to a scale of 5.12 kpc per arcsec and a lookback time of 4 Gyr. ",
        "conclusions": "These four BGG snapshots provide unprecedented observational evidence for the late ($z<0.5$) stellar assembly of massive spheroidal galaxies in clusters via dissipationless merging of massive progenitors.  Two of the four BGGs are part of merging pairs and another has a double nucleus with a separation of only 2.5~kpc.  When ordered by increasing stellar mass ($2.9-6.7\\times10^{11}M_{\\odot}$), the BGGs present a time sequence where galaxy-galaxy merging increases the BGG mass by $\\gtrsim40$\\% ($t_{MGR}\\lesssim2$ Gyr) and $V$-band luminosity by $\\sim0.4$ mag.  Our results are remarkably consistent with recent claims from semi-analytic models that brightest cluster galaxies tend to assemble at $z<0.5$ \\citep{delucia:07b}; such a model also naturally explains the observed number of brightest cluster galaxies in the local universe with multiple nuclei \\citep[$\\sim20-40$\\%;][]{hoessel:85,laine:03}.  Despite the merger activity, the BGGs are dominated by old ($>10^9$ Gyr) stellar populations and show little sign of ongoing or recent star formation.  The BGGs have rest-frame $(B-V)$ colors corresponding to mean stellar ages of $>3$ Gyr, and their tight scatter in color ($\\sigma_{BV}=0.032$) confirms that the BGGs formed their stars at $z>0.9$.  If the BGGs continue to age passively, their color scatter by $z\\sim0$ will be even smaller than that of the $\\sim100$ brightest early-type galaxies in Coma ($\\sigma_{BV}=0.020$ vs. 0.028).  Their composite optical spectrum shows no sign of star formation in the last $\\sim1.5$ Gyr.  While BGG-2 \\& 3 are weakly detected at \\mipsmu, the X-ray and optical data indicate that BGG-2's \\mipsmu~emission is due to AGN rather than star formation, and possibly for BGG-3 as well.  SG1120 thus provides unique and powerful confirmation of hierarchical formation on both galaxy and cluster scales because 1) SG1120 has massive galaxies assembling from galaxy-galaxy merging, and 2) it is a proto-cluster assembling from the merging of galaxy groups.  Our findings show that the most massive galaxies in groups and clusters form via dissipationless merging, and that the galaxy group environment is critical to forming the massive galaxies dominating local clusters."
    },
    "0806/0806.0025.txt": {
        "abstract": "We report the detection of an extrasolar planet of mass ratio $q \\sim 2 \\times 10^{-4}$ in microlensing event MOA-2007-BLG-192. The best fit microlensing model shows both the microlensing parallax and finite source effects, and these can be combined to obtain the lens masses of $M = 0.060 {+0.028\\atop -0.021}\\msun$ for the primary and $m = 3.3{+4.9\\atop -1.6} \\mearth$ for the planet. However, the observational coverage of the planetary deviation is sparse and incomplete, and the radius of the source was estimated without the benefit of a source star color measurement. As a result, the 2-$\\sigma$ limits on the mass ratio and finite source measurements are weak. Nevertheless, the microlensing parallax signal clearly favors a sub-stellar mass planetary host, and the measurement of finite source effects in the light curve supports this conclusion. Adaptive optics images taken with the Very Large Telescope (VLT) NACO instrument are consistent with a lens star that is either a brown dwarf or a star at the bottom of the main sequence. Follow-up VLT and/or Hubble Space Telescope (HST) observations will either confirm that the primary is a brown dwarf or detect the low-mass lens star and enable a precise determination of its mass. In either case, the lens star, MOA-2007-BLG-192L, is the lowest mass primary known to have a companion with a planetary mass ratio, and the planet, MOA-2007-BLG-192Lb, is probably the lowest mass exoplanet found to date, aside from the lowest mass pulsar planet. ",
        "introduction": "When the first extrasolar planets were discovered orbiting main sequence stars more than a decade ago \\citep{51peg,70vir,47uma}, the radial velocity surveys responsible for the discoveries focused their observations on stars of spectral type F, G, and K, because such stars offered the greatest planet detection sensitivity. However, observations of star forming regions indicate that stars of virtually all types have evidence of proto-planetary disks, suggesting that the first stages of the planet formation are practically independent of star type. The radial velocity surveys have since expanded their target lists to include stars ranging from a spectral type of mid-M to stars that are thought to have spectral type A when they were on the main sequence. Gas giant planets have been found in orbit around stars of all these types \\citep{butler-catalog,sato_plan_giant,johnson_Mdwarf_plan,psu_Gstar_pl,lovis_pl_evolved}, but they appear to be increasingly rare in orbit around low-mass stars.  In the past few years, microlensing surveys \\citep{exoplanet_book} have extended the range of sensitivity to cool,  ``super-Earth\" planets \\citep{ogle390,ogle169} with masses of $\\sim 10\\mearth$ in orbits beyond the ``snow-line\" \\citep{ida_lin,laughlin,kennedy-searth}, where the core accretion theory predicts that the most massive planets should form. These discoveries indicate that such low-mass planets are significantly more common than gas giants in orbit around  the stars of $\\simlt 1\\msun$ that are probed by the microlensing method. Because microlensing does not rely upon light from the planetary host star in order to detect the planet, its sensitivity extends to host star masses well below the $\\sim 0.25\\msun$ lower limit for current radial velocity surveys.  In this paper, we present the analysis of microlensing event MOA-2007-BLG-192, and show that the lens system is likely to be a low-mass planet orbiting a primary that is either a brown dwarf or a very low-mass main sequence star. The data are discussed in \\S~\\ref{sec-data}, and the planetary nature of the light curve is discussed in \\S~\\ref{sec-lc}. The uncertainties in the microlensing model parameters are discussed in \\S~\\ref{sec-param_un}, and in \\S~\\ref{sec-lens_prop}, we show that the microlensing parallax and finite source features of the light curve favor a sub-stellar mass primary and a very low-mass planet. Adaptive optics images from the VLT NACO instrument are only consistent with a primary lens mass that is a brown dwarf or a star at the bottom of the main sequence. In \\S~\\ref{sec-conclude}, we show how future observations with VLT/NACO can confirm this interpretation and determine the parameters of the planetary system more precisely. ",
        "conclusions": "\\label{sec-conclude}  MOA-2007-BLG-192 is the first planetary microlensing event to be discovered without follow-up observations of the light curve, and this discovery was made possible by the very wide field (2.2 sq. deg.) of the MOA-II telescope, which allows the entire Galactic bulge to be imaged hourly. This hourly coverage is sufficient to establish that the light curve deviation can only come from a planet and not a stellar binary, it does leave the planetary parameters less well constrained than would be the case for a high magnification event that had been discovered and announced significantly before peak magnification. As a result, there is significant uncertainty in some of the event parameters, including the source radius crossing time, $t_\\ast$, and the planetary mass ratio, $q$. Some of this uncertainty is due to two different types of light curve models that can explain the data: a cusp approach model and a caustic crossing model.  However, this event also has a significant microlensing parallax signal, which indicates that the planetary host star is likely to have a very low mass, with sub-stellar masses favored. A sub-stellar planetary host mass is favored even more when the marginal measurement of $t_\\ast$ is included in the analysis. However, a preliminary analysis of AO imaging from the VLT/NACO instrument indicates that there may be some excess stellar $J$-band flux at the location of the source star, although the significance of this excess is less than 2-$\\sigma$.  If this $J$-band excess is due to flux from the lens star, then it can be explained by a cusp approach model that has a $\\chi^2$ value larger than the best fit by $\\Delta\\chi^2 = 1.2$. The implied planet host star mass is $M \\approx 0.09\\msun$, and the implied planet mass is quite low, $m \\approx 1.4\\mearth$ because the best low $t_\\ast$ solutions that correspond to stellar mass host stars also have low planet mass ratios.  If the host star is a low-mass star instead of a brown dwarf, then we will have a further opportunity to improve our characterization of the lens system with follow-up HST observations \\citep{ml_plan_char}. These stellar mass lens models require a lens-source relative proper motion of $\\mu_{\\rm rel} \\simgt 5\\,$mas/yr, and this is enough to allow the detection of the separation in HST images taken in a few years. Although the lens and source stars will not be resolved, their separation will be large enough so that the blended lens+source image will be elongated, and the very stable PSF of HST will allow this elongation to be measured precisely when the separation is $\\simgt 15\\,$mas. Such a measurement, will pin down the direction of lens-source motion, which will help to restrict the remaining uncertainty in the microlensing parallax measurement. This would dramatically reduce the parameter uncertainties for this event and leave only one parameter, the mass ratio $q$, that will have a substantial uncertainty (although the uncertainty in $q$ will be slightly reduced with a more precise measurement of $t_\\ast = \\theta_\\ast/\\mu_{\\rm rel} $).  The only possibility that could seriously modify our main conclusion of a very low planetary host star mass would be if the apparent microlensing parallax signal were due to xallarap (or orbital motion of the source star). However, we argue in the Appendix that this is unlikely.  Our analysis indicates that the planetary host star, MOA-2007-BLG-192L, is likely to have a mass in the 0.02--$0.10\\,\\msun$ range (at 95\\% confidence), and it is the first such object known to have a companion with a planetary mass ratio ($q < 0.03$). Two brown dwarfs have previously been reported to have companions that could be of planetary mass \\citep{chauvin,chaha8}, but their mass ratios are $q \\sim 0.2$, which suggests that they did not form like the planetary systems around more massive stars. Thus, our discovery of MOA-2007-BLG-192Lb represents the first discovery of an extrasolar planet with a planetary mass ratio orbiting an extremely low-mass primary that is likely to be a brown dwarf, and with a mass of $m = 3.3{+4.9\\atop -1.6} \\mearth$ (or a 2-$\\sigma$ range of 1.0--$18\\,\\mearth$). Thus, the median estimated mass for MOA-2007-BLG-192Lb represents the lowest mass  for a planetary companion yet to be discovered, aside from the lowest mass of the pulsar planets \\citep{pulsar_planets}. This discovery suggests that planetary systems can form around stars of extremely low-mass \\citep{payne_BD_planets}, and confirms that microlensing is indeed sensitive to Earth-mass planets \\citep{em_planet}. It also indicates that Earth-mass planets should be able to form around very low-mass M-dwarfs, which should provide encouragement for programs that seek to find transiting Earth-mass planets in the habitable zone of M-dwarfs in order to study their atmospheres.  In fact, it is possible that MOA-2007-BLG-192Lb could have a habitable surface temperature itself, despite the fact that its host star or brown dwarf provides extremely feeble radiative heating. \\citet{stevenson99} has speculated that even a free floating Earth-mass planet could have a surface temperature that would allow liquid water even though the heating from internal radioactive decays provides a factor of $\\sim 10^4$ times less energy than the Earth receives from the Sun. The key point of Stevenson's argument was that such a free floating planet might retain a molecular Hydrogen atmosphere that could provide very strong insulation that would allow the surface temperature to remain above the melting point of water ice. If it was possible to detect nearby analogs to MOA-2007-BLG-192Lb, it would be worthwhile to attempt to study their spectra to see if they do have $H_2$ atmospheres that might allow warm surface temperatures.  Figure~\\ref{fig-mass_a} compares this new discovery (indicated by the red circle surrounding a white spot) to previous discoveries by microlensing and other methods. We should note that this discovery appears outside the predicted sensitivity range for ``ground-based microlensing\\rlap.\" This is largely because the ground-based calculations were done for planets with fixed mass ratios and separations, and they were added to this plot using the assumption that the ``typical\" planetary host mass was $0.3\\msun$. On a mass ratio plot, the MOA-2007-BLG-192Lb would appear much higher and closer to the ``ground-based microlensing\" curve.  In this plot, our new planet looks as if it might be a more massive version of Venus, but this is a bit misleading since such a low mass primary provides very much less heating than the Sun. If we want to consider the planet locations in the context of planet formation, then it is best to consider not the planetary semi-major axis, but the planetary semi-major axis divided by the ``snow-line\". In the context of planet formation theory \\citep{ida_lin,lecar_snowline,kennedy-searth,kennedy_snowline}, the ``snow-line\" is the region in the proto-planetary disk where it is just cold enough for water-ice to form. This is expected to increase the density of solids in the disk by a factor of $\\sim 5$, and it is where the most massive planets are expected to form, according to the core accretion theory. Since planets are expected to form early in a star's history, it is not the star's main sequence brightness that determines the location of the snow line. Instead it is the star's brightness at an age of $\\sim 1\\,$million years, when the stellar luminosity is thought to scale as $\\sim M^2$ \\citep{burrows93,burrows97} (G. Kennedy, C. Lada, private communications). Thus, we can estimate the distance of the ``snow-line\" to be $a_{\\rm snow} = 2.7 M/\\msun$, and with this definition we plot the known exoplanets as a function of mass and (semi-major axis)/(snow-line) in Figure~\\ref{fig-mass_snow}. Now with respect to the snow-line, MOA-2007-BLG-192Lb appears to be a much lower mass version of Uranus instead of a massive Venus. This figure also shows that, to date, only microlensing has been able to probe the region beyond the snow line, for planets of less than a Jupiter mass. Our discovery adds strength to the claim that low-mass planets are substantially more common at these separations around stars of less than a Solar mass \\citep{ogle169}."
    },
    "0806/0806.2482_arXiv.txt": {
        "abstract": "We present results from an extensive set of simulations of gravitational fragmentation in the presence of magnetic fields and ambipolar diffusion. The thin-sheet approximation is employed, with an ambient magnetic field that is oriented perpendicular to the plane of the sheet. Nonlinear development of fragmentation instability leads to substantial irregular structure and distributions of fragment spacings, fragment masses, shapes, and velocity patterns in model clouds. We study the effect of dimensionless free parameters that characterize the initial mass-to-flux ratio, neutral-ion coupling, and external pressure associated with the sheet. The average fragmentation spacing in the nonlinear phase of evolution is in excellent agreement with the prediction of linear perturbation theory. Both significantly subcritical {\\it and} highly supercritical clouds have average fragmentation scales $\\lamavg \\approx 2 \\pi Z_0$, where $Z_0$ is the initial half-thickness of the sheet. In contrast, the qualitatively unique transcritical modes can have $\\lamavg$ that is at least several times larger. Conversely, fragmentation dominated by external pressure can yield dense cluster formation with much smaller values of $\\lamavg$. The time scale for nonlinear growth and runaway collapse of the first core is $\\approx 10$ times the calculated growth time $\\taugm$ of the eigenmode with minimum growth time, when starting from a uniform background state with small-amplitude white-noise perturbations. Subcritical and transcritical models typically evolve on a significantly longer time scale than the supercritical models. Infall motions in the nonlinear fully-developed contracting cores are subsonic on the core scale in subcritical and transcritical clouds, but are somewhat supersonic in supercritical clouds. Core mass distributions are sharply peaked with a steep decline to large masses, consistent with the existence of a preferred mass scale for each unique set of dimensionless free parameters. However, a sum total of results for various initial mass-to-flux ratios yields a broad distribution reminiscent of observed core mass distributions. Core shapes are mostly near-circular in the plane of the sheet for subcritical clouds, but become progressively more elongated for clouds with increasing initial mass-to-flux ratio. Field lines above the cloud midplane remain closest to vertical in the ambipolar-drift driven core formation in subcritical clouds, and there is increasing amount of magnetic field curvature for clouds of increasing mass-to-flux ratio. Based on our results, we conclude that fragmentation spacings, magnitude of infall motions, core shapes, and, especially, the curvature of magnetic field morphology, may serve as indirect observational means of determining a cloud's ambient mass-to-flux ratio. ",
        "introduction": "\\label{s:intro} \\subsection{Molecular Cloud Cores}  Star formation occurs in dense cores within interstellar molecular clouds. The existence and properties of dense cores are well established by studies of molecular spectral line emission \\citep{Myer83b,Bens,Jiji}, submillimeter dust emission \\citep{Ward94,Andr,Kirk}, and infrared absorption \\citep{Bacm,Teix,Lada07}. The core formation process has been the subject of intense theoretical study for the past few decades. Ideas range from uninhibited gravitational fragmentation instability \\citep{Jean,Lars85,Lars03}, to ambipolar diffusion in a magnetically supported cloud \\citep{Mest,Mous78,Shu87}, to a very rapid fragmentation due to pre-existing turbulent flows \\citep{Pado,Kless,Gamm}. The last two ideas are related to scenarios for the nature of relatively low-density molecular cloud envelopes, which contain most of the mass of molecular clouds \\citep[see][]{Gold}. In one scenario, they are magnetically-dominated and evolve due to ambipolar diffusion, a gravitationally-driven redistribution of mass amongst magnetic flux tubes \\citep{Mous78}. The relatively low global star formation rate and efficiency in Galactic molecular clouds \\citep[see][]{McKe} is used to argue that molecular cloud envelopes have a subcritical mass-to-flux ratio \\citep[see e.g.][]{Shu99,Elme07}. Indirect empirical evidence based on the Chandrasekhar-Fermi (CF) method and velocity anisotropy measurements does imply that the low density ($n \\sim 100 \\cmc$) regions are subcritical \\citep{Cort,Heye}. The contrasting view is that the mass-to-flux ratio is supercritical on large scales, and that the star formation rate and efficiency are governed by supersonic turbulence \\citep{MacL}. This requires either a continual replenishment of rapidly dissipating turbulence, or a relatively rapid dispersal of the cloud envelopes. The focus of this paper is the fragmentation of embedded and relatively quiescent dense regions, and not on the nature of the larger envelopes. We study the effect of small-amplitude perturbations on a variety of cloud models with different mass-to-flux ratios, degrees of magnetic coupling, and external pressures.   For the purpose of this paper, we refer to any individual unit of star formation, which leads typically to a single or small multiple star system, as simply a ``core''. However, astronomers often subdivide this concept into two observational categories, that of ``prestellar core'' and ``prestellar condensation''. The former term is often used to describe the extended (mean density $n \\sim 10^5 \\cmc$, size $s \\sim 0.1$ pc, spacing $\\lambda \\sim 0.25$ pc) objects in regions of distributed star formation like the Taurus molecular cloud, while the latter term typically describes the more compact ($n \\ga 10^6-10^7 \\cmc$, $s \\sim 0.02-0.03$ pc, $\\lambda \\sim 0.03$ pc) objects that are located within cluster-forming cores in, for example, Ophiuchus, Serpens, Perseus, and Orion \\citep[see discussion in][]{Ward07}. Both the prestellar cores in e.g., Taurus, and the cluster-forming cores (which harbor multiple condensations) in the other clouds occupy only a very small fraction of the total volume and mass of their larger, more turbulent, molecular cloud complex \\citep[e.g.,][]{John,Gold}. A compendium of current data from spectral line emission, dust emission, and infrared absorption tends to show that cores exhibit central density concentration \\citep{Ward94,Andr} subsonic inward motions \\citep{Tafa,Will,Lee,Case} that often extend beyond the nominal core boundary, near-uniform gas temperatures \\citep{Bens}, subsonic internal turbulence \\citep{Myer83a,Full,Good98}, and near-critical magnetic field strengths, when detected by the Zeeman effect \\citep{Crut99,Bour} or inferred by the CF method \\citep{Crut04}.  The measured subsonic infall motions constitute indirect evidence for a force that mediates gravity. The mediative force may be due to the magnetic field, whose strength in dense regions is very close to the critical value for collapse \\citep{Crut99}. If the fragmenting region has a subcritical mass-to-flux ratio, then the formation of cores will be regulated by ambipolar diffusion, i.e. neutral molecules diffusing past ions that are tied to magnetic fields, and will occur on a time scale that may significantly exceed the local dynamical time for typical ionization fractions. Conversely, if the dense gas is supercritical, then the fragmentation takes place on a dynamical time scale. \\citet[hereafter BC04]{BC04} studied the nonlinear development of fragmentation instability in clouds that were initially either critical or decidedly supercritical. They found that supercritical fragmentation is characterized by somewhat supersonic motions on the core scale ($\\sim 0.1$ pc) while critical fragmentation is characterized by subsonic inward motions on those scales. The nonlinear fragmentation of decidedly subcritical clouds is also presented in this paper as part of a comprehensive parameter study of the effects of mass-to-flux ratio, initial ionization fraction, and external pressure. We believe that the three broad categories of subcritical, transcritical, and supercritical fragmentation should all occur in the interstellar medium, and even within separate regions of a single molecular cloud complex. Our results can help to distinguish which mode of fragmentation is occurring in a given observed region.  The linear theory of fragmentation of a partially ionized, magnetic thin sheet has been presented by \\citet[hereafter CB06]{CB06}. An important quantity is the dimensionless critical mass-to-flux ratio, $\\mu = 2 \\pi G^{1/2} \\sigma/B$, where $\\sigma$ is the column density of a sheet and $B$ is the strength of the magnetic field that is oriented perpendicular to the plane of the sheet. In the limit of flux-freezing, fragmentation can occur only if $\\mu > 1$, i.e. the sheet is supercritical. Conversely, gravitationally-driven fragmentation instability cannot occur at all if $\\mu < 1$, i.e. the sheet is subcritical. CB06 show that the inclusion of ambipolar drift in such a model means that fragmentation can occur for {\\it all mass-to-flux ratios}, but on widely varying time scales and length scales. A very important result of CB06 is their Fig. 2 (see also Fig. 1 in this paper), which demonstrates that transcritical ($\\mu \\approx 1$) fragmentation has a preferred scale that can be many times larger than $\\lammax$, the wavelength of maximum growth rate in the thermal (i.e. nonmagnetic) limit. The latter is actually the preferred scale for both highly supercritical {\\it and} highly subcritical clouds. This is because gravitational instability in subcritical clouds develops by ambipolar drift of neutrals past near-stationary magnetic field lines, and occurs on the ambipolar diffusion time scale rather than the dynamical time scale \\citep[see also][]{Lang78,Mous78}. The ``resonant'' transcritical modes with preferred scale $\\lamgm \\gg \\lammax$ occur due to a combination of ambipolar drift and field-line dragging; the latter leads to magnetic restoring forces that can stabilize the perturbations unless they are of the required large size. The time scale for ambipolar-diffusion mediated gravitational instability is also presented extensively in CB06. They showed that highly supercritical clouds undergo gravitational instability on the dynamical time $t_{\\rm d} \\approx Z/\\cs$, where $Z$ is the half-thickness (effectively the scale height) of the sheet and $\\cs$ is the isothermal sound speed. On the other hand, highly subcritical clouds undergo instability on the quasistatic ambipolar diffusion time $\\tau_{\\rm AD} \\approx 10\\,Z/\\cs$ for typical ionization fraction. Transcritical clouds undergo a hybrid instability on an intermediate time scale -- see also \\citet{Zwei} for a similar result. Also, we note that the inclusion of nonlinear fluctuations can reduce the ambipolar diffusion time under certain circumstances \\citep{Fatu,Zwei02}; we do not model such effects in this paper.  Our nonlinear simulations represent a significant extension of the parameter space of models from that presented by BC04. It similarly extends the parameter space studied by \\cite{Inde}, who presented nonaxisymmetric evolution of an infinitesimally thin subcritical sheet, including the effects of magnetic tension but ignoring magnetic pressure. Recent fully three-dimensional simulations by \\citet{Kudo07}, including magnetic fields and ambipolar diffusion, have confirmed the basic results of BC04. However, that paper presents three representative models, and an extensive parameter study remains computationally out of reach. Here, in addition to a parameter study, we carry out large numbers of simulations for {\\it each} unique set of parameters. This is done in order to compile statistics on core spacings, mass distributions, and shapes. All simulations end at the time of the runaway central collapse of the first core, hence we are compiling information about the early phases of star formation in a molecular cloud. The subsequent history of the molecular cloud, after it has been stirred up by the initial star formation, remains to be determined. This paper represents the product of a total of over 700 separate simulations for models with 14 distinct sets of dimensionless parameters.  An alternate approach to modeling fragmentation is to input turbulence directly into the dense thin-sheet model \\citep[e.g.,][]{Li04,Naka05}, rather than relegating its presence to an unmodeled envelope. This results in highly supersonic motions within the dense sheet itself, and more rapid formation of cores than in our models. We do not pursue that approach in the present study but leave it open to future investigation.  \\subsection{Relation to Global Cloud Structure} Our simulations represent an intermediate approach between attempting a global model of large-scale molecular cloud structure and of modeling the interior collapse of individual cores. Large-scale models need to account for the overall structure and self-gravity of molecular clouds and cannot be realistically studied using a periodic-box model. Truly global three-dimensional models are computationally expensive and still rarely attempted. Smoothed particle hydrodynamics (SPH) techniques have been used successfully to model entire cluster-forming regions \\citep{Bate,Bonn}, and can explain some important features of star formation like mass segregation, binary fraction, and the initial mass function. These models are not fully global in that they do not include the effect of the molecular cloud envelopes. They also do not include the effect of magnetic fields or feedback from outflows, so they cannot address the observed low global star formation efficiency (SFE), $1-5$\\% \\citep{Lada03}. However, the latest version of such models \\citep{Pric} does include flux-frozen magnetic fields and yields somewhat lower SFE's than the non-magnetic models. Another three-dimensional approach which is fully global and based on finite-difference or SPH techniques is modeling the formation of a molecular cloud from large-scale supersonic gas flows and the early evolution of the cloud \\citep{Vazq06,Vazq07,Heit,Henn}. These models yield an initially flattened cloud whose subsequent evolution is affected by several instabilities (thermal instability, thin-shell instability, and Kelvin-Helmholtz instability). The overall structure evolves away from a sheet-like configuration, but individual segments may be treated as such. Of the above mentioned studies, only the work of \\citet{Henn} includes the effect of magnetic fields, which may work to suppress some of the instabilities. It is not clear whether this is an important factor in that model. However, the smaller-scale simulation of \\citet{Naka08} does show the formation of a sheet-like structure due to the presence of a dynamically important magnetic field. Amongst our models, the cases with high bounding pressure may be most appropriate for the scenario of cloud formation due to colliding flows. Our low external pressure models are more appropriate to the alternate scenario where the flows that form a molecular cloud are driven by gravity and/or channeled by a large scale magnetic field, e.g. by the magneto-Jeans instability \\citep{Kim} or the Parker instability \\citep{Park}.  Despite the various theoretical emphases on flattened structure, we note that observations of molecular clouds reveal complex morphologies with projected shapes that may not look like globally flattened structures. Molecular clouds have been described variously in the literature as stratified objects supported by internally generated turbulence \\citep{McKe99}, or as fractal objects \\citep{Elme96} in which the internal pressure is not as relevant. Our local sheet model may be applicable to the dense subregions of the clouds (where stars actually form in weak or rich clusters) in either scenario. As an example from the first scenario described above, one-dimensional global models \\citep{Kudo03,Kudo06,Foli} of molecular clouds reveal that internally-driven turbulence yields large-amplitude motions in lower-density envelopes, while retaining transonic motions in embedded dense regions; the fragmentation of the latter may be described by our models. From an observational point of view, we may apply our models to dense star forming regions such as L1495 and HCl 2 in the Taurus molecular cloud \\citep[see][]{Gold}, the L1688 cluster-forming core in Ophiuchus \\citep{Mott}, or better yet to the Pipe Nebula \\citep{Muen}, which represents an even earlier stage of evolution, with a large number of relatively quiescent prestellar cores that are found in a dense elongated region (the ``stem'' of the Pipe). This region may represent the best available laboratory for the study of the early stage of star formation, before feedback from star formation has significantly modified a cloud's internal sructure and motions. There is also evidence that the stem of the Pipe Nebula is flattened along the direction of the mean magnetic field \\citep{Alve}. This property is similar to the better established result for the elongated structures in Taurus \\citep{Good90,Gold}. ",
        "conclusions": "\\label{s:disc}  Our simulations of the nonlinear development of gravitational instability under the influence of magnetic fields and ambipolar diffusion start from a background state of uniform column density and magnetic field strength. Small-amplitude white-noise perturbations initiate the evolution and eventually lead to the nonlinear growth of fragments. Averaging over a large number of simulations reveals that the average spacing of nonlinearly developed cores is essentially that predicted from the preferred fragmentation scales in linear perturbation theory (CB06). However, the time to reach fully developed runaway collapse is up to ten times longer than that of the eigenmode with minimum growth time $\\taugm$. The quantity $\\taugm$ itself varies from $\\approx Z_0/\\cs$ (essentially the free-fall time $\\approx 1/\\sqrt{G\\rhoni}$ for unpressured sheets) for highly supercritical models to $\\approx 10 Z_0/\\cs$ for highly subcritical models (for a typical neutral-ion coupling level). The times to reach runaway collapse vary widely amongst models with different mass-to-flux ratios, ionization fractions, and external pressures. For a cloud with $\\Nni=10^{22}\\cms$ and $T=10\\K$, the times to reach runaway growth of the first core ranges from 0.45 Myr to 9.53 Myr (see Table 2). Since our simulations start from a flat density background, these times represent {\\it upper limits} to the time that fragmentation might take for each set of parameters. However, an advantage of the uniform background density is that it allows for a self-consistent modeling of the entire core formation process, without questions about the origin of initially peaked density distributions used in earlier axisymmetric calculations \\citep[e.g.][]{CM93,BM94}.   In a medium with initial nonlinear perturbations, the time scales for all sets of parameters are indeed likely to be shorter. However, we believe that our calculated time scales are relevant if the corresponding dimensional values are obtained from higher starting values of column density brought about in certain regions by pre-existing (including turbulent) flows. For example, the Taurus molecular cloud has an overall background number column density $N \\approx (1-2) \\times 10^{21} \\cms$ but also contains embedded dark clouds with $N \\approx 5 \\times 10^{21} \\cms$, such as HCl 2 and L1495, within which there are small clusters of $\\sim 10-20$ YSO's and also many dense cores \\citep[see][]{Gome,Onis,Gold}. Our periodic model may be applied to such dark clouds that may themselves have been brought together by nonlinear flows originating in the larger cloud, external triggers, or an earlier phase of gravitational fragmentation. A second example application of our periodic model may be within the L1688 dark cloud in Ophiuchus, which has $N \\approx 10^{22}\\,\\cms$ \\citep{Mott}. The mean spacing of fragments in these two regions varies \\citep{Andr00}, with core edges measured to be at radii $\\la 5000\\, \\AU$ in L1688 but at $\\la 20000 \\, \\AU$ in the Taurus dark clouds \\citep{Andr00}. While some of the difference in spacing may be due to the different background column densities, other important aspects of spatial and kinematic structure may also arise due to different values of $\\mui$ and $\\Pexttil$ (and $\\tniitil$ to a lesser extent), as demonstrated in this paper.  Unlike a uniform density three-dimensional medium, our thin sheet actually has a preferred scale for gravitational fragmentation with a unique value for any given set of initial dimensionless parameters. Since our simulation region is always safely larger than this fragmentation scale, it is unlikely that the size of our system (in the $x$- and $y$-directions) influences the final outcome, as measured by fragment spacings, time scales to runaway collapse, and core mass distributions, for example. This is supported by our tests with runs at quadruple the size of the standard simulations. Incidentally, this is not the case for three-dimensional periodic box simulations, in which the fastest growing mode of gravitational instability is always that of the box size. We believe that a stratified medium is also a more physical and realistic starting point than a uniform three-dimensional medium. This is because the formation process of molecular clouds, or magnetic fields and nonlinear flows within it, will tend to set up compressed regions with a characteristic scale similar to the half-thickness $Z_0 \\approx \\csq/(\\pi G \\signi)$ of our adopted background reference sheet geometry. That scale is related to the Jeans scale and effectively determines the preferred fragmentation scale, as modified by magnetic field strength, ionization fraction, and external pressure.     \\bfig \\epsfig{file=f13.eps,width=\\linewidth,clip=} \\caption{Histogram of masses from models with $\\mui=0.5,0.8.1.0,1.1,2.0$ and $\\tniitil=0.2,\\Pexttil=0.1$. A weighted average of core masses is taken, so that an equal simulated area is assigned to each of models $1-5$. The cores are the same as the ones in the first five panels of Fig.~\\ref{masshist}. The bin width is 0.1. } \\label{totalmasshist} \\efig   The core mass distributions that we have compiled from our simulations are narrowly-peaked and do not match the broader distributions commonly observed in cluster-forming regions \\citep[e.g.][]{Mott}. What is the missing physics that can explain the discrepancy? The explanation that is closest to the spirit of our models is that real molecular clouds start their lives with an inhomogeneous distribution of physical quantities, including the mass-to-flux ratio. We expect that the values of $\\mu$ in a molecular cloud may fall within the observationally-established range of $[0.5,2]$, i.e. within a factor of two of the critical value in each direction. In that case, we can make a simple estimate of a core mass distribution by adding up the histograms for the five models with $\\mui=[0.5,0.8.1.0,1.1,2.0]$ shown in Fig.~\\ref{masshist}. Since the model with $\\mui=1.1$ is run fewer times but with a larger box size, we sample the number of cores necessary to give equal area weighting with the other models. The resulting histogram is presented in Fig.~\\ref{totalmasshist}. The broad variation of peaks in the individual histograms seen in Fig.~\\ref{masshist} leads to a smooth power-law tail in the high mass end of the composite histogram. Fig.~\\ref{totalmasshist} reveals a slope $d\\log N/d\\log M \\approx -2$ in this region, only somewhat steeper than the value $d\\log N/d\\log M \\approx -1.5$ measured for example by \\citet{Mott}. We note that this general mechanism, arising from an initially inhomogeneous distribution of mass-to-flux ratio, is an interesting new possibility for explaining the observed broad core mass distributions. An important point is that a relatively {\\it narrow} distribution of initial mass-to-flux ratios (factor of a few variation) can lead to a relatively broad distribution of fragment masses. This mechanism remains to be explored more generally and does not exclude the occurrence of other mechanisms like competitive accretion in a more global model \\citep{Bonn}, a temporal spread of core accretion lifetimes \\citep{Myer00,BJ04}, and turbulent fragmentation \\citep[e.g.][]{Pado, Kless, Gamm,Till}.   The richness of the physics revealed by our models of gravitational fragmentation up to runaway collapse of the first core, under conditions of varying mass-to-flux ratio, ionization level, and external pressure, pave the way for more extensive models in the future. Adding the effects of initial cloud turbulence, implementing a technique for integrating past the formation of the first generation of stars, and including some form of energy feedback from star formation, remain to be done. The addition of new and more complex effects will be facilitated by the fact that the thin-sheet approximation allows efficient calculation of the fragmentation process while retaining a high level of realism. In this context we point out that the recent fully three-dimensional fragmentation simulation of \\citet{Kudo07} with magnetic fields and ambipolar diffusion bears out the main physical results presented by BC04."
    },
    "0806/0806.1953_arXiv.txt": {
        "abstract": "We investigate multi-field inflationary scenarios with fields that drop out of the model in a staggered fashion. This feature is  natural in certain multi-field inflationary setups within string theory; for instance, it can manifest itself when fields are related to tachyons that condense, or inter-brane distances that become meaningless when branes annihilate. Considering  a separable potential, and promoting the number of fields to a smooth  time-dependent function, we derive the formalism to deal with these models at the background and perturbed level, providing general expressions for the scalar spectral index and the running.   We recover known results of e.g. a dynamically relaxing cosmological constant in the appropriate limits. We further show that isocurvature perturbations are suppressed during inflation, so that perturbations are adiabatic and nearly Gaussian. The resulting setup might be interpreted as a novel type of warm inflation, readily implemented within string theory and without many of the shortcomings associated with warm inflation.  To exemplify the applicability of the formalism we consider three concrete models: assisted inflation with exponential potentials as a simple toy model (a graceful exit becomes possible), inflation from multiple tachyons (a constant decay rate of the number of fields and negligible slow roll contributions turns out to be in good agreement with observations) and inflation from multiple M5-branes within M-theory (a narrow stacking of branes yields a consistent scenario). ",
        "introduction": "An inflationary epoch of the early universe is widely accepted as the most efficient mechanism to solve the flatness and the horizon problem, while providing a nearly scale invariant spectrum of scalar perturbations, in agreement with observations. However, the embedding of inflation driven by a single scalar field within string theory (presently, the only known self-consistent theory of quantum gravity) has proven to be challenging, since it requires an extremely flat potential. In addition, the presence of many dynamic fields in string theory renders single field models less appealing because all but one degree of freedom need to be already stabilized at the onset of inflation. Though far from being simple or generic, a partially successful implementation is the KKLMMT construction \\cite{Kachru:2003sx}.  To ameliorate the problem of fine tuning the potential, Liddle et.~al.~proposed assisted inflation in \\cite{Liddle:1998jc}  (see also \\cite{Malik:1998gy,Kanti:1999vt} and follow ups), a certain type of multi-field inflation whereby the presence of many fields increases the Hubble friction, allowing for steeper potentials to still drive slow roll inflation. One shortcoming, however, is the need for fine tuned initial conditions for particular potentials, especially in the absence of an attractor solution (for instance in N-flation \\cite{Dimopoulos:2005ac}; see \\cite{Calcagni:2007sb} for a more general stability analysis of assisted inflation). The presence of many degrees of freedom in string theory has, over the years, galvanized the emergence of a variety of models implementing assisted inflation; these include inflation from multiple tachyons \\cite{Majumdar:2003kd}, from multiple M5-branes within M-theory \\cite{Becker:2005sg} or from axions \\cite{Dimopoulos:2005ac} among others, see e.g. \\cite{Cline:2005ty,Ward:2007gs,Grimm:2007hs,Panigrahi:2007sq,Piao:2002vf,Mazumdar:2001mm} for a small selection.  To make predictions in these models, it is common to introduce a single effective degree of freedom, visualized as the length of the trajectory in field space, see e.g. the review \\cite{Wands:2007bd}. Utilizing this approach correctly recovers adiabatic perturbations, but not isocurvature/entropy modes. The latter ones can be pictured as perturbations perpendicular to the trajectory and develop whenever several degrees of freedom are present \\cite{Gordon:2000hv}. Although non-Gaussianities may be produced (see the review \\cite{Wands:2007bd} for details), it is possible to show that they are generically slow roll suppressed \\cite{Battefeld:2006sz,Battefeld:2007en} (see however \\cite{Byrnes:2008wi}). With this in mind, is it then feasible to discriminate a given multi-field model from its single field analog?  In this article, we investigate a, so far largely ignored (see however \\cite{Ashoorioon:2006wc}), discriminating property of multi-field models, namely the possibility that fields decay during inflation. By \"decay\" we mean that individual fields suddenly become obsolete, while their energy is converted into other forms, such as radiation. This is a widespread feature of models within string theory: consider, for instance, an association of inflatons with distances between branes that are located in some internal space. If a  brane annihilates during inflation, for example as a result of its dissolution into a boundary brane or a collision with an anti-brane, the distance to the just mentioned brane becomes meaningless. Of course, energy does not vanish, but is converted into a different type during the annihilation. A further example is inflation driven by tachyons, which can condense during inflation in a staggered fashion.  The disappearance of fields causes an additional decrease of the potential energy during inflation, which can be even more important than the reduction of  potential energy induced by slow roll. Hence, leading order corrections to observable parameters that are sensitive to the slow roll parameters, such as the scalar spectral index, are possible. Indeed, one can construct models of inflation entirely without a slow roll phase, similar to inflation without inflatons as proposed in \\cite{Watson:2006px}.  In order to compute these effects we assume a simple separable potential $W=\\sum_{A=1}^{\\mathcal{N}}V_A$ and promote the number of fields $\\mathcal{N}$ to a time dependent function. We further smooth out $\\mathcal{N}$ so that we can introduce a continuous decay rate $\\Gamma(t)\\equiv -\\dot{\\mathcal{N}}/\\mathcal{N}$. This approach is only justified if the number of fields and $\\Gamma$ are large enough so that several fields disappear in any given Hubble time. Further, any signal due to a rapid drop in the potential energy, such as a ringing in the power spectrum  or additional non-Gaussianities \\cite{Covi:2006ci,Ashoorioon:2006wc,Chen:2006xjb}, cannot be recovered by this approach. To guarantee energy-momentum conservation ($\\nabla_\\nu T^{\\mu \\nu}=0$), we are forced to introduce an additional component of the total energy momentum tensor that takes over the energy of the disappearing fields. The ratio of its energy density to the one in the remaining inflatons yields an additional small parameter $\\bar{\\varepsilon}$ appearing alongside the usual slow roll parameters. As a consequence, the equation of motion of the effective field is modified. The resulting set-up is reminiscent of warm inflation \\cite{Berera:1995ie}, and might indeed be seen as a new, less problematic implementation of warm inflation within string theory.  At the perturbed level, we consider adiabatic and entropy perturbations and show that the latter ones are suppressed in the models of interest. Focusing on the Mukhanov variable, we derive the scalar spectral index and its running for general decay rates and number of fields. We recover the usual slow roll result as well as \\cite{Watson:2006px} (a dynamically relaxing cosmological constant) in the appropriate limits, but we also find leading order corrections to the slow roll result in general. The running remains second order in small parameters and is, therefore, too small to be observed.  To exemplify the applicability of the formalism we make a detailed study of three concrete models; firstly, as an instructive toy model, we consider assisted inflation with exponential potentials and a nonzero, constant decay rate, which we insert by hand. $\\Gamma\\neq 0$ provides a graceful exit to inflation at the cost of requiring somewhat flatter potentials or more fields so that the model remains consistent with observations. If the decay rate is too big ($\\Gamma\\gtrsim 0.04 H_{inf}$) the scalar spectral index becomes unacceptably red, but all smaller rates work well. Secondly, we investigate inflation from multiple tachyons, as proposed in \\cite{Majumdar:2003kd}. Here, tachyons get displaced from the top of their potential by thermal or quantum mechanical fluctuations, causing them to condense during inflation. We find that the best motivated and least fine tuned setup consists of a constant decay rate and tachyons close to the top of their potential so that their slow roll evolution yields negligible contributions to observable parameters. The resulting spectral index is $n_s-1=3/N$, where $N\\approx 60$ is the number efolds at which we evaluate $n_s$. This is in good agreement with the observations of the CMBR, $n_s=0.960^{+0.014}_{-0.013}$ \\cite{Komatsu:2008hk}. Lastly, we look at inflation from multiple M-5 branes within M-theroy \\cite{Becker:2005sg}. Here, inflatons are associated with the inter-brane distance between M5-branes located within an orbifold $\\bf S_1/Z_2$. Whenever a brane comes close to a boundary brane it dissolves into the boundary via a small instanton transition. Given the model's parameter ranges, which encompass, for instance, the concrete potential and the maximal number of branes, we show that inflation comes to an end within a few efolds after the first outermost brane disintegrates. Hence, only an initial narrow stacking of branes  (involves fine tuning) permits a consistent scenario, and the usual slow roll expressions apply. Based upon these findings, cascade inflation as investigated in \\cite{Ashoorioon:2006wc} is irrelevant for the cosmological scales that are observed in the CMBR, given the parameters put forward in \\cite{Becker:2005sg}.  It is evident that whether or not corrections due to $\\Gamma\\neq 0$ are important is  model-dependent, ranging from being the primary ingredient, as in the tachyon case, to being not pertinent, as in the M5-brane case. Thus, the many implementations of multi-field inflation within string theory should be thoroughly re-investigated.  The outline of this paper is as follows: In section \\ref{sec:main}, we derive the formalism to deal with dacaying fields, at the background level (sec.~\\ref{sec:background}) and the perturbed one (sec.~\\ref{sec:spectralindex}). We compute the spectral index and the running in section \\ref{sec:spectralindex} and \\ref{sec:running}. Along the way, we draw comparisons to warm inflation in section \\ref{sec:warm} and comment on isocurvature modes as well as non-Gaussianities in section \\ref{sec:entropy}. We then shift gears and focus on applications,  that is we investigate concrete models: assisted inflation (sec.~\\ref{sec:staggered}), inflation from tachyons (sec.~\\ref{sec:ACDMM}) and inflation from multiple M5-branes (sec.~\\ref{sec:BBK}); summaries of the conclusions for each model can be found at the end of their respective subsection. We conclude in section \\ref{sec:concl}. ",
        "conclusions": "} In this study, we investigated the consequences of disappearing inflatons during multi-field inflation. Such a process occurs quite naturally in string theoretical models of multi-field inflation, for instance if the inflatons are associated with inter-brane distances and said branes start to annihilate during inflation, or if inflatons are identified with tachyons that condense in a staggered fashion.  Focusing on inflationary models with separable potentials, we promoted the number of fields to a time dependent, continuous variable and derived the general formalism at the background and perturbed level to allow for a non-zero decay rate. This approach is only valid if the number of fields is large and the decay rate is big enough so that at least a few fields drop out of the model in a given Hubble time.  To satisfy conservation of the energy momentum tensor, an additional component to the energy budget needs to be introduced, which takes over the energy of the decaying fields. A non-zero decay rate can cause leading order corrections in observable quantities such as the scalar spectral index or the running, for which we derived general expressions. We recovered standard slow roll results and the case of a dynamically relaxing cosmological constant of  \\cite{Watson:2006px} in the appropriate limits. In addition, we showed that isocurvature modes are generically suppressed during inflation for the models of interest; as a consequence, perturbations are adiabatic and nearly Gaussian. We also point out that staggered multi-field inflation might be seen as a new, less problematic type of warm inflation, which can be implemented within string theory.  Based on our formalism, we investigated three concrete models: assisted inflation with exponential potentials \\cite{Liddle:1998jc}, inflation from multiple tachyons \\cite{Majumdar:2003kd}, and inflation from multiple M5-branes \\cite{Becker:2005sg}. For assisted inflation, the presence of a non-zero decay rate offers a graceful exit to inflation, while remaining observationally viable. For inflation from multiple tachyons, the use of our formalism offers a reliable method to compute the scalar spectral index; we concluded that the setup with best computational control, minimal fine-tuning and good agreement with observations consists of a constant decay rate slightly below one, a few hundred tachyons and a negligible slow roll contribution.  For inflation from multiple M5-branes, we showed that the dissolution of M5-branes into boundary branes is quite rapid within the allowed parameter ranges. Therefore, inflation comes to an end within a few efolds after the first M5-branes disappear. Hence, the only self consistent setup employs a (fine-tuned) narrow stack of branes, whose number is constant around the pivot scale for which the scalar spectral index is measured, and the standard slow roll result is applicable.  We did not investigate further the connection to warm inflation, since it would go beyond the scope of this article, but it is surely an avenue worth investigating more thoroughly. Further, our selection of models is not extensive, and had the primary purpose of exemplifying the applicability of the formalism in a few concrete, distinct scenarios. Variations of the models discussed here (and also entirely different ones) are abundant in the literature and should be re-examined."
    },
    "0806/0806.3167_arXiv.txt": {
        "abstract": "{The \\textsc{Planck} satellite will observe the full sky at nine frequencies from 30 to 857~GHz. Temperature and polarization frequency maps made from these observations are prime deliverables of the \\textsc{Planck} mission. The goal of this paper is to examine the effects of four realistic instrument systematics in the 30~GHz frequency maps: non-axially-symmetric beams, sample integration, sorption cooler noise, and pointing errors. We simulated one year long observations of four 30~GHz detectors. The simulated timestreams contained cosmic microwave background (CMB) signal, foreground components (both galactic and extra-galactic), instrument noise (correlated and white), and the four instrument systematic effects. We made maps from the timelines and examined the magnitudes of the systematics effects in the maps and their angular power spectra. We also compared the maps of different mapmaking codes to see how they performed. We used five mapmaking codes (two destripers and three optimal codes). None of our mapmaking codes makes an attempt to deconvolve the beam from its output map. Therefore all our maps had similar smoothing due to beams and sample integration. This is a complicated smoothing, because every map pixel has its own effective beam. Temperature to polarization cross-coupling due to beam mismatch causes a detectable bias in the TE spectrum of the CMB map. The effects of cooler noise and pointing errors did not appear to be major concerns for the 30~GHz channel. The only essential difference found so far between mapmaking codes that affects accuracy (in terms of residual root-mean-square) is baseline length. All optimal codes give essentially indistinguishable results. A destriper gives the same result as the optimal codes when the baseline is set short enough (Madam). For longer baselines destripers (Springtide and Madam) require less computing resources but deliver a noisier map.} ",
        "introduction": "\\label{sec:intro}  Starting in 2003, {\\sc Planck} Working Group 3 (the ``CTP'' group) undertook a comparison of mapmaking codes in increasingly realistic situations. The approach to realism proceeded in four steps, named after the locations of working meetings of the group (Cambridge, Helsinki, Paris, and Trieste). Results from the Cambridge, Helsinki, and Paris steps have been presented in previous papers (Poutanen et al.~\\cite{Pou06}, Ashdown et al.~\\cite{Ash07a}, \\cite{Ash07b}). Here we present results from the Trieste simulations, designed to determine how mapmaking codes handled four aspects of real {\\sc Planck} data not included in previous simulations. The first was non-axially-symmetric beams. In previous simulations, we assumed that the beams on the sky were axially symmetric Gaussians. The second was the effect of detector sample integration, which introduces an effective smearing of the sky signal along the scanning direction. The third was ``cooler noise'', representing the effect of temperature fluctuations induced in the focal plane by the 20~K sorption cooler. The fourth was the pointing errors. In previous simulations, we assumed that the detector pointings were known without error in the mapmaking. In this paper we present the results of this latest round of simulations, which are realistic enough to allow us to draw some preliminary conclusions about mapmaking. We also outline additional work that must be done before final conclusions can be drawn.  The organization of the paper is as follows. In Section 2 we describe the simulations that produced the time-ordered data (TOD) streams that were inputs to our mapmaking. In Section 3 we give the inputs that were used in these simulations. In Section 4 we describe the mapmaking codes we used in this study. Section 4 details the changes that we made in those codes since our earlier Paris simulation round. Section 5 gives the results of our Trieste simulation round and the computational resource requirements of our mapmaking codes are listed in Section 6. Finally we give our conclusions and proposal for future mapmaking tests in Section 7. In Appendix A we describe an analytic model that we used in explaining the effects of beam mismatch in the CMB maps. Appendix A also shows how we can use this model to correct these effects from the observed spectrum. ",
        "conclusions": "In this paper we have presented results from a simulation that we designed to determine how mapmaking codes handle four aspects of real {\\sc Planck} data: non-axially-symmetric beams, detector sample integration, temperature fluctuations induced by the sorption cooler, and pointing errors. Five mapmaking codes of two basic types (destriping and optimal) were used in this study. We generated one year long streams of observations representing four LFI 30~GHz detectors. This simulation round (Trieste simulation) was number four in a series of mapmaking comparison studies that the {\\sc Planck} CTP group has undertaken.  None of our mapmaking codes made an attempt to deconvolve the beam effects from its output map. Therefore the smoothing effects of the beam and sample integration showed up similarly in all our maps. This is, however, a complicated smoothing, because every map pixel has its own effective beam. We also made a thorough examination of the temperature and polarization cross-couplings that arose from the mismatch of the beams. These effects were also similar in all our maps.  Our conclusions are: \\begin{enumerate} \\item Our studies showed that the temperature to polarization cross-coupling of CMB signal caused a detectable bias in the TE spectrum of the CMB map. In the EE spectrum the effect was small compared to the residual noise of the map. The $E\\rightarrow B$ cross-coupling of CMB produces a spurious signal whose magnitude in intermediate and large multipoles (at $\\ell \\geq 300$ in this case) seems to exceed the magnitude of the CMB $B$-mode signal that we might expect to detect with LFI 30~GHz detectors. However, in this range of multipoles the detection of the $B$-mode is difficult due to map noise anyway, and at low multipoles where the signal-to-noise ratio is higher the effect is small. Our study shows that the spectrum biases that $T\\rightarrow P$ and $E\\rightarrow B$ cross-couplings cause may lead to errors in the cosmological studies, if not corrected. We show a method in this paper that can correct part of these effects in the spectrum domain. Map-domain methods that address the beam convolution issues properly have been proposed by other authors (see Sect.~\\ref{subsubsec:correction}).  \\item Signal error (error that mapmaking couples to the output map from the small-scale signal structure) is the part of the output map where beam effects appear differently in the maps of different mapmaking codes. In optimal and short baseline Madam maps this error is larger than in the destriper maps of long baselines (Madam and Springtide). Signal error is, however, a small effect in a high-resolution map if we compare it to the residual noise of the map. Techniques to decrease the signal error were thoroughly examined in our earlier study (Ashdown et al.~\\cite{Ash07b}).  \\item Based on the results of this study, cooler noise or pointing errors of $\\leq 0\\farcm5$ do not appear as a major concern for the 30\\,GHz channel. \\end{enumerate}  Five mapmaking codes (two destripers and three optimal codes) have been developed and compared in four {\\sc Planck} CTP simulation rounds (Cambridge, Helsinki, Paris, and Trieste simulations). We give the following summary of these studies: \\begin{enumerate} \\item At the end of a long process, the only essential difference found so far between codes that affects accuracy---as assessed by the RMS residual between an input binned map and the output map---is baseline length, shown in Fig.~\\ref{fig:kevin}.  This assumes that all maps have been made at sufficiently high resolution that the effects of sub-pixel structure are small (see Ashdown et al.~\\cite{Ash07b}). (Optimal codes are somewhat more affected [in an RMS residual sense] by sub-pixel structure than are destripers, but this is not an issue at sufficiently high map resolution.) However, the fundamental difference between destriping and optimal codes is in their assumptions about, and handling of, noise.   Some of the future tests listed in Sect.~\\ref{subsec:future} will probe this difference more intensively than ones performed to date.  \\item All optimal codes give essentially indistinguishable results. Madam, a destriper with adjustable baseline length, gives the same result as the optimal codes when the baseline is set short enough. Reliable noise estimation is required for these codes.  \\item Resource requirements for the codes vary by more than an order of magnitude in the important categories (memory per processor, total memory, elapsed time, and total CPU-hours), as shown in Table~7 and Fig.~\\ref{fig:cpuh_vs_memory}.   Codes set up to deal with pointing information in multiple forms have a practical advantage, in that the resources required in memory, I/O speed, and processor speed can be matched to the resources available (memory, I/O speed, and CPU time).  \\item It seems inevitable that {\\sc Planck} will need both a fast and ``resource light'' mapmaking code for everyday use, and an ``ultimate accuracy'' but ``resource heavy'' code, for final data products. \\end{enumerate}  \\subsection{Future Tests}\\label{subsec:future}  Our final round of simulations was quite realistic compared to earlier rounds, nevertheless there are still many instrumental and mission effects, both subtle and blatant, that have not been included, some of which are listed below.  Further tests of mapmaking must be performed.  \\begin{itemize}  \\item Gaps in TOD.  \\item High-pass filtering of TOD (e.g., to reduce low frequency noise before mapmaking, to deconvolve the bolometer frequency response in HFI). High-pass filtering will make the noise matrix $\\mathbf{C}_n$ of Eq. (\\ref{eq:binned_map_in_map}) non-diagonal.  \\item Leakage of galactic temperature signal to the polarization maps due to the bandpass mismatch of the LFI radiometers.  \\item Beam sidelobes.  \\item Cross-polarization leakage of the detectors, and polarization angle errors.  \\item Other forms of random noise (e.g., ``popcorn noise'').  \\item Features in the noise power spectrum (e.g., microphonics, frequency spikes). Note that some of these features may make the noise matrix $\\mathbf{C}_n$ of Eq. (\\ref{eq:binned_map_in_map}) non-diagonal.  \\item Changes in the baseline signal level (e.g., from telescope emission, background loading of bolometers).  \\item Realistic pointing errors (errors in the pointing reconstruction)  \\end{itemize}  The scope of this paper was restricted to the mapmaking from the LFI 30~GHz data. The {\\sc Planck} mapmaking codes need, however, to be able to make maps from both LFI and HFI observations. Therefore the list includes mapmaking tests that are relevant for LFI, HFI or both.  Future tests should be run on both a destriping code and an optimal code.  As a practical matter, it will be efficient to run the less resource intensive destriping code first to wring out any problems in setting up the test.   It is no longer necessary to run tests on multiple implementations of destriping or optimal codes, which have been shown to give the same answer.  Generally speaking, destriping algorithms assume that the noise power spectrum is white at high frequency.  The shape of the low frequency power spectrum is not important, so long as the low frequency noise may be fit by a series of offsets (or baselines). The duration of these offsets varies from 1~second to 1~hour in the codes we have tested here, with a trade-off between accuracy and resource usage. For baseline durations longer than 1~minute, destripers need no information about the noise power spectrum.   By contrast, the optimal codes are not restricted to white noise at high frequency, but must be informed of the shape of the noise power spectrum.  Deconvolution of a bolometer time constant exercises a difference in these approaches.  Bolometers have a response lag to a change in the sky similar to the time for the beam to cross a fixed point.  This effect smears the sky image in the direction of the scan.  A time stream deconvolution filter, applied prior to mapmaking, can repair this effect on the sky signal, but at the cost of correlating the noise at small scales.  We have tested the effect of this with simulated HFI 100\\,GHz data, where the bolometer time constants are the longest, and where the effect is the most severe (G\\'orski et al.~\\cite{Gor08})."
    },
    "0806/0806.0602_arXiv.txt": {
        "abstract": "The hierarchical clustering inherent in $\\Lambda$-CDM cosmology seems to produce many of the observed characteristics of large-scale structure. But some glaring problems still remain, including the over-prediction (by a factor 10) of the number of dwarf galaxies within the virialized population of the local group. Several secondary effects have already been proposed to resolve this problem. It is still not clear, however, whether the principal solution rests with astrophysical processes, such as early feedback from supernovae, or possibly with as yet undetermined properties of the dark matter itself. In this paper, we carry out a detailed calculation of the dwarf halo evolution incorporating the effects of a hypothesized dark-matter decay, $D\\rightarrow D^\\prime+l$, where $D$ is the unstable particle, $D^\\prime$ is the more massive daughter particle and $l$ is the other, lighter (or possibly massless) daughter particle. This process preferentially heats the smaller haloes, expanding them during their evolution and reducing their present-day circular velocity. We find that this mechanism can account very well for the factor 4 deficit in the observed number of systems with velocity 10--20 km s$^{-1}$ compared to those predicted by the numerical simulations, if $\\Delta m/ m_{D^\\prime} \\sim 5-7\\times 10^{-5}$, where $\\Delta m$ is the mass difference between the initial and final states. The corresponding lifetime $\\tau$ cannot be longer than $\\sim 30$ Gyr, but may be as short as just a few Gyr. ",
        "introduction": "Observations of the cosmic microwave background (CMB) radiation with the Wilkinson Microwave Anisotropy Probe (WMAP) have facilitated the precision measurement of several cosmological parameters (Bennett et al. 2003, Spergel et al. 2003), including the mass-energy density of the Universe, $\\Omega$, which appears to be close (if not equal) to its critical value. Baryons contribute only about 4\\% of this; the rest is presumably in the form of dark matter (DM; roughly 22\\%) and dark energy ($\\sim$74\\%). In addition, WMAP's detection of early reionization also rules out the presence of a warm DM, so the non-baryonic component must be cold (CDM). Together with earlier observations by other finer scale CMB experiments, and with the Hubble Key Project (Mould et al. 2000), which provided the unprecedentedly accurate value $H=71\\pm6$ km s$^{-1}$ of the Hubble constant, this combined body of work has lead to a consensus that the Universe is adequately described by the so-called flat $\\Lambda$-CDM standard model, in which dark energy is the manifestation of a cosmological constant $\\Lambda$.  The existence of dark energy has been confirmed through the analysis of Type Ia supernova data (Riess et al. 1998; Perlmutter et al. 1999), providing even stronger evidence that this component of $\\Omega$ has negative pressure leading to an acceleration of the Universe in the current epoc. To be fair, however, it is not yet entirely clear whether these results require a cosmology with a true cosmological constant, in which matter and radiation become dominant looking back towards redshifts $z>$2--3, or whether the current acceleration might be due to a so-called scaling solution, in which the dark energy density scales with matter, and affects the formation of structure even at early times (see Melia 2007, 2008, and references cited therein).  Very little is known about dark matter, and almost nothing is understood about dark energy. Their nature is one of the biggest mysteries in contemporary physics. Yet their influence is evidently quite important in the formation of large scale structure (LSS). The hierarchical clustering inherent in $\\Lambda$-CDM seems to produce many of the characteristics observed in the local (e.g., Murali et al. 2002; Abadi et al. 2003) and high-redshift (e.g., Springel, Frenk, and White 2006) Universe. In hierarchical models, smaller dark matter haloes on average collapse earlier than larger ones, when the density of the universe was higher (e.g., Kravtsov et al. 1998). The current mass function, however, is determined not only by the halo formation history, but also by their merger rate which, over time, tends to deplete the dwarf-galaxy end of the distribution.  But some glaring problems still remain. Numerical simulations involving collisionless CDM predict dark haloes with steep cusps in their centre (Navarro et al. 1996), whereas most of the observed rotation curves of dwarf galaxies and low surface brightness galaxies indicate constant density cores (see de Blok et al. 2003, and references cited therein). A related problem is the over-prediction of the number of dwarf galaxies within the virialized population of the local group. CDM simulations over-predict the number of satellite galaxies orbiting a Milky Way-sized galaxy by a factor of 10 (Klypin et al. 1999, Moore et al. 1999, Diemand et al. 2007, and Simon \\& Geha 2007). Generally speaking, both of these problems may be described as a $\\Lambda$-CDM prediction of too much power on small scales.   Since in hierarchical models smaller galaxies merge to make larger ones, the number of remaining dwarf haloes is an important diagnostic to test both the hierarchical picture and the process of halo condensation in the evolving universe. Thus, the dwarf-halo deficit may be taken as an indicator of new physics associated with dark matter (and/or dark energy). For example, the small-scale power can be reduced by substituting warm dark matter for CDM. But as we have seen, WMAP observations have already ruled this possibility out.  Other simple attempts to fix the dwarf-halo deficit problem are not well motivated physically. Some involve altering the fundamental nature of dark matter by introducing self-interaction, or annihilation (see, e.g., Spergel and Steinhardt 2000; Kaplinghat, Knox, and Turner 2000; and Giraud 2001, among others). Without a proper indication from physical considerations, all of these models contain free parameters tunable to fit the observations. For example, the self-interacting dark matter scenario rests on the viability of a huge velocity-dependent cross section. Unfortunately, the implied large interaction rate is inconsistent with most weakly interacting, massive particle and axion theories (see, e.g., Hennawi and Ostriker 2002).  Of course, the dwarf-halo deficit may be due to reasons other than DM physics. Many of them may be invisible because they contain a very small amount of luminous matter, either because of early feedback from supernovae (Dekel and Silk 1986; Mac Low and Ferrara 1999), or because their baryonic gas was heated by the intergalactic ionizing background radiation (Rees 1986; Barkana and Loeb 1999). Others may have turned into high-velocity clouds in the Local Group (see, e.g., Blitz et al. 1999). The deficit may not even be real if the Universe is actually older than its current inferred age, which in reality is only the light-travel time to the cosmic horizon rather than the Big Bang (Melia 2007, 2008). In such a scenario, the dwarf haloes would have had more time to merge, depleting the lower end of the mass function and bringing it into better alignment with the observations.  In any case, it is still too early to tell whether or not the discrepancy between the cosmological simulations and observations really indicates a major problem for hierarchical models in $\\Lambda$-CDM. Several of the effects we have listed here may resolve at least part of the deficit problem. However, given that this is still an open question, observations of the halo mass function also allow us to explore non-astrophysical reasons for the discrepancy, with the goal of learning more about the nature of dark matter.  Our focus in this paper is the suggestion that DM particles may be unstable to decay (see, e.g., Davis et al. 1981; Turner et al. 1984; Turner 1985; Dicus and Teplitz 1986; Dekel and Piran 1987; Sciama 1990; Cen 2001a; S\\'anchez-Salcedo 2003). The impact of the interactions we describe above and/or decays is almost always to provide a mass-dependent expansion of the cusps and haloes to lower the core density and to reduce the number of small galaxies. However, attempts to couple these ideas to particle physics have been few and ambiguous, partly because these have been concerned more with global effects, rather than aimed at finding specific particle properties that may be consistent with the requirements to fix the problem. Our goal in this paper is to begin a more careful search for the properties DM particles must have in order to account for the deficit of dwarf haloes, if in the end their decay is indeed responsible for the observed effect.  Our approach here is closest in spirit to the work of Cen (2001a) and S\\'anchez-Salcedo (2003), though their papers had different goals. Cen's (2001a) primary interest was to demonstrate that if DM particles decay and become relativistic, they escape the virialized halo, whose remaining energy then exceeds that required to sustain virial equilibrium and forces it to expand. His suggestion was that the overproduction of dwarf haloes may be solved not by removing them, but by modifying them into failed, dark galaxies, in which star formation has been quenched due to the effects of evaporation and expansion. This is an intriguing idea, though not yet sufficiently developed to provide a useful probe into the properties of the particles themselves.  S\\'anchez-Salcedo's (2003) goal was to mitigate the disparity between the very steep central cusps in dark haloes of dwarf galaxies predicted by $\\Lambda$-CDM and the relatively flat distributions actually seen in these systems. He demonstrated that if DM decays into a relativistic, nonradiative light particle, plus a stable massive particle with some recoil velocity in the center-of-mass frame, the former escapes the bound system while the latter remains with an energy exceeding that of the parent, forcing the halo to expand.  In this paper, we introduce several new aspects of the DM-decay scenario, including the time-dependent and mass-dependent halo formation and hierarchical-merger rates in order to more accurately gauge the impact of heating on the present-day circular velocity distribution. There are too many aspects of the DM decay to consider in just one set of calculations, so we here restrict our attention to cases in which at least one of the decay products remains within the halo, maintaining its mass, though heating it with the liberated energy. Other regions of the DM particle phase space will be explored elsewhere.  In the next section, we summarize the circular-velocity data and demonstrate the nature of the dwarf-halo deficit problem. Then in \\S 3 we describe a technique for following the mass-dependent formation and destruction of haloes as the universe evolves. In \\S 4 we describe the DM-particle decays and how we incorporate the impact of this process into our calculational algorithm. We present the results of our calculations in \\S 5, and discuss them in \\S 6. ",
        "conclusions": "We have shown that the decay of unstable CDM particles can fully account for the deficit of dwarf galaxies in the local group, and have identified some of the particle properties required to achieve this result. In figure~6, the lifetime $\\tau$ and energy conversion fraction $\\chi$ must have values consistent with the thick black line between the two gray regions in order to reduce the number of dwarf haloes within the range of velocities 10--20 km s$^{-1}$ by a factor $\\sim$ 4, in agreement with the observations (Simon \\& Geha 2007). Broadly speaking, the DM decay model works very well as long as $\\chi \\sim 5-7\\times 10^{-5}$. The lifetime cannot be longer than $\\sim 30$ Gyr, and may be as short as only a few Gyr, depending on the precise value of $\\chi$. Note, however, that we have here restricted our analysis to cases in which the more massive decay particle $D^\\prime$ remains bound to the halo. We may find other regions of $\\chi-\\tau$ phase space that produce reasonable results when this restriction is removed.  We emphasize that this model can reduce power on small scales consistent with the observations without altering the number of Milky Way-sized galaxies. Very importantly, we have shown that although the expansion produced by these decays changes the circular velocity distribution, it does not change the halo mass function, at least not directly. It is beyond the scope of the present paper to seek the ultimate fate of dwarf haloes expanding to circular velocities below 10 km s$^{-1}$, which remains an open question, though several possibilities have been proposed by Cen (2001a, 2001b).  It should be pointed out, however, that in addition to expanding (preferentially) the smaller haloes, and thereby reducing their measurable velocity dispersion, DM decay would also have the effect of speeding up their rate of evaporation within the Milky Way's tidal field. In this way, dwarf haloes would be removed from the overall velocity distribution, not only due to their migration in velocity space towards the low end, but would also be removed entirely due to evaporation. This effect would not be evident with the larger haloes, for which DM decay would have little impact on their condensation (and hence on their velocity profiles).  Interestingly, some limits on DM-decay models have already been established based on limits provided by the cosmic $\\gamma$-ray background. Although it is beyond the scope of the present paper to consider the implications of our work on all possible DM scenarios, it is useful to see how coupling our astrophysically-motivated simulations to the various particle physics proposals could develop in the future. For example, in their consideration of WIMPs decaying to Kaluza-Klein gravitons and gravitinos, Cembranos et al. (2007) demonstrated that both the energy spectrum and flux of the observed diffuse MeV $\\gamma$-ray excess may be explained by decaying DM with $\\sim$MeV mass splittings. In this picture, a decay timescale of 10 Gyr would require a mass splitting of $\\sim 10$ MeV, for which the unstable DM particle would then have a mass $\\sim 1$ TeV within the context of our model. In a second example, Kribs and Rothstein (1997) placed bounds on long-lived primordial relics using measurements of the diffuse $\\gamma$-ray spectrum from EGRET and COMPTEL. They concluded that relics decaying predominantly through radiative channels are excluded for lifetimes between $10^{5}$ and $10^{15}$ years. Since the DM decay timescale in our model fits within this range, the radiative decay of relics such as these could not resolve the dwarf spheroidal problem.  We have kept our analysis semi-analytical in order to better gauge the impact of our assumptions and chosen parameter values. Of course, to get a more accurate set of results, one should couple the properties of decaying DM particles with a more realistic N-body simulation. We intend to carry out such a calculation in the near future and will report the results elsewhere.  \\vspace*{-0.3cm}"
    },
    "0806/0806.0572_arXiv.txt": {
        "abstract": "{The Parkes Galactic Meridian Survey (PGMS) is a $5^\\circ \\times 90^\\circ$ strip to map the polarized synchrotron emission along a Galactic meridian from the Galactic plane down to the south Galactic pole. The survey is carried out at the Parkes radio telescope at a frequency of 2.3~GHz with 30 adjacent 8-MHz bands which enable Faraday Rotation studies. The scientific goal is twofold: (1)  To probe the Galactic magnetism off the Galactic plane of which little is known so far. PGMS gives an insight into the Galactic magnetic field in the thick disc, halo, and disc-halo transition; (2) To study the synchrotron emission as foreground noise of the CMB Polarization, especially for the weak B-Mode which carries the signature of the primordial gravitational wave background left by the Inflation. PGMS observations have been recently concluded. In this contribution we present the survey along with first results.}  \\resumen{}  \\addkeyword{Cosmology: CMB} \\addkeyword{Galaxy: disk} \\addkeyword{Galaxy: halo} \\addkeyword{ISM: magnetic fields} \\addkeyword{polarization}   \\begin{document} ",
        "introduction": "\\label{sec:intro}  The magnetic field is an important component of the Interstellar Medium (ISM) and, as in almost any astrophysical context from stars to cosmology, the physics of the Galaxy cannot be studied without accounting for it.  Observations of other spiral galaxies show that the spiral arms are usually dominated by a turbulent or tangled component. A weaker coherent field aligned with the arms can be  present as well. In the inter-arm regions the regular component is more relevant and depicts  {\\it magnetic arms} with coherent scales up to the size of the disc (e.g. M51, Patrikeev et~al. 2006). The magnetic field in halos is less clear, instead. This is not only because of the weaker signal, which makes hard to detect the synchrotron emission, but also of the many different patterns found so far: from galaxies without evident halo field, to X-shaped fields centred at the galaxy centre (e.g. NGC 891, Krause 2007), and up to large almost spherical magnetic halos (e.g. NGC 4631, Hummel \\& Dettmar 1990).  Studying other galaxies can help understand the case of the Milky Way, but our position internal to the Galaxy makes it harder to figure out how the Galactic magnetic field is structured.  Probes to study magnetic fields and what is known to date about the Galactic field will be briefly reviewed in Sect.~\\ref{sec:milkyway}, while available data of synchrotron polarized emission and requirements for new surveys will be discussed in Sect.~\\ref{sec:need}. Finally, the Parkes Galactic Meridian Survey will be presented in Sect.~\\ref{sec:pgms}. ",
        "conclusions": ""
    },
    "0806/0806.1742_arXiv.txt": {
        "abstract": "The faint radio point sources that are unresolved in cosmic microwave background (CMB) anisotropy maps are likely to be a biased tracer of the large-scale structure dark matter distribution.  While the shot-noise contribution to the angular power spectrum of unresolved radio point sources is included either when optimally constructing the CMB angular power spectrum, as with WMAP data, or when extracting cosmological parameters, we suggest that clustering part of the point source power spectrum should also be included.  This is especially necessary at high frequencies above 150 GHz, where the clustering of far-IR sources is expected to dominate the shot-noise level of the angular power spectrum at tens of arcminute angular scales of both radio and sub-mm sources.  We make an estimate of source clustering of unresolved radio sources in both WMAP and ACBAR, and marginalize over the amplitude of source clustering in each CMB data set when model fitting for cosmological parameters.  For the combination of WMAP 5-year data and ACBAR, we find that the spectral index changes from the value of $0.963 \\pm 0.014$ to $0.959 \\pm 0.014$ (at $68 \\%$ c.l.) when the clustering power spectrum of point sources is included in model fits.  While we find that the differences are marginal with and without source clustering in current data, it may be necessary to account for source clustering with future datasets such as Planck, especially to properly model fit anisotropies at arcminute angular scales. If clustering is not accounted and point sources are modeled with a shot-noise only out to $l \\sim 2000$, the spectral index will be biased by about 1.5$\\sigma$. ",
        "introduction": "As discussed in a variety of papers, unresolved radio point sources are an important foreground in temperature anisotropy maps of the cosmic microwave background (CMB) \\cite{zotti,tof1, wright, tegmark}. With Wilkinson Microwave Anisotropy Probe (WMAP) data \\cite{bennett}, the difference in the CMB power spectra determined at various frequency channels and the cross power spectra between the channels, has allowed the unresolved point source contamination to be constrained with an amplitude $A_{\\rm ps}=(0.011 \\pm 0.001) \\mu$K$^2$-sr \\cite{nolta} for the foreground power spectrum, when scaled to the Q-band.  In the WMAP analysis, this point source amplitude is taken to be a constant in $C_l$, similar to the case of a shot-noise type power spectrum for unresolved point sources. This shot-noise, with an appropriate scaling in frequency, is then removed from each of the power spectra when constructing the final WMAP temperature anisotropy power spectrum \\cite{hinshaw}.  While the WMAP estimate on the point source correction is consistent with a point source power spectrum dominated by the shot-noise, this estimate is dominated by measurements relative to the Q-band \\cite{nolta}.  Since the WMAP temperature anisotropy power spectrum is based on V- and W-band data \\cite{dunkley}, and the point source correction has a larger uncertainty in V- and W- bands \\cite{nolta}, it is unclear if a simple shot-noise correction fully describes point sources in the WMAP temperature anisotropy power spectrum.  To account for uncertainty in the amplitude of point-source shot-noise in parameter estimates, the WMAP likelihood contains an additional marginalization of the uncertainty of $A_{\\rm ps}$, but the best-fit amplitude of point-sources remain fixed to the a priori determined value \\cite{hinshaw}.  We also note that alternative approaches have been considered to estimate the amplitude of point-sources \\cite{Huffenberger1,Huffenberger2}, though these works also concentrated on establishing the shot-noise correction.  Beyond the shot-noise, unresolved radio point sources are likely to have a clustered distribution on the sky as they are expected to be a biased tracer of the large-scale structure. Thus, the angular power spectrum of sources contains not just a shot-noise but also a clustering piece determined by the dark matter power spectrum, point source bias, and the redshift distribution. Existing calculations suggest that the shot-noise part of the power spectrum from bright, rare sources dominates clustering at low radio frequencies, especially when the flux threshold for point source removal is at the level of $\\sim$ 1 Jy \\cite{GonzalezNuevo:2004fj,tof1}.  Thus, the assumption of a shot-noise point source contribution to the angular power spectrum of temperature anisotropies is likely to be adequate for low-frequency bands of WMAP such as the Q-band, but may not be appropriate at high frequencies, such as the W band, whose data are used in the temperature power spectrum.  The approach using a shot-noise spectrum with an uncertainty that is marginalized over in the WMAP likelihood is is bit different from the approach advocated by the WMAP team to account for another foreground in CMB data involving the Sunyaev-Zel'dovich (SZ) effect from galaxy clusters. There, the angular power spectrum is estimated based on a model for the cluster distribution and gas properties \\cite{Komatsu}, with an overall uncertainty in the amplitude of the SZ power spectrum captured by a free parameter which is then freely varied as a nuisance parameter when best-fit cosmological parameter values and their uncertainties using a Markov-Chain Monte-Carlo (MCMC) code \\cite{lewis,spergel,dunkley}.  Since the clustering component of the angular power spectrum of unresolved radio sources may be important, it could be that simply including point-sources as a shot-noise correction in the V/W-band WMAP temperature anisotropy power spectrum results in biased estimates of cosmological parameters, especially for parameters like the spectral index of density perturbations, which has been discussed previously in the context of uncertainties related to point-source shot-noise amplitude \\cite{Huffenberger1}.  Moreover, current CMB analyses make use of the combination of datasets such as WMAP and ACBAR which have different treatments related to how point sources are accounted in the parameter fits. At the ACBAR frequency of 150 GHz \\cite{reichardt}, the clustering of sources may need to be included properly, especially given that the high angular resolution of ACBAR also allows removal of sources down to a lower flux density level, where the shot-noise associated with rare, bright sources may be subdominant. The clustering of point sources could also account for some fraction of the excess arcminute-scale anisotropies detected by ACBAR \\cite{tof2,Cooray3}.  Given the lack of adequate details related to the exact clustering power spectrum of radio sources in datasets such as WMAP and ACBAR, we make a general estimate of the angular power spectrum of radio point-sources and include the overall amplitude of clustering as an additional nuisance parameter to be included and marginalized over when estimating for cosmological parameters. For this, we make use of number counts at 95 GHz \\cite{sadler} and assume the redshift distribution of high-frequency radio sources follows the same distribution as estimated for NVSS \\cite{condon} sources at low frequencies \\cite{Ho}. As some of these assumptions are likely to be invalid to some extent, we do not fix the clustering spectrum to our model but allow the overall amplitude to vary and marginalize over that uncertainty when constraining cosmological parameter. Thus, the uncertainty in our predictions related to clustering of sources is unlikely to dominate and we confirm this by noting that the differences to best-fit parameter values with point source clustering included are not significant, especially for the case of combined WMAP and ACBAR data.  Moreover, our estimate of clustering is consistent with the allowed level of point source correction in V- and W-band combination of WMAP, as measured in terms of differences in the power spectra \\cite{nolta}.  While the differences in estimated cosmological parameters are small, a proper estimate of cosmological parameters with clustering included is useful for a proper statistical analysis on important scientific results such as on the extent to which the spectral index of density perturbations $n_s$ is different from the Harrison-Zel'dovich value at -1.  While the impact on current datasets is small, for future data such as Planck that probe down to smaller scales over a wide range of frequencies, we suggest that it will be necessary to account for clustering of point sources when model fitting cosmological parameters.  This paper is organized as follows: we first discuss the angular clustering power spectrum of radio points sources.  Section~III discusses model fits to recent CMB anisotropy data from WMAP and ACBAR. We discuss our results and conclude with a summary in Section~IV. ",
        "conclusions": "The faint radio point sources that are unresolved in cosmic microwave background (CMB) anisotropy maps are likely to be a biased tracer of the large-scale structure.  While the shot-noise contribution to the angular power spectrum of radio point sources has been considered so far when extracting cosmological parameters with CMB data, we have shown here that one should also allow for the possibility of source clustering. This is especially necessary at high frequencies where the clustering of sources is expected to dominate the shot-noise level of the angular power spectrum at tens of arcminute angular scales.  As we find, the differences seen by the WMAP team for V and W-band angular power spectra do allow point source clustering, though one can wrongly conclude clustering is unnecessary if lower frequency data are included.  Here, we have made an estimate of the clustering of unresolved radio sources in both WMAP and ACBAR by making use of existing counts at 95 GHz and by making several assumptions on the sources such as the redshift distribution.  To account for the uncertainty in modeling the clustering, we included an extra nuisance parameter for each dataset and have marginalized over this parameter when model fitting for cosmological parameters.  For the combination of WMAP 5-year data and ACBAR, we find that the spectral index changes from a mean value of $0.963 \\pm 0.014$ without point-source clustering to a value of $0.959 \\pm 0.014$ when the clustering of point sources are included in model fits, a difference of $0.3\\sigma$.  We also discussed the full parameter set with clustering of radio point sources and changes to additional parameters such as $dn_s/d\\ln k$ and the tensor-to-scalar ratio $r$.  While we find that the differences are marginal with and without source clustering in current data, we have suggested that it is necessary to account for source clustering with future datasets such as Planck, especially to properly model fit anisotropies at arcminute angular scales and using high-frequency data. For Planck, we find that simply including the point sources as a shot-noise only out to $l$ of 2000 for cosmological parameter estimation results in biases at the level of 1.5$\\sigma$. While we simply model Planck point sources with a single power spectrum, since at high frequencies both radio and far-IR  sources are expected to contribute, it may be necessary to return to a proper model of total unresolved source clustering in Planck in future.  \\begin{center} {\\bf Acknowledgments} \\end{center} This work was supported by NSF CAREER AST-0645427. We thank Mike Nolta for clarifying our questions related to the point source modeling by the WMAP team. AC thanks Dipartimento di Fisica and INFN, Universita' di Roma-La Sapienza and Aspen Center for Physics for hospitality while this research was completed. AA acknowledges partial support from a McCue fellowship from the UCI Center for Cosmology."
    },
    "0806/0806.3747_arXiv.txt": {
        "abstract": "We present the analysis of a Suzaku observation of the Ophiuchus galaxy cluster. We confirmed that the cluster has a cool core. While the temperature of the intracluster medium (ICM) decreases toward the center, the metal abundance increases. Except for the core ($r\\lesssim 50$~kpc), the cluster is hot ($\\sim 9$--10~keV) and is almost isothermal for $r\\lesssim 1$~Mpc; the latter contradicts a previous study. We do not detect the variation of the redshift of the ICM in the cluster; the upper limit of the velocity difference is $3000\\:\\rm km\\: s^{-1}$.  The iron line ratios in X-ray spectra indicate that the ICM has reached the ionization equilibrium state. From these results, we conclude that the Ophiuchus cluster is not a major merger cluster but one of the hottest clusters with a cool core. We obtain the upper limit of non-thermal emission from the cluster, which is consistent with both the recent claimed detection with INTEGRAL and the recent upper limits with the Swift/BAT. If the cluster has bright non-thermal emission as suggested by the INTEGRAL measurement, it is probably not due to a recent major cluster merger. ",
        "introduction": "The cold dark matter (CDM) model predicts that the typical mass of virialized objects in the Universe increases with time. This means that the most massive or hottest galaxy clusters ($kT\\gtrsim 10$~keV) should be the objects that have formed most recently. In fact, X-ray observations have shown that the hottest clusters are often `major-merger' clusters. These clusters are now growing through the mergers. For example, cluster 1E~0657--56 ($kT=14.8_{-1.2}^{+1.7}$~keV) has a clear bow shock associated with a merger; a smaller cluster is penetrating a larger cluster with a velocity of 3000--$4000\\rm\\: km\\: s^{-1}$ \\citep{mar02}. On the other hand, these hottest clusters are rarely `cool core clusters', which have a distinct cool core and are almost isothermal outside the core. This is probably because cool cores are destroyed by cluster mergers \\citep{ric01,fuj02}.  Cluster mergers may accelerate relativistic particles at shocks and in turbulent regions in the ICM \\citep{jaf77,rol81,sch87,sar99,tak00,fuj01,bru01,min01,ohn02,fuj03,ryu03,bru05,cas05,ino05,kan07}. With BeppoSAX, non-thermal hard X-ray emission from relativistic particles has been detected from several clusters such as the Coma cluster \\citep{fus04}. Most of them are actually merging clusters \\citep{nev04}. However, the detections are still controversial \\citep{ros04}, and confirmation with other instruments would be highly desirable.  The Suzaku HXD has an improved sensitivity for hard X-rays \\citep{tak07}, and attempts have been made to detect non-thermal hard X-ray emission from clusters. So far, no firm detections have been reported \\citep{kaw08}.  The Ophiuchus cluster is one of the hottest known clusters. It has been observed with several satellites, such as HEAO~1 \\citep{joh81}, EXOSAT \\citep{arn87}, Tenma \\citep{oku88}, and GINGA \\citep{kaf92}. Previous observations with ASCA and BeppoSAX showed that the temperature of the ICM is $\\sim 10$~keV \\citep{wat01,nev04}. In addition to the high temperature, this cluster has a few interesting features. First, it has been reported to have a complicated temperature distribution, in particular, the temperature of its western part reaches $\\sim 20$~keV \\citep{wat01}. This temperature is comparable to that observed in 1E~0657--56 \\citep{mar02}. Thus, this cluster may be undergoing a major merger that induces the motion of the ICM with a velocity of $\\gtrsim 3000$ km s$^{-1}$.  Second, non-thermal hard X-ray emission may have been detected from the cluster.  \\citet{nev04} searched for non-thermal emission with BeppoSAX and found some evidence at a $2\\sigma$ level. Recently, \\citet{eck08} detected non-thermal emission from the center of the cluster with INTEGRAL at a $6.4\\sigma$ level.  However, a recent long exposure with the Swift/BAT instrument failed to detect any nonthermal emission, and set an upper limit below the claimed INTEGRAL detection \\citep{aje08,oka08}.  In this paper, we report our observational results of the Ophiuchus cluster with Suzaku. Thanks to the low background of Suzaku \\citep{mit07}, we can obtain an X-ray spectrum even in regions far from the cluster center; we discuss the temperature and metal abundance distributions in the ICM. The high spectral resolution of the XIS \\citep{koy07} on board Suzaku also enables us to measure or limit the bulk motion of the ICM. Moreover, the hard X-ray detector, HXD, can constrain the non-thermal emission from the cluster, which can be compared with the INTEGRAL and Swift/BAT results.  In this paper, we assume cosmological parameters of $\\Omega_0=0.3$, $\\lambda_0=0.7$, and $H_0=70\\rm\\: km\\: s^{-1}\\: Mpc^{-1}$. Since the cluster redshift is $z=0.028$ \\citep{lah89}, one arc-minute corresponds to 33~kpc. All statistical errors are given at the 90\\% confidence level unless otherwise mentioned. ",
        "conclusions": "\\subsection{Comparison with ASCA Results}  Although the Suzaku results are consistent with the previous ASCA results for the central region of the cluster (table~\\ref{tab:spec}), they are not for the outer regions. The results of ASCA indicated that the Ophiuchus cluster has a large temperature variation \\citep{wat01}. In particular, the temperatures of the western and southern regions were reported to be extremely high ($kT\\sim 15$--20~keV). On the other hand, our Suzaku results indicate that the cluster is almost isothermal ($kT\\sim 9$--10~keV).  One possible cause for the difference would be a problem with the X-ray background for the ASCA observation.  While \\citet{wat01} used a background obtained from a source-free region near the cluster (H~1705--250), we constructed a background based on a model of the Galactic emission$^1$ plus a model for the hard CXB.  Since the western and southern regions are dim, the difference of the background could affect the results. However, we found that this is not likely to be the case. We observed the region around H~1705--250 with Suzaku with an exposure time of 15~ks. Then, we used the spectrum of this region as the background for a spectral fit of region W. In this case, the temperature is $kT=9.7_{-1.0}^{+0.9}$~keV, which is consistent with that in table~\\ref{tab:spec}. Another possibility is that the high temperature in these faint outer regions is due to stray light from the bright central regions of the cluster. For the X-ray telescopes carried on ASCA, such problem was much more serious for ASCA than for Suzaku \\citep{ser07}. Although \\citet{wat01} took account of the effect in their analysis, the correction might not have been sufficient.  Although we could not find the cause of the inconsistency, we believe that the results of Suzaku are more reliable, because the problem of the stray light has been much improved.   \\subsection{The Cool Core} \\label{sec:cc}  The existence of a cool core (section~\\ref{sec:depro}) and the isothermality of the cluster (section~\\ref{sec:spec}) indicate that the Ophiuchus cluster is a cool core cluster. While many of clusters with medium temperatures ($\\sim 5$~keV) are cool core clusters having short central cooling times, most of the hottest clusters with temperatures similar to the Ophiuchus ($\\gtrsim 10$~keV) do not seem to have cool cores (e.g. figure~18c,d in \\cite{ota04}). This is probably because the latter are major-merger clusters in which the core has been destroyed and the temperature has been boosted by the merger \\citep{ran02}. For example, the extended HIFLUGCS sample includes 106 nearby clusters \\citep{rei02}. Among them, there are only 8 clusters, including the Ophiuchus cluster, with temperatures larger than 9.0~keV \\citep{che07}. Five of them (A~1689, A~1914, A~2163, A~2319, and Triangulum Australis) do not have a prominent cool core \\citep{and04,gov04,whi00}. Although A~754 and A~2142 have a cool core, they are now undergoing a major merger \\citep{mar00,gov04}. Thus, the Ophiuchus cluster seems to be exceptional in having a very high temperature and a cool core with no strong evidence of a major merger. Since clusters with cool cores tend to be old (e.g. \\cite{fuj00b,ota06}), the Ophiuchus cluster would be exceptionally old given its large mass. Contrary to merging clusters in which star formation activities might be induced in the disk galaxies (e.g., \\cite{fuj99a,fuj99}), the environment in the Ophiuchus cluster might have been calm for a long time. Thus, it would be interesting to study the nature of the galaxies inside it.  \\subsection{Non-Thermal Hard X-ray Emission}  Most of the clusters from which non-thermal X-ray emission has been detected are major-merger clusters \\citep{nev04}. Since non-thermal synchrotron radio emission has often been observed from these clusters \\citep{gov04b}, the non-thermal hard X-ray emission could be attributed to inverse Compton scattering of photons by the relativistic electrons that are responsible for the radio emission. Since the cooling time of these electrons ($\\sim 10^8$~yr) is much shorter than the dynamical time-scale of a cluster ($\\sim 10^9$~yr; \\cite{sar99}), the electrons are likely to be primary electrons, which are now being directly accelerated at shocks or turbulence in the ICM induced by recent cluster mergers.  The hard X-ray excess detected by \\citet{eck08} from the Ophiuchus cluster is curious in this regard. Although a cluster merger does not always destroy a cool core (e.g. A~2142; \\cite{mar01}) and our temperature map (figure~\\ref{fig:TZ}a) is coarse, the X-ray image and spectra from this cluster do not show any strong evidence for a major merger. One possibility is that the Ophiuchus cluster is currently undergoing an early stage, small impact parameter (head-on) merger with the merger axis being close to our line of sight.  Then, the two merging subclusters would be projected on top of one another, and no azimuthal temperature variations would necessarily be expected.  By the same argument, one would not necessarily expect any azimuthal variations in the mean redshift. Although the likelihood of a head-on collision along the line of sight is small, this hypothesis would appear to be consistent with the apparent regularity of Ophiuchus, the presence of a cool core (albeit without any very cool gas), and the presence of nonthermal hard X-ray emission due to a merger shock or turbulence.  On the other hand, if Ophiuchus really is not a merging cluster, then the nonthermal hard X-rays cannot be attributed to primary electrons accelerated through a recent major cluster merger, although we cannot rule out particle-acceleration by turbulence in the core generated by a minor merger \\citep{fuj04,fuj05,maz08}. One possibility is that particles might have been accelerated by the central AGN rather than a cluster merger, because the hard X-ray emission is observed at the center of the Ophiuchus cluster \\citep{eck08}. Moreover, since recent strong AGN activities have not been observed at the cluster center \\citep{dun06}, the non-thermal emission from the cluster may come from secondary electrons that are generated through proton-proton interaction \\citep{pfr04}. In this case, the non-thermal emission can be observed even long after the AGN activities and proton acceleration cease, because the cooling time of the protons is longer than the age of a cluster. Recently, \\citet{fuj07c} proposed that the protons are accelerated at the forward shock around a cocoon created by an AGN outburst.  Another possible source of the non-thermal X-ray emission would be dark matter annihilation (\\cite{pro08}, see also \\cite{tot04}). In any case, a confirmation of the non-thermal emission in Ophiuchus with other instruments would also be very useful.  The Swift/BAT instrument observed Ophiuchus for 1.3 Msec, and failed to detect any non-thermal hard X-ray emission \\citep{aje08,oka08}. Based on a joint Chandra and Swift/BAT analysis (with Chandra used to constrain the thermal emission from the cluster), \\citet{aje08} found an upper limit of $5.2 \\times 10^{-13}\\rm\\:erg\\: s^{-1}\\: cm^{-2}$ (20--60 keV), which is inconsistent with the INTEGRAL detection at more than 2$\\sigma$.  We note that \\citet{eck08} fit the thermal emission from the cluster with with a single temperature model with a temperature of 8.5~keV.  We find that most of the cluster is hotter than this, in agreement with the results from \\citet{aje08} and \\citet{oka08}. The difference with \\citet{eck08} may be due to the cool core in Ophiuchus.  It is possible that a more complex thermal model may be needed for Ophiuchus, and might bring the INTEGRAL and Swift/BAT results into better agreement.  Given the present disagreement of the INTEGRAL and Swift/BAT spectral results, the lack of evidence for a merger, and the lack of a modern radio detection of a halo or relic, we view the non-thermal detection in Ophiuchus as tentative."
    },
    "0806/0806.1929_arXiv.txt": {
        "abstract": "We present an analysis of the energetics and particle content of the lobes of 24 radio galaxies at the cores of cooling clusters. The radio lobes in these systems have created visible cavities in the surrounding hot, X-ray-emitting gas, which allow direct measurement of the mechanical jet power of radio sources over six decades of radio luminosity, independently of the radio properties themselves. Using these measurements, we examine the ratio between radio power and total jet power (the radiative efficiency).  We find that jet (cavity) power increases with  radio synchrotron power approximately as $P_{\\rm jet} \\sim L_{\\rm radio}^{\\beta},$ where $0.35 \\le  \\beta \\le 0.70$ depending on the bandpass of measurement and state of the source.    However, the scatter about these relations caused by variations in radiative efficiency  spans more than four orders of magnitude. A number of factors contribute to this scatter including aging, entrainment, variations in magnetic field strengths, and the partitioning of energy between electrons and non-radiating heavy particles.   After accounting for variations in synchrotron break frequency (age), the scatter is reduced by $\\approx 50$\\% , yielding the most accurate scaling relation available between the lobe bolometric radio power and the jet (cavity) power.  We place limits on the magnetic field strengths and particle content of the radio lobes using a variety of X-ray constraints.   We find that the lobe magnetic field strengths vary between a few to several tens of microgauss depending on the age and dynamical state of the lobes.  If the cavities are maintained in pressure balance with their surroundings and are supported by internal fields and particles in equipartition, the ratio of energy in electrons to heavy particles ($k$) must vary widely from approximately unity to $4000$, consistent with  heavy (hadronic) jets.  The synchrotron ages are substantially smaller than the independently measured X-ray buoyancy ages. Low frequency radio observations provide the best tracer of AGN power integrated over the past several hundred million years, and may provide a useful proxy for X-ray cavities in systems where cavities cannot be directly detected but are likely to exist (e.g., in cooling-flow clusters). ",
        "introduction": "A source of heat has been invoked to explain the lack of gas cooling below 2 keV at the centers of galaxy clusters \\citep{pete01,kaas04}, which was expected in the classical cooling flow model \\citep{fabi94}. The same source may prevent the formation of excessively bright cluster central galaxies \\citep{bens03,bowe06,crot06,sija06}. The most plausible candidate is the energy injected by relativistic jets originating near supermassive black holes (SMBH) at the centers of the clusters \\citep{tabo93,binn95,tuck97}. When the jets interact with intracluster medium (ICM) two lobes of plasma are inflated, which contain relativistic electrons and magnetic fields seen in radio observations. Such interactions have been discussed extensively in the literature \\citep{sche74,blan74,bege89,kais97,hein98}.  The idea that the central active galactic nuclei (AGN) are the source of energy has been given impetus by \\emph{Chandra} X-ray Observatory images showing the keV gas being displaced by the radio sources harbored at the center of the cluster \\citep{mcna00,fabi00}. Some recent analysis of samples of systems with X-ray cavities show that the cavities are long lasting and that they provide enough energy to suppress cooling in clusters \\citep{birz04,dunn04,dunn05,raff06}.  By considering only the cavity enthalpy, \\citet{birz04} and \\citet{raff06} show that the current outburst in more than half of the systems in their sample of cavity systems are powerful enough to offset cooling, although the integrated energy is sufficient to offset cooling in all systems \\citep{mcna07}. However, the enthalpy is only a lower limit for the total energy injected with a cavity. Evidence is mounting for the existence of weak but powerful shocks in the hot thermal gas. Notable examples include MS 0735+7421 \\citep{mcna05}, Perseus \\citep{fabi06}, M87 \\citep{form05,form07}, Hydra A \\citep{nuls05a,wise07} among many others. Where they have been measured, shock energies are comparable to cavity enthalpies. In addition, the energy is still probably underestimated due to further adiabatic losses, undetected cavities, etc., the energy of which we cannot quantify \\citep[e.g.,][see McNamara \\& Nulsen 2007 for a review]{nuss06,binn07}.  Cavities provide a lower limit to the jet energy and are the best-available gauges of the average jet power in these systems; therefore, they can be used, together with the radio data, to estimate the radiative efficiencies and radio-lobe properties. The total energy in the radio lobes is the sum of magnetic and particle energies: \\begin{equation} E_{\\rm{tot}} = E_B + E_{\\rm p}=E_B+(1+k)E_e\\label{energy_eq}, \\end{equation} where $k$ is the ratio of the energy in heavy particles to that in the electrons ($E_{e}$) and $E_{B}$ is the energy in the magnetic field. The energies in the magnetic field and electrons are calculated as: \\begin{equation} \\label{eq:eb} E_{B}=\\frac{B^2}{8\\pi}\\phi V \\end{equation} \\begin{equation} E_{e}=B^{-3/2}L_{\\rm{rad}}c_{12}(\\alpha, \\nu_{1}, \\nu_{2}) ,\\label{Ee} \\end{equation} where $\\phi$ is the filling factor (the fraction of the volume filled by relativistic particles and magnetic field), $c_{12}(\\alpha, \\nu_{1}, \\nu_{2})$ is a constant as defined by \\citet{pach70} and $L_{\\rm{rad}}$ is the total radio luminosity. For the first time, the cavities allow one to estimate $E_{\\rm{tot}}$ of radio lobes to within a factor of a few (given the uncertainties in projection and the omission of shock energy). Additionally, cavities can be used to estimate the pressures (assuming they are in pressure balance with the surrounding ICM) and ages of the lobes (assuming that the cavities rise buoyantly or at the local sound speed). Together, these quantities break the degeneracies in the energy equation and allow meaningful constraints to be placed on the magnetic field strengths and particle content of the lobes.  From radio observations we know that relativistic electrons and magnetic field are present in the jets. Under the assumption of charge neutrality, there must be positively charged particles beside relativistic electrons, such as protons \\citep{cell93} or positrons \\citep{reyn96,ward98,kino04,dunn06b}. The positively charged particle content of the jets is observationally inaccessible and the existence of proton-electron jets or positron-electron jets is inferred using indirect arguments. However, the discovery of X-ray cavities has made it possible to quantify the particle content, by calculating the ratio between the total particle energy and the energy in relativistic electrons and positrons \\citep{fabi02,dunn04,dunn05}.  In this paper, we use VLA\\footnote{The VLA (Very Large Array) is a facility of the National Radio Astronomy Observatory (NRAO). The NRAO is a facility of the National Science Foundation operated under cooperative agreement by Associated Universities, Inc.} radio data at 4 different frequencies for a sample of 24 systems with X-ray cavities to investigate radiative efficiencies.  We derive synchrotron break frequencies and ages.  Using both X-ray and radio data for the cavities, we place limits on their magnetic field strengths and on their particle content (i.e., the energies in heavy-particles and in synchrotron radiating electrons and positrons). We investigate the effects that differences in the radio source ages or magnetic field strengths and particle content have on the efficiencies.  Using both the X-ray and radio data, we investigate the radio source ages by comparing synchrotron ages with buoyancy ages. In general, the synchrotron ages are lower limits on the radio source age \\citep{eile96, eile97,eile99} due to the inhomogeneities in the magnetic field, which are difficult to estimate \\citep{siah90, witt90}. The X-ray derived ages (e.g., buoyancy ages) assume that the bubbles detach early from the driving jet and rise buoyantly in the ICM. By finding discrepancies between the two ages, we can test the assumptions that go into their determination. Lastly, on a related topic, we are now able to determine more precisely whether a cavity is radio filled or ghost.  We adopt $H_{0}=70$ km s$^{-1}$ Mpc$^{-1}$, $\\Omega_{\\rm{M}}=0.3$, and $\\Omega_{\\Lambda}=0.7$ in all calculations throughout this paper. ",
        "conclusions": "In summary, using an X-ray and radio study of 24 systems with X-ray cavities, we found that these systems are in general poor radiators with a ratio of cavity (jet) power to radio power ranging from a few to a few thousands. We conclude that this scatter is probably due to several effects including aging and entrainment plus a combination of intrinsic differences in magnetic field strength and particle content. Accounting for the dependence of the radiative efficiencies on aging reduces the scatter (standard deviation) of the relationship between radio luminosity and average jet power by $\\approx 50$\\%, improving the accuracy with which one can determine the total jet power of the AGN using radio observations alone. However, age estimates are costly to obtain, as an accurate estimate of the break frequency requires observations at several different frequencies. When data at only a single frequency are available and in the absence of an estimate of the break frequency, we found that low-frequency (327 MHz) monochromatic radio power allows a more accurate estimate of the jet power than the higher-frequency 1400 MHz powers (the scatter at 327 MHz is $\\approx 20$\\% lower than at 1400 MHz). Therefore, low-frequency radio observations may be the cheapest strategy for studying feedback in radio galaxies.  Our data allow us to put limits on the particle content and, through comparisons of the buoyancy and synchrotron ages, to test the assumptions that go into their determination. We found that the two age estimates are weakly coupled. We argue that the synchrotron ages underestimate the true ages for the younger systems (with high break frequencies), which likely are actively being inflated by the radio source and which generally have a lower value of $k$. On the other hand, for the older systems (with low break frequencies), the cavity ages and synchrotron ages are closer to being in agreement, and $k$ is generally much higher.  Using both radio and X-ray data we placed limits on the magnetic field strengths and particle content in the lobes ($k,$ defined as the ratio between the heavy-particle energy and the electron energy). Assuming a lepton jet ($k=0$), we found that the magnetic field strengths required for pressure equilibrium are in general much larger than both the equipartition magnetic field strengths and those derived by requiring the synchrotron age to be equal to the buoyancy age.  Therefore, either equipartition does not apply in the lobes of the radio sources or there is additional pressure support from other sources. If we allow for heavy jets ($k>0$), equating the equipartition magnetic field strength to that required for pressure equilibrium gives values of $k$ between approximately unity and 4000, a range similar to that found by \\citet{dunn04} and \\citet{dunn05}. The range in $k$ may be due to the effects of aging and entrainment.  Our multiwavelength radio data allow us to better classify the cavities based on their radio emission. Motivated by the traditional classification system based on the presence or absence of 1400 MHz emission in the cavity, we used the break frequency to classify the cavities as ghost  ($\\nu_{\\rm C} \\lesssim 1400$ MHz) and active ($\\nu_{\\rm C} \\gtrsim 1400$ MHz). Based on this classification, only M84, M87, Hydra A, Cygnus A, and Centaurus have radio-filled cavities.  However, we note that since our systems have a continuous range in break frequencies, any separation of the cavities into discrete categories is somewhat arbitrary.  Lastly, by comparing the cavity sizes inferred from the X-ray maps with those inferred from 1400 MHz and 327 MHz radio maps, we found that the sizes from our 1400 MHz images often under-predict that cavity energy, whereas the lobes sizes measured from our 327 MHz images generally equal or exceed the X-ray cavity sizes. Where the low-frequency lobe emission appears to extend far beyond the X-ray cavities, there may exist faint outer cavities which would only be visible in very deep X-ray images (such as the outermost cavities in Hydra A).  The discovery of X-ray cavities in the hot atmosphere of clusters has made it possible to estimate the radiative efficiencies of the radio sources, which can be much lower than the theoretical estimates (as low as 0.0001). Therefore, most of the AGN's energy in such radio galaxies is dumped directly into the ICM through the X-ray bubbles and shocks, likely preventing the overcooling of gas below $\\sim 2$ keV which would lead to an over-abundance of extremely bright galaxies at the centers of clusters \\citep{sija06,crot06}."
    },
    "0806/0806.1432_arXiv.txt": {
        "abstract": "Spectral analysis by means of Non-LTE model-atmosphere techniques has arrived at a high level of sophistication: fully line-blanketed model atmospheres which consider opacities of all elements from H to Ni allow the reliable determination of photospheric parameters of hot, compact stars. Such models provide a crucial test of stellar evolutionary theory: recent abundance determinations of trace elements like, e.g., F, Ne, Mg, P, S, Ar, Fe, and Ni are suited to investigate on AGB nucleosynthesis. E.g., the strong Fe depletion found in hydrogen-deficient post-AGB stars is a clear indication of an efficient s-process on the AGB where Fe is transformed into Ni or even heavier trans iron-group elements. We present results of recent spectral analyses based on high-resolution UV observations of hot stars. ",
        "introduction": "\\label{sect:introduction}  In the last decades, our picture of post-AGB stellar evolution has been greatly improved. The ``standard'' evolution, i.e\\@. the hydrogen-rich sequence, has been understood in the early eighties of the last century by comparison of spectral analysis and evolutionary models. The spectral analysis of the hottest post-AGB stars with effective tempe\\-ratures ($T_\\mathrm{eff} > 100,000\\,\\mathrm{K}$) was hampered by the lack of appropriate model atmospheres which considered deviations from the local thermodynamic equilibrium (LTE) playing an important role in the photospheres of these stars. The development of such Non-LTE model atmospheres is briefly summarized in Sect.~\\ref{sect:models}.  Spectral analyses of hot post-AGB stars have shown then that about a quarter of these are hydrogen-deficient \\citep{werner06}. The ``born-again post-AGB star'' scenario by \\citet{iben83} is -- in general -- able to explain the evolution of these stars. A final thermal pulse (TP, re-ignition of the helium shell) brings the star back to the AGB and it has a second, helium-burning  post-AGB phase. However, the mechanism to dispose of the entire hydrogen-rich envelope was unclear. Present evolutionary calculations which consider the mixing and burning processes during the helium-shell flash in detail, are able to explain the hydrogen deficiency.  The amount of remaining hydrogen is depending on the particular time when the TP occurs. Still on the AGB (AGB Final Thermal Pulse, AFTP), the masses of the hydrogen-rich envelope and the helium-rich intershell layer (Fig.~\\ref{fig:structure}) are about equal, being roughly $ 10^{-2}\\,\\mathrm{M_\\odot}$. The TP mixing results in $\\approx 40\\,\\%$ hydrogen content (by mass) at the stellar surface. This is detectable in the observations. After the departure from the AGB, the mass of the hydrogen-rich envelope is much lower ($M_\\mathrm{H} \\approx 10^{-4}\\,\\mathrm{M_\\odot}$). If the TP occurs at still constant (high) luminosity (Late Thermal Pulse, LTP), i.e\\@. nuclear burning is still ``on'', the mixing during the flash reduces the surface hydrogen to $\\approx 1\\,\\%$ which is not detectable at the high surface gravity ($\\log g \\approx 5 - 8$ in cm/sec$^2$). A TP at already declining luminosity, i.e\\@. there is no entropy barrier due to the hydrogen-burning shell, will mix down the hydrogen-rich envelope to the bottom of the now re-ignited helium shell where hydrogen is burned and the star becomes hydrogen free.   \\begin{figure} \\includegraphics{rauch_01.eps}\\centering \\caption{Inner structure of an AGB star. Note that the convection zone from the surface to the bottom of the helium-burning shell is established during a TP only.} \\label{fig:structure} \\end{figure}   The mixing process during the TP brings intershell matter up to the stellar surface and, thus, allows a direct view on it. Spectral analysis allows then to conclude on details of both, mixing as well as nuclear processes like, e.g\\@. the s-process, in AGB stars. This is a crucial test for stellar evolutionary models. ",
        "conclusions": ""
    },
    "0806/0806.3279_arXiv.txt": {
        "abstract": "We revisit  the Milky Way satellite problem  using a semi-analytical model  of  galaxy formation  and  compare  the predicted  luminosity function   to   recent   result   from  the   SDSS.    With   cosmic photoionization, the  luminosity function can be  brought into broad agreement with  the data between $-15< M_{V}  <-2$. This improvement over  previous semi-analytical  model results  (e.g., Benson  et al. 2002)  is from  our adoption  of improved  models for  galaxy merger history and  galaxy merging time-scales.  The  very faint satellites ($M_{v}  >  -5$)  formed  in  halos  with  virial  temperature  over $10^{4}K$ (  mass around $10^{9} M_{\\odot}$ before accretion), but    their   baryon   content    are   strongly    suppressed   by photoionization. We  model the mass  evolution of the  subhalos, and compare the  predicted mass-to-light ratio  with the data.   We find that the measured total mass inside the luminous radii of satellites are about 5\\% of their present total dark matter mass. ",
        "introduction": "One of  the most mysterious of  our Milky-Way Galaxy  is the ``missing satellite problem'' that cold dark matter scenario predicts its number of  satellite galaxies  to  be  around hundred,  but  only dozens  are observed  (e.g.  Moore  et  al.   1999).  A  few  solutions have  been proposed for  this contradictory.  One  instant question is if  we are comparing the  right things, i.e,  how to relate the  observed stellar velocity  dispersion to  the  measured circular  velocity from  N-body simulation (Hayashi  et al.  1999). Another possible  solution is that photoionization suppress  the gas accretion  in small halos,  and only halos  formed  before  reionization  can form  stars.   (e.g.   Gnedin 2000). Also there is worry  about the incompleteness of observation as very faint satellites have too  low surface brightness to be observed. In recent years,  more faint satellites are observed  along with their kinematics  and  mass   measurements,  and  the  satellite  luminosity function is also well determined from the SDSS (Koposov et al.  2008). One  the  other  hand,  theoretical  modelling  of  dark  matter  halo formation history  (Cole et al.  2008) and galaxy  merging time-scales (Jiang et  al. 2008,  Boylan-Kolchin et al.   2008) are  also improved. Given  these  progress,  it  is  deserved  to  revisit  the  ``missing satellite problem'' and  there have been a few  papers to address this (e.g. Simon \\& Geha 2007).  Here we  use the  semi-analytical model of  galaxy formation  (Kang et al. 2005) to predict the luminosity function, mass-to-light ratios for satellite galaxies and compare them with the data. ",
        "conclusions": "we revisit  the ``missing satellite problem''  using a semi-analytical model of galaxy formation  combined with a high-resolution merger tree from  Monte-Carlo algorithm (Parkinson  et al.   2008).  we  model the luminosity function  and mass-to-light ratio for  the satellites.  The model  luminosity function  agrees  well with  the  recent results  of Koposov et al.   (2008) from the SDSS DR5  only if the photoionization effect is included to suppress the gas fraction in less massive halos. Our ability to produce  more luminous satellite galaxies than previous semi-analytical  models  (e.g.,  Benson  et  al.  2002)  is  from  the improvement  on  modelling of  halo  merger  tree  and galaxy  merging time-scales.  Very  faint satellites  ($M_{V}>-5$) form in  halos with virial temperature above $10^{4}K$, but their gas content are strongly suppressed by photoionization. In addition their number density favors an  isothermal density distribution  in the  Milky-Way.  We  model the mass evolution of subhalos using  the model of Giocoli et al.  (2008), and find  that the measured  total kinematic mass inside  the luminous radii of satellite galaxies are about 5\\% of their present dark matter mass."
    },
    "0806/0806.3753_arXiv.txt": {
        "abstract": "We investigate the impact of neutron capture rates near the $A=130$ peak on the $r$-process abundance pattern.  We show that these capture rates can alter the abundances of individual nuclear species, not only in the region of $A=130$ peak, but also throughout the abundance pattern.  We discuss the nonequilibrium processes that produce these abundance changes and determine which capture rates have the most significant impact. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0806/0806.4109.txt": {
        "abstract": "In this paper, we discuss the non-Gaussianity originated from the $\\alpha$-vacuum on the CMB anisotropy. For $\\alpha$-vacuum,  there exist correlation %are highly correlated background states between points in the acausal two patches of de Sitter spactime. Such kind of correlation can lead to large local form non-Gaussianity in $\\alpha$-vacuum. For the single field slow-roll inflationary scenario, the spacetime is in a quasi-de Sitter phase during the inflation. We will show that the $\\alpha$-vacuum in this case will lead to non-Gaussianity with distinguished feature, of a large local form and a very different shape. % nonGaussianity, %which has than Meanwhile, we find that the shape of non-Gaussianity %sensitively depends on the background of spacetime. De Sitter space %and inflationary background will give dramatically different %prediction of non-Gaussianity. ",
        "introduction": "In the standard hot big bang cosmology, there are several tough problems, including the flatness, isotropy and homogeneity, horizon  and topological defects problems. The hot big bang theory is unable to answer these problem in a natural way. Inflation, as an add-on, is remarkably successful in solving these problems. It gives a natural initial condition for our observed universe~\\cite{Guth81,Linde82,Steinhardt82}. Furthermore the quantum fluctuations during inflation seed wrinkles in the Cosmic Microwave Background (CMB), and today's large scale structure~\\cite{Mukhanov81,Guth82,Hawking82,Starobinsky82,Bardeen83}. As a result, inflation predicts a nearly scale invariant Gaussian CMB spectrum, which has been confirmed very well in the experiments\\cite{wmap}.  However, inflation as a successful scenario in the very early universe has its own difficulties. One of the problems with inflation is that there are too many inflationary models, which cannot be distinguished by the scalar power spectrum and power spectrum index from the CMB observation. It is essential to find more powerful signatures which could distinguish various models from each other. Moreover, inflation has some conceptual problems, cosmological singularity problem and Trans-Planckian physics being two of them. In a sense, inflation scenario is not really a fundamental theory. The trouble with it mainly comes from our ignorance of the physics of very early universe, which should be governed by a quantum gravity theory.  The rapid development of precise experiment cosmology  opens new windows to very early universe. The scalar spectral index and its running, the gravitational wave, non-Gaussianity and the isocurvature perturbation in the CMB~\\cite{wmap} are among the important probes. These probes will not only constrain a large amount of inflation models and  make the paradigm more clear, but also shed light on various other issues beyond usual inflation scenario. These issues include the initial conditions of inflation model, Trans-Planckian physics and alternative models to inflation, etc..  Among the probes, non-Gaussianity is one of the most important ones. It contains much information: magnitude, shape, sign, and even running. In principle, it could distinguish various inflation models. The deviation from the Gaussian distribution of the CMB in the WMAP observation is parameterized by $\\fnl$~\\cite{fnl}, \\be \\zeta= \\zeta_g + \\frac{3}{5}\\fnl \\zeta_g^2 \\ , \\ee where $\\zeta$ is the curvature perturbation in the uniform density slices, and the subscript $g$ denotes the Gaussian distribution. In the WMAP5 year data~\\cite{wmap},  two kinds of non-Gaussianity, local form and equilateral form, have been analysed \\be -9 < \\fnl^{local} < 111 (95\\%\\ CL) \\ , \\quad -151 < \\fnl^{equil} < 253 (95 \\% \\ CL) \\ . \\ee The central value of the local form non-Gaussianity is $51$. If the value of the local form non-Gaussianity is confirmed by the future experiments, such as Planck satellite, then it will be a great challenge to many inflation models, including the most well-studied single field inflation models. %because the local form non-Gaussianity is the order of the slow roll %parameter $\\fnl \\sim {\\cal O}(\\epsilon, \\eta)$~\\cite{NG}. Even in %the model of non-canonical action, such as DBI~\\cite{DBI} and %K-inflation~\\cite{K}, the sign of the equilateral form %non-Gaussianity is usually negative and the local form %non-Gaussianity is a small quantity~\\cite{Chen:2006nt,NG2}. Curvaton %mechanism~\\cite{Lyth:2001nq} and non-inflation ekyrotic %models~\\cite{ekpyrotic} could have large local form non-Gaussianity.  In fact, the non-Gaussianity in CMB spectrum may come from various sources during the evolution. The temperature fluctuation $\\frac{\\Delta T}{T}$ is the observable in CMB observation. During the inflation, the quantum fluctuation of the inflaton $\\delta \\phi$ are generated, and the modes of the fluctuation grow with the exponentially expanding universe. After the fluctuations leave the horizon, the decoherence effect makes the quantum fluctuations to be the classical ones. In the large scale, the physical freedom of scalar perturbation is curvature perturbation $\\zeta$. If we follow a mode, \\be {\\rm Initial \\ Condition \\above 0pt (Vacuum) } \\rightarrow \\delta\\phi \\rightarrow \\zeta \\rightarrow \\frac{\\Delta T}{ T} \\ .\\ee All the transformations are linear at first order, thus the temperature fluctuations are Gaussian. Meanwhile, any deviation from linearity in these transformations and the changes in the initial condition will influence the final observable. \\begin{itemize}   \\item  Let us first consider the last stage of the transformation $\\zeta \\rightarrow \\frac{\\Delta T}{ T}$. The fluctuations in the gravitational potential on the last scattering surface result in temperature fluctuations in the CMB, which is known as Sachs-Wolfe effect. The nonlinear Sachs-Wolfe effect generate $\\fnl$ of order one.  \\item  The curvature perturbation $\\zeta$ is conserved in the single field inflation, while in the multiple field case, the entropy perturbation change the evolution of $\\zeta$. It will suppress the perturbation conversion factor for $\\zeta$ in this process. That is why in the curvaton mechanism~\\cite{Lyth:2001nq} and new ekyrotic models~\\cite{ekpyrotic} the large local form non-Gaussianity is possible. (The other important reason for ekyrotic models generating large non-Gaussianity is that the slow-roll condition breaks down.)  \\item The primordial non-Gaussianity, which resides on the quantum fluctuation of the scalar field~\\cite{NG}, can be from the microphysics deep in the horizon. Since the observation requires the potential of the inflaton to be slow-roll,  the interaction of the inflaton is weak, and non-Gaussianity is only the order of the slow roll parameter $\\fnl \\sim {\\cal O}(\\epsilon, \\eta)$. The picture will change when the modified gravity and non-canonical action are considered, such as in ghost inflation~\\cite{ghost}, DBI inflation~\\cite{DBI}, and k-inflation~\\cite{K}. The non-linearity in the action can produce large equilateral form non-Gaussianity in the CMB. On the other hand, the back reaction argument \\cite{NG} explains why microphysics in the horizon cannot have large local form non-Gaussianity.  \\item The initial condition could be another important source of non-Gaussianity. One attempt is to consider the thermal vacuum in the inflationary cosmology~\\cite{thermal}, in which the equilateral and local form non-Gaussianity are both $\\gtrsim \\mathcal O (1)$ in some cases. In this paper, we will consider one parameter family of vacuum states, called $\\alpha$-vacuum, in de Sitter spacetime~\\cite{A,dscft} and quasi-de Sitter spacetime in inflation. In these vacua, there are correlations between points in the acausal two patches of de Sitter spacetime. We will show that $\\alpha$-vacuum can induce large local form non-Gaussianity. \\end{itemize}  As we know,  the standard treatment in scalar driven inflation scenario is based on semi-classical gravity, in which the background is described by classical Einstein Gravity and the perturbations are quantized in the background. During the slow roll period, since the evolution of the background is in a quasi-de-Sitter phase, the perturbations could be treated as the quantum field in a de-Sitter spacetime. In the expanding background of inflation, the quantum modes are stretching across the horizon. The modes observed in the CMB could be in the trans-Planckian region at the very early time. If the inflation lasts about 70 e-foldings, the perturbation which we observe in our Horizon, at that time, is deep inside the horizon. And the wavelength of these perturbation is smaller than the Planck scale. The semi-classical description is not applicable for the perturbation. It is necessary to consider stringy effect or some other quantum gravity effects on inflation. %One way is to define another vacuum. The trans-Planckian physics in inflation was first raised in \\cite{RHBrev}. As in black hole physics, an efficient way to count the Trans-Planckian effect is to modify the dispersion relations. Various dispersion relations and their physical implications have been widely studied\\cite{MB1}. %(Although they could get the large %correction of the correlation,) the parameters need be fine-tuned. Another way to discuss the trans-Planckian physics is based on the space-time uncertainty from quantum gravity, such as string theory. This noncommutative effect will impact the power spectrum and gravitational waves of CMB~\\cite{kempf,HoB} and also modify the non-Gaussianity\\cite{xue2}.  %Meanwhile, with the %symmetry of the spacetime SO(1,4), the vacuum is not the only one.   There is another attempt to  address the trans-Planckian physics in inflation, firstly suggested by Danielsson~\\cite{Dan1}. The new important ingredient in the discussion is the introduction of one parameter family of $\\alpha$-vacuum state  in the inflationary background. In de Sitter spacetime, the Bunch-Davies vacuum is the standard vacuum which is invariant under de Sitter space isometry group. However, due to the absence of globally defined time-like Killing vector, the vacuum in de-Sitter spacetime cannot be defined uniquely. Similarly, the choice of the vacuum in the inflationary background is subtle and may induce observable signature on CMB data. It was argued that the effective field theory and semi-classical gravity are applicable from the length of  new physical scale cutoff to the large scale of the universe. And it was also assumed that the modes were generated one by one at the Planck scale or new physics scale $\\Lambda$ such that the initial conditions are imposed at a mode-dependent time instead of in the infinite past. This is the motivation to introduction of $\\alpha$-vacuum in inflation. Its physical implications on inflation have been intensely studied. %The correction to the power %spectrum in the CMB is $\\sim \\mathcal O (\\frac{H}{\\Lambda})$. The order of the correction to the power spectrum has been discussed in~\\cite{Lowe1,Alberghi,Dan3,Kaloper2,Easther3,npc,Kaloper1,BM4,BM5,Dan4}. For example, in~\\cite{Kaloper1} using the method of  effective field theory the authors found that the correction was $\\sim \\mathcal O (\\frac{H^2}{\\Lambda^2})$, and in~\\cite{npc} the authors calculated the correction of power spectrum when the modes are initially created by adiabatic vacuum state, and found that the correction was $\\sim \\mathcal O (\\frac{H^3}{\\Lambda^3})$. The careful analysis of these different corrections can be seen in~\\cite{BM5}. For the non-Gaussianity from the trans-Planckian physics it was first roughly analyzed in~\\cite{Chen:2006nt}, and in~\\cite{Holman:2007na} its folded form was analyzed in the effective field theory. In this paper we follow the treatments in \\cite{Dan1,Easther3}. In \\cite{Easther3} the authors evaluated the $\\alpha$-vacuum effect in general single field inflation background, and found an oscillating dependence on the wavenumber $k$ for the power spectrum. The reason is that during inflation Hubble scale is not constant, and the coefficient for the $\\alpha$-vacuum state sensitively relies on $k$. We will show that this $k$-dependence lead to a distinguished feature of non-Gaussianity.    This paper is organized as follows: in Sec.~\\ref{chapter:alpha}, we first discuss the vacuum states in de Sitter space, the relation between Euclidean vacuum and $\\alpha$-vacuum. And then we introduce different correlation functions, and explain the property of the antipodal correlation in de Sitter space. In Sec.~\\ref{chapter:ADM} and Sec.~\\ref{chapter:three point}, we review the lagrangian formalism to compute the power spectrum and three point correlation of the curvature perturbation. Sec.~\\ref{chapter:shape} is the main result of this paper. We evaluate the local form and equilateral form non-Gaussianity for Euclidean vacuum and $\\alpha$-vacuum  in both de Sitter spacetime and inflationary backgrounds. We also draw the shapes of non-Gaussianity in each cases. Finally, we conclude in Sec.~\\ref{chapter:discussion}. ",
        "conclusions": "\\label{chapter:discussion}  %The development of cosmology and the understanding of the very %early universe need more and more precise data from cosmological %experiments. The observation of primordial gravitational wave and %non-Gaussianity are very useful probe to pin down many models of %inflation and other non-inflation models. The non-Gaussianity %contains many information, such as shape, sign, and even running. %If some certain magnitude of non-Gaussianity, such as the local %form non-Gaussianity, is confirmed in near future experiments, it %will challenge the most single field model of inflation. In this paper, we studied the physical implication of $\\alpha$-vacuum on the CMB non-Gaussianity. We found that the $\\alpha$-vacuum may lead to  large local form non-Gaussianity, and its signature %the phenomenon of non-Gaussianity in $\\alpha$-vacuum is distinct. In de-Sitter spacetime, the local form is large which is proportional to $k_2/k_3$, and in the inflationary background, the local form could still be large but there is oscillating character for $\\fnl$. Another distinctive feature of $\\alpha$-vacuum is that it leads to a divergent folded form non-Gaussianity . To illustrate the picture more clearly, we drew the shapes for different vacua and different backgrounds. We found that all of these figures have dramatically different feature. For the de Sitter spacetime the local form is large, and for the inflationary spacetime the local form not only large, but also oscillating. These are different from standard slow-roll inflation, where the local form is proportional to spectral index of scalar perturbation.  The observational feature from our study is different from \\cite{Chen:2006nt} which only considered the $\\alpha$-vacuum correction from the de Sitter background. %And in %\\cite{Holman:2007na} the authors considered the irrelevant operators %of the effective theory. For the perturbation modes blue-shifting in %the inflationary horizon, the irrelevant operators give dominant %effect as the trans-Planckian physics. The correction to the %non-Gaussianity origins from the stronger interaction at the %beginning of the inflation. %%However, we explain the trans-Planckian %physics using $\\alpha$-vacuum and the method of %Danielsson~\\cite{Dan1}. The large non-Gaussiantiy origins from %correlation from the different patches of de Sitter space. The local %and folded form non-Gaussianity are large and have distinct feature. % which dealt with the excited states in the %inflationary background and is diluted by the expansion. %It is different from our treatment, in which the $\\alpha$-vacuum %play the central role, even though the final trans-Plackian effects %on CMB spectrum is somehow similar. For example, in both cases, the %folded form non-Gaussianity is significant. %Thus it %deserves further theoretical and experimental analysis to understand %the issue. And even though our study has some overlaps with the one in \\cite{Holman:2007na}, our method and main results are different. In \\cite{Holman:2007na}, they used the effective field theory (EFT) method during inflation. Since the energy scale of inflation is very high, from the point of view of EFT, some higher order derivative terms will be not negligible, and the correction to the non-Gaussianity origins from the stronger interaction at the beginning of the inflation. Comparing with trans-Planckian physics from $\\alpha$-vacua, EFT method could be another branch in inflationary cosmology. Our treatment focused on the influence of the fluctuation vacuum, while EFT emphasized the action. From the view of experiments, due to different motivation, the final prediction of observation is not the same. \\cite{Holman:2007na} analyzed the folded form non-Gaussianities, which has  not been analyzed in WMAP5 yet. In our study, we concluded that the trans-Planckian effect would enhance the local form non-Gaussianities, which is the essential part of the WMAP data analysis.   There are various interesting issues to address on the implication of $\\alpha$-vacuum on inflation.  First of all, the loop effect in the $\\alpha$-vacuum is still an open question. Steven Weinberg has given a wonderful discussion about the loop effect in inflationary correlators~\\cite{loop1,loop2,Cogollo:2008bi}. In $\\alpha$-vacuum the problem is focused on how to renormalize the scalar perturbation in the loop diagram. In some papers, it has been argued that the $\\alpha$-vacuum is not well defined in de Sitter space due to the divergence~\\cite{Banks,Einhorn}. However in \\cite{Lowe2} using the Schwinger-Keldysh formalism a consistent renormalization method has been constructed to deal with the divergence.  The other discussion on this issue can be found in \\cite{collins}. It would be interesting to understand the loop effect and its physical implications in $\\alpha$-vacuum. %Phenomenal effect from the loop %contribution deepens the understanding for the quantum gravity in %(quasi)-de Sitter space.  Secondly, the non-Gaussianity in $\\alpha$-vacuum is sensitively dependent of the initial condition. In the case of single field inflation with higher derivative terms, the sound speed $c_s$ is not $1$. For example in the DBI inflation~\\cite{DBI} and K-inflation~\\cite{K}, the lagrangian is not canonical, and the sound speed $c_s \\ll 1$ in some situations. The sound horizon $c_s H^{-1}$ may be smaller than the length scale of new physics. In this case, the initial condition at the new physics scale is chosen in the place larger than sound Hubble scale. From the calculation of the non-Gaussianity, we know that in the path integral formalism the near Horizon crossing region impacts the results of non-Gaussianity. The divergence of the non-Gaussianity in the $\\alpha$-vacuum may not exist as pointed out in the paper~\\cite{Chen:2006nt}. This situation also appear in some special trans-Planckian physics, such as noncommutative inflation~\\cite{HoB}, in the IR region the effective string scale is smaller than the Hubble scale. %Meanwhile, %the effective field theory method dealing with the trans-Planckian %effect~\\cite{Holman:2007na} can give large but not divergent folded %form non-Gaussianity. %And also in the case %of the new physics scale of inflation is changing with the %background~\\cite{xue1}, since the noncommutative scale is the B %field in the D-brane, the non-Gaussianity will give different the %information about the new physics scale in (quasi)-de Sitter space.  Thirdly, the different initial condition of $\\alpha$-vacuum will dramatically change the correction of power spectrum~\\cite{npc,Kaloper1}, thus it predicts different non-Gaussianity. Meanwhile, the trans-Planckian dispersion relation and noncommutative geometry will also give different prediction for non-Gaussianity. %In~\\cite{Holman:2007na}, the authors study the %non-Gaussianity from trans-Plankian physics using effective field %theory.  Fourthly, it is also interesting to consider the physical implication of the $\\alpha$-vacuum in other inflationary models. One class of them is the multiple field inflation~\\cite{Nflation}. It has some very different signatures from the single field inflation: it has non-negligible gravitational wave and large local form non-Gaussianity~\\cite{LindeMuk,Bernardeau,Enqvist:2004ey,Gordon:2000hv}. Another class of inflation models is inspired by string theory. In particular, DBI inflation is a very remarkable scenario.  It may give large equilateral non-Gaussianity but no local form non-Gaussianity. It is worthwhile to discuss the $\\alpha$-vacuum effect in these inflationary models.  Finally, it could be expected that  the tri-spectrum of CMB from $\\alpha$-vacuum is different from the one in Euclidean vacuum. A more careful investigation would be valuable."
    },
    "0806/0806.4459_arXiv.txt": {
        "abstract": "Using the cut-and-paste procedure, we construct static and dynamic plane symmetric wormholes by surgically grafting together two spacetimes of plane symmetric vacuum solutions with a negative cosmological constant. These plane symmetric wormholes can be interpreted as domain walls connecting different universes, having planar topology, and upon compactification of one or two coordinates, cylindrical topology or toroidal topology, respectively. A stability analysis is carried out for the dynamic case by taking into account specific equations of state, and a linearized stability analysis around static solutions is also explored. It is found that thin shell wormholes made of a dark energy fluid or of a cosmological constant fluid are stable, while thin shell wormholes made of phantom energy are unstable. ",
        "introduction": "There has been much interest in traversable wormholes since the important Morris-Thorne article \\cite{Morris}. It was found that these geometries, which act as tunnels from one region of spacetime to another, possess a peculiar property, namely exotic matter, involving a stress-energy tensor that violates the null energy condition \\cite{Morris,MTY,Visser}, see also \\cite{LLQ} for a review and references therein. In fact, traversable wormholes violate all of the pointwise energy conditions and averaged energy conditions (see, e.g., \\cite{Lobo:2007zb}). As the violation of the energy conditions is a particularly problematic issue \\cite{Barcelo}, it is useful to minimize the usage of exotic matter. Recently, Visser  et al \\cite{Kar1,Kar2}, noting the fact that the energy conditions do not actually quantify the ``total amount'' of energy condition violating matter, developed a suitable measure for quantifying this notion by introducing a ``volume integral quantifier''. Although the null energy and averaged null energy conditions are always violated for wormhole spacetimes, they considered specific examples of spacetime geometries containing wormholes that are supported by arbitrarily small quantities of averaged null energy condition violating matter.  Another elegant way of minimizing the usage of exotic matter is to construct a simple class of wormhole solutions using the cut and paste technique \\cite{visser1,visser2} (see also \\cite{Visser}), in which the exotic matter is concentrated at the wormhole throat, i.e., a thin shell wormhole solution. The surface stresses of the exotic matter were determined by invoking the Darmois-Israel formalism \\cite{Israel}. These thin-shell wormholes are extremely useful as one may apply a stability analysis for the dynamical cases, either by choosing specific surface equations of state \\cite{eqstate1}, or by considering a linearized stability analysis around a static solution \\cite{Brady,Poisson,Ishak,Eiroa,Lobolinear,Lobolinear2}, in which a parametrization of the stability of equilibrium is defined, so that one does not have to specify a surface equation of state. The dynamical analysis in \\cite{Brady,Poisson,Ishak} was generalized by considering solutions with electrical charge \\cite{Eiroa}, solutions in the presence of a cosmological constant \\cite{Lobolinear}, and spherically symmetric dynamical solutions \\cite{Lobolinear2}. More recently, by using the cut and paste technique, several solutions were analyzed, such as, the dynamics of non rotating cylindrical thin-shell wormholes \\cite{Eiroa:2004at}, charged thin-shell Lorentzian wormholes in dilaton gravity \\cite{Eiroa:2005pc}, five dimensional thin-shell wormholes in Einstein-Maxwell theory with a Gauss-Bonnet term \\cite{Thibeault:2005ha}, solutions in higher dimensional Einstein-Maxwell theory \\cite{Rahaman:2006vg}, thin-shell wormholes associated with global cosmic strings \\cite{Bejarano:2006uj}, solutions in heterotic string theory \\cite{Rahaman:2006xb}, spherical thin-shell wormholes supported by a Chaplygin gas \\cite{Eiroa:2007qz}, a new class of thin-shell wormhole by surgically grafting two black hole solutions localized on a three brane in five dimensional gravity in the Randall-Sundrum scenario \\cite{Rahaman:2007bf}, and spherically symmetric thin-shell wormholes in a string cloud background in (3+1)-dimensional spacetime were also analyzed \\cite{Richarte:2007bx}. Other wormhole solutions have been analyzed in \\cite{Lobo:2006ue,Sushkov,Lobo-phantom,Eiroa:2008ky}.   One can go beyond thin shell solutions, and construct wormhole solutions by matching an interior wormhole to an exterior vacuum solution, at a junction surface. In particular, a thin shell around a traversable wormhole, with generic surface stresses  was analyzed in \\cite{Lobo-CQG}, a particular case of this, a wormhole with a zero surface energy density and constant redshift function, having been studied in \\cite{LLQ}. A similar analysis for the plane symmetric case, with a negative cosmological constant, was done in \\cite{LL}. The plane symmetric traversable wormhole is a natural extension of topological black hole solutions with a negative cosmological constant \\cite{lemos,zanchin,lemosreview,Lemos4}, upon addition of exotic matter. These plane symmetric wormholes can be interpreted as domain walls connecting different universes, having planar topology, and upon compactification of one or two coordinates, cylindrical topology or toroidal topology, respectively. The construction of these wormholes does not alter the topology of the background spacetime (i.e., spacetime is not multiply-connected), so that these solutions can instead be considered domain walls. Thus, in general, these wormhole solutions do not allow time travel.  From recent astronomical observations, it seems that we presently live in a world with a positive cosmological constant, $\\Lambda>0$. However, a spacetime with $\\Lambda<0$ is also of significant relevance, since it allows a consistent physical interpretation in the context of supergravity and superstring theories \\cite{lemosreview}. Indeed, one can enlarge general relativity into a gauged extended supergravity theory, in which the vacuum state has an energy density given by $\\Lambda=-3g^2/(4\\pi pG)$, where $g$ is the coupling constant of the theory, and $G$ is the gravitational constant. Thus, the vacuum of the theory is described by an anti-de Sitter spacetime. It is important to emphasize that if these theories are correct, they imply that the anti-de Sitter spacetime should be considered as a symmetric phase of the theory, although it must have been broken, since we do not presently live in a universe with $\\Lambda<0$. In addition, note that anti-de Sitter spacetimes have other interesting features: (i) they are one of the rare gravitational backgrounds yielding a consistent interaction with massless higher spins, (ii) they permit a consistent theory of strings in any dimension, (iii) they allow valid definitions for the mass, angular momentum and charge, and (iv) it has been conjectured that they has a direct correspondence with conformal field theory on the boundary of that space, the AdS-CFT conjecture, see, e.g., \\cite{Aharony:1999ti}. Moreover, even with its preference to negative cosmological constant scenarios, string theory can although in a contrived way, produce a landscape of positive cosmological constant universes \\cite{kklt}, indicating perhaps that one can transit bewteen both signs of the cosmological constant. So, it is certainly of interest to study astrophysical structures that appear within a positive cosmological constant scenario, as well as to study micro structures that may appear within a negative cosmological constant scenario. Here we analyze wormhole structures within a negative cosmological context. Within this context other structures, such as black holes with spherical, toroidal, and hyperbolical topologies, have been analyzed (see \\cite{lemosreview} for a review).   The paper is organized as follows. By using the cut-and-paste procedure for a thin shell, we construct static and dynamic wormholes by surgically grafting together two spacetimes of plane symmetric vacuum solutions, with planar, cylindrical or toroidal topologies, with a negative cosmological constant. This analysis is displayed in Section \\ref{Sec:II}. We shall also consider the specific case of static wormhole solutions, and consider several equations of state. In Section \\ref{Sec:III}, we consider a dynamical stability analysis. In particular, a stability analysis is carried out for the dynamic case by taking into account specific surface equations of state, and a linearized stability analysis around static solutions is also further explored. Finally, in Section \\ref{Conclusion}, we conclude. ",
        "conclusions": "Using the cut-and-paste technique, we have constructed an elegant and simple class of static and dynamic plane symmetric wormholes by surgically grafting together two spacetimes of plane symmetric vacuum solutions with a negative cosmological constant. We have further considered a dynamical stability analysis, first, by considering specific equations of state, and paid special attention to the dark energy equation of state. Second, a linearized stability analysis was also explored. Considering radial perturbations around a static solution, the respective stability regions were presented, and the specific case of a dark energy equation of state was also analyzed. It was found that thin shell wormholes made of a dark energy fluid or of a cosmological constant fluid are stable, while thin shell wormholes made of phantom energy are unstable. These plane symmetric wormholes may be viewed as domain walls connecting different universes. The construction of these wormholes does not alter the topology of the background spacetime, i.e., spacetime is not multiply-connected, so that, in general, these wormhole solutions do not allow time travel.      \\bigskip"
    },
    "0806/0806.2755_arXiv.txt": {
        "abstract": "\\noindent The problem of corrections to Einstein's equations arising from averaging of inhomogeneities (``backreaction'') in the cosmological context, has gained considerable attention recently. We present results of analysing cosmological perturbation theory in the framework of Zalaletdinov's fully covariant Macroscopic Gravity. We show that this framework can be adapted to the setting of cosmological perturbations in a manner which is free from gauge related ambiguities. We derive expressions for the backreaction which can be readily applied in \\emph{any} situation (not necessarily restricted to the linear perturbations considered here) where the \\emph{metric} can be brought to the perturbed FLRW form. In particular these expressions can be employed in toy models studying nonlinear structure formation, and possibly also in $N$-body simulations. Additionally, we present results of example calculations which show that the backreaction remains negligible well into the matter dominated era. ",
        "introduction": "\\label{sec:intro} \\noindent It is known \\cite{ellis} that in order to apply the equations of Einstein's General Relativity at the large length scales of interest in cosmology, it is necessary to first perform a smoothing or averaging operation, which will generate nontrivial corrections in these nonlinear equations. There is a debate in the literature concerning two basic questions regarding this smoothing operation : (a) How does one obtain consistent and observationally relevant variables associated with the corrections due to averaging or ``backreaction'', and (b) what is the magnitude of this backreaction?  This discussion has had a long history \\cite{avg,siegel,avg-bran}, although in recent times attention has been focussed on two promising candidates for a consistent \\emph{nonperturbative} averaging procedure, namely the spatial averaging of scalars due to Buchert \\cite{buchert} and the covariant approach due to Zalaletdinov \\cite{zala1,zala2}. Perhaps owing to its appealing simplicity of implementation, the approach due to Buchert has attracted a significant amount of attention both in the context of cosmological perturbation theory \\cite{pert-buchertavg,kolb,behrend} and with regards fully nonlinear calculations \\cite{avg-buchert,rasanen1,rasanen2}, although very interesting results have been derived using Zalaletdinov's covariant procedure as well \\cite{coley,vdhoogen}. The most fascinating aspect of these studies has been the possibility that the phenomenon of Dark Energy \\cite{kolb,rasanen2} and also Dark Matter \\cite{vdhoogen} might be attributed to the backreaction from averaging (see \\Cite{buchert-rev} for a recent review of the subject).  The physical relevance of the averaged variables defined in Buchert's formalism, has been questioned in the literature \\cite{wald}. It has been argued \\cite{wald,smalleffect} that effects of inhomogeneities which only perturbatively affect the (on average homogeneous) metric of the universe, \\emph{cannot} lead to effects large enough to account for the inferred acceleration of the universe from, say supernovae observations \\cite{SNe}. This has been countered by the argument \\cite{rasanen2} that perturbation theory may break down during the epoch of fully nonlinear structure formation. Recently it was shown \\cite{SphColl} in the context of a specific model of spherical collapse, that an explicit coordinate transformation could be found, which brought the metric to the perturbed Friedmann-Lema\\^itre-Robertson-Walker (FLRW) form, \\be ds^2 = -(1+2\\vphi)d\\tau^2 + a^2(\\tau)(1-2\\psi)d\\vec{x}^2\\,, \\label{eq1} \\ee satisfying all conditions required for a perturbation formalism in $\\vphi$ and $\\psi$ to hold, \\emph{even} in the regime of fully nonlinear collapse.  In such a context, it is important to have a consistent formalism, free from issues such as gauge artifacts, in which one has derived expressions for the backreaction which can be applied in a straightforward manner to any model in which the metric can be brought to the form \\eqref{eq1}. Since the Buchert framework by construction is best adapted to coordinates comoving with the matter, most applications using perturbation theory in this framework have focussed on the synchronous and comoving gauge, although recently Behrend et al. \\cite{behrend} have also performed calculations in the conformal Newtonian gauge. However, as pointed out elsewhere \\cite{SphColl}, calculations using the Buchert framework in perturbation theory necessarily face the ambiguity of dealing with \\emph{two} scale factors : one is the scale factor with which the perturbed FLRW metric is defined, and the other is the volume averaged scale factor defined by Buchert \\cite{buchert}, and it is not clear which of these scale factors is the observationally relevant one.  In this paper we shall deal with the covariant framework of Zalaletdinov's Macroscopic Gravity (MG), and its spatial averaging limit proposed in \\cite{spatavglim}. We will see that here, one has a well defined averaged metric with a corresponding \\emph{uniquely defined} scale factor. Further, we will argue that the structure of MG allows us to completely specify a consistent, observationally relevant averaging operation adapted for perturbations to the FLRW geometry. The final expressions obtained for the backreaction can be applied directly (or with some straightforward modifications), to \\emph{any} situation in which the metric can be brought to the form \\eqref{eq1} or a suitable generalisation thereof.  The organisation of this paper is as follows : \\Sec{sec:MG} presents a brief recap of the basic ideas underlying MG and some useful expressions from standard cosmological perturbation theory (PT), together with a proposal for practically estimating the effect of the backreaction on the time evolution of the FLRW scale factor. Sec. \\ref{sec:avgVP} presents details of the MG averaging procedure adapted to cosmological PT, including general expressions for the leading order backreaction terms, with a discussion of gauge related issues and the definition of the averaging operator. The heart of the paper is in \\Sec{sec:corrscal}, where we derive final expressions for the backreaction, both in real space and Fourier space, which can be directly utilised in model calculations. While these expressions use a few simplifying restrictions, these can be lifted if necessary in a completely straightforward manner. \\Sec{sec:exs} contains example calculations in first order PT, which show that the magnitude of the backreaction is, as expected, negligible compared to the homogeneous energy density of matter in the radiation dominated era and for a significant part of the matter dominated era. We conclude in \\Sec{sec:discuss} with some final comments. Throughout the paper, lower case Latin indices $a,b,c$... will refer to spacetime indices $0,1,2,3$, and upper case Latin indices $A,B,C,$... to spatial indices $1,2,3$. The speed of light $c$ is set to unity, and a prime refers to a derivative with respect to conformal time unless stated otherwise. ",
        "conclusions": "\\label{sec:discuss} \\noindent This paper has presented an analysis of cosmological perturbation theory (PT) in the fully covariant averaging framework of Zalaletdinov's Macroscopic Gravity (MG) \\cite{zala1,zala2} and its restriction to spatial averaging \\cite{spatavglim}. While this framework is generally covariant, the issue of gauge dependence in perturbation theory introduces certain subtleties in the problem. We have shown that, provided one takes seriously the idea that the cosmological background must be defined by an averaging procedure \\cite{zala-pert}, it is possible to attach a gauge invariant meaning to the averaging condition and the corresponding correlation objects which appear as corrections to the cosmological equations. While there remains considerable freedom in an explicit choice of the averaging operator (through a choice of the volume preserving gauge used in its definition), this freedom can be fixed by some additional requirements based on knowledge of cosmological PT in the standard implementations. In particular we have seen that properties of the conformal Newtonian or Poisson gauge can be used to motivate a fully specified choice of the averaging operation adapted to first order PT.  One prerequisite to the formulation of a consistent averaging framework in the context of perturbation theory, is the absence of perturbative fluctuations with arbitrarily large wavelengths, since such fluctuations would render meaningless the notion of recovering a homogeneous background on averaging. This problem also manifested itself in the correlation integrals \\eqref{corrscal-kspace}, which diverge in the presence of a finite amplitude of fluctuations as the wavenumber $k\\to0$. Accordingly, all the calculations of this paper have assumed that the \\emph{initial} power spectrum of metric fluctuations has a sharp cutoff at the scale corresponding to $k=H_0$, a hypothesis which is in fact supported by analysis of CMB data \\cite{cutoff-obsvns}.  The main purpose of this paper was to lay down the formalism of MG in a language most convenient from the point of view of cosmological PT. This was accomplished by writing \\eqns{avgVP-corrscal}, \\eqref{corrscal} and the Fourier space version \\eqref{corrscal-kspace} (with certain simplifying assumptions regarding vector and tensor perturbations which can if needed be relaxed in a completely straightforward manner). This was supplemented by calculations in the sCDM model \\cite{dodelson} (which is the flat FLRW model with radiation and Cold Dark Matter but no Dark Energy) with the additional simplification of ignoring the baryons. The analytical results of \\Sec{sec:exs:sub:eds} as well as the more detailed numerical results of \\Sec{sec:exs:sub:nobar} show that the correlation objects or ``backreaction'' remain negligibly small upto epochs corresponding to a scale factor of $a\\sim0.01$. While the calculations ignored corrections from quasi-linear and nonlinear scales, these are not expected to contribute dramatically to the correlations obtained here, an expectation which is justified by the work of Behrend et al. \\cite{behrend} and further by the calculation in \\Cite{NLcoll-MG} (see below).  We have seen that by using the framework of MG, we have completely bypassed the problem mentioned in the Introduction, which one faces when applying the Buchert framework to cosmological PT, namely of having to deal with \\emph{two} scale factors. Here, one has a single well-defined scale factor associated with the background metric, and its evolution can be obtained in an iterative fashion as described in \\Sec{sec:MG:sub:iter}. In practice, we saw that since the backreaction is small, convergence can be achieved by essentially a single calculation, at least in the context of first order PT.  This brings us to a final, and very important issue : What is the magnitude and behaviour of the backreaction in the fully nonlinear regime of structure formation? There are (at least) two possible avenues to approach this question. The first is to set up the problem in a manner which is suitable for $N$-body simulations. The iterative approach suggested earlier can presumably be adapted to full-fledged $N$-body codes as well. While this is a possibility worth pursuing, there is also a less involved (but correspondingly less realistic) way of determining the effect of nonlinearities on the backreaction, which is to study toy models of structure formation. As mentioned earlier, such a toy model of spherical collapse was recently presented in \\Cite{SphColl}, and it was shown by an explicit coordinate transformation, that the metric can be brought to the Newtonian form \\eqref{eq1}. Now, it is worth noting that while all the calculations of this paper assumed that first order PT is valid, the actual expressions for the correlation integrals in \\eqns{corrscal} \\emph{only} assume that the potentials $\\vphi$ and $\\psi$ satisfy $|\\vphi|,|\\psi|\\ll1$. The \\emph{dynamics} governing the potentials is irrelevant at this stage. This means that, as long as one is interested in leading order effects only,  the expressions in \\eqns{corrscal} \\emph{can be directly applied} to any model of structure formation where the \\emph{metric} can be brought to the conformal Newtonian form \\cite{note5-vectortensor}: in particular they can be applied to the model of \\Cite{SphColl}. This has been done in \\Cite{NLcoll-MG}, and one finds that even in the fully nonlinear regime, the effect of the backreaction remains negligible. In this context see also \\Cite{siegel} for a third approach.  To conclude, all our calculations and arguments seem to indicate that the averaging of perturbative inhomogeneities in a consistent manner appears to lead only to very small effects. This does not, however, mean that the effects do not exist. It remains to be seen whether any observable consequences of the backreaction may be detectable by future experiments. \\\\"
    },
    "0806/0806.3203_arXiv.txt": {
        "abstract": "{Accurate mass, radius, and abundance determinations from binaries provide important information on stellar evolution, fundamental to central fields in modern astrophysics and cosmology.} {Within the long-term Copenhagen Binary Project, we aim to obtain high-quality light curves and standard photometry for double-lined detached eclipsing binaries with late A, F, and G type main-sequence components, needed for the determination of accurate absolute dimensions and abundances, and for detailed comparisons with results from recent stellar evolutionary models.} {Between March 1985 and July 2007, we carried out photometric observations of \\AD, \\HW, \\SW, \\V636, \\VZ, and \\WZ\\ at the Str{\\\"o}mgren Automatic Telescope at ESO, La Silla.} {We obtained complete $uvby$ light curves, ephemerides, and standard $uvby\\beta$ indices for all six systems. For V636\\,Cen and HW\\,CMa, we present the first modern light curves, whereas for AD\\,Boo, SW\\,CMa, VZ\\,Hya, and WZ\\,Oph, they are both more accurate and more complete than earlier data. Due to a high orbital eccentricity ($e = 0.50$), combined with a low orbital inclination ($i = 84\\fdg7$), only one eclipse, close to periastron, occurs for HW\\,CMa. For the two other eccentric systems, V636\\,Cen ($e = 0.134$) and SW\\,CMa ($e = 0.316$), apsidal motion has been detected with periods of $5\\,270 \\pm 335$ and $14\\,900 \\pm 3\\,600$ years, respectively.} {} ",
        "introduction": "\\label{sec:intro}   Detached, double-lined eclipsing binaries are well-known as the main source of accurate (1--2\\% or even better) data on stellar masses and radii. During the last two-three decades, such results have been published for about 50 sy\\-stems within the main-sequence band (e.g. Andersen \\cite{ja91}; Ribas et al. \\cite{ir00a}; Southworth, private communication \\footnote{{\\scriptsize\\tt http://www.astro.keele.ac.uk/$\\sim$jkt/}}), but some mass intervals, especially around and below 1 $M_{\\sun}$, still need better coverage, as do parts of the main-sequence band at higher masses. A rich sample of accurate binary results is mandatory for detailed tests of stellar models in general, as well as for many specific astrophy\\-si\\-cal investigations (Ribas \\cite{ir04}, \\cite{ir06b}). Remarkable recent examples are the determination of the helium-to-metal enrichment ratio (Ribas et al. \\cite{ir00a}) and investigations of the mass dependence of convective core overshooting (Ribas et al. \\cite{ir00b}, Claret \\cite{claret07}).  In this paper, we present new accurate and complete $uvby$ light curves for six detached double-lined eclipsing binaries having late A, F, and G-type main-sequence components. In addition, new spectroscopic observations for radial velocities and abundance analyses have been obtained, and detailed analyses for all the systems will be published se\\-pa\\-ra\\-tely (Clausen et al. \\cite{avw08}, Clausen et al. in prep., Torres et al. in prep.).  When we decided to observe AD\\,Boo (F6+G0) and SW\\,CMa (A8), they were lacking modern photometric and spectroscopic observations. This was independently noticed by Lacy (\\cite{lacy97a}, \\cite{lacy97b}) who studied both sy\\-stems, and by Popper (\\cite{dmp98a}) who obtained spectroscopic orbits for AD\\,Boo, leading to masses for the components which differ somewhat from those by Lacy. For both sy\\-stems, a new analysis based on new light curve and radial velocity observations, supplemented with a spectroscopic abundance analysis, is desirable.  HW\\,CMa (A8) is only about $2 \\arcmin$ distant from SW\\,CMa and was discovered to be a double-lined spectroscopic binary du\\-ring the radial velocity observations of SW\\,CMa. It was later found to be eclip\\-sing (Liu et al. \\cite{liu92}), and we decided to obtain complete light curves and spectroscopic orbits.  V636\\,Cen (F8/G0V) was selected as part of a large-scale study of new systems with solar-type components (Clausen et al. \\cite{jvcetal01}). We want e.g. to investigate a serious dilemma that appears to be present in the compa\\-ri\\-son of predictions from current stellar models with fundamental properties of known $0.7$--$1.1$ $M_{\\sun}$ eclipsing binaries (Popper~\\cite{dmp97}, Clausen et al.~\\cite{granada99b}). Our results for V636\\,Cen support the presence of a real discrepancy - current models, adopting a fixed mixing length parameter calibrated on the Sun for envelope convection, do not predict identical ages for the two components (Larsen \\cite{al98}, Clausen et al. in prep.). A similar case, V1061\\,Cyg, was presented by Torres et al. (\\cite{wt06}). See also Ribas (\\cite{ir06a}) for eclipsing binaries with K and M type components and Morales et al. (\\cite{mrj07}) for single active late-K and M stars.  VZ\\,Hya (F5) and WZ\\,Oph (F8) are well known systems included by Popper (\\cite{dmp80}) in his critical review of stellar masses, but excluded by Andersen (\\cite{ja91}), since the available results did not fulfill his 2\\% criteria for masses and radii. We therefore decided to provide the new observations needed for the determination of more accurate dimensions.  In the following, we first give a general description of the observations, the basic photometric reduction, and the light curve formation (Sect. \\ref{sec:obs}), and we then present the results for the individual systems (Sect. \\ref{sec:adboo} -- \\ref{sec:wzoph}). Throughout the paper, the component eclipsed at the deeper eclipse at phase 0.0 is referred to as the primary, and the other as the secondary component. ",
        "conclusions": "\\label{sec:conclusions}  Complete $uvby$ light cuves and $uvby\\beta$ standard indices have been obtained for the six main-sequence eclipsing binaries \\AD, \\HW, \\SW, \\V636, \\VZ, and \\WZ. For \\V636\\ and \\HW, the light curves are the first modern ones obtained, and for the other systems they are both more complete and more accurate than earlier data. Ephemerides have been determined from all available times of minima, and apsidal motion has been detected and quantified for \\SW\\ and \\V636.  For \\AD, the new photometry indicates low-level surface activity, and spectroscopy supports that the cooler secondary component is the most likely candidate (Clausen et al. \\cite{avw08}). The new photometry and spectroscopy also suggests that \\WZ\\ is mildly active. The $uvby$ light curves of \\V636\\ clearly show variations due to activity at a moderate level, also confirmed by spectroscopy (Clausen et al. in prep.).  In parallel, new spectroscopic observations for accurate radial velocities and abundance determinations were obtained. The combined data yield accurate absolute dimensions and heavy-element abundances for all six systems and allow detailed tests of current stellar evolu\\-tio\\-na\\-ry mo\\-dels. Our analyses will be published in three separate papers (Clausen et al. \\cite{avw08}; the F-type systems \\AD, \\VZ, and \\WZ), (Clausen et al. in prep.; the G-type system \\V636); (Torres et al. in prep.; the late A-type systems \\HW\\ and \\SW)."
    },
    "0806/0806.3729_arXiv.txt": {
        "abstract": "Type Ia supernovae (SNe Ia) occur in both old, passive galaxies and active, star-forming galaxies. This fact, coupled with the strong dependence of SN Ia rate on star formation rate, suggests that SNe Ia form from stars with a wide range of ages.  Here we show that the rate of SN Ia explosions is about 1\\% of the stellar death rate, independent of star formation history. The dependence of SN Ia rate on star formation rate implies a delay time distribution proportional to $t^{-0.5\\pm0.2}$. The single degenerate channel for SNe Ia can be made to match the observed SN Ia rate -- SFR relation, but only if white dwarfs are converted to SNe Ia with uniform efficiency of $\\sim$1\\%, independent of mass. Since low-mass progenitors are expected to have lower conversion efficiencies than high mass progenitors, we conclude that some other progenitor scenario must be invoked to explain some, or perhaps all, SNe Ia. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0806/0806.1801_arXiv.txt": {
        "abstract": "{}% {Archival observations of infrared fine-structure lines of the young Galactic supernova remnant Cassiopeia~A allow us to test existing models and determine the physical parameters of various regions of the fast-moving knots, which are metal-dominated clouds of material ejected by the supernova explosion. }% {The fluxes of far-infrared [\\ion{O}{i}] and [\\ion{O}{iii}] lines are extracted from previously unpublished  archival ISO data. The archival \\emph{Spitzer} data are used to determine the fluxes of the O, Ne, Si, S, Ar, and Fe ion fine-structure lines originating in the fast-moving knots. The ratios of these line fluxes are used as plasma diagnostics. We also determine the infrared line flux ratios with respect to the optical [\\ion{O}{iii}] 5007~\\AA\\ line in the knots with previously measured reddening. Additionally, we analyze several optical and near-infrared observations of the fast-moving knots to obtain clearer insight into the post-shock photoionized region structure. }% {We show that the infrared oxygen line flux predictions of all existing theoretical models are correct only to within a factor of a several. Comparison of the model predictions shows that to reproduce the observations it is essential to include the effects of the electron conductivity and dust. Detailed analysis of the diagnostic line flux ratios of various ions allows us to qualitatively confirm the general model of fast-moving knot emission and determine observationally for the first time the physical conditions in the photoionized region after the shock. We infer from the [\\ion{O}{iii}] line flux ratios that the pre-shock cloud densities are higher than assumed in existing theoretical models and most probably correspond to several hundred particles per cm$^3$. We also determine the Cas~A luminosity in the infrared continuum and lines. We show that accounting for the charge exchange processes in the post-shock photoionized region allows us to reproduce most of the relevant spectral line ratios even in the frame of a single-temperature model of this region. We also estimate its plasma parameters, thickness, and carbon abundance. }% {} ",
        "introduction": "\\label{SecIntro}  The young Galactic supernova remnant Cassiopeia~A (Cas~A) is one of only a few objects where the inner composition of the supernova progenitor may be studied directly, as the explosion ejecta have not yet mixed with the surrounding medium. In places where these ejecta are being crossed by the reverse shock wave, they are observable as compact knots emitting optical forbidden lines of ions of oxygen, sulfur, argon, and other elements. Because of their high density, these ejecta knots are not strongly decelerated by the reverse shock and, as a result, have high proper motions that was the reason to name them as the fast-moving knots (FMKs). Studies of \\citet{Peimbert71} and \\citet{CheKirCasA,CheKirAbund} firmly established that the FMKs are an observational manifestation of the supernova ejecta.  In FMKs, we therefore deal with a very unusual situation in astrophysics, where oxygen comprises 80\\%-90\\% of all atoms in the plasma, other components being mostly Ne, Si, S, Ar, and Fe. The hydrogen and helium abundances are, in contrast, understood to be negligibly low. Optical observations have shown that some knots consist even of almost pure oxygen (so-called [\\ion{O}{iii}] filament, \\citet{CheKirAbund}).   Theoretical interpretation of the optical observations has proven that the plasma heated up to temperatures of several million degrees is rapidly cooling because of its high density and the anomalous chemical composition of the gas. At the same time, the ionization degree lags behind the rapid cooling rate. This results in a unique situation where atoms with strongly different ionization degrees coexist at low temperatures from $10^4$~K to several hundreds Kelvin. Because of this, as we now show, the lines of highly-charged ions, such as [\\ion{Si}{x}] and [\\ion{Ar}{v}], originate at temperatures of about $10^4$~K and lower% \\footnote{% The [\\ion{Si}{x}] line has also been detected in spectra of several active galactic nuclei \\citep{AGNSiX,AGNNIRAtlas}, where it is apparently formed in the photoionized region. Even so, the models predict line-forming region temperatures of around $10^5$~K \\citep{SyFerland97}. }.   Traditionally, theoretical models of fast-moving knot emission were compared only with optical line observations. Since there is only a very limited number of diagnostic line ratios in the optical spectrum, the FMK theoretical models are difficult to both construct and compare with optical observations. The most well-known models are those constructed by \\citet{Itoh81a,Itoh81b}, \\citet{BorkowskiO}, and \\citet{SD95}.  \\cite{CasAORL} illustrated one way of constraining the theoretical models using prospective observations of optical and near-infrared recombination lines of highly-charged oxygen ions.  In this paper, we demonstrate how existing archival infrared observations can discriminate between existing models and place strong constraints on the construction of future ones. We compare the model predictions of the ratios of far-infrared (FIR) fine-structure oxygen line fluxes to the flux in the optical [\\ion{O}{iii}] 5007~\\AA\\ line with their observational values, and show that all of the predictions are incorrect by a factor of several for one or several flux ratios. It should be remembered, however, that the theoretical models predict the deviations of ionic abundances from their collisional equilibrium values by several orders of magnitude. Therefore, inconsistencies of the observed order suggest that some corrections should be applied to the existing models, while the general picture described in the models is entirely correct.  The infrared fine-structure line flux ratios of other element (in addition to oxygen) ions are currently the most reliable tool available for direct studies of the physical parameters in the FMKs. We attempted to estimate these parameters using available observational data. As a result, parameters of the post-shock photoionized region were estimated observationally for the first time, confirming qualitatively the predictions of the theoretical models, but again highlighting some quantitative differences.  The FIR lines will provide even more important probe in the future, as several far-infrared observatories, such as the Herschel Space Observatory and SOFIA, begin their operation. These telescopes will provide far higher quality FIR data ofhigher sensitivity and angular resolution, allowing us to constrain the FMK models more tightly. Results of this present study are also important to understand the small-scale structure of other oxygen-rich supernova remnants, such as Puppis~A, N132D, and G292+1.8.   \\bigskip  The paper structure is the following. In the next section, we describe briefly the existing theoretical models of the FMKs and estimate oxygen infrared line fluxes in one of these models. Section~\\ref{SecIRObs} describes the archival infrared observations of Cas~A and their analysis. In Sect.~\\ref{SecAnalysis}, we estimate the physical conditions and abundances in different regions of FMKs from the observed line flux ratios. We also discuss these results and compare them with values predicted by the models. In Sect.~\\ref{SecPIR}, we analyze spectral line ratios to obtain a consistent model of the post-shock photoionized region. Section~\\ref{SecIRL} is devoted to the predictions of the mid- and far-infrared recombination line intensities. Finally, in Sect.~\\ref{SecConclusions} we present our conclusions. ",
        "conclusions": "\\label{SecConclusions}  We have performed an analysis of the supernova remnant Cas~A fast-moving knot infrared line intensities, by comparing observations with the predictions of various theoretical models, which describe the FMK emission originating from the reverse shock interaction with the pure-oxygen or oxygen-dominated clouds. For this comparison, we have analyzed archival observational data acquired by the ISO and \\emph{Spitzer} observatories. We conclude that accounting for the electron conductivity is essential to reproducing the observed line ratios. This is the reason why the \\citet{BorkowskiO} model BS-DC provides the most accurate description of the observed oxygen line relative fluxes, although it is only precise up to a factor of several (the [\\ion{O}{i}] lines are overpredicted partially because of the too low model pre-shock density).  This emphasizes the need for more sophisticated models describing shock propagation in the oxygen-dominated plasma. In this article, many line flux ratios are derived that will help us to construct these future theoretical models. Future far-infrared and optical observations of higher sensitivity and angular resolution will allow one to obtain far more information about the fast-moving knots once the models are constructed.  Analysis of the infrared lines of O, Ne, Si, S, Ar, and Fe ions have confirmed the existence of three regions contributing to the infrared and optical line emission: \\begin{itemize} \\item Lines of the intermediate- and highly-charged ions that arise in the post-shock rapidly cooling region (e.g., [\\ion{Ne}{v}], [\\ion{Ar}{v}], [\\ion{Si}{vi}], [\\ion{Si}{x}] lines). \\item Singly-charged ions that radiate mostly from the post-shock photo-ionized region (e.g., [\\ion{O}{i}], [\\ion{Si}{ii}], [\\ion{Ar}{ii}], [\\ion{Fe}{ii}] lines). \\item Weakly- and intermediate-charged ion lines originate in the pre-shock photoionized region as well. It gives major contribution to the emission of lines of [\\ion{Ne}{ii}], [\\ion{O}{iii}], [\\ion{S}{iii}], [\\ion{Ar}{v}], etc. \\end{itemize}  An analysis of the FIR observations jointly with the \\citet{HF96} and \\citet{CasANIR01} optical and near-infrared data has allowed us to construct a consistent model of the post-shock photoionized region. This became possible by accounting for the charge exchange process that diminishes oxygen ionization in the post-shock PIR. Observed high ($>0.9$) ionization of sulfur and low ($\\approx0.06$) ionization of oxygen confirms the importance of the charge exchange reactions in this region.  From the observed line flux ratios, it then follows that the post-shock PIR has a temperature $T_{\\rm e}\\approx (1500-2000)$~K, electron and ion densities of $n_{\\rm e}\\approx 5\\times10^5$~cm$^{-3}$ and $n_{\\rm t}\\approx 1\\times10^6$~cm$^{-3}$ and is approximately $(1-2)\\times10^{11}$~cm thick, assuming that oxygen constitutes about 70\\% of all atoms in the FMK plasma. Given the uncertainties of the model, this result is consistent with the thickness predicted in the \\citet{BorkowskiO} model.  We have constrained the helium abundance to be $n({\\rm He})/n({\\rm O}) < (0.3-0.6)$ and the carbon abundance to be about $n({\\rm C})/n({\\rm O}) \\approx 0.5$.  However, we recall that the obtained values of elemental abundances and ionization fractions in the post-shock PIR are affected by unknown biases, as our assumptions (plane-shock, single-temperature model) clearly do not adequately represent the structure of this region.  From the comparison of observed infrared line intensities with the theoretically expected values, we conclude that \\begin{itemize} \\item The theoretical models correctly describe the general structure of the FMKs; \\item Accounting for the electron conductivity brings the BS-DC model into far closer agreement with the observed oxygen FIR line ratios to the [\\ion{O}{iii}] 5007~\\AA\\ line than the BS-F and SD-200 models; \\item Accounting for the emission from the photoionized region before the shock front is essential to reproducing the FIR [\\ion{O}{iii}] line ratio; \\item The FIR [\\ion{O}{i}] line intensities in the \\citet{Itoh81b} and \\citet{BorkowskiO} models are partially overestimated because of the too low pre-shock atom number densities; \\item The pre-shock atom number density is at least 100~cm$^{-3}$ and may be as high as 300~cm$^{-3}$; \\item The high density of the post-shock PIR suggests that there is no significant magnetic-field pressure in this region; \\item The post-shock photoionized region truncation, as suggested by \\citet{Itoh86}, may be needed to explain the weakness of several lines produced by the final plasma recombination. \\end{itemize}  The infrared metal recombination lines are shown to be detectable by the planned far-infrared instruments and useful for the derivation of the plasma properties.   The high brightness of the [\\ion{S}{ii}] lines close to 10\\,300~\\AA\\ remains unexplained by our single-temperature model of the post-shock PIR."
    },
    "0806/0806.4086_arXiv.txt": {
        "abstract": "A new luminosity function for galaxies can be built starting from the product of two random variables X and Y represented by a gamma variate with argument 2 . The mean , the standard deviation and the distribution function  of this new distribution are computed. This new probability density function is assumed to  describe the mass distribution of galaxies. Through   a non linear rule of conversion from mass to luminosity a second new luminosity function  for galaxies is derived. The test  of reliability of these two luminosity functions was made on the Sloan Digital Sky  Survey (SDSS) in five different bands. The  Schechter function gives a  better fit with  respect to the two  new luminosity functions for galaxies here derived. ",
        "introduction": "Given  two independent non-negative random variables X and Y, their  product XY  represents an  active field of research. When X  and Y are  Student's $t$ random variables the product XY is applied  in the field of finance~\\cite{Su2006} . When  X and Y are n-Rayleigh distribution , the application can be the  wireless propagation research~\\cite{Salo2006}~. In Section~\\ref{secprobability} this  paper explores the product XY when X and Y are  gamma variate with argument 2. Section~\\ref{vorogalaxies} explores  the connection between  the Voronoi Diagrams and galaxies. Section~\\ref{sectionlum} reports two new luminosity functions for galaxies as deduced from the product XY. ",
        "conclusions": "The PDF of the product of two independent random variables $X$ and $Y$ as represented by two gamma variates with argument 2 has  been analytically derived.  The mean , the variance and the DF  of this new PDF are computed. As an application we assumed that the mass of the galaxies behaves in the same way  as this new PDF. The LF  for galaxies can therefore derived assuming a linear or  nonlinear relationship between  mass and luminosity: in the first case we have  a two parameter LF and in the second case  a three parameter LF; recall that the Schechter LF for galaxies has three parameters. The parameter $a$ that characterizes the non-linear relationship between mass and luminosity of the second LF  , see equation~(\\ref{equation_mia}) , is found to be around 1.  The comparison  between  the two  LF for galaxies here derived is   performed  on the SDSS data and can be done  only  by introducing an upper limit in magnitude , $M_{lim}$,  in the   five  bands analyzed. The three  tests of reliability here adopted show  that the Schechter function  always has a smaller $\\chi^2$ , $AIC$ and   $BIC$   with respect to the two new  LF for galaxies here derived , see Table~\\ref{datamia} and  Table~\\ref{datamiamassa}.  The theoretical number of galaxies  as a  function of the red-shift presents a maximum that is a function of $\\alpha$ and $f$   for the Schechter  function; conversely, when   the  first  new LF  here derived is considered  , the maximum is a function only of  $f$, see equation~(\\ref{maximumia}). The first new LF  for galaxies , once   implemented on a 3D Voronoi slice ,   allow us  to  reproduce the large scale structures of our universe."
    },
    "0806/0806.3459_arXiv.txt": {
        "abstract": "We utilize the local velocity dispersion function (VDF) of spheroids, together with their inferred age--distributions, to predict the VDF at higher redshifts ($0<z\\lsim 6$), under the assumption that (i) most of the stars in each nearby spheroid formed in a single episode, and (ii) the velocity dispersion \\sis\\ remained nearly constant afterward. We assume further that a supermassive black hole (BH) forms concurrently with the stars, and within $\\pm$ 1 Gyr of the formation of the potential well of the spheroid, and that the relation between the mass of the BH and host velocity dispersion maintains the form $M_{\\rm BH} \\propto \\sigma^{\\beta}$ with $\\beta\\approx 4$, but with the normalization allowed to evolve with redshift as $\\propto (1+z)^{\\alpha}$. We compute the BH mass function associated with the VDF at each redshift, and compare the accumulated total BH mass density with that inferred from the integrated quasar luminosity function (LF; the so--called So\\l tan argument).  This comparison is insensitive to the assumed duty cycle or Eddington ratio of quasar activity, and we find that the match between the two BH mass densities favors a relatively mild redshift evolution, with $\\alpha\\sim 0.26$, with a positive evolution as strong as $\\alpha\\gtrsim 1.3$ excluded at the 99\\% confidence level. A direct match between the characteristic BH mass in the VDF--based and quasar LF--based BH mass functions also yields a mean Eddington ratio of $\\lambda\\sim 0.5-1$ that is roughly constant within $0\\lesssim z \\lesssim 3$. A strong positive evolution in the \\mbh-\\sis\\ relation is still allowed by the data if galaxies increase, on average, their velocity dispersions since the moment of formation, due to dissipative processes. If we assume that the mean velocity dispersion of the host galaxies evolves as $\\sigma(z)=\\sigma(0)\\times(1+z)^{-\\gamma}$, we find a lower limit of $\\gamma\\gtrsim 0.23$ for $\\alpha\\gtrsim 1.5$. The latter estimate represents an interesting constraint for galaxy evolution models and can be tested through hydro simulations. This dissipative model, however, also implies a decreasing $\\lambda$ at higher $z$, at variance with several independent studies. ",
        "introduction": "\\label{sec|intro}  It has now been assessed that most, if not all, local galaxies have a supermassive black hole (BH) at their center, the mass of which is tightly correlated with the velocity dispersion \\sis$\\,$ and other bulk properties of the host galaxy (e.g., Ferrarese \\& Merritt 2000; Gebhardt et al. 2000). However, the sample of local galaxies for which the BH sphere of influence has been resolved amounts to only $\\sim 30$.  It is not clear how representative this small sample is of the whole BH population, and whether the correlations seen in the sample already held in the past.  Peng et al. (2006) have collected a sample of 31 lensed and 18 non-lensed Active Galactic Nuclei (AGN) at redshifts $z>1.7$. They measured rest-frame $R$-band luminosities from $H$-band fluxes and BH masses by applying virial relations based on emission line widths. They found that the BH-to-host galaxy luminosity at $z\\sim 2$ is about the same as the one at $z\\sim 0$. Therefore, once the observed rest-frame luminosity is dimmed through passive evolution to $z\\sim 0$, at fixed BH mass the ratio BH-to-host luminosity grows significantly, and the resulting BH-luminosity normalization is several times higher than the local one. Similar results were derived by McLure et al. (2006), who measured the BH-to-host galaxy mass ratio in a sample of radio-loud AGNs in the redshift range $0<z<2$ finding \\mbh/\\ms$\\propto (1+z)^2$. Shields et al. (2006) found that the CO emission lines in a sample of $z>3$ quasars is very narrow, suggesting bulge mass about an order of magnitude lower than measured in the local universe, at fixed BH mass (see also Coppin et al. 2008). Treu et al. (2007) found that the BH masses in a sample of 20 Seyferts galaxies at $z=0.36$ are offset by an amount of $\\Delta \\log$\\mbh$\\sim 0.5$ at fixed velocity dispersion, which implies an evolution of \\mbh/\\ms$\\propto (1+z)^{1.5\\pm 1.0}$, consistent with that derived by the previous works.  On the other hand, Lauer et al. (2007) have discussed several possible biases which may seriously affect these findings. At high redshifts a sample will be biased toward the most luminous AGNs and more massive BHs. Given the observed scatter in the local relations, especially significant in the \\mbh-host luminosity relation, these massive BHs will be preferentially associated with the less massive, but more numerous galaxies, yielding a false sign of evolution. When the cumulative mass density of AGNs is taken into account, several authors (e.g., Haiman et al. 2004; Marconi et al. 2004; Silverman et al. 2007; Shankar et al. 2008a, hereafter SWM) have shown that once rescaled by a simple constant, it provides a good match to the cosmological star formation rate density. De Zotti et al. (2006) and SWM have shown that the galaxy stellar mass function at $z\\sim 2$, mostly composed of massive early-type galaxies (e.g., Drory et al. 2005), converted into a BH mass density assuming a \\mbh/\\ms$\\,$ ratio 3-5 times higher than in the local universe, would imply a BH mass density already close, if not higher, than that inferred in the local universe, leaving no room for further accretion at $z\\lesssim 2$, where, in fact, a significant fraction of the total AGN energy output is produced. Recently, Ho et al. (2008) compiled a sample of 154 nearby ($z<0.1$) active galaxies showing substantial ongoing BH growth in the most actively accreting AGNs, where BH growth appears to be delayed with respect to the assembly of the host galaxy.  In this paper, we propose a simple, yet robust, way to constrain the degree of redshift evolution in the \\mbh-\\sis$\\,$ relation, that is relatively insensitive to assumptions that relate the SMBH population to quasars. We combine the measured velocity dispersion function (VDF) of local spheroids with a postulated power--law redshift--dependence of the \\mbh-\\sis$\\,$ relation. By comparing the resulting total BH mass density at each redshift with the same quantity inferred from integration of the active galactic nuclei (AGN) luminosity function (see So\\l tan 1982), we find the degree of evolution required in the \\mbh-\\sis$\\,$ relation to match these two independent estimates. This approach yields results based on the ``bulk'' of the active BHs at all redshifts, and is therefore relatively insensitive to possible biases which may affect studies performed on small samples of high--redshift luminous quasars (e.g., Lauer et al. 2007). After describing the sample used in our computations in \\S~\\ref{sec|data}, we proceed to derive our main results in \\S~\\ref{sec|results}. These results are discussed further in \\S~\\ref{sec|conclu}, where we also offer our conclusions. Throughout this paper we use the cosmological parameters $\\Omega_m=0.30$, $\\Omega_\\Lambda=0.70$, and $h\\equiv H_0/100\\, {\\rm km\\, s^{-1}\\, Mpc^{-1}}=0.7$, consistent with the three-- (Spergel et al. 2007) and five--year (Dunkley et al. 2008) data from the {\\it Wilkinson Microwave Anisotropy Probe (WMAP)}. ",
        "conclusions": "\\label{sec|conclu}  In this work we combined the local VDF with the stellar age distributions estimated by Bernardi et al. (2006), to compute the VDF at higher redshifts, $\\Phi(\\sigma,z)$.  In agreement with previous work, we find statistical evidence for downsizing, whereby the stellar populations in galaxies with higher velocity dispersions formed earlier, irrespective of the specific model we adopt for computing the galactic ages. We then computed the BH mass function associated with $\\Phi(\\sigma,z)$ at each redshift $z$, through a BH mass -- velocity dispersion relation whose normalization was allowed to evolve with redshift as $\\propto (1+z)^{\\alpha}$.  Our main underlying assumptions are that most of the growth of the central BH occurs simultaneously (within $\\pm$ 1 Gyr) with the formation of the host's potential well, and that the measured stellar ages represent this formation time to within a similar accuracy.  The BH mass density \\rhovdf$\\,$ inferred from the VDF can then be compared with the accumulated BH mass density implied by the time--integral of the AGN luminosity function, \\rhoz$\\,$.  We find significant evidence that the match between \\rhovdf$\\,$ and \\rhoz$\\,$ implies a relatively mild redshift evolution, with $\\alpha\\lesssim 0.30$, and with values of $\\alpha\\gtrsim 1.3$ excluded at 99\\% confidence. If a positive redshift evolution stronger than $\\alpha\\gtrsim 1$ were to be confirmed independently in the future, then this would be a robust indication that dissipative processes played a significant role in galaxy evolution, resulting in an increase in the velocity dispersion of the spheroid components of individual galaxies with cosmic time. However, we also find evidence that a dissipative model predicts a mean Eddington ratio decreasing with increasing redshift, at variance with several independent studies."
    },
    "0806/0806.0988_arXiv.txt": {
        "abstract": "M\\,87 is the only known non blazar radio galaxy to emit very high energy (VHE) gamma-rays. During a monitoring program of M\\,87, a rapid flare in VHE gamma-rays was detected by the MAGIC telescope in early 2008. The flux was found to be variable above 350\\,GeV on a timescale as short as 1 day at a significance level of $5.6\\sigma$. The highest measured flux reached 15\\% of the Crab Nebula flux. We observed several substantial changes of the flux level during the 13 day observing period. The flux at lower energies (150 -- 350\\,GeV), instead, is compatible with being constant. The energy spectrum can be described by a power law with a photon index of $2.30 \\pm 0.11_\\mathrm{stat} \\pm 0.20_\\mathrm{syst}$. The observed day-scale flux variability at VHE prefers the M\\,87 core as source of the emission and implies that either the emission region is very compact (just a few Schwarzschild radii) or the Doppler factor of the emitting blob is rather large in the case of a non expanding emission region. ",
        "introduction": "M\\,87 is a giant elliptical radio galaxy (RG) of Fanaroff-Riley-I-type \\citep[FR\\,I,][]{FanaroffRiley74} in the Virgo Cluster at a distance of 16\\,Mpc \\citep{Macri99}. It is powered by a supermassive black hole (BH) of $(3.2\\pm0.9)\\times10^{9} \\mathrm{M}_\\odot$ \\citep{Macchetto97}. The M\\,87 jet was the first-ever observed \\citep{Curtis18}, and due to the proximity of M\\,87, its morphological substructures can be resolved and a unique view of its innermost regions is possible. The jet, originating from the RG core, extends to 20'' (\\citealt{Marshall02}; equivalent to a 2\\,kpc projected linear distance). Several compact regions (``knots'') along its axis are resolved in the radio, optical, and X-ray regimes. These knots have similar morphologies in all wave bands, although the X-ray knots appear to be tens of pc closer to the core than the optical and radio knots \\citep{WilsonYang02}. The variable brightness of the knots may be due to several shock fronts in the jet, being responsible for particle acceleration and non-thermal emission. Superluminal motion of the knots has been observed in the optical \\citep{Biretta96} and radio \\citep{Forman07} wave bands, constraining the viewing angle of the jet to $<43^{\\circ}\\pm4^{\\circ}$.  The innermost resolved bright knot HST-1 is located at 0.82'' (64\\,pc) from the core and is the most prominent feature of the jet. HST-1 has shown many flares exceeding the luminosity of the M\\,87 core emission. Its X-ray brightness has increased by a factor $>50$ from 2000 to 2005 \\citep{Harris06}. A correlation between radio, optical, and X-ray luminosity points to a common origin of the emission. The measured superluminal motion in HST-1 is higher than in other knots, suggesting a viewing angle of $<19^{\\circ}$ for this part of the jet. The core itself is variable, too, and also shows a correlation between the emission levels from radio frequencies through X-rays \\citep{Perlman03}.  M\\,87 was not detected by EGRET. The first hint of very high energy (VHE; $>250$\\,GeV) $\\gamma$-ray emission was reported by \\citet{HEGRA-M87}, and later confirmed by \\citet{HESS-M87} and \\citet{VERITAS-M87}. The emission is variable on a timescale of years. The reported $\\approx 2$-day variability \\citep{HESS-M87} narrows down the size of the emission region to be on the order of the light-crossing time of the central BH. With its expected low accretion rate, the M\\,87 core radiation is not strong enough to attenuate significantly TeV $\\gamma$-rays even at 5 Schwarzschild radii ($R_S$) away from the BH \\citep{na07}. All this implies a production region in the immediate vicinity of the M\\,87 core. During later observations, no significant flux variation was found \\citep{VERITAS-M87}. An X-ray--VHE $\\gamma$-ray correlation is expected in most emission models, but was not unambiguously found so far. Whereas \\citet{HESS-M87} claim a hint of a correlation between the soft (0.3--10 keV) X-rays at HST-1 and the VHE $\\gamma$-rays, \\citet{VERITAS-M87} find a year-by-year correlation between the 2--10 keV X-ray flux of the M\\,87 core and the VHE $\\gamma$-ray emission instead, but do not observe a correlation between the two energy bands on shorter timescales.  The radio to X-ray emission of the jet is due to non-thermal synchrotron radiation of relativistic electrons in the jet. The observed knots and flares in M\\,87 point to a complicated morphology with several shock fronts producing these electrons. While all 23 currently known extragalactic VHE $\\gamma$-ray emitters\\footnote{See, e.g., http://www.mppmu.mpg.de/$\\sim$rwagner/sources/ for an up-to-date list} are blazars, M\\,87 is assumed to be a blazar not aligned to our line of sight \\citep{Tsv98}. If the observed VHE emission from M\\,87 is associated with the innermost part of its jet, then blazar emission models may hold. In blazars without prominent disk or broad-line features, the VHE emission is explained by inverse Compton processes involving the synchrotron photons and their parent electron population \\citep[synchrotron self-Compton models, e.g.][]{Maraschi92}. Alternatively, in hadronic models, interactions of a highly relativistic jet outflow with ambient matter \\citep{DarLaor,Beall99}, proton-induced cascades \\citep{PIC}, or synchrotron proton radiation \\citep{Muecke01,Aharonian00} may produce VHE $\\gamma$-rays. In such a scenario, M\\,87 might also account for parts of the observed ultra-high-energy cosmic rays \\citep{uhecr}. It should be noted, that for M\\,87 the location of the VHE emission is still uncertain. Specific emission models for high-energy processes close to the core \\citep{gpk05,ghi05,len07,tg08}, in the large-scale jet \\citep{sta03,hon07}, or in the vicinity of a BH \\citep{na07,riegeraha08} have been developed.  The MAGIC collaboration performed monitoring observations of M\\,87 starting from 2008 January, sharing the task with the VERITAS and H.E.S.S.\\ experiments. Here we report on MAGIC results from a subset of these data, revealing a variability timescale of M\\,87 of 1 day. The energy spectrum and light curves are discussed. ",
        "conclusions": "\\begin{figure} \\begin{center} \\includegraphics[width=0.875\\linewidth]{fig4.eps} \\caption{Differential energy spectra of M\\,87 divided into high (filled circles) and low (open circles) states. See text for the details. The best-fit functions, assuming power laws, are given by the black solid and red dashed-dotted curves, respectively.} \\label{fig:specHL} \\end{center} \\end{figure}  M87 is the only non blazar radio galaxy known to emit VHE $\\gamma$-rays and one of the best-studied extragalactic black-hole systems. To enable long-term studies and assess the variability timescales of M\\,87, the H.E.S.S., VERITAS, and MAGIC collaborations established a regular, shared monitoring of M\\,87 and agreed on mutual alerts in case of a significant detection. During the MAGIC observations, a strong signal of $8\\sigma$ significance was found on 2008 February 1, triggering the other IACTs as well as {\\it Swift} observations. For the first time, we assessed the energy spectrum below 250 GeV, where our observations can be well described by a power law that shows no hint of any flattening.  Our analysis revealed a variable (significance: 5.6\\,$\\sigma$) night-to-night $\\gamma$-ray flux above 350 GeV, while no variability was found in the 150--350\\,GeV range. We confirm the $E>730\\,\\mathrm{GeV}$ short-time variability of M\\,87 reported by \\citet{HESS-M87}. The observed variability timescale is on the order of or even below 1 day, restricting the emission region to a size of $R\\leq \\Delta t\\,c\\,\\delta = 2.6 \\times 10^{15}\\,\\mathrm{cm} = 2.6\\,\\delta R_S$. The Doppler factor $\\delta$ is only relevant for an emission region not expanding while traversing the jet. In case of an expanding-jet hypothesis, the \\textit{initial} radius of the expanding shell is given by $R^* < c\\,\\Delta t$. The emission can occur very close to the BH, provided that the ambient photon density is low enough as to allow the propagation of VHE $\\gamma$ rays. Otherwise the emission region must be located farther away from the BH. In the latter case, the variability constraints can be met only if the emitting plasma does not substantially expand while traversing the jet, or if it moves with $\\delta \\gtrsim 10$.  There exists no lower limit on the size of HST-1, and thus the flux variability cannot dismiss HST-1 as possible origin of the TeV flux. During the MAGIC observations, however, HST-1 was at a historically low X-ray flux level, whereas at the same time the luminosity of the M\\,87 core was at a historical maximum (D. Harris 2008, private communication). This strongly supports the core of M\\,87 as the VHE $\\gamma$-ray emission region.  Our data alone cannot put strong constraints on VHE $\\gamma$-ray emission models. The relatively hard VHE spectrum found for M\\,87 ($\\Gamma\\approx -2.3$) is not unique among the extragalactic VHE $\\gamma$-ray sources if one considers intrinsic spectra, i.e. EBL corrected. Also, we did not measure a high-energy spectral cut-off. The found marginal spectral hardening may be interpreted as a similarity to other blazars detected at VHE, where such hardening has often been observed.  Our results show that a dense TeV monitoring, as exercised by ground-based IACTs, has revealed highly interesting rapid flares in M\\,87. This fastest variability observed so far in TeV $\\gamma$-rays from M\\,87 restricts the size of the $\\gamma$-emission region to the order of $R_S$ of the central BH of M\\,87 and suggests the core of M\\,87 rather than HST-1 as the origin of the TeV $\\gamma$-rays. Results from the entire monitoring campaign, comprising data from other IACTs, will appear in a separate paper."
    },
    "0806/0806.3634_arXiv.txt": {
        "abstract": "{The Crab pulsar emits across a large part of the electromagnetic spectrum. Determining the time delay between the emission at different wavelengths will allow to better constrain the site and mechanism of the emission.We have simultaneously observed the Crab Pulsar in the optical with S-Cam, an instrument based on Superconducting Tunneling Junctions (STJs) with $\\mu$s time resolution and at 2 GHz using the Nan\\c cay radio telescope with an instrument doing coherent dedispersion and able to record giant pulses data.} {We have studied the delay between the radio and optical pulse using simultaneously obtained data therefore reducing possible uncertainties present in previous observations.} {We determined the arrival times of the (mean) optical and radio pulse and compared them using the tempo2 software package.} {We present the most accurate value for the optical-radio lag of 255$\\pm$21~$\\mu$s and suggest the likelihood of a spectral dependence to the excess optical emission asociated with giant radio pulses.} {} ",
        "introduction": "The Crab pulsar was the first pulsar detected as a consequence of its occasional bright pulses (now known as Giant Radio Pulses) at radio wavelengths (Staelin \\& Reifenstein \\cite{staelin:1968}; the first pulsar ever detected was PSR J1921+2153 Hewish et al.\\ \\cite{hewish:1968}). Later on pulsations have been detected at optical wavelengths (Cocke et al.; \\cite{cocke:1969}) and throughout the electromagnetic spectrum at X-rays and $\\gamma$-rays. One feature which is constant over the whole spectrum is the pulsed emission, but the details differ substantially. E.g.\\ the strength of the interpulse, the presence of a pre-cursor to the main pulse (at low radio frequencies) change with wavelength. However, the precise timing of the pulse over many orders of magnitude of energy imposes severe constraints on the emission regions (and mechanism). For example Romani \\& Yadigaroglu (\\cite{romani:1995}) have suggested that the pre-cursor originates at the polar cap, while the pulse and interpulse originate in the outer gap in the magnetosphere, with higher energy pulses being generated at significantly greater heights. This should then be reflected in the timing of the pulses at different energies. Small differences in pulse alignment allow to study the emission regions at relatively small scales (about 10 km per 0.001 phase in case of the Crab pulsar), and precise timing of pulsar light curves throughout the electromagnetic spectrum is thus a powerful tool to constrain theories of the spatial distribution of various emission regions: a time delay most naturally implies that the emission regions differ in position.  In recent years, it has become clear that the main and secondary pulses of the Crab Pulsar (PSR J0534+2200) are not aligned in time at different wavelengths. X-rays are leading the radio pulse by reported values of 344$\\pm$40~$\\mu$s (Rots et al.\\ \\cite{rots:2004}; RXTE data) and 280$\\pm$40~$\\mu$s (Kuiper et al.\\ \\cite{kuiper:2003}; INTEGRAL data) and $\\gamma$-rays are leading the radio pulse by 241$\\pm$29~$\\mu$s (Kuiper et al.\\ \\cite{kuiper:2003}; EGRET data. The uncertainty in this value does not include the EGRET absolute timing uncertainty of better than 100~$\\mu$s). At optical wavelengths, the observations presented a less coherent picture. Sanwal (1999) has reported a time shift of 140~$\\mu$s (optical leading the radio). The uncertainty in this value is 20~$\\mu$s in the determination of the optical peak and 75~$\\mu$s in the radio ephemeris. Shearer et al.\\ (\\cite{shearer:2003}) have reported a lead of 100$\\pm$20~$\\mu$s for simultaneous optical and radio observations of giant radio pulses. Golden et al.\\ (\\cite{golden:2000}) have reported that the optical pulse {\\it trails} the radio pulse by $\\sim$80$\\pm$60~$\\mu$s. Romani et al.\\ (\\cite{romani:2001}) conclude that the radio and optical peaks are coincident to better than 30~$\\mu$s, but their error excludes the uncertainty of the radio ephemeris (150~$\\mu$s). Oosterbroek et al.\\ (\\cite{oosterbroek:2006}, hereafter O06) have studied this issue in detail and found an optical lead of 273$\\pm$100~$\\mu$s. In this paper we will use simultaneous optical and radio observations at 2 GHz to further improve the accuracy of this value.  We also study the possible changes in the peak emission coincident with Giant Radio Pulses (hereafter GP) as previously reported by Shearer et al.\\ (\\cite{shearer:2003}). ",
        "conclusions": "\\begin{figure} \\centerline{ \\includegraphics[width=8cm,angle=-0]{9751f8.eps}} \\caption[]{The radio delay with respect to other frequencies. The S-Cam data point is presented in this paper, the RXTE value is from Rots et al.\\ (\\cite{rots:2004}), the INTEGRAL and EGRET points are taken from Kuiper et al.\\ (\\cite{kuiper:2003}). We have added a 100~$\\mu$s error bar to the EGRET value, because of uncertainties in the EGRET absolute time calibration (see Kuiper et al.\\ \\cite{kuiper:2003}).} \\label{fig:delays_em} \\end{figure}   We have performed simultaneous observations of the Crab pulsar in the optical (3500--8200 \\AA) and radio (2 GHz). The results provide the current best estimate of the optical lead (with respect to the radio peak) of 255$\\pm$21~$\\mu$s. The two major contributions to the uncertainty are the determination (or definition) of the peak in the optical, since the profile is asymmetric, and uncertainties in the DM value which translate into significant uncertainty in the correction of the radio data. Additionally a true jitter in the optical peak on different time scales might take place with an amplitude of 5--10~$\\mu$s. However, our value is an order of magnitude more precise than determined by O06 and fully consistent with the value reported there. The close match with their value (273$\\pm$100~$\\mu$s) must be fortuitious given the large uncertainties on the values, and can be considered a statistical coincidence. A comparison with previous results can be be found in O06, however due to the increased precision we can now conclude that the result of Shearer et al.\\ (\\cite{shearer:2003}), an optical lead of 100$\\pm$20~$\\mu$s, is not consistent with our result. Together with the result of Golden et al.\\ (\\cite{golden:2000}), an optical {\\it delay} of $\\sim$80$\\pm$60~$\\mu$s, which was already inconsistent with the results presented in O06, this raises the possibility of a long-term change in the radio-optical difference.   From a plot of the radio delay throughout the electromagnetic spectrum (see Fig.\\ \\ref{fig:delays_em}) it appears that this delay is approximately constant over more than 7 orders of magnitude. We note that the non-optical results have been derived from a comparison with, non-simultaneous Jodrell Bank observations, which results in larger error bars. This near-constancy of the delay with respect to the radio profile implies a similar emission mechanism and emission location from the optical to $\\gamma$-rays. This is indeed the case in modern models by e.g.\\ Takata \\& Chang (\\cite{takata:2007}).   Our results for the stability of the pulse profile (differences of up to $\\sim$10~$\\mu$s on a time scale of a day and even smaller variations on a time scale of less than a day; see Fig.\\ \\ref{fig:results}) are not consistent with the results from Karpov et al.\\ (\\cite{karpov:2007}) who present variations of up to $\\sim$150~$\\mu$s in a fraction of a day. However, this variation is only present in a small part of their data set (December 2005 and January 2006), while the remainder of their data (obtained in December 1994 and 1999, and November 2003) do not show substantial variations. This suggest that the behaviour observed in December-January 2005/2006 is an anomaly and that night-to-night variations in the Crab optical peak positions amounts to at most $\\sim$10~$\\mu$s.     Our value for the time delay between red and blue photons of -5.9$\\pm1.9$~$\\mu$s/1000 \\AA\\ can be compared to the results of Sanwal (\\cite{sanwal:1999}) who found a delay between the R and U band of -10$\\pm$4~$\\mu$s and Golden et al.\\ (\\cite{golden:2000}) who found that the peaks in B and V are coincident within 10~$\\mu$s. Taken together these three measurements indicate a small delay between blue and red photons. However changes in the FWHM of the pulse profile and the assymmetry of the profile and possible differences in the separation between the main and secondary peak make an interpretation complex.  Current models for high-energy pulsar emission (e.g.\\ Dyks \\& Rudak \\cite{dyks:2003}), describe the peaks as caustics: special relativity effects (aberration of emission directions and time of flight delays due to the finite speed of light) cause photons emitted at different altitudes to pile up at the same pulse phase.  It is difficult if not contrived to imagine how caustics originating from the trailing edge extending along the full length of the open field lines within the magnetosphere could produce phase bunching yielding a common main peak cusp spanning 7 orders of magnitude in photon energy from incoherent synchrotron emission, and less than 1\\% in phase difference from coherent radio emitting sources, spanning in this case some 12 orders in photon energy.  As such arguments are usually presented in the context of outer gap or slot gap model formalisms, these new results argue strongly for a localised source for the observed multiwavelength emission from the Crab pulsar.     That the radio and higher energy emission process are in someway linked is not in doubt, following on from the optical-GP association originally shown by Shearer et al. (\\cite{shearer:2003}). Our confirmation of an excess, albeit at low significance further strengthens the connection between these energy bands. Variations in the optical emission correlated with GPs have been argued to be a consequence of local density enhancements to the plasma stream (Shearer et al. \\cite{shearer:2003}). Recent studies of GPs from the Crab have indicated that these can and do originate as short-lived ($\\sim$ ns), narrowband nanoshots, which have been argued to have their origin in strong, highly localised plasma turbulence (Hankins \\& Eilek, \\cite{hankins:2007}).  That there is no change in phase or shape of the Crab's optical light curve morphology during a GP associated event, just flux, requires an extended region of plasma turbulence to produce the same phase-bunching effect following the caustic model $ansatz$, which is clearly implausible.  The measured time delay across the optical bands represents $\\leq$ 3.0 $\\times$ 10$^{-4}$ of phase, with some evidence of a jitter in the optical peak across several time scales occuring with a similar variation in terms of phase. Again, these effects are not associated with any measurable change in incident flux. Assuming a constant emission source, localised geometric constraints represent the most likely cause of such observational effects. Together these interpretations merit some revision of the fundamental approaches to modeling the Crab pulsar's emission processes, and there may be some use in trying to geometrically constrain emission regions within the magnetosphere first, and then address the local physical conditions that are consistent with the data we have available to us - rather than the other way round.    Finally the anticipated detection of excess optical emission associated with GPs extrapolating from the work of Shearer et al. (2003) was expected to be $<1\\sigma$ prior to our observations, so its marginal detection is somewhat surprising. We note that our passband ($\\sim$3500--$\\sim$8000 \\AA) is quite different from the passband (6000--7500 \\AA) of Shearer et al.\\ (\\cite{shearer:2003}) and might result in different results depending on the spectrum of the ``excess'' optical emission.  If our observed excess had been significant we could have constructed the optical spectrum obtained when a GP occurs and compare it with the reference spectrum when no GP is present. This is possible by using the intrinsic energy resolution of S-Cam (each photon is tagged with a time stamp, but also its energy, Martin et al.\\ \\cite{martin:2006}). Such an approach could reveal the spectrum of the excess emission and consequently would impose important further constraints on the emission mechanism. We regard the acquisition of deeper coincident radio-optical data as critically important to confirm the spectral dependence of such an excess."
    },
    "0806/0806.2859_arXiv.txt": {
        "abstract": "Using two {\\it Chandra} observations we have derived estimates of the dust distribution and distance to the eclipsing high mass X-ray binary (HMXB) Cen X-3 using the energy-resolved dust-scattered X-ray halo. By comparing the observed X-ray halos in 200 eV bands from 2--5 keV to the halo profiles predicted by the Weingartner \\& Draine interstellar grain model, we find that the vast majority ($\\approx$70\\%) of the dust along the line of sight to the system is located within about 300 pc of the Sun, although the halo measurements are insensitive to dust very close to the source. One of the \\ch~observations occurred during an egress from eclipse as the pulsar emerged from behind the mass-donating primary. By comparing model halo light curves during this transition to the halo measurements, a source distance of $5.7 \\pm 1.5$ kpc (68\\% confidence level) is estimated, although we find this result depends on the distribution of dust on very small scales. Nevertheless, this value is marginally inconsistent with the commonly accepted distance to Cen X-3 of 8 kpc. We also find that the energy scaling of the scattering optical depth predicted by the Weingartner \\& Draine interstellar grain model does not accurately represent the results determined by X-ray halo studies of Cen X-3. Relative to the model, there appears to be less scattering at low energies or more scattering at high energies in Cen X-3. ",
        "introduction": "X-ray halos, which appear as diffuse emission surrounding X-ray sources, are created by small-angle scattering of soft X-ray photons from dust grains in the interstellar medium (ISM). Their study can provide information on interstellar grain properties (density, morphology, composition) and on the spatial distribution along the line of sight. Given a dust distribution, variability in the X-ray halo can be used to geometrically derive estimates of the distances to X-ray sources based on the time delays of the photons scattered along the line of sight \\citep{ts73}. This method has been applied in a number of cases \\cite[e.g.,][]{predehl00,me06,audley06,xiang07}. The goal of this work is to estimate the dust distribution and distance to the eclipsing HMXB Cen X-3 using measurements of its dust-scattered halo with the \\cxo.  The theory of X-ray scattering and the production of X-ray halos has been described by a number of authors (e.g., Mathis \\& Lee 1991, Smith \\& Dwek 1998, Draine 2003). Various interstellar grain models have been proposed. Mathis et al. (1977) developed a model composed of silicate and graphite grains with a size distribution of $n(a) \\propto a^{-3.5}$, which reproduces the observed extinction of starlight. Weingartner \\& Drain (2001, hereafter WD01) produced a grain model that additionally accounts for the diffuse infrared and microwave emission emission from the ISM by including sufficient small carbonaceous grains and polycyclic aromatic hydrocarbons (PAHs). \\cite{zubko04} created a series of models by simultaneously fitting the extinction, infrared emission, and elemental abundance constraints, by including, for example, amorphous carbon particles, organic refractory material, water ice, and voids.  One of the many challenges in developing a viable model is that the characteristics of the dust may vary in different Galactic locations due to different local evolutionary histories, ISM phases, and abundances of elements. Given such difficulties, a preeminent interstellar grain model has not been established. For our purposes, we use the WD01 grain model, although we allow the energy dependence of the scattering cross section to be a free parameter.  The basic quantity that determines the characteristics of X-ray halos is the differential scattering cross section $d\\sigma/d \\Omega$. This can be calculated using the exact Mie solution or the simpler Rayleigh-Gans approximation. For the WD01 grain model, the differential scattering cross section as a function of scattering angle $\\phi$ can be approximated by the simple analytic form \\begin{equation} \\frac{d\\sigma}{d\\Omega} \\approx \\frac{\\sigma_{\\rm sca}}{\\pi \\phi^{2}_{{\\rm s},50}} \\left[1+\\left(\\frac{\\phi}{\\phi_{{\\rm s},50}}\\right)^{2}\\right]^{-2} \\end{equation} where $\\phi_{{\\rm s},50} \\approx 360\\arcsec~({\\rm keV}/E)$ is the median scattering angle as a function energy (Draine 2003). It is important to note that the following results are probably somewhat dependent on this chosen form of the differential scattering cross section.  Cen X-3 is one of the brightest accreting X-ray pulsars, and it is one of the six pulsars in which the observation of eclipses has permitted the determination of all orbital and stellar parameters. The rotation period of the pulsar is 4.8 s, and the binary orbital period is 2.1 days. The inclination of the orbital plane and the mass of the neutron star have been estimated to be $i \\sim 70^{\\circ}$ and $M = 1.21 \\pm 0.21$ \\Msun~(Ash et al. 1999). The optical counterpart of Cen X-3 has been identified as an early type star of radius $R \\sim 12$ \\Rsun~and mass $M \\sim 20$ \\Msun~\\citep{rapjoss83}. Accretion in the system probably occurs via an accretion disk (Bildsten 1997; Takeshima et al. 1992), although the mass transfer mechanism probably includes an X-ray excited wind \\citep{ds93}.  Although well-studied, the distance to Cen X-3 remains uncertain. Using a B0 II stellar classification model, an 8 kpc distance to Cen X-3 was obtained \\citep{krz74}, however, subsequent work indicated that the supergiant primary was of type O6\u2013-8 III (Hutchings et al. 1979). \\cite{day91} observed soft emission during two eclipses using {\\em EXOSAT} and attributed it to the dust-scattered halo, from which they derived a distance of $5.4 \\pm 0.3$ kpc. Discrepancies of $\\sim$2.5 kpc in distance estimates to Cen X-3 lead to nearly a factor of two difference in the inferred optical and X-ray luminosities.   Cen X-3 ($l=292\\fdg09, b=0\\fdg34$) is particularly suited for X-ray halo studies. The source is usually very bright, and the interstellar hydrogen column density along the line of sight is $1.20 \\times 10^{22}$ cm$^{-2}$ \\citep{dl90}, implying a large enough optical depth to result in an appreciable halo \\citep{ps95}, yet not so large as to include a substantial contribution due to multiply scattered photons. Moreover, the eclipsing nature of the binary system is advantageous because it should produce the largest fraction of variability in the X-ray halo, which makes a geometric distance measurement more direct.  Section 2 presents the general equations that characterize the shape and variability of X-ray halos. For our study of the X-ray halo of Cen X-3, two \\ch~High Energy Transmission Grating (HETG) observations are used; one observation took place outside of eclipse ($\\Phi_{\\rm orb} \\sim 0.4-0.6$) when there was little variability in the flux (hereafter, called the ``plateau'' phase of the orbit based on the shape of the light curve), and the other occurred during an egress from eclipse (\\S~\\ref{obs}). Because accurate halo measurements critically depend on careful subtraction of the instrumental point spread function (PSF), the shape and normalization of the PSF is parametrized using all \\ch/HETG observations of the unabsorbed sources Her X-1 and PKS 2155-304. These results are presented in the appendix to the paper. The method for deriving the dust distribution to Cen X-3 is discussed in \\S~\\ref{distributionsect}. With a dust distribution established, the large amplitude change in flux during eclipse egress is used to constrain the distance to the system (\\S~\\ref{distancesect}). A comparison of the inferred dust distribution to observations of CO emission, star counts, and interstellar extinction in the direction of Cen X-3 is presented in \\S~\\ref{compare}. We discuss the implications of our results in \\S~5, and we provide a brief summary in \\S~\\ref{summary}. ",
        "conclusions": "\\label{summary} \\begin{enumerate} \\item The vast majority of the dust along the line of sight to Cen X-3 is located within 300 pc ($x < 0.05$) of the Sun. The most likely location of the dust is the local Orion spur. However, we are not sensitive to dust located in the vicinity of Cen X-3 because data within $\\sim$3\\arcsec~of the central source are affected by pile-up.  \\item The geometric distance to Cen X-3 is $5.7 \\pm 1.5$ kpc (68\\% confidence level), although this result is not valid if a large fraction of the dust is concentrated within $\\Delta x =0.01$.  \\item The energy scaling of the scattering optical depth predicted by the WD01 interstellar grain model does not accurately represent the results determined by X-ray halo studies of Cen X-3. Relative to the WD01 model, there appears to be less scattering at low energies or more scattering at high energies in Cyg X-1, GX 13+1, and Cen X-3. \\end{enumerate}"
    },
    "0806/0806.2318_arXiv.txt": {
        "abstract": "IRAS04325+2402C is a low luminosity object located near a protostar in Taurus. We present new spatially-resolved mm observations, near-infrared spectroscopy, and Spitzer photometry that improve the constraints on the nature of this source. The object is clearly detected in our 1.3\\,mm interferometry map, allowing us to estimate the mass in a localized disk$+$envelope around it to be in the range of 0.001 to $0.01\\,M_{\\odot}$. Thus IRAS04325C is unlikely to accrete significantly more mass. The near-infrared spectrum cannot be explained with an extincted photosphere alone, but is consistent with a 0.03-0.1$\\,M_{\\odot}$ central source plus moderate veiling, seen in scattered light, confirming the edge-on nature of the disk. Based on K-band flux and spectral slope we conclude that a central object mass $\\gtrsim 0.1\\,M_{\\odot}$ is unlikely. Our comparison of the full spectral energy distribution, including new {\\it Spitzer} photometry, with radiative transfer models confirms the high inclination of the disk ($\\gtrsim 80\\arcdeg$), the very low mass of the central source, and the small amount of circumstellar material. IRAS04325C is one of the lowest mass objects with a resolved edge-on disk known to date, possibly a young brown dwarf, and a likely wide companion to a more massive star. With these combined properties, it represents a unique case to study the formation and early evolution of very low mass objects. ",
        "introduction": "\\label{intro}  The origin of brown dwarfs has been one of the most intensely debated subjects in cool stars research over the past decade. As of today, there is no consensus on the dominant mechanisms that determine the low-mass end of the stellar initial mass function. Recent modeling has provided predictions for observable properties in the very low mass (VLM) regime, which in principle would permit distinguishing between the proposed scenarios. In practice, however, it has turned out to be challenging to place firm limits on the formation theory for VLM sources \\citep[see reviews by][]{2007prpl.conf..459W,2007prpl.conf..443L}.  Of particular relevance for evaluating formation scenarios are the global properties of disks and envelopes. The combination of high-resolution imaging and analysis of the spectral energy distribution (SED) in the infrared/mm regime is generally acknowledged as the optimum source of information for disk/envelope parameters of young stellar sources. As of today, constraints on radii and masses of brown dwarf disks are sparse: Recent deep mm observations provided first limits on disk masses ($<0.001-0.003\\,M_{\\odot}$) and outer radii \\citep[10-100\\,AU;][]{2006ApJ...645.1498S,2003ApJ...593L..57K}. Only one substellar disk has been resolved with high-resolution imaging \\citep{2007ApJ...666.1219L}, with a derived radius of 20-40\\,AU.  In this context, the young source IRAS04325+2402C (hereafter: IRAS04325C) in the Taurus star forming complex may be of particular interest. HST/NICMOS images of the Class I source IRAS04325+2402 \\citep{1995ApJS..101..117K}, also known as L1535\\,IRS, revealed  IRAS04325C as a faint companion, at a separation of $8\\farcs1$ \\citep{1999AJ....118.1784H}\\footnote{Note that the primary itself is likely to be a subarcsecond binary, see \\citet{1999AJ....118.1784H}}. The object is resolved `into a double-lobed emission separated by a dark lane' \\citep{1999AJ....118.1784H}, resembling the images of edge-on T Tauri disks. Models including an edge-on disk and an envelope are able to reproduce the appearance of the images. From the near-infrared (NIR) photometry, the mass of this possible Class I object was estimated to be between 0.02 and 0.06\\,$M_{\\odot}$, with an upper limit of 0.25\\,$M_{\\odot}$. Based on statistical arguments, \\citet{2004A&A...427..651D} find that the system is likely to be a physically bound binary  (probability 98\\%).  Thus, IRAS04325C is a potential benchmark object to study substellar formation. In this Letter, we present new millimeter interferometry and NIR spectroscopy as well as an analysis of the resolved SED for this source, including new {\\it Spitzer} photometry, with the aim of improving the constraints on the properties of the system, particularly the mass of the central object and the disk$+$envelope. ",
        "conclusions": ""
    },
    "0806/0806.4544_arXiv.txt": {
        "abstract": "Applications of the theory of optimal design of experiments to radial velocity planet search surveys are considered. Different optimality criteria are discussed, basing on the Fisher, Shannon, and Kullback-Leibler informations. Algorithms of optimal scheduling of RV observations for two important practical problems are considered. The first problem is finding the time for future observations to yield the maximum improvement of the precision of exoplanetary orbital parameters and masses. The second problem is finding the most favourable time for distinguishing alternative orbital fits (the scheduling of discriminating observations).  These methods of optimal planning are demonstrated to be potentially efficient for multi-planet extrasolar systems, in particular for resonant ones. In these cases, the optimal dates of observations are often concentrated in quite narrow time segments. ",
        "introduction": "\\label{sec_intro} Since the discovery of the first extrasolar planet orbiting the main-sequence star 51~Pegasi \\citep{MayorQueloz95}, more than $300$ planets orbiting other stars were found, and about $30$ extrasolar systems containing at least two planets are known (see \\emph{The Extrasolar Planets Encyclopaedia} by J.~Schneider, {\\tt www.exoplanet.eu}). The majority of these planets was discovered using the radial velocity (hereafter RV) method. Many multi-planet systems show interesting dynamical behaviour like mean-motion resonances and apsidal corotation locks. The sharp understanding of dynamical regimes in these systems is of a great importance. For example, this information may provide significant constraints on the planetary migration processes \\citep{Beauge06}.  Unfortunately, orbital parameters and masses of many of the extrasolar planets are still poorly known due to a lack and/or a non-uniform sampling of RV observations. Estimations of parameters may be too uncertain or alternative models of the systems may be allowed. Unfortunately, such uncertainties are especially inherent in multi-planetary configurations. It seems that strictly justified strategies of allocating observations still are not widely used in the current practice of RV planet searches. The aim of the present work is to describe several mathematically strict ways of optimal planning of RV observations, based on the general theory of optimal design of experiments \\citep[e.g.,][chapter~10]{Ermakov,ErmZhig,Bard}, and with particular attention to refining orbits in multi-planet configurations. In Section~\\ref{sec_criteria}, several criteria of observing schedule optimality are discussed. In Section~\\ref{sec_planning}, algorithms of optimal planning of RV observations are described in details. In Section~\\ref{sec_apply}, the region of their applicability is discussed. In Section~\\ref{sec_example}, the efficiency of these algorithms is demonstrated using data for real planetary systems. ",
        "conclusions": "\\label{sec_conc} This paper demonstrates how the general tools of the theory of optimal design of experiments can be used for planning RV observations in planet search surveys. Two important practical problems were considered. In the first one, the observations are required to produce the largest improvement in the precision of estimations of exoplanetary parameters. In the second one, the observations are planned to yield maximum information for distinguishing between two alternative orbital models of an exoplanetary system.  These optimising tools are demonstrated using RV data for several real planetary systems. It is shown that these tools may significantly increase the efficiency of observations in RV planet search surveys. They would be especially useful for multi-planet extrasolar systems, in particular for systems containing planet pairs in a mean-motion resonance, for many of which the orbits are still poorly determined. They include, among others, the well-known systems of GJ876, HD82943, and possibly HD37124. The algorithms of optimal scheduling may also be useful in resolving ambiguities concerning planetary orbital configurations, e.g. the alias ambiguity for the HD208487 system."
    },
    "0806/0806.2681_arXiv.txt": {
        "abstract": "We study the impact of the capture and annihilation of Weakly Interacting Massive Particles (WIMPs) on the evolution of Pop III stars. With a suitable modification of the Geneva stellar evolution code, we study the evolution of 20 and 200 M$_\\odot$ stars in Dark Matter haloes with densities between 10$^{8}$ and $10^{11}$ GeV/cm$^3$ during the core H-burning phase, and, for selected cases, until the end of the core He-burning phase. We find that for WIMP densities higher than 5.3 $10^{10}(\\sigma^{SD}_p/10^{-38} \\mbox{ cm}^2)^{-1}$ GeV cm$^{-3}$ the core H-burning lifetime of $20 M_{\\odot}$ and $200 M_{\\odot}$ stars exceeds the age of the Universe, and stars are sustained only by WIMP annihilations. We determine the observational properties of these `frozen` objects and show that they can be searched for in the local Universe thanks to their anomalous mass-radius relation, which should allow unambiguous discrimination from normal stars. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0806/0806.2823_arXiv.txt": {
        "abstract": "We introduce a novel statistical way of analyzing the projected mass distribution in galaxy lenses based solely on the angular distribution of images in quads around the lens center. The method requires the knowledge of the lens center location, but the images' distances from the lens center are not used at all. If the images of a quad are numbered in order of arrival time, $\\theta_1$ through $\\theta_4$, and $\\theta_{ij}$ is the angle between images $i$ and $j$, then we define the 'bisector' plane whose axes are linear combinations of $\\theta_{23}$ and $\\theta_{14}$. The bisector plane of a given lens contains all the quads produced by the lens. We show empirically that all two-fold symmetric lenses with convex, i.e. non-wavy or petal-like isodensity contours are identical in the bisector plane of their quads. We also study lenses with twisting isodensity contours, lumpy substructure, etc. Our results suggest that to reproduce the general characteristics of the observed quad population, kpc-scale substructure must be a common feature of galaxy lenses. ",
        "introduction": "In the last decade or so gravitationally lensed QSOs, both doubles and quads, have been used mostly for the determination of the Hubble parameter (see \\citet{coles08} for the latest work, and summary of earlier results), and for the estimation of the mass distribution in the lensing galaxies. In this paper we will concentrate on the latter.  One can loosely divide the information on the lens mass distribution into two categories: radial and angular. Much attention has been paid in the literature to the sky-projected radial mass distribution in lenses because the slope of the density profile, and its variation with radius is a test of cosmological models \\citep{nfw96,nfw97}. The density profile slope in the central regions is also important because it is affected by the (adiabatic) contraction of dark matter halos in response to the collapsing baryons during galaxy formation \\citep{fsw05,fsb08,g08}.  The angular distribution of lensing mass, for example, the degree of ellipticity, the change in the ellipticity position angle with radius, etc. have received some attention as well \\citep{cdk08,sw06,ok04}, but mostly as ``nuisance'' parameters in determining the radial density profile or the Hubble constant.  It is somewhat ironic that the generally uninteresting ellipticity position angle can be unambiguously estimated by any reasonable modeling method, even by eye \\citep{sw03}, whereas, the more interesting density profile slope is often very uncertain because of the mass-sheet, or steepness degeneracy \\citep{fgs85,s00}.  The positions of lensed images of a quad or a double can also be looked at as consisting of angular and radial information. By radial information we mean the relative spread of images in distance from the lens center. The angular information is the angular separation of the images as viewed from the lens center. For example, in the Cloverleaf, H1413+117, and the Einstein Cross, Q2237+030 any two adjacent images are roughly $90^\\circ$ apart. In doubles, the two non-central images tend to be separated by $\\sim 150^\\circ-180^\\circ$.  Obviously there is no simple one-to-one relation between, say, the radial structure of the lensing mass and the radial distribution of lensed images. However, there are some qualitative connections between the two. For example, a steep lens mass distribution tends to produce quads with narrow radial spread of images, largely independent of the angular distribution of these images, or the ellipticity of the lensing mass. Conversely, if the lensing mass has a shallow density profile the images tend to have a wider radial spread. In the Appendix of this paper we show that one angular property of the lensing mass, its ellipticity position angle can be straightforwardly and rather precisely estimated from the angular positions of the four images of the quad (Section~\\ref{estimatingPA}).  The main work presented in this paper is loosely motivated by the preceding paragraph. Specifically, we ask what information about the lensing mass can be retrieved by looking solely at the angular distribution of lensed images around the lens center. ",
        "conclusions": "We introduce a novel way of analyzing the projected mass distribution in galaxy lenses that relies on the angular distribution of images in quads and doubles around the lens center. If the images of a quad are numbered in order of arrival, as $\\theta_1$, through $\\theta_4$, and $\\theta_{ij}$ is the angle between images $i$ and $j$ then we define the bisector plane whose axes are linear combinations of $\\theta_{23}$ and $\\theta_{14}$. We show empirically that all two-fold symmetric lenses with convex isodensity contours are identical when considered in the bisector plane. We derive an analytical expression for the boundary envelope of the allowed region, for a specific type of lens. These results concerning the invariance of the bisector plane for two-fold symmetric lenses is one of the main findings of the paper. It means, for example, that from the point of view of $\\theta_{23}$ and $\\theta_{14}$ of quads, a Pseudo Isothermal Elliptical Mass Distribution is identical to a circular lens, with any density profile plus an external shear.  This invariance of the bisector planes of two-fold symmetric lenses can be used to examine the structure of the real galaxy lenses. We conclude that the observed quad population was not produced by two-fold symmetric lenses.  We also look at three realistic types of non two-fold symmetric mass distributions, (1) galaxies with twisting isodensity contours, and elliptical galaxies with external shear axis, (2) galaxies with single substructure clumps, and (3) galaxies with two substructure clumps. It appears that only the last type of lenses is able to reproduce the real quad population. This of course does not mean that all galaxies with observed quads are of type (3), but it does suggest that kpc-scale substructure is a common feature in galaxy lenses.  To confirm and quantify this conclusion a much more detailed exploration of the parameter space of non two-fold symmetric lenses is needed. Such a study should also include potential sources of bias in the quads. For example, in this paper we have assumed that the real lenses represent a random sampling of the relevant region in the source plane; in other words, all sources have the same weights. This means that we have neglected magnification bias, which makes sources at certain source plane locations more magnified, and hence more likely to enter a magnitude limited sample. The bias is probably negligible for quads, since they are already highly magnified; after all, quads are closely related to Einstein rings. It is unlikely that there is a missing population of faint quads. However, the magnification bias could be an issue for the doubles, and will need to be taken into account in future work.  Two final notes are in order. First, the lumpy substructure we refer to here is different from that searched for using image flux anomalies, e.g. \\citet{m04}. In the latter case substructure lumps are small, and have to lie close to the line of sight to the images. Our substructure lumps are larger, kpc-sized, more extended and can live anywhere within the central several kpc of the galaxy lens center. Second, the varied and complex lumpy substructure that our analysis implies the lenses should have argues strongly for using non-parametric, or semi-parametric modeling techniques."
    },
    "0806/0806.2538_arXiv.txt": {
        "abstract": "We have undertaken a  pilot survey for faint QSOs in the UKIDSS Ultra Deep Survey Field using the KX selection technique. These observations exploit the very deep near-infrared and optical imaging of this field from UKIRT and Subaru to select candidate QSOs based on their $VJK$ colours and morphologies.  We determined redshifts for 426 candidates using the AAOmega spectrograph on the AAT in service time.  We identify 17 QSOs ($M_B\\ls -23$) in this pilot survey at $z=1.57$--3.29.  We combine our sample with an X-ray selected sample of QSOs in the same field (a large fraction of which also comply with our KX selection) to constrain the surface density of QSOs with $K\\leq 20$, deriving limits on the likely surface density of 85--150\\,deg$^{-2}$. We use the good image quality available from our near-infrared imaging to detect a spatially extended component of the QSO light which probably represents the host galaxies. We also use our sample to investigate routes to improve the selection of KX QSOs at faint limits in the face of the significant contamination by compact, foreground galaxies. The brightest examples from our combined QSO sample will be used in conjunction with a large VLT VIMOS spectroscopic survey of high redshift galaxies in this region to study the structures inhabited by gas, galaxies and growing super-massive black holes at high redshifts in the UKIDSS UDS. ",
        "introduction": "Simulations of the distribution of baryons in the Universe predict that much material is spread out in a ``cosmic web'' between the galaxies, groups and larger structures visible in conventional surveys.  The main components of this material are the Ly$\\alpha$ forest, which traces highly ionized hydrogen of low neutral column density and low chemical enrichment distributed through the lower density regions, and metal absorption lines which trace massive galaxy halos (and hence higher density regions) via their metal-enriched gas. To probe this material and investigate its 3-dimensional distribution we must exploit QSOs as bright background sources to trace the web through its absorption.  An optimal survey of the cosmic web would use a grid of distant bright QSOs as probes of the intervening matter, with the redshift and equivalent width of these absorbers providing information on the relative distribution and clustering of gas around galaxies and larger structures, as well as providing insights into the chemical enrichment and heating of this material (e.g.\\ Morris \\& Januzzi 2006).  Studies comparing the results from high-redshift galaxy surveys around individual bright QSOs have provided unique insights into the effect of star formation on the gas surrounding young galaxies (e.g.\\ Adelberger et al.\\ 2005).  However, these studies are observationally expensive -- it would be much more efficient if a large number of galaxies and QSOs could be compared within a single field as the number of absorber--galaxy pairs scales as N$^2$.  The difficulty is that the surface density of bright QSOs is low and hence faint QSOs are required to set up a dense grid of probes needed to sample the full range of cosmic structures in a single field (Prescott et al.\\ 2006).  Nevertheless, with enough QSOs in a single field the long integration time required to undertake absorption-line analysis on relatively faint background sources (AGN or galaxies) become worthwhile.  We are undertaking a highly-sampled survey of the cosmic web in a uniquely well-studied field: the UKIDSS (Lawrence et al.\\ 2007) Ultra Deep Survey (UDS) field. The UDS is the deepest, panoramic near-infrared observations so far undertaken with the goal of reaching $K\\sim 23$ and $J\\sim 25$ over $\\sim 0.8$ sq.\\ degrees (Foucaud et al.\\ 2007).  Equally deep $BV\\! RIz$ imaging is available across the whole field from Subaru/XMM Deep Survey (SXDS, Furusawa et al.\\ 2008) and this region has also been extensively studied at X-ray wavelengths as part of the same project (Ueda et al.\\ 2008). These data are being analysed to study the build-up of massive galaxies out to $z\\sim3$ and its variation with environment (Foucaud et al.\\ 2007; Lane et al.\\ 2007; Cirasulo et al.\\ 2007).  This field is also the target of a VLT Large Programme (PI: O.\\ Almaini) to use the VIMOS spectrograph to obtain redshifts for $\\sim 5$,000 high-redshift $K$-selected galaxies.  In this paper we present a parallel pilot survey to expand the size of the QSO sample in this field to allow us to relate these luminous galaxies and structures to their surrounding enriched and unenriched gas from a highly-sampled grid of sightlines. Combining these two unique datasets will yield powerful insights into the influence of galaxies and AGN (at the peak era in their activity) on low-density gas in their environments, material which is required for the long-term fueling of the star formation in galaxies.  To construct a highly-sampled grid of QSOs in this field we have used $VJK$-band imaging from UKIDSS/UDS to apply the KX method (Warren et al.\\ 2000; Croom et al.\\ 2001; Sharp et al.\\ 2002) to identify a sample of $\\sim$\\,800 $V<23$ potential AGN candidates in this region.  The good seeing available in the UDS imaging, 0.8$''$ FWHM across the whole field, yields relatively clean catalogs of point sources -- which are a mix of AGN, stars and compact galaxies.  The KX method relies on the power-law shape of AGN continuua across the $VJK$ bands, compared to the $H$-band ``bump'' arising from the opacity minimum in stellar atmospheres, to distinguish between those point sources which are AGN and those which are stars.  This color selection has the additional advantage that it provides an unbiased selection of QSOs independent of their dust reddening out to high redshifts (Warren et al.\\ 2000), compared to the classical ``UVX'' selected samples.  We use a cosmology with $H_{0}=70$\\,km\\,s$^{-1}$, $\\Omega_{M}=0.3$ and $\\Omega_{\\Lambda}=0.7$ in which 1$''$ corresponds to 8.2\\,kpc at $z=2.5$ and 8.5\\,kpc at $z=1.5$.  All quoted magnitudes are  Vega and as the galactic reddening in our field is low, $E(B-V)=0.02$, we have not applied any extinction corrections to the colours or magnitudes in this paper. ",
        "conclusions": "We present the results from a pilot survey for QSOs in the UKIDSS UDS field.  We employ the KX selection method of Warren et al.\\ (2000) to select a sample of faint candidate QSOs from the very deep optical and near-infrared imaging available for this field.  We have then followed up these sources with the 400-fibre AAOmega spectrograph in service time on the AAT to identify the high-redshift QSOs in the $\\sim 700$  KX-selected compact sources in our catalogue. We identify 17 QSOs with M$_B\\leq -23$ in this pilot survey. Supplemented by existing spectroscopy of X-ray selected AGN in this field (Akiyama et al.\\ 2008) this yields a combined QSO sample with a surface density of $\\sim 85$--150\\,degree$^{-2}$ at $K_{tot}\\leq 20$ (where the range reflects the uncertainty in the spectroscopic incompleteness in our sample).  This sample is roughly $4\\times$ deeper than previous KX-selected surveys over comparable areas (Jurek et al.\\ 2008) and $40\\times$ larger than the previous surveys at similar $K$-band depth (Croom et al.\\ 2001).  We have used the good image quality of the UKIDSS UDS  data to create composite $K$-band images of our QSO sample and a matched sample of stars.  The light profiles of the two composites show that the QSOs are more spatially extended.  The extended component comprises some 25--40\\% of the total $K$-band light of the QSO and has a spatial extent of $\\sim 2$--5\\,kpc.  This most likely represents the host galaxy of the QSO.  We analyse our sample to attempt to refine the KX selection technique to reduce the considerable contamination by low redshift, $z\\ls 0.6$, compact galaxies.  We find that colour or morphological selections focusing on bluer passbands (e.g.\\ an additional $(B-V)$ cut or limits on image FWHM in the $B$-band) would reduce the contamination in our parent sample.  However, these approaches would bias the resulting sample against redder or higher-redshift QSOs and so may not be desirable for certain applications where completeness is important.  Alternatively, if the goal is to maximise the number of high-redshift QSOs detected, then the KX selection using a more conservative FWHM limit, $\\leq 1.0''$ in the case of the $0.8''$-FWHM UDS imaging,  can significantly reduce the foreground contamination while at the same time yielding high completeness ($\\sim 90$\\%).  From our sample we cull a subsample of $\\sim 21$ QSOs at $z>2.5$ and $V<23$ which are suitable for an absorption line study (see Adelberger et al.\\ 2005) employing deep blue spectroscopy with VIMOS on VLT.  In addition, the full sample of QSOs (including those at lower redshifts) provide a map of the distribution of rapidly growing black holes within our survey volume, enabling us to relate these to the growth of mass-selected samples of galaxies and the surrounding absorption-line systems.  These luminous AGN are expected to have measurable feedback effects on the gas -- an issue of great interest to current models of galaxy formation (in which such forms of feedback are an essential part, Bower et al.\\ 2006)."
    },
    "0806/0806.2230_arXiv.txt": {
        "abstract": "A search for solar axions has been performed using an axion helioscope which is equipped with a 2.3-m long 4\\,T superconducting magnet, a gas container to hold dispersion-matching gas, PIN-photodiode X-ray detectors, and a telescope mount mechanism to track the sun. A mass region around $m_a = 1{\\rm\\,eV}$ was newly explored. From the absence of any evidence, analysis sets a limit on axion-photon coupling constant to be $\\gagg<\\mbox{5.6--13.4}\\times10^{-10}\\rm GeV^{-1}$ for the axion mass of $0.84<m_a<1.00\\rm\\,eV$ at 95\\% confidence level. It is the first result to search for the axion in the $\\gagg$-$m_a$ parameter region of the preferred axion models with a magnetic helioscope. ",
        "introduction": "Quantum chromodynamics (QCD) is the theory of the strong interactions. Although QCD has proven remarkably successful, there is a blemish called the strong $CP$ problem. The strong $CP$ problem is that the effective Lagrangian of QCD has $CP$ violating term but it is not observed, i.e., the experimental value of neutron electric dipole moment is smaller than expected by many orders of magnitude. Peccei and Quinn proposed an attractive solution to solve this problem \\cite{axion-bible1,axion-bible2,axion-bible3,axion-bible4,axion-bible5}. They introduced a new global $U(1)$ symmetry, Peccei--Quinn (PQ) symmetry. When PQ symmetry is spontaneously broken, a new effective term arises in QCD Lagrangian which cancels the $CP$ violation term. The solution also predicts a new pseudo Nambu--Goldstone boson, axion. The expected behavior of an axion is characterized mostly by the scaling factor of the PQ symmetry breaking, $f_a$, and so its mass, $m_a$, which is directly related to $f_a$ by $m_a=6\\times10^{15} [\\mathrm{eV}^2]/f_a$.  Axions are expected to be produced in stellar core through their coupling to photons with energies of order keV. Especially, the sun can be a powerful source of axions and the so-called `axion helioscope' technique may enable us to detect such axions directly \\cite{sikivie1983,bibber1989}.   The principle of the axion helioscope is illustrated in Fig.~\\ref{fig:principle}. Axions would be produced through the Primakoff process in the solar core. The differential flux of solar axions at the Earth is approximated by \\cite{bahcall2004,raffelt2005} \\begin{eqnarray} \\d \\Phi_a/\\d E&=&6.020\\times10^{10}[\\mathrm{cm^{-2}s^{-1}keV^{-1}}] \\nonumber\\\\ &&{}\\times\\left(\\gagg\\over10^{-10}\\mathrm{GeV}^{-1}\\right)^2 \\left( \\frac{E}{1\\,\\mathrm{keV}}\\right)^{2.481} \\exp \\left( -\\frac{E}{1.205\\,\\mathrm{keV}}\\right), \\label{eq:aflux} \\end{eqnarray} where $\\gagg$ is the axion-photon coupling constant% \\footnote{The formula is different from the one in our previous paper \\cite{sumico1997,sumico2000}, but both coincide numerically with good approximation.}% . Their average energy is 4.2\\,keV reflecting actually the core temperature of the sun($\\sim 3kT$), since low energy axion production is suppressed due to screening effects\\cite{raffelt1986}. Then, they would be coherently converted into X-rays through the inverse process in a strong magnetic field at a laboratory. The conversion rate is given by \\begin{equation} P_{a\\to\\gamma} = \\left| \\frac{g_{a\\gamma\\gamma}}{2} \\exp\\left[-\\int_{0}^{L} \\mathrm{d}z \\Gamma/2 \\right] \\times \\int_{0}^{L}dz B_\\bot\\exp \\left[i \\int_{0}^{z} \\mathrm{d}z^{\\prime}\\left( q - \\frac{i\\Gamma}{2}\\right) \\right]\\right|^2, \\label{eq:prob} \\end{equation} where $z$ and $z^{\\prime}$ are the coordinate along the incident solar axion, $B_\\bot$ is the strength of the transverse magnetic field, $L$ is the length of the field along $z$-axis, $\\Gamma$ is the X-ray absorption coefficient of helium, $q=(m_\\gamma^2-m_a^2)/2E$ is the momentum transfer by the virtual photon, and $m_\\gamma$ is the effective mass of the photon which equals zero in vacuum. Eq. (\\ref{eq:prob}) is reduced to \\begin{equation} P_{a\\to\\gamma} = \\left( \\frac{g_{a\\gamma\\gamma}B_\\bot L}{2}\\right)^2 \\left[ \\frac{\\sin(qL/2)}{qL/2}\\right]^2, \\label{eq:prob_plain} \\end{equation} in case $q$ and $B_\\bot$ is constant along $z$-axis and $\\Gamma=0$.  In 1997, the first phase measurement \\cite{sumico1997} was performed using an axion helioscope with a dedicated superconducting magnet which is identical to the one used in the present experiment, except that the gas container was absent and the conversion region was vacuum. Its sensitivity was limited to the axion mass region of $m_a<0.03\\rm\\,eV$ due to a loss of coherence by non-zero $q$ in Eq.~(\\ref{eq:prob_plain}).  If one can adjust $m_\\gamma$ to $m_a$, coherence will be restored for non-zero mass axions. This is achieved by filling the conversion region with gas. A photon in the X-ray region acquires a positive effective mass in a medium. In light gas, such as hydrogen or helium, it is well approximated by \\begin{equation} m_\\gamma=\\sqrt{4\\pi\\alpha N_e\\over m_e}, \\end{equation} where $\\alpha$ is the fine structure constant, $m_e$ is the electron mass, and $N_e$ is the number density of electrons. We adopted cold helium gas as a dispersion-matching medium. Here, light gas was preferred since it minimizes self absorption by gas. It is worth noting that helium remains at gas state even at 5--6\\,K, the operating temperature of our magnet. Since the bore of the magnet is limited in space, the easiest way is to keep the gas at the same temperature as the magnet. Moreover, axions as heavy as a few electronvolts can be reached with helium gas of only about one atmosphere at this temperature.  In this way, in 2000, the second phase measurement \\cite{sumico2000} was performed to search for sub-electronvolt axions. This experiment, together with the first phase measurement of 1997 \\cite{sumico1997} with vacuum conversion region, yielded an upper limit of $\\gagg<\\hbox{6.0--10.5}\\times10^{-10}\\rm GeV^{-1}$ (95\\% CL) for $m_a<0.27\\rm\\,eV$.   In this Letter, we will present the result of the third phase measurement in which we scanned the mass region between $0.84<m_a<1.00\\rm\\,eV$ using the upgraded apparatus to withstand higher pressure gas. ",
        "conclusions": "The axion mass around 1\\,eV has been scanned with an axion helioscope with cold helium gas as the dispersion-matching medium in the $4{\\rm\\,T}\\times2.3\\rm\\,m$ magnetic field, but no evidence for solar axions was seen. A new limit on $\\gagg$ shown in Fig. \\ref{fig:exclusion} was set for $0.84<m_a<1.00\\rm\\,eV$. It is the first result to search for the axion in the $\\gagg$-$m_a$ parameter region of the preferred axion models~\\cite{GUT_axion} with a magnetic helioscope.  When the complete unmanned operation is ready with the automatic cable handling system, the mass scan will be continued to cover still wider mass range around 1\\,eV.  \\begin{ack} The authors thank the former director general of KEK, Professor H. Sugawara, for his support in the beginning of the helioscope experiment. This research was partially supported by the Japanese Ministry of Education, Science, Sports and Culture, Grant-in-Aid for COE Research and Grant-in-Aid for Scientic Research (B), and also by the Matsuo Foundation. \\end{ack}"
    },
    "0806/0806.0235_arXiv.txt": {
        "abstract": "{} {Several frequency combinations are widely used in the analysis of stellar oscillations for comparisons between models and observations. In particular, the ``small separations'' can help constraining the stellar cores. We showed in a previous paper that they can change sign, in contradiction with the ``asymptotic theory'', and that this behaviour could correspond to signatures of convective and/or helium cores. Here we analyse this behaviour in detail by systematic modelling during stellar evolution.} {We computed evolutionary tracks for models with various masses (from 1.05 to 1.25 \\msol) and various chemical compositions, with and without overshooting. We computed the adiabatic oscillation frequencies of the models and analysed the evolution of their small separations along an evolutionary track.} {We find that, for all cases, the stars go through a stage during their evolution, where the small separations computed between degrees $\\ell=0$~ and ~$\\ell=2$ become negative in the observed range of frequencies. This behaviour is clearly related to the signature of a helium-rich core. We discuss the consequences for interpreting of the acoustic frequencies observed in solar-type stars.} {} ",
        "introduction": "Asteroseismology has proved to be an excellent tool for deriving stellar parameters more precisely than ever before: mass, age, radius, metallicity, helium abundance value, stellar gravity, effective temperature. This has been discussed recently in detail for the case of the exoplanet-host star $\\iota$~Hor (Vauclair et al. \\cite{vauclair08}). Spectroscopic observations lead to values of \\teff, $g$, and [Fe/H], which differ for different authors, according to the scale they use for the effective temperature determinations. The associated error bars are quite large. In this framework, observing the acoustic oscillations of the star and determining the mode frequencies is extremely helpful.  Once the oscillations of the star have been observed long enough (typically eight nights with HARPS), the analysis of the Fourier transform may lead to the determination of the ``large separations'', defined as $\\Delta\\nu_{n,\\ell}=\\nu_{n+1,\\ell}-\\nu_{n,\\ell}$ and to the identification of the modes, as discussed in Bouchy et al. (\\cite{bouchy05}). Meanwhile, various evolutionary tracks are computed with different input parameters (mass, chemical composition, presence or not of overshooting, etc). For a given set of parameters, only one model may reproduce the observed frequencies satisfactorily (Bazot et al. \\cite{bazot05}; Soriano et al. \\cite{soriano07}; Laymand \\& Vauclair \\cite{laymand07}; Vauclair et al. \\cite{vauclair08}). As shown in Vauclair et al. (\\cite{vauclair08}), the various models obtained in this way have similar gravities and ages. The other parameters are constrained with the help of the spectroscopic observational boxes.  Asteroseismology can also give information about the internal structure of the stars and, more specifically, about the regions where the sound velocity changes rapidly. This happens in various transition layers, such as the limits of convective regions or layers with strong helium gradients. The transition layers that occur in stellar outer regions (bottom of outer convective zones, diffusion-induced helium gradients) may be characterised owing to the reflexion of the acoustic waves on the stellar surface and on the region of rapidly varying sound velocity. This is generally studied with the help of the ``second differences'', defined as $\\delta_{2}\\nu_{n,\\ell}=\\nu_{n+1,\\ell}+\\nu_{n-1,\\ell}-2\\nu_{n,\\ell}$, which present oscillations with a period related to the acoustic depth of the transition layer (e.g. Gough \\cite{gough90}; Monteiro \\& Thompson \\cite{monteiro98}; Vauclair \\& Th\\'eado \\cite{vauclair04}; Castro \\& Vauclair \\cite{castro06}). On the other hand, the transition layers in the stellar internal regions are better characterised using the ``small separations'', defined as $\\delta\\nu_{n,\\ell}=\\nu_{n,\\ell}-\\nu_{n-1,\\ell+2}$ (e.g. Roxburgh \\& Vorontsov \\cite{roxburgh94}). In a series of papers, Roxburgh \\& Vorontsov (see Roxburgh \\& Vorontsov \\cite{roxburgh07} and references therein) show how the presence of a convective core could lead to oscillations in the small separations.  In a previous paper dedicated to the study of the exoplanet-host star HD 52265 (Soriano et al. \\cite{soriano07}), we showed that in some cases the small separations, which should be positive in first approximation, could become negative. We explained how this special behaviour was related to either a convective core or a helium core with abrupt frontiers, resulting from the presence of a convective core in the past history of the star.  In the present paper, we give a general analysis of this effect for solar-type stars, with or without overshooting. We focused our analysis on stars that have either a solar metallicity or that are overmetallic, which is the case of exoplanet-host stars. We show that, along their evolutionary tracks, all stars go through a stage where the small separations between the $\\ell= 0$ and $\\ell= 2$ modes become negative, whether near the end of the main-sequence or at the beginning of the subgiant branch. This leads to a special behaviour in the ``echelle diagram'', which is obtained by plotting the mode frequencies in ordinates and the same frequencies modulo the large separations in abscissae (e.g. Soriano et al. \\cite{soriano07}). Instead of remaining parallel, as predicted by the asymptotic theory (Tassoul \\cite{tassoul80}), a crossing point appears between the $\\ell=0$ and $\\ell=2$ curves. This may also happen between the $\\ell=1$ and $\\ell=3$ curves, but it is less frequent and we concentrate here on the $\\ell=0$ and $\\ell=2$ case.    \\begin{figure*} \\begin{center} \\includegraphics[angle=0,width=10cm]{fig1a.ps} \\includegraphics[angle=0,width=10cm]{fig1b.ps} \\includegraphics[angle=0,width=10cm]{fig1c.ps} \\end{center} \\caption{Evolutionary tracks computed for solar metallicity [Fe/H]$_{\\odot}$ (upper panel), overmetallicity [Fe/H]=0.30 with a helium abundance following the law of the chemical evolution of galaxies Y$_G$ (middle panel), and overmetallicity with a solar helium abundance \\ysol (lower panel). All these tracks are computed without overshooting. The masses represented are: 1.05, 1.10, 1.15, 1.20, and 1.25 \\msol.  The crosses represent the ``transition model'' for each evolutionary track, or first model from which we found negative small separations below 3.5 mHz. The characteristics of these models are displayed in Tables~\\ref{tab1}, \\ref{tab2}, and \\ref{tab3}.} \\label{fig1} \\end{figure*}   \\begin{figure*} \\begin{center} \\includegraphics[angle=0,width=10cm]{fig2a.ps} \\includegraphics[angle=0,width=10cm]{fig2b.ps} \\includegraphics[angle=0,width=10cm]{fig2c.ps} \\end{center} \\caption{Evolutionary tracks computed with the same hypothesis as in Fig.~\\ref{fig1} but with overshooting. Here, the overshooting parameter is fixed as $\\alpha_{ov}=0.20$. The masses represented are the same ones as in Fig.~\\ref{fig1}. On each evolutionary track, we have represented the ``transition model'' by crosses, or first model from which we have negative small separations below 3.5 mHz. The characteristics of these models are displayed in Tables~\\ref{tab4}, \\ref{tab5}, and \\ref{tab6}.} \\label{fig2} \\end{figure*}    For comparisons with observations, this behaviour becomes important if the crossing points correspond to frequencies below the cut-off frequency of the acoustic waves, which is the highest frequency for which the waves are reflected on the atmospheric regions. As this cut-off frequency depends on the characteristics of the stellar outer regions and may vary from star to star, we chose to study all the cases for which the crossing point occurs below 3.5 mHz. Along the evolutionary track, we computed the models for which this behaviour appears for the first time, for various masses, and give their ages. In the following, we refer to these special models as ``transition models''. As the apparition of this behaviour is related to the radius of the convective or helium core, negative small separations appear earlier when overshooting is introduced.  This particular behaviour of the small separations could be used to derive the size of stellar cores and the presence of overshooting for stars at the end of the main-sequence. A direct application to the exoplanet-host star $\\mu$ Arae will be given in a forthcoming paper. ",
        "conclusions": "As powerful as asteroseismology may be, it will always give much less precise results about stellar internal structure than helioseismology does for the Sun. While more than ten million modes can be identified for the Sun, only tens of modes may be observed for stars. In this framework, skilled recipes are needed to better characterise the stellar structures from the observed acoustic frequencies. Among the frequency combinations that have been proposed to test stellar interiors, the ``small separations'', which represent the frequency differences between modes $\\ell$, $n$, and $\\ell+2$, $n-1$, are known to give important hints about the stellar cores.  In the present paper, we have shown that the small separations, which should, in a first approximation, have a positive sign, may become negative during stellar evolution. Moreover, this behaviour occurs systematically for solar-type stars and may become observable at the end of the main-sequence or at the beginning of the subgiant branch for the $\\ell=0$~-~$\\ell=2$ separations. It also occurs for the $\\ell=1$~-~$\\ell=3$ separations, at higher frequencies that are generally outside the observing range. For this reason, we concentrated on the $\\ell=0$~-~$\\ell=2$ ones. We found that this effect is directly related to the presence of a helium core, itself induced by convection in the central stellar regions. Although convective cores by themselves lead to a significant decrease in the sound velocities in the central stellar regions, the resulting gradients are not sufficient to inverse the sign of the small separations, until enough helium has accumulated. At the end of the main sequence, or at the beginning of the subgiant branch, the chemical gradient induces a much larger sound velocity gradient than convection alone and leads to negative small separations in the observing frequency region. When the star lies at the beginning of the subgiant branch, the convective core has disappeared but the effect continues due to the remaining helium core. From then on, this behaviour goes on, until the stellar structure becomes strongly modified on the subgiant branch.  The occurrence of negative small separations stems from the modes of different $\\ell$ not seeing exactly the same stellar regions. In particular, while the $\\ell=0$ modes travel down to the stellar centre, the $\\ell=2$ modes travel between the stellar surface and a ``turning point'' that depends on its frequency. Modes with low frequencies may penetrate the stellar core, whereas modes with high frequencies stay outside. The $\\ell=2$ modes for which the turning point corresponds to the edge of the stellar core play a special role in this respect. In most cases, negative small separations begin to occur at the frequency of these modes.  The interest of this behaviour of the small separations is evident, as it may help for characterising convective and helium cores and for constraining overshooting, at least for stars at the end of the main sequence or the beginning of the subgiant branch. However, this may not be very simple in real life. When we observe the acoustic spectra of a solar-type star, the identification of the modes is not easy. That there may be a mode crossing point adds a new difficulty in this respect. Once an observed ``echelle diagram'' has been obtained, it may be difficult to decide whether it shows a crossing point or not.  We will discuss these observational versus theoretical points in a forthcoming paper, in which we will give a new analysis of the star $\\mu$ Arae (after Bazot et al. \\cite{bazot05}). We will show that models with crossing points for the $\\ell=0$~-~$\\ell=2$ modes can be obtained if overshooting is added to the convective core. However, only very specific models may fit the overall echelle diagram, so that the constraints are really strong. We will also show that, in the presence of rotational splitting, the modes can be separated without ambiguity, as $\\ell=0$ modes cannot show the splitting effects.  As a conclusion, the negative small separations are quite a helpful tool for characterising convective and helium cores, and they may give strong constraints for the possible overshooting at the edge of the convective cores."
    },
    "0806/0806.2189_arXiv.txt": {
        "abstract": "{} {We study the stellar and dust properties of a well-defined sample of local elliptical galaxies to investigate the relationship between host galaxy properties and nuclear activity.} {We select a complete sample of 45 ellipticals from the Palomar spectroscopic survey of nearby galaxies, which includes 20 low-luminosity active galactic nuclei classified as LINERs and 25 inactive galaxies. Using a stellar population synthesis method, we compare the derived stellar population properties of the LINER versus the inactive subsamples.  We also study the dust and stellar surface brightness distributions of the central regions of these galaxies using high-resolution images obtained with the {\\it Hubble Space Telescope}.} {Relative to the inactive subsample, ellipticals hosting LINERs share similar total optical and near-infrared luminosity, central stellar velocity dispersions, and nuclear stellar populations as judged from their luminosity-weighted ages and metallicities.  LINERs, on the other hand, have a larger fraction of core-type central surface brightness profiles and a much higher frequency of circumnuclear dust structures.} {Our results support the suggestion that LINERs are powered by low-luminosity AGNs rather than by young or intermediate-age stars. Nuclear activity in nearby elliptical galaxies seems to occur preferentially in those systems where sufficient cold interstellar material has managed to accumulate, perhaps via cooling condensations from hot gas.} ",
        "introduction": "Supermassive black holes (SMBHs) are ubiquitous in elliptical galaxies. The discovery of a tight correlation between black hole mass and the stellar velocity dispersion of the bulge of the host galaxy (\\citealt{Gebhardt}; \\citealt{Ferrarese}) suggests that the formation and growth of central black holes are an integral part of the formation of galactic bulges.  However,  the majority of SMBHs in nearby ellipticals are not very active. In the Palomar spectroscopic survey of nearby galaxies (Ho et al. 1995, 1997a), only about 50\\% of the ellipticals show detectable emission-line nuclei, most of which ($\\sim$87\\%) are classified as low-ionization nuclear emission-line regions (LINERs; \\citealt{Heckman}; see Ho 2008 for a review).  The physical origin of LINERs has been intensively debated. A variety of recent observations, however, provided convincing evidence that LINERs are low-luminosity active galactic nuclei (AGNs).  This includes the detection of broad H$\\alpha$ emission (Ho et al. 1997b), some of which show double-peaked profiles (Ho et al. 2000; Shields et al. 2000; Barth et al. 2001) or are polarized (\\citealt{Barth1, Barth2}), ultraviolet variability (Maoz et al. 2005), compact X-ray cores (e.g., Ho et al. 2001; Terashima \\& Wilson 2003), and compact radio cores (e.g., Nagar et al. 2005).  Ho (2008) suggests that nearly {\\it all}\\ LINERs are genuine AGNs with low accretion rates.  The low accretion rates of the SMBHs in nearby ellipticals probably stem from the lack of enough cold gas to fuel the central engine, as it is generally accepted that nearby elliptical galaxies contain mostly old stars and not much cold interstellar material. However, apart from the high detection rate of nuclear optical emission lines in spectroscopic surveys (Phillips et al. 1986; Ho et al. 1997a), H$\\alpha$ + [N~II] narrow-band imaging has shown that elliptical galaxies certainly have plenty of extended warm ($\\sim 10^4$ K) gas (e.g., \\citealt{Shields1991}). By coadding images from the {\\it Infrared Astronomy Satellite}\\ survey scans, \\citet{Knapp} show that $\\sim$50\\% of ellipticals contain cool dust. \\citet{Temi} find that in 16 elliptical galaxies observed with the {\\it Infrared Space Observatory}\\ the dust masses are on average 10 times larger than those previously estimated from {\\it Infrared Astronomy Satellite}\\ observations. Most recently, \\citet{Kaneda} even detected polycyclic aromatic hydrocarbon emission features in four elliptical galaxies with the {\\it Spitzer} Infrared Spectrograph. All these results suggest that the presence of a significant amount of interstellar material in elliptical galaxies is quite common, although its origin is still uncertain.  We attempt to determine the factor(s) that control the onset of nuclear activity by comparing the stellar population, central light distribution, and dust content for ellipticals with different levels of nuclear activity drawn from the Palomar spectroscopic survey of nearby galaxies (Ho et al. 1995, 1997a). This paper is organized as follows.  The sample and analysis method are described in Section 2; the results of the stellar population synthesis are presented in Section 3; the discussion and conclusions are given in Section 4. ",
        "conclusions": "Analyzing a statistically complete sample of 45 elliptical galaxies from the Palomar spectroscopic survey of nearby galaxies, comprising 20 LINERs and 25 inactive members, we show that galaxies hosting LINER nuclei, compared to systems lacking optical evidence for nuclear activity, have very similar stellar properties, both on global and nuclear scales.  Specifically, we find that the two galaxy subgroups have essentially indistinguishable total optical and near-infrared luminosities, central stellar velocity dispersions, and luminosity-weighted mean stellar ages and metallicities on nuclear (\\lax 300 pc) scales.  The lack of young or intermediate-age stars in LINERs strongly rule out starburst and post-starburst models for the excitation of LINER-like emission, further strengthening the proposition that LINERs are low-luminosity AGNs (Ho 2008).  Our results are qualitatively consistent with those from previous studies of the Palomar survey.  \\citet{Ho6} systematically investigated the stellar population of the emission-line nuclei in the Palomar survey and found that LINERs generally contain evolved stellar populations.  Other evidence concerning the old ages of the nuclear stellar population in nearby low-luminosity AGNs is summarized in Ho (2008).  Ho et al., however, did not derive detailed ages or metallicities from a population synthesis analysis.  This work is the first quantitative analysis of its kind for the Palomar survey.  On the other hand, studies of larger samples of LINERs drawn from the Sloan Digital Sky Survey reach somewhat different conclusions.  While LINERs on average have older stars than Seyferts (Kewley et al. 2006), Graves et al. (2007) and Schawinski et al.  (2007) show that LINERs have a tendency to exhibit somewhat younger populations compared to inactive early-type galaxies devoid of emission lines. {\\bf The difference may arise from the following two points. First, the sources from the Sloan Digital Sky Survey have redshifts $z \\approx 0.05-0.1$, much more distant than the Palomar galaxies. As discussed by Ho (2008), aperture effects render the interpretation of LINER emission in distant galaxies particularly ambiguous.  Second, the stellar population analysis method of this paper is different from that used in SDSS galaxies. However, the reported age differences are relatively small and may require more data than we have and/or a more detailed analysis to be checked.}  Surprisingly, {\\bf we find that} LINER host galaxies have a greater tendency to exhibit central core light profiles.  Lauer et al. (2005) suggest that core galaxies are more luminous, more massive, and may be older than power-law galaxies, seemingly at odds with our results.  Considering the modest size of our sample, however, and the large scatter in the statistical correlations between profile type and global properties, this disagreement is not too alarming.  By far the clearest distinction between active and inactive ellipticals is that the former has a much higher likelihood of containing circumnuclear dust. This finding is in excellent agreement with that of \\citet{Lauer05}, who also found a strong correlation between the incidence of nebular line emission---which in ellipticals invariably signifies a LINER classification (Ho 2008)---and central dust features: 90\\% of emission-line nuclei show detected central dust structures, to be compared with just 4\\% for the lineless systems.  \\citet{Ravindranath} similarly suggested a connection between H$\\alpha$ emission and dust features, and the tendency for LINERs to be more dusty.  Recently, Sim\\~{o}es Lopes et al. (2007) presented  strong evidence that circumnuclear dust is correlated with AGN activity in early-type galaxies.  Among 34 early-type galaxies classified as AGNs, all show circumnuclear dust structures, to be compared with only 26\\%  in the control sample of inactive systems.  Taken collectively, the above evidence suggests that ellipticals with weakly active, LINER nuclei are more gas-rich than their inactive counterparts. Given the lower stellar density in the cores, we actually expect {\\it less}\\ cold gas accumulation from stellar mass loss, not more, among the LINER hosts (see, e.g., Soria et al. 2006).  On the other hand, massive ellipticals are surrounded by hot X-ray halos (e.g., Jones et al. 2002), whose cooling time is sufficiently short that cooling flows should develop.  We speculate that the central dust features in core-type LINER ellipticals may partly originate from cooling condensations in the X-ray-emitting gas.  The ``extra'' cold gas and dust, in combination with direct spherical accretion from the hot gas itself, may be responsible for sustaining the weak nuclear activity. Perhaps elliptical galaxies with LINER nuclei identify the subclass of early-type galaxies in which this mode of accretion is significant."
    },
    "0806/0806.2376_arXiv.txt": {
        "abstract": "We study magnetospheric structure surrounding rotating magnetized neutron star with nonvanishing NUT (Newman-Tamburino-Unti) parameter. For the simplicity of calculations Goldreich-Julian charge density is analyzed for the aligned neutron star with zero inclination between magnetic field, gravitomagnetic field and rotation axis. From the system of Maxwell equations in spacetime of slowly rotating NUT star, second-order differential equation for electrostatic potential is derived. Analytical solution of this equation indicates the general relativistic modification of an accelerating electric field and charge density along the open field lines by the gravitomagnetic charge. The implication of this effect to the magnetospheric energy loss problem is underlined. In the second part of the paper we derive the equations of motion of test particles in magnetosphere of slowly rotating NUT star. Then we analyze particle motion in the polar cap and show that NUT parameter can significantly change conditions for particle acceleration. ",
        "introduction": "The theoretical study of radio pulsars can be traced back to the work of Goldreich $\\&$ Julian (1969) where the existence of magnetosphere with charge-separated plasma around rotating magnetized neutron stars has been suggested. A spinning magnetized neutron star generates huge potential differences between different parts of its surface. The cascade generation of electron-positron plasma in the polar cap region (Sturrock 1971; Ruderman $\\&$ Sutherland 1975) means that the magnetosphere of a neutron star is filled with plasma, screening the longitudinal electric field. This screening results in the corotation of plasma with a neutron star. Such a rotation is not possible outside the light cylinder, thus it forms essentially different groups of field lines: closed, i.e., those returning the stellar surface, and open, i.e., those crossing the light cylinder and going to infinity. As a result, plasma may leave the neutron star along the open field lines and it is generally thought that pulsar radio emission is produced in the open field line region well inside the light cylinder (the radius at which the corotation speed equals $c$).  Study of plasma modes along the field lines is boosted by the pioneering works of Goldreich $\\&$ Julian (1969), Sturrock (1971), Mestel (1971), Ruderman $\\&$ Sutherland (1975), and Arons $\\&$ Scharlemann (1979). The subsequent achievements and some new ideas are reviewed by numerous authors, for example, by Arons (1991), Michel (1991), Mestel (1992), and Muslimov $\\&$ Harding (1997). Although a self-consistent pulsar magnetosphere theory is yet to be developed, the analysis of plasma modes in the pulsar magnetosphere based on the above-mentioned papers provides firm grounds for the construction of such a model.  The existence of strong electromagnetic fields is one of the most important features of rotating neutron stars observed as pulsars. It was shown starting pioneering paper of~\\citet{d55} that the electric field is induced due to the rotation of highly magnetized star. The general relativistic effect of dragging of inertial frames is very important in pulsar magnetosphere~\\cite{bes90,mt90} and a source of additional electric field of general relativistic origin.  The field-aligned electric field is driven by deviation of space charge from the Goldreich-Julian charge density, which is determined by the magnetic field geometry. Therefore as it has been noted by several authors (e.g.~\\citet{bes90}, \\citet{mt90}, Muslimov \\& Harding (1997); Dyks et al (2001), Mofiz $\\&$ Ahmedov (2000)) the general relativistic frame dragging effects on the field geometry in plasma magnetosphere of rotating neutron stars is a first-order effect, which has to be carefully included in a self-consistent model of pulsar magnetospheric structure and associated electromagnetic radiation. It was first shown by \\citet{mt90} and independently by \\citet{bes90} that general relativistic effects due to frame dragging are crucial for formation of the field aligned electric field and particle acceleration in pulsar magnetosphere.  At present there is no any observational evidence for the existence of gravitomagnetic monopole, that is of exotic space-time, called NUT space (Newman, Unti and Tamburino~\\cite{nut63}). Therefore it is interesting to study the electromagnetic fields and processes in NUT space with the aim to get new tool for studying new important general relativistic effects which are associated with nondiagonal components of the metric tensor and have no Newtonian analogues.  \\citet{MusTsy92} initiated the detailed general relativistic derivation of magnetospheric electromagnetic fields around rotating magnetized neutron star. In this paper, we attempt to extend work of \\citet{MusHar97} by including NUT parameter in studying particle acceleration along the open field lines of a rotating neutron star. In \\S~\\ref{ms} general relativistic equations describing the electrodynamics of a rotating NUT star are formulated. The equations are rewritten in the frame of reference corotating with the rotating NUT star. A detailed analysis of Goldreich- Julian charge density is also performed. A general equation governing electrostatic potential in the magnetosphere is derived.  As it was shown by several authors, in particular by \\citet{ss03}, that the general relativistic effect on particle motion in pulsar magnetosphere can not be neglected. Due to this reason  in \\S~\\ref{acceleration} we are solving equations of motion for the charged particles for the region near the magnetic pole just above the stellar surface. From study of the particle motion along the field lines it has been shown that the effect of NUT parameter, coupling with the frame-dragging effect, plays noticeable role in the particle's dynamics. In \\S~\\ref{concl} we conclude our findings and discuss them for further investigations.  Throughout, we use a space-like signature $(-,+,+,+)$ and a system of units in which $G = 1 = c$ (However, for those expressions with an astrophysical application we have written the speed of light explicitly.). Latin indices run $1,2,3$ and Greek ones from $0$ to $3$. ",
        "conclusions": "\\label{concl}   We have considered astrophysical processes in the polar cap of pulsar magnetosphere in assumption of slowly rotating NUT space-time. In particular, general-relativistic corrections to the Goldreich-Julian charge density, electrostatic scalar potential and accelerating component of electric field being parallel to magnetic field lines in the polar cap region due to the presence of gravitomagnetic charge are found. The presence of NUT-parameter slightly modulates Goldreich-Julian charge density near the surface of the star. However as it is known the effective electric charge density i.e.  difference between Goldreich-Julian charge density $\\rho_{GJ}$  (being proportional in the case of flat space-time to $\\mathbf{\\Omega\\cdot B}$) and electric charge density (being proportional to $\\mathbf{B}$) in the star magnetosphere is responsible for the generation of electric field being parallel to magnetic field lines. This difference is equal to zero at the surface of the star and changes with the distance from it due to the fact that $\\rho$ can not compensate $\\rho_{GJ}$. General relativitistic terms arising from the dragging of inertial frames and presence of the gravitomagnetic charge give very important additional contribution to this difference. Both of these terms depend on the radial distance from the star as $1/r^3$ and have equally important influence on the value of accelerating electric field component generated in the magnetosphere near the surface of the neutron star.  These results are applied to find an expression for electromagnetic energy losses along the open magnetic field lines of the slowly rotating NUT star. It is found that in the case of non-vanishing NUT-parameter an additional important term to the coefficient of standard magneto-dipole energy losses expression appears. Comparison of effect of NUT-parameter with the already known effects shows that it can not be neglected. Obtained new dependence may be combined with astrophysical data on pulsar periods slow down and be useful in further investigations on possible detection of the gravitomagnetic monopole.  It is also shown that the presence of gravitomagnetic charge has influence on the conditions of particles motion in the polar cap region. From derived results it can be seen, that NUT-parameter modulates the period of oscillations of particle's momentum. As the effect, connected with NUT-parameter has the angular velocity of the star's rotation in denominator, it will be increased for the old stars, where the frame-dragging effect will be decreased. Thus possible candidates for verification of the existence of NUT parameter are old compact objects with extremely low rotation period."
    },
    "0806/0806.3693_arXiv.txt": {
        "abstract": "A complete manual search has been carried out of the list of 285423 objects, nearly all of them galaxies, identified in the COSMOS field that are brighter than $I=25$. Two certain and one highly probable new gravitational lenses are found, in addition to the lenses and candidate lens systems previously found by Faure et al. (2008). A further list of 112 candidate lens systems is presented. Few of these are likely to be true gravitational lens systems, most being star-forming rings or pairs of companion galaxies. It is possible to examine of order 10$^6$ objects by eye in a reasonable time, although reliable detection of lenses by such methods is likely to be possible only with high-resolution data. The loss of completeness involved in a rapid search is estimated as up to a factor of 2, depending on the morphology of the lens candidate. ",
        "introduction": "There are a number of reasons why gravitational lenses, systems in which a distant quasar or galaxy are multiply imaged by foreground galaxies, are important. Because lensing is a purely gravitational effect, the inverse problem can be solved in order to reconstruct the lensing mass distribution from the observed images, independently of the nature of the matter in the lens. Variable lensed sources can be monitored for time delays between variations of the images, leading to determination of the Hubble constant (Refsdal 1964, for reviews see e.g. Courbin 2003, Kochanek \\& Schechter 2004, Jackson 2007); non-correlated image variability allows us to study microlensing by stars in relatively high-redshift galaxies (e.g. Agol \\& Krolik 1999, Wambsganss 2001, Ofek \\& Maoz 2003, Morgan et al. 2008); lensing statistics, along with a cosmological model, allow constraints on galaxy evolution (e.g. Chae \\& Mao 2003, Ofek, Rix \\& Maoz 2003, Matsumoto \\& Futamase 2008); and the magnification induced by lenses allows us to study intrinsically fainter objects.  Since the discovery of the first gravitational lens by Walsh et al. (1979), over 100 lens systems have been discovered. The CASTLES compilation website (Kochanek et al. 2008) lists 100 cases in which quasars are multiply imaged by foreground galaxies. In addition, large numbers of lens systems in which the lensed sources are extended galaxies are now being produced by searches such as the SLACS survey (Bolton et al. 2006). Searches for lensed arcs in large-area sky surveys are now beginning to be successful; for example, Cabanac et al. 2007 report progress in the use of the Canada-France-Hawaii Telescope Legacy Survey for this purpose. Recently, Faure et al. (2008) report the use of the Cosmic Evolution Survey (COSMOS, Scoville et al. 2007) as a means of identifying gravitational lenses. COSMOS is a 2 square degree area in which HST/ACS imaging, together with space- and ground-based followup by telescopes at many different wavelengths, has identified over two million objects.  In the COSMOS survey there are 285423 sources brighter than 25th magnitude in the ACS $I$-band observations. Faure et al. (2008) considered a subset of 9452 of these objects. This sample was chosen to be the most likely to contain gravitational lens systems. The most likely lens systems are those at moderate redshift ($0.2<z<1.0$), with intrinsically high luminosity ($M_V<-20$) and those spectrally classified as early type galaxies. They find 20 good candidates, many of which are likely on morphological grounds to be gravitational lenses. In addition, they list 47 other objects in which detection of a single arc gives some indication of lensing by the primary galaxy.  In the present work, all of the COSMOS catalogue images are manually inspected. Two definite gravitational lenses are found, together with a third highly probable system. The search strategy is described in section 2, and the definite lens systems are presented in section 3. Finally in section 4, implications and future prospects are briefly discussed. ",
        "conclusions": "Th number of lens systems that we should expect in the COSMOS survey is likely to be higher than the number of secure lens systems as identified by eye. For example, Miralda-Escude \\& Lehar (1992), in their study of lensed arcs in optical surveys, predict an approximate number of 100 per square degree in a survey to a depth of $B=26$, which for blue objects roughly corresponds to the survey depth of the search undertaken in this work. Since the COSMOS footprint is about 1.6 square degrees, this suggests that over 100 lens systems should be present. In addition, Bolton et al. (2008) report a total of $\\sim$70 clearly identified lenses from the SLACS survey, which results from the study of $\\sim$50000 spectra of luminous red galaxies (LRGs). Such galaxies are more likely than the average to be gravitational lenses, because of their larger lensing cross-section; nevertheless, even if the lensing cross-section of an average galaxy is a factor of 10 lower than an LRG we would expect several dozen lens systems. Most of the systems are therefore likely to be concealed within the single-arc systems of Faure et al. (2008) or those of Figure 1.  Direct confirmation of such lenses is not easy, because it requires spectroscopy of $I=25$ arcs to determine redshifts. This is a non-trivial task with an 8-m class telescope, and a programme is under way to investigate the best candidates. In practice, however, it is likely to be impossible to be substantially complete for lenses in a large, blind optical survey such as COSMOS.  In the future, it is likely that arc detection algorithms will be applied to large-area surveys such as COSMOS and CFHTLS, and indeed steps in this direction are described by Cabanac et al. (2007) and by Seidel \\& Bartelmann (2008). The latest algorithm to be used is an automated robot which explicitly fits each image as a potential lens system and adjusts the model to maximise source plane flux (Marshall et al. 2008). Other methods include searches for multiple blue objects around likely lens galaxies (Belokurov et al. 2007).  Provided that high-resolution images are available, examination by eye of large samples is likely to be a valuable adjunct, both to provide example sets for arc detection algorithms, or even for neural network versions of these algorithms. The examination performed in this work is crude, but it is surprisingly quick to perform; in fact, a sample of a million sources could be processed by one investigator with about a month of dedicated effort. The disadvantage is that many forthcoming surveys will not have diffraction-limited resolution, making galaxy-scale lens systems much more difficult to detect. Despite this, manual inspection of galaxies from forthcoming medium-deep surveys, such as the VST and VISTA KIDS/VIKING surveys, each of which will image several hundred thousand galaxies, will be useful. Lensing by galaxies assisted by groups or clusters will be detectable; such inspections will be more incomplete for lens systems such as those presented here."
    },
    "0806/0806.2343.txt": {
        "abstract": "Spitzer-MIPS \\24m observations and ground-based optical imaging and spectroscopy of the rich galaxy cluster Abell 851 at $z = 0.41$ are used to derive and compare star formation rates from the mid-IR \\24m and from \\OII $\\lambda\\lambda$3727 emission.  Many cluster galaxies have SFR(\\24m) /SFR(\\OII) $>>$ 1, indicative of star formation in regions highly obscured by dust.  We focus on the substantial minority of A851 cluster members where strong Balmer absorption points to a starburst on a $10^{8-9}$ year timescale.  As is typical, two types of galaxies with strong Balmer absorption are found in A851: with optical emission (starforming), and without (post-starburst).  Our principal result is that the starforming variety, so-called e(a) galaxies, are mostly detected (9 out of 12) at \\24m --- for these we find typically SFR(\\24m)/SFR(\\OII) $\\sim 4$.  Strong Balmer absorption and high values of SFR(\\24m)/SFR(\\OII) indicate moderately active starbursts; both observations support the picture that e(a) galaxies are the active starbursts that feed the post-starburst population.  While \\24m detections are frequent with Balmer-strong objects (even 6 out of 18 of the supposedly ``post-starburst'' galaxies are detected), only 2 out of 7 of the continuously starforming `e(c)' galaxies (with weak Balmer absorption) are detected --- for them, SFR(\\24m)/SFR(\\OII) $\\sim 1$.  Their optical spectra resemble present-epoch spirals that dominate today's universe;  we strengthen this association by showing that SFR(\\24m)/SFR(\\OII) $\\sim$ 1 is the norm today.  That is, not just the amount of star formation but also its mode has evolved strongly from $z \\sim 0.4$ to the present.  We fit spectrophotometric models in order to measure the strength and duration of the bursts and to quantify the evolutionary sequence from active- to post-starburst.  Our results harden the evidence that moderately active starbursts are the defining feature of starforming cluster galaxies at $z \\sim 0.4$.   %\\vspace{1.0cm} ",
        "introduction": "A remarkable change in our view of galaxy evolution over the last few decades has been the realization that the properties of galaxies are changing on relatively short time scales, in cosmic terms.  The Butcher-Oemler (1978) effect, first at-odds with the then-prevailing picture that strong galaxy evolution occurred only within the first few billion years, is now part of the orthodoxy -- a large and still-growing body of evidence shows that the stellar populations and even the morphologies of many galaxies have undergone significant change, even within the last 4 billion years.  Butcher and Oemler's discovery of a substantial population of luminous starforming galaxies in intermediate-redshift clusters -- a great contrast with cluster populations today where only a small fraction of galaxies are still actively forming stars -- is now widely confirmed and accepted.  Less well known, but likely as important for understanding the causes of the Butcher-Oemler effect, is the discovery by Dressler and Gunn (1983) that most of these starforming galaxies in the cores of distant clusters show a starburst signature in their optical spectra: both starbursts and post-starbursts are common.  Later work has confirmed the ubiquity of this result (Couch \\& Sharples 1987, Wirth \\etal 1994, Poggianti \\etal 1999 -- P99, Tran \\etal 2003).\\footnote{Interested readers will appreciate a review of the often-cited rebuttal by Balogh \\etal (1999).  These authors actually found a 5\\% incidence of strong Balmer-line sources --- for a lower redshift cluster sample, with a higher threshold of Balmer strength --- quite consistent with the other work cited above.  Balogh \\etal discounted their detections as simple noise in their spectra, after determining the ``noise'' from the \\Hd distribution in their full sample.  Dressler \\etal (2004) show that the detections were in fact reliable -- and the noise greatly overestimated --- due to a difficulty with the bandpass measurement technique, which for small \\Hd equivalent width returns negative values --- see also Prochaska \\etal 2007.}  Data from these studies suggest that a significant fraction of field populations at $z \\sim 0.5$ is also undergoing moderate bursts of star formation  (Dressler \\etal 2004).  This behavior contrasts with today's population of field galaxies, for which continuous star forming is the dominant mode, and starbursts are rare.  This is yet further evidence of rapid evolution of the galaxy population, at least as important, since it applies to field and cluster galaxies alike.  In this paper we address one of the long-standing uncertainties of this subject --- the amount of dust-obscured star formation that has not been generally accounted-for in representative field and cluster populations at $z \\sim 0.5$. Although Balmer absorption is a reliable indicator of past starburst activity, a coherent picture requires evidence of the bursts themselves. This has been hampered by the possibility that much of the star formation is obscured by dust, and therefore not part of the tally made with optical-to-near-IR observations. The remarkable sensitivity of the Spitzer Space Telescope with its \\24m MIPS camera is allowing great progress in removing this uncertainty.  As part of a successful Spitzer program to observe two rich clusters of galaxies, we investigated archival MIPS data for Abell 851, a cluster that is well studied in the optical, x-ray, and radio, and for which a significant spectroscopic sample has been obtained and explored in a companion paper (Oemler et al. 2008, hereinafter Oem08).  Our purpose here is limited to the implications for the amount and character of star formation in distant clusters, specifically the subject  of ``starburst vs. continuously starforming'' for distant galaxies.  Issues of causes and mechanisms are beyond the scope of this paper, but Spitzer data for additional rich, distant clusters and their surrounding field populations will be able to advance these discussions considerably.  The paper is organized as follows.  In section 2 we lay out the issues, in Section 3 we address the ability of mid-infrared imaging to add to our knowledge of the star formation rates of distant galaxies, in Section 4 we present the optical and mid-IR data for A851, and Section 5 we give the results of combining these data sets.  In Section 6 we compare our detections of dusty starbursts in A851 with similar observations of present-epoch spirals.  Finally, in Section 7, we apply stellar population models to the optical spectra + MIPS data and discuss the results in terms of a consistent picture of the starburst activity in distant clusters. ",
        "conclusions": "%Galaxies with strong \\Hd and relatively weak \\OII really are dusty starbursts %Starbursts are an important part of the life of starforming galaxies  Our study of optical spectroscopy and Spitzer-MIPS observations of a spectroscopic sample of cluster members of A851 at $z = 0.42$ has yielded the following results: (1) Detection with MIPS at \\24m is a strong function of spectral type: almost all passive galaxies are undetected, and most normal starforming galaxies (type e(c), weak \\Hd absorption) are also undetected. However, some post-starburst galaxies (strong \\Hd) and {\\it most} e(a) galaxies (strong \\Hd {\\it and} \\OII emission) are detected.  The non-detection of e(c) galaxies as compared to the detection of most e(a) galaxies occurs across the full range of luminosity sampled --- in fact, the e(c)'s are systematically more massive in this sample --- so this difference in detection is not a selection effect. (2) For the e(a) galaxies, the SFR derived from \\24m is typically several times that inferred from \\OII, indicating a substantial amount of dust obscuration of ongoing star formation in these systems.  Ongoing star formation is seen in some galaxies classified as post-starburst, for which SFR(\\24m)/SFR(\\OII) is likely to be as much as a factor of 10, but this it represents a significant decline from the burst that came before.  (3) The high rates of star formation in e(a) galaxies --- factors of 2-3 above the past average for these systems --- qualifies them as moderate starbursts capable of accounting for the post-starbursts seen in A851, although for the specific, small samples we have here, there is a factor of 2--3 lower mass associated with the e(a) compared to the k+a galaxies. (4) A small sample of present-day starforming galaxies shows that these typically have SFR(IR)/SFR(optical) $\\sim$ 1.  We suggest that these galaxies are like the e(c) galaxies in A851, which would be consistent  with these objects having \\24m flux below the detection limit of these observations.  The assembly of these observations supports the model in which the e(a) galaxies and even some post-starburst galaxies are actually dust-obscured starbursts, with a factor of $\\sim$4 of the star formation hidden.  As we will show in a forthcoming paper, the significant fraction of e(a) spectra among field galaxies at intermediate-redshift suggests that starbursts are also relatively common among field galaxies at intermediate-redshift.  The rarity of such objects in the present-epoch universe indicates that star formation in ordinary spirals was more bursty --- more variable ---  as recently as 4 Gyr ago.  The extent and magnitude of this phenomenon will become increasingly clear as further Spizter-MIPS observations become available for intermediate-redshift clusters and the field."
    },
    "0806/0806.1005_arXiv.txt": {
        "abstract": "We determine characteristic timescales for the viscous damping of $r$-mode oscillations in rapidly rotating compact stars that contain quark matter.  We present results for the color-flavor-locked (CFL) phase of dense quark matter, in which the up, down and strange quarks are gapped, as well as the normal (ungapped) quark phase. While the ungapped quark phase supports a temperature window $10^8$ K $\\leq T\\leq 5\\times 10^9$ K where the $r$-mode is damped even for rapid rotation, the $r$-mode in a rapidly rotating pure CFL star is not damped in the temperature range $10^{10}$ K$\\leq T\\leq 10^{11}$ K.  Rotating hybrid stars with quark matter cores display an instability window whose width is determined by the amount of quark matter present, and they can have large spin frequencies outside this window.  Except at high temperatures $T\\geq 10^{10}$~K, the presence of a quark phase allows for larger critical frequencies and smaller spin-periods compared to rotating neutron stars. If low-mass X-ray binaries contain a large amount of ungapped or CFL quark matter, then our estimates of the $r$-mode instability suggest that there should be a population of rapidly rotating binaries at $\\nu\\gtrsim 1000$ Hz which have not yet been observed. ",
        "introduction": "\\label{sec_intro}    Neutron stars are astrophysical test-beds for the physics of strongly-interacting matter in regimes of density that are not presently accessible either in terrestrial experiments or by numerical studies of lattice-regularized QCD. Therefore, connecting theoretical advances in the knowledge of the equation of state (EoS) of dense matter to astrophysical observations of neutron stars is essential. The EoS, which plays an important role in determining a compact star's properties and in its evolution with time, is constrained by laboratory data only up to densities of about $2\\rho_N$, where $\\rho_N\\approx 2.6\\times 10^{14}\\ {\\rm g~cm^{-3}}$ is the saturation density of nuclear matter. At asymptotically high densities where perturbative calculations in QCD are reliable,  quark matter is conjectured to be in a color superconducting state commonly called the color-flavor-locked (CFL) phase~\\cite{Alford98}.  Meanwhile, models of low-energy QCD at moderate baryon chemical potentials $\\mu_B\\gtrsim 1$ GeV in a restricted temperature domain $T\\leq T_c\\sim 50$ MeV predict a rich patchwork of color superconducting phases in the QCD phase diagram~\\cite{Steiner02,Schmitt,Shovkovy,Sharma,Neumann},  with important consequences for the properties of self-bound quark stars or neutron stars whose core density is large enough for deconfined quark matter to exist there. We refer to the former as strange stars and the latter as hybrid stars throughout this work.   The effects of color superconductivity on the EoS are small, though potentially important for the surface structure of strange stars~\\cite{Lugones:2002va, JRS05,ARSS06}, or the nuclear-quark interface in hybrid stars~\\cite{ABPR05}. More interesting are its effects on transport phenomena, as pairing among the up, down and strange quarks breaks the $SU(3)$ global chiral symmetry spontaneously~\\cite{Alford98}, leading to (pseudo-) Nambu-Goldstone modes that determine the material response to external perturbations. See the recent review Ref.~\\cite{Alford:2007xm} for the effective theory of these Nambu-Goldstone bosons. The pairing pattern that leaves the maximal residual symmetry  in the three-flavor case is the  CFL phase. In this phase, the up, down and strange quarks of all three colors participate in Cooper  pairing and the order parameter reduces the QCD symmetry $SU(3)_L\\times SU(3)_R\\times SU(3)_C$ to the diagonal subgroup $SU(3)_{L+R+C}$ locking the color and flavor symmetries. Apart from the $8$ pseudo-Nambu-Goldstone bosons associated with the spontaneous breaking of chiral symmetry, other light modes in this phase are the exactly massless Nambu-Goldstone boson (the superfluid phonon), which is a long-wavelength fluctuation of baryon number, and the massless photon, which is effectively a combination of the vacuum photon and a gluon.  Studies of transport coefficients reveal that the thermal conductivity and specific heat of the CFL phase at temperatures much smaller than $T_c$, the critical temperature for pairing, is controlled by the phonon and the photon, while the electrical conductivity is determined by thermally excited electron-positron pairs~\\cite{Shovkovy:2002kv}. In the CFL phase the long-term cooling of the star is controlled by neutrino emission from the scattering and weak decay of phonons~\\cite{JPS}. These novel transport phenomena imply  that neutrino transport and neutron star cooling curves have the  potential to distinguish between the numerous possible patterns of Cooper pairing that might occur at stellar densities~\\cite{Jot,Blaschke1,Blaschke2,JRS06}.   Another set of transport coefficients that are modified by Cooper pairing are the bulk and shear viscosities. Bulk viscosity can lead to energy dissipation when the rate of chemical equilibration in matter is comparable to the driving frequency of the volume perturbation. Shear viscosity determines the relaxation of momentum components perpendicular to the direction of fluid flow. Together, they characterize the material response to compressional and shearing forces that are typically present in rotating and pulsating compact stars, and serve to damp out large-amplitude oscillatory modes of the fluid. Moreover, since the fluid inside a star is self-gravitating, these oscillations can couple to metric perturbations, causing the star to emit gravitational waves. Indeed, the primary physical interest in exploring viscosities of color superconducting quark matter is to assess their role in damping out the so-called $r$-mode, a particular oscillatory mode of the fluid that dissipates the star's rotational energy by efficiently coupling the angular momentum of the star to gravitational waves~\\cite {Andy, Friedman, Morsink}. The damping timescale, which depends sensitively on the viscosity, and hence on the low-energy modes of the color superconductor determines if the $r$-mode will be driven unstable. The unsuppressed growth of the $r$-mode eventually saturates due to non-linear effects~\\cite{Arras03} since energy can be transferred from $r$-modes to ``inertial'' oscillation modes which do not couple to gravitational waves. However, the saturation amplitude can still be large enough to emit gravitational wave signals that can be detected directly by Advanced LIGO~\\cite{AK}.   The observational context of $r$-modes is that neutron stars do not typically spin at rates near the maximum allowed frequency, the Kepler frequency $\\Omega_K$. $r$-modes offer a possible explanation of this fact: in rotating neutron stars, these modes lose energy through gravitational waves, which carry away angular momentum from the star and act as braking radiation. We mention two distinct settings in which $r$-modes can be instrumental in slowing the neutron star's rotation speed.   (a) In the newly-formed isolated neutron star in the aftermath of a core-collapse supernova: Whether $r$-modes are relevant for neutron stars  at birth is difficult to determine because the number of rapidly spinning pulsars that can be detected close to their birth is a  negligible fraction of the total detectable sample. Limited by this selection effect, most  inferences of rotation rates of neutron stars rely on a model of time evolution of radio pulsars from birth (neglecting any contribution from $r$-modes), with conflicting results on the birth spin-periods, ranging from milliseconds~\\cite{ARZ} to hundreds of milliseconds~\\cite{Kaspi}.  An independent method of determining the initial spin period based on correlating a supernova's observed X-ray luminosity (or upper limits thereof) with the rotational energy loss of the neutron star suggests that a population of millisecond pulsars at birth is ruled out~\\cite{Stella}. On the other hand, simulations of Type II supernovae including rotational effects predict relatively fast initial spin periods~\\cite{Ott06}, which is difficult to reconcile with observed young and slowly rotating pulsars (like the Vela and Crab pulsars).  Thus, $r$-modes may play an important role in neutron stars at birth, and the impact of $r$-modes depends critically on the composition of the neutron star core. The critical temperature for color superconductivity being of the order of a proto-neutron star's temperature, it is important to understand the $r$-mode in the quark superfluid as it can determine the initial spin-down evolution of the strange star or hybrid star.   (b) In old neutron stars spun-up by accretion in binary systems: Millisecond pulsars (typically, pulsars with spin-periods $\\sim 10$~ms) are old and most of them are observed to be in low-mass X-ray binary systems (LMXBs)~\\cite{Lorimer}. The fast rotation is thought to result from spin-up due to steady accretion from a low-mass companion star. This picture has recently received confirmation through the observation of pulsations in LMXBs. Spin frequencies in LMXBs, as determined either from X-ray burst oscillations or from kHz quasi-periodic oscillations, are typically between 300-640 Hz, well below the maximum rotational rate (see Ref.~\\cite{Lamb} for a recent review). The upper limit on the rotation rate of LMXBs has been recently revised by the detections of objects spinning at 716 Hz ~\\cite{Hessels06} and 1122 Hz~\\cite{Kaaret07}. However, the statistical significance of the latter observation is not very strong. One possibility is that accretion generates a mass quadrupole moment which results in gravitational waves and this prevents the star from rotating too quickly~\\cite{Lars}. Upper limits on the generation of gravitational waves from this mechanism have already been generated by LIGO~\\cite{Ben}. $r$-mode oscillations can be an alternative explanation for the limiting frequency. The observation of these fast-rotating objects provides an important constraint on the extent to which $r$-mode oscillations can spin down neutron stars.  Since inferred temperatures of an LMXB's interior are much less than $T_c$ for a quark superfluid, it is pertinent to ask how quark matter inside an accreting neutron star affects the $r$-mode and the star's rotation speed, assuming the quark core is not spun out of existence.   Motivated by these reasons, in this article, we investigate the effect of  quark matter on $r$-modes in compact stars. In section~\\ref{rmode}, we calculate the  $r$-mode frequency for perturbed rotating fluids that obey chosen nuclear or quark matter EoS. In section~\\ref{viscosity}, we review the viscosities of neutron matter, ungapped and gapped CFL quark matter, which are inputs in a calculation of the damping timescale of the $r$-mode. In section~\\ref{damping}, we present results for $r$-mode damping times in strange and hybrid stars, and discuss the critical spin-frequency in section ~\\ref{critical}. In section~\\ref{summary},  we remark on observational consequences for such compact stars and present our conclusions. ",
        "conclusions": "\\label{summary}   We have calculated the frequency and viscous damping timescale of the $m=2$ $r$-mode for compact stars made of hadronic matter, quark matter, and for hybrid stars. To our knowledge, this is the first attempt at studying $r$-mode characteristics for pure or hybrid stars with color superconducting quark matter,  particularly the CFL phase. In absolute terms, the $r$-mode frequency is about 50\\% larger for strange stars in comparison to neutron stars of the same mass with a polytropic EoS, assuming both are rotating at their respective Kepler frequency. The ratio of the $r$-mode frequency to the star's rotation frequency is larger for strange stars and hybrid stars; this is due to the softer equation of state.  It is also evident that EoS effects on the mode frequency are smaller for softer equations of state.   The dominant bulk and shear viscosities of the normal hadronic phase, the ungapped quark phase as well as the gapped CFL quark phase, which was derived recently, are used to estimate viscous damping timescales for homogeneous and hybrid stars. The current results confirm established results for neutron stars; viz. that since neutron stars are stable against the $r$-mode instability only at very high ($T\\gtrsim$ 10$^{10}$ K) or very low temperatures ($T\\lesssim$ 10$^{6}$ K), they would be spun down rapidly by the $r$-mode instability shortly after their birth at MeV temperatures, when large neutrino losses cause rapid cooling. The predicted critical frequency for LMXBs with hadronic neutron stars is close to  or inside the region given by observations and LMXBs. A more careful analysis taking the effect of superfluidity in neutron matter into account would provide a clearer picture on the presence of $r$-modes assuming no quark matter is present.   Strange stars with non-superfluid quark matter display a stability window between $10^{8}$~K~$<T<5\\times 10^{9}$~K where they can spin at a substantial fraction of the Kepler frequency. If LMXBs contain strange stars, then on statistical grounds, we should expect to observe some of them spinning close to their Kepler frequency, unless some mechanism not connected to $r$-modes is responsible for limiting their frequency or the quark phase is more complicated than considered here (ungapped or CFL). If newly-born superconducting strange stars occur in nature, they may still spin down rapidly at birth, provided the cooling timescale is slower than the spin-down timescale. A consistent analysis of spin-down and cooling is currently under investigation and will be reported in separate work. Hybrid stars tend to be more resilient against the $r$-mode instability, except for a temperature range which is dependent on the quark EoS. This effect is not surprising: in hybrid stars, $r$-modes which would normally be undamped in hadronic matter are damped by the quark core and vice versa. We conclude that hybrid proto-neutron stars are not likely to be damped by $r$-modes at all. This work also has important implications for LMXBs if they are hybrid stars containing a significant amount of deconfined quark matter in the ungapped or CFL phase, or if LMXBs are strange stars with a thin accreted hadronic crust. Based on the estimates of viscous damping timescales, the presence of quark matter in LMXBs implies that either they must contain rapidly rotating stars that have not yet been observed, or have a more complicated phase structure. A similar explanation regarding the LMXB rotating at 1122 Hz was put forth recently~\\cite{Parenti07}.   To make the results more concrete, future work will have to address several assumptions more closely. We assumed that the modified Urca process provides the dominant contribution to bulk viscosity while  neutron-neutron scattering determines the shear viscosity. In practice, in the core of neutron stars, for $T\\leq T_c\\sim 10^8$ K, neutron matter is  likely to be in a triplet-paired superfluid state. This will suppress the  modified Urca rate~\\cite{Haensel:2001}, rendering it insignificant at $T\\ll T_c$.  Singlet pairing among neutrons and/or protons has a larger critical temperature $T_c\\sim 5\\times 10^{9}$ K~\\cite{Yakovlev01} and can occur in the  crust with similar consequences. Under these conditions, it is more realistic to include alternate sources of viscosity along the lines of Ref.~\\cite{Alpar}. It would also be useful to explicitly check the contributions from  the additional viscosities $\\zeta_1$ and $\\zeta_3$ that are inherent in a 2-fluid model of superfluidity.  We also assume that in the CFL phase, the shear viscosity is from phonons alone, whereas below $T\\sim 10^{10}$ K, the damping timescales are not reliable as the phonon mean free path exceeds the stellar size. At these low temperatures, other contributions such as from photons would be more relevant and need to be explicitly computed and included for an accurate treatment of $r$-modes in CFL matter.  In the current work, we considered a simplified  model for a neutron/strange star in this work, ignoring realistic features  such as a crustal layer, rubbing friction at the quark-hadron interface and possibly the role of other contributions to the bulk and shear viscosity at very high or very low temperatures. Neutron star crusts composed of hadronic matter typically damp $r$-mode oscillations further. Early work in Ref.~\\cite{Bildsten00} suggested that damping timescales were decreased by as much as a factor of $10^5-10^7$ because of the additional  viscosity provided by the neutron star crust. This conclusion has been softened in some later works: Ref.~\\cite{Levin01} concluded that the associated decrease in the damping timescale was a factor of $10^2-10^3$ smaller, and Ref.~\\cite{Glampedakis06} has pointed out that the damping timescale is sensitive to the crustal composition. The final result may well be that $r$-modes are not as important a mechanism as believed for  spinning down fast-rotating compact objects. Ref.~\\cite{Lin07} has also stressed the importance of updating the multi-component nature of equation of state in order to fully assess the effects of superfluidity. However, this does not change the central conclusion, viz.,  if $r$-modes do play a role in spinning down neutron stars, the present work demonstrates that this effect is much less efficient in neutron stars containing ungapped or completely gapped quark matter."
    },
    "0806/0806.0415_arXiv.txt": {
        "abstract": "Continuous observations of sunspot penumbrae with the Solar Optical Telescope aboard \\textit{Hinode} clearly show that the outer boundary of the penumbra fluctuates around its averaged position. The penumbral outer boundary moves inward when granules appear in the outer penumbra. We discover that such granules appear one after another while moving magnetic features (MMFs) are separating from the penumbral ``spines'' (penumbral features that have stronger and more vertical fields than those of their surroundings). These granules that appear in the outer penumbra often merge with bright features inside the penumbra that move with the spines as they elongate toward the moat region. This suggests that convective motions around the penumbral outer boundary are related to the disintegration of magnetic flux in the sunspot. We also find that dark penumbral filaments frequently elongate into the moat region in the vicinity of MMFs that detach from penumbral spines. Such elongating dark penumbral filaments correspond to nearly horizontal fields extending from the penumbra. Pairs of MMFs with positive and negative polarities are sometimes observed along the elongating dark penumbral filaments. This strongly supports the notion that such elongating dark penumbral filaments have magnetic fields with a ``sea serpent''-like structure. Evershed flows, which are associated with the penumbral horizontal fields, may be related to the detachment of the MMFs from the penumbral spines, as well as to the formation of the MMFs along the dark penumbral filaments that elongate into the moat region. ",
        "introduction": "Sunspot penumbrae consist of many dark and bright radial filaments. The inclination of penumbral magnetic fields fluctuates with azimuthal angle around the spot center \\citep{Degenhardt1991, Title1993, Lites1993, Stanchfield1997, Westendorp2001a, Westendorp2001b, Bello2005, Langhans2005}. The bright penumbral filaments have magnetic fields that are more vertical than those of the dark penumbral filaments in the outer penumbra. The vertical and horizontal components of the penumbral fields have been called a ``penumbral spine'' and a ``penumbral intra-spine,'' respectively \\citep{Lites1993}. Radial flows and radial motions are also observed in the penumbra. The Evershed flow is a horizontal outward flow along the horizontal component of the penumbral magnetic fields \\citep{Degenhardt1991, Title1993, Lites1993, Rimmele1995, Stanchfield1997, Westendorp2001a, Westendorp2001b, Bellot2003, Bellot2004, Borrero2004, Borrero2005, Cabrera2006, Cabrera2007, Cabrera2008}. Evershed flows with supersonic velocities are sometimes inferred around the outer boundary of the penumbra \\citep{del2001}, and are observed directly \\citep{Shimizu2008b}. Penumbral grains, which are patchy bright features, move toward the umbra in the inner penumbra, and the penumbral grains in the outer penumbra show both inward and outward motions \\citep[e.g.][]{Sobotka1999, Sobotka2001, Bovelet2003}. The Evershed flow appears to originate at the inward-moving penumbral grains in the inner penumbra \\citep{Rimmele2006, Ichimoto2007}. Various models have been proposed to explain the complex magnetic and velocity fields of penumbrae \\citep[e.g.,][]{Solanki1993, Schlichenmaier1998, Thomas2002, Spruit2006, Heinemann2007}.  Radial outward flows are dominant in the moat region that surrounds mature sunspots. Moving magnetic features (MMFs) are small magnetic elements that have sizes of typically less than 2$\\arcsec$ in the moat region \\citep{Sheeley1969, Vrabec1971, Harvey1973}. MMFs mostly appear around the penumbral outer boundary and then move almost radially outward in the moat region for mature and decaying sunspots \\citep{Harvey1973, Ryutova1998, Zhang2003, Sainz2005, Ravindra2006}. The lifetimes of MMFs range from a few minutes to 10 hr \\citep{Vrabec1974, Zhang2003, Hagenarr2005}. Those MMFs with longer lifetimes may reach the outer boundary of the moat region. Recent observations suggest that MMFs located on lines extrapolated outward from the horizontal component of the penumbral magnetic fields into the moat region correspond to the photospheric intersections of undulating horizontal fields extending from the penumbra \\citep{Sainz2005, Cabrera2006, Cabrera2007, Kubo2007a}. This supports the ``sea serpent'' concept for MMFs, whereby the horizontal penumbral fields undulate into and out of the solar surface \\citep{Harvey1973, Schlichenmaier2002}. Many authors have proposed that the fluctuating penumbral fields and the subsequent MMFs are produced by Evershed flows \\citep[e.g.,][]{Ryutova1998, Martinez2002, Thomas2002, Schlichenmaier2002, Zhang2003, Zhang2007b, Cabrera2006}. On the other hand, MMFs located on the lines extrapolated from the penumbral spines have magnetic fields that are more vertical than those of their surroundings \\citep{Kubo2007a}. Such MMFs appear to be formed as a result of the detachment of a fraction of the magnetic field from the penumbral spines.  The net magnetic flux transported by all the MMFs exceeds the flux-loss rate of the sunspot \\citep{Martinez2002, Kubo2007a}. Hence, it has been suggested that the MMFs detached from the penumbral spines alone are the agent that removes magnetic flux from the sunspot \\citep[e.g.,][]{Shine2001, Martinez2002, Thomas2002, Weiss2004}. Those MMFs might carry away an amount of magnetic flux sufficient to account for the flux loss of the sunspot \\citep{Kubo2007a}. If so, the formation of those MMFs that detach from the penumbral spines is more important to the process of sunspot decay than that of other types of MMFs.  \\citet{Kubo2007b} found that bright features (granules) appear where an MMF is separating from the penumbral spine, but only one event was examined in detail in that work. Using observations of two mature sunspots made by the Solar Optical Telescope (SOT; Tsuneta et al. 2008; Suematsu et al. 2008; T. D. Tarbell et al. 2008 in preparation; Ichimoto et al. 2008; Shimizu et al. 2008a) aboard the \\textit{Hinode} satellite \\citep{Kosugi2007}, we build statistics of the comparison between MMFs detaching from the spines and both bright and dark features around the penumbral outer boundary. The \\textit{Hinode} SOT reveals the evolution of fine structures in both the penumbra and the moat region from continuous observations with a spatial resolution of $0.2\\arcsec-0.3\\arcsec$. ",
        "conclusions": "The SOT allows us to perform continuous observations with high spatial resolution for a duration longer than the typical lifetime of MMFs. The results regarding flux removal from the sunspot in this study are summarized as follows (see also Fig.~\\ref{summary_fig}):  \\begin{enumerate} \\item G-band bright features in the outer penumbra are located at the elongating spines or at the MMFs detaching from the spines. G-band bright features moving both outward and inward can be found in the middle and outer penumbra. The penumbral spines elongate toward the moat region in concert with the outward motions of the penumbral bright features. \\item Granules appear in the outer penumbra at locations where MMFs that have the same polarity as the sunspot separate from the penumbral spines. Such granules often merge with bright features moving along the elongating spines. \\item Dark penumbral filaments frequently elongate into the moat region, and longer ones reach the outer moat region. The elongating dark penumbral filaments coincide with filamentary structures that have strong horizontal fields that extend outward from the penumbra. We present clear examples of bipolar MMFs indicating a serpentine field with multiple intersections with the solar photosphere. MMFs with positive and negative polarities are located along the filamentary structures with strong horizontal fields. \\item The MMFs that are detached from the penumbral spines are usually observed in the neighborhood of the elongating dark penumbral filaments. \\end{enumerate}  The total flux transport rate by the MMFs that are detached from the penumbral spines is greater than the flux-loss rate of the sunspot \\citep{Kubo2007a}. Therefore, items 1 and 2 above confirm that convection in the outer penumbra is related to the flux removal of the sunspot \\citep{Kubo2007b}. However, the magnetic energy density of the penumbral fields is larger than the kinetic energy density of the penumbral bright features at the photospheric surface. Converging and downward flows are inferred in the subsurface convection zone around a stable sunspot \\citep{Zhao2001}, and such flows might be essential to the formation and sustenance of sunspots \\citep{Meyer1974, Parker1979, Parker1992}. The appearance of the bright features moving outward in the outer penumbra suggests subsurface upwelling and diverging flows. Such flows would work against the stabilization of the sunspot, and recent MHD simulations showed that magnetic flux is carried away from the spot by subsurface diverging flows \\citep{Heinemann2007}. \\citet{Meyer1974} suggested that small flux tubes diffuse out into the moat region by small scale convection beneath the sunspot. The fact that the bright features and granules appear in the outer penumbra at the location of the elongating spines (that subsequently become MMFs) supports this idea.  Another possibility is that Evershed flows detach penumbral flux and advect it into the moat region. Evershed flows with supersonic velocities are often observed around the penumbral outer boundary \\citep{Shimizu2008b}. Supersonic flows are also observed in connection with horizontal fields extending from the penumbra \\citep{Kubo2007b}. The kinetic energy density of such supersonic flows is larger than the magnetic energy density of the magnetic fields in the outer penumbra. Nevertheless, the \\textit{Hinode} SP observations of vector magnetic fields with 0.3$\\arcsec$ resolution have confirmed that the Evershed flows are radial outward flows along those penumbral filaments that have nearly horizontal fields \\citep{Ichimoto2007}. Detachment of the MMFs is observed when the penumbral horizontal fields elongate into the moat region. This means that Evershed flows are related not only to the formation of the bipolar MMFs located along the extending penumbral horizontal fields, but also to the detachment of the MMFs from the penumbral spines. The Evershed flows originate at penumbral bright features that are moving inward in the inner penumbra \\citep{Rimmele2006, Ichimoto2007}, and the detachment of MMFs from the penumbra is observed with penumbral bright features that are moving outward in the outer penumbra. This suggests the possibility that both Evershed flows and the detachment of MMFs originate from the activity of convection below the sunspot penumbra. The relationship between Evershed flows and sunspot decay also should be investigated using Doppler and magnetic field measurements of both high spatial and high temporal resolution."
    },
    "0806/0806.2410_arXiv.txt": {
        "abstract": "{ We present initial results of an ESO-VLT large programme (AMAZE) aimed at determining the evolution of the mass-metallicity relation at z$>$3 by means of deep near-IR spectroscopy. Gas metallicities are measured, for an initial sample of nine star forming galaxies at z$\\sim$3.5, by means of optical nebular lines redshifted into the near-IR. Stellar masses are accurately determined by using Spitzer-IRAC data, which sample the rest-frame near-IR stellar light in these distant galaxies. When compared with previous surveys, the mass-metallicity relation inferred at z$\\sim$3.5 shows an evolution much stronger than observed at lower redshifts. The evolution is prominent even in massive galaxies, indicating that z$\\sim$3 is an epoch of major action in terms of star formation and metal enrichment also for massive systems. There are also indications that the metallicity evolution of low mass galaxies is stronger relative to high mass systems, an effect which can be considered the chemical version of the galaxy downsizing. The mass-metallicity relation observed at z$\\sim$3.5 is difficult to reconcile with the predictions of some hierarchical evolutionary models. Such discrepancies suggest that at z$>$3 galaxies are assembled mostly with relatively un-evolved sub-units, i.e. small galaxies with low star formation efficiency. The bulk of the star formation and metallicity evolution probably occurs once small galaxies are already assembled into bigger systems. } ",
        "introduction": "\\label{sec_intro}  The connection between galaxy mass and metallicity has been known for a long time, starting with the seminal work of \\cite{lequeux79}. Given the difficulty in obtaining reliable galaxy masses several authors have resorted in using the (optical) luminosity. In particular, various works have reported a clear correlation between blue luminosity and metallicity, in the sense that more luminous galaxies are characterized by higher metallicities \\citep[e.g.][]{garnett87,skillman89,brodie91,zaritsky94}. A major step forward has been recently achieved by \\cite{tremonti04}, who used optical photometric and spectroscopic data of a sample of $\\sim$53,000 galaxies from the SDSS to determine the mass-metallicity relation of local galaxies (z$\\sim$0.1). Their work clearly showed that the primary physical parameter driving the correlation with the gas metallicity is the (stellar) mass of galaxies and not their luminosity. While \\cite{tremonti04} focused on the gas metallicity, a similar relation by using the SDSS survey was found by \\cite{gallazzi06} for what concerns the stellar metallicity.  Various physical processes may be responsible for the mass-metallicity relation. One possibility is that outflows, generated by starburst winds, eject metal-enriched gas into the IGM preferentially out of low-mass galaxies (due to the shallow gravitational potential well), making their enrichment less effective than in massive systems \\citep[e.g.][]{tremonti04, delucia04,finlator08}. An alternative scenario is that low mass systems are still at an early evolutionary stage and have still to convert most of their gas into stars, hence they are poorly metal-enriched relative massive galaxies (which are instead already evolved). This is the so-called ``galaxy downsizing'' scenario, supported by various observational evidences \\citep[e.g.][]{juneau05,feulner05,franceschini06,asari07,perez07}, where massive galaxies formed most of their stars rapidly and at high redshift, while low mass systems are characterized by a slower evolution, which extends to low redshift. Finally \\cite{koppen07} ascribes the mass-metallicity relation to variations of the IMF high-mass cutoff in different star forming environments.  The relative role of these processes in shaping the mass-metallicity relation is debated. It is likely that each of them contributes at least to some extent, since observational evidences have been found for all of them. Each of these factors (outflows/feedback, downsizing, IMF) has profound implications on the evolution of galaxies. Therefore, it is clear that the mass-metallicity relation contains a wealth of information useful to constrain models of galaxy formation and evolution. Indeed, any model of galaxy evolution is now required to match the mass--metallicity relation observed locally \\citep[e.g.][]{kobayashi07,brooks07,derossi07,dave07, dalcanton07,delucia04,tissera05,bouche06,bouche07,koppen07,cidfernandes07,finlator08,tassis08}. However, different models predict different evolutionary patterns of the mass-metallicity relation as a function of redshift, and observational data are required to test and discriminate among them.  Observational constraints of the mass-metallicity relation have been obtained up to z$\\sim$2.2 thanks to various deep surveys \\citep{savaglio05,liang06,erb06}. Additional observational studies have investigated the evolution of the luminosity-metallicity relation or, more generally, the metallicity of high-z star forming galaxies \\citep{kobulnicky00,kobulnicky03,kobulnicky04,maier04,maier05,maier06,forster06}. Note that all of these studies refer to the gas metallicity, while limited work has been done on the stellar metallicity of high redshift sources \\citep{rix04,demello04,halliday08,mehlert06} due to difficulties in obtaining high S/N spectra on the stellar continuum. The general observational result is that the mass-metallicity relation (as well as the luminosity-metallicity relation) evolves, in the sense that on average higher redshift galaxies are characterized by lower metallicities (at a given mass). Whether the relation evolves also in terms of its shape is still matter of debate. Theoretical models can generally cope with the observed evolution, within both the theoretical and observational uncertainties \\citep[e.g.][]{kobayashi07,derossi07,brooks07, finlator08}.  At z$\\ge$3, except for a few individual studies, little work has been currently done for what concerns the mass-metallicity relation. \\cite{pettini01} have measured the metallicity for a small sample of Lyman Break Galaxies (LBG) at z$\\sim$3, but without investigating the mass-metallicity relation. The metallicity evolution has been investigated in DLA systems \\citep{prochaska03,kulkarni05,akerman05}. However, a study of the mass-metallicity relation for high-z absorption systems has not been properly performed, due to difficulties in measuring the associated stellar masses. The closest result is the finding of a relation between metallicity and velocity dispersion (probably related to the mass) inferred by the width of the absorption systems \\citep{ledoux06,prochaska07}.  We have undertaken a large observing programme at ESO-VLT aimed at carefully determining the mass-metallicity relation at z$>$3 for a sizeable sample of galaxies. The final goal is to obtain a detailed description of the evolution of the mass-metallicity evolution through the cosmic epochs, and therefore constrain galaxy evolutionary scenarios. In this paper we present preliminary results obtained by such a program, and discuss the relevant implications for our understanding of the galaxy evolution at high redshift. Throughout the paper we adopt the following cosmological parameters: $\\rm H_0=71~km~s^{-1}~Mpc^{-1}$, $\\rm \\Omega _{\\Lambda}=0.73$ and $\\rm \\Omega _m=0.27$ \\citep{spergel03}. ",
        "conclusions": "We have presented initial results of the AMAZE project, an ESO large programme aimed at determining the mass-metallicity relation of star forming galaxies at z$>$3. Near-IR spectra are being obtained with SINFONI, the VLT near-IR integral field spectrometer, for a sample of Lyman Break Galaxies at 3$<$z$<$5.  Gas metallicities are inferred by using a combination of diagnostics involving nebular lines observable in the H and K bands. To have a metallicity scale consistent with the results obtained by previous surveys at lower redshifts, we derived new accurate calibrations of various strong-line metallicity diagnostics spanning the wide range $\\rm 7.1<12+log(O/H)<9.1$.  AGNs (which would affect and make unusable the metallicity diagnostics) were carefully removed through a multiwavelength approach using X-ray, optical and mid-IR data. In particular, the use of mid-IR (Spitzer-MIPS) data allows us to discard even heavily obscured, Compton thick AGNs.  Stellar masses are inferred by fitting multi-band photometric data with galaxy templates. Within this context crucial is the use of Spitzer-IRAC data, which sample the rest-frame near-IR stellar light at 3$<$z$<$5.  In this paper we have presented results from an initial sample of 9 LBGs at 3.1$<$z$<$3.7. Emission lines required to constrain the gas metallicity are detected in all sources. From the mass-metallicity relation at z$\\sim$3.5 inferred from this initial sample we obtain the following results:  \\begin{enumerate}  \\item When compared with the results of previous surveys at lower redshift (recalibrated to match our metallicity and mass calibrations), we obtain evidence for a clear evolution of the mass-metallicity relation as a function of redshift. Since at different redshifts we are sampling different populations of galaxies, the observed evolution should {\\it not} be considered the evolutionary pattern of individual galaxies, but as the evolution of the mass-metallicity relation of the dominant population of star forming galaxies at each epoch. The evolution of the mass-metallicity relation is faster at 2.2$<$z$<$3.5 than at later epochs (z$<$2.2). This result indicates that z$\\sim$3 is an epoch of major action for the enrichment of galaxies, even in high mass systems  \\item There are indications that the metallicity evolution is not constant with mass: low mass galaxies evolve more strongly than massive systems. If confirmed, this result can be considered as the chemical version of the galaxy downsizing, i.e. a scenario where high mass galaxies reach high metallicities at high redshift, on short timescales, while low mass systems enrich their ISM over a prolonged period of time, extending to the current epoch. However, this finding is subject to uncertainties due to low statistics at low masses, and must await additional data to be confirmed.  \\item The observed mass-metallicity relation at z$\\sim$3 is difficult to reconcile with the simulations by some hierarchical models, which predict metallicities higher than observed. If aperture effects are taken into account then the discrepancy is probably even higher. The main problem of these models seems to be that galaxies are assembled once they are already evolved, from the stellar and chemical point of view. Our results suggest that galaxies at z$>$3 are assembled mostly with relatively un-evolved small galaxies, whose star formation efficiency is low. Most of the chemical evolution (hence most of the star formation) must occurs once small galaxies are already assembled into bigger ones. This implies that most of the merging occurs before most of the star formation. Indeed, models and simulations where a strong feedback keeps star formation low in the assembling small galaxies (preventing them to evolve strongly from the stellar and chemical point of view), provide a much better description of the mass-metallicity relation at z$\\sim$3.5.  The finding that galaxies at z$>$3 are mostly assembled with un-evolved sub-units is not necessarily in contrast with models of ``dry-merging'', i.e. models where galaxy assembly occurs through systems that are significantly evolved and with little residual gas. Indeed, dry merging may be the main mode of galaxy evolution at lower redshifts (z$<$3) without being in conflict with our findings at z$>$3.  \\end{enumerate}"
    },
    "0806/0806.0623_arXiv.txt": {
        "abstract": "s{We study a model where two scalar fields, that are subdominant during inflation, decay into radiation some time after inflation has ended but before primordial nucleosynthesis. Perturbations of these two curvaton fields can be responsible for the primordial curvature perturbation. We write down the full non-linear equations that relate the primordial perturbation to the curvaton perturbations on large scales, and solve them in a sudden-decay approximation. We calculate the power spectrum of the primordial perturbation, and finally go to second order to find the non-linearity parameter, $\\fNL$. {\\em Not surprisingly}, we find large positive values of $\\fNL$ if the energy densities of the curvatons are sub-dominant when they decay, as in the single curvaton case. But we also find {\\em a novel effect}, which can be present only in multi-curvaton models: {\\em $\\fNL$ becomes large even if the curvatons dominate the total energy density} in the case when the inhomogeneous radiation produced by the first curvaton decay is diluted by the decay of a second nearly homogeneous curvaton. The minimum value $\\min(\\fNL)=-5/4$ which we find is the same as in the single-curvaton case. Using (non-)Gaussianity observations, Planck can be able to distinguish between single-field inflation and curvaton model. Hence it is important to derive theoretical predictions for curvaton model. From particle physics point of view it is more natural to assume multiple scalar fields (rather than just one ``curvaton'' in addition to inflaton). Our work updates the theoretical predictions of curvaton model to this case. } ",
        "introduction": "Theories beyond the standard model often contain a large number of scalar fields in addition to the standard-model fields. In the very early universe it is natural to expect the initial values of these fields to be displaced from the minimum of their potential. If they are displaced by more than the Planck scale, then they can drive a period of inflation. Otherwise they will oscillate about the minimum of their potential once the Hubble rate, $H$, drops below their effective mass. An oscillating massive field has the equation of state (averaged over several oscillations) of a pressureless fluid. Thus the energy density of a weakly interacting massive field tends to grow relative to radiation in the early universe. Such fields must therefore decay before the primordial nucleosynthesis era to avoid spoiling the standard, successful hot big bang model. If the energy density of a late-decaying scalar is non-negligible when it decays, then any inhomogeneity in its energy density will be transfered to the primordial radiation.\\cite{Mollerach:1989hu} % Therefore, the observed perturbations in the Universe can solely result from these inhomogeneities and the perturbations in the inflaton decay products (i.e., pre-existing radiation) can be negligible. This is the curvaton scenario for the origin of structure.\\cite{Enqvist:2001zp} % We focus on this {\\em pure} curvaton model.\\footnote{{\\em Before their decay}, the curvatons represent pure isocurvature degrees of freedom which are ruled out by data.\\cite{Enqvist:2000hp} Nevertheless, the pure curvaton model generates {\\em adiabatic primordial density perturbations}, if all the species are in thermal equilibrium and the baryon asymmetry is generated after the curvatons decay. However, if inflaton perturbation was not fully negligible and the curvaton(s) decayed into radiation and cold dark mater (CDM), then a correlated mixture of adiabatic and CDM isocurvature primordial perturbations would follow. These possibilities have been studied extensively, e.g., in Refs.~\\cite{Malik:2002jb}.   % It should be noted that an admixture of subdominant isocurvature (up to 5-10\\%) and predominant adiabatic primordial perturbation can fit well~\\cite{KurkiSuonio:2004mn} the current data. } ",
        "conclusions": ""
    },
    "0806/0806.3766_arXiv.txt": {
        "abstract": "As part of a multiwavelength study of the unusual radio supernova remnant DA 495, we present observations made with the {\\em Chandra X-ray Observatory}. Imaging and spectroscopic analysis confirms the previously detected X-ray source at the heart of the annular radio nebula, establishing the radiative properties of two key emission components: a soft unresolved source with a blackbody temperature of 1 MK consistent with a neutron star, surrounded by a nonthermal nebula 40$''$ in diameter exhibiting a power-law spectrum with photon index $\\Gamma = 1.6\\pm0.3$, typical of a pulsar wind nebula. The implied spin-down luminosity of the neutron star, assuming a conversion efficiency to nebular flux appropriate to Vela-like pulsars, is $\\sim 10^{35}$ ergs~s$^{-1}$, again typical of objects a few tens of kyr old. Morphologically, the nebular flux is slightly enhanced along a direction, in projection on the sky, independently demonstrated to be of significance in radio polarization observations; we argue that this represents the orientation of the pulsar spin axis. At smaller scales, a narrow X-ray feature is seen extending out 5$''$ from the point source, a distance consistent with the sizes of resolved wind termination shocks around many Vela-like pulsars. Finally, we argue based on synchrotron lifetimes in the estimated nebular magnetic field that DA 495 represents a rare pulsar wind nebula in which electromagnetic flux makes up a significant part, together with particle flux, of the neutron star's wind, and that this high magnetization factor may account for the nebula's low luminosity. ",
        "introduction": "With the {\\em Chandra X-ray Observatory}'s first glimpse of the Crab pulsar \\citep{2000ApJ...536L..81W}, a new window was opened onto the interactions of neutron stars with their surroundings. Rotation-powered pulsars drive outflows of particles and electromagnetic energy into the immediate environment; depending primarily on the pulsar's age, this environment might simply be the interstellar medium or, in the case of young, energetic pulsars, the circumstellar debris of the supernova explosion that accompanied the birth of the neutron star. The astrophysical value of these interactions lies in their uniqueness as a probe of the content of pulsar winds, which originate in the incompletely understood magnetospheres of neutron stars, and of the physics of particle acceleration and shock mixing of relativistic and non-relativistic fluids. In contrast to the optical and radio observations that provided early insights into these phenomena, X-ray studies, in particular with {\\em Chandra}'s unequalled angular resolution, sample the properties of energetic and short-lived particles nearest the accelerating engines that produce them.  In the class of filled-center supernova remnants (SNRs), the bulk of the radiation from a remnant can be traced back energetically to the rapidly rotating neutron star at its core, with emission from the shocked neutron star wind often defining a pulsar wind nebula (PWN). Shell-type remnants, by contrast, expend the energy of the supernova explosion's blast wave. Evidence for the transfer of a neutron star's rotational kinetic energy to its surroundings is provided by the axisymmetric structures---arcs, rings, tori, and jets---seen in high-resolution X-ray images of the environs of a number of pulsars. Together with amorphous nonthermal emission also found around energetic pulsars, these structures are brought to light through synchrotron radiation by the pulsar's ultrarelativistic particle wind where it is confined and perturbed by the surrounding medium; see, e.g., \\citet{2006ARA&A..44...17G} for a review of PWNs. The morphological and spectral properties of PWNs carry valuable information about the physical mechanisms of energy transfer, the local conditions (particle densities, flow velocities, magnetic field strengths, etc.), and the nature of the particle-accelerating central engine.  Although its identification as a pulsar wind nebula is now virtually certain, the DA 495 radio nebula (G65.7+1.2; \\citealp{1983AJ.....88.1810L}) is highly unusual: its morphology is distinctly annular, with a radial profile that rises rapidly from a central emission deficit (the ``hole'') then falls gradually with increasing distance from the center; see Figure~\\ref{fig:radio} for a contour map, and \\citet{klr+07} for a detailed discussion of the radio properties. We described in an earlier paper \\citep[hereafter, ASLK04]{2004ApJ...610L.101A} X-ray observations of DA 495 that revealed the central engine apparently powering the radio nebula, the {\\em ROSAT} source 1WGA J1952.2+2925, and provided some evidence for extended X-ray emission consistent with a PWN. To confirm the existence of a high-energy nebula, and because lingering doubts remained about a possible ``thick shell'' thermal origin for the radio annulus \\citep{1989JApA...10..161V}, we set out to perform with \\cxo\\ an X-ray observation sensitive to compact nonthermal structures as a means of strengthening the identification of DA 495 as a PWN, through the detection in particular of jets, toroidal structures, or other markers of large-scale axial symmetry, such as those seen around the Crab, Vela, and similar pulsars. ",
        "conclusions": "The soft spectrum of the unresolved source near the center of the DA 495 radio nebula, its implied blackbody temperature and luminosity, and its lack of variability on timescales of a few seconds through the 25 ks duration of the observation are highly suggestive of emission from the surface of an isolated neutron star. Taken together with the properties of the surrounding extended, nonthermal emission, no doubt remains that DA 495 is a pulsar wind nebula, driven by the rotation-powered magnetosphere of a neutron star. Its unusual radio characteristics can be explained as the result of aging in the absence of any evident interaction with the supernova explosion's reverse shock, such that DA 495 presents a valuable example of a pristine PWN at an advanced age."
    },
    "0806/0806.3263.txt": {
        "abstract": "% The statistics of catalogued quadruple stars consisting of two  binaries (hierarchy 2+2)  is studied in comparison  with triple stars,   with  respective  sample   sizes  of   81  and   724.   Seven representative quadruple systems are  discussed in greater detail. The main conclusions are: % (i) Quadruple  systems of $\\varepsilon$~Lyr type with similar masses and inner  periods are common, in 42\\% of  the sample the outer mass ratio is  above 0.5 and the inner periods differ  by less than 10 times. % (ii) The distributions of the inner periods in triple and quadruple stars are similar and bimodal. The inner mass ratios do not correlate with the inner periods. % (iii)  The statistics of  outer periods  and  mass ratios  in triples and quadruples  are different. The median outer  mass ratio in triples is 0.39  independently of the outer period,  which has a smooth distribution.  In  contrast,  the  outer  periods  of  25\\%  quadruples concentrate in the narrow range from 10\\,yr to 100\\,yr, the outer mass ratios of these tight quadruples are above 0.6 and their two inner periods are similar to each other. % (iv) The outer and inner mass ratios in triple and quadruple stars are not mutually correlated. In 13\\% of quadruples both inner mass ratios are above 0.85 (double twins). % (v) The inner and outer  orbital angular momenta and periods in triple and  quadruple  systems  with  inner  periods above  30\\,d  show  some correlation, the  ratio of outer-to-inner periods  is mostly comprised between  5 and $10^4$.  In the  systems with  small period  ratios the directions of the  orbital spins are correlated, while  in the systems with large ratios they are not. % The properties of multiple stars  do not correspond to the products of dynamical decay  of small clusters,  hence the N-body dynamics  is not the  dominant  process  of   their  formation.   On  the  other  hand, rotationally-driven  (cascade)   fragmentation  possibly  followed  by migration  of  inner and/or  outer  orbits  to  shorter periods  is  a promising  scenario to  explain  the origin  of  triple and  quadruple stars. ",
        "introduction": "\\label{sec:intro}   Formation of binary and multiple stars is a subject of active research and  debate,  still  remaining   one  of  the  major  unsolved  issues \\citep{IAU200}.  Reliable data on the statistics of binary and multiple stars  in different  environments  are essential  for comparison  with theory and further progress.  So  far, binary stars have received some attention,  but  the  statistics  of higher  multiplicities  is  still uncertain.   The aim  of this  paper is  to review  the  statistics of multiple stars, focusing on quadruple systems, and to relate it to the formation scenarios, whenever possible.  Multiple stellar systems are  often considered individually, hence the formation  of each  object can  be explained  by a  unique  chain of events. Yet, multiples  are not rare ``freaks'' in  the stellar world. At  least  8\\% of  solar-type  stars  have  three or  more  components \\citep{Tok08}.  The  nearest star, $\\alpha$~Cen,  is triple.  Multiple systems are normal products  of star formation, and the star-formation theory will eventually be able to model their statistical properties.  A textbook example  of a quadruple system is  $\\varepsilon$~Lyr -- two wide  visual binaries  on a  still  wider common  orbit. Two  striking properties of this system are the similarity of the masses of all four components and comparable periods of the inner sub-systems, suggesting that the  formation process was  not completely random.  We  show that such quadruples  with similar components and inner  periods are indeed typical.  One reason why  the statistics of multiple stars  remains murky is the difficulty of  accounting for the  observational selection. Catalogued multiples  result from  random  discoveries using  a  wide variety  of techniques, so  quantifying the observational bias  seems hopeless.  A better  approach would  be to  study volume-limited  samples,  but the number of  multiples even  among nearby G-dwarfs  is small,  while the incompleteness of  higher-order hierarchies is  still significant. So, at present,  the only way to  look at the statistics  of $N>2$ stellar systems  is  by   using  catalogues.   \\citet{Fekel81}  pioneered  the statistical approach  to triple stars.  The number  of known multiples has significantly increased since then, making it possible to take the first look at  systems with more than three stars.   About half of the known quadruples were discovered during  the last 20 years as a result of  systematic radial velocity  programs and  high-resolution imaging, the  latest ones  \\citep{Shkolnik}  are not  included  in the  present study.   Here  we use  the updated  Multiple Star  Catalogue \\citep[][hereafter MSC]{MSC} to study  the statistics of quadruple stars  composed of two close  pairs orbiting each  other.  The  problem of  the observational selection is partially addressed by comparing with the triple systems, assuming  that  the  selection  effects  in both  cases  are  similar. Although the exact form of this selection filter remains undetermined, we can make  reasonable assumptions about its behaviour  to figure out which statistical features are real.   A catalogue such as MSC gives a distorted,  but  still  useful  information  on  the  true  underlying statistics of multiple stars.   We begin with a short  review of multiple-star formation processes and their  predictions for  the  statistics (Sect.~\\ref{sec:form}).   Some prototypical  multiple stars are  presented in  Sect.~\\ref{sec:sys} to introduce relevant terminology and to give a feeling of the data.  The samples   of   quadruple   and   triple   stars   are   described   in Sect.~\\ref{sec:samp}  together   with  the  selection   effects.   The comparative   statistics   is   presented   in   Sect.~\\ref{sec:comp}. Section~\\ref{sec:sum}  summarises  our   findings  and  discusses  the implications for multiple-star formation.   %--------------------------------------------------------------- ",
        "conclusions": "\\label{sec:sum}  The properties of catalogued triple and quadruple systems are: % \\begin{enumerate} \\item The distributions  of periods in the inner  sub-systems of quadruple and triple  stars are similar  and bimodal. They are  different from the period  distribution of pure binaries.  There  is no correlation between inner mass ratios and inner periods.  \\item The distributions of periods and mass ratios in the outer sub-systems of quadruple and triple stars  are different.  The median outer mass ratio in  triples is 0.39. It  does not depend on  the outer period, which has a smooth distribution.   In contrast, the outer periods of 25\\%  quadruples concentrate  in  the narrow  range  from 10\\,yr  to 100\\,yr, all  these tight quadruples have $q_{L4}>  0.6$ (possibly a selection effect), while the periods of their two inner binaries are similar.  \\item Quadruple  systems of  $\\varepsilon$~Lyr type  with  comparable inner periods and component  masses are common.  In 42\\%  of the quadruples the outer mass  ratios are above 0.5 and the  inner periods differ by less than 1\\,dex.  \\item The outer and inner mass ratios in triple and quadruple stars are not mutually correlated. In 13\\% of quadruples both inner mass ratios are above 0.85 (double twins).  \\item The  values  of the  inner  and  outer  orbital angular  momenta  (or periods) in  wide ($P_S  > 30$\\,d) triple  and quadruple  systems show some  correlation,  the  ratio  of outer-to-inner  periods  is  mostly comprised  between 5  and $10^4$.  In  the systems  with small  period ratios  (low  hierarchy)  the  directions  of the  orbital  spins  are correlated, while in the systems with large ratios they are not.  \\end{enumerate}  %------------------------------------------------ %"
    },
    "0806/0806.4831_arXiv.txt": {
        "abstract": "{} {The knowledge of physical conditions in young radio sources is important for defining the framework of models describing radio source evolution. We investigate whether young radio sources are in equipartition (i.e. minimum energy) conditions by comparing the equipartition magnetic fields of 5 High Frequency Peakers (HFP) with values directly inferred from the spectral peak assumed to be produced by synchrotron self absorption.} {Multi-frequency VLBA observations of 5 HFPs were carried out in both the optically thick and thin part of the spectrum to determine the spectral shape and angular size of the components for which individual radio spectra were obtained.} { We find that the magnetic fields measured using observations agree well with those obtained by assuming equipartition, which implies that these sources are in minimum energy condition and the turnover in their spectra is due probably to SSA. In two source components, we found that the peak of the spectrum is caused by absorption of a thermal plasma instead of being due to SSA. The magnetic fields found in the various components range from 10 to 100 mG. In the presence of such high magnetic fields, electron populations with rather low $\\gamma$ emit in the GHz-regime. In one source, we detect low-surface brightness extended emission at low frequency located $\\sim$ 30 mas ($\\sim$ 50 pc) from the main source. This feature may be related to either an earlier episode of radio activity or a discontinuous start of the radio activity (sputtering). By comparing our data with previous VLBA observations, we estimate the hotspot advance speed to be in the range 0.1-0.7$c$ and kinematic ages of a few hundred years.} {} ",
        "introduction": "In a scenario where radio sources grow in a self-similar way, the evolution of each radio object originated by an Active Galactic Nucleus (AGN) depends on its linear size. It is established that the population of intrinsically compact radio sources represents an early stage of radio source evolution. In such a context, the empirical anti-correlation found between the linear size and the turnover frequency (O'Dea \\& Baum \\cite{odea97}) implies that the youngest sources - with ages of about $\\sim$ 100 - 1000 years - must be sought among those whose turnover frequency occurs at a few GHz, which are termed ``High Frequency Peakers'' (HFPs, Dallacasa et al. \\cite{dd00}; Dallacasa \\cite{dd03}).\\\\ Several evolution models (e.g. Fanti et al. \\cite{cf95}; Readhead et al. \\cite{read96}; Snellen et al. \\cite{sn00}) were proposed to describe how the physical parameters (i.e. luminosity, linear size and velocity) evolve as the radio emission grows within the host galaxy. Such evolution models are based on the assumption of minimum energy content, corresponding to a near equipartition of energy between the radiating particles and the magnetic field (Pacholczyk \\cite{pacho70}) usually referred to as ``equipartition''. However, there is no {\\it a priori} reason for believing that magnetic fields in radio sources are in equipartition. It is thus important to test the reliability of the equipartition at various stages of the radio source evolution; the study of newly born radio sources provides useful information about the physical parameters at the beginning of source growth. If the spectral peak is due to pure synchrotron self-absorption (SSA) as claimed by Snellen et al. (\\cite{sn00}), we have an independent way of estimating the magnetic field by measuring the frequency at which the radio emission changes from being opaque to transparent. In this case, the magnetic field $H$ can be measured directly from the spectral peak parameters, namely the peak frequency $\\nu_{p}$, peak flux density $S_{p}$, and the source component angular sizes $\\theta_{maj}$ and $\\theta_{min}$ (Kellermann \\& Pauliny-Toth \\cite{kpt81}):  \\begin{equation} H \\sim f(\\alpha)^{-5} \\theta_{maj}^{2} \\theta_{min}^{2} \\nu_{p}^{5} S_{p}^{-2}(1+z)^{-1} \\label{h_synchro} \\end{equation}  \\noindent The function $f(\\alpha)$ only weakly depends on $\\alpha$, and for $\\alpha = 0.5$, $f(\\alpha) \\sim 8$ (Kellermann \\& Pauliny-Toth \\cite{kpt81})\\footnote{Throughout this paper, the radio spectrum is described by the power law $S \\propto \\nu^{- \\alpha}$.}. The main difficulty in applying this method has been the uncertainty in determining source component parameters at the turnover frequency. Since the magnetic field depends strongly on these observable quantities, the estimates of the field strength derived so far are not very accurate. In particular, the uncertainties in the component size, a parameter that may be difficult to measure, contribute to the limited accuracy of the magnetic field estimates. Furthermore, magnetic fields should not be measured for the entire source but for its individual sub-structures, which are more likely to represent ``homogeneous'' components such as those studied in analytic models.\\\\ Scott \\& Readhead (\\cite{sr77}) and Readhead (\\cite{read94}) computed the magnetic field for sources of low-frequency spectral turnovers close in value to the observing frequency and found that the magnetic fields inferred directly from the spectrum were within a factor of 16 of the equipartition values. However, there are no systematic studies of sources with spectra peaking at higher frequencies, i.e. objects younger than those in Scott \\& Readhead (\\cite{sr77}). \\\\ Spangler et al. (\\cite{spangler83}) studied the magnetic field in the GHz-peaked spectrum (GPS) sources DA\\,193 (J0555+3948) and B1848+283, which were both classified as HFP by Dallacasa et al. (\\cite{dd00}). They found that B1848+283 was in equipartition, while DA\\,193 was not. On the other hand, Orienti \\& Dallacasa (\\cite{mo08a}) investigated the magnetic field in the very young GPS object J1459+3337 and did not find any evidence of departure from equipartition. It can be considered to be one of the youngest objects known so far.\\\\ However, it is possible that the turnover in the spectrum is caused by free-free absorption (FFA) in the external ionized medium, as found for a number of young radio sources (e.g. Kameno et al. \\cite{kameno00}; Marr et al. \\cite{marr01}). In this case, the magnetic field inferred from the peak would be physically meaningless.\\\\ We determine the magnetic fields of the components of 5 High Frequency Peaker radio sources from the Bright HFP sample (Dallacasa et al. \\cite{dd00}). The turnover frequency has a value of a few GHz, corresponding to frequencies sampled well by VLBA observations; this implies that the magnetic field strength of these objects can be readily determined from spectral parameters and we can infer whether such objects are in equipartition.   Throughout this paper, we assume the following cosmology: $H_{0} = 71\\, {\\rm km\\, s^{-1}\\, Mpc^{-1}}$, $\\Omega_{\\rm M} = 0.27$ and $\\Omega_{\\Lambda} = 0.73$, in a flat Universe.\\\\ ",
        "conclusions": "We have presented the results of multi-frequency VLBA observations of 5 young High Frequency Peakers. By assuming that the turnover of the radio spectrum is due to synchrotron self-absorption, we derived the magnetic field of each source sub-component by means of observable quantities, namely the peak frequency, peak flux-density, and the angular size. The magnetic fields calculated in this way are usually in good agreement with those obtained by assuming equipartition, with two exceptions. This suggests that in general young radio sources are in minimum energy conditions and their spectral turnover is caused by SSA. The values found for the magnetic fields are in the range of 10 mG up to $\\sim$ 160 mG, if we consider also the GPS source J1459+3337.\\\\ In two sub-components, we found evidence that the turnover of the spectrum is probably due to absorption by a thermal plasma located in front of the component, instead of SSA. The other components of these sources do not show any evidence of departure from SSA, suggesting that the absorber is an inhomogeneous ambient medium with different free-free opacities along the lines of sight to the two lobes. \\\\ In young radio sources, the presence of such high magnetic fields implies that even electron populations with small $\\gamma$ (e.g. $\\gamma \\sim 200$) radiate at the ``high'' frequencies sampled by the available surveys. This may provide a further explanation of the high number counts of young radio sources with respect to the larger ones.\\\\ In HFP J1511+0518 at 1.7 GHz, we unambiguously detected a low-surface brightness extended emission not aligned with the main source structure. Such a feature may be the relic of a previous epoch of radio activity that occurred not long ago ($< 10^{4}$ years). This may be consistent with a scenario of a discontinuous start (sputtering) of the radio activity.\\\\ The analysis of the hot-spot separation velocity provided kinematic ages of a few hundreds of years, confirming that the targets can be considered young radio sources.\\\\"
    },
    "0806/0806.1735_arXiv.txt": {
        "abstract": "Precision data from cosmology suggest neutrinos stream freely and hence interact very weakly around the epoch of recombination. We study this issue in a simple framework where neutrinos recouple instantaneously and stop streaming freely at a redshift $z_i$. The latest cosmological data imply $z_i\\alt 1500$, the exact constraint depending somewhat on the assumed prior on $z_i$. This bound translates into a limit on the coupling strength between neutrinos and majoron-like particles $\\phi$, implying $\\tau \\agt 1 \\times 10^{10} \\, {\\rm s} \\,(m_2/50\\,{\\rm meV})^3$ for the decay $\\nu_2 \\to \\nu_1+\\phi$. ",
        "introduction": "\\label{sec:introduction}  With the advent of high-precision cosmology it has become feasible to probe progressively more detailed aspects of the cosmic neutrino background radiation~\\cite{Hannestad:2006zg, Lesgourgues:2006nd}. In the standard model, neutrinos provide relativistic energy density which influences the cosmic microwave background (CMB) radiation mainly via the early Integrated Sachs--Wolfe (ISW) effect and the matter fluctuation spectrum via the relation between neutrino energy density and the epoch of matter--radiation equality. The existence of a cosmological background of relativistic energy density has been unambiguously detected in the \\hbox{WMAP-5} data~\\cite{Komatsu:2008hk} and was already previously detected using the combination of CMB and Large Scale Structure (LSS) data~\\cite{Hannestad:2001hn, Hannestad:2005jj, Ichikawa:2006vm, Mangano:2006ur, Hamann:2007pi}. Furthermore, cosmological data provide a restrictive upper bound on the sum of neutrino masses of 0.2--1~eV, depending on the specific choice of data sets and model space~\\cite{Zunckel:2006mt, Cirelli:2006kt, Goobar:2006xz, Kristiansen:2006xu, Seljak:2006bg, Fogli:2006yq, Hannestad:2003xv, Hannestad:2006zg, Hannestad:2007tu}.  The present level of precision allows us to turn to more subtle issues. For example, it is timely to probe the possibility that neutrinos have non-standard interactions where one case in point is an interaction with the Nambu-Goldstone boson of a new, broken $U(1)$ symmetry as in majoron models~\\cite{Gelmini:1980re, Chikashige:1980ui, Schechter:1981cv}. Such an interaction would recouple the neutrinos to each other at some ``interaction redshift'' $z_i$, whereas at earlier epochs they would behave in the same way as standard-model neutrinos. For the cases of interest, this recoupling occurs much later than the electroweak decoupling. Therefore, in the limit of relativistic neutrinos the total energy density in the combined fluid of neutrinos and majorons is conserved, preventing any direct impact on cosmological observables.  However, neutrinos lose their free-streaming property if the interaction is sufficiently strong. As a consequence, any anisotropic stress components in the Boltzmann hierarchy are suppressed, effectively truncating the Boltzmann hierarchy at first order, equivalent to the equations for a perfect fluid \\cite{Raffelt:1987ah, AtrioBarandela:1996ur, Beacom:2004yd, Hannestad:2004qu, Hannestad:2005ex, Friedland:2007vv, Sawyer:2006ju, Bell:2005dr, DeBernardis:2008ys}. (See Ref.~\\cite{Ma:1995ey} for a detailed description of the Boltzmann hierarchy.)  The impact of neutrino free streaming on cosmological observables was recently studied in Ref.~\\cite{DeBernardis:2008ys}. The fit parameter was the effective viscosity $c_{\\rm vis}^2$, taken to be independent of redshift. Our study is complementary in that we assume that $c_{\\rm vis}^2$ drops instantaneously at $z_i$ from the free-streaming value 1/3 to the perfect-fluid value~0\\symbolfootnote[2]{A paper by the same authors with a more extensive treatment of this issue is in preparation.}. Our conclusion that neutrinos should stream freely around the epoch of recombination is perfectly consistent with Ref.~\\cite{DeBernardis:2008ys}. However, our approach lends itself more directly to an interpretation in terms of a specific interaction model where the recoupling redshift is related to a dimensionless coupling constant~$g$. Therefore, we can translate our limits on $z_i$ into limits on~$g$.  From flavor oscillation experiments we know that neutrinos have masses which therefore are unavoidable cosmological fit parameters.  The usual cosmological limits on the sum of neutrino masses imply that any single mass eigenstate should obey $m\\alt0.2$--0.3~eV so that all neutrinos would be relativistic around the recombination epoch. Treating them as massless is therefore a reasonable approximation for the simple problem addressed here. On the other hand, a strong majoron-type interaction can lead to the annihilation of ``heavy'' neutrinos into majorons (``neutrinoless universe''~\\cite{Beacom:2004yd}). Such scenarios lead to a complicated evolution of the neutrino-majoron fluid that we are not investigating, although it would have a strong impact on cosmological observables. In any event, our constraint on the free-streaming nature of the relevant radiation at recombination does not depend on the physical nature of the radiation.  Eventually the KATRIN experiment, unless it detects a significant neutrino mass, will constrain the neutrino mass scale to $m\\alt0.2$~eV \\cite{Drexlin:2005zt}. Such a bound would imply that neutrinos can not have disappeared at the recombination epoch and our constraint indeed applies to neutrinos. In this sense the anticipated KATRIN limit will strengthen the case for translating our limit on $z_i$ into a limit on exotic neutrino interactions.  We begin in Sec.~\\ref{sec:models} with a description of our model space, data sets, and statistical methodology. In Sec.~\\ref{sec:recoupling} we provide our bounds on $z_i$ that are translated, in Sec.~\\ref{sec:g}, into limits on a neutrino-majoron coupling strength $g$. We conclude in Sec.~\\ref{sec:discussion}. ",
        "conclusions": "\\label{sec:discussion}  We have updated bounds on the coupling strength between neutrinos and a new, light pseudo-scalar, $\\phi$, using the latest cosmological data. Performing a slightly more refined calculation than in an earlier paper by two of us~\\cite{Hannestad:2005ex} we find essentially unchanged constraints.  One way to improve this limit in future is by actually detecting neutrino hot dark matter in cosmological precision data. In decay scenarios involving massless pseudoscalars and for a mass of 50~meV, the lifetime limit would improve by some six orders of magnitude~\\cite{serpico}.  Our more general conclusion is that neutrinos which are strongly interacting around recombination are strongly disfavored by data. The present data strongly support the conclusion that the cosmic neutrino background (i)~exists around the epoch of recombination and (ii)~its fluctuations do have an anisotropic stress component. In future CMB data alone will likely suffice to reach the same or better sensitivity so that the bound on $g$ can be expected to improve significantly \\cite{Friedland:2007vv, DeBernardis:2008ys}.  \\vspace*{-0.5cm}"
    },
    "0806/0806.1673_arXiv.txt": {
        "abstract": "The atomic hydrogen (HI) 21~cm line measures the gas content within and around galaxies, traces the dark matter potential and probes volumes and objects that other surveys do not. Over the next decade, 21~cm line science will exploit new technologies, especially focal plane and aperture arrays, and will see the deployment of Epoch of Reionization/Dark Age detection experiments and Square Kilometer Array (SKA) precursor instruments.  Several experiments designed to detect and eventually to characterize the reionization history of the intergalactic medium should deliver first results within two--three years time. Although ''precision cosmology'' surveys of HI in galaxies at $z \\sim1$ to $3$ require the full collecting area of the SKA, a coherent program of HI line science making use of the unique capabilities of both the existing facilities and the novel ones demonstrated by the SKA precursors will teach us how many gas rich galaxies there really are and where they reside and will yield fundamental insight into how galaxies accrete gas, form stars and interact with their environment. ",
        "introduction": "Huge surveys being undertaken at optical and near-IR wavelengths today are cataloguing and characterizing many millions of galaxies, feeding their colors and structural properties into sophisticated  models that describe their evolution over cosmic time. In contrast to the wealth of detail available for stellar components of galaxies, quantitative measures of the gas content are limited to only a few tens of thousands of objects, and of the gas distribution within them, to only hundreds. Yet, galaxies and the stars within them form out of collapsing gas clouds and both the amount of gas and its distribution within and around galaxies play critical roles in the story of galaxy evolution.  But, how important really is the HI 21~cm line? Sometimes at topical meetings such as this one, the attendees are mainly ``the usual suspects'' and the elder participants spend most of their time remembering the ``good old days''. Perhaps to unsettle any tendency towards scientific smugness about the long tradition of 21~cm line studies, Riccardo Giovanelli, right at the beginning of this conference, reminded us that HI contributes a ``piffling'' amount even to the overall baryon budget in the universe. If that were not enough to deflate our HI egos, Todd Tripp rightly pointed out that Ly$\\alpha$ absorption lines are a factor of $10^6$ times more sensitive to low column density gas than are typical HI maps, and numerous speakers reminded us that nearly every map we make in HI pales in spatial detail to the minutely elegant structures gleaned from the pretty images made with ``real'' telescopes. And, as if to render the conference pointless, Leo Blitz questioned whether HI is really relevant to studies of star formation, since stars form out of H$_2$, not HI. With this self-skepticism as a backdrop to this conference, speakers and participants engaged in healthy discussion of the relative importance of HI and other techniques to cosmology and galaxy studies. The often-lively discussion was not just of what we do and don't know, but most importantly, what we {\\it might} know in the next 5-10 years. It was that focus on peering into the HI crystal ball that was most appealing: where is HI science headed and how can we expedite the voyage? In this review, I focus on the outcome of those discussions, especially as they relate to the the US-based facilities. ",
        "conclusions": "At this workshop, we discussed what are the topics in extragalactic science for which HI is relevant and powerful and what real progress could we make in the near future. We agreed that the long term objective was more collecting area, but we also discussed how real {\\it transformational} progress could be made within the next decade. HI studies are {\\it vital} to the understanding of galaxy evolution and provide quantitative clues that {\\it no other wavelength} can. To maximize the potential scientific payoff, we should, as a community, undertake {\\it a set of coordinated large surveys} that are {\\it fully integrated} with efforts to study the stellar and gaseous components of galaxies via other techniques; some {\\it astronomical socialism} will produce datasets which can be analyzed democratically. I do not think I am misrepresenting the general sentiment of the assembled group in claiming that a strategy for HI science over the next 5--10 years has three key elements: (1) undertaking of the best science with {\\it all} of the SKA ``precursors'' to nail down the role of HI (and dark matter!) in the local universe and to begin exploration of the evolution of that role over cosmic time; (2) development of innovative technologies deployed on the existing and precursor facilities which will lead us toward the SKA; (3) engaging the broadest community to understand and challenge our understanding of how gas fuels star formation and galaxy evolution.  The third element is the one least often discussed. Since the HI facilities we can use are few in number and shared with many other users, the undertaking of large survey programs must meet ``legacy'' standards but, for the health of our science, they must involve a broad cross-section of the community. Unless we are ready to turn over survey programs to a few individuals, we must figure out how to engage large groups in their undertaking without sacrificing quality within fiscal and practical boundaries. Our challenge cannot just be ``what to do'', it also must include ``how to do it''. And ``how'' demands ``strategy'', ``balance'' and ``inclusiveness''.  Neutral hydrogen may make up a ``piffling'' amount of the total matter in the universe, but its physical simplicity and the information it conveys about the gas content, distribution, kinematics and dynamics opens a powerful window on the evolution of galaxies, of the hierarchical structures in which they reside and of the universe itself.    \\begin{theacknowledgments} My participation in ALFALFA is supported by NSF grant AST-0607007 and by the Brinson Foundation. The Arecibo Observatory is part of the National Astronomy and Ionosphere Center which is operated by Cornell University under a cooperative agreement with the National Science Foundation. I thank the staff of the Arecibo Observatory for their ever warm and generous attention in support of this conference and especially Robert Minchin for inspiring and organizing it.  \\end{theacknowledgments}"
    },
    "0806/0806.1445_arXiv.txt": {
        "abstract": "{} {We measured radial metallicity profiles for a sample of $\\approx 50$ hot, intermediate redshift galaxy clusters, selected from the \\emph{XMM-Newton} archive.} {As in our previous paper, we used background modeling rather than background subtraction, and the Cash statistic rather than the $\\chi^2$. This method requires a careful characterization of all background components. We also performed montecarlo simulations to assess systematic effects.} {The mean metallicity profile shows a peak in the center, and gently declines out to 0.2~$R_{180}$. Beyond 0.2~$R_{180}$ the metallicity is $\\approx 0.2$ solar and, at variance with recently published expectations based on simulations, consistent with being flat. We find no evidence of profile evolution from $z = 0.1$ to $z = 0.3$. When comparing our mean profile to those obtained by recent works with \\emph{BeppoSAX} and \\emph{Chandra}, we find remarkable agreement over the entire radial range.} {} ",
        "introduction": "\\label{sec: intro} Clusters of galaxies are the most massive gravitationally bound systems in the universe. They are permeated by the hot, X-ray emitting intra-cluster medium (ICM), which represents the dominant baryonic component. The mean metallicity of the ICM is found to be roughly half of the solar value \\citep{lowe04}, therefore a substantial fraction of plasma is not of primordial origin. As heavy elements are only produced in stars, the processed material must have been ejected by cluster galaxies into the ICM. In rich clusters the plasma reaches temperatures of several $10^7$~K and emits X-rays mainly via thermal bremsstrahlung. At such temperatures, most elements are either fully ionized or in a high ionization state. The strong emission lines due to the transitions to the n=1 level of the H-like and He-like ions of Iron around 6.7~keV (rest frame energy) are the most prominent features in X-ray spectra of hot clusters. In fact, for hot clusters as those belonging to our sample, the measurement of the metallicity is a measurement of the iron abundance.  Spatially resolved analysis of metal abundance in clusters has become possible only recently, first with \\emph{ASCA} and \\emph{BeppoSAX} and then with \\emph{Chandra} and \\emph{XMM-Newton}. These measurements (e.g. \\citealp{grandi01,grandi04,tamura04,baldi07}) have shown that abundance gradients are common features in clusters of galaxies: most clusters show a peak of metallicity in the center and a gentle decline outward.  In this paper we present radial metallicity profiles from the same sample of hot intermediate redshift clusters analyzed in our previous paper \\citep{leccardi08}, hereafter Paper~I. We used the same novel data analysis technique, which is extensively described in Paper~I and summarized in Sect.~\\ref{sec: analysis}. In Sect.~\\ref{sec: mean prof} we show the mean metallicity profile, in Sect.~\\ref{sec: evolution} search for an indication of the evolution of the metallicity with redshift, and in Sects.~\\ref{sec: sim} and \\ref{sec: comp} compare our mean profile with those obtained from hydrodynamic simulations and from previous observational works. In the Appendix~\\ref{app} we describe a systematic effect, which could affect the measurement of the metallicity especially for the case of low metallicity, statistically poor spectra.  Quoted confidence intervals are 68\\% for one interesting parameter (i.e. $\\Delta$C~=~1), unless otherwise stated. All results are given assuming a $\\Lambda$CDM cosmology with $\\Omega_\\mathrm{m} = 0.3$, $\\Omega_\\Lambda = 0.7$, and $H_0$~=~70~km~s$^{-1}$~Mpc$^{-1}$. ",
        "conclusions": "\\label{sec: concl} We analyzed a sample of $\\approx 50$ hot, intermediate redshift galaxy clusters to measure their radial properties, in this paper we focused on metallicity profiles. Our main results are summarized as follows. \\begin{itemize} \\item The mean metallicity is $0.45 \\; Z_\\odot$ in the center and decreases out to $\\approx 0.2 \\; R_{180}$; beyond $0.2 \\; R_{180}$ the metallicity is consistent with being flat at $0.23\\pm0.01 \\; Z_\\odot$. \\item The profiles show a large scatter, which is mostly of statistical origin. In the central regions the scatter (i.e. 22~$\\pm$~2\\%) is also related to the presence of cool core clusters, in the outer regions it (i.e. 14~$\\pm$~8\\%) is comparable to systematics. \\item There is no evidence of profile evolution from $z = 0.1$ to $z = 0.3$. \\item We find a mean metallicity $\\approx 20\\%$ lower than found by \\cite{maughan08} using \\emph{Chandra}, when analyzing roughly the same regions for clusters in the same redshift range. \\item We point out the existence of a systematic effect, affecting in particular distant clusters, which, if unaccounted for, could cause a false detection of evolution. \\item When comparing our mean profile to the one derived from hydrodynamic simulations by \\cite{fabjan08}, we find differences at small and large radii. In particular, while the profile obtained by \\cite{fabjan08} shows a constant decrease, ours is consistent with being flat beyond $\\approx 0.2 \\; R_{180}$. \\item When comparing our mean profile to those obtained by recent works with \\emph{BeppoSAX} \\citep{grandi04} and \\emph{Chandra} \\citep{baldi07}, we find remarkable agreement over the entire radial range. \\end{itemize}  Our results have been obtained using the same novel data analysis technique as described in Paper~I; i.e., the background modeling rather than the background subtraction, and the Cash statistic rather than the $\\chi^2$."
    },
    "0806/0806.2393_arXiv.txt": {
        "abstract": "We measure the correlation between the arrival directions of the highest energy cosmic rays detected by the Pierre Auger Observatory with the position of the galaxies in the \\hi\\ Parkes All Sky Survey (HIPASS) catalogue, weighted for their \\hi\\ flux and Auger exposure. The use of this absorption--free catalogue, complete also along the galactic plane, allows us to use all the Auger events. The correlation is significant, being 86.2\\% for the entire sample of \\hi\\ galaxies, and becoming 99\\% when considering the richest galaxies in HI content, or 98\\% with those lying between 40--55 Mpc. We interpret this result as the evidence that spiral galaxies are the hosts of the producers of UHECR and we briefly discuss classical (i.e energetic and distant) long Gamma Ray Burst (GRBs), short GRBs, as well as newly born or late flaring magnetars as possible sources of the Auger events. With the caveat that these events are still very few, and that the theoretical uncertainties are conspicuous, we found that newly born magnetars are the best candidates.  If so, they could also be associated with sub--energetic, spectrally soft, nearby, long GRBs. We finally discuss why there is a clustering of Auger events in the direction on the radio--galaxy Cen A and an absence of events in the direction of the radio--galaxy M87. ",
        "introduction": "The origin of ultra--high energy cosmic rays (UHECR), exceeding 10 EeV (1 EeV=$10^{18}$ eV) has been a mystery for decades, but the recent findings of the large area detectors, such as AGASA (Ohoka et al. 1997), HIRes (Abu--Zayyad et al. 2000), and especially the Pierre Auger Southern Observatory (Abraham et al. 2004), began to disclose crucial clues about the association of the highest energy events with cosmic sources. The Auger collaboration (Abraham et al. 2007) found a positive correlation between the arrival directions of UHECR with energies grater than 57 EeV and nearby AGNs (in the optical catalogue of Veron--Cetty \\& Veron 2006). Although this result has not been confirmed by HIRes (Abbasi et al. 2008) and it has been criticised by Gorbunov et al. (2008), it received an important confirmation by George et al.  (2008), who considered a complete sample of nearby hard X--ray emitting AGNs detected by the BAT instrument onboard {\\it Swift}. This sample is much less affected by absorption than any optical sample although, to identify as such an AGN, one relies on optical identification. Moreover, George et al. (2008) found a correlation not simply with the AGN locations, but by weighting them for the X--ray flux and the Auger exposure.  This association, if real, is surprising, since the large majority of the correlating AGNs are radio--quiet, a class of objects not showing, in their electromagnetic spectrum, any sign of non--thermal high energy emission: no radio--quiet AGN was detected by the EGRET instrument onboard the {\\it Compton Gamma Ray Observatory} (Hartman et al. 1999).  Therefore they must accelerate particles (protons, nuclei, and presumably their accompanying electrons) to ultra--high energies without any noticeable radiative emission from these very same particles.  Radio--loud AGNs, instead, together with Gamma Ray Bursts (of both the long and short category) do show high energy non--thermal emission, and have been considered for a long time better candidates as UHECR sources (Vietri 1995; Waxman 1995; Milgrom \\& Usov 1995; Wang, Razzaque \\& M\\'esz\\'aros 2008; Murase et al. 2008; Torres \\& Anchordoqui 2004 and Dermer 2007 for reviews, and Nagar \\& Matulich 2008 and  Moskalenko et al. 2008 for the possible association of the AUGER events with radio--loud AGNs). Note also that some short GRBs could be due to the giant flares of highly magnetised neutron stars (``magnetars'', as the 27 Dec 2004 event from 1806--20; Borkowski et al. 2004; Hurley et al. 2005; Terasawa et al. 2005), and that, at birth, a fastly spinning magnetar can be much more energetic than when, later, it produces giant flares (Arons 2003).   The possibility that GRBs and magnetars are the sites of production of UHECR would directly imply the direct association of these events with (normal) galaxies. In this case the found association of UHECR with nearby AGNs might then be due to the fact that local AGNs just trace the distribution of galaxies. The aim of the present paper is to test this possibility directly correlating the locations of the ultra--high energy Auger events with a well defined, complete, and possibly absorption--free sample of galaxies.  For this purpose we use the sample of \\hi\\ emitting galaxies, compiled using the Parkes 64--m radio telescope (Barnes et al.  2001; Staveley--Smith et al. 1996), which is conveniently located in the south hemisphere, as the Auger observatory.  The entire sample covers the portion of the sky visible by Auger, making it possible to use, for the correlation analysis, all the 27 UHECR events with energies larger than 57 EeV detected by Auger, without excluding the galactic plane, as is instead necessary when dealing with AGNs or optically selected galaxies. Note that the presence of neutral hydrogen strongly favours spirals (or, more generally, gas--rich galaxies) with respect to elliptical galaxies.  We use a cosmology with $\\rm{h}_0=\\Omega_\\Lambda=0.7$ and $\\Omega_{\\rm M}=0.3$. ",
        "conclusions": "We have correlated the cosmic rays with $E>57$ EeV detected by the Auger Observatory with a complete, absorption--free sample of \\hi\\ selected galaxies. We found a significant correlation when correlating the \\hi\\ {\\it flux} of galaxies of our sample.  When considering the largest 95HIPASS catalogue and the 27 UHECRs we find a weak correlation (probability of 72\\%), while a larger significance (87.8\\%) is reached if we consider the most complete 99HICAT sample of galaxies (though with 25 UHECRs). These probabilities are maximised by cutting the \\hi\\ sample in distance or luminosity bins: it becomes 99\\% when considering the 500 most luminous (or most \\hi\\ massive) galaxies (1/4 of the sample), and 98\\% when considering the 500 galaxies lying between 38 and 54 Mpc, where the Centaurus cluster of galaxies is.  Thus there is the possibility that the UHECRs coming from the direction of Cen A are instead coming from the more distant Centaurus cluster. Galaxies of this cluster are richer in \\hi\\ than Virgo galaxies, explaining why there is no UHECR event from the direction of Virgo.  This sample is formed by \\hi\\ emitting galaxies, therefore it is biased against ellipticals.  The found correlation with these galaxies, per se, is not disproving the found correlation with AGNs (Abraham et al. 2007, 2008; George et al. 2008), since they also trace the local distribution of matter, as spiral galaxies do. On the other hand, it opens up the possibility, on equal foot, that UHECRs are produced by GRBs or newly born magnetars (see also Singh et al. 2004 who used AGASA events). With the caveat that it is premature, with so few events and big theoretical uncertainties, to draw strong conclusions, we have pointed out that although classical (i.e. energetic) long GRBs and short GRBs have difficulties in producing the required UHECR flux, newly born magnetars can.  If so, they could also be a subclass of {\\it long} GRBs, possibly sub--energetic and relatively nearby, powered by fastly spinning, newborn magnetars. The future increased statistics of UHECRs arrival directions will help to discriminate among the different proposed progenitors, especially if there will be (or not) an excess of events close to the radio core and/or lobes of Cen A."
    },
    "0806/0806.2861_arXiv.txt": {
        "abstract": "Previous models of galactic disk heating in interactions invoke restrictive assumptions not necessarily valid in modern $\\Lambda$CDM contexts: that satellites are rigid and orbits are circular, with slow decay over many orbital periods from dynamical friction. This leads to a linear scaling of disk heating with satellite mass: disk heights and velocity dispersions scale $\\propto M_{\\rm sat}/M_{\\rm disk}$. In turn, observed disk thicknesses present strong constraints on merger histories: the implication for the Milky Way is that $<5\\%$ of its mass could come from mergers since $z\\sim2$, in conflict with cosmological predictions. More realistically, satellites merge on nearly radial orbits, and once near the disk, resonant interactions efficiently remove angular momentum while tidal stripping removes mass, leading to rapid merger/destruction in a couple of free-fall plunges. Under these conditions the proper heating efficiency is non-linear in mass ratio, $\\propto (M_{\\rm sat}/M_{\\rm disk})^{2}$. We derive the scaling of disk scale heights and velocity dispersions as a function of mass ratio and disk gas content in this regime, and show that this accurately describes the results of simulations with appropriate ``live'' halos and disks. Under realistic circumstances, we show that disk heating in minor mergers is suppressed by an order of magnitude relative to the expectations of previous analyses. We show that the Milky Way disk could have experienced $\\sim5-10$ independent 1:10 mass-ratio mergers since $z\\sim2$, in agreement with cosmological models. Because the realistic heating rates are non-linear in mass, the predicted heating is dominated by the more stochastic, rare low mass-ratio mergers, and the existence of populations with little or no thick disk does not require fundamental modifications to the cosmology. This also leads to important differences in the predicted isophotal shapes of bulge-disk systems along the Hubble sequence. ",
        "introduction": "\\label{sec:intro}  In the widely-accepted $\\Lambda$CDM cosmology, it is expected that mergers and interactions between galaxies are crucial to galaxy assembly because structure grows hierarchically \\citep[e.g.][]{whiterees78}.  On this basis, it has been proposed that spheroids are produced when galaxies merge \\citep[making bulges in disks and elliptical galaxies;][]{toomre72,toomre77}.  Both theory \\citep[][and references therein]{ostrikertremaine75,maller:sph.merger.rates,stewart:mw.minor.accretion, fakhouri:halo.merger.rates} and observation \\citep[e.g.][]{lin:merger.fraction,woods:tidal.triggering,barton:triggered.sf, woods:minor.mergers} indicate that ``minor'' mergers of mass ratios $\\lesssim$\\,1:10 are commonplace, and that a large fraction ($\\sim1/2$) of the $\\sim L_{\\ast}$ galaxy population has experienced a ``major'' merger (mass ratio $\\lesssim$\\,1:3) since $z\\sim2-3$ \\citep{bell:merger.fraction, bridge:merger.fractions,kartaltepe:pair.fractions, lotz:merger.fraction,lin:mergers.by.type}.  The ubiquity of mergers has raised the possibility that cold, thin disks of galaxies may not survive to $z=0$ in hierarchical models. From numerical simulations \\citep[e.g.][]{barnes:disk.halo.mergers,barnes:disk.disk.mergers, hernquist:bulgeless.mergers,hernquist:bulge.mergers} it has been known for some time that major mergers destroy {\\it stellar} disks entirely and cause a large fraction of disk gas to lose angular momentum through gravitational torques \\citep[e.g.][]{barnes.hernquist.91,barneshernquist96} together yielding dense, quasi-spherical bulges \\citep[at least in typical situations such at those at low redshift, where the gas fractions are not large; see e.g.][]{hopkins:disk.survival}.  Moreover, although minor mergers (at least those with mass ratios $\\lesssim$\\,10:1) likely do not destroy disks, they are almost an order of magnitude more frequent than major mergers and as such may pose an even more severe problem for disk survival.  In the $\\Lambda$CDM cosmology it is unlikely than any disk with a significant stellar age has existed $\\sim5-10$\\,Gyr without experiencing a merger of mass ratio 10:1 or larger.  Analytic arguments \\citep{ostrikertremaine75,tothostriker:disk.heating, sellwood:resonant.disk.thickening} and early generations of numerical simulations \\citep{quinn.84,quinn86:dynfric.on.sats, hernquist.89,quinn93.minor.mergers, hernquist.mihos:minor.mergers, walker:disk.fragility.minor.merger,naab:minor.mergers} indicate that gas-poor minor mergers can convert a large fraction of a stellar disk into bulge and dynamically heat the disk. The observed thinness of galactic disks suggests that this may be a severe problem: \\citet{tothostriker:disk.heating} argued that large disks such as that in the Milky Way could not have undergone a merger of mass ratio $\\lesssim10:1$ in the last $\\sim10\\,$Gyr. More recently e.g.\\ \\citet{gilmore02:last.mw.merger.from.thick.disk,wyse06:thick.disk.stars} and \\citet{stewart:mw.minor.accretion} emphasized that the tension between these constraints and the expectation in CDM models that a number of such mergers should occur implies either a problem with our understanding of hierarchical disk formation or a challenge to the concordance cosmological model.  Given the many successes of the $\\Lambda$CDM model and observations that disks do experience (and survive) a large number of mergers, it seems likely that the main uncertainty is our understanding of disk galaxy evolution.  This has focused attention on the problem of forming disks in a cosmological context.  A number of studies have argued that detailed modeling of gas physics, star formation, and feedback, along with high numerical resolution, are required to simulate disk formation and understand how disks survive their violent merger histories \\citep{weil98:cooling.suppression.key.to.disks, sommerlarsen99:disk.sne.fb,sommerlarsen03:disk.sne.fb, thackercouchman00,thackercouchman01,abadi03:disk.structure, governato04:resolution.fx,governato:disk.formation, robertson:cosmological.disk.formation, okamoto:feedback.vs.disk.morphology,scannapieco:fb.disk.sims}. But there has been relatively less focus on whether these simulations produce disks -- even where the disks avoid being converted into bulges -- that are sufficiently thin relative to observational constraints. Interestingly, more recent attempts to model disk heating in a cosmological context \\citep[see e.g.][]{benson:heating.model, kazantzidis:mw.merger.hist.sim} suggest that disk heating may be significantly less efficient than has been assumed.  Indeed, a closer examination of this question reveals that much of the conventional wisdom and the implied constraints on the merger history of massive disks (and the Milky Way in particular) is fundamentally premised on historical assumptions that do not apply given what we now know about the $\\Lambda$CDM framework.  For example, implicit in the analysis of \\citet{tothostriker:disk.heating} and others is the assumption that the orbits of merging satellites can be described as being nearly circular, slowly spiraling in owing to dynamical friction, while the satellite remains relatively intact. Early simulations often adopted these choices -- modeling the satellites as rigid point masses on almost perfectly circular orbits, including relatively low-mass and smooth dark matter halos (such that the background density is low, and the in-spiral of the satellite therefore relatively slow and without much tidal stripping), and/or using rigid (as opposed to full N-body or ``live'') background disk/bulge/halo potentials. Moreover, these calculations were almost always dissipationless, and did not allow for the role of gas in making the system more robust against mergers and heating.  Under these conditions, it is straightforward to show that the disk heating owing to a merger is linear in the mass of the satellite.  This fundamentally underlies the constraint above: because the scaling is linear, it is possible to transform an observed disk thickness or velocity dispersion into an integral constraint on the allowed merger history of a galaxy.  Given cosmological constraints and numerical capabilities at the time, these were reasonable assumptions.  However, with improvements in simulations, yielding the ability to perform true cosmological simulations with resolved subhalo populations, and better understanding of cosmological parameters and estimates of halo masses and sizes, it is clear that the assumptions underlying these early works are not appropriate.  For example, it is now well-established that realistic cosmological orbits are quite radial, not circular. This is true on all mass scales, from large dark matter substructures which populate halos \\citep{vandenbosch:no.orbital.circularization.due.to.dyn.frict, benson:cosmo.orbits,khochfar:cosmo.orbits} to unresolved background velocity flows in the diffuse stellar and dark matter halo \\citep{cole.lacey:halo.structure,abadi06:stellar.halos}.  Given that realistic galaxy halos exhibit considerable substructure with large subhalo populations, it is generally true that systems which pass near a disk -- those of interest for mergers and disk heating -- will be those which can transfer angular momentum and collide with the disk on highly radial passages, rather than circular in-spirals \\citep[][]{font:sats.on.large.orbits,benson:heating.model}.  Moreover, properly accounting for live halos and disks leads to qualitatively different angular momentum transfer in the final stages of galaxy mergers: once a satellite reaches a radius comparable to the disk, resonances between the orbital motion and internal motions of the disk and halo remove angular momentum efficiently and lead to a rapid, radial coalescence in a couple of dynamical times (rather than the large number of orbital times required for a dynamical friction in-spiral in the case of mass ratios $\\gtrsim$\\,1:10).  Furthermore, allowing halos to have the large masses expected in modern cosmologies and resolving disks with N-body approaches, with proper live secondaries, leads to efficient tidal stripping of the satellite once it begins to pass through a disk, transferring orbital energy to the central bulge and extended dark matter halo (preferentially the hottest components of the system, rather than the disk).  The historical picture of ``sinking'' satellites on slowly decaying, circular orbits is based in part on an old view of the orbits of the Large and Small Magellanic Clouds (LMC and SMC). Recent observations, however, indicate that the LMC and SMC are on a high-velocity passage near pericenter and are only marginally bound to the Milky Way \\citep{kallivayalil:smc.proper.motion, kallivayalil:lmc.proper.motion, piatek:kallivayalil.reanalysis}. \\citet{besla:mc.orbits} used these results and modern cosmological models of the Milky Way halo to infer their orbital motion, and argued that they are in fact on highly eccentric radial orbits and are just now executing their first passage. Radial orbits are not uncommon in the local group: \\citet{howley08:ngc205.orbit} recently demonstrated that NGC 205 (an M31 dwarf) is also likely on a radial passage near its escape velocity; Leo I is similar \\citep{mateo:leo1.orbit, sohn:leo1.kinematics,munoz:leo1.kinematics}. This, and the arguments above, suggest that constraints regarding the merger history of the Milky Way and disk heating should be revisited in light of modern cosmological expectations.  One might think that the difference between the conventional lore on satellite heating and allowing for more realistic orbits would amount to a simple numerical correction -- i.e.\\ that the differences outlined above can be subsumed into an appropriate correction of the efficiency of disk heating, but will not fundamentally alter the scaling laws derived in the past (namely that disk heating is linear in merger mass).  In what follows, we show that this is not the case.  At a basic level, the linear scaling resulting from earlier assumptions is pathological, and allowing for moderate deviations in line with what is expected in a $\\Lambda$CDM cosmology yields a {\\em qualitatively} different scaling and a {\\em non-linear} dependence of disk heating on satellite mass.  This implies more than an order-of-magnitude suppression of disk heating in intermediate mass-ratio minor mergers, and essentially no contribution to disk heating from very minor mergers (unlike if the heating were linear, in which case there would be equal contributions to disk thickness from all logarithmic intervals in satellite and subhalo masses).  We show that the scaling we derive, non-linear (second-order) in the satellite mass, is dynamically inevitable and borne out in numerical simulations.  We outline the physics of interest in \\S~\\ref{sec:why}, explaining why cosmologically realistic merging orbits (even when nearly circular initially) tend to be highly radial by the time of coalescence and how this differs from the assumption of decaying circular orbits (\\S~\\ref{sec:orbits}), demonstrating how this changes the qualitative scaling of disk heating with mass ratio (\\S~\\ref{sec:heating}), deriving full analytic expressions for the disk thickness and vertical velocity dispersion as a function of merger and disk properties for the radial cases of interest (\\S~\\ref{sec:equations}), and testing that these equations accurately describe the results of full numerical simulations (\\S~\\ref{sec:tests}).  We discuss the accuracy and robustness of these approximations in \\S~\\ref{sec:robustness}, examining when they are or are not appropriate.  In \\S~\\ref{sec:consequences} we consider the consequences of the different scalings owing to radial versus circular orbits, in terms of the allowed merger histories of the Milky Way and other massive disk galaxies (\\S~\\ref{sec:consequences:mergerhistory}) and the shapes of galaxies along the Hubble sequence (\\S~\\ref{sec:consequences:shapes}). Finally, we summarize our results and discuss some of their cosmological implications and applications to other models and observations in \\S~\\ref{sec:discussion}.   \\breaker ",
        "conclusions": "\\label{sec:discussion}  We have shown that allowing for realistic satellite orbits -- in particular the fact that orbits are not exactly circular, and that satellites tend to be increasingly stripped as they execute multiple passages through a disk -- has dramatic implications for the efficiency and cosmological evolution of disk ``heating'' via minor mergers and galaxy-galaxy interactions.  In previous generations of analyses \\citep[e.g.][]{tothostriker:disk.heating}, several assumptions were made implicitly: that a satellite would enter a disk on a slowly-decaying circular orbit owing to dynamical friction, and that the satellite was rigid. In addition, much of this work focused on dissipationless (gas-free) cases. Under these conditions, it is straightforward to derive the fractional disk heating owing to a merger of a satellite of mass $M_{2}$ with a primary of mass $M_{1}$, and to show that this heating (and correspondingly the induced increase in disk scale height and vertical velocity dispersion) scale linearly in the mass ratio, $\\propto M_{2}/M_{1}$.  This fundamental premise has largely informed the literature on the subject of disk heating and empirical constraints on merger histories from the thickness and vertical profiles of observed disks. However, we show that it is no longer the case when the simplifying assumptions are relaxed \\citep[as in current, more detailed simulations and cosmological models; e.g.][]{velazquezwhite:disk.heating, ardi:subhalo.mergers,benson:heating.model, hayashichiba:subhalo.mergers, kazantzidis:mw.merger.hist.sim}. These assumptions are fine when one considers a system that is effectively a point particle at large radius in an extended (but relatively low-density) background halo, on a moderately circular orbit, such that there are no close passages of the satellite and primary and the entire regime of the orbit has a characteristic frequency much larger than the internal characteristic orbital frequency of either the primary or secondary -- in the regime where the satellite is far from the galaxy and has just been accreted near the halo virial radius, for example. And in a strictly instantaneous sense, they are not bad. But in a global sense, in terms of estimating the effects on a galactic disk from a merger with initial mass ratio $M_{1}/M_{2}$, these assumptions break down, and the result is not just a numerical correction, but a qualitative change in the character of the heating.  In modern $\\Lambda$CDM cosmologies, orbits tend to be highly radial, especially those orbits that actually pass near or through the disk. Circular orbits efficiently lose angular momentum to interactions and other substructures in halo formation, and even when they survive, the dynamical friction times required for the orbit to decay from near the virial radius are so large (of order the Hubble time) that they tend to survive far away from the central galaxy rather than merge with and/or heat the disk. Moreover, once the galaxy gets sufficiently close to the disk that it might effect some heating, the characteristic orbital frequencies are (by necessity) similar to the characteristic internal frequencies of the disk and halo. Because some of the galaxy will be in near-resonance with the secondary orbit, this allows much more efficient transfer of angular momentum from the orbital motion into internal motions, and even initially circular orbits become radial in just one or two passages. Because the orbital angular momentum is transferred quickly in a passage, the secondary cannot gently ``spiral in'' but must return from apocenter in a radial ``plunge'' -- the secondary effectively free-falls inwards (weakly decelerated by e.g.\\ dynamical friction at this point) and ``punches through'' the primary, where it executes a rapidly damping oscillation through the primary on nearly radial orbits. At this stage, violent relaxation scatters stars off the time-dependent potential and transfers the orbital energy to the central stellar distribution (the bulge or whatever the hottest stars are at the center of the galaxy), damping the orbit and completing coalescence in a couple of dynamical times. At the same time, once the secondary begins to penetrate the primary disk, it experiences rapid tidal mass loss and soon most of the mass has been tidally added to the extended halo and cannot further heat the disk.  For these reasons, the most useful approximation for disk heating in realistic cosmological situations is not a gradually decaying circular orbit of a rigid satellite, but rather something like the impact approximation: the heating is dominated by a single or a couple of radial oscillations at the initial mass, on a radial free-fall trajectory. We show that in this case, the heating is {\\em not} linear in the mass ratio $M_{2}/M_{1}$. Rather, the effective heating is second-order in the mass ratio, $\\propto (M_{2}/M_{1})^{2}$. We show that -- again for the reasons above -- this is a much more robust scaling.  We demonstrate this by comparing our analytic predictions to a variety of numerical simulations, which include a range in secondary mass ratios, assumed structural properties of the disk, halo, bulge, and satellite, gas content, redshift, orbital parameters, and orbital anisotropy and angular momentum content. Remarkably, these different simulations, from various independent groups, all agree with our analytic predicted scalings and with each other (when properly normalized) to within a factor $\\sim2-3$. It is straightforward to show that a linear scaling in mass ratio is not a good description of the simulation results (even if the normalization of the linear scaling is free -- i.e.\\ the scaling itself must be second-order).  This provides a physical reason why simulations increasingly find, contrary to historical expectations, that systems can undergo cosmologically expected merger histories and still yield relatively thin disks at $z=0$ in line with observations \\citep[see e.g.][]{velazquezwhite:disk.heating, ardi:subhalo.mergers,abadi03:disk.structure, benson:heating.model, robertson:cosmological.disk.formation,robertson:disk.formation, okamoto:feedback.vs.disk.morphology,hayashichiba:subhalo.mergers, governato:disk.formation,kazantzidis:mw.merger.hist.sim, younger:minor.mergers,villaloboshelmi:minor.mergers}. Many of these studies found that idealized numerical experiments yield less efficient disk heating than that predicted by \\citet{tothostriker:disk.heating}; the key contribution of our work is to outline the physical reasons for this difference and to demonstrate that these differences are {\\em not} a matter of e.g.\\ the normalization or numerical prefactors involved in an estimate of disk heating: heating is second-order in mass ratio ($\\mu$).  This {\\em qualitatively} changes the disk heating paradigm. Because the mass spectrum of mergers (in terms of halo mass) goes as ${\\rm d}N/{\\rm d}\\ln{\\mu}\\propto \\mu^{-1}$, if the heating per encounter were in fact linear ($\\propto \\mu$), then there would be more or less equal contributions to disk heating and thick disk formation from each logarithmic interval in mass ratio. Observationally, this would imply that no more than a few dwarf galaxies with properties similar to those of e.g.\\ the rather small Sagittarius dwarf \\citep[see][]{helmi04:sgr.stream,fellhauer06:sgr.stream.orbit} could have been accreted or destroyed by the Milky Way in the last 10\\,Gyr, in conflict with cosmological expectations and making the well-known ``missing satellite'' problem even more difficult to resolve.  In contrast, if the heating is second-order in $\\mu$, then the net heating is dominated by the few large mergers of relatively large mass ratios. This behavior has in fact been recognized in cosmological simulations, and yields better agreement with observations \\citep[see][]{ abadi03:disk.structure, robertson:cosmological.disk.formation,benson:heating.model, okamoto:feedback.vs.disk.morphology, governato:disk.formation,kazantzidis:mw.merger.hist.sim,read:thick.disk.cosmo.sims}. It also fundamentally enables the existence of very thin disks: since mergers of relatively large mass ratios $\\mu\\gtrsim0.1$ are rare, they are somewhat stochastic -- i.e.\\ some systems which have gone sufficient time or been sufficiently gas rich since their last $\\sim$\\,1:10 merger will exist and can therefore have very thin disks ($H/R\\lesssim0.1$) similar to the ``ultra-thin'' or ``super-thin'' disks observed \\citep{seth:disk.scale.heights, yoachim:disk.scale.heights}. If instead the same net heating came from mergers of mass ratios $\\mu\\sim0.001-0.01$ (where the expected number of mergers since $z\\sim1$ is of order several hundred per galaxy; making the existence of any galaxies without such mergers in their recent past an improbable event) as from mergers with mass ratio $\\mu\\sim0.1-1$, then it would be almost impossible to contrive realistic cosmological merger histories that could give rise to ultra-thin disks.  There are a number of other indirect implications of this result. For example, we demonstrate that this is of significant importance for the edge-on morphologies of disk plus bulge systems. Given some bulge formation in mergers, whether the induced heating in the surviving disk goes linearly or second-order in $\\mu$ produces a different relation between scale heights and bulge-to-disk ratios and, correspondingly, different morphologies in observed early-type disks. If realistic scale heights increased as efficiently as expected in the circular orbit, rigid secondary approximation, then disks of intermediate $B/T$ would necessarily be thicker than observed -- in fact, the Hubble sequence, viewed edge-on, would no longer have a characteristic ``bulge plus disk'' morphology and would instead appear as a sequence in ellipticity of characteristically boxy isophotal shapes (because a sufficiently thick disk appears in edge-on projection as somewhat rectangular). In contrast, disks appear to remain thin even in systems with significant (or dominant) bulges, giving rise to characteristically ``disky'' (``lemon-shaped'') isophotal shapes. The non-linearity of disk heating can explain why galactic morphological components form bimodal families (bulges and disks), as opposed to a sequence of increasingly ``spheroidal'' disks.  Applying these recalibrated models to the Milky Way and other massive disk galaxies, the implied constraints on the merger history from the thick and thin disk are dramatically changed and fully consistent with cosmologically expected merger histories. In detail, applying our scaling, calibrated against numerical simulations, to the observed age-velocity dispersion relation in the solar neighborhood, implies that the Milky Way could have survived as many as $\\sim5-10$ minor (mass ratio $\\sim$\\,1:10) mergers in the last 10\\,Gyr. If some of the disk thickness owes to other processes or there were some more violent events, this number might be a bit smaller, but if the Milky Way were more gas-rich in the past, then this number would in fact go up by a factor $1/(1-\\fgas)$. In either case, the median number of such mergers expected for a Milky-Way mass halo is only $\\sim3$ (with $\\sim90\\%$ of Milky-Way like systems having $<5-6$ such mergers in this interval). The observed vertical velocity dispersions of stars (and implied disk thickness) of the Milky Way is completely consistent with the cosmologically expected amount of minor merging in the last 10\\,Gyr, and in fact these constraints alone permit the Milky Way, far from being especially quiescent, to be in the upper cosmological quartile in terms of number of minor mergers since $z\\sim1-2$. Translating these constraints to a ``maximum'' mass ratio merger implies that the Milky Way could even have survived mergers of mass ratio $\\sim$\\,1:4-1:3 since $z\\sim2$, without producing too much heating or thick disk.  Because the heating is second order in mass ratio, there is almost no measurable heating effected by encounters with mass ratios $\\mu \\ll 0.1$, so there is no overall constraint based on the thickness of the disk or internal velocity dispersions on how much {\\em total} mass has been added to the Milky Way via merging. In contrast to the $<5\\%$ growth since $z\\sim2$ inferred by \\citet{tothostriker:disk.heating}, we find that the Milky Way could have accumulated up to $\\sim25-50\\%$ of its present mass in the last 10\\,Gyr by 1:10 mergers alone ($\\sim5-10$ individual 1:10 mergers). More interesting, because again smaller mass ratio mergers heat only to second order in $\\mu$, {\\em the Milky Way could easily, in fact, have grown by factors of several since $z\\sim2$}, provided this growth mostly came from relatively small mass ratio mergers $\\mu \\lesssim$\\,1:10. This is critical for our understanding of galaxy formation: at some level of sufficiently small mass ratios below $\\sim$\\,1:10-1:100, the distinction between accretion and mergers is increasingly blurred. Essentially all galaxies are expected to gain a large fraction of their mass and grow by large factors (up to an order of magnitude) in the last $\\sim10\\,$Gyr from such accretion events. If these small events heated as efficiently as major events (e.g.\\ if $\\sim30$ separate 1:100 mergers were as damaging to a disk as a 1:3 merger), there would be a fundamental difficulty for CDM models. In fact, owing to the inefficiency of disk heating, there is no tension here.  This is not, of course, to say that problems of disk formation and survival in a $\\Lambda$CDM Universe have been solved. There are still well-known problems related to the loss of angular momentum in disks (owing in part to over-merging), that tend to yield (in cosmological simulations) disks that are too concentrated and have too much bulge relative to observations. A great deal of theoretical effort has gone into studying this \\citep{weil98:cooling.suppression.key.to.disks, sommerlarsen99:disk.sne.fb,sommerlarsen03:disk.sne.fb, thackercouchman00,thackercouchman01,abadi03:disk.structure, governato04:resolution.fx,governato:disk.formation, robertson:cosmological.disk.formation, okamoto:feedback.vs.disk.morphology,scannapieco:fb.disk.sims}, with increasing agreement that some combination of improved resolution and proper implementation of feedback from supernovae and stellar winds are critical for the formation of realistic disks. These issues are not addressed or affected by our analysis here.  However, it is increasingly clear (as resolution improves such that thick disks can be properly resolved in the vertical direction) that these simulations do {\\em not} generally have the problem of producing disks that are much too thick relative to observations. In this paper, we have shown that there is a good physical reason for this, and that it is not a peculiarity of a couple of simulations or of an especially quiescent merger history, but rather is a robust and quite general expectation even for systems with violent merger histories and even in systems that will be bulge-dominated at $z=0$. We have derived a new, more generally applicable scaling of disk height with merger mass ratio, and shown that it qualitatively changes the expectations for how disk heating should proceed in a $\\Lambda$CDM universe, and generically leads to the expectation that disks are quite robust to being heated by minor mergers. Although there are still considerable improvements in models needed to explain when and how disks form and why the distribution of bulge-to-disk ratios is what is observed at $z=0$, there is no tension between the models and the heating and thickness of disks. In short, we have argued that there is no ``disk heating'' problem in $\\Lambda$CDM cosmologies, once a proper accounting of disk heating for realistic cosmological orbits and conditions is considered."
    },
    "0806/0806.0113_arXiv.txt": {
        "abstract": "{} {We present a catalog of point $\\gamma$-ray sources detected by the EGRET detector aboard the Compton Gamma Ray Observatory. We have used the whole $\\gamma$-ray dataset of reprocessed photons at energies above 100 MeV together with new Galactic interstellar emission models based on recent $CO$, $HI$, dark gas, and interstellar radiation field data. Two different assumptions have been used for the cosmic-ray distribution in the Galaxy to explore the resulting systematic uncertainties in source detection and characterization.} {We have used the same 2-dimensional maximum-likelihood detection method as for the 3rd EGRET catalogue.} {The revised catalogue lists 188 sources, 14 of which are marked as confused, compared to the 271 entries of the 3rd EGRET (3EG) catalogue. 107 former sources have not been confirmed because of the additional structure in the interstellar background. The vast majority of them were unidentified and marked as possibly extended or confused in the 3EG catalogue. In particular, we do not confirm most of the 3EG sources associated with the local clouds of the Gould Belt. Alternatively, we find 30 new sources with no 3EG counterpart. The new error circles for the confirmed 3EG sources largely overlap the previous ones, but several counterparts of particular interest that had been discussed in the litterature, such as Sgr A*, radiogalaxies and several microquasars are now found outside the error circles. We have cross-correlated the source positions with a large number of radio pulsars, pulsar wind nebulae, supernova remnants, OB associations, blazars and flat radiosources and we find a surprising large number of sources (87) at all latitudes with no counterpart among the potential $\\gamma$-ray emitters.} {} ",
        "introduction": "The Energetic Gamma-Ray Experiment Telescope (EGRET), which operated on board the Compton-Gamma Ray Observatory from April 1991 to May 2000, detected photons in the 20 MeV to 30 GeV range. The observation program made use of the large instrumental field of view (25$^{\\circ}$ in radius) to cover the whole sky and for in-depth studies of specific regions. The resulting exposure and flux sensitivity to point sources are therefore not uniform across the sky. The sensitivity threshold also varies because of the intense background emission that arises from cosmic-ray interactions with the interstellar gas and photon fields in the Milky Way. The minimum flux that EGRET could detect steeply rises with decreasing Galactic latitude. In order to detect point sources and assess their significance in these varying conditions, a 2-dimensional maximum-likelihood method using binned maps had been developed for the COS-B data (\\cite{pollock81}) and implemented for the EGRET one (\\cite{mattox96}). A first catalog using this method was published after 1.5 years of data (\\cite{fichtel94}), followed by the second one (\\cite{thompson95}) and its supplement (\\cite{thompson96}) after 3 years of data. Lamb \\& Macomb (1997) presented a catalog of sources detected above 1 GeV. The last EGRET catalog (hereafter 3EG, \\cite{hartman99}) comprised reprocessed data from April 1991 to October 1995 with the interstellar emission model from Hunter et al. (1997) and extragalactic background from Sreekumar et al. (1998). This version contained 271 point sources including a solar flare, the Large Magellanic Cloud, five pulsars, one radiogalaxy detection (Cen A), 66 high-confidence identifications of blazars (BL Lac objects and flat-spectrum radio quasars), and 27 lower-confidence blazar identifications. Because of the wide tails of the instrument point-spread function, seven potential artifacts were noted around the brightest sources and many sources were marked as confused or possibly extended.  The 3EG catalogue also contained 170 sources with no attractive counterpart at lower energy. About 130 of them remain unidentified as of today (see Grenier (2004) and references therein). Candidate counterparts that have been searched for include pulsars and their wind nebulae, supernova remnants, massive stars, X-ray binaries and microquasars, blazars and nearby radiogalaxies, luminous infrared and starburst galaxies, and galaxy clusters. It was also noticed (\\cite{grenier95,grenier00,gehrels00}) that the most stable unidentified sources are significantly correlated with the nearby Gould Belt, a system of massive stars and interstellar clouds that surrounds the Sun at a distance of hundreds of parsecs. The offset position of the Sun with respect to the Belt centre and the Belt inclination of 17$^{\\circ}$ to the Galactic plane indeed provides a useful spatial signature across the sky (\\cite{perrot03}).  EGRET went on observing for another 4.5 years after the 4 cycles used for the 3EG work. Its sensitivity was reduced because of the ageing gas in the spark chamber, but it gathered nearly ten percent more photons and saw several new variable sources. Several authors (\\cite{nolan03,sowards05}), however, noticed discrepancies between their studies and at least five 3EG sources. They failed to confirm sources and found others. The whole $\\gamma$-ray dataset and final instrument response functions have also been significantly reprocessed by the EGRET team in 2001. Furthermore, the spatial coverage of the $CO$ surveys has reached higher latitudes since 1999, finding new small $CO$ clouds (\\cite{dame01}). In parallel, new $HI$ surveys (\\cite{kalberla05}) have been completed to correct for the significant contamination of stray radiation in the older ones. Finally, an additional 'dark' gas component has been found in the Gould Belt clouds that significantly increases their mass and spatial extent (\\cite{grenier05}). The additional mass is structured into large envelopes around the dense $CO$ cores. They do not follow the $HI$ and $CO$ maps commonly used to trace atomic and molecular column-densities. So, the dark gas provides both $\\gamma$-ray intensity and structure that were not accounted for in the 3EG background model.  For all these reasons and in preparation of the new GLAST mission, it was necessary to revise the interstellar background model and to apply the EGRET detection method to the full nine years of data to build a new catalogue of sources above 100 MeV. In order to study the systematic uncertainties induced on source locations and fluxes by our limited knowledge of the intense interstellar background, we have applied the analysis to two different background models exploiting the same new interstellar data, but using different approaches to constrain the cosmic-ray gradient across the Galaxy. ",
        "conclusions": "We have searched for point-like sources in the reprocessed EGRET data from cycle 1 to 9 using new interstellar background models based on the most recent $HI$, $CO$, and dark gas data, as well as two different assumptions for the cosmic-ray distribution (the GALPROP diffusion model or a radial emissivity gradient fitted to the diffuse EGRET data). We have used the 3EG tools, likelihood method, procedure and significance threshold to detect sources, but have expanded the search to 3 different energy bands (above 100 MeV, 0.3-1 GeV, and above 1 GeV). The resulting number of detected sources has decreased by more than a third. Many unidentified sources, in particular among those spatially associated with the Gould Belt, are not confirmed as significant excesses. Their emission can be explained by the additional interstellar emission and its structure. Several interesting counterparts to 3EG sources, such as radiogalaxies, massive binaries, and microquasars, are now found outside the 95\\% confidence region. We have cross-correlated the new source positions with large pulsar, supernova remant, pulsar wind nebulae, OB associations, and radiosource catalogues, yet half the sample has no attractive counterpart among the potential $\\gamma$-ray emitters. 30 new possible $\\gamma$-ray sources have also been found.  This EGR catalog will be available in fits format at the Strasbourg astronomical Data Center (CDS) and in ASCII format at http://www.aim.univ-paris7.fr/EGRET\\_catalogue/home.html"
    },
    "0806/0806.2116_arXiv.txt": {
        "abstract": "The coincidence problem of late cosmic acceleration constitutes a serious riddle with regard to our understanding of the evolution of the Universe. Here we argue that this problem may someday be solved -or better understood- by expressing the Hubble expansion rate as a function of the ratio of densities (dark matter/dark energy) and observationally determining the said rate in terms of the redshift. ",
        "introduction": "The coincidence problem of late cosmic acceleration, i.e., the fact that the density values of dark matter, $\\rho_{m}$, and dark energy, $\\rho_{\\phi}$, are of the same order precisely today, constitutes a serious challenge to our understanding of the evolution of the Universe. The otherwise rather successful vacuum-cold dark matter ($\\Lambda$CDM) model offers no explanation. In that model, the cosmological constant must be fine-tuned many orders of magnitude for the ratio $r \\equiv \\rho_{m}/\\rho_{\\phi}$ to be of order unity nowadays. This partly explains why many authors turned to models in which the dark energy density is no longer a constant but evolves with expansion -see e.g. \\cite{edmund} for references.  One way to alleviate the coincidence problem is to provide a mechanism by virtue of which  $r$ tends to a constant at late times \\cite{int1, int2} or at least makes it vary slower than the scale factor at present time \\cite{softk}. This can be achieved by introducing an interaction between dark matter and dark energy regardless if the latter obeys the dominant energy condition  or not \\cite{int3}.   Most cosmological models assume, for the sake of simplicity, that matter and dark energy interact only gravitationally. In the absence of an underlying symmetry that would suppress a matter-dark energy coupling (interaction) there is no {\\em a priori} reason to dismiss it. Further, the coupling is not only likely but inevitable \\cite{jerome} and its introduction is not more arbitrary than assuming it to vanish.  Moreover, it has been forcefully argued that the interaction reveals itself  when one applies the Layzer-Irvine equation \\cite{layzer} to galaxy clusters \\cite{elcio}. Likewise, the cross-correlation of galaxy catalogs with cosmic microwave background (CMB) maps seems to observationally favor the interaction \\cite{isw}. Nevertheless, it still remains hypothetical and more abundant, varied, and accurate measurements are needed before its existence (or non-existence) can be established observationally.  Cosmological models where dark matter (DM) and dark energy (DE) do not evolve separately but interact with each other were first introduced to justify the small value of the cosmological constant \\cite{wetterich} and currently there is a growing body of literature on the subject -see, e.g. \\cite{list} and references therein. Further, in some holographic models of dark energy the accelerated expansion can be traced to the interaction \\cite{wd}. Recently, various proposals at the fundamental level, including field Lagrangians, have been advanced to account for the coupling \\cite{piazza}.  Obviously, a mechanism that makes  $r$ tend  to a constant today or decrease its rate to a lower value than the scale factor expansion rate ameliorates the coincidence problem significantly, but it does not solve it in full. To do so the said mechanism must also achieve $r_{0} \\sim {\\cal O}(1)$. We are not aware of any mechanism able to fulfill the latter.  It is the view of the present authors that, on the one hand, $r_{0}$ ought to be understood -at least, for the time being- as an input parameter much in the same way as other key observational quantities -say, the present value of the cosmic background radiation temperature, the cosmological constant $H_{0}$, the age of the Universe $t_{0}$, or the ratio between the number of baryons and photons. And, on the other hand, its value must be closely related to $H_{0}$ or, for that matter, to $t_{0}$. Accordingly, if one successfully express $r$ in terms of Hubble rate, $H$, the problem of explaining the $r_{0}$ value reduces to the problem of explaining the present value of $H$. This is interesting because while we are still lacking model-independent data of $r$ at different redshifts, two sets of observational data values -albeit scarce- of $H$ at various redshifts are now available. The first data set was obtained by computing differential ages of passively evolving galaxies in the redshift range $0.1 < z < 1.8 \\;$ \\cite{simon}. The other data set is based on measurements of 192 supernovae type Ia (SN Ia) and 30 radio galaxies up to redshift $1.2$ \\cite{daly1}.  Before proceeding, we wish to emphasize that, to the best of our knowledge, no model-independent data of the evolution of the energy densities of matter and dark energy do exist. To obtain such evolution (e.g., with respect to redshift) one must specify some cosmological model. It is to say, the field equations together with the equation of state of dark energy and some assumption about whether the components interact or not. If they are assumed to interact, one must also provide an expression for the interaction. ",
        "conclusions": "} The current value of the ratio $r = \\rho_{m}/\\rho_{\\phi} \\,$ is of order unity (more precisely, about $3/7 \\,$ -see e.g. \\cite{spergel}). This represents the ``coincidence problem\" that the $\\Lambda$CDM model seems unable to account for. Nowadays, $r_{0}$ cannot be derived from first principles and like some other key cosmic quantities it has to be considered an input parameter. It seems obvious, nevertheless, that in Einstein relativity it must be closely tied to the present value of the Hubble expansion rate, $H_{0}$, another quantity whose value we cannot explain but that contrary to $r_{0}$  it does not mean any puzzle. Therefore, a possible avenue to someday solve the coincidence problem may well be to express $r$ as a function of $H$ (or vice versa) in models such that $r$ at late times either tends to constant or varies slowly enough (i.e., slower than in the $\\Lambda$CDM model). Thus, the problem of explaining $r_{0}$ reduces to the less acute problem of explaining $H_{0}$.  In this paper we have taken a first step in that direction. Specifically, we have considered two phenomenological models, previously studied in the literature, in which dark energy and dark matter do not evolve separately but interact with each other in such a way that either $r$ stays constant at late times (case $(i)$) or varies very slowly (case $(ii)$) and, in both instances, we have related $H$ to $r$. Since model-independent data relating these two quantities are not currently available we have expressed $r$ in terms of $z$ [Eqs. (\\ref{rz1}) and (\\ref{rfz}), respectively] and resorted to the observational data of Refs. \\cite{simon} and \\cite{daly1} that link $H$ to $z$. Regrettably, given the scarcity and dispersion of the data (cf. Figs. \\ref{fig:hz1} and \\ref{fig:hz2}) they are unable to discriminate between models .   Nevertheless, it is fair to say  that in order not to spoil the well established primeval nucleosynthesis scenario, $\\Omega_{\\phi}$ should not exceed $5\\%$ at that early epoch -see Bean, Hansen, and Melchiorri  \\cite{rbean}. At any rate, as Figs. \\ref{fig:rvsz1} and \\ref{fig:rvsz2} show, the models considered in this work are consistent with this constraint which, in any case, was determined using specific non-interacting dark energy models.   It is to be hoped that in the not far future abundant and accurate observational data, in an extended redshift range, may help determine $H$ in terms of $z$  and  solve (or, at least, help to better understand) the coincidence problem of late acceleration. Yet, it may well be that the $\\Lambda$CDM model comes to fit the data better than any interacting model. In such a case, we should conclude that the coincidence problem is just that: a coincidence and, therefore, no longer a problem."
    },
    "0806/0806.0325_arXiv.txt": {
        "abstract": "We present a morphological study of the two richest superclusters from the 2dF Galaxy Redshift Survey (SCL126, the Sloan Great Wall, and SCL9, the Sculptor supercluster). We use Minkowski functionals, shapefinders, and galaxy group information to study the substructure of these superclusters as formed by different populations of galaxies. We compare the properties of grouped and isolated galaxies in the core region and in the outskirts of superclusters.  The fourth Minkowski functional $V_3$ and the morphological signature $K_1$- $K_2$ show a crossover from low-density morphology (outskirts of supercluster) to high-density morphology (core of supercluster) at mass fraction $m_f \\approx 0.7$. The galaxy content and the morphology of the galaxy populations in supercluster   cores and outskirts is different. The core regions contain a larger fraction of early type, red galaxies, and richer groups than the outskirts of superclusters. In the core and outskirt regions   the fine structure of the two prominent superclusters as delineated by galaxies from different populations also differs. The values of the fourth Minkowski functional $V_3$ show that in the supercluster SCL126 the population of early type, red galaxies is more clumpy than the population of late type, blue galaxies, especially in the outskirts of the supercluster. In the contrary, in the supercluster SCL9, the clumpiness of the spatial distribution of galaxies of different type and color is quite similar in the outskirts of the supercluster, while in the core region the clumpiness of the late type, blue galaxy population is larger than the clumpiness of the early type, red galaxy population.  Our results suggest that both local (group/cluster) and global (supercluster) environments are important in forming galaxy morphologies and colors (and determining the star formation activity).  The differences between the superclusters indicate that these superclusters have different evolutional histories. ",
        "introduction": "\\label{sec:intro}    Huge superclusters which may contain several tens of rich (Abell) clusters are the largest coherent systems in the Universe with characteristic dimensions of up to 100~\\Mpc. \\footnote{$h$ is the Hubble constant in units of 100~km~s$^{-1}$~Mpc$^{-1}$.} As they are very large, and dynamical evolution takes place at a slower rate for larger scales, superclusters have retained memory of the initial conditions of their formation, and of the early evolution of structure \\citep{kofman87}.  While we might be able to explain the structure and properties of most (average) superclusters, explaining rich superclusters is still a challenge. To start with, even their existence is not well explained by the main contemporary structure modeling tool, numerical simulations. There is a number of rich superclusters in our close cosmological neighbourhood -- we shall list them in a moment -- but the comparison of the luminosity functions of observed and simulated  (Millennium) superclusters shows that the fraction of really rich superclusters in simulations is much lower than in observed smaples \\citep{e06}. The extreme cases of observed objects usually provide the most stringent tests for theories; this motivates the need for a detailed understanding of the richest superclusters.  The richest relatively close superclusters are the Shapley Supercluster \\citep[][and references therein]{proust06} and the Horologium--Reticulum Supercluster \\citep{rose02,fleenor05,e2003}.  The two superclusters studied in this paper, the Sloan Great Wall and the Sculptor supercluster, belong also to this category of very rich superclusters.  The formation of rich superclusters had to begin earlier than smaller structures; they are sites of early star and galaxy formation \\citep[e.g.][]{mob05}, and the first places where systems of galaxies form \\citep[e.g.][]{ven,ouch}. Thus future deep surveys like the ALHAMBRA Deep survey \\citep{moles05} are likely to detect (core) regions of rich superclusters at very high redshifts. The supercluster environment affects the properties of groups and clusters located there \\citep{e2003,pl04}. Rich superclusters contain high density cores which are absent in poor superclusters \\citep[][hereafter Paper III]{e2007a}. The fraction of X-ray clusters in rich superclusters is larger than in poor superclusters \\citep[][hereafter E01]{e2001}, and the core regions of the richest superclusters may contain merging X-ray clusters \\citep{rose02,bar00}. The richest superclusters are more filamentary, less compact and more asymmetrical than poor superclusters \\citep[][hereafter Paper II]{e07b}. Moreover, as we noted above, the fraction of very rich superclusters in observed catalogues is larger than models predict \\citep{e06}.  In the present paper we continue the study of superclusters selected on the basis of the 2dF Galaxy Redshift Survey. Our paper is devoted to a detailed study of the two richest superclusters in the 2dF Galaxy Redshift Survey. We chose them from the catalogue of superclusters of the 2dFGRS by \\citet[][hereafter Paper I]{e07c}.  These are: the supercluster SCL126 in the Northern sky, and the supercluster SCL9 (the Sculptor supercluster) in the Southern sky (see Fig.~\\ref{fig:silvestre}).  \\begin{figure*}[ht] \\centering \\resizebox{0.9\\textwidth}{!}{\\includegraphics*{f1.eps}} \\caption{The nearest, almost complete part of the 2dF GRS survey (up to 400 \\Mpc). The two richest superclusters in this survey are shown. The disk showing the orientation of the galactic plane helps to allocate the 2dF survey in space; SCL 126 is located towards the Galactic North pole (left, red color), SCL 9 is at the right side of the figure (blue color). } \\label{fig:silvestre} \\end{figure*}  The supercluster SCL126 is the most prominent supercluster defined by Abell clusters in the Northern 2dFGRS  (SCL126 in the catalog of superclusters by E01, N152 in Paper I) in the direction of the Virgo constellation. This supercluster has also been called the Sloan Great Wall \\citep{hoyle02,vogeley04,gott05,nichol06}.  The presence of this supercluster affects the measurements of the correlation function \\citep{croton04}, and of the genus and Minkowski functionals of the SDSS and 2dF redshift surveys \\citep{park05,saar06}. The ``meatball'' shift in the measurements of the topology in the SDSS data is partly due to this supercluster \\citep{gott06}.  The Sculptor supercluster, the most prominent supercluster in the Southern strip of the 2dFGRS, is among the three richest superclusters in E01; it contains 25 Abell clusters, six of these are also X-ray clusters. \\citet{zap04} found evidence about the presence of warm-hot diffuse gas there, which is associated with the inter-cluster galaxy distribution in this supercluster.  In a recent paper \\citep[][hereafter RI]{e07d} we studied the morphology of rich superclusters (their shape and internal structure), using Minkowski functionals and shapefinders. Our calculations in RI showed that the morphology of the richest superclusters from the 2dFGRS is different from each other: the supercluster SCL126 resembles a very rich filament (wall) with a high density core region, while the supercluster SCL9 can be described as a collection of spiders (a multi-spider), consisting of a large number of cores connected by relatively thin filaments.   The main aim of the present paper is to understand whether the differences in the overall morphology of these two rich superclusters under study are also reflected by their fine structure as determined by the distribution of galaxies of different luminosity, color and spectral type in core region and in outskirts of superclusters.  The first study to show that in a supercluster early and late type galaxies trace the structure of the supercluster in a different manner was performed by \\citet{gio86} who demonstrated that in the Perseus supercluster, elliptical galaxies are mainly located along the central body of the supercluster, while spiral galaxies are distributed in the outer regions of the supercluster. The presence of a large-scale segregation of galaxies of different type in nearby superclusters was shown also by \\citet{e87}. In Paper III we showed that rich superclusters have a larger fraction of  red, non-star-forming galaxies than poor superclusters. Recently, galaxy populations have been studied in core regions of some very rich superclusters (\\citet{hai06} -- in the Shapley supercluster, \\citet{pr05} -- in the Pisces-Cetus supercluster). These studies showed that rich clusters in the core regions of superclusters contain a large fraction of passive galaxies, while actively star forming galaxies are located between the clusters.  Thus studies of rich superclusters which contain a large variety of environments and possibly a variety of evolutionary phases give us a possibility to study the properties of cosmic structures and the properties of galaxies therein in a consistent way, helping us to understand the role of environment in galaxy evolution.   One method to quantify the  structure of superclusters is to use the Minkowski functionals. This type of study has been called morphometry by \\citet{hik03}. Minkowski functionals, genus and shapefinders, defined on the basis of Minkowski functionals have been used earlier to study the 3D topology of the large scale structure from 2dF and SDSS surveys \\citep{saar06,hik03,park05,gott06,jam06} and to characterize the morphology of superclusters \\citep{sah98,sheth03, sss04,bpr01,kbp02,bas03,bas06} from observations and simulations. These studies concern only the ``outer'' shapes of superclusters and do not treat their substructure. We expanded this approach in Paper RI by using the Minkowski functionals and shapefinders to analyze the full density distribution in superclusters, at all density levels.   The fourth Minkowski functional $V_3$ (the Euler characteristic) gives us the number of isolated clumps (or voids) in the region \\citep{saar06}, meaning that we can use it to study the clumpiness of the galaxy distribution inside superclusters -- the fine structure of superclusters. We calculate the fourth Minkowski functional $V_3$ for galaxies of different populations in superclusters for a range of threshold densities, starting with the lowest density used to determine superclusters, up to the peak density in the supercluster core.  This analysis shows how the fourth Minkowski functional can be used in studies of the fine structure of superclusters as delineated by different galaxy populations. This study is of an exploratory nature since this is the first time when this method are used for studies of the fine structure of galaxy populations of individual superclusters.  Employing Minkowski functionals, we can see in detail how the morphology of superclusters is traced by galaxies of different type.  In our analysis we use also the shapefinders calculated on the basis of the Minkowski functionals. In addition, we compare the galaxy content of groups of different richness in core regions and in outskirts of superclusters.   The paper is composed as follows. In Section 2 we describe the galaxy data, the supercluster catalogue and the data on the richest superclusters. In Sect. 3 we compare the overall galaxy content of the superclusters. In Sect. 4 we describe the use of the fourth Minkowski functional (the Euler characteristic) to study the fine structure of the superclusters as delineated by different galaxy populations. In Sect. 5 we study the galaxy content in the core regions and in the outskirts of the superclusters, and compare galaxy populations in groups of various richness. In the last sections (6 and 7) we discuss our results and give the conclusions. In the Appendix we give a definition of the Minkowski functionals and shapefinders, and describe different kernels used to calculate the density fields of superclusters. ",
        "conclusions": "We have presented a morphological study of the two richest superclusters from the 2dF Galaxy Redshift Survey. We studied the internal structure and galaxy populations of  these superclusters.  Our main conclusions are the following.  \\begin{itemize}  \\item{} The values of the fourth Minkowski functional $V_3$, which contain information about both the local and global morphology, show   the fine structure of superclusters as determined by galaxies from different populations. The fourth Minkowski functional $V_3$ and the morphological signature $K_1$-$K_2$ show a crossover from low-density morphology (outskirts of supercluster) to high- density morphology (core of supercluster) at a mass fraction $m_f \\approx 0.7$.   \\item{} In the supercluster SCL126, the functional $V_3$ shows that the number of clumps in the distribution of red   galaxies is larger than the number of clumps determined by blue  galaxies, especially in the outer regions of the supercluster, where the maximum values of $V_3$ are, correspondingly, 8 for red galaxies and 5 for blue galaxies. Thus in the outskirts of the supercluster, blue galaxies form a more homogeneous population than  red galaxies. In the core region, the clumpiness of galaxy populations is of the same order.  \\item{} In the supercluster SCL9, the values of $V_3$ are large for both early type and red galaxies, and late type and blue galaxies. In the outskirts of the supercluster the differences between the clumpiness of galaxy populations are small (the values of $V_3$ are the same, 7--8 for both early and late type galaxies), while in the core region of this supercluster the clumpiness of the late type, blue galaxy population is larger (the maximum value of $V_3$ is 17) than the clumpiness of the early type, red galaxy population, where the maximum $V_3 = 12$.  \\item{} Our superclusters contain galaxies with mixed classifications: these galaxies show spectra typical to late type galaxies and red colors, or they are early type galaxies with blue colors. These galaxies are mostly located at intermediate densities in the outskirts of our superclusters, being members of poor groups or do not belonging to any group. Due to these galaxies, at intermediate mass fractions the curves of the fourth Minkowski functional $V_3$ for early type galaxies and for red galaxies, as well as for late type galaxies and blue galaxies, differ in some details.   \\item{} Groups in high-density cores of superclusters are richer than in lower-density (outer) regions of superclusters. In high-density cores, groups contain relatively more  early-type, red galaxies than the groups in lower-density regions in superclusters, where there are more late type, blue galaxies.  Therefore, both the richness of a group and its galaxy content depend on the large scale environment where it resides.   \\item{} In cores of superclusters, the fractions of early type, red galaxies is larger than these fractions in outskirts. In the supercluster SCL126 the morphological segregation of red and blue galaxies is stronger than in the supercluster SCL9.  In SCL126, the most luminous galaxies in rich groups have larger luminosities than most luminous galaxies in poor groups, while in SCL9 the luminosity of the brightest galaxies in rich and poor groups is comparable.   \\item{} The differences in overall morphology, fine structure and galaxy content of the supercluster suggest that there are differences in their evolutional history which affect their present-day properties.  \\end{itemize}  Our study shows the importance of the role of superclusters as a high density environment that affects the evolution and the present-day properties of their member galaxies and the groups/clusters of galaxies that constitute the supercluster.  The forthcoming Planck satellite observations will determine the anisotropy of the cosmic background radiation with unprecedented accuracy and angular resolution. As a by-product, Planck measurements will provide an all-sky survey of massive clusters via the Sunyaev-Zeldovich (SZ) effect. For the Planck project, detailed information of supercluster properties is important, helping to correlate the SZ-signals with the imprints of local superclusters. As a continuation of the present work, we are preparing supercluster catalogues for the Planck community."
    },
    "0806/0806.2320_arXiv.txt": {
        "abstract": "We constrain the physical nature of dark matter using the newly identified massive merging galaxy cluster {\\bbullet}. As was previously shown by the example of the Bullet Cluster (1E0657-56), such systems are ideal laboratories for detecting isolated dark matter, and distinguishing between cold dark matter (CDM) and other scenarios (e.g.\\ self-interacting dark matter, alternative gravity theories). {\\bbullet} consists of two merging subclusters of similar richness at $z=0.586$.  We measure the distribution of X-ray emitting gas from Chandra X-ray data and find it to be clearly displaced from the distribution of galaxies. A strong (information from highly distorted arcs) and weak (using weakly distorted background galaxies) gravitational lensing analysis based on Hubble Space Telescope observations and Keck arc spectroscopy confirms that the subclusters have near-equal mass. The total mass distribution in each of the subclusters is clearly offset (at $>4\\sigma$ significance) from the peak of the hot X-ray emitting gas (the main baryonic component), but aligned with the distribution of galaxies. We measure the fractions of mass in hot gas ($0.09^{+0.07}_{-0.03}$) and stars ($0.010^{+0.007}_{-0.004}$), consistent with those of typical clusters, finding that dark matter is the dominant contributor to the gravitational field.  Under the assumption that the subclusters experienced a head-on collision in the plane of the sky, we obtain an order-of-magnitude estimate of the dark matter self-interaction cross-section of $\\sigma/m < 4\\cs$, re-affirming the results from the Bullet Cluster on the collisionless nature of dark matter. ",
        "introduction": "\\label{sec:intro} One of the most important quests in cosmology is to understand the formation and evolution of galaxies and galaxy clusters. While the standard paradigm of cold dark matter (CDM) within a $\\Lambda$CDM cosmology works well overall, discrepancies between the predictions of this model and observations remain (such as the abundance of substructure in halos, and the shapes of dark matter profiles, especially in dwarf systems). These issues can be resolved in several ways, for example by invoking warm and/or self-interacting dark matter.  Mergers of galaxy clusters are excellent laboratories to test such ideas. Not only are cluster mergers the most energetic events since the Big Bang, but, as demonstrated by a recent study of the Bullet Cluster (1E0657$-$56; \\citealp{clowe04,bradac06,clowe06}), they provide a unique avenue by which to probe the self-interaction cross-section of dark matter particles and test alternative gravity theories.  In clusters with recent major merger activity, the positions of the dark matter and main baryonic component (the X-ray emitting gas) can become temporarily separated.  This occurs because the gas is collisional and experiences ram pressure, whereas galaxies and (presumably) dark matter are effectively collisionless. A merging direction perpendicular to the line-of-sight maximizes the apparent separation, allowing observers to independently study the main baryonic and dark matter components. In the example of the Bullet Cluster, \\citet{clowe04}, \\citet{bradac06}, and \\citet{clowe06} used X-ray and gravitational lensing data to derive mass maps for both the baryonic and dark matter components, clearly showing the presence and dominance of dark matter. In agreement with the standard CDM paradigm, these observations indicate that dark matter has to be in a near-collisionless form, with a scattering cross-section $\\sigma/m < 0.7\\cs$ \\citep{randall07}.  A recent study of the merging cluster A520 by \\citet{mahdavi07} rests less easily within this simple picture. For A520, \\citet{mahdavi07} find a central mass peak that contains gas and dark matter but no significant galaxy concentration. These authors argue that such a configuration could occur if dark matter is collisional with a self-interaction cross-section of $3.8 \\pm 1.1\\cs$, a result that is inconsistent with the Bullet Cluster limits.  However, the statistical significance of the result for A520 is not high and the dynamics of the cluster are more complicated than {\\bullet}, with A520 appearing to be in the late stages of multiple, complex merger events.  It has also been argued that the optical and X-ray observations of {\\bullet} might be explained without standard CDM, by instead modifying the theory of gravity (and thus the effects of gravitational lensing). Models using the modified Newtonian Dynamics (MOND) theory \\citep[e.g.][]{milgrom83,bekenstein04} have, however, so far achieved only limited success in reproducing the weak lensing observations (e.g.\\ \\citealp{angus06, brownstein07}) and are unable to explain the observed strong lensing. Moreover, the application of MOND models to galaxy clusters in general still requires the presence of large masses of invisible matter, which are typically argued to be a mixture of massive neutrinos and cold baryons \\citep[e.g.][]{angus07}.  Clearly, a satisfactory resolution to this debate requires further observations of massive, merging clusters; preferably those with simple, two-body configurations. Using deep, ground-based optical imaging and short, snapshot Chandra exposures of clusters in the all-sky, X-ray flux-limited MAssive Cluster Survey (MACS) \\citep{ebeling01, ebeling07}, we have identified the X-ray brightest, most massive, obvious post-merging clusters in the redshift range $0.3<z<0.7$. From this search, {\\bbullet} emerges as a massive, merging cluster with an apparently simple geometry, consisting of two large subclusters of similar richness, both at redshift $z=0.586$, colliding in approximately the plane of the sky.  Owing to their high masses and favorable redshifts, MACS clusters are ideally suited to gravitational lensing studies. Here we use new, multi-color Hubble Space Telescope (HST) data to determine a projected mass map for {\\bbullet}, using both strong and weak lensing information. When combined with the Chandra data, this allows us to map the distributions of dark matter, hot X-ray emitting gas, and stars. The dark matter is resolved into two main clumps, centered on the galaxy populations.  The Chandra data reveal a single, clear X-ray peak located between the two optical/lensing subclusters. The observed offset between the locations of the hot gas and galaxies/lensing mass is in good agreement with the expectations for standard collisionless cold dark matter.  Both {\\bullet} and {\\bbullet} provide direct evidence for the existence of dark matter and can be used to estimate the self-interaction cross-section.  This paper is structured as follows. Section~\\ref{sec:data} describes the optical and X-ray data and reduction procedures. Section 3 discusses the dynamics of the cluster, based on optical spectroscopy. Section~\\ref{sec:swdata} presents the strong and weak gravitational lensing analysis and the mass reconstruction based on these data. Section~\\ref{sect:xray_analysis} discusses the X-ray data analysis. Section~\\ref{sec:stellarmass} details the results on the stellar and X-ray gas mass fractions. In Section~\\ref{sec:bardm} we present our main result on the relative distributions of dark and baryonic matter in the cluster, and determine rough limits on the dark matter self-interaction cross-section. Our conclusions are summarized in Section~\\ref{sec:conclusions}.  Throughout the paper we assume a $\\Lambda$CDM cosmology with $\\Omega_{\\rm m} = 0.3$, $\\Omega_{\\Lambda} = 0.7$, and Hubble constant $H_0 = 70 {\\rm \\ km \\: s^{-1}\\:\\mbox{Mpc}^{-1}}$. At a redshift $z=0.586$, the length scale is 6.61 kpc/arcsec. ",
        "conclusions": "\\label{sec:conclusions}  We have presented new data for the post-merging galaxy cluster {\\bbullet}. This system exhibits many similar properties to the Bullet Cluster {\\bullet}, although it does not contain a low-entropy, high density hydrodynamical `bullet'. HST and Chandra data have allowed us to derive a high-resolution total mass map, as well as the X-ray gas mass and stellar mass distributions in the cluster. Our main conclusions are as follows: \\begin{enumerate} \\item The total mass distribution in {\\bbullet} shows two distinct peaks. Both are clearly offset from the main baryonic component (the hot gas) as traced by the X-ray emission, at $>4\\sigma$ significance (for each of the peaks). There is no significant offset between the peaks of the total mass distribution and the galaxies. \\item We have measured the projected, enclosed mass of the three main mass constituents (dark matter, gas, and stars) within 500 kpc of the X-ray center. The gas-to-total mass ratio ($9^{+7}_{-3}\\%$) and the mass-to-light ratio ($1.0^{+0.7}_{-0.4}\\%$) are typical for a massive cluster. \\item The mass-to-light ratios $M/L_{\\rm K}$ ($< 300\\mbox{kpc}) = 70^{+40}_{-50}$ and $100^{+40}_{-60}$ for the SE and NW peak, and the stellar-to-total mass ratios ($1.0^{+0.7}_{-0.4}\\%$) in these regions are consistent with values for other massive clusters, arguing that there was no significant loss of dark matter from either subcluster during the collision. \\item The majority of the mass is spatially coincident with the galaxies which implies, just as for the Bullet Cluster, that the cluster must be dominated by a relatively collisionless form of dark matter. We obtain an estimate for the self-interaction cross-section $\\sigma/m < 4\\cs$. \\end{enumerate}  {\\bbullet} is a remarkable cluster with properties that allow us to improve our understanding of the details of cluster formation and evolution, and the nature of dark matter. Our results can be further improved with additional strong lensing information to better constrain the normalization of the mass model and better X-ray data to improve the determination of the merger speed, and hence improve the limits on the dark matter cross-section. Discovering more clusters like {\\bbullet} and {\\bullet} will also improve the systematic uncertainties on dark matter self-interaction cross-section inherent to the study of a small number of these objects. \\vspace{0.3cm}"
    },
    "0806/0806.2334_arXiv.txt": {
        "abstract": "We introduce the concept of multi-point propagators between linear cosmic fields and their nonlinear counterparts in the context of cosmological perturbation theory.  Such functions express how a non-linearly evolved Fourier mode depends on the full ensemble of modes in the initial density field. We identify and resum the dominant diagrams in the large-$k$ limit, showing explicitly that multi-point propagators decay into the nonlinear regime at the same rate as the two-point propagator. These analytic results generalize the large-$k$ limit behavior of the two-point propagator to arbitrary order. We measure the three-point propagator  as a function of triangle shape in numerical simulations and confirm the results of our high-$k$ resummation.  We show that any $n-$point spectrum can be reconstructed from multi-point propagators, which leads to a physical connection between nonlinear corrections to the power spectrum at small scales and higher-order correlations at large scales. As a first application of these results, we calculate the reduced bispectrum at one-loop in renormalized perturbation theory and show that we can predict the decrease in its dependence on triangle shape at redshift zero, when standard perturbation theory is least successful. ",
        "introduction": "Detailed understanding of the development of gravitational instability  is one of the central issues for the study of structure formation in observational cosmology. Although the details of the formation of objects are affected at small scales by the presence of baryonic matter, in the context of a dark matter dominated universe it is thought that the global properties of the large-scale matter distribution is determined to a large extent by the collisionless dynamics of dark matter. Therefore, the Vlasov equation, i.e. the collisionless limit of the Boltzmann equation, describes the dynamics we are interested in (see e.g.~\\cite{1980lssu.book.....P,2002PhR...367....1B} for details). This is for instance what pure N-body cosmological simulations attempt to solve. At large scales, where orbit crossing is insignificant, the Vlasov equation reduces to the dynamics a pressureless perfect fluid.  Although the problem is well posed, its resolution beyond the linear regime is still largely an open question. Some ideas have been developed in the literature (see e.g.~\\cite{2002PhR...367....1B} and references therein) but to make predictions beyond the linear regime for  the density power spectrum, the standard approach in the literature has been to resort to semi-analytic prescriptions obtained by fitting to numerical simulations. These fitting formulae, e.g.~\\cite{1996MNRAS.280L..19P,2003MNRAS.341.1311S}, originate either from the near universal transform advocated in \\cite{1991ApJ...374L...1H} or are based on an even more empirical construction, the halo model~\\cite{2002PhR...372....1C}.  These prescriptions however offer predictions for the power spectrum with relatively low accuracy, at the level of 5-10\\%, and are insecure in cases of non standard cosmological models. With the advent of precision observations of the large-scale structure, in particular in the context of the search for dark energy signatures in the growth of structure, this is clearly insufficient~\\cite{2005APh....23..369H}. It then becomes crucial to build reliable, and well controlled, prescriptions for describing the nonlinear evolution of the power spectrum and higher order correlation functions such as the bispectrum. For the latter there is not much published work on the expected behavior beyond tree order, although fitting formulae analogous to the power spectrum case have been proposed~\\cite{2001MNRAS.325.1312S,2005ApJ...619..667H,2007arXiv0707.1594P}. Recently, a promising new approach to fitting measurements of statistics in N-body simulations and interpolating to other cosmologies has been developed in~\\cite{2007astro.ph..2348H}.  With the Renormalized Perturbation Theory (hereafter RPT) formalism introduced in \\cite{2006PhRvD..73f3519C},  a new approach emerged for the construction of perturbation schemes applicable to such  observables. It has been successfully applied to the two-point propagator~\\cite{2006PhRvD..73f3520C}, and to the density power spectrum in the weakly non-linear regime~\\cite{2008PhRvD..77b3533C}. These results show in particular how the acoustic peaks are affected by the mode-coupling effects in the growth of structure. This reformulation of the gravitational perturbation theory also finds alternative approaches~\\cite{2007PhRvD..75d3514M,2007A&A...465..725V,2007JCAP...06..026M,2008ApJ...674..617T,2007arXiv0711.2521M}, which give further insights into resummation schemes introduced by RPT.  The aim of this article is to show that some of the core results that have been originally discovered for the two-point propagator in~\\cite{2006PhRvD..73f3520C} can also be obtained in the case of multi-point propagators, thus extending exact known results regarding the evolution of cosmic gravitational instability.  The plan of the paper is the following. We first succinctly recall the general formalism and the results that have already been obtained at the two-point level~\\cite{2006PhRvD..73f3520C}, involving the cross-correlation between fields at the initial and final conditions. As shown in previous studies, the behavior of this correlator at large-$k$ is of crucial importance and provides us with solid ground for the expansion of the correlation functions of interest~\\cite{2008PhRvD..77b3533C}. We then show that the three-point propagator, namely the cross-correlation between the final density field and two of the initial modes, can also be computed explicitly in the high-$k$ limit. This quantity reflects how the final density is sensitive to the two mode-mode couplings. We then extend this result to propagators of arbitrary number of points. This is the central result of this paper.  To contrast our results with numerical simulations, we extend the method developed in~\\cite{2006PhRvD..73f3520C} to measure multi-point propagators. This allows us to compute the three-point propagator in numerical simulations against our predictions in the large-$k$ regime. We finally comment on the relevance of these results for the construction of power spectrum and bispectrum resummation schemes, and provide a first application to the case of the bispectrum. ",
        "conclusions": "We generalized the notion of two-point propagator to arbitrary multi-point propagators. We derived the one-loop corrections to the three-point propagator, and were able to resum the dominant class of diagrams contributing to multi-point propagators in the large-$k$ limit, extending what was derived before  for the two-point propagator in~\\cite{2006PhRvD..73f3520C}. Our results show that multi-point propagators decay with the same rate as the two-point propagator in the large-$k$ limit and are proportional to their tree-level values. We have extended the algorithm presented in~\\cite{2006PhRvD..73f3520C} to measure the three-point propagator in numerical simulations, and showed that our high-$k$ limit resummation agrees well with both two and three-point propagators at nonlinear scales.  Furthermore, we showed that the multi-point propagators can be considered building blocks out of which one can construct arbitrary $n-$point spectra. In addition, nonlinear corrections to the power spectrum depend on the initial spectrum through convolutions with multi-point propagators, which in turn determine the large-scale (tree-level) polyspectra. This gives a direct physical connection between nonlinear corrections to the power spectrum at small scales and higher-order correlations at large scales.  We used our results on the resummation of multi-point propagators to reconstruct the power spectrum and bispectrum up to one-loop level, and showed that the predictions for the reduced bispectrum agree well with measurements in numerical simulations. This is encouraging news for the calculation of higher-order correlation functions in renormalized perturbation theories."
    },
    "0806/0806.0331_arXiv.txt": {
        "abstract": "Gravitational lensing affects observed cosmological correlation functions because observed images do not coincide with true source locations. We treat this universal effect in a general way here, deriving a single formula that can be used to determine to what extent this effect distorts any correlation function. We then apply the general formula to the correlation functions of galaxies, of the 21-cm radiation field, and of the CMB. ",
        "introduction": "Gravitational lensing has emerged as a powerful cosmological tool. The spectacular images provided by strong lensing~\\cite{1996ApJ...471..643K,Kneib:2001qy,Bradac:2006er} together with the robust parameter determination anticipated with weak lensing~\\cite{Kaiser:1991qi,Mellier:1998pk,Bartelmann:1999yn,Hu:1999ek,Huterer:2001yu,Abazajian:2002ck,Takada:2003ef,Refregier:2003xe} are just two of the exciting developments in the field. When astronomers look for lensing, they have demonstrated that they can find it. Lensing can also contaminate observations of correlation functions because the images of objects are displaced from the true source positions. For example, it has long been known that lensing smoothes out the spectrum of the cosmic microwave background (CMB)~\\cite{Seljak:1995ve,Lewis:2006fu,Smith:2007rg}. The effect of lensing on galaxy-galaxy correlations has been re-examined by several groups recently~\\cite{Moessner:1997vm,Matsubara2000,Hui:2007cu,Vallinotto:2007mf,LoVerde:2007ke,Hui:2007tm,Schmidt:2008mb}. And a number of groups are anticipating the importance of lensing on future observations of neutral hydrogen via long wavelength measurements of the redshifted 21 cm line~\\cite{Mandel:2005xh,Zhang:2006fc,Lewis:2007kz}.  In this work we present a unified treatment of the lensing contamination of correlation functions. First, we show that weak lensing contaminates \\textit{any} cosmological correlation function. Whereas previous studies have quantified its effect on some specific correlation functions (for instance, CMB temperature and polarization \\cite{Lewis:2006fu} and baryon oscillations through the galaxy correlation function \\cite{Vallinotto:2007mf}), we focus here on the \\textit{generality} of the result. This general framework allows us to understand when and for what correlation functions weak lensing represents a sizeable or a negligible effect. By working in real space, we also hope to provide physical intuition for some of the properties of the lensing effects, which is harder to obtain from a pure $k$-space calculation. Second, we consider the dependence of the lensing effect on the orientation of the source separation with respect to the line of sight. We also include the lensing-induced time delay as the longitudinal companion effect to the transverse lensing deflections. Finally, we show how the general framework presented here can be applied to different correlation functions: we retrieve previous results for the CMB, supplement and correct previous galaxy-galaxy results~\\cite{Vallinotto:2007mf}, and apply the formalism to the anisotropy induced by lensing of 21-cm radiation.  The paper is organized as follows. In Sec.~\\ref{Sec:GenRes3D} a general formalism to calculate the weak lensing contribution to cosmological correlations is introduced along with a more physical understanding of the different quantities involved. In Sec.~\\ref{Sec:Apps} the general formalism developed is applied to three different cases: correlation of galaxies, of the cosmic microwave background and of the 21-cm radiation. Finally, Sec.~\\ref{Sec:Conclusions} contains some discussion together with concluding remarks. The calculations are relegated to the appendices. ",
        "conclusions": "\\label{Sec:Conclusions}  Gravitational lensing affects cosmological correlation functions in two ways. First, since the geodesics which photons travel on are sensitive to the distribution of matter between source and observer and since the latter is locally inhomogeneous thanks to structure formation, weak lensing acts by displacing the sources' observed (as opposed to true) positions. The correlation function that is measured from observed data therefore also includes the contribution from these lensing deflections. Second, weak lensing also acts on observations by focusing or defocusing the geodesics' congruences, thus altering the observed brightness of a source. This latter effect, known as \\textit{magnification bias}, is generally larger than the former \\cite{Matsubara2000,Hui:2007cu} but it applies only to pointlike sources (see e.g. \\cite{Hui:2007cu, Vallinotto:2007mf, Schmidt:2008mb} for studies of the effect of magnification bias on two- and three-point correlation functions). Both these effects are small but potentially relevant in the present age of precision cosmology. In this work we focused on the first of these effects and we derived a \\textit{general} perturbative formula that can be used to quantify it.  The results of the present work depend only on the following two assumptions: that photons travel on geodesics and that higher order corrections to the correlation function (arising from higher order correlators of the lensing induced displacements) can safely be neglected. The first of this assumption has two consequences of opposite nature. The positive consequence is that the result obtained in this work is general and applies to \\textit{any} cosmological correlation function, regardless of the nature of the source and of the physical observables that are being measured and correlated. Moreover, the correction terms that appear in Eq.~(\\ref{GenRes:LH}) and that quantify the effect of lensing depend only on the power spectrum of matter: given a specific cosmological model they need to be evaluated only once and they can then be applied to any correlation function. As such, they also represent a map that will tell whether weak lensing will have a relevant or an irrelevant role in the measurement of a given correlation function \\textit{before} the actual measurement is carried out. The negative consequence, on the other hand, is that this lensing-induced distortions represent a theoretical systematic that will \\textit{always} be present regardless of the precision of the instruments used to carry out the observation. In other words, even if ``perfect observations'' free of any systematics could be carried out, this lensing-induced noise would still creep into the correlation function that would be calculated using those data. This lensing-induced modification of the correlation may be avoided only by reconstructing a map of the lensing potential (e.g., \\cite{Hirata:2002jy}). Another avenue to a possible disentanglement of this lensing distortion from the observed correlation function -- at least for sources that are not confined to a fixed redshift -- could be to exploit the angular dependence of the effect as seen in Fig.~\\ref{fig:deltaxi}. Furthermore, since the lensing modification depends on the derivatives of the correlation function, correlation functions that are rapidly oscillating will in general be most affected by lensing, i.e. the features will be somewhat smoothed. Conversely, in quantities that are intrinsically uncorrelated, no correlation will be induced by lensing deflections.  The second assumption entering the derivation of Eq.~(\\ref{GenRes:LH}) is that terms of third and higher order are discarded when the exponential of Eq.~(\\ref{GenRes:expcorrfunc}) is expanded. As shown in Appendix \\ref{App1}, this assumption corresponds to Taylor expanding the coordinate dependence of the physical observables and to retain only terms up to second order, which in turn corresponds to neglecting all contributions arising from the non-Gaussianity of the displacements' PDF. Despite the fact that under the same assumptions the general formalism outlined in Appendix \\ref{App:Exact} can be used to calculate the lensing distortion arising from the sum of \\textit{all} the terms appearing in the exponential of Eq.~(\\ref{GenRes:expcorrfunc}), it is however unclear to what extent this might really represent an improvement. If on one hand the sum of all the terms appearing in the exponential may give important contributions on small scales, it is also true that on such small scales departures from gaussianity of the displacements' PDF -- induced by the non-linear growth of structure at low redshift -- could play a relevant role. The range of applicability of such a non-perturbative method and the gain in precision that it would allow on small scales are the focus of a current project."
    },
    "0806/0806.2758_arXiv.txt": {
        "abstract": "We investigate the relations among the stellar continuum-subtracted 8~$\\mu$m polycyclic aromatic hydrocarbon (PAH 8~$\\mu$m) emission, 24~$\\mu$m hot dust emission, and 160~$\\mu$m cold dust emission in fifteen nearby face-on spiral galaxies in the {\\it Spitzer} Infrared Nearby Galaxies Survey sample.  The relation between PAH 8 and 24~$\\mu$m emission measured in $\\sim2$~kpc regions is found to exhibit a significant amount of scatter, and strong spatial variations are observed in the (PAH 8~$\\mu$m)/24~$\\mu$m surface brightness ratio.  In particular, the (PAH 8~$\\mu$m)/24~$\\mu$m surface brightness ratio is observed to be high in the diffuse interstellar medium and low in bright star-forming regions and other locations with high 24~$\\mu$m surface brightness.  PAH 8~$\\mu$m emission is found to be well-correlated with 160~$\\mu$m emission on spatial scales of $\\sim2$~kpc, and the (PAH 8~$\\mu$m)/160~$\\mu$m surface brightness ratio is generally observed to increase as the 160~$\\mu$m surface brightness increases.  These results suggest that the PAHs are associated with the diffuse, cold dust that produces most of the 160~$\\mu$m emission in these galaxies, and the variations in the (PAH 8~$\\mu$m)/160~$\\mu$m ratio may generally be indicative of either the intensity or the spectrum of the interstellar radiation field that is heating both the PAHs and the diffuse interstellar dust. ",
        "introduction": "}  Observations of nearby spiral galaxies with the Infrared Space Observatory (ISO) demonstrated that polycyclic aromatic hydrocarbon (PAH) spectral feature emission at $\\sim8$~$\\mu$m is closely associated with hot dust emission at 15~$\\mu$m, and the 8~$\\mu$m emission was also shown to be correlated with other star formation tracers \\citep{retal01, frsc04}.  However, observations with the {\\it Spitzer} Space Telescope \\citep{wetal04} have demonstrated that the correlation of PAH emission at 8~$\\mu$m to hot dust emission at 24~$\\mu$m was not necessarily a one-to-one correlation within individual galaxies.  \\citet{hraetal04}, \\citet{betal06}, and \\citet{getal08} demonstrated that, on scales of hundreds of parsecs, continuum-subtracted PAH emission at 8~$\\mu$m (henceforth referred to as PAH 8~$\\mu$m emission) does not accurately trace 24~$\\mu$m emission.  In particular, PAH 8~$\\mu$m emission seems to appear in shell-like features around star-forming regions, while 24~$\\mu$m emission peaks within star-forming regions.  Using images of M51, \\citet{cetal05} demonstrated that PAH 8~$\\mu$m emission compared to other tracers of star formation, including 24~$\\mu$m emission, is disproportionately low within star-forming regions and disproportionately high in diffuse regions.  The results from \\citet{cetal07} show that the PAH 8~$\\mu$m/24~$\\mu$m emission ratio measured in H{\\small II} regions is partly dependent on the aperture used for measuring the flux densities, which also suggest that a significant fraction of PAH 8~$\\mu$m emission originates from outside the H{\\small II} regions.  \\citet{tetal07} also demonstrated that the PAH 8~$\\mu$m/24~$\\mu$m ratio of H{\\small II} regions in NGC~7331 may depend on infrared, H$\\alpha$, or ultraviolet surface brightness after corrections for dust extinction have been applied, which implied that PAH 8~$\\mu$m emission is suppressed within strong star-forming regions.  \\citet{eetal05}, \\citet{detal05}, \\citet{ddbetal07}, and \\citet{eetal08} have demonstrated that the global 8~$\\mu$m/24~$\\mu$m emission ratio may decrease in regions with lower metallicity, and \\citet{cetal07} also demonstrated that this ratio for H{\\small II} regions in other galaxies depends on metallicity.  Although these studies based on {\\it Spitzer} data have shown explicitly or implicitly that the relation between 8~$\\mu$m PAH emission and 24~$\\mu$m emission on $\\ltsim 1$~kpc scales is not a one-to-one relation, no study has yet explored variation in the (PAH 8~$\\mu$m)/24~$\\mu$m ratio between diffuse and star-forming regions within a broad range of spiral galaxies.  While most research has focused on comparisons between PAH emission and hot dust emission, some additional studies have compared PAH emission to cold dust emission at wavelengths longer than 100~$\\mu$m. Using ISO data, \\citet{mll99} and \\citet{hkb02} demonstrated that PAH emission at 8~$\\mu$m is correlated with large grain emission, and \\citet{hkb02} even argued that the correlation between 8~$\\mu$m emission and cold dust emission at 850~$\\mu$m was much stronger than the correlation between 8 and 15~$\\mu$m emission.  Using {\\it Spitzer} data, \\citet{betal06} demonstrated that PAH 8~$\\mu$m emission was correlated with 160~$\\mu$m cold dust emission in NGC~4631, but the ratio of (PAH 8~$\\mu$m)/160~$\\mu$m emission varied.  Although the (PAH 8~$\\mu$m)/160~$\\mu$m ratio appeared related to the infrared surface brightness in NGC~4631, \\citet{betal06} used bright infrared sources in the outer disc to argue that the variations in the (PAH 8~$\\mu$m)/160~$\\mu$m ratio are instead dependent on radius.  Given that the ratio of PAH emission to dust emission at longer wavelengths or the ratio of PAHs to total dust mass may vary with metallicity \\citep{eetal05, detal05, setal07, ddbetal07, eetal08} and that metallicity is generally expected to vary with radius within spiral galaxies \\citep[e.g.][J. Moustakas et al. 2008, in preparation]{s71, ws83, ve92, zkh94, vetal98, petal04}, the results from NGC~4631 suggested that the variations in (PAH 8~$\\mu$m)/160~$\\mu$m ratio within that galaxy could be related to metallicity.  Unfortunately, NGC~4631 is viewed edge-on, so local variations in the observed (PAH 8~$\\mu$m)/160~$\\mu$m ratio that are not dependent on radius may be suppressed by the line-of-sight integrations through the disc.  Variations in the (PAH 8~$\\mu$m)/160~$\\mu$m ratio and its dependence on both infrared surface brightness and radius need to be studied further using spiral galaxies with orientations closer to face-on.  In this paper, we examine the relation of PAH 8~$\\mu$m emission to both 24~$\\mu$m hot dust emission and 160~$\\mu$m cold dust emission in a set of fifteen face-on spiral galaxies observed by {\\it Spitzer} as part of the {\\it Spitzer} Infrared Nearby Galaxies Survey (SINGS) legacy project \\citep{kabetal03}.  The basic goal is to understand how PAH 8~$\\mu$m emission is related to dust emission at other wavelengths.  Section~\\ref{s_data} provides basic information on the wave bands used in the analysis, the observations and data reductions, the sample, and the preparation of the data for the analysis.  The comparison between PAH 8 and 24~$\\mu$m emission is presented in Section~\\ref{s_comp_pah24}, and the comparison between PAH 8 and 160~$\\mu$m emission is presented in Section~\\ref{s_comp_pah160}.  The analysis is then followed by a discussion in Section~\\ref{s_discuss} and conclusions in Section~\\ref{s_conclusions}. ",
        "conclusions": "\\label{s_conclusions}  We have shown that the relation between PAH 8 and 24~$\\mu$m emission in this sample of spiral galaxies exhibit significant scatter on scales of $\\sim2$~kpc.  In particular, we have shown that the PAH 8~$\\mu$m emission is relatively weak compared to 24~$\\mu$m emission in star-forming regions, but the PAH 8~$\\mu$m emission is relatively strong in the diffuse ISM.  In some cases, the (PAH 8~$\\mu$m)/24~$\\mu$m ratio may exhibit variations greater than a factor of 2 at a given 24~$\\mu$m surface brightness.  We argue that two mechanisms are responsible for these variations: enhancements in the 24~$\\mu$m emission relative to PAH emission in regions with very strong radiation fields with high-energy photons and the corresponding reduction of PAH 8~$\\mu$m emission.  We have also shown that the PAH 8~$\\mu$m emission is associated with 160~$\\mu$m emission on scales of $\\sim2$~kpc within these spiral galaxies and that the (PAH 8~$\\mu$m)/160~$\\mu$m ratio appears to be a function of 160~$\\mu$m surface brightness.  The scatter in the relation is only at the 10\\%-20\\% level for a given surface brightness, and the intrinsic scatter measured in the relation between the (PAH 8~$\\mu$m)/160~$\\mu$m ratio and 160~$\\mu$m surface brightness is notably lower than for the relation between the (PAH 8~$\\mu$m)/24~$\\mu$m ratio and 24~$\\mu$m surface brightness. While the (PAH 8~$\\mu$m)/160~$\\mu$m ratio appears to decrease monotonically with radius in some galaxies, the presence of peaks in the (PAH 8~$\\mu$m)/160~$\\mu$m ratio corresponding to large scale structure indicates that surface brightness may be a more important factor, and statistical tests of the best fit lines between the (PAH 8~$\\mu$m)/160~$\\mu$m ratio and either 160~$\\mu$m surface brightness or radius showed that the dependence with 160~$\\mu$m surface brightness was more significant.  The strong correlation between the (PAH 8~$\\mu$m)/160~$\\mu$m ratio and 160~$\\mu$m emission and the results from the comparison of PAH 8 and 24~$\\mu$m emission suggest that most of the PAHs are located in the diffuse interstellar medium with the dust grains that produce the majority of the 160~$\\mu$m emission.  We therefore suggest that the (PAH 8~$\\mu$m)/160~$\\mu$m ratio may be indicative of the intensity of the interstellar radiation field that heats the diffuse interstellar dust in these galaxies.  The results here indicate that PAH emission should be used very cautiously as a tracer of star formation on kiloparsec scales if it should be used at all.  Instead, the PAH emission may be more indicative of the distribution of diffuse dust within nearby galaxies, although PAH emission still may be affected by metallicity \\citep{eetal05, detal05, cetal07, eetal08}, and the results here suggest that PAH emission might be inhibited in the diffuse ISM if the radiation field is very high.  Follow-up observations with the Herschel Space Observatory and the James Clerk Maxwell Telescope will allow for studying the correlation between PAH and $\\sim20$~K dust emission on smaller spatial scales in nearby galaxies, thus placing further constraints on the relation between PAHs and cool dust."
    },
    "0806/0806.0088_arXiv.txt": {
        "abstract": " ",
        "introduction": "\\label{intro} It is well established observationally and theoretically that many young stars in the solar neighborhood are surrounded by disks, i.e., flattened rotating concentrations of gas and dust (see \\cite{Greaves05} for a review). In the early phases, up to a few Myr after collapse of the parent cloud, the disks are rich in gas and dust inherited from the cloud core. Once accretion onto the star stops and the disk becomes less turbulent, the $\\sim$0.1 $\\mu$m interstellar grains coagulate to larger and larger particles and settle to the midplane, eventually forming km-sized planetesimals which interact gravitationally to form protoplanets (e.g., \\cite{Weiden97}). In this period, between a few and 10 Myr, also most of the gas may be dissipated from the disk although the precise time scale for gas removal is still uncertain. Dust disks are observed around much older (up to a few Gyr) stars as well, but the grains in these disks are the `debris' produced by collisions and fragmentation of planetesimals and are no longer the original interstellar particles (e.g., \\cite{Lagrange00}). This paper is concerned with young, gas-rich disks up to the transitional stage when the disk becomes optically thin to UV radiation.  Young stars are known to be powerful emitters of UV radiation, with typical intensities at the disk surface that are orders of magnitude higher than the average interstellar radiation field \\cite{Herbig86}. For some sources, most of the flux may be contained in the Lyman $\\alpha$ line \\cite{Bergin03}. These UV photons play a very important role in the physical and chemical structure of the disk, especially in the surface and intermediate layers.  Many disks are thought to have a flaring structure, in which the disk height $H$ and opening angle $H/R$ increase with radius $R$ (e.g., \\cite{Chiang97,Alessio98}). In such a geometry the disk surface intercepts significantly more UV radiation from the star than in a flat geometry, heating the surface layers and increasing the scale height further when hydrostatic equilibrium is assumed in the vertical direction.  Because the density in the upper layers is below that at which gas-grain coupling becomes effective, the gas temperature can be higher than the dust temperature due to photoelectric heating, affecting both the structure, chemistry and line formation \\cite{Jonkheid04,Kamp04a}.  In the inner disk ($<10$AU), the temperature can even become so high (several thousand K) that photoevaporation becomes effective \\cite{Gorti04}.  The main effect of the UV radiation on the chemistry is through photodissociation and photoionization of molecules and atoms in the surface layers.  These processes remain effective until most photons have been absorbed by the dust and the photorates have fallen below those of other chemical reactions, typically at a depth of a few magnitudes of extinction. The chemical structure of the outer disk has been modeled by various groups in both the radial and vertical direction (e.g., \\cite{Willacy00,Aikawa02,Markwick02,vanZ03}).  In the vertical direction, the chemical structure resembles that of a photon-dominated region (PDR), with a transition from C$^+$ to C and CO around $A_V\\approx$1 mag and radicals such as CN and C$_2$H having a significant abundance in the intermediate layers. In the cold mid-plane ($<20$~K), most molecules are frozen out onto the grains. This layered chemical structure is consistent with observations of ions and radicals with abundance ratios that are similar or higher than those found in PDRs (e.g., \\cite{Dutrey97,Thi04}). If the UV radiation is energetic enough to dissociate CO, isotope selective processes may occur leading to fractionation of CO isotopes which can eventually be incorporated into meteorites \\cite{Lyons05}.  Finally, UV photons can affect the chemistry through photodesorption of ices brought to the surface by vertical mixing in the disk \\cite{Willacy00}.  The above discussion illustrates the need for accurate photodissociation cross sections to describe the disk chemistry.  In this paper, photodissociation rates relevant for disks exposed to different types of radiation fields are presented, including calculations of cross sections for species not considered before. Also, a critical evaluation of available cross sections at the Lyman $\\alpha$ wavelength is given. The resulting chemical structure is illustrated for the transitional disk around the Herbig Ae/Be star HD 141569A. ",
        "conclusions": "The main results of this work are as follows.  \\begin{itemize}  \\item{New photodissociation cross sections are presented for several species not considered before. Also, the first critical evaluation of cross sections at Lyman $\\alpha$ is given.}  \\item{The photodissociation of ions is potentially important in (inner) disk chemistry where the radiation field can be $10^7$ times more intense than the standard ISRF. However, the photorates of many ions, especially those containing oxygen, are found to be slow compared with those of neutrals.}  \\item{Photorates have been computed for a range of radiation fields appropriate for Herbig Ae/Be and T Tauri stars. Molecules such as H$_2$, CO, N$_2$ and CN, which are photodissociated only below 1200 \\AA, have rates that are more than 5 orders of magnitude decreased for a $T_{BB}=4000$ K blackbody field.}  \\item{The exponent $\\gamma$ characterizing the depth-dependence of the photorates decreases with $T_{BB}$ but the effect is not large in the 4000--30000 K range for most species. For larger $\\mu$m-sized grains, the depth dependence of the photorates becomes much shallower, with $\\gamma$ falling in a narrow range of 0.4--0.6 for all species.}  \\item{The example of the HD 141569A disk illustrates that both the shape of the stellar radiation field and the size of the grains affect the chemistry.  Both need to be treated correctly to derive quantitative conclusions about the gas mass and chemistry from observed lines. Even a small amount of PAHs or small grains can significantly affect dissociation and ionization rates, as well as the H$_2$ formation rate.}  \\end{itemize}"
    },
    "0806/0806.3873_arXiv.txt": {
        "abstract": "{ Eccentric accretion disks in superoutbursting cataclysmic and other binary systems. } { We study the development of finite eccentricity in  accretion disks in close binary systems using a grid-based numerical scheme. We perform detailed parameter studies to explore the dependence on viscosity, disk aspect ratio, the inclusion of a mass-transfer stream and the role of the  boundary conditions. } { Using a two-dimensional grid-based scheme we study the instability of accretion disks in close binary systems that causes them to attain a quasi-steady state with finite eccentricity. Mass ratios $0.05 \\le q \\le 0.3$ appropriate to superoutbursting cataclysmic binary systems are considered.} { Our grid-based scheme enables us to study the development of eccentric disks for disk aspect ratio $h$ in the range $0.01-0.06$ and dimensionless kinematic viscosity $\\nu$ in the range $3.3 \\times 10^{-6} - 10^{-4}.$ Previous studies using particle-based methods were limited to the largest values for these parameters on account of their diffusive nature. Instability to the formation of a precessing eccentric disk that attains a quasi-steady state with mean eccentricity in the range $0.3 - 0.5$ occurs readily. The shortest growth times are $\\sim 15$ binary orbits for the largest viscosities and the instability mechanism is for the most part consistent with the mode-coupling mechanism associated with the 3:1 resonance proposed by Lubow. However, the results are sensitive to the treatment of the inner boundary and to the incorporation of the mass-transfer stream. In the presence of a stream we found a critical viscosity below which the disk remains circular. } {Eccentric disks readily develop in close binary systems with $ 0.05 \\le q \\le 0.3.$ Incorporation of a mass-transfer stream tends to impart stability for small enough viscosity (or, equivalently, mass-transfer rate through the disk) and to assist in obtaining a  prograde precession rate that is in agreement with observations. For the larger $q$ the location of the 3:1 resonance is pushed outwards towards the Roche lobe where higher-order mode couplings and nonlinearity occur. It is likely that three-dimensional simulations that properly resolve the disk's vertical structure are required  to make significant progress in this case. } ",
        "introduction": "\\label{sec:introduction}  SU Ursae Majoris systems are a class of cataclysmic variable that in addition to undergoing ordinary outbursts undergo less frequent larger superoutbursts. During these, their  light curves show superhumps which have a period generally slightly longer than the orbital period \\citep [for a discussion of relevant observations see for example][] {{2002A&A...389..478A},{2005PASP..117.1204P}}.  A natural explanation of this phenomenon is that it is associated with a rigidly precessing eccentric accretion disk. Superoutbursts and superhumps are found to be limited to cataclysmic binaries with mass ratio $q < 0.35.$ Assuming that the difference between the superhump period  and the orbital period represents the solid body precession of an eccentric disk, \\citet{{2005PASP..117.1204P}} discuss the  establishment of the expected correlation of this  period with $q.$  Studies of the light curve and effects on it associated with the  mass-transfer stream impacting the disk have enabled an eccentricity of $0.38\\pm 0.1$ to be  found for the disk in OY Car by \\citet{1992A&A...263..147H}, who also argue that the normal superhumps are produced by tidal stresses after the maximally extended disk edge approaches the companion. Similar results  for IY UMa, which has a similar mass ratio, were obtained by \\citet{2001MNRAS.324..529R}. These authors also noted the potential importance of the impacting accretion or mass-transfer stream on the eccentric disk with its varying outer radius  for understanding  the late superhump light.  These observations motivate an interest in developing good numerical schemes for simulating accretion disks undergoing tidal interaction with a secondary binary component.  Although, starting  with \\citet{1988MNRAS.232...35W}, there have been many simulations of eccentric disks using particle-based methods \\citep[e.g.][ and references therein]{1997MNRAS.289..889K, 1998MNRAS.297..323M, 2007MNRAS.378..785S} there has been almost no grid-based work to date that is applicable to close binary systems, even though such an approach, apart from constituting a check, is promising in the context of being able to gain access to a wider region of parameter space that may be difficult to access with particle-based methods. \\citet{1994A&A...288..807H} performed grid-based simulations, but without including disk viscosity and was only able to obtain an eccentric disk when  only the $m=3$ component of the companion's tidal potential was retained. Subsequent grid-based simulations by \\citet{1999MNRAS.304..687S} and \\citet{2000AcA....50..163K} also did not show clear evidence for the onset of an eccentric instability. \\citet{2005A&AS..432..757P} solved the linearized equations for the $m=1$ mode of a free disk without a binary companion  and considered the possible development of a parametric instability when the vertical structure was incorporated.  A study of eccentric disks also has application to other areas of astrophysics where disks and orbiting companions interact. Thus \\citet{2006A&A...447..369K} have  studied  eccentric disks in the context of the very small mass ratios appropriate to disk--planet interactions where their existence may have important consequences for theories of planet formation.  Motivated by these considerations, the aim of this paper is to  make an extensive study of the dynamics of two-dimensional disks that lie in the orbital plane of a binary system, which consists of a primary star and a lower-mass secondary star in a circular orbit. The disk is centred on the primary star in a binary system with mass ratio in the range $0.05 \\le q \\le 0.3$ thought to be relevant for superoutbursting cataclysmic binaries. We perform grid-based simulations of disks that are initialised with circular motion, in order to study the onset of an eccentric instability and the properties of the resulting quasi-stationary eccentric disk. Use of  a  grid-based rather than a particle-based method enables consideration of disks with a wider range of viscosities and more realistic aspect ratios in the range $0.01-0.05$ that are presently difficult to access for particle-based methods such as SPH due to their inherently diffusive nature \\citep[but see][]{2004A&A...418..325S}. We establish numerical convergence and explore the dependence of the results on the physical parameters of the disk. We also consider variations of the inner boundary conditions and the presence of a mass-transfer stream. Due to the global nature of the eccentric disk mode, results are found to be sensitive to these. When a stream is present we are able to find a transition to stability when the viscosity is reduced below a critical value. We also find that incorporation of a mass-transfer stream is important for obtaining values for the disk precession frequency that are compatible with observations.  The plan of the paper is as follows.  In Sect.~\\ref{sec:hydro-model} we describe the hydrodynamic model. We then go on to describe the numerical methods we use in Sect.~\\ref{subsect:numerics}. Numerical simulations showing the growth of eccentricity for a typical disk model are then presented in Sect.~\\ref{sec:standard}. It is found that a quasi-stationary eccentric disk with mean eccentricity in the range $0.3-0.5$ is often produced. We explore the dependence of this phenomenon on disk viscosity, aspect ratio and the inner boundary condition in Sect.~\\ref{sec:physical}, showing that the development of an eccentric disk is favoured at high viscosity but, because it is a global phenomenon, it is sensitive to the treatment of the disk inner boundary. The quasi-stationary eccentric disk is found to precess as a rigid body. In this section we also study the dependence of the disk precession rate on the disk aspect ratio and boundary conditions showing that the precession rate is prograde for cool enough disks. We then  investigate the dependence of the tendency of the disk to become  eccentric on the binary mass ratio in Sect.~\\ref{subsec:mass} and discuss numerical resolution demonstrating the numerical convergence of our results  in Sect.~\\ref{subsec:resolution}.  In Sect.~\\ref{Instability mechanism} we explore the instability mechanism and establish the importance of the $m=3$ component of the tidal potential for causing the instability through the mode-coupling mechanism proposed by \\citet{1991ApJ...381..259L}. Comparison with a linear analysis of eccentric modes is made in Sect.~\\ref{sec:linear}.  We go on to investigate the effect of a mass-transfer stream in Sect.~\\ref{sec:stream}. This is shown to make an important contribution to the disk precession rate and also to cause the disk to become stable to developing eccentricity for sufficiently small viscosity. Finally in Sect.~\\ref{Concl} we summarize and discuss our results. ",
        "conclusions": "In this paper we have performed grid-based simulations of eccentric disks in close binary systems. Adopting a  grid-based rather than a particle-based method enables consideration of a wider and more appropriate range of physical parameters. We have been able to consider disk aspect ratios in the range $0.01 \\le h \\le 0.05$ and dimensionless viscosities in the range $3.3 \\times 10^{-6} \\le \\nu  \\le 10^{-4},$ whereas particle-based methods are limited to the upper  limits of these ranges. We note the relationship between $\\nu$ and the  $\\alpha$ parameter of \\citet{1973A&A....24..337S} $\\alpha = \\nu/h^2.$ Thus $\\nu=10^{-4}$  corresponds to $\\alpha \\sim  0.1$ for $h=0.03.$ We have also considered the effects of different inner boundary conditions, the effects of a mass-transfer stream and considered and demonstrated numerical convergence. However, due to heavy computational demands the grid-based simulations carried out in this first survey  have been limited to two dimensions and simply adopted a locally isothermal equation of state. We summarize the main results of our simulations below.  \\subsection{Instability to the formation of an eccentric disk}  We find instability to the formation of an eccentric disk in the mass range $0.05 \\le q \\le 0.3$. Explorative calculations have shown instability even for larger mass ratios up to $q=1$ in these two-dimensional simulations, see Fig.~\\ref{fig:q-mult1}. The most unstable cases are those with a reflecting inner boundary, large viscosity, small aspect ratio and no mass-transfer stream. An initially non-eccentric disk  quickly enters a linear phase of exponential growth of an unstable mode.   For $q=0.1$ we found growth times ranging between $162$ and $15$ binary orbits  for $\\nu$  in the range  $3.3 \\times 10^{-6}-10^{-4}$ with $h$ in the range  $0.02-0.05.$ A quasi-stationary precessing eccentric disk is attained in all  cases in about four growth times. In the most unstable cases (large $\\nu$ and small $h$) this corresponds to about $60$ binary orbits  or about $6$ days for superoutbursting cataclysmic binaries. This is consistent with reported superhump development times of a few days.   The mean eccentricities of the quasi-stationary disks are typically in the range $0.3-0.6$ when there is no mass-transfer stream and an inner reflecting boundary. This is in the upper part of the range $0.38\\pm 0.1$  found for OY Car \\citet{1992A&A...263..147H} and for IY Ma by \\citet{2001MNRAS.324..529R}.   However, it is important to note that results are sensitive to the inclusion of a mass-transfer stream and the treatment of the inner boundary condition. For the largest viscosities, use of an inner outflow condition rather than a reflecting boundary condition reduces the final mean  eccentricity from $0.5$ to $0.3.$ Incorporation of a mass-transfer stream does not affect the growth of eccentricity for $\\nu = 10^{-4}$ but we are able to see a transition to stability to formation of an eccentric disk at about $\\nu = 3\\times 10^{-5}$ in this case independently of the inner boundary condition. This is also consistent with the idea that the effective viscosity and consequently the mass-transfer rate through the disk is larger during the outburst phases of cataclysmic binaries during which superhumps are reported \\citep{2005PASP..117.1204P}.   \\subsection{Quasi-steady state of the eccentric disk}  Once the simulated disks go through an instability that results in them becoming eccentric, they typically attain a characteristic quasi-steady state. This was illustrated in Fig. \\ref{fig:m02f-inert}. Viewed from an inertial frame the secondary has a close approach to the slowly precessing disk apocentre approximately once every orbit, pulling out a tidal tail. This later impacts back on the disk which is not strongly affected by the secondary until the next close approach. Thus  almost the entire  tidal dissipation may occur as a result of these events. In this context we note that tidal dissipation has been proposed as the source of the luminosity variations during the main superhump phase while it has been proposed that effects due to modulation by the mass-transfer stream become significant during the late superhump phase \\citep{{1992A&A...263..147H},{2001MNRAS.324..529R}, {2005PASP..117.1204P}}.   \\subsection{Instability mechanism}  We have investigated the mechanism responsible for the excitation of the disk eccentricity and found this to be basically consistent with the mode-coupling mechanism of \\citet{1991ApJ...381..259L}. Neglecting the slow precession of the disk, viewed from the inertial frame the  eccentricity of the disk may be considered to arise  from a stationary surface density perturbation with azimuthal mode number $m=1.$ In addition the disk responds directly  to the tidal potential of the companion producing a surface density response that corotates with it and consists of a sum of Fourier components  with  $m= 1,2,3,\\dots$ with amplitudes that scale with the mass ratio, $q.$ Nonlinear coupling between the response to the tidal potential and the $m=1$ mode causes the excitation of secondary density waves with  pattern speeds $m\\omega/(m-1),$ $\\ m= 2, 3, 4,\\dots$ which exceed that of the binary and are thus associated with a tidal  energy and angular momentum transfer  to it.  The back reaction on the $m=1$ mode causes the eccentricity to grow. To work effectively the secondary waves should be associated with an inner Lindblad resonance. The most favourable case corresponds to $m=3$ and the disk gas orbiting with $\\Omega = 3\\omega,$ at a 3:1 commensurability. In this case  the Fourier  component of the tidal potential with $m=3$ plays a key role.  The results of simulations with $q=0.1$ supported this view with growth being almost absent when the $m=3$ component was removed. However, for $q=0.2$ growth still occurred in this situation  albeit at a reduced rate. This is apparently because higher-order couplings generated through the other Fourier components result in an effective forcing with $m=3.$ Being of higher order, this is more effective at larger $q.$ But it should be noted that as $q$ increases, the $3:1$ commensurability is driven towards the Roche lobe thus becoming increasingly nonlinear  and affected by factors such as the outer boundary condition and any stream inflow. As we indicate below, it is possible that dissipative effects not included in our two-dimensional simulations play a more important role. For these reasons we concentrated on the case $q=0.1$ in this  article.    \\subsection{Disk precession rates and the dependence on disk aspect ratio} The main determinants of the disk precession rate were found to be the disk aspect ratio and the impact of a mass-transfer stream if present. Results for $0.01 \\le h \\le 0.05$  for $q=0.1$ are summarized in Fig. \\ref{fig:hr-om}.  The quasi-steady  disk precession rate $\\dot\\varpi\\subscr{d}$ was found to decrease from positive values at small $h$ (prograde precession) to attain negative values at larger $h.$ This is because, as expected, the pressure contribution tends to produce retrograde precession. When no mass-transfer stream was present the transition from prograde to retrograde occurred for $h \\sim 0.025.$ We also determined the precession rates by independently solving the linear mode problem using the average surface density profile and in view of the strong nonlinearity present obtained surprising agreement with $h$ at the transition being $\\sim 0.03.$  An  approximate parabolic fit to the data points for the case with no mass-transfer stream was found to be $\\dot\\varpi\\subscr{d} = 0.022 -39h^2.$ This indicates that $\\dot\\varpi\\subscr{d} \\sim 0.02$ as indicated by observations \\citep{2005PASP..117.1204P} only for $h \\lsim 0.01.$ In view of the fact that values for $h$ estimated from accretion disk modelling fall in the range $0.01 \\le h \\le 0.03$ \\citep{{2007MNRAS.378..785S}, {2006MNRAS.368.1123G}}, this is rather small. Note that \\citet{2006MNRAS.368.1123G}  also obtained similarly small estimates  from their one-dimensional linear formalism.  However, the simulations for which the disk was impacted by a mass-transfer stream show the transition for $h \\approx 0.04$ and indicate the possibility of obtaining the observed precession rates for values of $h$ estimated from accretion disk modelling. Thus our results demonstrate the importance of the mass-transfer stream in this context as well as in the context of obtaining a transition from instability to stability at a reasonable value for the  disk viscosity.  On the other hand, it might be argued that the significance of the mass-transfer stream would be reduced during a superoutburst, when the accretion rate through the disk exceeds the mass-transfer rate from the companion.  In addition we comment that there is sensitivity to the inner boundary condition such that changing from a reflecting condition to a dissipative or free outflow condition also acts to make the precession frequency move in the  prograde direction.  \\subsection{Comparison with particle-based simulations} The most recent relevant SPH simulations have been performed by \\citet{2007MNRAS.378..785S}. Both two- and three-dimensional simulations have been considered. Several important issues should be noted. Such simulations are inevitably highly diffusive ($\\alpha \\ge 0.1$) and thick ($h \\ge 0.05$) so that parameter studies at smaller $h$ and smaller $\\alpha$ of the type considered here cannot be carried out. Comparison has to be limited to cases with large $\\nu$ and $h.$ The SPH simulations are performed with $\\sim 10^5$ particles and also with a mass-transfer stream. For a disk with $h =0.05$ one can estimate that there is only about one smoothing length per scaleheight in a three-dimensional simulation so that the vertical thickness is poorly resolved. The three-dimensional SPH  simulations tend to  produce a weaker phenomenon than that reported here with almost no  activity for $q \\gsim  0.2$ and mean disk eccentricities at the small end of the range implied by observation. In addition very long growth times approaching thousands of orbits are generally reported.  On the other hand   \\citet{2005PASP..117.1204P} indicate that the superhump phenomenon occurs for $q \\le 0.35$ and that it is associated with high mass-transfer  events during superoutbursts and a short development time requiring growth times $\\sim 20$ binary orbits (see above). It may also occur in a quasi-steady manner  in systems such as nova-like variables which are also believed to have a high mass-transfer rate which in this case is steady. This is also consistent with our finding that stability to eccentricity generation occurs at low viscosities and accordingly low  mass-transfer rates through the disk when a mass-transfer stream is present.  The reported two-dimensional SPH simulations are more in line with those reported here. This suggests that the three-dimensional case is more dissipative  though  the adequacy of the resolution in the vertical direction should also be an issue. It is likely that inclusion of effects arising from the vertical structure of the disk could result  in  processes leading to additional dissipation  \\citep{2005A&AS..432..757P} but these would need to be simulated with adequate resolution. The two-dimensional grid-based simulations reported here are possibly too unstable to forming eccentric disks, particularly at larger $q$. However, in those cases, the 3:1 resonance driving the instability is pushed towards the Roche lobe  where higher-order mode  couplings and increasing nonlinearity occur  and numerical studies should include a treatment of the vertical direction if this region is to be correctly dealt with.  Therefore an important  direction for future work will be to  develop three-dimensional grid-based simulations which resolve the disk vertical structure and address the issue of the importance of possible  additional dissipation modes."
    },
    "0806/0806.4454_arXiv.txt": {
        "abstract": "{The vertical stratification of the magnetic field strength in sunspot penumbrae is investigated by means of spectropolarimetric observations at high spatial resolution from the Hinode spacecraft. Assuming that the magnetic field changes linearly with optical depth we find that, in those regions where the magnetic field is more inclined and the Evershed flow is strongest (penumbral intraspines), the magnetic field can either increase or decrease with depth. Allowing more degrees of freedom to the magnetic field stratification reveals that the magnetic field initially decreases from $\\log\\tau_5 = -3$ until $\\log\\tau_5 \\simeq -1.0$, but increases again below that. The presence of strong magnetic fields near the continuum is at odds with the existence of regions void of magnetic fields at, or right below, the $\\tau_5=1$ level in the penumbra. However, they are compatible with the presence of a horizontal flux-tube-like field embedded in a magnetic atmosphere.} ",
        "introduction": "It is now widely accepted that the horizontal structure of the sunspot penumbra is composed of two magnetic components (Solanki 2003; Bellot Rubio 2003). One of them possesses a somewhat inclined ($\\simeq 40-50^{\\circ}$ with respect to the vertical direction to the solar surface) and strong ($\\sim 2000$ G) magnetic field, whereas the other is characterized by a weaker and more horizontal one (Lites et al. 1993; R\\\"uedi et al. 1998; Bellot Rubio et al. 2004; Borrero et al. 2004, 2005). Traditionally, these two magnetic components have been identified with a horizontal flux tube, that carries the Evershed flow, and is embedded in a more vertical background magnetic field: {\\it uncombed} model (Solanki \\& Montavon 1993; Schlichenmaier et al. 1998; Borrero 2007). Recently, this view has been challenged by Spruit \\& Scharmer (2006) and Scharmer \\& Spruit (2006), who propose instead that the penumbra is formed by magnetic field-free plumes (connected to the underlying convection zone) that pierce the penumbral magnetic field from beneath. This is the so-called  {\\it gappy} penumbral model.  So long as these two different magnetic structures (weak/horizontal and strong/vertical) have remained spatially (horizontally) unresolved, distinguishing between the uncombed and gappy penumbral scenarios has not been possible. However, with the new spectropolarimeter on board of the Japanese spacecraft Hinode (Kosugi et al. 2007; Shimizu et al. 2007) it is now possible to obtain high spatial resolution ($\\simeq 0.32$\") observations of the sunspot penumbra. This could be sufficient to distinguish between the uncombed and gappy models, since they both postulate the existence of flux tubes or field-free gaps that are about 200-300 km in diameter (Mart{\\'\\i}nez Pillet 2000; Spruit \\& Scharmer 2006). This feature is particularly interesting, because the {\\it uncombed} and {\\it gappy} models predict very different vertical stratifications in the magnetic field strength across the weak/horizontal magnetic field component: which is identified with an embedded flux tube in the uncombed model, but with a field-free gap in the gappy model. In the latter, the magnetic field decreases monotonically with depth, whereas the former possesses a magnetic field that decreases with depth only initially, since once the boundary of the flux tube is reached, the magnetic field can either decrease or increase depending upon the strength of the magnetic field inside the tube.  In this paper we will focus on obtaining the vertical stratification of the magnetic field for penumbral filaments (where the magnetic field is more horizontal and weaker) using high spatial resolution spectropolarimetric observations from Hinode, in order to establish which penumbral model is more realistic. The observations are described in Section 2. Section 3 describes our data analysis and results from our inversion technique.  Section 4 compares our findings with the predictions made by the uncombed and gappy penumbral models. In Section 5 we make a thourough investigation of the effects of the scattered light. Finally, Section 6 summarizes our findings. ",
        "conclusions": "%  The {\\it uncombed} model postulates that penumbral intraspines are characterized by the presence of horizontal flux tubes embedded in a surrounding atmosphere that possesses an inclined magnetic field. According to this model, looking along these regions should reveal a magnetic field that smoothly decreases at first, but once the flux tube contribution starts, the field strength could either increase or decrease. Alternatively, the {\\it gappy} penumbral model postulates that instraspines correspond instead to regions where convective field-free gaps penetrate the penumbral field. In this case the magnetic field strength should also decrease with optical depth at first, but suffer a much larger drop once the line-of-sight crosses the field-free gap.  In order to differentiate between these two models, we have used polarimetric data at very high spatial resolution, recorded with the spectropolarimeter on-board of the Japanese spacecraft Hinode, to investigate the depth variation of the magnetic field strength in the penumbra. We have selected a large number ($\\sim$7500) of pixels that are representative of weak and horizontal magnetic fields (i.e.: penumbral intraspines) carrying strong Evershed flows. From the inversion of the Stokes profiles at these locations we find that the magnetic field strength can either increase or decrease with optical depth. A more detailed inversion of the average Stokes vector over the selected pixels, shows that the magnetic field initially decreases, between $\\log\\tau_5 \\in [-3,-0.7]$, but increases thereafter until $\\log\\tau_5 = 0$.  The {\\it gappy} penumbral model can explain a smoothly decreasing magnetic field strength only if the $\\tau_5=1$ level is formed above the field free gap, otherwise a much more sudden decrease would be observed as the line-of-sight penetrates the field-free plasma. A partial solution to this problem can be found if we assume that the gap is not fully evacuated of magnetic field. However, it offers no explanation for about 66 \\% of the selected pixels, where an increasing magnetic field strength with optical depth is observed. The absence of field-free gaps, as indicated by the inversion, does not in itself imply that there is no form of convection present in the penumbra, but rather suggests that the convective energy transport takes places in the presence of a magnetic field (see Zakharov et al. 2008; Rempel \\& Sch\\\"ussler 2008). An example is the roll convection proposed by Danielson (1961).  All inferred stratifications are compatible with the scenario proposed by the {\\it uncombed} model. A magnetic field that decreases smoothly towards the interior of the Photosphere can be explained by either a flux tube (of any field strength) whose upper boundary layer lies below $\\tau_5=1$ or, if the upper boundary is above $\\tau_5=1$, with a magnetic field inside the flux tube that is weaker than the magnetic field a few hundred kilometers above. In addition, a magnetic field that increases towards the interior of the Photosphere is compatible with a flux tube with an upper boundary layer above $\\tau_5=1$, but with a stronger magnetic field than the one above. This very same configuration can explain a magnetic field that first decreases and then increases with optical depth, as inferred from the inversion of averaged intraspinal profiles.  We have also studied the effects of the scattered light in our inversions. We have seen that any inaccuracies in its treatment are unlikely to be a source of error in the stratification of the magnetic field strength. It would be very desirable to make a robust confirmation of our findings, namely $dB/d\\tau > 0$ in the outer penumbra, for a larger number of sunspots at different heliocentric angles and including also the disk-ward side of the penumbra. A natural extension of this work would be to use the Fe I lines at 1.56 $\\mu$m (which are formed deeper in the Photosphere) to confirm the absence of field-free regions around $\\tau_5=1$. Unfortunately, no such observations exist at the spatial resolution needed to resolve penumbral intraspines ($\\simeq$ 0.4\"). Indeed, some studies at slightly lower resolution (0.6-0.7\") have been presented by Cabrera Solana et al. (2008), who used simultaneous observations of Fe I 630 nm and 1.56 $\\mu$m recorded with the TIP (Mart{\\'\\i}nez Pillet et al. 1999) and POLIS (Schmidt et al. 2003) instruments. They found that in the outer penumbra, the horizontal magnetic field component (carrying the Evershed flow) was no longer weaker than the more vertical one. This can be used as an independent confirmation of our work, where we routinely find $dB/d\\tau > 0$ at large radial distances from the center of the sunspot. In addition, flux tubes with stronger magnetic field than the one of the environment in which they are embedded are also necessary to explain certain aspects of the net circular polarization observed in the outer penumbra of sunspots (Tritschler et al. 2007; Ichimoto et al. 2008)."
    },
    "0806/0806.1876_arXiv.txt": {
        "abstract": "A primordial magnetic field in the early universe will cause Faraday rotation of the linear polarization of the cosmic microwave background generated via Compton scattering at the surface of last scattering. This rotation induces a non-zero parity-odd (B-mode) polarization component. The Wilkinson Microwave Anisotropy Probe (WMAP) 5-year data puts an upper limit on the magnitude of the B-polarization power spectrum; assuming that the B-polarization signal is totally due to the Faraday rotation effect, the upper limits on the comoving amplitude of a primordial stochastic magnetic field range from $6 \\times 10^{-8}$  to $2 \\times 10^{-6}$ G on a comoving length scale of 1 Mpc, depending on the power spectrum of the magnetic field. ",
        "introduction": "The cosmic microwave background temperature and polarization anisotropy measurements presented in the 5-year data of the Wilkinson Microwave Anisotropy Probe (WMAP) provide the most comprehensive characterization to date of the perturbations present in the early universe at the time of photon-baryon decoupling \\cite{WMAP_5_meas}. The microwave background polarization field is conventionally decomposed into a sum of parity-even (``E-type'') and parity-odd (``B-type'') parts; since B-type polarization is absent in a cosmological scenario involving only scalar density perturbations, it serves as a powerful test of other cosmological physics, including primordial tensor and vector metric perturbations \\cite{EBdecomp}. Another source of B-type polarization is Faraday rotation of the orientation of linear polarization due to the presence of a cosmological magnetic field: if a purely E-type polarization pattern is everywhere rotated by $\\pi/4$, it becomes a purely B-type polarization. Smaller rotation angles generate a component of B-type polarization from an initial pure E-type polarization field.  Faraday rotation of polarized radiation is in general a powerful tool to estimate magnetic field amplitudes and spectra (see \\cite{Vallee} for a recent review). In the cosmological context, Kosowsky and Loeb \\cite{kosowsky96} proposed microwave background Faraday rotation to probe a homogeneous primordial magnetic field; this issue has been re-addressed by several authors for the homogeneous field \\cite{giovannini,Demianski:2007fz}, as well as for a stochastic field \\cite{dolgov,kklr05,giovannini05,GK08}. For a homogeneous field, the same polarization rotation effect additionally induces non-zero parity-odd cross correlations between microwave temperature and B-polarization anisotropies and E-polarization and B-polarizations anisotropies \\cite{ferreira} and correlations between different multipoles \\cite{DKY98} which are absent for standard cosmological scenarios. Primordial magnetic fields are interesting because they could serve as seed fields for the observed magnetic fields in galaxies and galaxy clusters \\cite{Widrow}.  Similar effects (microwave background birefringence and parity-odd anisotropy spectra) may also have non-magnetic explanations. In particular, microwave background polarization plane rotation occurs in models with violation of parity symmetry and/or Lorentz invariance in the early universe \\cite{reh}. Parity-odd polarization anisotropy spectra appear due to gravitational lensing \\cite{seljak} or non-zero primordial magnetic helicity \\cite{pogosian}.  Here we use WMAP 5-year B-polarization power spectrum upper limits \\cite{WMAP_5_spectra} to place limits from Faraday rotation on a stochastic primordial magnetic field. We assume that the cosmological magnetic field was generated during or prior to the early radiation-dominated epoch. The high conductivity of the primordial plasma  results in a ``frozen-in'' magnetic field, fixing the temporal dependence to be the simple scaling ${\\mathbf B}({\\mathbf x}, \\eta)={\\mathbf B}({\\mathbf x})/a^2$ where $a$ is the scale factor and $\\eta$ conformal time. Throughout this paper, ${\\bf B}$ represents the comoving value of the magnetic field. We also normalize the scale factor $a$ by setting $a_0 = 1$ today, and we employ cosmological units with $\\hbar = 1 = c$ and gaussian units for electromagnetic quantities. ",
        "conclusions": "To compare the data with the theory prediction, we calculate $C_{l}^{BB}$ for $B_{\\lambda} = 1$ nG, $\\lambda = 1$ Mpc, $n_B$ between -2.9 and -1 in steps of 0.005, and $l$ values from 2 to 1000. Our range of spectral indices is based on plausible magnetic field generation mechanisms \\cite{ratra} and spectral indices which are not excluded by other bounds; see Refs.~\\cite{n-limits,grasso}. The input $C_{l}^{EE}$ values up to $l=4000$ were calculated using CMBFAST \\cite{cmb} for the best-fit WMAP cosmological model. The obtained set of $C_{l}^{BB}$ is then further rescaled as $B_\\lambda^2$ and $\\lambda^{n_{B}+3}$ to obtain theoretical $C_{l}^{BB}$ for different values of magnetic field amplitude and scale, respectively. WMAP data \\cite{WMAP2} covers five frequency bands with central frequencies ranging from 23 GHz to 94 GHz. We use the measured B-polarization power spectrum limits for the three highest-frequency WMAP bands centered at 41, 61, and 94 GHz (Q, V, and W); our limits could be improved by including information from the lowest two frequency channels, but these are dominated by synchrotron foregrounds. We use only data for $l>32$, for which we can treat measurement errors for different $l$ values as uncorrelated. Most of the constraining power of the data comes from multipoles between $l=200$ and $l=500$, as displayed in Fig.~1, where the theoretical signal begins to be comparable to the errors on the data points; neglecting the data with $l<32$ only marginally affects the limits obtained here.   \\begin{figure} \\includegraphics[width=8.6cm]{limit_color_rev.pdf} \\caption{The 68\\% and 95\\% C.L. limit bands on the $B_\\lambda$ as function of $n_B$ for $\\lambda$ = 1 Mpc.} \\label{fig:bLimit_cl} \\end{figure}  The theoretical  $C_l^{BB}$ power spectra are compared with the WMAP data using $\\chi^2$ statistics. We conservatively include an extra 40\\% error on theoretical $C_l^{BB}$ values to compensate for using the nominal frequency center of each WMAP band rather than the detailed frequency response of each band. Theoretical uncertainties on calculated values of $C_{l}^{EE}$ are neglected, since these are much smaller than the experimental uncertainties on WMAP data. As the data is consistent with the hypothesis of $C_l^{BB}=0$ (see Fig. \\ref{fig:comparison_nb29_5nG}), we set limits on values of $B_\\lambda$ as a function of $n_B$ and $\\lambda$. The 68\\% and 95\\% confidence limit bands on $B_\\lambda$ as a function of $n_B$ are given in Fig.~\\ref{fig:bLimit_cl}. We also display in Fig.~\\ref{fig:bLimit_vs_lambda} the 95\\% confidence upper limits on $B_\\lambda$ for fixed values of $n_B$ ranging from -2.9 to -1.0 as functions of $\\lambda$.  These results presume that B-polarization signal is totally due to the Faraday rotation effect. Other possible source of B-polarization, such as inflationary gravitational waves \\cite{EBdecomp}, gravitational lensing \\cite{seljak} or non-zero primordial magnetic helicity \\cite{pogosian}, will strengthen the magnetic field limits derived here.  \\begin{figure} \\includegraphics[width=8.6cm]{b_vs_lambda_rev.pdf} \\caption{95\\% C.L.\\ upper limits on $B_\\lambda$ for fixed values of $n_B=-2.9$, -2.5, -2.0, -1.5, and -1.0 as functions of $\\lambda$.} \\label{fig:bLimit_vs_lambda} \\end{figure}  A primordial magnetic field on a scale $\\lambda=100$ Mpc must have an amplitude $B_\\lambda$ of less than $10^{-7}$ G for any power spectrum between $n_B=-2.9$ and $n_B=-1$. These limits are weaker that those obtained recently  in Refs.\\ \\cite{GK08}. Improvements on this limit using the Faraday rotation signal will be challenging, requiring substantially more sensitive polarization measurements at frequencies of 50 GHz or below; most current or planned high-sensitivity polarization experiments use bolometer detection technology, which tend to lose sensitivity in this frequency range. Upcoming measurements from the LFI instrument on Planck will be of roughly comparable polarization sensitivity to the WMAP measurements at similar frequencies. One notable effort is the QUIET experiment \\cite{samtleben07}, which uses coherent detector technology at frequencies of 40 and 90 GHz. It is eventually anticipated to measure the BB polarization power spectrum with an improvement in sensitivity over WMAP by a factor of $10^4$.  The lowest frequency of 40 GHz, compared to 23 GHz for WMAP, makes the Faraday polarization power spectrum smaller by a factor of 7.3 for the same magnetic field; such an experiment could thus place limits on the amplitude of a primordial field which are a factor of approximately $(10^4/7.3)^{1/2}=37$ more stringent than those here. A detailed analysis for this experiment \\cite{ferreira08} projects limits on a homogeneous primordial field below $10^{-10}$ G; however, primordial field limits will always be stronger than the stochastic field considered here, because Faraday rotation of a primordial field also generates nonzero TB and EB polarization cross-correlations which are larger than the corresponding BB power spectrum. Polarized foreground emission is also a potentially difficult systematic limit to these measurements. Other upcoming polarization experiments aimed at detecting the BB polarization from inflation are likely to improve on the Faraday rotation limits here, though the Faraday rotation power spectrum amplitude decreases by a factor of  230 between WMAP's 23 GHz channel and a likely 90 GHz lowest-frequency bolometer channel.  The cosmological magnetic field itself generates microwave background anisotropies, particularly via the tensor and vector perturbations it induces \\cite{mkk} (see Refs.~\\cite{grasso} for overviews of magnetized cosmological perturbations). Here we have neglected these perturbations, limiting ourself to considering the Faraday rotation effect alone, which provides a distinctive signature of magnetic fields nearly independent of other properties of the universe. We use the direct observational data without any priors.  Our magnetic field limits are weaker than those arising from  CMB temperature maps \\cite{lewis04,GK08} or current bounds on a homogeneous magnetic field due to correlations between different $l$ modes \\cite{chen} and CMB temperature non-gaussianity \\cite{ss09}. The strongest future limits on magnetic fields in the early universe will likely come from limits on the temperature and polarization fluctuations generated by magnetic field-induced vector perturbations \\cite{mkk}; we will address these current limits from WMAP data elsewhere."
    },
    "0806/0806.4381_arXiv.txt": {
        "abstract": "We correct the observed Milky Way satellite luminosity function for luminosity bias using published completeness limits for the Sloan Digital Sky Survey DR5. Assuming that the  spatial distribution of Milky Way satellites tracks the subhalos found in the \\emph{Via Lactea} $\\Lambda$CDM N-body simulation, we show that there should be between $\\sim 300$ and $\\sim 600$ satellites within 400 kpc of the Sun that are brighter than the faintest known dwarf galaxies, and that there may be as many as $\\sim 1000$, depending on assumptions.   By taking into account completeness limits, we show that the radial distribution of known Milky Way dwarfs is consistent with our assumption that the full satellite population tracks that of subhalos.     These results alleviate the primary worries associated with the so-called missing satellites problem in CDM.  We show that future, deep wide-field surveys such as SkyMapper, the Dark Energy Survey (DES), PanSTARRS, and the Large Synoptic Survey Telescope (LSST) will deliver a complete census of ultra-faint dwarf satellites out to the Milky Way virial radius, offer new limits on the free-streaming scale of dark matter, and provide unprecedented constraints on the low-luminosity threshold of galaxy formation. ",
        "introduction": "\\label{sec:intro} It is well-established from simulations in the $\\Lambda$CDM concordance cosmology that galaxy halos are formed by the merging of smaller halos \\citep[e.g.][and references therein]{stewart08}.  As a result of their high formation redshifts, and correspondingly high densities, a large number of self-bound dark matter subhalos should survive the merging process and exist within the dark matter halos of $L_*$ galaxies like the Milky Way \\citep{kl99ms,moo99ms,ZB03}.   A direct confirmation of this prediction has yet to occur.  Surveys of the halos of the Milky Way \\citep[e.g.][]{will05wI,belfos06,walsh07boo,bel08LeoV} and Andromeda \\citep[e.g.][]{martin06,majewski07,irwin08,mcconn08newands} have revealed only $\\sim 20$ luminous dwarf satellites around each galaxy, approximately 1 order of magnitude fewer than the expected number of subhalos that are thought to be massive enough to form stars \\citep{kl99ms, moo99ms,die07ms,str07redef}. Of course, the mismatch is much more extreme (larger than a factor of $\\sim 10^{10}$) when compared to the full mass function of CDM subhalos,  which is expected to rise as $\\sim 1/M$ down to the small-scale clustering cutoff for the dark matter particle,  $M_{cut} \\ll 1 M_\\odot$ \\citep {Hofmann:2001bi,bertschinger06,profumo06,jk08}  Astrophysical solutions to this missing satellites problem (MSP) include a reduction in the ability of small halos to accrete gas after reionization \\citep{BKW,somer02reion,ben02reion}, and tidal stripping, which shrinks the mass of halos \\emph{after} they have formed stars \\citep{KR04}.      Although there are some promising new techniques in development for studying the formation of low-mass galaxies within a CDM context \\citep{Robertson_Kravtsov,kauf07,orban08dsfh}, current models have trouble explaining many details, particularly regarding the lowest luminosity dwarfs (see below).  Other more exotic solutions rely on dark matter particles that are not cold \\citep{HD00,K05,Cembranos:2005us,SKB} or non-standard models of inflation \\citep{ZB03}, which produce small-scale power-spectrum cutoffs at $M_{cut} \\sim 10^6-10^8 \\Msun$, well above the cutoff scale in standard $\\Lambda$CDM.    In principle, each of these scenarios leaves its mark on the properties of satellite galaxies, although multiple scenarios fit the current data. Hence, precisely determining the shape and normalization of the lowest end of the luminosity function is critical  to understanding how faint galaxies form and how the efficiency of star formation is suppressed in the smallest dark matter halos.  The Sloan Digital Sky Survey (SDSS; \\citealt{SDSSDR5}) has revolutionized our view of the Milky Way and its environment, doubling the number of known dwarf spheroidal (dSph) galaxies over the last several years  \\citep[e.g.][]{will05wI,z06,z06b,belfos06,grillmair06,walsh07boo,Irwin07}.  Many of these new dSph galaxies are ultra-faint, with luminosities as low as $\\sim 1000 L_\\odot$, faint enough to evade detection in surveys with limits sufficient for detecting most previously known Local Group dwarfs \\citep{whiting07photolims}.  In addition to providing fainter detections, the homogeneous form of the SDSS allows for a much better understanding of the statistics of detection.  Unfortunately, given the inherent faintness of the newly discovered dSph's and the magnitude-limited nature of SDSS for such objects, a derivation of the full luminosity function of satellites within the Milky Way halo must include a substantial correction for more distant undetectable satellites.  \\citet{kop07} provided an important step in this process by performing simulations in order to quantify the detection limits of the SDSS and estimated the luminosity function by applying these limits to some simple radial distribution functions, finding $\\sim 70$ satellites and a satellite luminosity function consistent with a single power law of the form $dN/dL_V \\sim L^{-1.25}$.  Our aim is to take the detection limits for SDSS Data Release 5 (DR5) as determined by \\citet{kop07} and combine them with a CDM-motivated satellite distribution. In order to provide a theoretically-motivated estimate for the total number of Milky Way satellite galaxies, we adopt the distributions of subhalos in the \\emph{Via Lactea} simulation \\citep{die06vl}, and assume these to be hosts of satellite galaxies. Note that for the purposes of this paper, we define a  ``galaxy'' as a stellar system that is embedded in a dark matter halo.   The organization of this paper is as follows: In \\S \\ref{sec:app} we describe the overall strategy and sources of data used for this correction, as well as discussing the validity of the assumptions underlying these data. \\S \\ref{sec:corr} describes in detail how the correction is performed and presents the resulting luminosity functions for a number of possible scenarios, while \\S \\ref{sec:future} discusses the prospects for detecting the as-yet unseen satellites that the correction predicts.  \\S \\ref{sec:disc} discusses some of the caveats associated with the technique used for this correction, as well as the cosmological implications of the presence of this many satellite galaxies.  In \\S \\ref{sec:conc} we draw some final conclusions. ",
        "conclusions": "\\label{sec:conc} The goal of this paper has been to provide reasonable, cosmologically-motivated corrections to the observed luminosity function.   Our primary aim is to motivate future searches for faint dwarf galaxies and to explore the status of the missing satellites problem in light of the most recent discoveries.   By combing completeness limits for the SDSS \\citep{kop07} with the spatial distribution of subhalos in \\emph{Via Lactea}, we have shown that there are likely between $\\sim 300$ and $\\sim 600$ satellites brighter than Boo II within the Milky Way virial radius (Figure \\ref{fig:moneyplot}),  and that the total count may be as large as $\\sim 2000$, depending on assumptions (Table \\ref{tab:res}, Figure \\ref{fig:res}). We also showed that the observed satellites are indeed consistent with tracing the radial distribution of subhalos, provided completeness limits are taken into account (Figure \\ref{fig:raddist} and Table \\ref{tab:radKS}).  Moreover, we argued that future large sky surveys like LSST, DES, PanSTARRS, and SkyMapper should be able to see these satellites if they do exist, and thereby provide unprecedented constraints on the nature of galaxy formation in tiny halos.  While this correction predicts that nearly all of the undetected satellite are faint ($M_V>-7$) and consequently have low surface brightness, it is important to note that it  is not clear how this maps onto satellites' $v_{\\rm max}$.  While a possible test of this correction's result lies in selecting sub-samples of the observed satellite population, $v_{\\rm max}$ is difficult to constrain in the known satellites \\citep{str07redef}, and hence this exercise is suspect until simulations can resolve subhalos well enough to compare to observables such as the integrated mass within 300 pc (see \\S \\ref{sec:impl}).  There are two major points to take away from this correction: \\begin{itemize}   \\item  As it was first formulated, the MSP referred to the mismatch between the then $\\sim 10$ known dwarf satellite galaxies of the Milky Way and Andromeda, and the expected count of $\\sim 100-500$ subhalos with $v_{\\rm max} \\ge 10$ \\kms \\citep{kl99ms,moo99ms}.    Our results suggest that the recent discoveries of ultra-faint dwarfs about the Milky Way are consistent with a total population of $\\sim 500$ satellites, once we take into account the completeness limits of the SDSS.  In this sense, the primary worries associated with the MSP in CDM are alleviated.  Nonetheless, it is critical that searches for these faint galaxies be undertaken, as the assumptions of this correction must be tested.  \\item  The shape of the faint end of the satellite luminosity function is not yet constrained well enough to deeply understand the theoretical implications.  There exists yet a large parameter space in galaxy formation theory that will fall inside the error bars of Figure \\ref{fig:moneyplot}, and an even larger parameter space of models that are viable if our caveats and scenario possibilities are considered.  Our results strongly suggest that the luminosity function continues to rise to the faint end, with a faint end slope in our fiducial scenario given by $dN/dM_V = 10^{(0.35\\pm0.08) M_V + 3.42 \\pm 0.35}$, or $dN/dL_V \\propto L^\\alpha$, with $\\alpha=-1.9 \\pm 0.2$.  This is substantially steeper than the result of \\citet{kop07}, possibly the result of using a $\\Lambda$CDM-motivated subhalo distribution instead of the analytic profiles used in that paper.  But caution is advised in reading anything into the details of the shape, as nearly anything could be hiding within the faintest few bins when all the scenarios are considered.  This is apparent from Figure \\ref{fig:LSSTlfs}, where the Skymapper line appears as a power law over the range shown, while the actual luminosity function is certainly not.  Furthermore, depending on which scenario is used, the slope ($\\alpha$) can vary anywhere from -1.16 to -2.15 (see last column of Table \\ref{tab:res}).  \\end{itemize}  Fortunately, future deep large sky surveys will detect very faint satellites out to much larger distances and hence firmly observe the complete luminosity function out beyond the Milky Way virial radius (see Section \\ref{sec:future}).  With these data, it will be possible to provide stringent limits both on cosmology and galaxy formation scenarios (see Section \\ref{sec:impl}).  Nonetheless, the current data are not deep enough, and until the new survey data are available, there will be no way to put the spectre of the MSP completely to rest."
    },
    "0806/0806.4038_arXiv.txt": {
        "abstract": "{} {In order to deduce properties of dust in astrophysical environments where dust growth through aggregation is important, knowledge of the way aggregated particles interact with radiation, and what information is encoded in the thermal radiation they emit, is needed. The emission characteristics are determined by the size and structure of the aggregate and the composition and shape of the constituents. We thus aim at performing computations of compositionally inhomogeneous aggregates composed of irregularly shaped constituents. In addition we aim at developing an empirical recipe to compute the optical properties of such aggregates in a fast and accurate manner.} {We performed {CDA} computations for aggregates of irregularly shaped particles with various compositions. The constituents of the aggregate are assumed to be in the Rayleigh regime (i.e. they are much smaller than the wavelength of radiation inside and outside the constituents), and in addition we assume that the dominant interaction of the aggregate constituents is through dipole-dipole interactions. We computed the spectral structure of the emission efficiency in the $10\\,\\mu$m region for aggregates with 30\\% amorphous carbon and 70\\% silicates {by volume} with various fractions of crystalline and amorphous components.} {We find that the spectral appearance of the various components of the aggregate are very different and depend on their abundances. Most notably, materials that have a very low abundance appear spectroscopically as if they were in very small grains, while more abundant materials appear, spectroscopically to reside in larger grains. We construct a fast empirical approximate method, {based on the idea of an effective medium approximation,} to construct the spectra for these aggregates which almost perfectly reproduces the more exact computations. This new method is fast enough to be easily implemented in fitting procedures trying to deduce the dust characteristics from astronomical observations.} {} ",
        "introduction": "Dust grains in many astronomical environments are assumed to be aggregates of small particles of different sizes and compositions. This is especially true for cometary dust and dust in disks around young stars. The dust grains in these environments are formed by coagulation of smaller particles, leading to aggregated structures. In order to interpret observations of light scattered, emitted or absorbed by these composite particles, a method is required that takes the detailed properties of the particles into account.  Spectral resonances in the refractive index of the particle material translate into resonances in the absorption (and emission) spectra. Besides composition, the spectral location of such resonances is also influenced by the shape and size of the particles. Spectra computed for homogeneous spherical particles may often display resonances with strengths and spectral locations not found in the spectra of irregularly shaped particles \\citep[see e.g.][]{2001A&A...378..228F, 2003A&A...404...35M} or astronomical observations \\citep[see e.g.][]{2002A&A...390..533H, 2003A&A...401..577B, 2003ApJ...595..522M}. Much progress has been made on computational methods to calculate the optical properties of aggregates of spheres \\citep[see e.g.][]{1996JOSAA..13.2266M}. However, aggregates of homogeneous spherical particles {are still likely to display resonances typical for homogeneous spheres, and thus behave poorly when comparing to observations of natural particles. Test computations we performed confirm this.} Therefore, a method is required for computing the optical properties of aggregates of irregularly shaped constituents. {So far, an efficient and widely employable computational technique is still lacking.}  The aggregates expected to be present in astronomical environments like cometary comae or protoplanetary disks are generally assumed to be composites of various materials. {Although the inhomogeneous nature of the particles can be taken into account in several ways} \\citep[see e.g.][]{2003ApJ...595..522M, 2006A&A...456..535S, 2007A&A...463.1189K}, such inhomogeneous structures are usually modeled by assuming separated homogeneous particles of various compositions \\citep[see e.g.][]{2006ApJ...646.1024H, 2005Icar..179..158M, 2005A&A...437..189V, 2003A&A...401..577B, 1999P&SS...47..773B}. Although widely used, the validity of this approach is not obvious. In addition it is difficult to interpret the sizes and shapes derived for the dust grains in this way. Therefore, in this paper we study the effects of aggregation of various materials in single aggregates on the spectral appearance of these particles. In particular we will look at how the pronounced spectral structure of crystalline silicates is displayed when they are embedded in a larger aggregate.  In section~\\ref{sec:computational approach} we outline the computational approach we employ to compute the absorption spectra of inhomogeneous aggregates of irregularly shaped particles. Results of calculations are presented in section~\\ref{sec:results}. An efficient empirical method to compute the spectra for these aggregates using an effective medium approach is presented in section~\\ref{sec:effective medium}. Finally, in section~\\ref{sec:conclusions} we present the conclusions of our study. ",
        "conclusions": "\\label{sec:conclusions}  We have studied the spectral behavior of inhomogeneous aggregates of irregularly shaped constituents. We find that the spectral appearance of such aggregates is a complex interplay between the abundances of the materials and the size of the aggregate. Also the shape and structure of the aggregate plays an important part in the spectral appearance of such particles.  The most striking conclusion of our computations is that the {apparent particle size of a dust component, i.e. the size derived from the infrared spectrum}, is a strong function of its abundance. Materials with a relatively low abundance will appear spectroscopically smaller than materials with a relatively high abundance.  We present a very efficient, fast and accurate approximate method to compute the opacities of the inhomogeneous aggregates using an empirical effective medium approach. We refer to this approach as the Aggregate Polarizability Mixing Rule (APMR). This approach can be used to study the effects of inhomogeneity of particles of larger sizes than can be computed using more exact methods given the large computational demand that these methods put on the computations for increasing particle size.  Often, the analysis of infrared spectra of dusty environments is performed by assuming homogeneous, solid grains with various compositions and a limited number of grain sizes. The APMR approximate method is fast enough to be easily implemented in fitting algorithms trying to deduce the compositions and sizes of dust particles in these astrophysical environments. Furthermore, we propose to use the APMR to gauge the results obtained by previous studies and further investigate the effects of grain aggregation on the analysis of infrared spectra. Results of such an analysis will be presented in a separate paper."
    },
    "0806/0806.3104_arXiv.txt": {
        "abstract": "By adding a weak magnetic field to a spherically symmetric fluid configuration that caricatures a stalled shock in the post-bounce supernova environment, we explore the capacity of the stationary accretion shock instability (SASI) to generate magnetic fields. The SASI develops upon perturbation of  the initial condition, and the ensuing flow generates---{\\em in the absence of rotation}---dynamically significant magnetic fields ($\\sim 10^{15}$ G) on a time scale that is relevant for the explosion mechanism of core-collapse supernovae. We describe our model, present some recent results, and discuss their potential relevance for supernova models. ",
        "introduction": "Magnetic fields in the context of core-collapse supernovae---the final evolutionary stage of massive stars---have recently attracted considerable interest (see \\cite{shibata_etal_2006,burrows_etal_2007}, and references therein) due, in part, to the generically observed asphericity of supernova explosions \\cite{wang_etal_2001}.  In particular, simulations of rapidly rotating progenitors followed for several hundred milliseconds past core bounce demonstrate the ability of magnetic fields to drive and collimate bipolar explosions \\cite{burrows_etal_2007}: Differential rotation near the surface of the rapidly rotating proto-neutron star (PNS) winds and organizes poloidal magnetic fields into `magnetic towers' along the rotation axis, consisting mainly of toroidal fields. Although this mechanism is compelling, current models rely on rotation rates---and an initial field strength and topology---that are incompatible with stellar evolution calculations that take magnetic effects into account \\cite{heger_etal_2005}.  It is often assumed that the magneto-rotational instability (MRI), which may well operate in core-collapse supernovae \\cite{akiyama_etal_2003}, will bridge the gap and make the winding mechanism work. Capturing the MRI in numerical simulations requires higher spatial resolution than has been possible in realistic supernova simulations. The MRI has been observed in a simplified model of stellar collapse \\cite{shibata_etal_2006}, but the potential impact of the MRI (and magnetic fields generally) on the core-collapse supernova explosion mechanism remains unclear.  A different mechanism that may account for aspherical supernovae is the SASI, a hydrodynamical instability of the stalled supernova shock wave.  Idealized two-dimensional (2D) and three-dimensional (3D) simulations have shown the SASI to have a dramatic impact on the postshock flow \\cite{blondin_etal_2003,blondin_mezzacappa_2007}:  The $\\ell=1$ (sloshing) mode is seen to dominate 2D simulations, while the $m=1$ (spiral) mode dominates in 3D and may generate significant pulsar spin.  The SASI is also seen to play an important role in the dynamics of more realistic models \\cite{buras_etal_2006,burrows_etal_2006,bruenn_etal_2006,scheck_etal_2008}.  Given its impact on the postshock flow, it is natural to explore the influence of the SASI on magnetic fields (and vice-versa) in the supernova environment.  By extending an axisymmetric model \\cite{blondin_etal_2003} to include magnetic fields we have initiated an investigation of the evolution of magnetic fields when driven by the SASI:  Starting from a spherically symmetric, non-rotating configuration resembling a stalled supernova shock wave, we find that axisymmetric SASI-induced flows are able to amplify magnetic fields to dynamically significant levels within the time frame a supernova is expected to reach explosive conditions. ",
        "conclusions": "By considering an idealized model of a stalled supernova shock wave we have found that the SASI-induced flow beneath the shock is able to generate dynamically significant magnetic fields, in excess of $10^{15}$~G, on a time scale relevant to the supernova explosion, independent of the initial field strength. The influence of the SASI on the evolution of magnetic fields in core-collapse supernovae will need to be investigated in more realistic models, but our idealized models offer a proof of principle that (1) magnetic fields may be amplified to significant strengths {\\it even in the absence of rotation} and (2) compression may serve, at least in axisymmetry, as an effective mechanism for rapid field amplification to dynamically significant levels, contrary to conclusions reached in past MHD collapse calculations. Moreover, current 2D models \\cite{buras_etal_2006,burrows_etal_2006,bruenn_etal_2006,scheck_etal_2008} demonstrate that the SASI is common to core-collapse supernova models.  Our results then suggest that SASI magnetic field amplification and, in particular, the formation of dynamically significant fields may be common as well, and not confined to special cases---for example, to cases of rapid progenitor rotation. The key question now is: What happens in three dimensions? The amplification we document here relies on the continual focusing of fluid flow toward the polar axis, in part an artifact of our assumed axisymmetry. Hydrodynamics simulations of the SASI in 3D yield a significantly different and more complex postshock flow \\cite{blondin_mezzacappa_2007}, whose consequences for stellar core magnetic field amplification remain to be investigated.  \\ack  This work was supported by Oak Ridge National Laboratory (ORNL), managed by UT-Battelle, LLC, for the DoE under contract DE-AC05-00OR22725.  This research used resources of the National Center for Computational Sciences at ORNL, which is supported by the Office of Science of the U.S. Department of Energy under Contract No. DEAC05-00OR22725."
    },
    "0806/0806.2722_arXiv.txt": {
        "abstract": "We describe studies of star formation in various galaxies using primarily observations from the Wise Observatory. In addition to surface photometry in the broad band UBVRI, we also use a set of narrow-band H$\\alpha$ filters tuned to different redshifts to isolate the emission line. With these observational data, and using models of evolutionary stellar populations, we unravel the star formation histories of the galaxies and connect them to other parameters, such as the galaxy environment. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0806/0806.0511_arXiv.txt": {
        "abstract": "According to unified schemes of Active Galactic Nuclei (AGN), the central engine is surrounded by dusty, optically thick clouds in a toroidal structure. We have recently developed a formalism that for the first time takes proper account of the clumpy nature of the AGN torus. We now provide a detailed report of our findings in a two-paper series. Here we present our general formalism for radiative transfer in clumpy media and construct its building blocks for the AGN problem --- the source functions of individual dusty clouds heated by the AGN radiation field. We show that a fundamental difference from smooth density distributions is that in a clumpy medium, a large range of dust temperatures coexist at the same distance from the radiation central source. This distinct property explains the low dust temperatures found close to the nucleus of NGC1068 in 10 \\mic\\ interferometric observations. We find that irrespective of the overall geometry, a clumpy dust distribution shows only moderate variation in its spectral energy distribution, and the 10\\mic\\ absorption feature is never deep. Furthermore, the X-ray attenuating column density is widely scattered around the column density that characterizes the IR emission. All of these properties are characteristic of AGN observations. The assembly of clouds into AGN tori and comparison with observations is presented in the companion paper. ",
        "introduction": "Although there are numerous AGN classes, a unified scheme has been emerging steadily \\citep[e.g.,][]{Ski93, Ski02, Urry95}. The nuclear activity is powered by a supermassive black hole and its accretion disk, and this central engine is surrounded by a dusty toroidal structure. Much of the observed diversity is simply explained as the result of viewing this axisymmetric geometry from different angles. The torus provides anisotropic obscuration of the central region so that sources viewed face-on are recognized as ``type 1'', those observed edge-on are ``type 2''. From basic considerations, \\cite{Krolik88} conclude that the torus is likely to consist of a large number of individually very optically thick dusty clouds. Indeed, \\cite{Tristram07} find that their VLTI interferometic observations of the Circinus AGN ``provide strong evidence for a clumpy or filamentary dust structure''. A fundamental difference between clumpy and smooth density distributions is that radiation can propagate freely between different regions of an optically thick medium when it is clumpy, but not otherwise. However, because of the difficulties in handling clumpy media, models of the torus IR emission, beginning with \\cite{PK92}, utilized smooth density distributions instead. \\cite{RR95} noted the importance of incorporating clumpiness for realistic modeling, but did not offer a formalism for handling clumpy media.  We have recently developed such a formalism and presented initial reports of our modeling of AGN clumpy tori in \\cite*{NIE02}, \\cite*{Elitzur04}, and \\cite{Elitzur06, Elitzur07}. We now offer a detailed exposition of our formalism and its application to AGN observations. Because of the wealth of details, the presentation is broken into two parts. In this paper we describe the clumpiness formalism and construct the source functions for the emission from individual optically thick dusty clouds, the building blocks of the AGN torus. In a companion paper \\cite[part II hereafter]{AGN2} we present the assembly of these building blocks into a torus and the application to AGN observations. ",
        "conclusions": "The formalism presented here is general and can be applied to any clumpy distribution. In this study we employed it to construct the source functions of dusty clouds. As a small object, a cloud is primarily characterized by its overall optical depth \\tV. Two additional properties could potentially affect the cloud emission. One is surface smoothness. A rough, fractal-like surface can be expected to reduce the efficiency of scattering by absorbing some scattered photons and transferring them to the thermal pool. The other is shape. This factor could be studies by averaging ellipsoidal clouds over all orientations. The shape parameter would then correspond to the ellipsoid axial ratio as the cloud varied from a spherical shape to extreme elongation; the clouds constructed here can be considered representative of extremely elongated shapes. The isotropy of the external radiation field of indirectly heated clouds enabled us to explore in this case the two extremes of cloud geometrical shape: since the radiative transfer problem for a spherical cloud retains the spherical symmetry in an isotropic external field, it too can be solved by \\D. The source function is then found from $\\Si = F_\\lambda/\\Omega$, where $F_\\lambda$ and $\\Omega$ are, respectively, the flux and solid angle of the sphere at a large distance. We calculated \\Si\\ for spherical clouds imbedded in the same radiation field used in the calculations described in \\S\\ref{sec:Indirect}. The solution for a sphere with uniform density depends only on the overall optical depth \\tV, allowing direct comparison with the clouds constructed by averaging slabs of the same \\tV. We found no significant differences between the two extreme cases. We are currently studying also the direct illumination of spherical clouds, and will report our full findings in a separate publication (Kimball, Ivezi\\'c \\& Elitzur, in preparation).  Both shape and surface properties appear to be of secondary importance. The modeling of the source function presented here in which the clouds are characterized purely by \\tV\\ does seem to capture the essence of emission from a single dusty cloud, although definitive conclusions will have to await detailed comparisons with the calculations for externally illuminated spheres.  \\subsection{Dust temperature in clumpy media}  Dust temperature distributions are profoundly different in clumpy and smooth media. In smooth density distributions, dust temperature and distance from the AGN are uniquely related to each other---given the distance, the dust temperature is known, and vice versa. In contrast, as shown in \\S\\ref{sec:Tdust}, in a clumpy medium \\begin{itemize} \\item different dust temperatures coexist at the same distance from the AGN  \\item the same dust temperature occurs at different distances --- the dark side of a cloud close to the AGN can be as warm as the bright side of a farther cloud \\end{itemize} \\cite{Schartmann05} find that different components of the multi-grain mix employed in their torus modeling can have different temperatures at the same location, therefore the concept of ``dust temperature'' is ill defined at the microscopic level in such regions. Still, the fundamental differences between the temperature structures of smooth and clumpy media apply to the individual components of the grain mix and to the mixture average, as well as to the common, proper dust temperature when all the components equilibrate to the same one. These differences have profound implications for the torus emission, explaining the low dust temperatures found so close to the nuclei of NGC1068 \\citep{Jaffe04, Poncelet06} and Circinus \\citep{Tristram07}. We discuss these implications further in part II.  \\subsection{X-rays vs IR}  IR flux measurements collect the emission from the entire torus area on the plane of the sky. This flux is determined by the average number of clouds along all radial rays through the torus. In contrast, X-ray attenuation is controlled by the clouds along just one particular ray, the line of sight to the AGN. Since X-rays are absorbed by dust-free as well as dusty material, X-ray absorbing clouds will generally outnumber the torus clouds in any given direction. In fact, X-ray absorption in AGN could be dominated by the dust-free clouds (see paper II). But even for the dusty portion of the column, the number of X-ray absorbing clouds can differ substantially from the torus average. As an example, the appendix table \\ref{Table} presents a tabulation of the Poisson distribution for \\N\\ = 5, which is representative of AGN tori as shown in part II. More than 80\\% of the paths will have a number of clouds different from 5 in this case, and the probability to encounter just 1 cloud or as many as 9 is a full 20\\% of the probability to encounter the average 5.  Two type 2 sources with similar cloud properties and the same average \\N\\ will have an identical IR appearance, yet the X-ray absorbing columns in each torus could still differ by an order of magnitude. This can be expected to introduce a large scatter in X-ray observations.  Each spectral regime responds to large variations in cloud optical depth in an entirely different way. The IR emission depends on \\tV\\ through the probability for photon escape and the cloud source function (see eq.\\ \\ref{eq:IC}). Both factors saturate when \\tV\\ exceeds \\about 100. From eq.\\ \\ref{eq:Pesc}, $\\Pesc = e^{-\\cal N}$ when $\\tau_\\lambda \\gg 1$. Since this condition is met at all relevant wavelengths when $\\tV \\ga 50$, \\Pesc\\ becomes independent of \\tV. Similarly for \\Sl: Because each cloud is heated from outside, only its surface is heated significantly when \\tV\\ is large. Increasing \\tV\\ further only adds cool material, thus \\Sl\\ saturates for all relevant $\\lambda$ (similar to standard black-body emission). Indeed, figure \\ref{Fig:Sdir-tau} displays model results that cover three orders of magnitude of clump optical depth, yet the SEDs show only moderate variation, saturating altogether when $\\tV \\ga 100$. Even at $\\tV < 100$, significant spectral variety is mostly limited to $\\lambda \\la 10$ \\mic\\ for clouds along the line of sight to the AGN. In contrast, the \\tV-dependence of X-ray absorption is markedly different. Individual torus clouds are optically thin to X-rays, because the optical depth for Thomson scattering is only \\about 2\\x\\E{-3}\\tV, therefore the overall optical depth for X-ray absorption is \\N\\tV\\ (see \\S\\ref{sec:formalism}). It increases linearly with \\tV, in contrast with the saturated response in the IR regime.  The great disparity between the two spectral regions is expected to further increase the scatter in torus X-ray properties among AGN with similar IR emission. It also may help explain why the SEDs show only moderate variations in the infrared that are not well correlated with the X-ray absorbing columns \\citep[e.g.,][]{Silva04}.  \\subsection{The 10\\mic\\ feature}  In contrast with ultra-luminous infrared galaxies (ULIRGs),  AGN observations do not provide any example of deep 10\\mic\\ silicate absorption feature \\citep{Hao07}. This behavior conflicts with the results of smooth-density models \\citep[e.g.,][]{PK92} but is a natural consequence of clumpiness: as is evident from figures \\ref{Fig:Sdir-tau} and \\ref{Fig:Sdir-all}, single clumps never produce extremely deep features. This behavior reflects the flat slab temperature profile (see fig.\\ \\ref{Fig:Td-tau}) and is a realistic depiction of the situation inside an actual cloud heated from outside by a distant source. As noted by \\cite{Levenson07}, the different behavior of the 10\\mic\\ absorption in ULIRGs and AGN indicates that the dust distribution is smooth in the former and clumpy in the latter \\citep[see also][]{Spoon07}.  Clumpiness suffices by itself to explain the modest depth of the 10 \\mic\\ absorption feature in AGN. The complete behavior of the feature, including the transition from emission to absorption, involves an intricate interplay between the relative contributions of clouds at different locations and their mutual shadowing. This behavior displays a complex pattern that depends on the actual geometry of the cloud distribution. A detailed, quantitative analysis of the 10\\mic\\ feature in clumpy tori is performed in part II.   \\subsection{Conclusions}  Clumpy media differ from continuous ones in a number of fundamental ways. The low dust temperatures found close to the nucleus of NGC1068 contradict basic physical principles in smooth density distributions but arise naturally in clumpy ones. Two additional puzzling features of IR emission from AGN are simply hallmarks of clumpy dust distributions, independent of the distribution geometry: Even if the cloud configurations were spherical or irregular rather than toroidal, (1) the SED still would show only a moderate range incommensurate with the variation of X-ray attenuation, and (2) the 10\\mic\\ absorption feature would never be deep. Understanding the full range of 10\\mic\\ features observed in AGN spectra requires considerations of the actual geometry of the cloud distribution, though. This problem is addressed in part II together with other implications of clumpy tori to AGN observations."
    },
    "0806/0806.0382_arXiv.txt": {
        "abstract": "Our \\Chandra\\ X-ray Observatory archival study of intracluster entropy in a sample of \\clnum\\ galaxy clusters shows that \\halpha\\ and radio emission from the brightest cluster galaxy are much more pronounced when the cluster's core gas entropy is $\\la 30 \\ent$. The prevalence of \\halpha\\ emission below this threshold indicates that it marks a dichotomy between clusters that can harbor multiphase gas and star formation in their cores and those that cannot. The fact that strong central radio emission also appears below this boundary suggests that AGN feedback turns on when the intracluster medium starts to condense, strengthening the case for AGN feedback as the mechanism that limits star formation in the Universe's most luminous galaxies. ",
        "introduction": "\\label{sec:intro}  In recent years the ``cooling flow problem'' has been the focus of intense scrutiny as the solutions have broad impact on our theories of galaxy formation (see \\citealt{cfreview} for a review). Current models predict that the most massive galaxies in the Universe -- brightest cluster galaxies (BCGs) -- should be bluer and more massive than observations find, unless AGN feedback intervenes to stop late-time star formation \\citep{bower06, croton06, saro06}. X-ray observations of galaxy clusters have given this hypothesis considerable traction. From the properties of X-ray cavities in the intracluster medium (ICM), \\cite{birzan04} concluded that AGN feedback provides the necessary energy to retard cooling in the cores of clusters (see \\citealt{mcnamrev} for a review). This result suggests that, under the right conditions, AGN are capable of quenching star formation by heating the surrounding ICM.  If AGN feedback is indeed responsible for regulating star formation in cluster cores, then the radio and star-forming properties of galaxy clusters should be related to the distribution of ICM specific entropy\\footnote{In this paper we quantify entropy in terms of the adiabatic constant $K = kTn_e^{-2/3}$.}. In previous observational work (see \\citealt{radioquiet, d06, accept}), we have focused on ICM entropy as a means for understanding the cooling and heating processes in clusters because it is a more fundamental property of the ICM than temperature or density alone \\citep{voitbryan,voitreview}. ICM temperature mainly reflects the depth and shape of the dark matter potential well, while entropy depends more directly on the history of heating and cooling within the cluster and determines the density distribution of gas within that potential.  We have therefore undertaken a large \\Chandra\\ archival project to study how the entropy structure of clusters correlates with other cluster properties. \\cite{accept} presents the radial entropy profiles we have measured for a sample of \\clnum\\ clusters taken from the \\Chandra\\ Data Archive. We have named this project the Archive of Chandra Cluster Entropy Profile Tables, or \\accept\\ for short. To characterize the ICM entropy distributions of the clusters, we fit the equation $K(r) = K_0 +K_{100}(r/100 \\kpc)^{\\alpha}$ to each entropy profile. In this equation, \\khun\\ is the normalization of the power-law component at 100 kpc and we refer to \\kna\\ as the central entropy. Bear in mind, however, that \\kna\\ is not necessarily the minimum core entropy or the entropy at $r=0$, nor is it the gas entropy which would be measured immediately around the AGN or in a BCG X-ray coronae. Instead, \\kna\\ represents the typical excess of core entropy above the best fitting power-law found at larger radii. \\cite{accept} shows that \\kna\\ is non-zero for almost all clusters in our sample.  In this letter we present the results of exploring the relationship between the expected by-products of cooling, \\eg\\ \\halpha\\ emission, star formation, and AGN activity, and the \\kna\\ values of clusters in our survey. To determine the activity level of feedback in cluster cores, we selected two readily available observables: \\halpha\\ and radio emission. We have found that there is a critical entropy level below which \\halpha\\ and radio emission are often present, while above this threshold these emission sources are much fainter and in most cases undetected. Our results suggest that the formation of thermal instabilities in the ICM and initiation of processes such as star formation and AGN activity are closely connected to core entropy, and we suspect that the sharp entropy threshold we have found arises from thermal conduction (\\citealt{conduction}, in press).  This letter proceeds in the following manner: In \\S\\ref{sec:data} we cover the basics of our data analysis. The entropy-\\halpha\\ relationship is discussed in \\S\\ref{sec:sf}, while the entropy-radio relationship is discussed in \\S\\ref{sec:agn}. A brief summary is provided in \\S\\ref{sec:diss}. For this letter we have assumed a flat \\LCDM\\ Universe with cosmogony $\\OM=0.3$, $\\OL=0.7$, and $\\Hn=70\\km\\ps\\pMpc$. All uncertainties are 90\\% confidence. ",
        "conclusions": "\\label{sec:diss}  We have presented a comparison of ICM central entropy values and measures of BCG \\halpha\\ and radio emission for a \\Chandra\\ archival sample of galaxy clusters. We find that below a characteristic central entropy threshold of $\\kna \\approx 30 \\ent$, \\halpha\\ and bright radio emission are more likely to be detected, while above this threshold \\halpha\\ is not detected and radio emission, if detected at all, is significantly fainter. The mean \\kna\\ for clusters with and without \\halpha\\ detections are \\fha\\ and \\nfha, respectively. For clusters at $z < 0.2$ with BCG radio emission the mean \\frad, while for BCGs with only upper limits, the mean \\nfrad. While other mechanisms can produce \\halpha\\ or radio emission besides star formation and AGN, if one assumes the \\halpha\\ and radio emission are coming from these two feedback sources, then our results suggest the development of multiphase gas in cluster cores (which can fuel both star formation and AGN) is strongly coupled to ICM entropy.\\\\\\\\   We were supported in this work through CXO grants AR-6016X, AR-4017A, and NASA LTSA program NNG-05GD82G. The CXC is operated by the SAO for and on behalf of NASA under contract NAS8-03060."
    },
    "0806/0806.4042_arXiv.txt": {
        "abstract": "{% } {% We seek to find a shapelet-based scheme for deconvolving galaxy images from the PSF which leads to unbiased shear measurements. } {% Based on the analytic formulation of convolution in shapelet space, we construct a procedure to recover the unconvolved shapelet coefficients under the assumption that the PSF is perfectly known. Using specific simulations, we test this approach and compare it to other published approaches. } {% We show that convolution in shapelet space leads to a shapelet model of order $n_{max}^h = n_{max}^g + n_{max}^f$ with $n_{max}^f$ and $n_{max}^g$ being the maximum orders of the intrinsic galaxy and the PSF models, respectively. Deconvolution is hence a transformation which maps a certain number of convolved coefficients onto a generally smaller number of deconvolved coefficients. By inferring the latter number from data, we construct the maximum-likelihood solution for this transformation and obtain unbiased shear estimates with a remarkable amount of noise reduction compared to established approaches. This finding is particularly valid for complicated PSF models and low $S/N$ images, which renders our approach suitable for typical weak-lensing conditions. }{} ",
        "introduction": "Shapelets have been proposed as an orthonormal set of two-dimensional functions to quantify shapes of galaxy images \\citep{shapeletsI}. They have several convenient mathematical properties suggesting their use in measurements of weak gravitational lensing \\citep{shapeletsII}. Their application to data with low signal-to-noise ratio, however, is hampered by a number of problems. First, their enormous flexibility allows shapelets to represent random noise patterns well, which generates their tendency to escape into high orders and to overfit noisy images. Second, this same property not only affects the order, but also the scale of the best-fitting shapelet model. The scales of the shapelets decomposing noisy images is thus likely too high. These drawbacks are inherent properties of the shapelet method and have to be addressed for an effective usage in astronomical image processing. At first sight, it seems natural to limit the order of the shapelets from above to avoid overfitting, and to limit also the scale from above to prevent shapelets from creeping into the noise. On the other hand, convolution with the PSF alters the shape and spatial extent of imaged objects, and thus leaves an imprint on the maximum shapelet order and the scale size.  Guided by these considerations, we address the following question: How should the maximum order and the spatial scale of a shapelet decomposition be determined in cases where PSF convolution significantly modifies the object's appearance? This question aims at applications of shapelets in weak-lensing measurements, where the situation is particularly delicate because typically very small and noisy images are convolved with structured, incompletely known PSF kernels with scales similar to those of the images. How can deconvolution schemes be constructed in this case in order to find significant and unbiased shear estimates?  We show in Sect. \\ref{sec:convolution} and in the Appendix that mathematical sum rules exist for the shapelet orders and the squared spatial scales of original image, kernel, and convolved image. We then proceed in Sect. \\ref{sec:deconvolution} to devise an algorithm respecting these sum rules as well as possible, which leads to a deconvolution scheme based on the (possibly weighed) pseudo-inverse of a rectangular rather than the inverse of a quadratic convolution matrix. In Sect. \\ref{sec:benchmark}, we demonstrate by means of simulations with different signal-to-noise levels and PSF kernels that our algorithm does indeed perform very well, and leads in most realistic cases to substantially improved results compared to previously proposed methods. Our conclusions are summarized in Sect. \\ref{sec:conclusions}. ",
        "conclusions": "\\label{sec:conclusions} Based on an analytic consideration of shapelet convolution, we have studied algorithms for the PSF deconvolution of galaxy images in shapelet space. The starting point are sum rules for shapelet convolution, showing that the intrinsic shapelet orders of PSF and image add in the convolved image, and that the squares of their scales are also added. We suggest an algorithm respecting these sum rules as well as possible in presence of noise, whose central step is the deconvolution of the convolved image with the pseudo-inverse of the convolution matrix. Applications to simulated images have shown that our algorithm performs very well and in many cases noticeably better than previously suggested methods. We identify three main reasons for the improved performance: \\begin{itemize} \\item As the sum rule for the shapelet order shows, the convolution transports power to higher shapelet modes. The mean signal-to-noise ratio of the convolved coefficients is thus reduced. Our reduction of the order during the deconvolution increases the signal-to-noise without the need of calibration. \\item The sum rules typically require a much higher shapelet order than the $\\chi^2$ minimization, in particular if the PSF model is structured. Many of the high-order coefficients are thus highly insignificant. The significance of the coefficients is re-established by the reduction to the order $\\bar{n}_{max}^f$, which is chosen such that the shapelet expansion contains only significant coefficients. Fortunately, $\\bar{n}_{max}^f$ depends mainly on the signal-to-noise ratio of the galaxy, but only weakly: many different galaxy shapes can be deconvolved with the same $\\bar{n}_{max}^f$. Thus, binning the galaxies into broad signal-to-noise bins will suffice. \\item The STEP-2 project \\citep{Massey07.1} has shown that shear measurements generally depend strongly on the galaxy brightness. Our algorithm seems to have the advantage of lowering the influence of pixel noise in a well-defined manner. \\end{itemize} We shall proceed to study the performance of our algorithm in shear measurements under realistic conditions."
    },
    "0806/0806.1983_arXiv.txt": {
        "abstract": "The triaxial-octupole Y$_{32}$ correlation in atomic nuclei has long been expected to exist but experimental evidence has not been clear. We find, in order to explain the very low-lying 2$^-$ bands in the transfermium mass region, that this exotic effect may manifest itself in superheavy elements. Favorable conditions for producing triaxial-octupole correlations are shown to be present in the deformed single-particle spectrum, which is further supported by quantitative Reflection Asymmetric Shell Model calculations. It is predicted that the strong nonaxial-octupole effect may persist up to the element 108. Our result thus represents the first concrete example of spontaneous breaking of both axial and reflection symmetries in the heaviest nuclear systems. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0806/0806.3986_arXiv.txt": {
        "abstract": "We investigate the radio emission of $\\sim$ 185,000 quiescent (optically unclassifiable) galaxies selected from the Sloan Digital Sky Survey (SDSS).  By median-stacking FIRST cutouts centered on the optically-selected sources, we are able to reach flux densities down to the 10s of $\\mu$Jy.  The quiescent galaxy sample is composed of two subgroups inhabiting vastly different regimes: those targeted for the SDSS MAIN Galaxy Sample ($\\sim$55\\%), and those targeted for the Luminous Red Galaxy (LRG) sample ($\\sim$45\\%).  To investigate the star-formation of these quiescent galaxies, we calibrate a radio-SFR conversion using a third sample of star-forming galaxies.  We confirm a tight power-law dependence for the star-forming (SF) sample, where L$_{1.4 GHz}$ $\\sim$ (SFR)$^{1.37}$.  Comparing this SFR-indicator with indicators in the optical and UV, we derive conflicting SFR estimates for the MAIN sample quiescent galaxies.  These radio-derived SFRs intersect those calculated using the 4000-\\AA\\ break (D4000) around an SFR of 1 M$_{\\sun}$ yr$^{-1}$  and agree to within a factor of 3 over the range of SFRs.  However, we find that the radio-derived SFRs are too high relative to the SFRs estimated for similar populations of galaxies using analysis of UV emission, implying either contamination of the radio by Active Galactic Nuclei (AGN) or incomplete dust modeling.  If AGN activity is dominant in these galaxies, then a relation between AGN radio luminosity and galaxy mass is required to explain the observed trends.  For the LRGs, on the other hand, we find the radio luminosity to be relatively high (compared to the SF galaxies) and independent of SFR as derived from D4000, indicating an AGN component dominates their radio emission.  AGN-based radio emission often implies the existence of radio jets, providing evidence of a mechanism for low-level feedback in these quiescent LRGs. ",
        "introduction": "For normal galaxies -- those not deriving their radio luminosity from a nuclear ``monster\" -- the radio emission is a probe of recent star-formation \\citep{con92}.  In particular, the emission below 30 GHz is thought to be largely dominated by synchrotron radiation from ultrarelativistic electrons accelerated in supernova remnants.  As most of the remnants responsible for this non-thermal emission result from massive, short-lived stars, the resulting radio emission tracks recent star-formation.  Further evidence in favor of this claim comes from the existence of a tight radio/far-infrared (FIR) correlation, which has been found to hold over several orders of magnitude in radio luminosity \\citep{har75, cram98, yun01}.  The complex astrophysics connecting the integrated radio emission to supernova explosions has led to a historical preference for other methods.  Normal galaxies, with quite low radio flux densities, are also usually below the sensitivity of typical wide-area radio surveys.  Alternative indicators of star-formation rate (SFR) include continuum emission in both the near-ultraviolet and the far-infrared \\citep{don84, har73}.  Another common diagnostic, referred to later in this paper, is the use of integrated emission-line fluxes \\citep{coh76, ken83}.  (For additional indicators, see the excellent review by \\citet{ken98}.)  However, star-formation is not the only process that produces radio emission; Active Galactic Nuclei (AGN) can also be strong emitters in the radio.  This ambiguity has led to much debate, particularly for the less well-studied sub-mJy population.  On the one hand, there is a widely held belief that the population of sub-mJy radio sources is dominated by star-forming galaxies \\citep{con84, win85, kro85, wal97}. For example, using an infrared-selected sample, \\citet{boy07} found the radio/FIR correlation persisted down to $\\mu$Jy flux densities, and they conclude that the radio emission of these sources is dominated by star-formation  (with redshift still unknown).  However, the hypothesis that there are a significant number of AGN at these low flux densities has recently received some attention as well.  \\citet{jar04}, in a theoretical paper, suggested that the flattening in the observed counts of the normalized local luminosity function below S $\\sim$ 1 mJy is due to a ``radio-quiet\" AGN population.  In a follow-up observational study, \\citet{sim06} concluded that radio-quiet AGN comprise a significant population below $\\sim$ 300 $\\mu$Jy, and that they may, in fact, dominate.  Studying the general properties of the low-luminosity radio population, however, has proven to be difficult; the perennial trade-off between coverage area and depth discourages surveys that are both deep and wide.  In this paper, we take advantage of the overlapping coverage of two surveys, FIRST (Faint Images of the Radio Sky at Twenty Centimeters) and the SDSS (Sloan Digital Sky Survey), to study the faint radio emission of a large sample of optically-normal (``quiescent\") galaxies.  By selecting galaxies with no strong optical emission lines from star-formation or AGN activity, we construct a sample that is likely to be inherently radio-quiet in nature.  Any radio emission we do find, then, must result from either extremely low-level SF and/or AGN activity, or activity that is obscured in the optical. By observing how this emission behaves as a function of other SFR indicators, we will attempt to discern the origin (SF or AGN) of the radio flux, and, in the absence of AGN activity, the radio emission will give us a handle on the amount of star-formation occurring in these galaxies.  \\citet{sal07} found that galaxies with no H$\\alpha$ detectable in their SDSS spectra were mostly normal early-type galaxies with predominantly red optical colors, though some have relatively blue UV to optical colors (see also: Rich et al. 2005).  They calculated that some of these galaxies had UV-derived SFRs as high as 1 M$_{\\sun}$ yr$^{-1}$, and a significant portion of those showed structure, like traces of faint spiral arms, or disturbed light profiles.  They identified an evolutionary sequence from star-formation to quiescence via nuclear activity.   Using a sample of nearby, early-type galaxies in the MOSES project (Morphologically Selected Ellipticals in SDSS), \\citet{schaw07} reached a similar conclusion.  \\citet{yan06} looked at optically-quiescent galaxies with red rest-frame colors and found that they were virtually indistinguishable in the color-magnitude-concentration space from galaxies classified as low-ionization nuclear emission-line regions (LINERs).  Still, very little has been done on the radio properties of optically-quiescent galaxies, and we intend to use this relatively unexplored wavelength regime to add to our knowledge of such galaxies.  For the analysis of such faint radio sources, we have utilized the stacking technique presented in \\citet{whi07}, hereafter referred to as Paper I.  This technique significantly reduces the noise in the combined images, allowing the investigation of median radio properties for sources significantly below the FIRST detection threshold without requiring countless more hours of telescope time.  In addition, stacking all of the fields guarantees that the radio data will not be restricted by survey limits.  In de Vries et al. (2007; hereafter, Paper II) we applied this method to the analysis of low-luminosity AGN, concluding that radio emission from star-formation can dominate over that due to the AGN emission in composite systems.  We begin by  describing the radio data and the details of the sample selection (\\S \\ref{data}).  We then go into the details of the stacking method (\\S \\ref{stack}).  In \\S \\ref{results} we present our analysis and a detailed discussion of the results; \\S \\ref{conclusions} contains our concluding remarks.  Throughout this paper we use the standard FRW cosmology of H$_0$ = 70 km s$^{-1}$ Mpc$^{-1}$, $\\Omega_{\\Lambda}$ = 0.7, and $\\Omega_{M}$ = 0.3. ",
        "conclusions": "\\label{conclusions}  We have investigated the radio emission from a sample of $\\sim$ 185,000 optically-selected quiescent galaxies from the Sloan Digital Sky Survey.  By stacking FIRST cutouts centered on the galaxy positions, we reach median radio flux densities down to the 10s of $\\mu$Jy.  We calibrate the relation between 1.4 GHz radio emission and optically-derived SFR using a sample of star-forming galaxies, allowing us to look for star-formation in their quiescent counterparts.  The main findings of this study are summarized below.   \\begin{itemize}  \\item We find a tight correlation between radio luminosity and SFR for the star-forming galaxy sample, with L$_{1.4GHz}$ going as (SFR)$^{1.37}$.  \\item The quiescent galaxy sample consists of two subgroups: those targeted as part of the MAIN Galaxy Sample, and those targeted as part of the Luminous Red Galaxy (LRG) Sample.  The former show partial overlap with the star-forming sample on a plot of L$_{R}$ vs. optically-derived SFR, while the latter show a relatively constant, high value of median radio luminosity that does not vary with SFR.  \\item   For the LRGs, the correspondingly high rates of radio-derived star-formation conflict with those estimated by \\citet{brinch04} via calibration of the 4000-\\AA\\ break by two orders of magnitude on a plot of derived-SFR vs. M$_z$.  This result, combined with the relative constancy of L$_{R}$ with D4000 SFR, suggests that the radio emission of these quiescent LRGs derives not from star-formation but from an AGN component.  This implies the possible presence of radio jets and a mechanism for low-level feedback in these galaxies.  \\item The MAIN sample galaxies, on the other hand, exhibit a L$_R$-SFR trend which intersects the SF sample around 1 M$_{\\sun}$ yr$^{-1}$ and deviates by no more than a factor of 3 over the whole range of SFRs.  This agreement as well as the dependence of L$_{R}$ with SFR suggests that the  L$_{R}$ might arise in part from star-formation (particularly in so far as the AGN of Paper II show no such dependence.)  If we assume the radio emission of these galaxies originates from star-formation, the deviation at high/low values of SFR might then be due to the different diagnostic measures used to estimate SFR for the two samples (i.e., 4000-\\AA\\ break vs. emission-lines).  However, we cannot rule out the possibility that the radio emission may have a large or even dominant AGN component, with a L$_R$-mass relation causing the observed correlation.  \\item Composite spectra of the MAIN sample show Balmer absorption strong enough to plausibly swallow any H$\\alpha$ emission present.  This is supported by some [NII] emission even in the absence of H$\\alpha$.  The [OII] emission, however, increases only slightly with D4000 SFR, giving reason to doubt the validity of D4000 (and thus radio emission) as a proxy of star-formation in these galaxies until we understand why the line emission is so highly-suppressed.  \\item The MAIN quiescent sample shows a correlation between radio luminosity and z-band absolute magnitude of the form: log($L_{1.4}$) = (-0.92 $\\pm$ 0.02)$M_z$ + (7.6 $\\pm$ 0.4).  The SFRs implied by radio emission reach above 10 M$_{\\sun}$ yr$^{-1}$, yet lie just below the objects classified as ``star-forming\" for a given absolute magnitude.   If we assume the radio emission stems from SF, one possible interpretation is that the strength of the optical emission lines needed to constitute a ``detection\" varies with galaxy luminosity, and using our radio-stacking method allows us to probe low levels of star-formation despite SF dilution in the optical.  If we assume the radio emission stems from AGN activity, then both the D4000 and radio-derived SFRs are too high (for separate reasons), and the observed trend implies a relation between AGN radio luminosity and galaxy mass.  \\item UV-derived SFRs from GALEX are 2-3 orders of magnitude lower than the SFRs predicted by radio-stacking for the MAIN quiescent galaxy sample.  This indicates possible AGN contamination in the radio emission of the quiescent galaxies, incomplete modeling of the dust attenuation for the UV emission, or both.  The different processes at play obviously need to be better understood before we can choose the most useful SFR diagnostic tool for these optically-quiescent galaxies.   \\end{itemize}  In short, we use a novel stacking technique to provide strong evidence that quiescent LRGs harbor AGN.    This result has clear implications for galaxy formation scenarios, providing observational support to the existence of a mechanism for low-level feedback in these galaxies.  The results for the quiescent MAIN sample galaxies are more nebulous, with the discrepancies between the different SF indicators clearly highlighting how much we still have to learn about these persistently unclassifiable galaxies.  Overall, we believe the technique utilized here shows great promise as a tool to investigate further the radio emission of various sub-populations in the vast community of radio-quiet galaxies."
    },
    "0806/0806.3381_arXiv.txt": {
        "abstract": "We analyze macroscopic effects of TeV-scale black holes, such as could possibly be produced at the LHC, in what is regarded as an extremely hypothetical scenario in which they are stable and, if trapped inside Earth, begin to accrete matter.  We examine a wide variety of TeV-scale gravity scenarios, basing the resulting accretion models on first-principles, basic, and well-tested physical laws.  These scenarios fall into two classes, depending on whether accretion could have any macroscopic effect on the Earth at times shorter than the Sun's natural lifetime.  We argue that cases with such effect at shorter times than the solar lifetime are ruled out, since in these scenarios black holes produced by cosmic rays impinging on much denser white dwarfs and neutron stars would then catalyze their decay on timescales incompatible with their known lifetimes.  We also comment on relevant lifetimes for astronomical objects that capture primordial black holes.  In short, this study finds no basis for concerns that TeV-scale black holes from the LHC could pose a risk to Earth on time scales shorter than the Earth's natural lifetime. Indeed, conservative arguments based on detailed calculations and the best-available scientific knowledge, including solid astronomical data, conclude, from multiple perspectives, that there is no risk of any significance whatsoever from such black holes. ",
        "introduction": "One of the most spectacular manifestations of nature realizing certain extra-dimensional scenarios\\cite{ArkaniHamed:1998rs,Antoniadis:1998ig,RSi,GKP} could be the production of microscopic black holes at the LHC~\\cite{Giddings:2001bu,Dimopoulos:2001hw}.\\footnote{For reviews with more references see \\cite{Giddings:2001ih,Giddings:2007nr}.} These are expected to undergo prompt, quasi-thermal, Hawking~\\cite{Hawking:1974sw} decay to large multiplicities of elementary particles, leading to very characteristic final states. It has been suggested~\\cite{Unruh:2004zk,Vilkovisky:2005db}, however, that black hole decay via Hawking radiation may not be a universal feature, and could for example depend on the details of the Planck-scale degrees of freedom.  While this suggestion is not based on any complete microphysical picture, and furthermore appears contradictory to basic quantum-mechanical principles\\footnote{See, for example, \\cite{'tHooft:1996tq}.},  it does raise a possible question about stability of microscopic black holes that might be produced at the LHC, in TeV-scale gravity scenarios.  This in turn has led some to express concern about the fate of their evolution: could their accretion pose any threat to the Earth? This is the question addressed in this paper.  The structure of this work, and a succinct summary of our findings, are outlined here. We begin our work by reviewing what are widely regarded as quite convincing arguments for the robustness of the exceedingly rapid decay of microscopic black holes. We then discuss the arguments asserting that macroscopic consequences of new and unknown particles can be ruled out by the lack of evidence for their effects from production by cosmic rays hitting the surface of the Earth or other astronomical bodies. We argue here that charged black holes will lose enough energy to stop when traversing the Earth or the Sun, via standard electromagnetic processes. Since black holes would be typically produced by the collision of quark pairs, whether in cosmic-ray interactions or at the LHC, they would often be initially charged. To the extent that no mechanism leads to their neutralization, the cosmic-ray based argument for their being harmless is therefore robust. Their neutralization through the Schwinger mechanism proceeds  according to quantum principles like those underlying Hawking radiation.  There is therefore no concrete framework where neutralization occurs without Hawking decay taking place as well, leading to a likely contradiction in assuming that stable black holes must be neutral.  We nonetheless make the hypothesis that this odd situation could occur, and analyze the possible effects of such neutral and stable black holes, beginning with a review of some essential features of gravity and black holes in $D> 4$ dimensions, including both large- and warped-extra-dimensions models.  Next, we develop the formalism to describe the evolution of such black holes trapped inside the Earth or inside dense objects such as white dwarfs and neutron stars.  We introduce and discuss accretion scenarios that apply within nuclear, atomic and macroscopic matter, evaluating the time scales corresponding to various phases in the evolution of a growing black hole.  We establish upper and lower limits to the rate at which accretion can take place, building on very basic principles such as conservation laws and classical and quantum dynamics. This is possible since, in order for accretion to become macroscopic, it is necessary that a black hole exerts its influence over distance scales much larger than its event horizon. Such black holes are very small, and their accretion power, if limited to absorbing particles that have impact parameters of the order of the Schwarzschild radius, is typically insufficient to cause macroscopic growth. At large distances the physical processes become independent of the short-distance properties of the black hole, which only acts through its attractive potential, and as a mass sink. In this long-distance domain, classical or quantum dynamics are well-tested, making the study of accretion possible, independently of our detailed knowledge or ignorance about the physics inside the black hole.  We investigate two concrete scenarios, for various configurations of the extra dimensions: growth of hypothetical LHC-produced stable black holes within the Earth, and growth of such black holes produced by cosmic rays impinging on dense stars, such as white dwarfs or neutron stars.  In both cases, we make conservative assumptions, namely fastest possible growth in Earth, and slowest possible growth in a dense star.  The first scenario shows that, if the radius of crossover from higher-dimensional gravity to four-dimensional gravity is less than about 200 \\AA, the natural lifetime of the solar system is too short to allow significant growth of stable black holes that may be captured inside the Earth.  In contrast, in a scenario where the crossover radius exceeds the 200 \\AA\\ scale, accretion times could be shorter than the solar time scale.  In this case, however, examination of the latter dense-star scenario then produces an  argument that, given observational data setting the lifetimes of such objects at a billion years or more, such stable black holes cannot in fact exist for a crossover radius greater than $\\simeq15$\\AA.  These arguments thus conclude with the exclusion of any relevant consequences for Earth due to the evolution of black holes possibly produced by the LHC.  While the main line of the argument is relatively straightforward, being based on well-established macroscopic properties of matter, several issues require an in-depth study in order to provide a robust basis to our conclusions. When studying accretion, for example, we need to consider the possibility of effects such as the Eddington limit, where radiation emitted during the infall of matter would slow down the accretion via radiative pressure, constraining the accretion rate of a black hole and thus possibly slowing down its growth inside Earth or a dense star. The estimates of the number of black holes produced by cosmic rays need to rely on a solid understanding of the latter's spectra and composition. Here we therefore consider the worse-case scenarios where, contrary to mounting experimental evidence, the cosmic primaries are dominantly heavy nuclei rather than protons. The loss of acceptance for the highest-energy cosmic ray flux hitting dense stars, due to the intense surface magnetic fields, is also an important element that we analyze. Finally, the slow-down and stopping of relativistic black holes inside dense stars requires a detailed study of energy loss in gravitational scattering and absorption processes.  We note that our bounds are conservative.  In particular, at each point where we have encountered an uncertainty, we have replaced it by a conservative or ``worst case\" assumption.  For this reason, our bounds can likely be improved through further work removing such uncertainties. Furthermore, we do not exclude that there could be other independent arguments, based for example on astrophysical or cosmological production of microscopic black holes, leading to the exclusion of stable black holes, or of their macroscopic effects.  In outline, the next section summarizes existing arguments (and their extension) against risk from TeV-scale black holes, namely the robustness of quantum black hole decay, and constraints from cosmic rays impinging on objects in the solar system.  Section three then summarizes aspects of TeV-scale gravity scenarios and the corresponding black holes.  Section four discusses black hole accretion in Earth.  Section five gives an in-depth analysis of the question of stopping of neutral black holes, via gravitational interactions.  Section six describes production of black holes by cosmic rays impinging on white dwarfs.  Section seven describes the resulting accretion of a white dwarf, and thus derives constraints on hypothetical TeV-scale scenarios that otherwise might have been of concern.  Section eight describes similar constraints arising from black hole production on, and accretion of, neutron stars.  A summary and conclusions appear in section eight.  Many of the technical details are treated in a series of  Appendices. Appendix~\\ref{app:bondi} summarizes macroscopic Bondi accretion, which is the canonical framework to deal with the flow of matter into a black hole. For its use in our context, Bondi accretion is generalized here to higher-dimensional gravitational fields.  Appendix~\\ref{app:rerad} discusses the effects of radiation emitted during accretion, and  examines the radiative transport conditions relevant to the question of an Eddington limit for the growth of microscopic black holes. In Appendix~\\ref{app:scattering} we derive the main equations governing the particle scattering and capture in the field of a $D$-dimensional black hole,  and Appendix~\\ref{app:accretion} proves few facts required by the study of stopping by white dwarfs. Appendix~\\ref{app:prod} reviews the formalism for the calculation of black hole production at the LHC and in the collisions of cosmic rays, including the case of high-energy neutrinos.  Here we discuss the impact of different assumptions about the composition of the high-energy cosmic rays, we evaluate the production rates under the various assumptions, and establish the most conservative lower limits on such rates. We also discuss the production properties of black holes produced by cosmic rays, to serve as initial conditions for their slow-down  in dense stars. In Appendix~\\ref{app:boundBH} we calculate the probability that black holes produced at the LHC are trapped in the Earth's gravitational field. We argue that, independently of our general conclusion that trapped black holes would not accrete macroscopically, the velocity spectrum of black holes produced during the LHC lifetime in allowed extra-dimensional scenarios is such that they would  typically escape the Earth's attraction. Appendix~\\ref{app:synch} discusses the impact of magnetic fields on the penetration of cosmic rays down to the surface of dense stars.  Appendix~\\ref{app:background} addresses the production of black holes by cosmic rays hitting ``background'' objects, such as neutron star companions, the interstellar medium, or even dark matter. Here we discuss the conditions for the companion of a neutron star in X-ray binary systems to efficiently act as a beam-dump for cosmic rays, leading to the production of black holes and their capture in the neutron star, unimpeded by the presence of magnetic fields.  A final Appendix contains useful conversion factors and physical constants that are used throughout the paper. ",
        "conclusions": "In this paper we have studied accretion of hypothetical stable TeV-scale  black holes in two primary contexts: the Earth, and compact stars -- white dwarfs and neutron stars.  For Earth, we identified two main evolution domains: that where the black hole's gravitational range of influence is less than the atomic scale, and that where it is greater.  An important distinction occurs depending on where the crossover radius $R_C$ to four-dimensional behavior lies relative to the atomic scale.  In particular, if $R_C$ is in the subatomic regime, evolution is four-dimensional at both subatomic and macroscopic scales; this case includes unwarped scenarios with $D\\geq8$.  In this case we have argued that this evolution occurs on times longer than the expected natural solar lifetime, in two different ways: via a microscopic argument, and via a macroscopic, hydrodynamic argument.  In both approaches, we used conservative assumptions, leading to the largest accretion rates and to fastest evolution.  At the end of the first phase, at times ${\\cal O}(10^{11})$ years, the mass of the black hole is still small, with a mass of less than a megaton. Such statements extend to more general warped scenarios, and crossover scales up to $\\simeq200$\\AA\\ lead to accretion times longer than the Earth's natural lifetime.   On the other hand, those cases where $R_C\\gsim 15$~\\AA\\ have been treated by deriving bounds from white dwarfs, and also from neutron stars.  In particular, we have argued that in such scenarios cosmic rays will produce black holes on such astronomical objects, and that these objects will stop even these very high-momentum black holes.  We then studied accretion, showing that accreting black holes will disrupt such objects on time scales short as compared to their observed lifetimes.  In particular, we found a general relationship (\\ref{accrat}) between accretion times for Earth and for white dwarfs, which, when combined with white dwarf ages exceeding $10^9$ years, provides a very strong constraint.  Thus, the implication of these arguments is that such scenarios, where Earth could be disrupted on time scales short as compared to its natural lifetime, are ruled out.   We summarize here the origin of our constraints, as a function of dimensionality $D$ and of black hole mass $M$. \\begin{description} \\item{$D=5$:} The bounds on evolution scales on Earth for $D=5$ in the case of maximum allowed crossover radius $R_C$ are quite short. This is a result of the higher-dimensional force law extending well into the macroscopic regime. On the other hand,  the greater interactivity of $D=5$ black holes makes it possible for those produced by cosmic rays to get promptly trapped in both white dwarfs and neutron stars. In white dwarfs they are produced abundantly, with build-up time scales of the order of few thousand years even assuming a cosmic ray composition of 100\\% Fe, and at the largest  Planck mass of interest. After being produced and trapped, in these extreme scenarios they quickly accrete to masses comparable to that of  the star on time scales that can be short when measured in years. This would make it impossible for any white dwarf with a mass of the order of one solar masses to have survived longer than  a few thousand years, contrary to  observations.  Scenarios with increased warping have correspondingly lower $R_C$ and longer accretion times.  In particular, once $R_C\\lsim 200$\\AA, accretion times on Earth exceed its future lifespan. \\item{$D=6$:} The bounds on evolution times on Earth for $D=6$ are of the order tens of thousand years, thus short on geological time scales. Once again the main reason is the large extra-dimensional radius, and the high capture rates. As in the $D=5$ cases, such black holes produced by cosmic rays can be stopped inside dense stars. The production rates are even larger than in $D=5$, and the star accretion time scales for unwarped $D=6$ are comparable to the maximum-$R_C$ version of that scenario.  With increased warping, again $R_C$ decreases and accretion times increase. \\item{$D=7$:} The bound on time scale for their macroscopic evolution on Earth is in the range of 6--80 billion years, depending on the black hole mass. Furthermore, $D=7$ black holes would be produced plentifully by cosmic rays on white dwarfs, and be stopped inside their surface, if the white dwarf mass is larger than 1.1 solar masses. The evolution times would be longer than in $D=5$ and 6, reflecting the lower growth rate that keeps them microscopic for billion years inside the Earth. But within the very conservative estimate of 80 million years (for a $M=\\Msun$ star, and shorter in the more massive cases) their accretion process of the white dwarf would be completed. Massive white dwarfs older than a few hundred million years would therefore be ruled out in these scenarios, once again contrary to observation.  Again, warping decreases $R_C$, thus increasing accretion times. \\item{$D\\ge 8$:} For these black holes the bound on evolution time on Earth is extremely long, with times of size 100 billion years. This is due to the radius of the extra dimensions being smaller than 1~\\AA, thus forcing most of the evolution to take place in 4 dimensions, where gravity is a totally negligible force. Warping only magnifies this effect.  In spite of this slow growth, these black holes would still grow fast enough inside a neutron star to consume it within about ten million years. The significant production rates on neutron stars when $D\\ge 8$, and the existence of billion-year old X-ray binary systems, provide therefore additional evidence that such black holes either do not exist, or decay promptly. \\end{description}  We also note that these bounds likely extend in case other objects are imagined that could result from high-energy collisions in the relevant energy ranges, that have weak-scale cross sections, and that could threaten the long-term stability of matter.   We conclude by first summarizing the conditions needed for our bounds to be necessary to rule out a possible risk.  In order for our bounds to have relevance, a sequence of unlikely things would have to be true.  First, TeV scale gravity, with a Planck scale no higher than a few TeV, would have to be correct, so that black holes can be produced at LHC.  Most workers consider this to be a fascinating possibility, but also a somewhat unlikely possibility.  Second, black hole radiance, which has been deeply studied from a number of theoretical perspectives, would have to be wrong, {\\it and} more general quantum mechanical arguments for black hole instability would have to be wrong.  Most workers consider this to be an exceedingly improbable, if not impossible, scenario. Finally, one would need a mechanism to shut off the quantum effects responsible for Hawking radiation, but still leave intact either the quantum effects responsible for Schwinger discharge, or some other neutralization mechanism that acts to discharge the resulting stable black holes.  It is very difficult to conceive of a consistent physical framework that provides such a mechanism.  In the event that all these conditions are  satisfied, one can turn to the considerations of this paper to assess the possible impact on Earth.  This paper has argued that in order for such a scenario to have an impact on Earth at time scales short as compared to the natural lifetime of the solar system, in the five billion year range, the configuration of extra dimensions would have to be such that gravity doesn't transition to four-dimensional behavior until around the $200$\\AA\\ scale.  This apparently requires additional fine-tuning, reducing the likelihood even further. But beyond that, this paper has argued that such scenarios are ruled out by the longevity of known white dwarfs, on billion-year time scales.  In such a scenario, cosmic ray-produced black holes should have catalyzed white dwarf destruction on significantly shorter time scales.  Moreover, decay of observed neutron stars would also have been catalyzed, unless both of two unlikely possibilities are realized, namely that the composition of ultrahigh-energy cosmic ray primaries is dominantly heavy elements, and ultrahigh energy cosmic ray neutrinos either are not produced, or have suppressed gravitational interactions with partons. To summarize, the present study argues for the following {\\it additional} layers of safety, beyond those that would have to fail to make this study relevant:  \\begin{description} \\item[{\\bf 1.}] {\\bf Only in scenarios such that the crossover scale to four-dimensional gravity is larger than about $ 200$~\\AA\\ does one have significant accretion at times short as compared to the natural lifetime of Earth.  This is a-priori unlikely, due to the additional fine-tuning required to realize such a TeV-scale gravity scenario.}  \\item[{\\bf 2.}]{\\bf In these scenarios where the bound on black hole accretion time on Earth is short as compared to natural time scales, white dwarfs would likewise be accreted, on much shorter time scales, in contradiction to observation.}  \\item[{\\bf 3.}]{\\bf Unless cosmic rays have dominantly a very heavy composition, and moreover either the expected neutrino flux doesn't exist or has unusual gravitational couplings to hadronic matter, neutron star decay would likewise be catalyzed on time scales contradicting observation.} \\end{description}  In short, this study finds no basis for concerns that TeV-scale black holes from the LHC could pose a risk to Earth on time scales shorter than  the Earth's  natural lifetime.  Indeed, conservative arguments based on detailed calculations and the best-available scientific knowledge, including solid astronomical data, conclude, from multiple perspectives, that there is no risk of any significance whatsoever from such black holes.  \\vskip.2in  \\noindent{\\bf Acknowledgements} We are grateful to many colleagues who helped us in the course of the nine months of this project, providing valuable guidance and advice on the many facets of our study, helping us identify the key issues and pointing us to the relevant literature sources.  Among these: J. Arons, J. Ellis, M. Fairbairn, G. Giudice, G. Horowitz, D. Ida, Y. Kanti, G. Landsberg, D. Marolf, J. March-Russell, K.-Y. Oda, S. Park, J. Polchinski, T. Rizzo, S. Rychkov, M. Salaris, M. Srednicki, N. Toro, I. Tkachev, M. Vietri, U. Wiedemann and T. Wiseman.  The input of L. Anchordoqui, A. Dar, A. De Rujula, S. Dodelson, D. Hooper and G. Gelmini helped us clarify several issues related to the cosmic ray composition, and that of L. Balents and D. Scalapino, questions of atomic structure relevant to subatomic accretion.  We would particularly like to thank O. Blaes for a number of important discussions on Bondi and Eddington accretion, K. Shen and G. Schmidt for guidance on white dwarf structure and populations, and L. Bildsten both for many crucial discussions and for supplying information about neutron-star binary systems.  We would also like to thank J. Hartle for earlier collaboration on this work, and for many useful discussions.  The work of SBG was supported in part by the U.S. Department of Energy under Contract DE-FG02-91ER40618, and by grant RFPI-06-18 from the Foundational Questions Institute (fqxi.org).  \\appendix"
    },
    "0806/0806.0898_arXiv.txt": {
        "abstract": "A method to derive stationary solutions of the relativistic Vlasov-Maxwell system is explored. In the non-relativistic case, a method using the Hermite polynomial series to describe the deviation from the Maxwell-Boltzmann distribution is found to be successful in deriving a few stationary solutions including two dimensional one. Instead of the Hermite polynomial series, two special orthogonal polynomial series, which are appropriate to expand the deviation from the Maxwell-J\\\"uttner distribution, are introduced in this paper. By applying this method, a new two-dimensional equilibrium is derived, which may provide an initial setup for investigations of three-dimensional relativistic collisionless reconnection of magnetic fields. ",
        "introduction": "INTRODUCTION} Relativistic plasmas play important roles in high energy phenomena. In nature, gamma-ray burst afterglow is considered to be a phenomenon resulting from interaction of ultrarelativistic plasmas with a Lorentz factor $\\gamma\\simeq 10^2-10^3$ \\cite{SP90,Paczynski}. Detailed observations of high energy astrophysical phenomena, such as supernova remnants, gamma-ray bursts, and active galactic nuclei, reveal that non-thermal components in their spectrum comprise of radiation from relativistic particles. In recent laser experiments, interaction of laser and plasmas is realized in highly relativistic regime. These phenomena often involve sufficiently rarefied plasma that the effect of collisions is negligible. In order to give plausible explanations for such phenomena, many authors have devoted their energies to investigations into properties of relativistic plasmas, i.e., instabilities, radiation, and stationary equilibria.  Here I propose a new method to obtain stationary equilibrium configurations of collisionless plasmas described by the relativistic Vlasov-Maxwell system. Such solutions are of great interest because they give self-consistent configurations of relativistic plasmas, which are used as an initial setup for studies on the relativistic collisionless reconnection of magnetic fields \\cite{parker,sweet,petschek}. In some recent numerical simulations of the relativistic collisionless magnetic reconnection \\cite{ZH01,J04}, relativistic extensions of the harris sheet equilibrium \\cite{harris} are used as the initial setup. While stationary solutions of the non-relativistic Vlasov-Maxwell system are well explored \\cite{harris,bennett,BGK57,mahajana,mahajanb,AP99,BBSS01,mottez,CMPC05,MP07}, exploration of the relativistic Vlasov-Maxwell equilibria is not sufficient. Although Ref. \\cite{Braasch} discussed the existence of such solutions and derived a special class of them, studies on the derivation of the concrete expressions of the relativistic Vlasov-Maxwell equilibria are still rare.  In the previous paper \\cite{SS08}, a novel method to construct stationary solutions of the non-relativistic Vlasov-Maxwell system is proposed. In the method, the key concept is to describe the deviation of the distribution function from the Maxwell-Boltzmann by orthogonal polynomial series. The Hermite polynomial series is chosen for the purpose, because the weight function is Gaussian, which is equivalent to the Maxwell-Boltzmann distribution. Ref. \\cite{SS08} reveals that this choice is appropriate to recover some well-known equilibria (the Harris sheet and the Bennet pinch) and derive a new equilibrium. In this paper, I will extend the previous method to deal with relativistic plasmas. Instead of the Maxwell-Boltzmann distribution, I use  a relativistic extension of the Maxwell-Boltzmann distribution, i.e., the Maxwell-J\\\"uttner distribution \\cite{J11} as the weight function of these polynomial series. To the best of my knowledge, there is no appropriate orthogonal polynomial series that describes the deviation from the Maxwell-J\\\"uttner distribution as the Hermite polynomial does for non-relativistic plasmas. So I introduce two orthogonal polynomials in this paper. Applying the method, I can derive a new two-dimensional equilibrium, which is a relativistic extension of the equilibrium proposed in the previous paper. This equilibrium may provide an initial setup for numerical simulations of the relativistic collisionless reconnection.  This paper is organized as follows. In Sec.\\ref{form}, I describe the procedure of the method. In Sec.\\ref{ex}, I apply the method to derive a stationary solution that I can treat it analytically. I conclude this paper in Sec.\\ref{conc}. Appendix gives some algebraic relations used in this paper. ",
        "conclusions": ""
    },
    "0806/0806.0851_arXiv.txt": {
        "abstract": "{We present time series photometry for six partial transits of GJ\\,436b obtained with the Fine Guidance Sensor instrument on the \\textit{Hubble Space Telescope} (\\textit{HST}). Our analysis of these data yields independent estimates of the host star's radius $R_{\\star} = 0.505^{+0.029}_{-0.020}\\,R_{\\sun}$, and the planet's orbital period $P = 2.643882^{+0.000060}_{-0.000058}$\\,d, orbital inclination $i = 85.80\\degr\\ ^{+0.21\\degr}_{-0.25\\degr}$, mean central transit time $T_{c} = 2454455.279241^{+0.00026}_{-0.00025}$\\,HJD, and radius $R_{p} = 4.90^{+0.45}_{-0.33}\\,R_{\\oplus}$. The radius we determine for the planet is larger than the previous findings from analyses of an infrared light curve obtained with the \\textit{Spitzer Space Telescope}. Although this discrepancy has a 92\\% formal significance (1.7$\\sigma$), it might be indicative of systematic errors that still influence the analyses of even the highest-precision transit light curves. Comparisons of all the measured radii to theoretical models suggest that GJ\\,436b has a H/He envelope of $\\sim$ 10\\% by mass. We point out the similarities in structure between this planet and Uranus and Neptune and discuss possible parallels between these planets' formation environments and dynamical evolution. We also find that the transit times for GJ\\,436b are constant to within 10\\,s over the 11 planetary orbits that the \\textit{HST} data span. However, the ensemble of published values exhibits a long-term drift and our mean transit time is 128\\,s later than that expected from the \\textit{Spitzer} ephemeris. The sparseness of the currently available data hinders distinguishing between an error in the orbital period or perturbations arising from an additional object in the system as the cause of the apparent trend. Assuming the drift is due to an error in the orbital period we obtain an improved estimate for it of $P = 2.643904 \\pm 0.000005$\\,d. This value and our measured transit times will serve as important benchmarks in future studies of the GJ\\,436 system.} ",
        "introduction": "Transiting exoplanets present a unique opportunity to evaluate planet formation and evolution models. The principle information obtainable from the study of such planets are their masses, radii, and orbits. This complete data set cannot be obtained for non-transiting planets using any other methods with current technology. The radius of a planet for a given mass and insolation is a strong constraint for theoretical models and yields information about the planet's structure. Additionally, carefully designed observations of transiting exoplanets can yield data that were, up until recently, widely thought to only be the purview of future direct imaging studies. Such data include, but are not limited to, exoplanets' chemical compositions, thermal emission, albedos, and energy redistributions \\citep[for a recent review of the field see][]{charbonneau07}.  The recent discovery by \\citet{gillon07b} that the planet GJ\\,436b, which was originally identified with Doppler spectroscopy by \\citet{butler04}, transits its host star is an important advance in this rapidly burgeoning field. It is the first member of the emerging class of planets known as ``Hot Neptunes'' (i.e. planets having masses similar to that of Neptune and orbital semimajor axes $\\sim$ 0.01 -- 0.10\\,AU) that has been observed to transit. Therefore, detailed observational study of an ice giant outside our solar system can now be carried out.  Since the reported detection of transits for GJ\\,436b there have been a flurry of follow-up observations and analyses of the system. So far, published results include complete analyses (i.e. data reduction and light curve modeling) by two separate groups of a single transit light curve and secondary eclipse light curve obtained at 8\\,$\\mu$m using the InfraRed Array Camera (IRAC) on the \\textit{Spitzer Space Telescope} (hereafter \\textit{Spitzer} for short). \\citet[][hereafter G07]{gillon07a} and \\citet[][hereafter D07]{deming07} analyzed the transit light curve and determined a precise mass, radius, and orbit for the planet. Both groups obtained consistent results and both concluded that the planet must have a gaseous H/He envelope similar to Neptune based on comparisons of its radius with theoretical models. Additionally, \\citet[][hereafter S08]{southworth08} has recently analyzed the reduced \\textit{Spitzer} light curve presented by G07 to give yet another estimate of the planet's radius, but did not use the result to make deductions about the planet's structure. The radius of GJ\\,436b given by S08 is smaller than that ascertained by G07 and D07, although the results are consistent within the quoted uncertainties.  D07 and \\citet{demory07} carried out complete analyses of the \\textit{Spitzer} secondary eclipse light curve for GJ\\,436b and were able to measure the planet-to-star flux ratio, and thus deduce the brightness temperature of the planet at the bandpass of the observations. The two groups found nearly identical results, but differed in their interpretation. D07 suggested that the measured temperature was higher than expected from a simple model of the planet and was indicative of tidal heating arising from the planet's non-circular orbit. Comparing the measured brightness temperature to more sophisticated models, \\citet{demory07} found that a model with inefficient redistribution of energy from the day to night side could reproduce the observations.  In this paper we present an analysis of a transit light curve for GJ\\,436b obtained with the Fine Guidance Sensor (FGS) instrument on the \\textit{Hubble Space Telescope} (\\textit{HST}). Following the announcement of transits for GJ\\,436b we were granted Director's Discretionary time to carry out these observations. The motivation was to use the exceptional photometric capabilities of the \\textit{HST} \\citep[e.g.][]{benedict98} to make a precise, independent determination of the radius and orbital ephemeris for the planet. Such measurements are of critical importance because they are the foundation on which the illuminating follow-up studies that are only possible for transiting planets are based. In the case of GJ\\,436b specifically, an independent measurement of the radius is a test of previous assertions that the planet has a gaseous H/He envelope like Neptune and allows robust interpretation of secondary eclipse measurements. Additionally, the measured ephemeris will be a benchmark for future transit timing studies aimed at detecting additional planets in the system. ",
        "conclusions": "\\subsection{The planet's characteristics} Our determined mass and radius for GJ\\,436b is placed in context with the previous results of G07, D07, and S08, the solar system ice giants, and the theoretical mass-radius relationships of \\citet{fortney07} and \\citet{baraffe08} in Fig.~\\ref{f2}. The radius we find differs by $1.82\\sigma$ (93\\% significance), $1.54\\sigma$ (87\\% significance), and $2.39\\sigma$ (98\\% significance) from the values of G07, D07, and S08 respectively. To collapse the previous results into a single value we first average the G07 and S08 results because they are from analyses of the same reduced data. Then we average this value with the D07 value to give $R_{p}$ = 4.22\\,$R_{\\oplus}$. We conservatively adopt a symmetric error in this value of 0.21\\, $R_{\\oplus}$, which is the largest error bar among the three results. The standard deviation of the three values is 0.15\\,$R_{\\oplus}$. The radius we find for GJ\\,436b is larger by $1.74\\sigma$ (92\\% significance) than this representative of the previously determined values.  The \\textit{Spitzer} data have an advantage in that the effect of limb darkening is almost non-existent due to the much longer wavelength of the bandpass used. However, we have carefully considered the effect of limb darkening in our analysis through specific calculation of it in the FGS bandpass using a realistic model atmosphere. Changing the parameters of the model atmosphere we used within reasonable limits or using a different functional form to fit the specific intensities (i.e. quadratic) does not significantly affect our results. Therefore, our determined parameters are relatively robust against uncertainties in the limb darkening.  \\begin{figure} \\resizebox{\\hsize}{!}{\\includegraphics{fig3.eps}} \\caption{Determined mass and radius for GJ\\,436b from G07 (triangle), D07 (square), S08 (circle), and this paper (star). The values for Uranus and Neptune are indicated by the ``U'' and ``N'' characters respectively. The lines are theoretical relationships for ice giants from \\citet{fortney07}, where the solid line is for a planet with a 10\\% by mass H/He envelope and the dotted line for a planet without an envelope. The shaded region delineates the possible range of values from \\citet{baraffe08} models for an ice giant with a 10\\% H/He envelope and age ranging from 1 -- 7\\,Gyr. } \\label{f2} \\end{figure}  Our data have a similar level of precision and time sampling as the \\textit{Spitzer} data. The larger uncertainties in the parameters from our analysis stems from the additional degrees of freedom introduced by fitting for the relative flux normalizations. However, as we have shown above, this does not mean the results are biased and is unlikely to be the reason for the planet radius discrepancy.  The larger radius we determine for the planet cannot construed as a detection of the planet's effective radius dependency on wavelength arising from spectral features \\citep[e.g.][]{charbonneau02, barman07, pont08}, although our bandpass does cover the Na doublet near 600\\,nm \\citep{fortney03, barman07}. One reason for this is that we find a planet to star radii ratio consistent with G07, D07, and S08. If we fix the star's radius to the average of the values given by those authors we obtain a planet radius nearly exactly the same. Furthermore, our data are not sensitive to the likely size of the effect, which is probably at most a few percent in our wide bandpass.  A plausible explanation for the differences between our results and the previous ones is stellar activity. As mentioned above, \\citet{demory07} report that GJ\\,436 exhibits visual flux variations with semiamplitude $\\sim$ 1\\%, and period $\\sim$ 50\\,d. These flux variations are consistent with periodic variations of emission in the Ca II H \\& K lines, which presumably are modulated by the star's rotation.  Breaking the degeneracy between the planet and host star radii, as we have attempted to do, requires high-precision coverage of the transit ingress and egress \\citep{charbonneau07}. It is possible that during the different visits the planet obscured regions with different brightnesses arising from the spots on the star's surface. If this happened during the ingress and/or egress then the morphology of the light curve was altered and our model was inadequate, which could result in incorrect determined parameters. It is also possible that if the planet obscured an unusually brighter or dimmer region of the star during the other parts of the transit then the fitting would be negatively affected due to the need to phase fold the partial transit observations and determine the relative flux normalizations.  Another possible explanation for the discrepancy is that we, G07, D07, and S08 have simply underestimated our uncertainties. In addition to analyzing the \\textit{Spitzer} light curve for GJ\\,436b, S08 presented an interesting study of light curves for other transiting planet systems. One conclusion he reached was that analyses of even the highest-precision data sets for a given system can yield significantly different parameters owing to unidentified systematic errors in the data. In this context it is important to keep in mind that the G07, D07, and S08 results are not independent as they utilize the same data, albeit reduced independently by two of the groups and analyzed separately by all three. If there is a systematic error inherent to the \\textit{Spitzer} data then the three groups would have obtained still similar, but systematically incorrect parameters. Formally, the difference between our planet radius and the representative \\textit{Spitzer} value is significant. However, it could be that this difference is only representative of the true uncertainty in the analyses.  In the absence of any hard evidence that our result is spurious, we suggest that the radius of GJ\\,436b is larger than previously thought. If we assume that a planet's radius dependence on the mass fraction of its envelope is linear in this regime, then extrapolation of the \\citet{fortney07} models to our measured radius gives a H/He envelope mass fraction of 15\\%. The weighted average of our radius value and the representative value of the \\textit{Spitzer} results is 4.48 $\\pm$ 0.17\\,$R_{\\oplus}$, for which the \\citet{fortney07} models suggests an envelope mass fraction of 10\\% $\\pm$ 3\\%. The \\citet{baraffe08} models are fairly consistent with those of \\citet{fortney07}, although they imply a slightly lower envelope mass fraction. \\citet{baraffe08} also emphasize the impact of the evolutionary cooling on the radius of a planet like GJ\\,436b and this adds to the ambiguity of specific interpretations because the age of GJ\\,436 is quite uncertain. More theoretical and observational work is obviously needed to improve our understanding of GJ\\,436b and, by extension, general models of planet formation and evolution.  From an observational standpoint, obtaining and analyzing an additional high-precision transit light curve of GJ\\,436b with a different instrument than already used would still be valuable. Such data would preferably be obtained in the near-infrared to reduce the impact of stellar variability and uncertainties in the limb darkening. Another interesting observational study would be to make a direct radius measurement of the host star with interferometry, similar to what has been done for the transiting planet host star HD 189733 \\citep{baines07}. We have found a similar ratio of the planet and star radii as the previous studies, and so an additional constraint on the stellar radius would give a tighter constraint on the planet radius.  What can be definitively ascertained from comparison of the observational results and theoretical models is that GJ\\,436b has a significant H/He envelope similar to Uranus and Neptune. Such similarities in structure suggest that they formed in a similar environment. Additionally, GJ\\,436b and the solar system ice giants are similar in that they likely could not have formed in their current locations. Uranus and Neptune are thought to be too far away from the Sun \\citep{levison01}, while GJ\\,436b is likely too close to its host star. One hypothesis for the formation and evolution of Uranus and Neptune is that they were originally in the same region of the protoplanetary disk as Jupiter and Saturn \\citep{thommes99}. The later two planets accreted gas much more quickly and scattered the ice giants to-be outward towards their current locations, which limited their growth. Uranus and Neptune are therefore really gas giant planet cores that did not accrete gas quickly enough to complete formation. In contrast, the structure of GJ\\,436b likely does not require a comparably violent past because it orbits a M dwarf. \\citet{laughlin04} have shown that formation of gas giants around low mass stars is severely hindered because these stars are expected to have correspondingly low surface density protoplanetary disks and the dynamical timescale of orbiting bodies is longer. \\citet{laughlin04} predict a dearth of Jovian mass planets due to the resulting slow accretion of gas, but plenty of Neptune and lower mass planets around M dwarfs. In this context GJ\\,436b could be considered a failed gas giant akin to Uranus and Neptune.  Furthermore, GJ\\,436b could have experienced scattering due to another body like Uranus and Neptune even though this probably wasn't required to limit its growth into a larger body. The evidence for this comes from its current location, which is far away from its likely formation site; and its observed orbital eccentricity ($e$ = 0.16), which is in direct contradiction to predictions of tidal circularization theories \\citep{maness07}. A major dynamical interaction event after formation leading to inward scattering is an explanation that unifies both of these properties into a single evolutionary picture. This interaction could have occurred with the outer object in the system that is indicated by the long term trend in GJ\\,436's radial velocities. An alternative explanation for GJ\\,436b's orbital properties is disk interaction leading to migration and eccentricity excitation \\citep[e.g.][]{goldreich80, goldreich03}. Continued observations of the GJ\\,436 system are needed to detect and characterize all of the objects it contains. Complementary theoretical studies based on the system census and the likelihood that GJ\\,436 had a low-mass disk are also needed to specifically assess the plausibility of the conceivable evolutionary scenarios.  \\subsection{Transit timings} The results of our analysis also allow us to search for deviations in transit times that might arise from the gravitational perturbations of another planet in the system. The times determined for each of the five \\textit{HST} partial transit observations, which are given in Table~\\ref{t4}, exhibit a maximum deviation from the mean ephemeris of 10\\,s. The average uncertainty on these times is 145\\,s so the deviations are fully consistent with regular transit times over this observational period.  To search for possible long-term variations we compare our mean transit time with the published transit times from two other epochs. The previous results come from the initial discovery data presented by \\citet{gillon07b}, and the analyses of the \\textit{Spitzer} light curve by G07, D07, and S08 that were discussed above. It should be noted that the time stamps of the \\textit{Spitzer} photometry presented by G07, which was also analyzed by S08, are too late by 33\\,s (M. Gillon private communication). The corrected central transit times for the G07 and S08 results are 2454280.78148 and 2454280.78174\\,HJD respectively. We utilized these corrected times for all the invesigations described below.  The transit times determined by D07 and G07 are quite consistent, but disagree with the S08 determined value by more than expected from the given uncertainties ($\\sim 2\\sigma$). To compare the \\textit{Spitzer} transit times with the other data we collapse the G07, D07, and S08 times into a single value by first averaging the G07 and S08 results, and then averaging this with the D07 result. We adopted the largest given uncertainty from the three analyses (0.00016\\,d from D07) as the error in the collapsed transit times.  A plot of the observed transit time residuals from our ephemeris (transit time and orbital period) for the three available epochs is shown in Fig.~\\ref{f3}. Our determined orbital period is consistent with, but 30\\% more precise than, the value found by \\citet{maness07} from analyzing the radial velocities of GJ\\,436 only. We find a difference of 128\\,s between the representative value from the \\textit{Spitzer} observation and our own. The deviations show a nearly even trend with orbital epoch when also considering the lower precision time given by \\citet{gillon07b}. This could be a result of an error in the period because the deviations are well within the range expected by the uncertainty in this parameter. Alternatively, the transit timing deviations could be due to perturbations of the host star arising from an additional object in the system.  \\begin{figure} \\resizebox{\\hsize}{!}{\\includegraphics{fig2.eps}} \\caption{Observed transit time deviations from our ephemeris. The data are from \\citet[][square]{gillon07b}, an average of the G07, D07, and S08 values (triangle), and this paper (star). For the later two values the error bars are smaller than the size of the points. The dashed lines give the $\\pm 1\\sigma$ range allowed by the uncertainty in the orbital period propagated backwards in time. The dotted line shows the trend expected due to the revised value of the period based on fitting the transit times.} \\label{f3} \\end{figure}  \\citet{maness07} found strong evidence for a long-term slope in the GJ\\,436 radial velocities (1.4\\,m\\,s$^{-1}$\\,yr$^{-1}$), which suggests it has an additional outer companion. The nature of this companion is currently unconstrained, but the observed long-term acceleration of the star means that it has a second degree of motion on top of its regular orbital motion due to the known transiting planet. Therefore, it is reasonable, and even expected, that the observed transit times of GJ\\,436b are varying due to the changing position of the host star. What would be seen in such a case is a curvature in the transit timing offsets rather than the simple linear trend that would be due to an error in the orbital period.  Another possible way to detect perturbations with the transit timings is to look for deviations between the observed transit times and those predicted by orbital parameters found from modeling the radial velocities only. The key in this case is that the radial velocity orbit is relative to the system barycenter, which is invariant, and the transit times are relative to the position of the star, which should be changing. It is for this reason that we chose not to model our transit light curve simultaneously with the radial velocities.  \\citet{ribas08} have recently suggested that GJ\\,436 hosts an additional ``Super-Earth'' planet in a near 2:1 resonance with the transiting planet. Such a planet would likely have a strong influence on the observed eclipse (both transit and secondary) times for planet b due to dynamical interactions. In this context, our finding a long-term trend in the transit timings could be a result of orbital precession caused by such interactions. Ultimately, a detailed dynamical study of the model proposed by \\citet{ribas08} in comparison with the available eclipse times and radial velocities should be undertaken to evaluate their claim. This is beyond the scope of the current paper, but we plan to present such a study in the future.  If we assume that the transit timing deviations are due to an error in the orbital period then we may calculate a new, more precise value by leveraging the multiple orbits that have elapsed. Using our transit time and the average of the values reported by G07, D07, and S08 we find $P = 2.643904 \\pm 0.000005$\\,d, which is an order of magnitude more precise than the value we find by analyzing the \\textit{HST} transit photometry alone. Adopting this period also brings the earlier transit time reported by \\citet{gillon07b} in line with the other two, but this is not very significant due to the relatively low-precision of the value. After submitting this paper we became aware of recent transit times measured by \\citet{alonso08} and amateur astronomers\\footnote{See http://brucegary.net/AXA/GJ436/gj436.htm}. These data support our finding that the transit times for GJ\\,436b are drifting from the \\textit{Spitzer} ephemeris. The transit times we have determined will provide an additional benchmark for future studies aimed at detecting or setting limits on additional objects in the GJ\\,436 system."
    },
    "0806/0806.0256_arXiv.txt": {
        "abstract": "The StarScan machine at the U.S.~Naval Observatory (USNO) completed measuring photographic astrograph plates to allow determination of proper motions for the USNO CCD Astrograph Catalog (UCAC) program. All applicable 1940 AGK2 plates, about 2200 Hamburg Zone Astrograph plates, 900 Black Birch (USNO Twin Astrograph) plates, and 300 Lick Astrograph plates have been measured. StarScan comprises of a CCD camera, telecentric lens, air-bearing granite table, stepper motor screws, and Heidenhain scales to operate in a step-stare mode. The repeatability of StarScan measures is about 0.2 $\\mu$m. The CCD mapping as well as the global table coordinate system has been calibrated using a special dot calibration plate and the overall accuracy of StarScan $x,y$ data is derived to be 0.5 $\\mu$m. Application to real photographic plate data shows that position information of at least 0.65 $\\mu$m accuracy can be extracted from course grain 103a-type emulsion astrometric plates. Transformations between ``direct\" and ``reverse\" measures of fine grain emulsion plate measures are obtained on the 0.3 $\\mu$m level per well exposed stellar image and coordinate, which is at the limit of the StarScan machine. ",
        "introduction": "This paper introduces the current StarScan hardware, and gives details about its operation, calibration, and its scientific program. This effort is part of the IAU Task Force on Preservation and Digitization of Photographic Plates (PDPP) \\footnote{www.lizardhollow.net/PDPP.htm} to measure astronomical plates worldwide. Reductions of these astrograph data is in progress and astrometric results will be presented in an upcoming paper.  The StarScan machine in its original set-up was acquired by the U.S.~Naval Observatory (USNO), Washington in the 1970s. A projection system \\citep{tac1} was used to determine centroids of stellar images in a semi-automatic mode using a list of target stars to be measured. The Twin Astrograph Catalog (TAC) based on over 4000 plates resulted from these measures \\citep{tac1, tac2}. In the early 1990's StarScann operations stopped and only due to its extremely large and heavy granite table the machine was not thrown out, instead was sealed off in a small room in its original building with a new access door to the outside.  Starting in 1998 \\citep{baas} StarScan was completely refurbished switching to the current operation mode with an imaging CCD camera and step-stare scanning of entire plates. The hardware change was performed by L.~Winter (contractor), T.~Rafferty (USNO, now retired) and the USNO instrument shop. The new operations software was developed by E.~Holdenried (USNO, now retired), and the raw data reduction pipeline was adopted from a similar program at the Hamburg Observatory (L.~Winter).  In section 2 of this paper the current science program is described, aiming at early epoch star positions to derive proper motions. Section 3 presents details on the hardware including the 2006 upgrade and section 4 describes the calibration process from set-up of dedicated measures to interpretation of results. Section 5 reveals external errors derived from measures of real astrometric plates, and section 6 summarizes the astrometric performance of StarScan. ",
        "conclusions": "For well exposed stellar images StarScan has a repeatability error of 0.2$\\mu$m per coordinate. Using overlapping areas in the digitization process the mapping parameters of the CCD camera with respect to the $x,y$-table coordinate system are calibrated for each plate and orientation measure to the same precision. The overall accuracy of StarScan measures is limited by the systematic errors of the $x,y$-table and the stage motions. Using a dot calibration plate specifically manufactured for this purpose, 2-dimensional correction data are obtained to calibrate the $x,y$-table coordinate system. The high precision of the StarScan measures allows one to see variations in this calibration pattern which depend at least on time, maybe on other parameters as well. The uncalibrated machine performs at an accuracy level of about 1 $\\mu$m RMS with a parabola-like distortion for the $\\Delta y$ as a function of $y$ with $\\pm$ 1.5 $\\mu$m amplitude being the largest error contribution. After applying the 2-dim calibration corrections StarScan performs on the 0.3 $\\mu$m per coordinate accuracy level.  By comparing different photographic plates of the same field it has been demonstrated that astrometric useful information at an accuracy level of at least 0.65 $\\mu$m can be extracted from the data of course grain emulsions exposed at dedicated astrographs. Comparisons of ``direct\" and ``reverse\" measures of fine grain emulsions show errors of 0.3 $\\mu$m per single star image and measure.  This is only an upper limit of the plate data quality, because of StarScan measuring machine errors close to the same level. A higher accuracy than StarScan currently can provide is recommended for fine grain astrometric plates and such a new machine was just delivered to the Royal Observatory Belgium \\citep{d4d}.  It is recommended to measure the dot calibration plate often to verify and improve the 2-dim $xy$-table corrections. Measuring a plate in 2 orientations, rotated by $180^{\\circ}$, is mandatory, at least for StarScan and when aiming at high accuracy metric results, in order to correct for the non-linear magnitude equation from the measure process and to monitor quality by identifying occasional problems in the measure process."
    },
    "0806/0806.0609.txt": {
        "abstract": "We present an analysis of wind-blown, parsec-sized, mid-infrared bubbles and associated star-formation using GLIMPSE/IRAC, MIPSGAL/MIPS and MAGPIS/VLA surveys.  Three bubbles from the Churchwell et al. (2006) catalog were selected. The relative distribution of the ionized gas (based on 20 cm emission), PAH emission (based on 8 $\\mu$m, 5.8 $\\mu$m and lack of 4.5 $\\mu$m emission) and hot dust (24 $\\mu$m emission) are compared. At the center of each bubble there is a region containing ionized gas and hot dust, surrounded by PAHs. We identify the likely source(s) of the stellar wind and ionizing flux producing each bubble based upon SED fitting to numerical hot stellar photosphere models.  Candidate YSOs are also identified using SED fitting, including several sites of possible triggered star formation. ",
        "introduction": "Mid-infrared (MIR) imaging surveys of the Galactic plane such as the Mid-Course Space Experiment (MSX; Price et al.  2001) and the Galactic Legacy Infrared Midplane Survey Extraordinaire (GLIMPSE; Churchwell et al. 2001; Benjamin et al. 2003) have revealed a large number of bubbles in the disk of the Galaxy.  The bubbles have bright 8 $\\mu$m shells that enclose bright 24 $\\mu$m interiors.  ISO (Infrared Space Observatory; Kessler et al.  1996) and MSX were the first to reveal the existence of these objects (including the association with triggered star-formation, Deharveng et al.  2005), but the GLIMPSE survey (with ten times better spatial resolution and a hundred times the sensitivity than MSX) has enabled the detection of many more bubbles and has enabled the refinement of their morphological structure.  In fact, the GLIMPSE survey has shown that bubbles are a hallmark of the diffuse emission in the Galactic plane.  Small and faint bubbles are apparent only in the GLIMPSE survey due to the lower resolution and sensitivity of MSX.  Churchwell et al.  (2006; 2007) have cataloged almost 600 bubbles. An average of $\\gtrsim$5 bubbles per square degree were found within 20$^\\circ$ of the Galactic center.  These bubbles are identified as complete or partial rings in the GLIMPSE images, which Churchwell et al.  (2006) argue are two dimensional representations of three dimensional bubbles.  Churchwell et al. (2006) showed that the bubbles are distributed in longitude and latitude like O and B stars, $>$25\\% of which are coincident with known radio HII regions, and $\\sim$13\\% enclose known stellar clusters.  About 1/3 of the bubbles are produced by O stars.  Approximately two thirds of the sample have small angular diameters, which Churchwell et al.  (2006) propose to be mostly physically small bubbles produced by late B type stars whose UV photon fluxes are adequate to excite a shell of PAH emission but not adequate to produce a detectable radio HII region.  Churchwell et al.  (2006; 2007) tabulated bubble morphological properties such as angular diameter, shell thickness, and eccentricity.  %Here, we will examine in some detail the diffuse emission properties of dust %in the IRAC and MIPS bands toward three bubbles and compare their IR and radio %continuum emission distributions.  We show that apparent bubble morphology is %a strong function of wavelength and examine the implications of this for the %excitation and photo-destruction of polycyclic aromatic hydrocarbons (PAHs) %and thermal dust emission.  Bubbles/HII regions identified by MIR emission present a new and powerful tool to study the interaction of young hot stars with their environment.  Young stars impact the ambient interstellar medium (ISM) by heating the dust, ionizing the gas (in the case of O and early B stars), and heating an expanding bubble.  Bubbles can be produced by stellar winds and over-pressure by ionization and heating by stellar UV radiation.  Bubble sizes and morphologies are determined by the density structure of the ambient ISM, the luminosity of the central star(s) responsible for producing the bubble, the stellar wind luminosity and the motion of the central star(s) relative to the ambient ISM.  The morphologies of MIR dust bubbles/HII regions reveal important information about the strength and directionalities of stellar winds and the structure and physical properties of the ambient ISM into which the bubbles are expanding.  Pillars, scalloping, and sharpness of the inner boundaries of bubbles defined by 8$\\mu$m emission provide unique insights into the hydrodynamics, photo-ionization, and evaporation of gas and sublimation of dust in expanding bubbles, and stellar mass loss rates during their evolution. For example, based on GLIMPSE-MIR morphology of ultra-compact (UC) HII regions, Hoare et al.  (2007) and references therein have found an unusually high fraction of cometary shapes. They argue that the morphology of the PAH-traced photo-dissociated region (PDR), combined with the radio free-free emission, is suggestive of a champagne flow.  MIR bubbles that surround HII regions ionized by O stars and B supergiants (i.e. those that have strong stellar winds) have been modeled by several groups.  Analytic evolution models have been developed by Castro, McCray, and Weaver (1975), Weaver et al. (1977), Whitworth et al. (1994), and Capriotti and Kozminsky (2001). These models predict that wind-blown bubbles around O stars and B supergiants should have the following distinct regions: a hypersonic stellar-wind-evacuated central cavity surrounding the central star where densities are very low; a hot, low-density, shocked stellar wind region surrounding the central evacuated cavity that occupies most of the volume of the HII region; a thin conduction zone where temperature sharply decreases and density rapidly increases (with radius); a thin, dense shell of shocked H$^+$ gas surrounded by a very thin shell of non-shocked H$^+$ gas at about 10$^4$ K. The outer thin shell of 10$^4$ K H$^+$ gas is the region classically thought to represent the photo-ionized HII region. This basic picture is supported by detailed numerical evolution models calculated for a 35 M$_\\odot$ (Freyer, Hensler, \\& Yorke, 2006; hereafter FHY06), a 60 M$_\\odot$ (Freyer, Henseler, and Yorke, 2003), and an 85 M$_\\odot$ star (Kr\\\"oger, Hensler, \\& Freyer, 2006, and Kr\\\"oger, Freyer, Hensler, \\& Yorke, 2007).  The numerical models, however, show that the sharp boundaries predicted by the analytic models probably are not sharp due to turbulent mixing.  The consequence of this is that clumps of cool 10$^4$ K gas are predicted to be found mixed with the hot 10$^7$ K gas.  The thickness of the 10$^4$ K shell is predicted to be a function both of age and luminosity of the central star; the lower the wind luminosity the thicker the cool ionized shell is.  Also the outer boundaries of the HII regions are predicted to be quite jagged due to instabilities produced when a dense medium expands into a much lower density interstellar medium (ISM).  We will compare some of the model predictions with our observations in \\S 5. All theoretical evolution models up to now have omitted the possible role of dust in HII regions, which new MIR observations clearly demonstrate is present.  Recent observations are providing a deeper understanding of the relative distributions of thermal dust, PAHs, ionized gas, and PDRs.  Peeters et al. (2005) review the ISO spectroscopic observations of compact and evolved HII regions and the PDRs surrounding them. They find gas temperatures in the PDRs between 100 K (NGC 2024) and 200 K - 500 K (S106IR) and densities between 10$^4$ cm$^{-3}$ (W75N) and 10$^6$ cm$^{-3}$ (NGC 2024).  They find broad agreement between these results and models of the illumination of PDRs by the appropriate nearby hot star(s). Kassis et al. (2006), on the other hand, report ground-based observations of the Orion bar that indicate PAHs are present just outside the HII region. They show that [Ne II] 12.81 $\\mu$m emission traces the geometry of the PDR in the bar. They also conclude that there is not a sharp transition between neutral and ionized PAH emission within the PDR.  Each of the four Spitzer/IRAC (Infrared Array Camera) bandpasses are dominated by different emission processes in the neighborhood of a hot star (see Draine 2003 and Peeters et al. 2003).  The brightest objects in the 3.6$\\mu$m band are stars, but this band also has contributions from a weak, diffuse PAH feature at 3.3$\\mu$m, and possibly from scattered star light.  The 4.5$\\mu$m band samples diffuse emission in the Br$\\alpha$ and Pf$\\beta$ lines (from HII regions), H$_2$ v=0-0, S(9), S(10), and S(11) and CO v=1-0 ro-vib lines at 4.665$\\mu$m (from shocked molecular gas).  The 4.5$\\mu$m band contains no PAH features and the brightest sources in this band are also stars.  The 5.8$\\mu$m band contains a strong PAH feature at 6.2$\\mu$m which can dominate the diffuse emission except very close to O stars where PAHs are destroyed (see \\S 4.1). Near hot O stars, strong contributions from thermally emitting dust plus a small contribution from stochastically heated small grains probably dominate the diffuse emission in the 5.8$\\mu$m band.  The 8.0$\\mu$m band contains two very strong PAH features at 7.7 and 8.6$\\mu$m which dominate the diffuse emission in this band, although near hot stars it may be dominated by thermal dust emission with little or no PAH emission.  Within ionized zones, Hoare (1990) and Hoare et al. (1991) have shown that trapped Lyman $\\alpha$ heated dust can maintain T $>$ 100 K throughout the ionized region.  Because the emission process that dominates in each IRAC band depends on the environment, band ratios provide a powerful tool to measure the extent of each environment.  The main focus of this study is to use wavelength-dependent distributions of MIR diffuse dust emission and high resolution radio continuum emission to determine the PAH destruction radius of three bubbles and to trace the relative sizes and locations of 1) the HII region, 2) the hot thermally emitting dust, and 3) the location and extent of the PDR region associated with each bubble.  In section 2 we introduce the different data sets used for this analysis. In section 3, we discuss results of the observations toward three bubbles. In section 4, we estimate the PAH destruction radius and dust temperature. In section 5, we examine the stellar population within each of the bubbles to identify the ionizing star(s) and any young stellar objects (YSOs) associated with the bubbles. Conclusions are summarized in section 6. ",
        "conclusions": "Based on MIR and 20 cm observations of three bubbles, we conclude the following:\\\\ $\\bullet$ At the center of each bubble there is a region containing ionized gas and hot dust delineated by radio free-free emission and 24 $\\mu$m emission, respectively.\\\\ $\\bullet$ At the center of N10 and N21 (but not N49), the hot dust is also traced by 5.8 and 8.0 $\\mu$m emission.\\\\ %$\\bullet$ The dust is, at most, 150 K.\\\\ $\\bullet$ Based on the decreased ratio of 8.0 $\\mu$m/4.5 $\\mu$m, 5.8 $\\mu$m/4.5 $\\mu$m and the lack of a similar decrease in 8.0 $\\mu$m/3.6 $\\mu$m, we conclude that inside the 8 $\\mu$m shell PAHs are destroyed by hard, direct stellar UV radiation.\\\\ $\\bullet$ Based on the increased 8.0 $\\mu$m/4.5 $\\mu$m, 5.8 $\\mu$m/4.5 $\\mu$m, the bright 8.0 $\\mu$m ring emission is dominated by PAH emission that defines the PDR region around each bubble.\\\\ $\\bullet$ We have identified YSO candidates and probable ionizing sources for each bubble. This was accomplished by employing the SED model fitter developed by Robitaille et al. (2007) to fit model SEDs of YSOs and hot stellar photospheres (Kurucz 1993) to the observed fluxes.\\\\ $\\bullet$ Based on morphology and environment, several of the identified YSO candidates in N10 and N49 appear to be triggered by expansion of the bubbles.\\\\ $\\bullet$ The wind-blown cavity at the center of N49 appears to be produced by a central 05 V star, the hottest ionizing star observed here. This bubble appears to be dominated by the wind from the O5 star and have a dynamical age of $\\ge$10$^5$ yrs.\\\\ %  $\\bullet$ When the ambient density is scaled to 10$^{5}$ %cm$^{-3}$, the numerical models of FHY06 are consistent with the measured %radii of N49 and with the age (10$^5$ yrs) necessary for YSOs to be %triggered by the bubble expansion.\\\\"
    },
    "0806/0806.1120_arXiv.txt": {
        "abstract": "{}{Relations connecting gamma ray burst quantities can be used to constrain cosmographic parameters of the Hubble law  at medium-high redshifts.}{We consider a sample of 27 gamma ray bursts to construct the luminosity distance to  redshift relation and derive the values of the parameters $q_0$, $j_0$, and $s_0$. The analysis is compared with other methods in the literature. }{ gamma gay bursts, if calibrated by SNeIa, seem reliable as distance indicators and give cosmographic parameters in agreement with the $\\Lambda$CDM model.}{} ",
        "introduction": "A class of very accurate standard candles, the supernovae Ia (SNeIa), has been highly developed in the last two decades \\citep{Branch}; however, these objects are hardly detectable at redshifts higher than $z$ = 1.7, so the study of more distant regions of the Universe leads to the needing to implement more powerful standard candles. The problem becomes particularly crucial at intermediate redshift, $z$ = 6 - 7, where, up to now, not very well - defined distance indicators are available.  In the last years, several efforts have been made in order to implement gamma ray bursts (GRBs), the most powerful explosions in the Universe, as  standard candles, and several interesting results have  recently been achieved  (e.g. Amati et al 2008; Basilakos $\\&$ Perivolaropoulos 2008 and references therein). Considering the standard model of such objects, the GRB phenomenon should originate from the black hole formation and reach  huge amounts of energy (up to $10^{54} erg$). These events are observed at considerable distances, so there are several efforts to frame them into the standard of cosmological distance ladder.  In the literature, several more - detailed models  give account for the GRB formation, e.g. Meszaros 2006; Ruffini et al 2008, but, up to now, none of them is intrinsically capable of connecting all the observable quantities. For this reason, GRBs cannot be used as standard candles. Despite  this shortcoming, there are several observational correlations among the photometric and  spectral properties of GRBs.  These features allow  use of GRBs as distance indicators \\citep{Schaefer}, even when they cannot be fully ''enrolled\" in the class of standard candles. In particular, it is possible to connect the peak energy of GRBs, $E_p$, with the isotropic energy released in the burst, $E_{iso}$, and with the rest frame jet break - time of the afterglow optical light curve, measured in days, $t_b$, \\citep{Liang2}:  \\begin{equation}\\label{eq:no1} \\log{E_{iso}}=a + b_1 \\log{\\frac{E_p (1+z)}{300keV}} + b_2 \\log{\\frac{t_b}{(1+z)1day}} \\end{equation} where $a$ and  $b_i$, with $i=1,2$, are calibration constants.  Another interesting result is the relation given by Ghirlanda et al. \\citep{Ghirlanda}. It connects the peak energy $E_p$ with the collimation-corrected energy, or the energy release of a GRB jet, $E_{\\gamma}$, where  \\begin{equation} E_{\\gamma} = ( 1 - \\cos{\\theta_{jet}} ) E_{iso}, \\end{equation} with $\\theta_{jet}$ the jet opening angle, given by \\citep{Sari}:  \\begin{equation} \\theta_{jet} = 0.163\\left(\\frac{t_b}{1 + z}\\right)^{3/8}\\left(\\frac{n_0\\eta_{\\gamma}}{E_{iso,52}}\\right)^{1/8}, \\end{equation} where $E_{iso,52} = E_{iso}/10^{52}$ ergs,  $n_0$ is the circumburst particle density in 1 cm$^{-3}$, and $\\eta_{\\gamma}$ the radiative efficiency. The Ghirlanda et al. relation is  \\begin{equation}\\label{eq:no2} \\log{E_{\\gamma}} = a + b \\log{\\frac{E_p}{300 keV}}, \\end{equation} where $a$ and $b$ are two calibration constants.  These two relations are used the most   in constraining cosmology due to their relatively small scatter,  interestingly very tight in the Ghirlanda et al. one, and the sufficient number of data points available.  In \\citep{Schaefer},  an example of the  discrepancy between data and theoretical curves is shown for these two relations. It is worth noticing that the calibration of the above relations is necessary to avoid the circularity problem: all the relations need to be calibrated for every set of cosmological parameters. Indeed, all GRB distances, obtained only in a photometric way, are strictly dependent on the cosmological parameters since there is no low-redshift  set of GRBs to achieve a cosmology-independent calibration.  Recently, Liang et al. \\citep{Liang}  present a calibration method (Liang thereafter) for several GRB relations, included the above relations (\\ref{eq:no1}) and (\\ref{eq:no2}), in a cosmology-independent way using the SNeIa. In fact, the SNeIa are very accurate standard candles, but their range is limited up to $z \\approx$ 1.7; hence, assuming that relations (\\ref{eq:no1}) and (\\ref{eq:no2}) work at any $z$ and that, at the same redshift, GRBs and SNeIa  have the same luminosity distance, it becomes possible, in principle, to calibrate GRB relations at low redshifts. The calibration parameters are shown in Table \\ref{table:no1}. \\begin{table} \\caption{Parameter values obtained by \\citep{Liang}} % \\label{table:no1} % \\centering % \\begin{tabular}{c c c} % \\hline\\hline % Relation & a & b \\\\ % \\hline % $E_{\\gamma} - E_p$ & 52.26 $\\pm$ 0.09 & 1.69 $\\pm$ 0.11  \\\\ % $E_{iso}-E_p - t_b$ & 52.83 $\\pm$ 0.10 & 2.28 $\\pm$ 0.30 \\\\ &  & -1.07 $\\pm$ 0.21  \\\\ \\hline % \\end{tabular} \\end{table} For the  $E_{iso}-E_p-t_b$ relation, the $b$-values in the first line is $b_1$ and in the second line is $b_2$.  When our working-relations are calibrated  with the Liang method, we can compute the luminosity distance $d_l$ from the well-known relation between $d_l$ and the energy-flux ratio of the distance indicators in consideration. Afterwards, we can use a formulation given by Visser \\citep{Visser1}, where the luminosity distance $d_l$ is related to the cosmographic parameters \\citep{Weinberg} by means of a Taylor series expansion for the same $d_l$. Such an analysis works very well at low and intermediate redshifts, since very good classes of standard candles are available there. Besides, it is useful to constrain alternative theories of gravity, as shown in Capozziello et al. 2008. Since we are calibrating GRBs by SNeIa (in the SNeIa redshift range, the $d_l$ Taylor series analysis works very well), the method could also be extended  to the next step (intermediate-high redshifts) where GRBs are expected to be suitable distance indicators. This working hypothesis could be useful  in order to link low and high redshift ranges and then fully probe $d_l$. However, it is clear that such a Taylor expansion, derived for low redshifts, can be problematic for fitting GRBs at high redshifts. Here, we consider it  a viable methodological approach to link GRBs to SNeIa.   The aim of this work is  to achieve the cosmographic parameters \\citep{Weinberg} using the above GRB relations and then to test the cosmological density parameters in a $\\Lambda$CDM model. The only assumption that we make here is that the Universe is described by a Friedmann-Robertson-Walker geometry and the scale factor of the universe $a(t)$ can be expanded in a Taylor series (Sect.2). In Sect.3, after considering a sample of 27 GRBs, we use a best-fit analysis to derive  the cosmographic parameters discussed in the previous section, adopting the so - called Chevallier, Polarsky, Linder parameterization for the equation of state (EoS). Discussion and conclusions are given in Sect.4. ",
        "conclusions": "Starting from  some relations connecting the observable quantities of GRBs, we have used a sample of 27 GRBs  to derive the luminosity distance - redshift diagram of the Hubble law. The relations  conveniently calibrated by SNeIa  to make them independent of any cosmological models.  We have taken the  Hubble law into account in the Taylor series form, assuming the luminosity distance $d_l$ as a redshift function whose coefficients are combinations of the cosmographic parameters $H_0$, $q_0$, $j_0$, and $s_0$. The aim was to evaluate such parameters starting from the GRB data. A direct analysis of the fits leads to the conclusion that, in the error range, the SNeIa results can also be extended  at higher redshifts \\citep{Visser3}. Besides, such results agree with the $\\Lambda$CDM model according to Eqs. (\\ref{eq:no3}),(\\ref{eq:no4}),(\\ref{eq:no5}). In particular, the value of the parameter  $q_0$ that we found is in agreement with the observed $\\Omega_M$ (see Table \\ref{table:no4}).  \\begin{figure} \\includegraphics[width=10.4cm, height=8cm]{adjdlz.eps} \\caption{Luminosity distance - redshift diagram and the residuals of the $d_l(z)$ GGL fit. Note the discrepancy at high-$z$. The dotted lines are  the bounds predicted at $68\\%$ confidence level.} \\label{fig:no3} \\end{figure}  \\begin{figure} \\includegraphics[width=10.4cm, height=8cm]{adjlndlz.eps} \\caption{Logarithmic version of the luminosity distance relation versus redshift and the residuals of the $ln[d_l/zMpc]$ LZ fit. In this version of the Hubble series, the discrepancy is the same at every $z$. The dotted lines are  the bounds predicted at $95\\%$  confidence level. } \\label{fig:no4} \\end{figure}   \\begin{table} \\caption{Cosmological density parameters} % \\label{table:no4} % \\centering % \\begin{tabular}{c c c} % \\hline\\hline % Fit & $\\Omega_M$ & $\\Omega_{\\Lambda}$ \\\\ % \\hline % $d_l(z)$ LZ & $0.04 \\pm 0.03$ & $0.65 \\pm 0.73$  \\\\ % $d_l(z)$ GGL & $0.46 \\pm 0.43$ & $0.54 \\pm 2.82$  \\\\ $\\ln [d_l/z Mpc]$ LZ & $0.37 \\pm 0.31$ & $0.63 \\pm 1.13$ \\\\ $\\ln [d_l/z Mpc]$ GGL & $0.28 \\pm 0.30$ & $0.72 \\pm 1.09$ \\\\ \\hline % \\end{tabular} \\end{table}   However, the sample  we used is quite poor at high redshifts and, in some sense, this justifies the use of the method of Taylor series which works very well at low redshifts.  In particular, at $z > 6$, we only have  one GRB, GRB050904 (see Fig. \\ref{fig:no3}). This GRB is very important in the fit results because it affects the trend of the fits. For this reason we need some richer sample at medium-high redshifts  to constrain  the results better. However, if we had richer samples at high redshifts, the Taylor series analysis would fail to constrain cosmological models since an exact, and not approximated, $d_l(z)$ expression is needed in that case. The best constraint, however, would be an absolute relation between several, GRB observables which would make the GRBs a powerful standard candle at intermediate-high redshift.    Considering these preliminary results, it seems that cosmography by GRBs could be a useful tool in constraining self-consistent cosmological models even if, up to now,  GRBs are not standard candles in the proper sense.  \\paragraph{Acknowledgements.} We thank the referee for the useful suggestions that improved the paper."
    },
    "0806/0806.4019_arXiv.txt": {
        "abstract": "We present maps of the cosmic large-scale structure around the twelve most distant galaxy clusters from the Massive Cluster Survey (MACS) as traced by the projected surface density of galaxies on the cluster red sequence. Taken with the Suprime-Cam wide-field camera on the Subaru telescope, the images used in this study cover a $27\\times 27$ arcmin$^2$ area around each cluster, corresponding to $10 \\times 10$ Mpc$^2$ at the median redshift of $z=0.55$ of our sample. We directly detect satellite clusters and filaments extending over the full size of our imaging data in the majority of the clusters studied, supporting the picture of mass accretion via infall along filaments suggested by numerical simulations of the growth of clusters and the evolution of large-scale structure. A comparison of the galaxy distribution near the cluster cores with the X-ray surface brightness as observed with Chandra reveals, in several cases, significant offsets between the gas and galaxy distribution, indicative of ongoing merger events. The respective systems are ideally suited for studies of the dynamical properties of gas, galaxies, and dark matter. In addition, the large-scale filaments viewed at high contrast in these MACS clusters are prime targets for the direct detection and study of the warm-hot intergalactic medium (WHIM). ",
        "introduction": "As the largest gravitationally bound constituents of matter, galaxy clusters provide important clues about the distribution and fundamental properties of galaxies, gas, and dark matter. Studying their characteristics over a range of redshifts can yield crucial insights into the overall structure and evolution of the cosmos. Theoretical considerations and the results of numerical simulations (e.g.\\ Colberg et al.\\ 2000; Evrard et al.\\ 2002) predict that clusters grow through accretion and merger events and reside within filamentary structures extending up to $10/h$ Mpc in size. While growth through cluster mergers is a relatively well studied phenomenon, the existence and relevance of filaments for cluster and galaxy evolution remains much less well understood, primarily because of observational challenges created by the very large scales and low densities involved.  Numerical simulations (e.g.\\ Cen \\& Ostriker 1999, Dav\\'{e} et al.\\ 1999) suggest that a large fraction of the baryons in the local universe takes the form of highly diffuse gas at a temperature of $10^{5} - 10^{7} K$ -- the so-called ``Warm-Hot Intergalactic Medium\" (WHIM). The temperature and diffuse nature of the WHIM makes it very difficult to detect directly, the currently most promising techniques being very deep observations at extreme ultraviolet (EUV) and soft X-ray wavelengths, to detect high-ionization metal absorption lines and diffuse bremsstrahlung emission, respectively (e.g. Soltan, Freyberg \\& Hasinger 2005, Nicastro et al.\\ 2005). Since simulations also predict that the WHIM is distributed along large-scale filaments around galaxy clusters (Dav\\'{e} et al.\\ 2001) and may be detectable in observations of the Sunyaev-Zel'dovich Effect (Hallman et al.\\ 2007), promising targets for the direct detection and subsequent study of the WHIM could be unveiled by a systematic survey of the large-scale structure surrounding massive galaxy clusters.  Large optical surveys of the nearby universe such as 2dF and SDSS have allowed detailed investigations of the galaxies within the vast cosmic web consisting of clusters, superclusters, voids, walls, and filamentary structures (e.g.\\ Peacock et al.\\ 2001; Zandivarez, Merch\\'{a}n \\& Padilla 2003; Erdo\\^{g}du et al.\\ 2004; Einasto et al.\\ 2007). Because of the limited depth of the underlying imaging and spectroscopic data obtained by these surveys, studies of low-density environments have focused on structures at relatively low redshift ($z \\sim 0.1-0.2$). A search for filamentary structure around the vertices of the ``cosmic web'' would, however, be particularly rewarding at higher redshifts where the contrast is greatly enhanced by the smaller angular size of such features. In addition, observations at higher redshift would establish an important baseline for evolutionary studies of this environment. Galaxy clusters at higher redshifts are expected to be less evolved and reside within the densest regions of the precursors of the structure seen locally. A large fraction of the matter, dark and luminous, in these structures will eventually become part of virialized cluster cores through merging and infall along the filaments. Hence a multitude of astrophysical topics can be addressed by studying the large-scale surroundings of distant clusters, including the impact of environment on star formation in galaxies, the relative importance of gradual infall along filaments vs. merging for both galaxy and cluster evolution, and the existence and properties of the WHIM.  The study described in the following looks for evidence of filamentary structure in a complete sample of twelve moderately distant, very massive galaxy clusters in order to establish observationally the frequency and prevalance of such structure around the densest nodes of the cosmic web. By targeting galaxy clusters at intermediate redshift ($z \\sim 0.5$), we observe filaments on angular scales of typically 10 arcmin, which makes them accessible to efficient observational sampling in wavebands from the optical to X-rays. The high mass, intermediate redshift, and compact angular scale of these systems make them particularly promising targets for the detection of WHIM emission. Previous studies of the large-scale environment of clusters at similar redshift (e.g. Kodama et al.\\ 2001) have either focused on a single object or on systems of relatively low mass. Since more massive clusters are connected to more, and denser, filaments than less massive systems (Colberg, Krughoff, \\& Connolly 2005, Colberg 2007), a systematic survey of the environment of a statistically complete sample of the most massive clusters at these redshifts is advantageous in the context of the science goals outlined above.  The structure of this paper is as follows. In \\S2 we describe the selection of our cluster sample. \\S3 describes the observations and data reduction as well as the methods used in our analysis of the data. We present and discuss our results in \\S4. \\S5 contains our conclusions and identifies areas for future work. Throughout the paper we assume a cosmology with $ \\Omega_M = 0.3, \\Omega_\\Lambda = 0.7$, and $h_0 = 0.7$. ",
        "conclusions": "Figure 2 shows, for each of the twelve clusters in our sample, contours of the adaptively smoothed, projected surface density of the red-sequence selected galaxies as obtained with {\\sc Asmooth}. By design, all features within the bold outline are significant at the $3\\sigma$ confidence level. Figure 3 shows the same data in a scaled representation, but only for the cluster cores and with X-ray surface brightness countours overlaid.  \\subsection{Optical Morphology}  All but three of the twelve clusters show evidence for at least one satellite cluster, the exceptions being MACS\\,J0911.2$+$1746 (which, however, consists of two close subclusters), MACS\\,J1423.8+2404, and MACS\\,J2129.4$-$0741. The prominent satellite cluster almost 10 arcmin southwest of the core of MACS\\,J0018+1626 (A.K.A. Cl0016+16, Spinrad 1980) has been previously identified (Connolly et al.\\ 1996) but the filament connecting the two systems is reported here for the first time. More than half of all systems feature more than one apparent satellite (MACS\\,J0018+1626, MACS\\,J0025.4--1222, MACS\\,J0257.1--2335, MACS\\,J0454.1$-$0300, MACS\\,J0647.7+7015, MACS\\,J0717.5+3745, MACS\\,J2214.9$-$1359) and an equal number show well-defined filamentary structure (MACS\\,J0018+1626, MACS\\,J0257.1--2335, MACS\\,J0454.1$-$0300, MACS\\,J0647.7+7015, MACS\\,J0717.5+3745, MACS\\,J2129.4$-$0741, MACS\\,J2214.9$-$1359). The filament extending to the southeast of the main cluster in MACS\\,J0717.5+4735 was previously identified by Ebeling et al.\\ (2004) who point out that its orientation matches that of the merger axis of the cluster proper, suggesting that the system has grown for many Gyr via infall of matter along this filament. We find similar alignments between filaments and the galaxy-density contours of the cluster core for almost all other systems exhibiting filamentary structure (seven out of twelve as listed above), providing strong observational evidence that the growth of most clusters occurs along a single preferred direction, be it through major mergers and/or steady accretion, and that these ``assembly lines'' are stable over remarkably long time scales. Even one cluster without prominent large-scale filaments (MACS\\,J0025.4$-$1222) exhibits an elongation of the outermost contours of the galaxy density near the main cluster that appears to point toward a satellite cluster, suggesting that an active filament may have been present in its direction at earlier times.   The most complex systems in terms of large-scale structure in this sample are MACS\\,J0257.1--2335, MACS\\,J0647.7+7015, and MACS\\,J2214.9--1359. Each of these three features at least three collapsed cluster cores separated by several Mpc, as well as two nearly orthogonal axes of merging and infall along filaments. The large-scale morphology of the galaxy distribution in these fields thus strongly suggests that the respective MACS clusters occupy intersections of multiple filaments in the cosmic web. The detection of several filaments connecting satellite clusters to the main cluster provides strong observational support for the results of numerical simulations (Colberg et al.\\ 2005; Colberg 2007) which suggest (a) that cluster pairs with very close separations ($<5$h$^{-1}$ Mpc) are always connected by filaments, and (b) that, statistically, the number of filaments increases with cluster mass.  A closer look at the prevalence of substructure in the cluster cores -- indicative of an ongoing merger -- and the strength of the filamentary extensions connecting these cores to their large-scale surroundings provides a fascinating snapshot of the wide variety in the dynamics and geometry of cluster growth. MACS\\,J2214.9--1359, for instance, exhibits only subtle signs of disturbance in its core region (primarily a 90-degree rotation in the mild ellipticity of the galaxy-density isophotes from the very core to the outer regions) suggesting a state of advanced relaxation since the last major merger. On larger scales, however, faint filaments still trace the paths of past and future infall of matter. By stark contrast, the disturbed core of MACS\\,J0647.7+7015 is clearly the scene of ongoing merging, and the strength of the filemantary connection to a series of satellite clusters to the southeast of the main cluster suggest that this phase of violent growth is to continue for many Gyr to come. Viewed as a whole, our twelve snapshots of the large-scale structure around the most massive galaxy clusters at $z\\sim 0.55$ thus corroborate work by Colberg et al.\\ (1999) whose numerical simulations show strong correlations between infall patterns and the large-scale structure surrounding massive clusters, with preferred directions for the assembly of clusters being established by tidal forces. Our observational results strongly suggest that, although multiple filaments are often present, one of them usually dominates in funneling matter onto the cluster occupying the respective vertex of the cosmic web.  Finally, a minority of our cluster targets exhibit noticeably elongated galaxy-density contours in their core regions (MACS\\,J0744.8+3927, MACS\\,J1149.5+2223, MACS\\,J1423.8+2404), but little sign of substantial filaments on larger scales. Since the orientation of major filaments is random with respect to our line of sight it is possible that significant large-scale structure is present but not visible in projection. Indeed extensive spectroscopic follow-up of the galaxy population of these systems finds strong evidence for line-of-sight substructure in one of these three, MACS\\,J1149.5+2223 (Smith et al.\\ 2008, in preparation). At least one cluster in our sample can, however, be considered to be near virialization and safe from major mergers for some time, namely MACS\\,J1423.8+2404 (see also the following section).  \\subsection{X-ray Emission}  Although the small field of view, compared to SuprimeCam, of ACIS-I allows us to study only the X-ray emission from the main clusters shown in Fig. 2, fresh insights about the dynamical state of these clusters can be obtained by comparing the X-ray surface-brightness contours to the optical morphology of the main cluster (Fig. 3). In all cases we find the orientation of any elongation in the X-ray contours around the cluster core to be well aligned with the direction of infall as traced by the large-scale galaxy distribution. This is expected since the intra-cluster gas responds to the same gravitational forces as the galaxies and should thus, globally, show similar bulk motions.  There is a fundamental difference between gas and galaxies, however, in that the former is a strongly interacting viscous medium, whereas the latter are, on cluster scales, weakly or non-interacting point-like objects. As a consequence, gas and galaxies can show noticeable segregations in non-equilibrium situations like major cluster mergers. Indeed we find two systems for which the centroids of the gas and the galaxy distributions in the cluster cores are noticeably displaced from each other, namely MACS\\,J0025.4--1222 and MACS\\,J0717.5+3745, with the displacement corresponding to roughly 200--300 kpc in the plane of the sky. In either case the X-ray emission peaks between two well separated galaxy concentrations. A third, less obvious, example of the same phenomenon is MACS\\,J0911.2+1746, a cluster consisting of two components of very different mass about 1 Mpc apart. A comparison of the galaxy and gas distributions around the less massive component shows the centroid of the X-ray emission to be significantly displaced from the peak in the galaxy surface density. The cause is, most likely, the same for all three systems mentioned, namely an ongoing merger in which the galaxy populations of the two merging clusters have already passed through each other with little interaction, whereas the collision of the dense gas in the cluster cores creates shocks through which kinetic energy from the bulk motion of the gas is partly converted into thermal energy and partly redirected into transverse velocity components.  Similar cases of a clear segregation between gas and galaxies have been observed in other clusters and have helped greatly to improve our understanding of cluster mergers. A particularly spectacular example is the ``bullet cluster'' 1E0657--558 (Markevitch et al.\\ 2002) in which the fortuitous alignment of the merger axis with the plane of the sky enabled remarkably detailed and quantitative studies of the dynamics of the merging process. It remains, so far, the only system to allow direct constraints to be placed on the dark matter self-interaction cross section (Markevitch et al.\\ 2004, Clowe et al.\\ 2006). The potential implications of significant offsets between the centroids of the gas, galaxy, and dark-matter distribution in clusters are thus far-reaching, and further follow-up studies, in particular a gravitational-lensing analysis of the total-mass distribution, of MACS\\,J0025.4--1222, MACS\\,J0717.5+3745, and MACS\\,J0911.2+1746 may offer a rare opportunity to test the results obtained for 1E0657--558 on the nature and fundamental properties of dark matter.  Two of our 12 MACS clusters (MACS\\,J0744.8+3927 and MACS\\,J1423.8+2404) deserve special mention as the only ones showing obvious cooling cores in our X-ray images. MACS\\,J1423.8+2404 is also the sole cluster that shows neither obvious large-scale structure nor significant substructure in the core region. Although, again, structure along the line of sight cannot be ruled out conclusively, the system's optical and X-ray appearance suggests that it is close to virialization (see also Ebeling et al.\\ 2007). While a cooling core is also clearly present in MACS\\,J0744.8+3927, the large-scale environment of this system and thus its overall assembly history and future are, at present, difficult to assess owing to the limited depth of the existing SuprimeCam data for this most distant cluster in our sample."
    },
    "0806/0806.1646_arXiv.txt": {
        "abstract": "{The formation of planetesimals in protoplanetary disks due to collisional sticking of smaller dust aggregates has to face at least two severe obstacles, namely the rapid loss of material due to radial inward drift and particle fragmentation due to destructive collisions. We present a scenario to circumvent these two hurdles. Our dust evolution model involves two main mechanisms. First, we consider a disk with a dead zone. In an almost laminar region close to the midplane, the relative velocities of the turbulent particles are comparatively small, which decreases the probability of destructive particle collisions. Second, turbulence is not the only source of violent relative particle velocities, because high radial drift speeds can also lead to boulder fragmentation. For this reason, we focus additionally on the snow line. Evaporation fronts can be associated with gas pressure maxima in which radial drift basically vanishes. This implies that particle fragmentation becomes even less likely. Our simulation results suggest that particles can overcome the fragmentation barrier. We find that boulders of several $10^2$~m can form within only a few thousand years.} ",
        "introduction": "The process of planetary formation is widely accepted to be initiated in circumstellar disks by the collisional coagulation of dust grains to larger objects \\citep{h06,natta07}. Evidence of grain growth to particle sizes beyond those typically observed in the interstellar medium is provided by mid-infrared spectroscopy of disks around young stars \\citep{bouwman01,boekelaebe03,apai04,kessler07,sic}. Analyses of millimeter interferometry data of disks implies that large populations of dust grains exist with particle radii up to several millimeters \\citep{Tes03,Wil05,Rod06}.  Models of dust particle growth in protoplanetary disks predict coagulation, which supports the idea of planetesimal formation by particle hit and stick mechanisms. Numerical calculations support the formation of cm-sized particles within a few $10^3$ orbital timescales in the inner parts of the disk \\citep{Wei80}. After particles have grown to larger sizes, they can also be affected by other processes described by alternative planet formation scenarios. For example, large grains experience vertical settling, which involves the formation of a dense midplane dust layer \\citep{garlin2004}. This layer can be affected by gravitational instability \\citep{GolWar73,Wei79,YouShu02,schr04,garlin2004} -- an issue that is still highly debated. Due to the overwhelming amount of observational data for protoplanetary disks available, models of dust particle growth have attracted more attention \\citep{tan05,DD05,nomura06,allesio06,ormel07_2,ciesla07}. All of these particle evolution models can explain observational disk features such as, for example, the group I/II disk classification \\citep{meeus01,dd04_} or the disappearance of infrared excess in the spectra of disks with ages higher than a few Myrs \\citep{hai01,carp05}, although, there are still many open observational riddles that these models will hopefully answer in the near future.  Models of protoplanetary disks do not only help to interpret observational data; they also identify the serious obstacles to planetesimal formation \\citep{you04,domppv07}. One of these obstacles is particle fragmentation. Due to the high particle relative velocities in disks that can reach 100~m/s \\citep{Wei77,ormel07}, particle collisions lead to particle destruction instead of particle growth \\citep{BluWur99,brauer2,anners08}. Depending on the disk model, dust particle growth is inhibited significantly by high-speed impacts of particles around a meter size. Even for times as long as 1~Myr, solid particles are unable to overcome the fragmentation barrier and the formation of planetesimals as precursors for Earth-like or Jovian planets, therefore, poses a major problem. Despite this theoretically predicted inability for planetesimals to form, the orbital decay of dust particles can, nevertheless, reproduce the observed evolution in the spectral energy distribution \\citep{hai01,carp05}. The presence of a growth barrier is, therefore, not in contradiction with observations.  Several mechanisms in protoplanetary disks can produce planetesimal formation. \\cite{andersnat07} showed that the non-linear feedback of dust onto gas can lead to the rapid formation of gravitationally bound clumps of dust, which subsequently form Ceres-size bodies. The dust particles, however, must have already grown to some meters in size before this scenario can take place. Another possibility for solving the formation problem is particle trapping in gas pressure maxima \\citep{bargesomm95,KlaHen97}. In gas pressure bumps, relative particle velocities are substantially lower and, hence, the collision between large bodies is more likely to lead to particle growth than to particle disruption.  Interesting means of gas pressure maxima and dust particle trapping was presented by \\cite{kretke07}. This particle retention mechanism requires the presence of an evaporation front, for example the snow line, which acts in the following way. As we travel through the snow line in a direction away from the central star, the dust-to-gas ratio increases suddenly. This jump in dust density affects the strength of the magneto-rotational turbulence, since the amount of free electrons in the disks strongly depends on the dust density. With increasing dust density, the amount of turbulence in the disk decreases \\citep{sano00,ilg106}. We assume a constant mass accretion rate throughout the disk. The gas surface density then has to be higher in low than in high disk turbulent regions for the gas mass accretion to remain at a similar value. The sudden increase in the dust density could therefore also reproduce a sudden increase in gas densities. For certain accretion rates, \\cite{kretke07} found the occurrence of a local gas density bump in which solid particles tend to accumulate. In their simulations, dust particle retention reproduced very high surface dust densities of the order of several $10^3$~g/cm$^2$, which again raises concern about gravitational instability.  However, even for a particle retention mechanism similar to that described above, the growth of solid material towards larger sizes remains an open issue. Gas pressure maxima may decrease relative radial particle velocities triggering coagulation; the most severe reason for violent particle fragmentation is not however radial drift but the turbulent nature of the protostellar disk itself. Relative turbulent particle velocities are approximately $10$~m/s \\citep{ormel07,voelk80} and particle sticking at these high speeds is extremely unlikely. Nevertheless, \\cite{ciesla07} found that particles can grow to several 10~m in radius if we consider layered MRI active disks. Under certain circumstances, MRI is only active in the upper layers of the disk, while the disk midplane is almost laminar. Since most of the larger grains are located around the midplane where the disk is quiescent, dust growth is not inhibited by the high speed collisions that produce particle fragmentation.  We combine three ingredients for a planetesimal formation model. We consider dust particle growth, particle fragmentation and radial motion \\citep{brauer2} around the snow line \\citep{kretke07} in a layered MRI active protoplanetary disk \\citep{ciesla07}. The inclusion of a snow line in our simulations provides a particle retention mechanism which almost erases the radial drift velocities favouring particle growth. We adopt a layered MRI-driven disk to circumvent particle fragmentation in the midplane of the disk due to turbulent motions of the gas. We investigate whether solid particles can overcome the fragmentation barrier and produce larger objects, which represent the possible precursors to the planets. We study the influence of two parameters, namely the gas accretion rate $\\dot{M}_{\\mathrm{acc}}$ and the critical threshold velocity for fragmentation $v_{\\mathrm{f}}$. ",
        "conclusions": "Even though particle fragmentation by high speed collisions is a severe obstacle to the formation of planetesimals, we have shown that dust particle growth to the size of enormous boulders is possible under certain circumstances. If we consider dust particle coagulation in the presence of an evaporation front, solid particles grow to almost km-sizes within only a few thousand years. The protoplanetary disk must also contain a dead zone in which turbulence -- the main source of relative particle velocities -- is almost entirely absent. Our conclusion is that planetesimal formation due to dust particle agglomeration is a possible mechanism to form large bodies in protoplanetary disks."
    },
    "0806/0806.3369_arXiv.txt": {
        "abstract": "{The study of high redshift Tully-Fisher relations (TFRs) is limited by the use of long slit spectrographs, rest frame B band and star formation selected galaxies.  } % {We try to circumvent these issues by using integral field spectroscopy (SINFONI), by studying the rest frame K band and stellar mass TFR, and by selecting targets without a bias to strongly star forming galaxies. In this paper, we demonstrate our methods on our best case. This galaxy, F257, at $z=2.03$, was selecte from a sample of  candidate high redshift large disk galaxies in the Hubble Deep Field South that were selected with photometric and morphological criteria. } {We used SINFONI at the VLT to obtain an integral field spectrum of the H$\\alpha$ line and hence a velocity field and rotation curve. We also use UBVIJHK+IRAC band photometry to determine a stellar photometric mass.    } {We find that F257 is indistinguishable from local late type galaxies in many respects: it has a regular velocity field, increasing velocity disperion towards its center, its rotation curve flattens at 1-2 disk scale lengths, it has the same specific angular momentum as local disks, its properties are consistent with the local K band TFR. Although mainly rotationally supported, its gas component is  dynamically heated with respect to local galaxies ($V/\\sigma_{z}\\sim4$) and it is offset from the local stellar mass TFR at the $2\\sigma$ level.  But, this offset depends on the SED modeling parameters. In particular, for a 2-component star formation history (SFH), F257 is in agreement with the local stellar mass TFR. F257 is then a nearly ($\\sim 75\\%$) maximum disk. The dynamical properties of F257 are more like those of local galaxies than those of any other galaxy at similar redshift observed to date. However, the gas-to-stellar mass ratio is unusally large: 2.5. } {} ",
        "introduction": "Starting with density fluctuations in the early universe, dark matter halos grow through gravity. Galaxy formation is described by semi-analytical models or hydrodynamical simulations. These may include physical processes such as gas cooling, star formation, feedback from stellar winds, supernovae and/or AGN, chemical evolution and stellar populations (Okamoto \\cite{okamoto} and references therein).  The disk formation epoch in earlier generations of models is relatively late and mostly at $z \\lesssim 1$. These early models also predicted that these disks are much smaller than observed in local spiral galaxies and  that their specific angular momenta are only  20\\% of that of their halo's (Navarro, Frenk \\& White \\cite{navarro}). More recent models produce more realistic size disks (Mo et al. \\cite{mo1998}, Portinari \\& Sommer-Larsen \\cite{portinari}, Governato et al. \\cite{governato}) and place the formation epoch somewhat earlier: $0.4<z<2.0$ (Buchalter, Jimenez \\& Kamionkowski \\cite{buchalter}). Reproducing the specific angular momentum of local disk galaxies remains a challenge (e.g. Portinari \\& Sommer-Larsen \\cite{portinari}).   More than 30 years after the first publication of what is now known as the Tully-Fisher relation (TFR) (Tully \\& Fisher \\cite{tullyfisher}), the origin of the TFR is still subject of debate. Some authors explain the TFR as a consequence of self-regulated star formation in disks with different masses (e.g. Silk \\cite{silk}). In other models, the TFR is a direct consequence of the cosmological equivalence between mass and circular velocity resulting from the finite age of the universe which imposes a maximum radius from with matter can accrete to form a disk (e.g. Mo et al. \\cite{mo1998}, Steinmetz \\& Navarro \\cite{steinmetz}). Steinmetz \\& Navarro (\\cite{steinmetz}) show that although the slope of the TFR can be naturally explained by hierarchical models, the normalization and evolution of the TFR depend strongly on the star formation prescriptions and the cosmological parameters. Recently, Courteau et al. (\\cite{courteau}) showed that two long known, different theoretical predictions for two different slopes of the TFR are in fact related.   Disk galaxy scaling relations, like the TFR, put strong constraints on galaxy formation scenarios, and in particular on the connection of the dark and baryonic component (van den Bosch \\cite{bosch}). The zero-point of the TFR constrains the SFH and the cosmological parameters that control the universal baryon fraction and the time of assembly of galaxies of different mass (Steinmetz \\& Navarro \\cite{steinmetz}). Many authors struggle to reproduce the zero-point of the TFR (e.g. Portinari \\& Sommer-Larsen \\cite{portinari}, Steinmetz \\& Navarro \\cite{steinmetz}). The slope of the TFR and the steepening of the slope towards longer wavelengths can only be explained by supernova feedback (van den Bosch \\cite{bosch}) although others claim it is a natural result of the hierarchical models (Steinmetz \\& Navarro \\cite{steinmetz}). The small scatter in the TFR is explained by scatter in halo spin parameters and formation redshifts (e.g. van den Bosch \\cite{bosch}).   Even the most recent galaxy formation and evolution models are not able to simultaneously reproduce for example the slope, zero-point and scatter of the TFR and the zero-point and shape of the luminosity function (e.g. Courteau et al. \\cite{courteau} and references therein).    The TFR defines a tight correlation between the luminosity and the flat rotation curve velocity of spiral galaxies. If luminosity is converted into stellar mass (which requires detailed modelling of the stellar spectral energy distribution (SED)), the tight correlation remains: the stellar mass TFR. Semi-analytic models predict evolution in the slope, zero-point and scatter properties of the TFR and the stellar mass TFR (e.g. Steinmetz \\& Navarro \\cite{steinmetz}, Buchalter, Jimenez \\& Kamionkowski \\cite{buchalter}, Portinari \\& Sommer-Larsen \\cite{portinari}). A key element is the relative evolution of the  dark and stellar mass components, which is a direct test of the stellar mass buildup within a dark halo of a given mass.     Critical tests of these theories can be made by observing the evolution of disks backward in time. Rotation curves (RCs) of field galaxies out to $z = 1$ measured with the Keck telescope and Very Large Telescope (VLT) are found to be similar to those of nearby galaxies in both shape and amplitude implying that some massive disks were already in place at $z \\sim 1$ (Vogt et al. \\cite{vogt1996}, \\cite{vogt1997}, Barden et al. \\cite{barden}, Swinbank et al. \\cite{swinbank}). Observations of $z \\sim 2$ galaxies have revealed ordered rotation, but only the rising part of the rotation curve (e.g. the lensed sources of Lemoine-Busserolle et al. \\cite{lemoine}). Van Starkenburg et al. (\\cite{vanstarkenburg}) found ordered rotation in a $z = 2.46$ galaxy, with an asymmetric line profile indicating a central starburst and weaker starformation in the outer parts. Ordered rotation has been found out to $z \\sim 3.2$ by Moorwood et al. (\\cite{moorwood00}, \\cite{moorwood03}) and by Nesvadba et al. (\\cite{nesvadba6}), the latter again on a lensed source.  Also, several attempts to measure size evolution and the TFR at redshift $0.5 - 1.5$ have been made over the past few years with contradictory claims about disk evolution. For instance, Mo et al. (\\cite{mo1998}) used previously acquired data (Vogt et al.\\cite{vogt1996}, \\cite{vogt1997}; Schade et al. \\cite{schade}; Lilly et al. \\cite{lilly}) to conclude that $z \\sim 1 $  galaxies are smaller than local galaxies at a fixed rest-frame B-magnitude or rotation velocity, consistent with hierarchical models. A brightening with redshift of disk galaxies at fixed rotation speeds was also detected by Barden et al. (\\cite{barden}). However, Simard et al. (\\cite{simard}) and Ziegler et al. (\\cite{ziegler}) have shown that massive galaxies at $z < 1$ are consistent with having no evolution in surface brightness, once selection effects are taken into account.  The largest source of uncertainty in those studies, all done with long slit spectra, are the velocity measurements. In recent years, it has become possible to do integral field spectroscopy at near infrared wavelengths, for example with SINFONI at the VLT (Eisenhauer et al. \\cite{eisenhauer}, Bonnet et al. \\cite{bonnet}). Integral field spectroscopy has significant advantages over long slit spectroscopy. First, long slit measurements may suffer from slit misalignment. More fundamentally, the photometric and kinematic axes may be different (due to star forming regions or intrinsic shape of galaxy). Kinematic disturbances are harder or impossible to identify in slit spectra. Integral field spectroscopy does not have all these slit related problems. Finally, integral field spectroscopy makes for a fairer comparison with low-$z$ Tully-Fisher data, which are typically derived from full imaging of the HI velocity field with radio aperture synthesis. Integral field spectroscopy thus presents a major improvement over slit spectroscopy, provided significant resolution over the rotating disk is achieved. Recently, F\\\"orster Schreiber et al. (\\cite{forster}) and Bouch\\'e et al. (\\cite{bouche}) used SINFONI to obtain two-dimensional velocity fields of a number of galaxies at $z \\sim 2$ with different selection criteria, clearly demonstrating the power of an efficient integral-field spectrograph such as SINFONI for this type of work.   A second major problem that plagued previous high-$z$ TFR studies is the use of rest-frame B-band, which is very sensitive to small amounts of recent star formation and extinction by dust. There are now a number of deep surveys available with comprehensive photometric coverage ideal for studying the high redshift TFR in rest frame near-IR luminosities. We used the HDF-S and MS1054 fields which includes recent Spitzer IRAC data (Labb\\'e et al. \\cite{irac}). This allows us to use rest-frame near-IR luminosities, which define much tighter TFRs than rest-frame optical luminosities (e.g. Verheijen \\cite{verheijen}). Rest-frame K band luminosities as derived from Spitzer data are much less sensitive to the effects of dust and recent starformation than rest frame optical or UV luminosities and will provide us with more accurate stellar masses.    A third limitation of previous studies is that many of these target starformation-selected galaxies (e.g., selected on rest-frame UV or bright emission lines). This compromises attempts to use these galaxies to trace the evolution of the TFR (see van Starkenburg et al. \\cite{vanstarkenburg} for a detailed discussion), since local TFR samples are composed of relatively quiescent galaxies. Although a certain amount of star formation is necessary to observe the velocity fields of high redshift galaxies, there are several ways to partly circumvent this problem.    First, by selecting from ultra-deep near-IR imaging surveys, such as the FIRES data of the HDF-S and MS1054 fields\\footnote{The FIRES data are publicly available, see http://www.strw.leidenuniv.nl/\\~{}fires}(Labb\\'e et al. \\cite{fireshdfs}, Franx et al. \\cite{franx}, F\\\"orster-Schreiber et al. \\cite{firesms1054}), we can select galaxies in the rest frame optical instead of the rest frame UV. The FIRES data is sufficiently deep to select galaxies in the rest-frame optical out to $z \\sim 4$ and the depth on the Hubble Deep Field South (HDF-S) is sufficient to detect non-evolving ($0.6L_*$  at $z = 0$) spirals out to $z = 3$.   A second improvement is to preselect massive galaxies. Contrary to low mass galaxies, their $\\log(L)$ (and hence offset from the TFR) is not or less affected by the luminosity of  a starburst  (van Starkenburg et al. \\cite{vanstarkenburg}). We observed several targets in this category in two redshift ranges, which will be discussed in two separate papers (van Starkenburg et al. \\cite{prep1}, \\cite{prep2}).   Finally, it is now possible to select high redshift galaxies that morphologically resemble local disks. A search for such disks in the FIRE survey of the  2\\farcm5 x 2\\farcm5 HDF-S field resulted in a surprisingly large number (6) of candidate large disks (Labb\\'e et al. \\cite{labbe}). These galaxies, at redshifts 1.4 to 3.0 (most of them spectroscopically confirmed), have exponential profiles with effective radii ranging from 0\\farcs65 to 0\\farcs9 (or $5.0 - 7.5 h^{-1}_{70}$ kpc) and extending over 2-3 effective radii and appear remarkably regular in the near infrared (rest frame optical). Four are well fit by just exponential disks and two require an additional $r^{1/4}$  bulge  component, contributing about 40\\% of the light. At shorter wavelengths, they have more extended and irregular morphologies. They have strong rest-frame optical breaks in their centres and evidence for intense star formation in the outer parts, reminiscent of local spirals, with blue star-forming disks and red bulges. Their sizes are comparable to the Milky Way and those of L* disks in the local universe. Specifically, these sizes are much larger than those of typical LBGs ($1 - 2 h^{-1}_{70}$ kpc; Giavalisco et al. \\cite{giavalisco}; Lowenthal et al. \\cite{lowenthal})) and their sizes are in the upper end of the range found by Trujillo et al. (\\cite{trujillo}) for all galaxies at similar $z$ in the two FIRES fields. They are bright in K (K$ = 19.7 -  21.1$) and have large stellar masses ($M > 10^{11}\\textrm{ M}_\\odot$). They constitute half of the most rest-frame luminous galaxies ($L_V > 6 \\times 10^{10} h^{-2}_{70} L_\\odot $) and their number density exceeds model predictions (e.g., Mo et al., \\cite{mo1998}) by a factor of two. Obviously, the first and foremost question that needs to be resolved is whether these objects are truly rotationally supported disks, and this can only be tested by obtaining kinematic data.    In this paper, we present SINFONI observations of one of the large disks of the Labb\\'e et al. sample (id 257 in the FIRES HDF-S catalog). These observations were part of a larger project of SINFONI observations of mass and/or disk selected high redshift galaxies, which will be presented in two separate papers (van Starkenburg et al. \\cite{prep1}, \\cite{prep2}). Here, we will discuss our methods as used on our best case.   We observed F257 with SINFONI during the night of August 15, 2006. We selected F257 as our primary target for this observing run because it is the second brightest K band source ($K=20.25$) of the large disks in the HDF-S (Labb\\'e et al. \\cite{labbe})  that  has a known spectroscopic redshift ($z=2.03$, Vanzella et al. \\cite{vanzella}) with H$\\alpha$ at an observable wavelength with SINFONI. (We had observed the brightest K band target fullfilling this criteria during an earlier observing run, see van Starkenburg et al. \\cite{prep2}). Its effective radius is  0\\farcs74 in K band, the total radial extent  above the K-band detection threshold is two effective radii or 3-3.5 scale lengths ($R_d=0\\farcs44$). The inclination ($\\sim 54\\degr$, calculated from the ellipticity of Labb\\'e et al. \\cite{labbe}) is suited for a velocity field analysis. F257 shows a clumpy distribution in the HST F814W filter (see Fig. \\ref{images}a) which corresponds to rest frame UV. The rest frame optical (observed K band) image shows a smooth exponential profile (see Fig. \\ref{images}c). This difference cannot be attributed to the different PSF of the F814W and the K band image: the F814W image smoothed to the K band PSF still shows the clumps. Labb\\'e et al. (\\cite{labbe}) fitted an exponential profile without the need for a bulge component.  NICMOS H band shows an intermediate picture (see Fig. \\ref{images}b). The S\\'ersic index as determined by Trujillo et al. (\\cite{trujillo}) on the FIRES (ISAAC) H band image is 0.76. Wuyts et al. (\\cite{wuyts}) made an extensive study of the SEDs of high redshift galaxies in HDF-S and found for F257 $M_* = 3.5^{ +0.5}_{ -0.3} \\times 10^{10} \\textrm{ M}_\\odot $, a mass weighted age of $\\sim160 \\textrm{ Myr}$ and $A_V=1.0$ using a Salpeter initial mass function (IMF), solar metallicity, a $\\tau$ 300 model for the star formation history (SFH) and Bruzual \\& Charlot (\\cite{bruzual}) models. A number of $z\\sim2$ galaxies have been observed with SINFONI and other IFUs (Genzel et al. \\cite{genzel}, F\\\"orster-Schreiber et al. \\cite{forster}, Bouch\\'e et al. \\cite{bouche}, Law et al. \\cite{law}). F257 is not different in terms of colors, SFR or age, and satisfies the conventional U-dropout criterium. It is the only galaxy selected as large disk galaxy, but the other galaxies were selected from catalogs with insufficient resolution to identify morphological large disk galaxies at optical rest frame wavelengths.   \\begin{figure}  \\includegraphics[width=8cm]{9632f1a.ps} \\includegraphics[width=8cm]{9632f1b.ps} \\includegraphics[width=8cm]{9632f1c.ps}   \\caption{From top to bottom: F814W image, NICMOS H band image and Ks image from FIRES data. Image size is always 3\\arcsec times 4\\arcsec and the images have been rotated so that east is to the top and north is to the right. The (K band) center from the FIRES data is marked by the white crosses.} \\label{images}  \\end{figure}      Throughout this paper, we use $H_0 = 70 \\textrm{ km s}^{-1}\\textrm{Mpc}^{-1}$, $\\Omega_M=0.3$, $\\Omega_\\Lambda=0.7$ and Vega magnitudes. ",
        "conclusions": "Labb\\'e et al. (\\cite{labbe}) presented photometric evidence that F257 is a large disk galaxy at $z=2.03$. We confirm the disk nature of this galaxy by its kinematical properties and find that is in many respects similar to local disk galaxies. Specifically, we find: \\begin{enumerate} \\item The velocity field is consistent with a rotating disk galaxy, a tilted ring model leaves small random residuals and no evidence for non circular motions; \\item There is no evidence for an AGN in the form of broad lines or a large [\\ion{N}{ii}]/H$\\alpha$ ratio; \\item The [\\ion{N}{ii}]/H$\\alpha$ ratio peaks in the centre; \\item The disk is primarily rotation supported: $V/\\sigma_z\\sim3.9$; \\item The dispersion increases towards the center; \\item The RC slowly rises like that of a late-type spiral in the local universe and then flattens. There is no evidence for dark matter within $r < 12 \\textrm{kpc}$; \\item F257 has a specific angular momentum and spin parameter comparable to those of local galaxies; \\item F257 is consistent with the local K band TFR; \\item F257 lies below the local stellar mass TFR at the $2\\sigma$ level and is consistent with the local baryonic mass TFR. A 2-component model for the SFH increases the stellar mass so that F257 is consistent with the stellar and baryonic mass TFR and makes F257 a nearly maximum disk.  \\end{enumerate} Despite its very young age, high star formation rate, gas mass and clumpy star formation, F257 is very similar to local late type spiral galaxies. Its dynamical properties are more like those of local galaxies than those of other galaxies observed at similar redshift.  These findings put constraints on the evolution of F257 from $z=2.03$ to the present. If the stellar mass increases so that F257 becomes consistent with the local stellar mass TFR, the K band mass-to-light ratio must also increase so that that the total K luminosity does not change. Infalling gas must have the same specific angular momentum as the galaxy, i.e. it must fall in from the outer parts of the present halo."
    },
    "0806/0806.3643_arXiv.txt": {
        "abstract": "{We use a sample of galaxies from the Sloan Digital Sky Survey (SDSS) to search for correlations between the $\\lambda$ spin parameter and the environment and mass of galaxies. In order to calculate the total value of $\\lambda$ for each observed galaxy, we employed a simple model of the dynamical structure of the galaxies which allows a rough estimate of the value of $\\lambda$ using only readily obtainable observables from the luminous galaxies. Use of a large volume limited sample (upwards of 11,000) allows reliable inferences of mean values and dispersions of $\\lambda$ distributions. We find, in agreement with some N-body cosmological simulations, no significant dependence of $\\lambda$ on the environmental density of the galaxies. For the case of mass, our results show a marked correlation with $\\lambda$, in the sense that low mass galaxies present both higher mean values of $\\lambda$ and associated dispersions, than high mass galaxies. This last direct empirical result, at odds with expectations from N-body cosmological simulations, provides interesting constrain on the mechanisms of galaxy formation and acquisition of angular momentum, a valuable test for cosmological models. }    \\addkeyword{galaxies:formation} \\addkeyword{galaxies: statistics} \\addkeyword{galaxies: structure} \\addkeyword{cosmology: observations}   \\begin{document} ",
        "introduction": "\\label{sec:intro}   One of the most studied parameters in numerical simulations of formation and evolution of galaxies, is the $\\lambda$ spin parameter, first introduced by Peebles in 1969, where it was used to test the gravitational instability picture, as a theory for the origin of galactic angular momentum. Since then, there have been several studies using the $\\lambda$ parameter as an indispensable tool of analysis, or characterizing its distribution in cosmological N-body simulations. In general, the distribution of this parameter, coming from numerical simulations, is well fitted by a log-normal distribution, characterized by two parameters; $\\lambda_{0}$, the most probable value, and $\\sigma_{\\lambda}$, which accounts for the spread of the distribution. In a recent work, Shaw et al. (2006) showed a compilation of results for estimates of these parameters, arising from numerical studies performed by several authors, lying in the ranges: $0.03<\\lambda_{0}<0.05$ and $0.48<\\sigma_{\\lambda}<0.66$, and some of us in a previous work (Hernandez et al. 2007), using the same sample of 11 597 galaxies from the Sloan Digital Sky Survey (SDSS) we treat here, obtained the distribution of this parameter, with results well fitted by a log-normal function with parameters $\\lambda_{0}=0.04\\pm 0.005$ and $\\sigma_{\\lambda}=0.51 \\pm 0.05$, for the first time derived from a statistical sample of real galaxies.  With the recent explosion of high resolution cosmological simulations, studies concerning the spin parameter begun to explore dependencies of the $\\lambda$ parameter with other physical parameters, mainly the total mass of the galaxy and the environment in which the galaxy is immersed. The results are as varied as the techniques employed to perform and analyse the simulations, but  there is good agreement that for the case of a $\\lambda$-environment relation, it should be very weak or inexistent (Bett et al. 2007, Maccio et al 2007). The case of the $\\lambda$-mass relation presents more discrepancies between results coming from different groups, some claiming the existence of a weak relation (Jang-Codell \\& Hernquist 2001, Bett et al. 2007) and others not finding any (Maccio et al. 2007). The aim of the present work is to derive distributions of the spin parameter as functions of mass and environment, using a first order estimate of galactic halo $\\lambda$ applied to a large sample of galaxies from the Sloan Digital Sky Survey (SDSS), in order to constrain and infrom theories of angular momentum acquisition and galaxy formation. ",
        "conclusions": ""
    },
    "0806/0806.3705_arXiv.txt": {
        "abstract": "{% The cosmological origin of Gamma-Ray Bursts (GRBs) has been firmly established, with redshifts up to $z = 6.29$. GRBs are possible candidates to be used as ``distance indicators'' to test cosmological models in a redshift range hardly achievable by other cosmological probes. Asserting the validity of the empirical relations among GRB observables is now crucial for their calibration. }{% Motivated by the relation proposed by Amati and collaborators, we look within the ``fireshell'' model for a relation between the peak energy $E_p$ of the $\\nu F_\\nu$ total time-integrated spectrum of the afterglow and the total energy of the afterglow $E_{aft}$, which in our model encompasses and extends the prompt emission. }{% The fit within the fireshell model, as for the ``canonical'' GRB050315, uses the complete arrival time coverage given by the Swift satellite. It is performed simultaneously, self-consistently and recursively in the four BAT energy bands ($15$--$25$ keV, $25$--$50$ keV, $50$--$100$ keV and $100$-$150$ keV) as well as in the XRT one ($0.2$--$10$ keV). It uniquely determines the two free parameters characterizing the GRB source, the total energy $E^{e^\\pm}_{tot}$ of the $e^{\\pm}$ plasma and its baryon loading $B$, as well as the effective CircumBurst medium (CBM) distribution. We can then build two sets of ``gedanken'' GRBs varying the total energy of the electron-positron plasma $E^{e^\\pm}_{tot}$ and keeping the same baryon loading $B$ of GRB050315. The first set assumes for the effective CBM density the one obtained in the fit of GRB050315. The second set assumes instead a constant CBM density equal to the average value of the GRB050315 prompt phase. }{% For the first set of ``gedanken'' GRBs we find a relation $E_p\\propto (E_{aft})^a $, with $a = 0.45 \\pm 0.01$, whose slope strictly agrees with the Amati one. Such a relation, in the limit $B \\to 10^{-2}$, coincides with the Amati one. Instead, in the second set of ``gedanken'' GRBs no correlation is found. }{% Our analysis excludes the Proper-GRB (P-GRB) from the prompt emission, extends all the way to the latest afterglow phases and is independent on the assumed cosmological model, since all ``gedanken'' GRBs are at the same redshift. The Amati relation, on the other hand, includes also the P-GRB, focuses on the prompt emission only, and is therefore influenced by the instrumental threshold which fixes the end of the prompt emission, and depends on the assumed cosmology. This may well explain the intrinsic scatter observed in the Amati relation.} ",
        "introduction": "The detection of Gamma-Ray Bursts (GRBs) up to very high redshifts \\citep[up to $z = 6.29$, see][]{2005A&A...443L...1T}, their high observed rate of one every few days, and the progress in the theoretical understanding of these sources make them useful as cosmological tools, complementary to Supernovae Ia which are observed only up to $z = 1.7$ \\citep{2001ARA&A..39...67L,2001ApJ...560...49R}. One of the hottest topics on GRBs is constituted by the possible existence of empirical relations between GRB observables \\citep{2002A&A...390...81A,2004ApJ...616..331G,2004ApJ...609..935Y,2005ApJ...633..611L,2006MNRAS.370..185F,2008arXiv0805.0377A}, which may lead, if confirmed, to use GRBs as tracers of models of universe. The first empirical relation, discovered analyzing the \\emph{BeppoSAX} so-called ``long'' bursts with known redshift, was the ``Amati relation'' \\citep{2002A&A...390...81A}. It was found that the isotropic-equivalent radiated energy of the prompt emission $E_{iso}$ is correlated with the cosmological rest-frame $\\nu F_{\\nu}$ spectrum peak energy $E_{p,i}$: $E_{p,i}\\propto (E_{iso})^{a}$, with $a = 0.52 \\pm 0.06$ \\citep{2002A&A...390...81A}. The existence of the Amati relation has been confirmed by studying a sample of GRBs discovered by Swift, with $a = 0.49^{+0.06}_{-0.05}$ \\citep{2006ApJ...636L..73S,2006MNRAS.372..233A}.  Swift has given for the first time the possibility to obtain high quality data in selected energy bands from the GRB trigger time all the way to the latest afterglow phases \\citep{2004ApJ...611.1005G}. This has given the opportunity to apply our theoretical ``fireshell'' model obtaining detailed values for its two free parameters, the total energy $E^{e^\\pm}_{tot}$ and the baryon loading $B$ of the fireshell, as well as the values of the effective density and filamentary structure of the CircumBurst Medium (CBM). This allowed to compute multi-band light curves and spectra, both instantaneous and time-integrated, compared with selected GRB sources, like e.g. GRB050315.  In the ``fireshell'' model $E^{e^\\pm}_{tot}$ comprises two different components: the Proper GRB (P-GRB), with energy $E_{P-GRB}$, emitted at the moment when the $e^+e^-$-driven accelerating baryonic matter reaches transparency, and the following afterglow phase, with energy $E_{aft}$, with the decelerating baryons interacting with the CBM \\citep{2001ApJ...555L.113R}. These two phases are clearly distinguishable by their relative intensity and temporal separation in arrival time. We have: \\begin{equation} E^{e^\\pm}_{tot} = E_{P-GRB} + E_{aft}\\, . \\label{Etot} \\end{equation} What is usually called the ``prompt emission'' corresponds within the fireshell model to the P-GRB together with the peak of the afterglow \\citep[see below, e.g.][and references therein]{2001ApJ...555L.113R,2006ApJ...645L.109R,2007AIPC..910...55R,2007A&A...471L..29D,2007A&A...474L..13B,2008AIPC..966...16C,2008AIPC..966...12B}.  Among the crucial issues raised by the Amati relation there are its theoretical explanation and its possible dependence on the assumed cosmological model. We have examined a set of ``gedanken'' GRBs, all at the same cosmological redshift of GRB050315. Such a set assumes the same fireshell baryon loading and effective CBM distribution as GRB050315 and each ``gedanken'' GRB differs from the others uniquely by the value of its total energy $E^{e^\\pm}_{tot}$. We have then considered a second set of ``gedanken'' GRBs, differing from the previous one by assuming a constant effective CBM density instead of the one inferred for GRB050315. In both these sets we looked for a relation between the isotropic-equivalent radiated energy of the \\emph{entire} afterglow $E_{aft}$, and the corresponding time-integrated $\\nu F_{\\nu}$ spectrum peak energy $E_p$: \\begin{equation} E_p\\propto (E_{aft})^a\\, . \\label{al} \\end{equation}  In section \\ref{model} we recall the main features of the ``fireshell'' model. In section \\ref{fit} we recall the main features of the GRB050315 fitting procedure. In section \\ref{correlation} we then present the derivation of the $E_p$ -- $E_{aft}$ relation for two sets of ``gedanken'' GRBs. The results of this analysis and their discussion are shown in section \\ref{results}. In section \\ref{conclusions} we present the conclusions, strongly confirming the validity of the Amati relation. ",
        "conclusions": "The Swift high quality data, giving for the first time gapless and multiwavelength coverage from the GRB trigger all the way to the latest afterglow phases, have allowed to make a complete fit of the GRB050315 multiband light curves based on our fireshell model. We have so fixed the free parameters describing the source and we determined the instantaneous and time-integrated spectra during the entire afterglow.  Starting from this, we have examined two sets of ``gedanken'' GRBs, constructed at a fixed cosmological redshift. The first set assumes the same fireshell baryon loading and effective CBM distribution as GRB050315 and each ``gedanken'' GRB differs from the others uniquely by the value of its total energy $E^{e^\\pm}_{tot}$. The second set assumes a constant effective CBM density $\\sim 1$ particle/cm$^3$ instead of the one inferred for GRB050315.  Recalling that the ``canonical'' GRB light curve in the fireshell model is composed of two well separated components, the P-GRB and the entire afterglow, we have looked in both sets for a relation between the isotropic-equivalent radiated energy of the \\emph{entire} afterglow $E_{aft}$ and the corresponding time-integrated $\\nu F_{\\nu}$ spectrum peak energy $E_p$: $E_p\\propto (E_{aft})^a$. In doing so, we have assumed that the Amati relation is directly linked to the interaction between the accelerated baryons and the CBM. The P-GRBs, which originates from the fireshell transparency, in our approach do not fulfill the Amati relation. Consequently, also the short GRBs, which have a vanishing afterglow with respect to the P-GRB, should not fulfill the Amati relation. This last point is supported by the observational evidence \\citep{2006MNRAS.372..233A}.  We notice that the first set of ``gedanken'' GRBs fulfills very well the $E_p\\propto (E_{aft})^a$ relation with $a = 0.45 \\pm 0.01$. This slope strongly agrees with the Amati relation. On the contrary, for the second set no relation between $E_p$ and $E_{aft}$ seems to hold. We conclude that the Amati relation originates from the detailed structure of the effective CBM.  Turning now to the analogies and differences between our $E_p$ -- $E_{aft}$ relation and the Amati one, our analysis excludes the Proper-GRB (P-GRB) from the prompt emission, extends all the way to the latest afterglow phases and is independent on the assumed cosmological model, since all ``gedanken'' GRBs are at the same redshift. The Amati relation, on the other hand, includes also the P-GRB, focuses on the prompt emission only, and is therefore influenced by the instrumental threshold which fixes the end of the prompt emission, and depends on the assumed cosmology. This may well explain the intrinsic scatter observed in the Amati relation \\citep{2006MNRAS.372..233A}. Our theoretical work is a first unavoidable step to support the usage of the empirical Amati relation for measuring the cosmological parameters."
    },
    "0806/0806.1700_arXiv.txt": {
        "abstract": "An increase in the central density of a neutron star may trigger a phase transition from hadronic matter to deconfined quark matter in the core, causing it to collapse to a more compact hybrid-star configuration. We present a study of this, building on previous work by \\citet{lin_06_a}. We follow them in considering a supersonic phase transition and using a simplified equation of state, but our calculations are general relativistic (using 2D simulations in the conformally flat approximation) as compared with their 3D Newtonian treatment. We also improved the treatment of the initial phase transformation, avoiding the introduction of artificial convection. As before, we find that the emitted gravitational-wave spectrum is dominated by the fundamental quasi-radial and quadrupolar pulsation modes but the strain amplitudes are much smaller than suggested previously, which is disappointing for the detection prospects. However, we see significantly smaller damping and observe a nonlinear mode resonance which substantially enhances the emission in some cases. We explain the damping mechanisms operating, giving a different view from the previous work. Finally, we discuss the detectability of the gravitational waves, showing that the signal-to-noise ratio for current or second generation interferometers could be high enough to detect such events in our Galaxy, although third generation detectors would be needed to observe them out to the Virgo cluster, which would be necessary for having a reasonable event rate. ",
        "introduction": "\\label{section:introduction}  The existence of compact stellar objects partially or totally consisting of matter in a deconfined quark phase was already predicted long ago \\citep{bodmer_71_a, itoh_70_a, witten_84_a}. Such stars are thought likely to originate as a result of the conversion of purely hadronic matter in the interior of a neutron star (NS) into a deconfined quark matter phase when the density exceeds a certain threshold \\citep[for a review see, e.g.][]{weber_99_a, glendenning_02_a}. We focus here on cases where the conversion occurs only in the core of the NS, while the outer parts remain unchanged. The theory of dense nuclear matter predicts that such compact stars [hereafter we will use the term ``hybrid quark star'' (HQS) when referring to these objects] would generally be more compact than the progenitor standard NS, and their equilibrium radii would be smaller by up to 20\\%. The potential energy $ W $ released by the phase transition is expected to be of order \\begin{displaymath} \\label{eqq:energy_release} \\Delta W \\sim \\frac{M^2}{R} \\frac{\\Delta R}{R} \\sim 10^{52} \\mathrm{\\ erg}, \\label{eq:DW} \\end{displaymath} where $ R $ and $ M $ are the typical radius and mass of the NS, respectively, and $ \\Delta R $ is the decrease in radius.  A first-order phase transition is expected to be the most interesting case as far as the dynamics and structure of the star are concerned. Such a transition would proceed by the conversion of initially metastable hadronic matter in the core into the new deconfined quark phase \\citep{zdunik_07_a, zdunik_08_a}. The metastable phase could be formed as the central density of the NS increases due to mass accretion, spin-down or cooling. This could happen soon after the birth of the NS in a supernova or it could occur for an old NS accreting from a binary companion. The first of these channels is likely to give the higher event rate but the second is also interesting. In the widely accepted scenario for the formation of millisecond pulsars, where an old NS is spun up by accretion from a binary companion, the amount of spin-up would be directly related to the amount of mass accreted \\citep{burderi_99_a}, meaning that there would be a population of rapidly rotating NSs with rather high mass (and hence high central densities). Recent observational data seems to confirm that a significant proportion of millisecond pulsars do indeed tend to be high-mass objects \\citep{freire_08_a} which would then be candidates for undergoing (or having undergone) a phase-transition-induced collapse of the type which we are discussing here. While some significant proportion of the potential energy release given by Eq.~(\\ref{eq:DW}) would probably go into neutrino emission, a significant proportion might also go into pulsations of the newly-formed HQS, if the conversion process to the new phase is rapid enough, and this could be an interesting source of gravitational waves \\citep[GWs; see][]{marranghello_02_a, miniutti_03_a, lin_06_a}. If detected, these GW signals, and in particular the identification of quasi-normal mode frequencies in their spectrum, could help to constrain the properties of matter at the high densities encountered here. For non-rotating cold NSs with various compositions, the related theory of asteroseismology has already been formulated in recent years \\citep{andersson_98_a, kokkotas_01_a, benhar_04_a}.  A problem for making any detailed studies of phase-transition-induced collapse is that the description of the physics of the phase transition remains very uncertain and controversial \\citep[for a recent review see][]{horvath_07_a}. \\citet{drago_07_a} discussed possible modes of burning of hadronic matter into quark matter in the framework of relativistic hydrodynamics using a microphysical equation of state (EoS). They found that the conversion process always corresponds to a deflagration and never to a detonation. They also argued that hydrodynamical instabilities can develop at the burning front. They estimated the corresponding increase in the propagation velocity of the phase transition and noticed that, although the increase is significant, it is not sufficient to transform the deflagration into a detonation in essentially all realistic scenarios. On the other hand, \\citet{bhattacharyya_06_a} considered the transition as being a two-step process, in which the hadronic matter is first converted to two-flavour $ u $ and $ d $ quark matter, which then subsequently transforms into strange quark matter ($ u $, $ d $ and $ s $ quarks) in a second step. They used the relativistic hydrodynamic equations to calculate the propagation velocity of the first front and found that, in this first stage, a detonation wave develops in the hadronic matter. After this front passes through, leaving behind two-flavour matter, a second front is generated which transforms the two-flavour matter to three-flavour matter via weak interactions. The timescale for the second conversion is $ \\sim 100 \\mathrm{\\ s} $, while that for the first step is about $ 1 \\mathrm{\\ ms} $.  Against this background, \\citeauthor{lin_06_a} (\\citeyear{lin_06_a}; hereafter referred to as LCCS) carried out a first study of GW emission by a phase-transition-induced collapse of a rotating NS to a HQS, using a very simplified model for the microphysics and treating the phase transition as occurring instantaneously. On the basis of 3D calculations using Newtonian gravity and hydrodynamics, they obtained waveforms of the emitted GWs for several collapse models, and found that the typical predicted dimensionless GW strain $ h $ ranged from $ 3 $ to $ 15 \\times 10^{-23} $ for a source at a distance of $ 10 \\mathrm{\\ Mpc} $. The corresponding energy $ E_\\mathrm{gw} $ carried away by the GWs was found to be between $ 0.3 $ and $ 2.8 \\times 10^{51} \\mathrm{\\ ergs} $. They also determined the modes of stellar pulsation excited by the collapse and showed that the spectrum of the emitted GWs was dominated by the fundamental quasi-radial and quadrupolar pulsation modes of the final star. They suggested that the damping of the stellar pulsations observed in their calculations was due to the production of shock waves leading to the development of differential rotation, which proceeds on a timescale of about $ 5 \\mathrm{\\ ms} $ for typical collapse models.  The study by LCCS treated the conversion process as being instantaneous, and this was mimicked by replacing, at the initial time, the EoS describing the hadronic nuclear matter in the core (with which the initial equilibrium model had been built) by one describing a central zone of deconfined quarks surrounded by a region of mixed phase. The material outward of this remained in the original hadronic phase. In a first-order phase transition that proceeds via a detonation, the conversion front propagates supersonically with respect to the matter ahead of it. The sound speed in the stellar interior typically has a value of the order of $ 0.3 $ to $ 0.5\\,c$. Assuming that the quark-matter core has a radius of about $ R \\sim 5 \\mathrm{\\ km} $, a rough estimate for the timescale of a supersonic conversion gives $ \\tau \\sim R / v \\sim 0.05 \\mathrm{\\ ms} $. Clearly, this value is not much smaller than the dynamical timescale of a NS and so one does not know, a priori, the impact that the properties of the conversion would have for the subsequent stellar dynamics taking place on similar timescales. Therefore, one should first check how the dynamics of the star after the phase transition would depend on the finite propagation velocity of the front, and we do this here.  LCCS represented the hadronic matter by means of a polytropic EoS, initially having an adiabatic index $ \\gamma = 2 $. When the phase transition was triggered (by changing the EoS in the central regions, reducing the pressure support), they also replaced the original hadronic EoS by a softer ideal-gas type of EoS (with $ \\gamma $ then ranging from $ 1.95 $ to $ 1.75 $ depending on the model), which artificially lowered the pressure \\emph{outside} the deconfined quark matter core as well [see their Eq.~(44)]. The lowering of the adiabatic index also in the outer regions leads to an increased release of gravitational binding energy and is hard to motivate on physical grounds.  Although the study by LCCS was an important step forward, it was clearly using an extremely simplified model for the physics of the NS matter and of the phase transition (as well as not including emission processes other than GW emission) and one must question how closely it represents such processes occurring in real NSs. Since the full problem is complex and involves some input physics which is very incompletely known at present, the only way to make progress at this stage is, indeed, to work in terms of simplified models but making improvements where possible. The strategy of this paper is to take the work of LCCS as a starting point and then to take some steps forward in the direction of including further aspects. In particular, the effects of general relativity (GR) certainly play an important role in such studies: for instance, the total rest mass of a typical Newtonian NS may be almost twice as large as that of its counterpart in GR with the same central density and the same EoS. The impact of GR is even more pronounced if the star is near to the maximum mass limit. Therefore, one expects the Newtonian and GR descriptions of NS collapse to differ significantly, thus necessitating a proper treatment of the GR effects.  In this study we extend the previous work of LCCS in a number of ways: (i) we take into account the effects of GR; (ii) we modify the EoS of the stellar matter only in regions which undergo the phase transition and take care to avoid introducing spurious convection; (iii) we consider a larger set of physical models for different values of the stellar parameters and different properties of the stellar matter; (iv) we study the effect of introducing a finite timescale for the initial phase transformation which destabilises the core rather than treating it as occurring instantaneously. Our model for the HQS is based on that of LCCS. The numerical hydrodynamics simulations are here performed in axisymmetry using the conformal flatness approximation to GR. While our model remains extremely simplified, we believe that these modifications do represent valid and significant steps forward. We feel, however, that we should caution against making further elaborate extensions of the hydrodynamics (e.g.\\ to 3D GR hydrodynamics or GR magneto-hydrodynamics) until such time as one has the possibility of including a greatly improved treatment of the microphysics and emission processes so that the models come closer to those of real NSs.  This article is organized as follows: in Section~\\ref{section:numerics} we summarize the numerical methods used, while in Section~\\ref{section:models} we introduce the models investigated. In Section~\\ref{subsection:comparison} we validate our code by performing a comparison with the Newtonian models of LCCS. In Section~\\ref{section:gr_collapse_models} we discuss the results of our simulations of GR models and their dependence on the various parameters, and in Section~\\ref{section:conclusion} we conclude with a summary. In Appendix~\\ref{appendix:conversion_velocity} we examine the impact on the collapse dynamics of considering a finite timescale for the initial phase transformation, and in Appendix~\\ref{appendix:damping_times} we present our method for determining the damping times for the emitted gravitational radiation waveforms.  Unless otherwise noted, we choose geometrical units for all physical quantities by setting the speed of light and the gravitational constant to one, $ c = G = 1 $. Latin indices run from 1 to 3, Greek indices from 0 to 3. ",
        "conclusions": "\\label{section:conclusion}  In this paper, we have described numerical simulations of the phase-transition-induced collapse of rotating neutron stars to become hybrid quark stars. The simulations were made using a code which solves the general relativistic hydrodynamic equations in axisymmetry and within the conformally flat approximation. The initial stellar models were taken as being rapidly rotating $ \\gamma = 2 $ polytropes, while during the evolution we used an equation of state composed of three parts, depending on density: a normal hadronic-matter phase, a mixed phase of normal matter together with deconfined quark matter, and a pure quark-matter phase. The hadronic component of the equation of state was described with an ideal-gas type of equation of state, while for the deconfined quark matter phase we used the MIT bag model.  To validate our code, we first repeated several of the Newtonian simulations performed previously by \\citet{lin_06_a}. We found that the differential rotation which develops in these models during the post-bounce phase is almost entirely due to strong transient convection which arises because the way in which they treated the onset of the phase transition leads to an artificial negative specific-entropy gradient. We argue that their conclusion about there being a causal link between the large amplitude post-bounce pulsations and the creation of differential rotation was a misinterpretation of the results. We also suggest that a significant part of the damping of the pulsations which they observed was a consequence of the numerical dissipation present in their calculations, rather than being related to differential rotation, although a part of this damping was related to the other physical processes which we have discussed here and which were present in both calculations.  Having clarified the dynamics of the collapse in the Newtonian framework, we then investigated the corresponding situation within general relativity, using a modified prescription for triggering the phase transition and initiating the collapse. Here we change the equation of state only in the regions where the phase transition has taken place and leave it unchanged elsewhere. We recognize that our procedure for describing the phase transition remains extremely idealized but we believe that it represents a step forward.  Despite this difference in the way in which the phase transition is triggered, we have not found any major qualitative differences in the waveforms produced when comparing the relativistic simulations with the earlier Newtonian ones. Also in the general relativistic simulations, the waveforms produced are mainly composed of the fundamental $ l = 0 $ quasi-radial $ F $-mode and the fundamental $ l = 2 $ quadrupolar $ {}^{2\\!}f $-mode. However, in contrast to the Newtonian models, the $ F $-mode is at a lower frequency than the $ {}^{2\\!}f $-mode as a consequence of the different density profiles. In addition to these modes, a nonlinear self-coupling of the $ F $-mode at twice the original frequency, the $ 2 \\cdot F $-mode, is strongly excited due to the violent nature of the collapse. Although qualitatively similar to their Newtonian counterparts, the relativistic models exhibit quantitative differences in their dynamics. In order to investigate these, we have considered a set of 23 different models organized in several sequences. In each of these sequences, only one of the characteristic parameters of the models was allowed to vary.  The main trends observed were as follows. For the sequence with constant initial central rest-mass density, the maximum gravitational wave strain $ |h|_\\mathrm{max} $ increases monotonically with the rotation rate (except for some models where the waveform is strongly altered by mode resonances). For the constant rest mass sequence, on the other hand, we observe first an increase of $ |h|_\\mathrm{max} $ with the rotation rate and then a decrease for very rapid rotation. For the sequence with constant rotation period but varying rest mass, we see $ |h|_\\mathrm{max} $ increasing monotonically with the rest mass, which is a different behaviour from that seen for the Newtonian models, where $ |h|_\\mathrm{max} $ first increases and then decreases again as the rest mass is increased. The reason for this difference may simply be that having a decreasing part of the curve would require progenitor neutron stars with rest masses beyond the upper limit for general-relativistic models. Finally, changing the equation of state in the mixed phase has straightforward consequences: the central rest-mass density at bounce, the amplitude of the post-bounce oscillations, and the maximum gravitational wave strain all increase as the overall pressure in the mixed phase is reduced.  Other points to arise include the following. Firstly, the influence of rotation on the frequencies of the $ F $ and $ {}^{2\\!}f $-modes agrees well with what was found by \\citet{dimmelmeier_06_a} for pulsations of equilibrium neutron stars, suggesting that studies of linear pulsations of equilibrium models can (at least in some cases) correctly predict the properties of the eigenfrequencies also when the pulsations are excited in a dynamical situation. Secondly, in some models the emission of gravitational radiation is considerably enhanced by the nonlinear resonance between the $ 2 \\cdot F $-mode and the $ {}^{2\\!}f $-mode. When the frequencies of these two modes are sufficiently close, the weakly radiating quasi-radial $ 2 \\cdot F $-mode transfers some of its kinetic energy to the strongly radiating quadrupolar $ {}^{2\\!}f $-mode, leading to a considerable increase in $ |h|_\\mathrm{max} $. Thirdly, we have proposed a simple explanation for the strong damping of the post-bounce pulsations seen for a sub-set of our models. Our analysis reveals that that these models are all both rotating close to the Kepler limit and also undergoing large-amplitude post-bounce pulsations, resulting in significant mass shedding from the stellar surface close to the equator. As already discussed by \\citet{stergioulas_04_a}, this ejection of loosely bound matter is very efficient in the damping quasi-radial pulsations.  In order to assess the prospects for the detection of phase-transition-induced collapse events by gravitational wave interferometers, special attention has been paid to making an accurate calculation of the gravitational wave emission resulting from this scenario. We find that the dimensionless gravitational wave strain $ h $ from a source at a distance of $ 10 \\mathrm{\\ Mpc} $ ranges between $ 0.1 $ to $ 2.4 \\times 10^{-23} $ for all of the models considered and that the total energy emitted in gravitational waves during first $ 50 \\mathrm{\\ ms} $ of evolution is between $ 10^{-6} $ and $ 10^{-4} M_\\odot c^2 $, corresponding to $ 2 \\times 10^{48} $ and $ 2 \\times 10^{50} \\mathrm{\\ erg}$, respectively. These numbers are considerably smaller than those presented by \\citet{lin_06_a} for their Newtonian calculations and so are disappointing for the prospects of detecting these sources. The damping times for the post-bounce oscillations, as computed from the gravitational radiation waveform, range from $ 8 $ to $ 687 \\mathrm{\\ ms} $ for the $ F $-mode, and from $ 18 $ to $ 130 \\mathrm{\\ ms} $ for the $ {}^{2\\!}f $-mode. For all of the models considered, we have also calculated the characteristic frequency $ f_\\mathrm{c} $, the characteristic strain $ h_\\mathrm{c} $, and the signal-to-noise ratio for current and future detectors. For current detectors such as LIGO or VIRGO, all of the models (except one) have a signal-to-noise ratio above $ 1 $ for a source at $ 10 \\mathrm{\\ kpc} $. For the advanced LIGO detector, the signal-to-noise ratio rises to well above $ 10 $ for a source at $ 10 \\mathrm{\\ kpc} $ for all of the models and reaches almost $ 400 $ when mode resonance is active. However, for detecting sources within the Virgo cluster at a distance of $ 20 \\mathrm{\\ Mpc} $, which is probably necessary in order to reach a realistic event rate, third generation detectors would be needed.  In conclusion, we note that while our study represents an improvement over previous ones, it still lacks a number of very important aspects which would be necessary for a properly realistic modelling of these objects. Firstly, there is the treatment of the phase transformation itself which remains extremely crude, containing no detailed picture of the transformation of the material, the local heat input or the emission of neutrinos or photons. A consistent treatment of radiative transfer is likely to be essential for following the cooling phase of the newly formed hybrid quark star and could highlight that the radiative losses would produce differences in the specific-entropy stratification and hence trigger real convection. Also, in our discussion, we have been considering the phase transition as occurring by means of a detonation; the conclusions would be drastically altered if it takes place instead via a slow deflagration. A second aspect concerns the treatment of the standard neutron-star matter: while using a gamma-law equation of state to model this can be reasonable for some simplified calculations, it gives an extremely poor approximation to the complex physics actually occurring in real neutron stars. Thirdly, we have not been considering the influence of magnetic fields which are not only expected to affect the dynamics, but could also lead to a dramatic modification of the phase transition process itself \\citep{lugones_02_a}. It is clear that future studies will need to take these aspects into account."
    },
    "0806/0806.3019_arXiv.txt": {
        "abstract": "This paper summaries the status of a large project to improve distance scales of various classes of variable stars. This is being carried out by a large group in Cape Town, Japan, England and the USA. The results are illustrated by giving the distances to the Large Magellanic Cloud and the Galactic Centre ($R_{o}$) as well as the value of the Hubble Constant, $H_{o}$, based on our current results. The classes of variables considered are; Classical Cepheids, Type II Cepheids, RR Lyrae stars, O- and C- type Miras. ",
        "introduction": "In studying the structure and dynamics of our own Galaxy as well as the distances and distributions of galaxies in the nearby Universe we need accurate distance tracers. These tracers should be relatively common, easily recognizable, intrinsically luminous and of well defined absolute magnitudes. Certain classes of variable stars fit all these criteria. It is well known, for instance that the brightnesses of classical Cepheids in nearby galaxies are used to calibrate the luminosities of supernovae of type Ia, and that this leads to a determination of the Hubble constant ($H_{o}$). But other classes of variable stars are becoming of increasing importance in setting distance scales. This is not only because they provide, or should provide, a confirmation of the classical Cepheid scale when applied to the same object (e.g. the LMC). There are many systems in which classical Cepheids are not present (e.g. globular clusters, the Galactic bulge, dwarf spheroidal galaxies etc.) and for which other classes of variable stars can provide distances. In addition, a knowledge of the luminosities of variable stars is essential if we are to understand their structure and evolution.  A group of us have been interested for some years in the calibration of the luminosities of variable stars. The present paper is a short review of the current status of this work. The principal workers in Cape Town, in recent years, have been; Patricia Whitelock, Dave Laney, John Menzies and the writer; In Japan, Noriyuki Matsunaga and his many colleagues; In England, Floor van Leeuwen; and in the United States, Fritz Benedict and a large group of colleagues as well as Tom Kinman. A full list of co-workers will be found in the references at the end of this paper. ",
        "conclusions": "The above discussion shows that the distance scales we obtain for the various classes of variable stars are quite consistent with one another. Of particular interest is that this scale leads to a modulus of the LMC somewhat smaller than conventionally assumed. The results also lead to a distance to the Galactic Centre which is also smaller than that adopted in recent times but it agrees well with the value derived from the motion of a star round the central blackhole. Work is in progress to further strengthen these distance scales."
    },
    "0806/0806.4914_arXiv.txt": {
        "abstract": "We calculated bulk viscosity due to non-equilibrium weak processes in superfluid nucleon-hyperon matter of neutron stars. For that, the dissipative relativistic hydrodynamics, formulated in paper \\cite{gusakov07} for superfluid mixtures, was extended to the case when both nucleons and hyperons are superfluid. It was demonstrated that in the most general case (when neutrons, protons, $\\Lambda$, and $\\Sigma^{-}$ hyperons are superfluid), non-equilibrium weak processes generate {\\it sixteen} bulk viscosity coefficients, with only {\\it three} of them being independent. In addition, we corrected an inaccuracy in a widely used formula for the bulk viscosity of non-superfluid nucleon-hyperon matter. ",
        "introduction": "\\label{1} It is well known (see, e.g., \\cite{Lindblom95,ak01,andersson03,andersson06}), that neutron stars can be unstable with respect to the emission of gravitational waves. The matter in a pulsating neutron star is not (even locally) in chemical equilibrium. The relaxation towards chemical equilibrium is accompanied by the dissipation of pulsation energy. This process is one of the most important dissipative processes, suppressing the growth of gravitational-wave instability. It can be described by the introduction of an effective bulk viscosity in the hydrodynamic equations (see, e.g., \\cite{ll87}).  The bulk viscosity due to non-equilibrium weak processes in neutron stars was calculated by a number of authors (for a review see, e.g., \\cite{dsw07}). It strongly depends on the composition of stellar matter. In this paper we consider the matter of inner layers of neutron stars, composed of electrons, muons, neutrons, protons, $\\Lambda$, and $\\Sigma^{-}$ hyperons (nucleon-hyperon matter). A calculation of the bulk viscosity for nucleon-hyperon matter is complicated by the fact that baryons in such matter can be superfluid \\cite{yls99,ls01,yp04,bb98}.  The bulk viscosity of superfluid matter was calculated in a number of papers (see, e.g., \\cite{hly00,hly01,hly02,lo02,no06,ssr06}). In these papers only one bulk viscosity coefficient was studied, analogous to that in non-superfluid hydrodynamics. The effects of superfluidity were taken into account only in calculating the reaction rates. However, it is known that there are several bulk viscosity coefficients in hydrodynamics of superfluid liquid \\cite{putterman74,ll87,khalatnikov89,gusakov07}.  In a recent paper \\cite{gusakov07} {\\it four} bulk viscosity coefficients were calculated for the matter composed of superfluid neutrons, superfluid protons and electrons. It was shown that taking into account three additional bulk viscosity coefficients results in a significant decrease of characteristic damping times of sound modes (approximately, by a factor of three).  In this paper we extend the results of Ref. \\cite{gusakov07} to the case of nucleon-hyperon matter. In particular, we demonstrate how the dissipative hydrodynamics \\cite{gusakov07} should be modified to describe a possible presence of superfluid hyperons. Next, we show that in the most general case, when baryons of any species are superfluid, non-equilibrium processes of mutual transformations of particles generate {\\it sixteen} bulk viscosity coefficients, only three of them being independent. In addition, we correct an inaccuracy in the expression for the bulk viscosity of non-superfluid nucleon-hyperon matter made in Ref. \\cite{lo02} and spreaded widely in the literature (see, e.g., \\cite{vanD04,drago05,mohit06,debarati06,pan06,debarati07}). We calculate the bulk viscosity correctly and compare our results with those of Ref. \\cite{lo02}.  The paper is organized as follows. In section II we calculate the bulk viscosity of non-superfluid nucleon-hyperon matter. In section III we calculate and analyze all {\\it sixteen} bulk viscosity coefficients describing dissipation in superfluid nucleon-hyperon mixtures; the relations between these coefficients are also discussed. Section IV presents a summary. ",
        "conclusions": "In this paper we have analyzed the bulk viscosity due to non-equilibrium particle transformations in superfluid nucleon-hyperon matter of neutron stars. Our approach is similar to that used in Ref.\\ \\cite{gusakov07} for the case of superfluid {\\it npe} matter.  Our main results are as follows:  ({\\it i}) We have demonstrated that the expression for the bulk viscosity of normal (non-superfluid) nucleon-hyperon matter, widely used in the literature, is inaccurate. We have presented the correct derivation of the bulk viscosity and compared it with the results of Ref.\\ \\cite{lo02}. Numerically, both formulae give almost similar results (within a few tens of percent).  ({\\it ii}) We have extended the hydrodynamics of superfluid mixtures reported in Refs.\\ \\cite{ga06,gusakov07} to allow for a possible presence of superfluid hyperons. We have determined the general form of dissipative terms entering the equations of this hydrodynamics and showed that generally (when all baryon species are superfluid), it contains {\\it sixteen} bulk viscosity coefficients.  ({\\it iii}) We have calculated and analyzed the sixteen bulk viscosity coefficients assuming they are generated by non-equilibrium reactions (\\ref{s})--(\\ref{ll}) of particle mutual transformations. We have shown that only {\\it three} of them are independent. All other 13 bulk viscosities can be expressed through these three using Eqs.\\ (\\ref{onsager11}) and (\\ref{Sp})--(\\ref{additional2}). In addition, we have demonstrated that Eq.\\ (\\ref{xi22}) for the bulk viscosity coefficient $\\xi_2$ formally coincides with Eq.\\ (\\ref{ksi}) for that in normal matter (however, the reaction rates $\\lambda_l$ ($l=1,\\ldots,4$) are affected by superfluidity).  Our results can be important for the studies of dynamical instabilities in pulsating superfluid neutron stars, especially for the studies of the r-mode instability. They can also be important for modeling of the thermal evolution of pulsating neutron stars and for analyzing rotochemical and gravitochemical heating of millisecond pulsars with superfluid nucleon-hyperon cores (for non-superfluid $npe$ matter of neutron stars, this problem was considered in Refs.\\ \\cite{reisenegger95, fr05, rjfk06, jrf06}).  Let us notice, that to start such an analysis one needs to know the matrix $Y_{ik}$, which is the most important ingredient in the hydrodynamics of superfluid nucleon-hyperon mixture. To our best knowledge, it has not been considered in the literature so far. We are planning to fill this gap and present its extensive calculations in a subsequent publication."
    },
    "0806/0806.1522_arXiv.txt": {
        "abstract": "We present new results characterizing cosmological shocks within adaptive mesh refinement N-Body/hydrodynamic simulations that are used to predict non-thermal components of large-scale structure.  This represents the first study of shocks using adaptive mesh refinement.  We propose a modified algorithm for finding shocks from those used on unigrid simulations that reduces the shock frequency of low Mach number shocks by a factor of $\\sim3$.  We then apply our new technique to a large, $(512~Mpc/h)^3$, cosmological volume and study the shock Mach number~($\\Mach$) distribution as a function of pre-shock temperature, density, and redshift.  Because of the large volume of the simulation, we have superb statistics that results from having thousands of galaxy clusters.  We find that the Mach number evolution can be interpreted as a method to visualize large-scale structure formation. Shocks with $\\Mach<5$ typically trace mergers and complex flows, while $5<\\Mach<20$ and $\\Mach>20$ generally follow accretion onto filaments and galaxy clusters, respectively.  By applying results from nonlinear diffusive shock acceleration models using the first-order Fermi process, we calculate the amount of kinetic energy that is converted into cosmic ray protons.  The acceleration of cosmic ray protons is large enough that in order to use galaxy clusters as cosmological probes, the dynamic response of the gas to the cosmic rays must be included in future numerical simulations. ",
        "introduction": "What determines the thermal history of galaxy clusters?  On large scales, it is governed by the in-fall of material onto dark matter halos and the conversion of gravitational potential energy into thermal energy.  This process occurs through the heating of the gas via strong accretion shocks surrounding galaxy clusters and filaments \\citep{Ryu:2003aa, Miniati:2001aa, Pfrommer:2006aa, Pavlidou:2006aa}.  Once inside collapsed structures, complex flows associated with the merging of subhalos continue to create moderate-strength shocks that allow the halos to virialize.  Because of this, shocks encode information about structure formation and its thermal effects on the gas.  Cosmological shocks affect three important realms of structure formation and leave feedback on the surrounding medium.  First, shocks thermalize the incoming gas, providing much of the pressure support in baryons.  This process is the basis upon which clusters are able to virialize.  Additionally, the thermalization of gas at the standing accretion shocks surrounding large-scale filaments produces the warm-hot intercluster medium (WHIM) at temperatures of $10^5 K - 10^7 K$ \\citep{Dave:1999aa,Cen:1999aa}.  The history of the mass flux through these shocks describes the evolution of gas in the WHIM phase \\citep{Pfrommer:2008aa}.  Second, the strength of the outer accretion shocks onto halos, characterized by the Mach number, is dependent upon the mass of the gravitating object.  This is because the higher mass generates larger acceleration and velocity in the diffuse gas while the sound speed of the upstream gas is uniform for previously unshocked gas, corresponding to a temperature of $\\approx10^4K$.  This temperature floor is created by the reionization from stars.  Thus, the Mach number of accretion shocks can be used as an independent measure of cluster mass.  This could conceivably be a powerful new tool for cluster mass estimation if we are able to observe the accretion shock with radio observations \\citep[e.g.][]{Giacintucci:2008aa}.  Finally, because these shocks are collisionless features whose interactions in the hot plasma are mediated by electromagnetic fields, it is possible for a portion of the thermal distribution of particles to be accelerated and transformed into non-thermal populations of cosmic rays (CRs) through the process of diffusive shock acceleration \\citep[DSA; e.g.][]{Drury:1986aa,Blandford:1987aa}.  This process results in a fraction of the kinetic energy of shocking gas being converted into both thermal and non-thermal components \\citep{Kang:2002aa,Kang:2005aa,Kang:2007aa}.  The cosmic ray electron populations are likely sources of radio halos and radio relics in galaxy clusters \\citep{Pfrommer:2008aa,Kim:1989aa, Giovannini:1993aa}, while the cosmic ray protons may be sources of $\\gamma$-ray emission through their interactions with gas protons.  If a significant portion of the gas pressure resides in cosmic rays, then it will likely affect gas mass fraction estimates as well as the assumption of hydrostatic equilibrium.  Because of the importance of these mass estimates in measuring dark energy, we must include the underlying physics in order to perform precision cosmology \\citep{Allen:2008aa}.  To date, studies of cosmological shocks have included observational, theoretical, and numerical techniques.  Observationally, the majority of the work surrounding cosmic shocks are related to radio relics, of which only a few have been studied in depth \\citep[e.g.][]{Rottgering:1997aa, Clarke:2006aa, Orru:2007aa}.  Using the spectral index of the non-thermal particle distribution, we can infer a Mach number if the acceleration is due to first order Fermi acceleration \\citep{Giacintucci:2008aa}.  Additionally, GLAST will begin observing $\\gamma$-rays and will likely see signatures from galaxy clusters due to hadronic cosmic ray interactions with pions \\citep{Pfrommer:2008ab}.  On the theoretical side, the majority of analyses are based upon manipulating the Press-Schechter formalism \\citep{Press:1974aa,Sheth:1999aa} to deduce first the mass function of accreting objects and then their interactions with infalling material.  \\citet{Pavlidou:2006aa} extended these analyses to calculate the energy and mass flux through accretion shocks.  Furthermore, several analytical attempts have been made to describe merger shocks, including those by \\citet{Fujita:2001aa} and \\citet{Gabici:2003aa,Gabici:2003ab}.  However, it is quite difficult to account for the complex morphologies that arise during structure formation using purely analytical frameworks.  For this reason, multiple numerical techniques have been developed using hydrodynamical simulations.  There have been numerical studies of shocks using both Eulerian ``single-grid\" codes \\citep[e.g.][]{Ryu:2003aa, Kang:2007aa} as well as smoothed particle hydrodynamics (SPH) codes \\citep{Pfrommer:2006aa,Pfrommer:2007aa,Pfrommer:2008aa,Pfrommer:2008ab}.  There are advantages and disadvantages of both methods.  In previous work using grid-based codes, shocks were analyzed during post-processing by examining temperature jumps for a given point in time \\citep{Ryu:2003aa}.  However, it was impossible to cover the spatial dynamic range needed to describe both the complex flow within halos and their coupling to large scale structures because of the use of a uniform grid.  Therefore, even the largest simulations, with $1024^3$ cells in a $100 h^{-1} Mpc$ volume, has a comoving spatial resolution of only $97.7 h^{-1} kpc$ \\citep{Ryu:2003aa}.  There have been recent attempts at prescribing hybrid models to study turbulent generation, but full resolution convergence of the results have still not yet been achieved \\citep{Ryu:2008aa}.  The advantage of a grid code is its superb shock-capturing algorithms that do not rely on the use of artificial viscosity when using higher-order methods \\citep{OShea:2005ab}.  In contrast, SPH codes are implicitly adaptive in space due to the Lagrangian nature of the method, e.g. high density regions are resolved by a larger number of particles than low-density regions. This approach conserves hydrodynamic quantities exactly when they are advected with the flow. However, because the properties of the gas are determined by a weighted average over neighboring particles, formally discontinuous shocks are spread over a length determined by the smoothing length. Additionally, SPH relies on artificial viscosity to dissipate flows and produce the correct amount of entropy. Because of these restrictions, \\citet{Pfrommer:2006aa} developed a method that is able to identify shocks by examining the time-evolution of the entropy of individual SPH particles. Comparing the instantaneous entropy injection rate to the characteristic time it takes a particle to cross the broadened shock surface, they are able to identify and calculate the instantaneous Mach number of shocks while remembering the pre-shock conditions. Therefore, the analysis can be performed on-the-fly and shock quantities were traced along with the usual hydrodynamic properties.  These calculations use calibrations against ``shock tube\" simulations to derive the correct relationship between entropy injection rate and Mach number, which may vary with respect to different artificial viscosity implementations.  To address all of the problems listed above, we have developed a novel numerical algorithm capable of detecting and identifying shocks in the 3-D adaptive mesh refinement (AMR) grid-based code, $Enzo$.  The use of AMR allows us to analyze unprecedented dynamic ranges with an advanced hydrodynamic code that is able to capture shocks exceedingly well.  In this work we explore simulations with dynamic range of up to $2^{16}=65,536$, but we are not limited from going further in future work.  Because of the complexity of the structure of AMR simulations, it was necessary to develop several new numerical algorithms to identify the shocks.  This shock-finding analysis algorithm will be presented and compared to previous methods \\citep[e.g.][]{Ryu:2003aa, Pfrommer:2006aa}.  In order to validate and quantify the robustness of our method, we carry out a resolution study that includes both mass and spatial resolutions that vary by factors of $16$ and $64$, respectively.  We also propose a new method of characterizing shocks by their pre-shock overdensity and temperature.  This then allows analysis that goes beyond the traditional internal vs. external (of filaments/clusters) shocks classification suggested by \\citet{Ryu:2003aa}.  By refining the temperature and density ranges examined, we are able to identify shocks that reside in voids, filaments, and halos.  Additionally, by using temperature cuts, we can identify the population of gas that is being shocked into the warm-hot intercluster medium (WHIM).  After calculating the shock structure in a given simulation, we are able to compute the amount of shock kinetic energy that is transferred to high-energy cosmic ray protons through diffusive shock acceleration.  While the surface area of the large scale structure shocks are dominated by low pre-shock temperature and density, the bulk of the cosmic ray energy generation occurs in the centers of collapsed structures.  Since stronger shocks will produce harder spectra \\citep{Blandford:1987aa}, we expect that the strong accretion shocks could be the source of high energy cosmic rays.  In Section 2 we describe the numerical methods used for both the cosmological simulations as well as the analysis of the shock-finding process.  In Section 3, we compare our algorithm to that of \\citet{Ryu:2003aa} and test it using 3-D ``shock tube\" tests.  Section 4 contains the main results of analyzing a large, $(512~Mpc/h)^3$, cosmological simulation with a peak spatial resolution of $7.8\\ kpc/h$. Section 5 describes the effects of spatial and mass resolution on the shock populations and cosmic ray acceleration.  In Section 6, we discuss the limitations of our analysis, and in Section 7 we summarize our findings and discuss potential future directions. ",
        "conclusions": "Our study of cosmological shocks has resulted in several advances in both scientific understanding and numerical algorithms.  We now summarize the key findings: \\begin{itemize} \\item We have developed a novel numerical scheme that is capable of detecting and accurately characterizing the Mach number of shocks in an adaptive mesh refinement simulation.  This method has relaxed the previous restriction of using a coordinated axis-based approach and now allows us to accurately characterize shocks that have any orientation with respect to the coordinate grid.  \\item Using our new shock-finding technique on a unigrid cosmological simulation that is identical to the highest-resolution calculation in \\citet{Ryu:2003aa}, we have shown that previous methods resulted in an overestimate of the number low Mach number shocks by a factor of $\\sim3$ due to confusion of the direction of shock propagation, and that this underestimate is consistent with using shock-finding algorithms that only sweep along coordinate axes.  \\item We have analyzed the largest AMR cosmological simulation to date that includes adiabatic gas physics, the ``Santa Fe Light Cone.\"  This simulation has an effective spatial dynamic range of $65,536$ and resolves both large scale structure and small-scale features within galaxy clusters.  Whereas previous studies were able to study on the order of 10 high mass clusters, we have thousands within a single simulation volume.  Our study of this simulation has led to a new technique for conceptualizing structure formation because we are able to analyze the evolution of three different populations of shocks: cluster accretion, filament accretion, and internal merger and complex flow shocks.  \\item  By applying the results of 1-D Diffusive Shock Acceleration models, we calculate the amount of kinetic energy at shock fronts that is used to accelerate cosmic rays, and find it to be in agreement with previous studies.  These cosmic rays will make up a significant fraction of the total pressure in the intracluster medium and therefore their dynamical effects need to be studied.  \\item We have performed a resolution study that varies both the dark matter particle mass and peak spatial resolution.  From the results of this study, we believe that the spatial refinement in the ``Santa Fe Light Cone\" simulation is adequate.  The state of mass resolution convergence is less clear, suggesting that for future studies we should focus on higher mass resolution.  \\end{itemize}  While our numerical technique of characterizing shocks has been proven to be robust, our results are still somewhat limited by the physics.  We have not yet included potentially important effects such as radiative cooling, star formation and feedback, AGN heating, or a photoionizing UV background.  These physics will be included in future work.  Second, our results are based upon a post-processing of the simulation output.  Ideally, the shocks would be identified in an ``on-the-fly\" manner during simulation runtime.  Additionally, the cosmic ray acceleration would be traced in a self-consistent manner that allowed for a back-reaction on the gas.  Attempts at tracing the cosmic ray pressure have been made by \\citet{Pfrommer:2006aa,Pfrommer:2007aa} using an SPH code, and we will be working towards the same goal in the near future within $Enzo$.  Finally, the acceleration of cosmic rays is still dependent on the underlying magnetic field strength and orientation.  Cosmological MHD has been implemented within $Enzo$, and in the near future we will include magnetic fields and their coupling to cosmic rays within a cosmological AMR volume."
    },
    "0806/0806.1714_arXiv.txt": {
        "abstract": "Starting in the 1970s, the Arecibo 305m telescope has made seminal contributions in the field of extragalactic spectroscopy. With the Gregorian upgrade completed in the late  1990s, the telescope acquired a field of view. Population of that field of view with a seven--feed array at L--band (ALFA) increased by nearly one order of magnitude its survey speed. As a result, much of the extragalactic astronomy time of the telescope is now allocated to survey projects, which are briefly discussed. The next technical development stage for the 305m telescope is foreseen as that of a 40 beam system that would take advantage of phased array technology: {\\bf AO40}. This would further speed up the survey performance of the telescope. It is shown how the figure of merit for survey speed of AO40 would be comparable with that of SKA--precursor facilities, planned for operation in the next decade. A number of scientifically desirable new surveys that would become possible with AO40 are briefly discussed. ",
        "introduction": "The 21 cm line of neutral Hydrogen was first detected in 1953 in a galaxy other than our own. The Arecibo Observatory entered operations a decade later, but it took another decade and the replacement of the 305m telescope primary mirror with a finer surface, for the field of extragalactic spectroscopy to become established at that facility. In the following years, the Arecibo telecope became a protagonist in the field of extragalactic spectroscopy, with fundamental contributions to the understanding of galactic structure, the predominance of dark matter in galaxies and the case of nature vs. nurture in their evolution; the large scale structure of the Universe, its basic parameters and peculiar velocity field. In the 1990s, a second upgrade of the 305m telescope corrected for the spherical aberration of its primary, delivering an extended field of view at its focus. The exploitation of that field of view by populating it with multifeed arrays opened a new era at Arecibo: that of surveys which would cover large fractions of the 13,000 square degrees of the accessible sky. Currently about to start or already underway, extragalactic spectroscopy surveys that use the ALFA feed array will deliver results of unprecedented scope and impact in the field. A progressively larger fraction of telescope time is allocated to survey work. ",
        "conclusions": ""
    },
    "0806/0806.3461_arXiv.txt": {
        "abstract": "The abundance and distribution of collapsed objects such as galaxy clusters will become an important tool to investigate the nature of dark energy and dark matter. Number counts of very massive objects are sensitive not only to the equation of state of dark energy, which parametrizes the smooth component of its pressure, but also to the sound speed of dark energy as well, which determines the amount of pressure in inhomogeneous and collapsed structures. Since the evolution of these structures must be followed well into the nonlinear regime, and a fully relativistic framework for this regime does not exist yet, we compare two approximate schemes: the widely used spherical collapse model, and the pseudo-Newtonian approach. We show that both approximation schemes convey identical equations for the density contrast, when the pressure perturbation of dark energy is parametrized in terms of an effective sound speed. We also make a comparison of these approximate approaches to general relativity in the linearized regime, which lends some support to the approximations. ",
        "introduction": "We now have overwhelming evidence that the Universe is accelerating, possibly under the influence of some type of negative-pressure substance -- dark energy (DE) \\cite{Frieman:2008sn,Albrecht:2006um,Seljak:2006bg}. However, even though DE may be directly responsible for this enhanced expansion, it is widely believed that the direct impact of perturbations in DE density and pressure on structure formation is very weak. This is strictly correct only for a cosmological constant model of DE, which does not have perturbations.  For most scalar field models of DE, this component remains very homogeneous even on galaxy and cluster scales. Heuristically, this can be understood as follows. In these models the scalar field can not have relaxed to its minimum energy state and one must require that the time scale for the variation of the field is longer than the Hubble time, implying a very flat potential. Therefore, the scalar field must be extraordinarily light,  $m < H_0$, where $H_0$ is the Hubble parameter today. The mass of the scalar field sets the scale for its spatial variation and hence one usually expects small perturbations in the scalar field for scales $\\lambda < 1/m$ (the Compton wavelength), which are of the order of the Hubble radius. However, this argument may not apply to more general models of dark energy.     If our only concern is the evolution of the background, then the role of dark energy in the evolution of dark matter perturbations is completely determined by its equation of state $w=p_e/\\rho_e$, where $p_e$ is the homogeneous pressure and $\\rho_e$ is the homogeneous energy density of dark energy \\cite{Peebles:2002gy,Padmanabhan:2003gd,Sahni:2004ai}. At this level, dark energy affects structure formation indirectly because, as it starts to dominate the background, very large structures are ripped apart by the ensuing accelerated expansion \\cite{Linder:2005in,Liberato:2006un}.   However, dark energy can influence structure formation in an additional manner. If it is a dynamical field or fluid, then dark energy must possess inhomogeneities, and these perturbations will interact gravitationally both with themselves and with clumps of dark matter \\cite{Coble:1996te}. This means that, unless dark energy is just a cosmological constant, it will both feel and create local gravitational potentials.  Although the effect of these inhomogeneities in the dark energy component becomes small as $w \\rightarrow -1$, in many models with $w \\neq -1$ it can be non-negligible when evolved in the nonlinear regime \\cite{Basilakos:2003bi, Klypin:2003ug, Linder:2003dr, Lokas:2003cj, Kuhlen:2004rw, Mota:2004pa, Nunes:2004wn, Nunes:2005fn, Horellou:2005qc, Manera:2005ct, Dutta:2006pn, Abramo:2007iu, Mota:2007zn, Abramo:2007mv}. Since the effects of dark energy perturbations on the cosmic microwave background are quite small (see, e.g. \\cite{Abramo:2004ji}), structure formation is the only remaining probe of the nature of dark energy on small and intermediate scales.  Nevertheless, a fully relativistic method to treat nonlinear perturbations is not available. When there is a pressure ingredient the nonlinear relativistic equations take a very complicated form. The Lema\\^itre-Tolman-Bondi (LTB) model \\cite{Lemaitre:1933gd,Tolman:1934za,Bondi:1947av} is the closest one can get to a working formalism, but it only works if matter is pressureless. The problem is not with the gravity side of the equations, but with the nonlinear evolution of matter and the relativistic treatment of pressure.  In this respect, the only well-studied models with inhomogeneous dark energy are those involving canonical scalar fields \\cite{Mota:2004pa, Nunes:2004wn, Nunes:2005fn, Horellou:2005qc, Manera:2005ct, Dutta:2006pn}, for which the equations of motion and the pressure follow directly from a given Lagrangian. For these models the free parameters are the scalar potential and some set of initial conditions. In this approach, the equation of state, the density perturbations and the pressure perturbations are derived quantities. Hence, a more kinematical  and model-independent approach to structure formation, closer in spirit to the homogeneous description of dark energy in terms of a parametrized equation of state $w(z)$, is sorely lacking.   There are two very different approximations to full-blown general relativity that have been frequently used. They are the spherical collapse (SC) model \\cite{Gunn:1972sv, Padmanabhan, Fosalba:1997tn} and the pseudo-Newtonian (PN) approach \\cite{Lima:1996at, Hwang:1997xt, Reis:2003fs, Hwang:2005xt}. We have recently used these approximations in the nonlinear regime in order to show that the ``effective equation of state'' of dark energy inside a collapsed region could be very different from its background value \\cite{Abramo:2007mv}.   In this work we show that, even though the underlying assumptions for either approach are rather different, they yield exactly the same nonperturbative equations as long as the pressure perturbations are treated in the same way. They also have an important advantage: they allow for a completely parametrized approach to dark energy. Furthermore, we compare the growth of perturbations in the linear regime with a linearized relativistic analysis and show that they are similar, lending support to the approximations.  This paper is organized as follows. In Section 2 we review both the PN and SC approaches and show that they are equivalent. In Section 3 we study the linear evolution of perturbations in DE in this approximation. The linear evolution of perturbations in a universe with a 2 component fluid is studied in general relativity in Section 4. We present a comparison between the relativistic analysis and the approximate analysis in the linear regime in Section 5. Section 6 concludes. ",
        "conclusions": "The study of perturbations in dark energy has received a great deal of attention recently. DE perturbations have the potential to alter the process of large scale structure formation in the universe. The existence of DE perturbations can in principle be tested in future surveys such as the Dark Energy Survey (DES) \\cite{Abbott:2005bi} and EUCLID \\cite{Euclid:2008eu}. These future observations may help to distinguish among different models of DE.  Structure formation occurs during the nonlinear stages of the evolution of perturbations. Unfortunately, there is no rigorous analytical description of this nonlinear stage in full GR. Aproximation methods must be used. Possibly the most trusted method is N-body simulations, but due to its very intensive computing requirements, it is not practical when one wants to study different models. Furthermore, N-body simulations employ newtonian physics and do not allow for the possibility of DE fluctuations.  In this paper we study two different approximation schemes, namely the Spherical Collapse and Pseudo-Newtonian approaches. The advantage of these schemes is that DE can be fully characterized by 2 functions: the equation-of-state parameter $w(z)$ and the effective speed of sound $c_{{\\rm eff}}(z)$. We show that, under a minimal set of assumptions, it is possible to \\ translate one approach into the other, rendering them completely equivalent. In order to compare these approximations with GR, we study perturbations in the linearized regime with all approaches. When the assumptions about the pressure perturbations are the same both in GR and PN/SC we find that the fluctuactions present the same qualitative behaviour, lending support to the approximations. However, in order to establish more firmly the validity of the approximations in the nonlinear regime, a comparison should be made with some nonperturbative model in GR, such as an extended LTB class of models, including fluids with pressure. Work along this direction is in progress."
    },
    "0806/0806.4594_arXiv.txt": {
        "abstract": "We study higher curvature corrections to the scalar spectral index, the tensor spectral index, the tensor-to-scalar ratio, and the polarization of gravitational waves. We find that the higher curvature corrections can not be negligible in the dynamics of the scalar field, although they are energetically negligible. Indeed, it turns out that the tensor-to-scalar ratio could be enhanced and the tensor spectral index could be blue due to the Gauss-Bonnet term. We estimate the degree of circular polarization  of gravitational waves generated during the  slow-roll inflation. We argue that the circular polarization can be observable with the help both of the Gauss-Bonnet and parity violating terms. We also present several examples to reveal observational implications of higher curvature corrections for chaotic inflationary models. ",
        "introduction": "According to the Wilkinson Microwave Anisotropy Probe (WMAP) data~\\cite{Komatsu08} and other observations, it seems  incontestable that the large scale structure of the universe stems from quantum fluctuations during the slow-roll inflation. Considering the accuracy of future observations, we expect higher order corrections to the simplest inflation model become important. Hence, it would be worth to study the corrections due to higher curvature terms to the slow-roll inflation. To the leading order, the corrections come from the Gauss-Bonnet and the axion-like parity violating coupling terms~\\cite{Weinberg08}. In fact, the same type of corrections can be derived from a particular superstring theory~\\cite{Antoniadis:1993jc}. At present, however, we have not fully understood such corrections based on superstring theory. Hence, we have to resort to the phenomenological approach. We hope the understanding of higher curvature corrections to the inflationary scenario provides a clue to the study of superstring theory.  There have been many works concerning the Gauss-Bonnet and the parity violating corrections to the Einstein gravity. It is  known that Gauss-Bonnet term can provide non-singular cosmological solutions~\\cite{Antoniadis:1993jc, Kawai:1998bn} and also provide alternative explanation for dark-energy~\\cite{Nojiri:2005vv,neupane06a,Neupane08,neupane06b}. There are many works on scalar and tensor perturbations in Gauss-Bonnet cosmology~\\cite{Kawai:1999xn, Kawai:1999pw, Hwang:1999gf, Cartier:2001is, Kawai:1998ab, Soda:1998tr, Kawai:1997mf,Gasperini:1997up,Cartier:2001gc,Leith:2007bu,Guo07, Vazquez:2008wb}. The parity violating term has been considered in conjunction with the  circular polarization of gravitational waves~\\cite{Lue:1998mq, Choi:1999zy, Alexander:2004wk, Lyth:2005jf, Satoh08}. However, since no one has systematically investigated  both terms simultaneously in the context of the slow-roll inflation, the combined effect on the slow-roll inflation is not so clear. Moreover, observational implications have not been discussed so far. Hence, a systematic study of higher curvature corrections to slow-roll inflationary scenario is desired.  In this paper,  we study the slow-roll inflation with higher curvature corrections and derive the formula  for the observables in terms of slow-roll parameters. It is stressed that the corrections could be relevant to observables in spite of the fact that its energetic contribution can be negligible. The reason is as follows. The Gauss-Bonnet term behaves as an effective  potential term in the scalar field equation, though it does not contribute to the Friedmann equation in the slow-roll approximation. It reminds us of a hybrid model. That is why  the effect could alter the tensor-to-scalar ratio. There is also an effect of the parity violating term on the gravitational waves as the leading effect. In fact, the circular polarization will be produced during the slow-roll inflation, which never occurs in the conventional inflationary scenario.  The organization of the paper is as follows: In section II, we define slow-roll parameters in the inflationary scenario with higher curvature corrections. In section III, we calculate the scalar spectral index. In section IV, we obtain the tensor spectral index and the tensor-to-scalar ratio. Here, we show the possibility of the enhancement of the tensor-to-scalar ratio due to the Gauss-Bonnet term. We also estimate the degree of the circular polarization of the primordial gravitational waves. In section V, the concrete examples are presented as an illustration. The final section is devoted to the conclusion. ",
        "conclusions": "We have studied the slow-roll inflationary scenario with the leading corrections, namely, the Gauss-Bonnet and the parity violating terms. We have obtained the higher curvature corrections to the scalar spectral index, the tensor spectral index, and the tensor-to-scalar ratio. We found that the higher curvature corrections can not be negligible in the dynamics of the scalar field, although they are negligible in the Friedmann equation. It turned out that the tensor-to-scalar ratio could be enhanced because of the modified dynamics of the scalar field. We have  found the tensor spectral index could be blue due to the Gauss-Bonnet term. We have also calculated the degree of circular polarization  of gravitational waves generated during the  slow-roll inflation. We discussed that the circular polarization can be observable due to the Gauss-Bonnet and parity violating terms. We also presented several examples to show observational implications of higher curvature corrections for chaotic inflationary models. We revealed that the observables  are sensitive to the higher curvature corrections. Interestingly, there are some models which shows the slowly climbing up of the potential. Although it is an exotic possibility, the phenomena arises within the slow-roll approximation. The climbing up inflationary model deserves further investigation.  We can extend the present analysis in various ways. Although we have considered only the scalar field, we could have incorporated  the vector field into the system~\\cite{Dimopoulos:2008rf,Yokoyama:2008xw,Kanno:2008gn}, which leads to more general higher derivative corrections. However, the most important issue is to determine the coupling functions based on the fundamental physics such as the superstring theory. We leave these issues for future work."
    },
    "0806/0806.1464_arXiv.txt": {
        "abstract": "The origin of the dichotomy of radio loudness among quasars can  be explained using recent findings that the mass of the central supermassive black hole (SMBH) in extended radio-loud quasars is systematically a few times that of their counterparts in radio-quiet quasars.  This sensitive dependence of radio jet ejection upon SMBH mass probably arises from the blockage of jets by the presence of substantial quantities of gas tidally stripped from stars by the central BH.  This disruptive gas, however, will only be available around BHs with masses less than $M_c ~\\gtrsim ~10^8\\Msun$, for which the tidal disruption radius lies outside the SMBH's event horizon.  Consequently, we find that AGN with $M_{BH} > M_c$ can successfully launch jets with a wide range of powers, thus producing radio-loud quasars.  The great majority of jets launched by less massive BHs, however, will be truncated in the vicinity of the SMBH due to mass loading from this stellar debris. This scenario also can naturally explain the remarkable dearth of extended radio structures in quasars showing broad absorption line spectra. ",
        "introduction": "The bi-modality in the distribution of the radio-to-optical flux ratio, $R$, was first revealed % using optically selected samples (e.g., Kellermann et al. 1989). A similar picture emerged from the plot of radio flux versus [O III] line intensity for QSOs (the latter marks the accretion power of the central engine), which showed that radio loud quasars (RLQs) are typically $\\sim 10^4$ times more radio luminous than radio quiet quasars (RQQs) at a given $L_{[O III]}$ (e.g., Miller, Rawlings \\& Saunders 1990;  Xu et al.\\ 1999). However, a somewhat ambiguous view of this dichotomy emerged from more recent investigations based on deep, large sky area radio surveys (FIRST, Becker et al.\\  1995; NVSS, Condon et al.\\ 1998) and extensive optical surveys (SDSS, York et al.\\  2000; 2dF, Croom et al.\\ 2001). Thus, while a radio loud fraction of $5-10\\%$ was estimated in some studies (e.g., Ivezi{\\'c} et al.\\ 2002, 2004), other papers reported no radio-loudness bimodality (White et al.\\ 2000; Cirasuolo et al.\\ 2003). Laor (2003) has argued that these negative results are probably an artifact of the lower sensitivity of the FIRST radio survey to extended radio emission, which, unlike the core emission, is unbeamed, and hence a more reliable indicator of the jet power. In fact, a double peaked distribution of $R$, is seen in an ``image stacking\" analysis of the FIRST data for QSOs (White et al.\\ 2007).  It is unlikely that the host galaxy morphology or the large-scale environment play a crucial role in the radio loudness dichotomy since in both these respects, optically bright RQQs are very similar to RLQs (Metcalf \\& Magliocchetti 2006, hereafter MM, and references therein). Therefore, it has been suggested that the main difference between RLQs and RQQs lies in the availability of SMBH spin energy to power the jets in RLQs (e.g., % Wilson \\& Colbert 1995; Hughes \\& Blandford 2003). In the scheme of Sikora, Stawarz \\& Lasota (2007), radio powerful jets, driven by the SMBH spin, are realized intermittently whenever a strong collimation is externally imposed on them by the MHD wind from the accretion disk (see, also, Nipoti et al., 2005).  A key observational clue  for the radio  dichotomy may stem from the inferred mass-related duality of radio loudness, in that powerful radio sources are only associated with galactic nuclei containing central black holes of masses $M_{BH} > 10^8 M_{\\sun}$ (Laor 2003; Dunlop et al.\\ 2003; Chiaberge, Capetti \\& Macchetto 2005; Sikora et al.\\ 2007). Note that this clue does not exclude the spin of the black hole or accretion disk from being the basic mechanism underlying % jet formation. However, for radio detection the jet must be able to emerge from the nuclear region out to parsec scale. An ``aborted\" jet scenario recently invoked by Ghisellini, Haardt \\& Matt (2004) envisions that radio-quiet AGN too eject relativistic particle jets but they are terminated due to the impact of infalling shells of ``heavy\" material that are ejected by the AGN intermittently. A likely case of an intermittently aborted jet comes from the X-ray monitoring of the RQQ PG 1407+265 (Gallo 2006) which is believed to possess a highly Doppler boosted compact radio jet (Blundell et al.\\ 2003).  While nuclear jet abortion may offer an attractive route to explaining radio quietness, can it be reconciled with the {\\it dichotomy} of radio loudness of quasars, which itself appears to be sensitively linked to $M_{BH}$? Specifically, it needs to be understood how the jet abortion scenario ties in with the observed systematic mass difference between the central SMBHs of RLQs and RQQs and the finding that all RLQs have $M_{BH} > 10^8 M_{\\sun}$ (Laor 2003; see above). Several recent studies using different techniques and based on large-sky radio and optical surveys have revealed a consistent mass excess of the central black-holes powering RLQs vis-a-vis those in RQQs (Dunlop et al.\\ 2003; MM; Jarvis \\& McLure 2006; Hyv{\\\"o}nen et al.\\ 2007). The determinations of $M_{BH}$ in these studies are based both on optical spectroscopy of the nuclei and accurate photometry of host galaxies using the Hubble Space Telescope. Moreover, care was taken  that the samples of RQQs and RLQs are matched in redshift and optical luminosity. Thus, even though it is only by a modest factor (1.4 to 4), the presence of a mass excess of the BHs in RLQs over those in RQQs has been found in independent studies.  Clearly, if this is the main factor behind the radio dichotomy, the underlying physical process must be very sensitive to $M_{BH}$.  A possible link of the SMBH critical mass, $(M_c)$, to the radio dichotomy, via tidal disruptions of stars, was noted in Laor (2000) but discounted on the ground that any influence of the tidal debris would be too short lived as the debris would be quickly sucked into the SMBH on a dynamical time scale (sub-day). Below we argue that even the modest difference found between the BH masses in RLQs and RQQs could play a critical role in aborting the nascent jets in RQQs. ",
        "conclusions": ""
    },
    "0806/0806.2521_arXiv.txt": {
        "abstract": "We consider a cosmological scenario in which the expansion of the Universe is dominated by phantom dark energy and black holes which condense out of the latter component. The mass of  black holes decreases via Hawking evaporation and by accretion of phantom fluid but new black holes arise continuously whence the overall evolution can be rather complex. We study the corresponding dynamical system to unravel this evolution and single out scenarios where the big rip singularity does not occur. ",
        "introduction": "Phantom dark energy fields are characterized by violating the dominant energy condition, $\\rho + p > 0$. Thereby the conservation equation, $\\dot{\\rho} + 3 H (\\rho + p) = 0$, has the striking consequence that the energy density increases with expansion \\cite{caldwell1,caldwell2}. In the simplest case of a constant ratio $w \\equiv p/\\rho$ one has $\\rho \\propto a^{-3(1+w)}$, where $1+w < 0$, while the scale factor obeys $a(t) \\propto (t_{*} - t)^{-n}$ with $n = - 2/[3(1+w)]$ and $t \\leq t_{*}$, being $t_{*}$ the ``big rip\" time (at which the scale factor diverges). However, as we shall show below, the big rip may be avoided if black holes are produced out of the phantom fluid at a sufficiently high rate. One may think of three different mechanisms by which phantom energy goes to produce black holes.  $(i)$ In phantom dominated universes more and more energy is continuously pumped into any arbitrary spatial three-volume of size $R$. Hence, the latter will increase as $a$ while the mass, $M$, inside it will go up as as $a^{m+3}$ -with $m = 3\\,|1+w|$. Therefore, the ratio $M/R$ will augment with expansion whenever $m > -2$ (i.e., $w < -1/3$, dark energy in general). As a consequence of the energy being pumped faster than the volume can expand, the latter will eventually contain enough energy to become a black hole. This is bound to occur as soon as $ \\rho \\, R^{3} \\geq R/(2 G)$ (i.e., when $R \\geq (2G\\, \\rho)^{-1/2}$).  $(ii)$ As statistical mechanics tells us, equilibrium thermal fluctuations obey $<(\\delta E)^{2}> = k_{B} \\, C_{V} \\, T^{2}$, where $\\delta E = E - <E>$ is the energy fluctuation at any given point of the system around the mean value $<E>$, $C_{V}$ the system heat capacity at constant volume, $k_{B}$ Boltzmann's constant, and $< ... >$ denotes ensemble average \\cite{textbook}. As demonstrated in \\cite{db}, unlike normal matter,  dark energy gets hotter with expansion according to the law $T \\propto a^{-3w}$. The latter follows from integrating the temperature evolution equation $\\dot{T}/T = -3H (\\partial p/\\partial \\rho)_{n}$ which, on its turn, can be derived from Gibbs' equation $T \\, dS = d(\\rho/n) + p \\, d(1/n)$ and the condition that the entropy be a state function, i.e., $\\partial^{2}S/(\\partial T \\, \\partial n) = \\partial^{2}S/(\\partial n \\, \\partial T)$ -see \\cite{db} for details. Therefore, it is natural to expect that the phantom fluid will be very hot at the time when it starts to dominate the expansion. Since the phantom temperature grows unbounded so it will the energy fluctuations as well -a straightforward calculation yields $<(\\delta E)^{2}> \\propto a^{-3(1+2w)}$. Eventually they will be big enough to collapse and fall within their Schwarzschild radius. (This simple mechanism was applied to hot thermal radiation well before phantom energy was introduced \\cite{piran-wald}).  $(iii)$ As demonstrated by Gross {\\it et al.} \\cite{quantumtunnelling}  thermal radiation can give rise to a copious  production of black holes of mass $ T^{-1}$ whenever $T$ is sufficiently high. In this case, black holes nucleate because the small perturbations around the Schwarzschild instanton involve a negative mode. This produces an imaginary part of the free energy that can be interpreted as an instability of the hot radiation to quantum tunnel into black holes. A simpler (but less rigorous) treatment can be found in Ref. \\cite{Kapusta}. One may speculate that black holes might also come into existence by quantum tunnelling of hot phantom fluid similarly as hot radiation did in the very early Universe.  Notice that mechanisms $(i)$ and $(iii)$ differ from the conventional gravitational collapse at zero temperature -the latter does not produce black holes when the fluid has $w < -1/3$, see e.g. \\cite{cai-wang}.  Obviously, one may wonder whether if phantom can produce particles other than black holes. In actual fact, there is no reason why dark energy in general (not just phantom) should not be coupled to other forms of matter. However, its coupling to baryonic matter is highly constrained by measurements of local gravity \\cite{localgravity} but not so to dark matter. In any case, we will not consider such possibility because it would introduce an additional variable in our system of equations below (Eqs. (\\ref{bhconsvbis})- (\\ref{dotm1})) and would increase greatly the complexity of the analysis.   The number black holes in the first generation (assuming that all those initially formed arise simultaneously) will be $N_{i} = (a/R)_{i}^{3} = C\\,  a ^{3-9(1+w)/2}$ with $C = [2 G \\rho_{0} \\, a_{0}^{3(1+w)}]^{3/2}$ -here the subscript zero signals the instant at which phantom dark energy starts to overwhelmingly dominate all other forms of energy-, and we may, reasonably, expect that they will constitute a pressureless fluid. Later on, more black holes may be formed but if much of the phantom energy has gone into black holes, then the second generation will not come instantly (as there will be less available phantom energy than at $t = t_{i}$).  Further, the black holes will accrete phantom energy and lose mass at a rate $\\dot {M} = -16 \\pi M^{2} \\, \\dot{\\phi} ^{2}$ regardless of the phantom potential, $V(\\phi)$ \\cite{babichev}. (Notice that for scalar phantom fields $\\rho + p = - \\dot{\\phi} ^{2}$). Additional mass will be lost to Hawking radiation \\cite{hradiation}.  The system of equations governing this complex scenario is \\\\ \\be \\dot{\\rho}_{bh} + 3H \\rho_{bh} = \\Gamma \\, \\rho_{x} + 16 \\pi \\, n M^{2} \\, (1+w)  \\rho_{x} \\, - n\\, \\frac{\\alpha}{M^{2}}\\, , \\label{bhconsv} \\ee \\be \\dot{\\rho}_{x} + 3(1+w)H\\rho_{x} = - \\Gamma \\, \\rho_{x} - 16 \\pi \\, n M^{2} \\,  (1+w) \\rho_{x} \\, , \\label{phconsv} \\ee \\be \\dot{\\rho}_{\\gamma} + 4 H \\rho_{\\gamma} = n \\frac{\\alpha}{M^{2}}\\, , \\label{gammaconsv} \\ee \\be 3 H^{2} = \\rho_{x} + \\rho_{bh}+ \\rho_{\\gamma} \\, , \\label{friedmann} \\ee \\be \\dot{M} = 16 \\pi \\, M^{2} (1+w) \\rho_{x} - \\frac{\\alpha}{M^{2}} \\label{dotm} \\ee \\\\ (subscripts $bh$, $x$, and $\\gamma$ stand for black hole, phantom, and radiation components, respectively, and we use units in which $8\\pi G = c = \\hbar = 1$). Here, $\\Gamma =$ constant $>0$ denotes the rate of black hole formation, $n = N/a^{3}$ the number density of black holes, $\\alpha$ a positive constant, and $H = \\dot{a}/a$ the Hubble function. The first term in (\\ref{dotm}) corresponds to the mass loss rate of a single black hole via phantom accretion; the second term to spontaneous Hawking radiation \\cite{donpage}. For simplicity, we assume that the black holes solely emit relativistic particles (i.e., a fluid with equation of state $p_{\\gamma} = \\rho_{\\gamma}/3$). This explains Eq. (\\ref{gammaconsv}). The second term on the right of Eq. (\\ref{phconsv})  ensures the energy conservation of the overall phantom plus black hole fluid in the process of phantom accretion -the mass loss of black holes must go into phantom energy.  At first glance, depending on the values assumed by the different parameters, $\\Gamma$, $w$, and $\\alpha$, two very different outcomes seem possible: $(i)$ a big rip singularity, if the phantom energy eventually gets the upper hand; and $(ii)$ a quasi-equilibrium situation between the black hole and phantom fluids, if none of them comes to dominate the expansion -in this case the  big rip cannot be guaranteed since the overall equation of state will not be constant.   There are five equations and, seemingly, six unknowns, namely, ($\\rho_{bh}$, $\\rho_{x}$, $\\rho_{\\gamma}$, $H$, $n$ and $M$) but, in actual fact only five ($\\rho_{bh}$, $\\rho_{x}$, $\\rho_{\\gamma}$, $a$, and $M$) as $n$ can be written as $\\rho_{bh}/M$.  To make things easier, we shall further assume that the black holes are massive enough to safely neglect Hawking evaporation and that, initially, there was no radiation present. The latter assumption is not much unrealistic, the fast expansion redshifts away radiation and dust matter very quickly. (The black-hole fluid also has the equation of state of dust but, as said above, black holes are continuously created out of phantom). Thus, we can dispense with Eq. (\\ref{gammaconsv}) and the last term on the right hand side of Eqs. (\\ref{bhconsv}) and (\\ref{dotm}).  One may  wonder whether the black holes may coalesce leading to bigger black holes and produce a big amount of relativistic particles. We believe we can safely ignore this possibility since the fast expansion renders the chances of black holes encounters highly unlikely.  Accordingly, the system of equations reduces to \\\\ \\be \\dot{\\rho}_{bh} + 3H \\rho_{bh} = \\Gamma \\rho_{x} + 16 \\pi \\, (1+w) M \\, \\rho_{bh} \\, \\rho_{x}  \\, , \\label{bhconsvbis} \\ee \\be \\dot{\\rho}_{x} + 3(1+w)H \\rho_{x} = - \\Gamma \\rho_{x} - 16 \\pi \\, M (1+w) \\rho_{bh} \\, \\rho_{x}  \\, , \\label{phconsvbis} \\ee \\be 3 H^{2} = \\rho_{x} + \\rho_{bh} \\, , \\label{friedmann1} \\ee \\be \\dot{M} = 16 \\pi \\, M^{2} (1+w)\\rho_{x}\\,  . \\label{dotm1} \\ee \\\\ Thus,  we may choose the  three unknowns, $\\rho_{bh}$, $\\rho_{x}$ and $M$  -since $H$ is linked to $\\rho_{bh}$ and $\\rho_{x}$ by the constraint Eq. (\\ref{friedmann1}).  Let us assume a phantom dominated universe (i.e., no other energy component enters the picture initially). Sooner or later a first generation of black holes will arise; then, the question arises: ``will new black holes condense out of phantom (second generation) before the first generation practically disappears eaten by the phantom fluid and consequently, the Universe become forever dominated by a mixture of phantom and black holes, or the black holes will disappear before they can dispute the phantom the energy dominance of cosmic expansion?\"  To ascertain this we shall apply, in the next Section, the general theory of dynamical systems  \\cite{sansone} to the above set of equations (\\ref{bhconsvbis})--(\\ref{dotm1}). We shall  analyze the corresponding critical points at the finite region as well as at infinity. As we will see, due to the existence of a critical point in the finite region of the phase portrait (connected with the creation rate, $\\Gamma$, of new black holes), there exist solutions that instead of ending up at the big rip tend asymptotically to a Minkowski spacetime. Hence, in some cases, depending on the initial conditions, the big rip singularity can be avoided thanks to the formation of black holes out of the phantom fluid.  Before going any further, it is fair to say that, in reality, no one knows for certain if phantom fluids have a place in Nature: they may suffer from quantum instabilities \\cite{instability}, although certain phantom models based in low--energy  effective string theory may avoid them \\cite{piazza}. On the other hand, observationally they are slightly more favored than otherwise -though, admittedly, this support has dwindled away in the last couple of years. In view of this unsettled situation, we believe worthwhile to explore the main possible consequences its actual existence may bring about on cosmic evolution. ",
        "conclusions": "In this paper, we briefly raised the  point that black holes may be produced in phantom dominated universes  by three different mechanisms:   $(i)$ via energy accumulation in any given spatial three-volume, $(ii)$ gravitational collapse of huge thermal fluctuations, and $(iii)$ quantum tunnelling of very dense and hot phantom fluid \\cite{quantumtunnelling}. In this regard, phantom fluids might be viewed as ``black holes factories\". However, while these processes look rather plausible the corresponding calculations are pending.  Nevertheless, after accepting that black holes may be produced by any of the sketched mechanisms, we studied the dynamical system associated to this scenario. Our main finding is that because of the existence of a critical saddle point in the finite region of the plane $(\\rho_{x},\\rho_T)$ -given by Eqs. (\\ref{obvious})- the big rip singularity (a generic feature of phantom-dominated universes) is no longer unavoidable. This critical point lies in the region of positive densities as $ w < 0$. Moreover, since $\\rho_T \\geq \\rho_x$  and $\\rho_T/\\rho_x = - w $, this point is located in the physical region ($\\rho_T \\geq \\rho_x > 0$) when $w < - 1$, that is, for  phantom equations of state. \\par In the said plane $(\\rho_{x},\\rho_T)$, there is an attractor at the origin, also the finite critical point described above (a saddle), and two critical points at infinity: one situated at the axis $\\rho_T$ (a repeller), and another at the straight line $\\rho_T = \\rho_x$ (an attractor). The finite critical point, connected with $\\Gamma > 0$, is a divider between the solutions that go from $\\rho_x \\rightarrow \\infty$ to the origin and from $\\rho_x \\rightarrow \\infty$ to the big rip. Those solutions that pass through this critical point, depending on the initial conditions, can either go to the origin (i.e.,  implying that the big rip is avoided), or to the critical point at infinity with $\\rho_x = \\rho_T$ (i.e., big rip) -see top panel of Fig. \\ref{fig:phase1}. \\par We are now in conditions to fix more precisely the possible scenario described above. Let us consider the relations for the finite critical point, Eqs. (\\ref{obvious}). Bearing in mind that the Universe is nearly spatially flat ($\\Omega_{T} =1$), and conceding that this critical point may lie not very far from the present expansion era (the Universe began accelerating recently), we find $\\Gamma = - 3\\sqrt{3}\\, (1 +w)H_0$, and $\\Omega_{x0} = -w^{-1}$. The Wilkinson microwave anisotropy probe (WMAP) \\cite{wmap} data are consistent with $w \\simeq - 1.1$. Hence, for the big rip to be attained the present dark energy density parameter must fulfill $\\Omega_{x0}  \\gsim \\, 0.9$. Since WMAP indicates $\\Omega_{x0} \\simeq 0.7$, our Universe may well avoid the big rip singularity (modulo $w$ is really a constant). This also implies $\\Gamma \\sim 0.5 \\,H_0$. \\par Our system is a three--dimensional one and obviously there are other dimensions. In the plane $(\\rho_{x},M)$ all points at infinity are singular; projecting this onto the Poincar\\'{e} sphere we obtain, in that plane, a singular circle around the origin. But, since the origin is an attractor, all the points in this circle are repellers. \\par As said above, some solutions can evade the big rip. Obviously, the latter become unavoidable if $\\Gamma = 0$ (no black hole production), since in this case the critical point at the finite region coincides with the origin.  There are also other situations in which the big rip can be avoided. For instance, when phantom dark energy corresponds to the generalized Chaplygin gas proposed in \\cite{pedro}, or when wormholes intervene \\cite{jimenez}, or when the curvature scalar gets very large and quantum effects become dominant \\cite{elizalde}. Nevertheless, to the best of our knowledge, the present scenario was never considered in the literature.  By contrast, as noted by Barrow \\cite{jdb}, there are situations in which finite-time future singularities can arise even if the fluid filling the Universe obeys $\\rho > 0$ and $\\rho + 3p >0$, i.e., under very mild conditions. We do not consider them here.  Admittedly, it can be argued that in view of the various simplifying assumptions, our treatment is not much realistic. In the first place, the rate $\\Gamma$ is not expected to be a constant, it will likely vary with expansion and depend on quantities like $w$ and $M$. Secondly, we have implicitly considered that all black holes are formed simultaneously with the same mass -a flat spectrum. It would be more natural to assume the number of black holes produced varies with mass and time. Further, as noted earlier, black hole spontaneous radiance should be included. Clearly, these features ought to be incorporated in future, more realistic, treatments. Nonetheless, we believe this small, first, step may lead the way to more ambitious undertakings."
    },
    "0806/0806.1591_arXiv.txt": {
        "abstract": "We study the angular resolution of the gravitational wave detector LISA and show that numerical relativity can drastically improve the accuracy of position location for coalescing Super Massive Black Hole (SMBH) binaries. For systems with total redshifted mass above $10^7\\,  M_{\\odot}$, LISA will mainly see the merger and ring-down of the gravitational wave (GW) signal, which can now be computed numerically using the full Einstein equations. Using numerical waveforms that also include about ten GW cycles of inspiral, we improve inspiral-only position estimates by an order of magnitude. We show that LISA localizes half of all such systems at $z=1$ to better than 3 arcminutes and the best 20\\% to within one arcminute. This will give excellent prospects for identifying the host galaxy. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0806/0806.1072_arXiv.txt": {
        "abstract": "We study the effects of rotation on the neutrino emission from Population III (Pop III) stars by performing a series of two-dimensional rotational collapse simulations of Pop III stellar cores. Our results show that rotation enhances the neutrino luminosities and the average energies of emitted neutrinos. This is because the thermalized inner core, which is the dominant neutrino source from Pop III stars, can be enlarged, due to rotational flattening, enough to extend the inner core outside the neutrinospheres.  This is in sharp contrast to the case of spherical collapse, in which the case of inner core shrinks deeper inside the neutrinospheres before black hole formation, which hinders the efficient neutrino emission.  In the case of rotational core-collapse, the emitted neutrino energies are found to become larger in the vicinity near the pole than the ones near the equatorial plane.  These factors make the emergent neutrino spectrum broader and harder than the spherical collapse case. By computing the overall neutrino signals produced by the ensemble of individual rotating Pop III stars, we find that the amplitudes of the relic neutrinos, depending on their star formation rates, can dominate over the contributions from ordinary core-collapse supernovae below a few MeV.  A detection of this signal could be an important tool to probe star formation history in the early universe. ",
        "introduction": "One of the most important goals in modern cosmology is to understand how the first stars formed at the end of the dark ages, and how they transformed the initially simple, homogeneous universe into a state of ever increasing complexity \\citep[e.g.,][]{bark01,brom04,loeb06}.  The first stars, so called Population III (Pop III), are predicted to have been predominantly very massive with $M\\gtrsim 100M_\\odot$ \\citep[e.g.,][]{naka01,abel02, brom02b,yosh06,oshe07}.  During their evolutions, the central core is thought to experience an electron-positron pair creation instability after carbon burning, which reduces the thermal energy of the core and eventually triggers gravitational core-collapse.  For stellar masses less than $\\sim 260 M_\\odot$, rapid nuclear burning releases a large amount of energy sufficient to entirely disrupt the star as pair-instability supernovae.  More massive stars, which also encounter pair-instability, are so tightly bound that the fusion of oxygen is unable to reverse infall.  Such stars are thought to collapse into black holes (BHs) \\citep{bond84,fryer01}, which we investigate in this paper.  Information about the formation and evolution of the Pop III stars might be observationally obtained from reionization \\citep{alva06}, the infrared background \\citep{dwek05}, nucleosynthesis yields \\citep{heger02,iwam05}, gamma-ray bursts \\citep{schn02,brom06}, and gravitational waves \\citep{buon05,suwa07b}.  However at present, it is still unclear whether or not one can obtain clear information about Pop III stars directly from observations.  Recently, neutrinos, which should have been emitted in large numbers at the phase of core-collapse, are expected to be new eyes to unveil the mysteries of the first stars.  Since the direct detection is unlikely because of their large distances, a diffuse background as relic neutrinos could potentially be observable.  In an analogous context, the relic neutrino background from ordinary core-collapse supernovae have been studied elaborately \\citep{bisn84,tota96,kapl00,ando03,ando04}. Since the physics involved in relic neutrinos covers a very wide range of topics, e.g., cosmic star formation rate and the collapse dynamics, detecting the relic neutrinos or even setting limits on their flux are expected to give us quite useful and unique implications not only for the understanding of the collapse physics itself, but also for various fields of astrophysics and cosmology.  The detectability of the relic neutrinos from Pop III stars has been discussed by several papers. \\cite{iocco05} estimated the amplitudes of the relic neutrinos produced during the nuclear burning phases as well as the core collapse phase, assuming a baryonic fraction of Pop III stars of $10^{-3}$ and monochromatic progenitors with a mass of $300 M_\\odot$.  \\cite{daig05} investigated the relic neutrino background produced by an early burst of Pop III stars taking into account a ``normal mode'' of the star formation at low redshift.  The framework of hierarchical structure formation, on which their computation relies, is based on the cosmic star formation histories, which are constrained by the observed star formation rate at redshift $z\\lesssim 6$, the observed chemical abundances in damped Lyman $\\alpha$ absorbers and in the IGM, and allow for an early reionization of the Universe at $z\\sim 10-20$.  Regardless of the differences in the treatment of the star formation histories, these studies concluded that the typical energy of the relic neutrino background is lowered as low as MeV or sub-MeV due to the cosmological redshift so that the detection is out of the question with presently known experimental techniques. However it should be noted that in their studies, the neutrino flux and energy emitted from the Pop III stars are set by hand due to the lack of Pop III star collapse simulations.  For more precise estimates of relic neutrinos from Pop III stars, we have to perform the hydrodynamical simulations.  So far, there have been only a few simulations studying the gravitational collapse of Pop III stars with BH formation including appropriate microphysical treatment \\citep{fryer01, naka06,suwa07a}.  \\cite{naka06} performed a spherically symmetric general relativistic simulation in a wide range of masses ($300-10000 M_\\odot$), in which the state-of-the-art neutrino physics are taken into account.  Their detailed calculations revealed the properties of the emergent neutrino spectrum, and based on that, they estimated the relic neutrino background and found that the detection for the spherically collapsing Pop III stars is difficult for the currently operating detectors.  Meanwhile current wisdom tells us that very metal poor stars lose very small amounts of mass and angular momentum through radiatively driven stellar winds so that Pop III stars may collapse with large angular momentum \\citep{hege03a}.  So stellar rotation might vary the spectrum of emitted neutrinos and affect its potential observability, which is the motivation of this study.  \\cite{fryer01} investigated the collapse of a rotating Pop III star of mass $300 M_\\odot$ with gray neutrino transport simulations and discussed effects of rotation on the emitted luminosities.  They, however, investigated only one rotational model. Considering the uncertainty of the angular momentum distributions of the Pop III stars, a more broad study is required.  More recently, we implemented a series of two-dimensional magnetorotational core-collapse simulations of Population III stars \\citep{suwa07a}, in which our interest was placed on the MHD effects on the formation of the jets.  In this study, we focus on the effects of rotation on the neutrino emission during the core-collapse phase of Pop III stars. For this purpose, we carry out two-dimensional simulations of rotational core-collapse of Pop III stars, changing the initial strength of the angular momentum in a parametric manner.  By so doing, we systematically investigate the dependence of the luminosity, the average energy, and the spectrum of neutrinos upon the rotation strength.  We also pay attention to the anisotropy of the emergent neutrino energies, that is, the neutrino emission at polar direction and in the equatorial plane.  Using the spectrum of single Pop III stars, we calculate the number flux of the relic neutrino background by summing up the contribution from individual Pop III stars and discuss their detectability.  This paper is organized as follows: In the next section, we briefly introduce our numerical methods and the initial models.  In \\S 3 and \\S4, we present the numerical results. We start with the outline of the dynamics and neutrino emission in the case of spherical collapse (\\S\\ref{sec:sph}) and then move on to discuss the effects of rotation on the dynamics and the neutrino luminosity (\\S\\ref{sec:lumi}), the neutrino energies (\\S\\ref{sec:ene}) and the spectrum (\\S\\ref{sec:spe}).  \\S 5 is devoted to summary and discussion. ",
        "conclusions": "We performed a series of 2D hydrodynamic simulations of rotational collapse of Pop III stars. We computed 8 models, changing the initial strength of the angular momentum in a parametric manner. We studied systematically the deformation of neutrinospheres and investigated the degree of anisotropic neutrino emission from the point of view of the neutrino luminosity and the average energy, and their impacts on the neutrino spectrum. Then we found the following, \\begin{enumerate} \\item The neutrino luminosity increases for more rapidly rotating cores. This is because the inner core, which is the dominant neutrino source, is deformed due to the rotation and can extend outside the neutrinospheres.  In particular, $\\nu_X$ is most sensitive to the rotation because its neutrinosphere is located at the smallest radius and nearest to the inner core.  The total energy emitted by neutrinos also increases with the strength of rotation. \\item The average energy of emitted neutrinos gets larger as rotation becomes faster. In addition, the energy is larger in the polar direction than in the equatorial plane due to the rotation. The neutrino spectrum gets broader and harder for rotating models. As a result, the number spectrum of the relic neutrinos also gets harder so that the relic neutrinos from Pop III stars overwhelm those from ordinary core-collapse supernova below $\\sim 7$ MeV, provided a fraction $f_{III}=0.1$ of all baryons forms Pop III stars. \\end{enumerate}  As for our numerical computations, we shall address a few of our assumptions.  In this study, we mimicked the neutrino transfer by the leakage scheme.  Although the scheme is a radical simplification, we checked that in the case of spherical collapse, the obtained integral quantities such as the total energy emitted by neutrinos and their average energy are very close to the results by \\citet{naka06}, in which the Boltzmann neutrino transport was solved.  We ignored the angle dependence of the neutrino emission, assuming that the neutrino track is radial. This approximation will break down just before BH formation because the neutrinosphere becomes oblate and the neutrino emission would be highly aspherical. Nonetheless, the multi-angle and multi-dimensional neutrino transfer is still computationally prohibitive \\citep[recently reported by][in the framework of core-collapse supernova]{ott08}.  The angle dependence is expected to make the neutrino spectrum harder. This is because the neutrinos produced in the vicinity of the equatorial plane might be able to escape to the polar direction even if they cannot escape radially, thus it should be taken into account.  Furthermore the Newtonian simulation in this paper, albeit with the general relativistic corrections, is nothing but an idealized study to describe the dynamics with BH formation. It is by no means definitive.  In order to explore these phenomena in more detail will require fully general relativistic simulations \\citep[e.g.,][for BH formation of massive stars]{shib03,seki05,liu07} but with the microphysical ingredients at least at the same level of this study, which is beyond the scope of this paper.   Bearing these caveats in mind, we state some discussions and speculations about the relic neutrinos from Pop III stars.  In estimating the relic neutrino background, we employed a parameter $f_{III}$, the baryon fraction of Pop III stars.  Because this parameter is entirely unknown so far, we employed three values of $10^{-1}$, $10^{-3}$, and $10^{-5}$.  The largest value is obtained for the explanation of the infrared background excess by UV photons from Pop III stars \\citep{sant02,dwek05}. So our estimation should be interpreted as giving an upper bound.  As for the star formation history, we employed only the delta function one because the baryon fraction dominates the uncertainty and the other ones such as the evolution of star formation rate are minor corrections \\citep[see,][for the discussion of the dependence on the star formation history]{iocco05,daig05,naka06}. In addition, we assume that all Pop III stars have the same mass of $300 M_\\odot$ for simplicity. The dependence of the progenitor masses should be also clarified, although the qualitative features induced by rotation studied here will not change.   The predicted foregrounds for the relic neutrinos in the relevant energy range computed here are the neutrinos produced in the Sun and the anti-neutrinos from nuclear reactors.  These (anti-)neutrinos, however, could be removable in principle because the distribution of these foregrounds are not isotropic.  Other possible foregrounds, which are isotropic, are extra-galactic {\\it neutrinos} produced by thermonuclear reactions in stars \\citep{porc04}.  On the other hand, as for {\\it anti-neutrinos} the contribution of Pop III stars are dominant in the range below a few MeV.  Thus we speculate that the detection of this signature could be the first direct probe of this as yet unseen metal-free first generation of stars."
    },
    "0806/0806.1558_arXiv.txt": {
        "abstract": "Results of a multi-color study of the variability of the magnetic cataclysmic variable BY Cam are presented. The observations were obtained at the Korean 1.8m and Ukrainian 2.6m, 1.2m and 38-cm telescopes in 2003-2005, 56 observational runs cover 189 hours. The variations of the mean brightness in different colors are correlated with a slope dR/dV=1.29(4), where the number in brackets denotes the error estimates in the last digits. For individual runs, this slope is much smaller ranging from 0.98(3) to 1.24(3), with a mean value of 1.11(1). Near the maximum, the slope becomes smaller for some nights, indicating more \"blue\" spectral energy distribution, whereas the night-to-night variability has an \"infrared\" character. For the simultaneous UBVRI photometry, the slopes increase with wavelength from dU/dR=0.23(1) to dI/dR=1.18(1). Such wavelength dependence is opposite to that observed in non-magnetic cataclysmic variables, in an agreement to the model of  cyclotron emission. The principal component analysis shows two (with a third at the limit of detection) components of variablitity with different spectral energy distribution, which possibly correspond to different regions of emission. The scalegram analysis shows a highest peak corresponding to the 200-min spin variability, its quarter and to the 30-min and 8-min QPOs. The amplitudes of all these components are dependent on wavelength and luminosity state. The light curves were fitted by a statistically optimal trigonometrical polynomial (up to 4-th order) to take into account a 4-hump structure. The dependences of these parameters on the phase of the beat period and on mean brightness are discussed. The amplitude of spin variations increases with an increasing wavelength and with decreasing brightness. The linear ephemeris based on 46 mean minima for 2003--2005  is HJD $2453213.010(3)+0.137123(3)E.$ The extensive tables of the original observations and of results of analysis are published in an electronic form. The nearby star GSC 4081-1562 was found to be an eclipsing red variable. ",
        "introduction": "Cataclysmic variables are close binaries with a red dwarf filling its Roche lobe and a white dwarf. Accretion structure is strongly dependent on the characteristics of the flow as well as on the magnetic field of the compact star (see e.g. monographs by Warner [55] and Hellier [20]).  The object H0538+608 = BY Cam was discovered as an X-ray source by Forman et al. [18]. Remillard et al. [44] classified it as an AM Herculis-type system. Mason et al. [34, 35] provided near simultaneous polarimetry and phase-resolved spectroscopy and noticed extreme variations in the accretion geometry. Andronov and Fuhrmann [11] suggested that these variations are due to a small asynchronism between the spin and orbital rotation of the magnetic white dwarfs. So there are slow periodic changes of the magnetic axis of the white dwarf (and corresponding changes of the accretion structure) in respect to the red dwarf with a period of 14\\fd52 [50]. These changes are sometimes called \"idlings\" (or, alternately, \"scrollings\", \"slewings\"). In classical polars, the white dwarf is synchronized with the orbital motion, but the orientation of the magnetic axis may undergo cyclic changes in respect to the line of centers of the rotating cataclysmic binary system. Such changes were called \"swingings\" (Andronov [2]), contrary to \"idlings\" in the asynchronous polars caused by different orbital and spin periods of the white dwarf. With changing orientation of the magnetic axis in respect to the secondary (donor) star, also changes the location and other characteristics of accretion columns above the photosphere of the white dwarf.  From the cyclotron humps in the infrared spectra, Cropper et al. [17] estimated magnetic field strength of the white dwarf of 41MGs, which is larger than in AM Her itself.  This puts the object between relatively numerous groups of DQ Her -- type stars (or intermediate polars) with rapidly rotating white dwarf and the AM Her -- type stars (or classical polars) which are in synchronism [39]. Another asynchronous system was Nova Cyg 1975 = V1500 Cyg [49], thus the desynchronization may be caused by a Nova outburst. Two more recent asynchronous polars are known (V1432 Aql and CI Ind). Sometimes these stars are called as the BY Cam -- type stars (e.g. Mason [33], Honeycutt and Kafka [23]).  According to the last \"live\" electronic version of the catalogue (dated 01.09.2006) by Ritter and Kolb [46], the number of classical, asynchronous and intermediate polars is 77, 4 and 38 (from a total number of cataclysmic variables of 641), respectively. So asynchronous polars are relatively rare objects. To study BY Cam, Silber et al. [50] have organized the \"Noah Project\" with a large flow of data obtained during more than 40 nights. Mason et al. [36] have succeeded to explain contradictions between different periods using the model of alternately changing active poles of the white dwarf. In this paper, we report results from the \"Noah-2\" project initiated to study different types of variability of this interesting object. During the \"Noah-2\" project, we have obtained 56 nightly runs of observations, more than 40 nights initially planned for the \"Noah-1\" project, and more attention is paid to the color information. The extensive tables of individual observations and the model parameters are presented electronically. ",
        "conclusions": ""
    },
    "0806/0806.3077_arXiv.txt": {
        "abstract": "It has been recently shown that quantum  Boltzmann equations may be relevant for leptogenesis. Quantum effects, which lead to a time-dependent CP asymmetry, have been shown to be particularly important for resonant leptogenesis when the asymmetry is generated by the decay of two nearly degenerate states. In this work we investigate the impact of the use of quantum Boltzmann equations in the framework ``soft leptogenesis'' in which supersymmetry soft-breaking terms give a small mass splitting between the CP-even and CP-odd right-handed sneutrino states of a single generation and provide the CP-violating phase to generate the lepton asymmetry. ",
        "introduction": "The discovery of neutrino oscillations makes leptogenesis a very attractive solution to the baryon asymmetry problem \\cite{fy,leptoreview}. In the {\\sl standard} type I seesaw  framework \\cite{ss}, the singlet heavy neutrinos have  lepton number violating Majorana masses and when decay out of equilibrium produce dynamically a lepton asymmetry which  is partially converted into a baryon asymmetry due to fast sphaleron processes.  For a hierarchical spectrum of right-handed (RH) neutrinos, successful leptogenesis requires generically quite heavy singlet neutrino masses~\\cite{di}, of order $M>2.4 (0.4)\\times 10^9$~GeV for vanishing (thermal) initial neutrino densities~\\cite{di,Mbound} (although flavour effects \\cite{flavour1,flavour2,db2,oscar} and/or extended scenarios \\cite{db1,ma} may affect this limit). Low-energy supersymmetry can be invoked to naturally stabilize  the hierarchy between this new scale and the electroweak one. This, however, introduces a certain conflict between the gravitino bound on the reheat temperature and the thermal production of RH neutrinos \\cite{gravi}. A way out of this conflict is provided by resonant leptogenesis~\\cite{PU1,PU,west}. In this scenario RH neutrinos are nearly degenerate in mass which makes the self energy contributions to the  CP asymmetries resonantly enhanced and allowing leptogenesis to be possible at much lower temperatures.  Once supersymmetry has been introduced, leptogenesis is induced also in singlet sneutrino decays.  If supersymmetry is not broken, the order of magnitude of the asymmetry and the basic mechanism are the same as in the non-supersymmetric case. However, as shown in Refs.\\cite{soft1,soft2}, supersymmetry-breaking terms can play an important role in the lepton asymmetry generated in sneutrino decays because they induce effects which are essentially different from the neutrino ones.  In brief, soft supersymmetry-breaking terms involving the singlet sneutrinos remove the mass degeneracy between the two real sneutrino states of a single neutrino generation, and provide new sources of lepton number and CP violation. As a consequence, the mixing between the sneutrino states generates a CP asymmetry in their decays. At zero temperature and at lowest order in the soft supersymmetry-breaking  couplings, the asymmetries generated in the sneutrino decays into fermions and scalars cancel out. However, thermal effects break this cancellation and once they are included the asymmetry can be sizable. In particular it is large for a RH neutrino mass scale relatively low, in the range $10^{5}-10^{8}$ GeV, well below the reheat temperature limits, what solves the cosmological gravitino problem. This scenario has been termed ``soft leptogenesis'', since the soft terms and not flavour physics provide the necessary mass splitting and CP-violating phase.  In general, soft leptogenesis induced by CP violation in mixing as discussed above has the drawback that in order to generate enough asymmetry the lepton-violating  soft bilinear coupling, responsible for the sneutrino mass splitting, has to be unconventionally small~\\cite{soft1,soft2,ourflasoft}. In this case, as for the case of resonant leptogenesis, the sneutrino self energy contributions to the  CP asymmetries are resonantly enhanced \\footnote{ Considering the possibility of CP violation also in decay and in the interference of mixing and decay of the sneutrinos~\\cite{soft3}, as well as extended scenarios~\\cite{ourinvsoft,softothers} may alleviate the unconventionally-small-$B$ problem.}  Till recently, the dynamics of thermal leptogenesis (both for the standard see-saw case, as well as for the soft leptogenesis scenario) has been studied using  the  approach of classical Boltzmann equations (BE). The possibility of using quantum Boltzmann equations (QBE) was first discussed in Ref.~\\cite{qbebuch} and it has been recently derived in detail in Ref.~\\cite{riottoqbe1}. In  Ref.~\\cite{riottoqbe1}  QBE were obtained starting from  the  non-equilibrium quantum field theory based on the Closed Time-Path   formulation. They differ from the classical BE in that  they contain integrals over the past times unlike in  the classical kinetic theory in which the scattering terms do not include any integral over the past history of the system which is equivalent to assume that any collision in the plasma does not depend upon the previous ones. Quantitatively the most important consequence is that  the CP asymmetry acquires an additional time-dependent piece, with its value at a given instant depending  upon the previous history of the system.  If the time variation of the CP asymmetry is shorter than the relaxation time of the particles abundances, the solutions to the quantum and the classical Boltzmann equations are expected to differ only by terms of the order of the ratio of the time-scale of the CP asymmetry to the relaxation time-scale of the distribution. This is typically the case in thermal leptogenesis with hierarchical RH neutrinos . However, as discussed in Refs.~\\cite{riottoqbe2,riottoqbe3}, in the resonant leptogenesis scenario, $(M_j-M_i)$ is of the order of the decay rate of the RH neutrinos.  As a consequence the typical time-scale to build up coherently the time-dependent CP asymmetry, which is of the order of $(M_j-M_i)^{-1}$, can be larger than the time-scale for the change of the abundance of the RH neutrinos. This, as shown in Refs.~\\cite{riottoqbe2,riottoqbe3}, leads to quantitative differences between the classical and the quantum approach in the case of resonant leptogenesis and, in particular, in the weak washout regime they enhance the produced  asymmetry.  Motivated by these results and the fact that in soft leptogenesis the CP asymmetry is produced resonantly, we perform a detailed study of the role of quantum effects in the soft leptogenesis scenario.  Our results show that because of the thermal nature of soft leptogenesis, the dependence of the quantum effects on the washout regime for soft leptogenesis is quantitatively different than in the see-saw resonant scenario.  In particular in the weak washout regime quantum effects do not enhance but suppress the produced baryon asymmetry. Quantum effects are most quantitatively important for extremely degenerate sneutrinos (that is far away from the resonant condition), $\\Delta M \\ll \\Gamma_{\\widetilde N}$, and in the strong washout regime they can lead to an enhancement of the produced asymmetry,  as well as change of sign, of the produced asymmetry. But altogether, for a given $M$ the required values of the lepton violating soft bilinear term $B$ to achieve successful leptogenesis are not substantially modified. ",
        "conclusions": "In this article  we have performed a detailed study of the role of quantum effects in the soft leptogenesis scenario. We have studied the effects on the produced asymmetry as a function of the washout parameter $m_{eff}$ and the degeneracy parameter $R=2\\Delta M/\\Gamma_{\\widetilde N}$ Our results show that, because of the thermal nature of soft supersymmetry, the characteristic time for the building of the asymmetry is larger than in the see-saw resonant leptogenesis which leads to quantitative differences on the dependence of the effect on the washout regime between the two scenarios.  In particular, in the weak washout regime, quantum effects do not enhance but suppress the produced lepton asymmetry in soft leptogenesis. Quantum effects are most quantitatively important for extremely degenerate sneutrinos $\\Delta M \\ll \\Gamma_{\\widetilde N}$. In this case and in the strong washout regime quantum effects can enhance the absolute value of the produced asymmetry as well as induce a change of its sign. But altogether, our results show that the required values of the Majorana mass $M$ and the lepton violating soft bilinear coefficient $B$ to achieve successful leptogenesis are not substantially modified."
    },
    "0806/0806.4657_arXiv.txt": {
        "abstract": " ",
        "introduction": "Dynamical interaction between the SMC, the LMC, and the Galaxy have long been considered to be closely associated not only with the evolution  of the SMC (e.g., Yoshizawa \\& Noguchi 2003; YN03) but also with the formation of the MS (e.g., Putman et al. 1998) and the MB (e.g., Muller et al. 2004). The observed structural, kinematical, and chemical properties of stars, star clusters,  and gas in the SMC  have been extensively discussed in variously different aspects of the SMC, such as the star formation history (e.g., Russell \\& Dopita 1992;  Pagel \\& Tautvaisiene 1998; Piatti et al. 2007), the total dynamical mass (e.g.,  Dopita et al. 1985; Stanimirovi\\'c et al.  2004, S04), and the interaction history with the LMC (YN03; Bekki \\& Chiba 2007). One of the key, longstanding problems related to the SMC as a galaxy is to understand how the past dynamical interaction between the SMC, the LMC, and the galaxy could have influenced the chemical and dynamical evolution and the star formation history of the SMC (e.g., Westerlund 1997).    In order to solve the above mentioned key problem, we need to know not only three-dimensional (3D) space motions of the Magellanic Clouds (MCs) but also dynamical and chemical properties of stars with different ages and metallicities. Recent proper motion measurements of the MCs and numerical studies of the orbits of the MCs based on the proper motion measurements have suggested that the MCs can be  a unbound with each other and they are unbound to the Galaxy (Kallivayalil et al. 2006, K06; Besla et al. 2007). The two-dimensional (2D) distributions of stellar populations with different ages and metallicities derived from AGB stars have revealed star formation histories in different local regions of the SMC (Cioni et al. 2006). Observational studies of stellar populations of the SMC  based on their color-magnitude diagrams have revealed  not only the 2D structures dependent on ages and metallicities but also a more precise global star formation history of the SMC (e.g., Harris \\& Zaritsky  2004). Thus, these observations on 2D structures of different stellar populations, if compared with the corresponding theoretical studies, can provide a vital clue to the above-mentioned key question of the SMC.     Previous theoretical models, however, did not provide useful predictions on the 2D (or 3D) structures of stars with different ages and metallicities owing to the strong limitations in their numerical simulations. YN03 discussed the star formation history of the SMC and the origin of structural and kinematical properties of the MS. They however did not discuss chemical and dynamical properties of stars and gas {\\it within the SMC}. Although previous one-zone chemical evolution models discussed the long-term chemical evolution history of the SMC (e.g., Spite et al. 1988; Tsujimoto et al. 1995; Idiart et al. 2007), they did not  discuss dynamical properties of the SMC. Thus no theoretical studies have been done which attempt to investigate {\\it both dynamical and chemical properties}  of the SMC and thereby provide reasonable explanations for the above-mentioned observations. Given that only chemodynamical simulations can provide predictions on 3D-dynamical and chemical properties in a self-consistent manner, it is vital for theoretical studies to investigate the above-mentioned key observational properties based on chemodynamical simulations of the SMC.   The purpose of this paper is thus to numerically investigate physical properties of the SMC in a fully self-consistent manner and thereby try to provide answers for the above-mentioned unresolved problems. We mainly discuss the origin of structural, kinematical, and chemical properties of the SMC's stellar and gaseous components derived  in previous observations (e.g., Suntzeff et al. 1986;  Torres \\& Caranza 1987; Hardy et al. 1989; Russell \\& Dopita 1992; Hatzidimitiriou et al. 1993; S04). We also discuss the latest observations that previous theoretical works did not discuss:  the structure of the stellar halo (No\\\"el \\& Gallart 2007), kinematical properties of older stellar populations (e.g.,  Harris \\& Zarisky 2006), possibly purely gaseous streams  in the MB (Harris  2007), and bifurcated structures and kinematics of the MC (Matthews \\& Staveley-Smith 2007). We however do not intend to discuss some important aspects of the SMC evolution, such as the long-term evolution of the MCs and the formation of globular cluster (GCs), which were discussed to some extent in our previous papers (Bekki et al. 2004; Bekki \\& Chiba 2005; paper I, Bekki 2007) in the preset study. We plan to discuss these in our paper III (Bekki \\& Chiba 2008a) based on more sophisticated and self-consistent models not only for the MCs but also for the Galaxy.     The layout of this paper is as follows. In \\S 2, we summarize our numerical models used in the present study and describe the methods for analyzing structure and kinematics of the simulated SMC. In \\S 3, we present numerical results on the time evolution of morphology, metallicity distribution, and dynamical properties of the SMC. In \\S 4, we discuss the above-mentioned   outstanding issues related to formation and evolution of the SMC. The conclusions of the preset study are given in \\S 5. Since the origin of the MS has been already discussed extensively in previous numerical simulations (e.g.,  Gardiner \\& Noguchi 1996, GN; Connors et al. 2006, C06), we briefly discuss it in the present paper. The present paper is still based on similar approaches used in previous numerical works (e.g., simple analytic models for dynamical friction and the fixed isothermal density profile for the Galactic dark matter halo). The more sophisticated models in which the LMC, the SMC, and the Galaxy are represented by fully consistent N-body models will be described in our forthcoming papers (Bekki \\& Chiba 2008a).          \\begin{figure}[h] \\begin{center} \\includegraphics[scale=0.6, angle=0]{f1.eps} \\caption { Orbital evolution of the MCs for the models C1 and C2 with the GN (upper) and the K06 (lower) velocity types. The time evolution of distances between MCs (thick solid), the LMC and the Galaxy (thin solid), and the SMC and the Galaxy (dotted) are separately shown. Here $M_{\\rm LMC}=2.0 \\times 10^{10} {\\rm M}_{\\odot}$ and $M_{\\rm SMC}=3.0 \\times 10^{9} {\\rm M}_{\\odot}$ are assumed for the two models. } \\label{figexample} \\end{center} \\end{figure} ",
        "conclusions": "We have investigated structural, kinematical, and chemical properties of stars and gas in the SMC interacting with the LMC  and the Galaxy based on chemodynamical simulations. We have adopted  a new dwarf spheroidal model with an extended HI gas disk for the SMC and investigated the chemical and dynamical evolution and the star formation history of the SMC for the last $\\sim 2.7$ Gyr. The main results are summarized as follows.  (1) The final spatial distribution of the old stars projected onto the sky  is  more  spherical even after the strong LMC-SMC-Galaxy interaction whereas that of the new stars is   significantly flattened and appears to be a bar. This difference in spatial distributions between the old and new stars is due to the fact that new stars are formed in the thin  HI disk of the SMC during the  tidal interaction. These result suggest that the observed ``bar'' is mostly composed of new stars recently formed in the SMC. Some fraction  of old stars initially in the SMC  are stripped to form the outer stellar halo of the SMC. The distribution of the  stellar halo projected onto the sky appears to be slightly asymmetric in the north-south direction, which possibly reflects the strong tidal interaction with the LMC about 0.2 Gyr ago.      (2) Old stars have a  line-of-sight velocity dispersion ($\\sigma$) of $\\sim 30$ km s$^{-1}$ and  weak  rotation with the maximum rotational velocity ($V$) of less than 20 km s$^{-1}$ and show asymmetry in the profiles. New stars have  a smaller $\\sigma$ than old ones  and a significant amount of rotation ($V/\\sigma >1$). These kinematical differences between old and new stars do not depend on model parameters. It should be stressed that the derived $V$ of old stars are highly likely to be overestimated by the stream motions of the stars along the tidal tails of the SMC.   (3) The HI gas shows velocity dispersions  of $ \\sigma \\sim 10-40 $ km s$^{-1}$ depending on the distances from the SMC's center, a high maximum rotational velocity  ($V \\sim 50$ km s$^{-1}$), and a  spatial distribution similar to that of new stars. The derived higher velocity dispersion of gas is due to the tidal disturbance of gas by the LMC and the Galaxy.    (4) The mean metallicity of stars can increase from ${\\rm [Fe/H]} \\sim -0.80$ to ${\\rm [Fe/H]} \\sim -0.56$ for $y_{\\rm met}=0.004$. The stellar populations show a weak AMR in the sense that younger stars are more likely to be  more metal-rich. The final stellar [Fe/H] and the AMR depend strongly on $y_{\\rm met}$ (e.g., the final metallicity is ${\\rm [Fe/H]} \\sim -0.66$ for $y_{\\rm met}=0.002$). More metal-rich stars are more likely to be located in the inner regions of the SMC, irrespective of model parameters.   (5) Although the MB can be formed in most  models, the simulated MB inevitably contains old stars. This is mainly because some fraction of old stars initially in the spheroidal component of the SMC are  tidally stripped to be distributed in the MB region. The mean surface density of the stars in the MB is different in different models and ranges from $ \\sim 6 \\times 10^4 {\\rm M}_{\\odot} {\\rm deg}^{-2}$ to $ \\sim 3 \\times 10^6 {\\rm M}_{\\odot} {\\rm deg}^{-2}$. If these numbers are inconsistent with observations, these results mean that the orbital evolution of the MCs for the last 0.2 Gyr in our models needs to be significantly modified so as to reduce significantly the stellar mass in the MB.   (6) The disk model shows flattened spatial distributions and higher maximum rotational velocities ($V>40$ km s$^{-1}$) in old stars, which appear to be inconsistent with observations. Thus the dwarf spheroidal model can better explain observations than the disk one for a given orbit of the SMC in terms of structures and kinematics of old stars in the SMC. Chemical enrichment processes are not so  different between the dwarf and the disk models.   (7)  The simulated MS, which is composed mostly of HI gas, shows bifurcation in its  structure and kinematics. The models based on the GN velocity type  can much better explain distributions and kinematics of HI gas in the MB than those based on the K06 one that are derived from the latest HST proper motion measurements. Since we did not investigate the MS formation for all of the possible unbound orbits of the MCs, it is  unclear whether all of the unbound orbits for the MCs consistent with  K06 are highly unlikely to explain the MS formation. These results imply that the physical interpretations of the HST proper motion measurements need to be reconsidered in the context of the MS formation. We  have suggested that future abundance studies of the MS through QSO absorption lines can provide valuable information on formation processes and early chemical evolution of the outer HI disk of the SMC.   (8) The SMC is suggested to initially have a large core ($>5$ kpc) in its dark matter distribution and a larger total mass ($M_{\\rm SMC} >3.0 \\times 10^9 {\\rm M}_{\\odot}$) so that the observed kinematical properties of HI in the present SMC can be better explained. The more massive SMC with a low-density dark matter halo is suggested to explain better observations of the inner profile of the HI rotation curve and of the total gas mass in the MS and the LAs.  The present study adopted the ``classical'' orbits of the MCs (GN) rather than the new one (K06) in order to compare results of the present new SMC model with those of the previous ones (e.g., Yoshizawa \\& Noguchi 2003;  C06) that also adopted the classical orbits. We thus have not solved the problems related to the inconsistency between the predicted velocities of the MCs and the observed ones (K06). In our forthcoming papers, we discuss how to solve the problems based on more sophisticated numerical simulations."
    },
    "0806/0806.4269.txt": {
        "abstract": "% Context {} % Aims {It is widely accepted that the progenitors of core collapse SNe are young massive stars and therefore their host galaxies are mostly spiral or irregular galaxies dominated by a young stellar population. Surprisingly, among morphologically classified hosts of core collapse SNe, we find 22 cases where the host has been classified as an Elliptical or S0 galaxy.} % Methods {To clarify this apparent contradiction, we carry out a detailed morphological study and an extensive literature search for additional information on the sample objects.} % Results {Our results are as follows:  1.\tOf 22 ``early type'' objects, 17 are in fact misclassified spiral galaxies, one is a misclassified irregular, and one is a misclassified ring galaxy.  2. Of the 3 objects maintaining the early type classification, one (NGC2768) is a suspected merger remnant, another (NGC4589) is definitely a merger, and the third (NGC2274) is in close interaction. The presence of some amount of young stellar population in these galaxies is therefore not unexpected.} % Conclusions {These results confirm the presence of a limited, but significant, number of core collapse SNe in galaxies generally classified of early type. In all cases, anyway, there are independent indicators of the presence in host galaxies of recent star formation due to merging or gravitational interaction.} ",
        "introduction": "As their cosmological importance as distance indicators continues to grow, a detailed understanding of the origin of supernovae (hereafter SNe) is becoming one of the more perplexing problems in modern astronomy. Especially in the last two decades, improvements in observational capabilities and the broader spectral range of SNe studies have resulted in an unexpected diversity of the observed properties of SNe leading to the introduction of new types and subtypes. These are likely to be related to different physical parameters of the stellar progenitors of SNe (possibly leading to observable differences in the details of their explosion mechanisms), and to differences in the properties of their immediate environment (e.g., \\citealp{TBP07}). Current limits on our understanding of these issues make it difficult to come to firm conclusions on the role of SNe in several fields like the chemical evolution of galaxies, the production of the stellar remnants and the origin of cosmic rays. Perhaps most vexing are the doubts that remain on the reliability  with which SNe can be used as distance indicators up to cosmological distances.  It is widely accepted that SNe result from two major explosion mechanisms: core collapse in massive stars ($M~\\geq~8M_\\odot$) and thermonuclear runaway of in a WD approaching the Chandrasekar limit by accretion from a companion. The focus of this paper is on SNe formed via core collapse (hereafter CCSNe). Different configurations of the progenitor at the moment of explosion, different energetics associated with the event, and possible interaction of the ejecta with circumstellar material produce a large variety of characteristics, presently classified as SNe of types Ib and Ic and several sub-types of type~II. Spectacular evidence for the core collapse nature of SNe II comes from the detection of neutrinos of SN1987A in LMC but indication that CCSNe have massive progenitors also comes from about a dozen cases of direct identification of core collapse SN progenitors (e.g., \\citealp{vD05,HSC06}), from theoretical modeling (e.g., \\citealp{WW95,HWF03}), and  from studies of the environment of these SNe. Their association with spiral arms in galaxies (e.g., \\citealp{MvdB76,BTF94}), with massive star-forming regions (e.g., \\citealp{vD92}), and with local young stellar population (e.g., \\citealp{vDPB99}) is well known. This is consistent with the dependence of the rate of explosions with the morphological types of the host galaxies. Many studies have concluded that types II and Ib/c SNe occur only in spiral and irregular galaxies, and mostly in those of the latest types (e.g., \\citealp{CET99}). In other words different studies show  that the progenitor stars of CCSNe are young massive stars born in recent episodes of star formation in their host galaxies (e.g., \\citealp{H03}), with possible differences in their nature (e.g., \\citealp{vD+03}). Conversely, the apparent certainty of the association between CCSNe and young massive stars has led to use CCSNe as tracers of recent star formation in galaxies (e.g., \\citealp{Nav+01,Petr+05}).  It is therefore puzzling that several authors have reported the presence of CCSNe in E or S0 galaxies. In three articles, \\citeauthor*{vdBLF02} (\\citeyear{vdBLF02,vdBLF03,vdBLF05}, hereafter vdBLF) discussed the morphologies of host galaxies in the Lick Observatory Supernova Search (LOSS; e.g., \\citealp{FLT+01}) and report 8 cases of E and S0 galaxies harboring CCSNe. In a new examination of the objects listed in the latest version (October 30, 2007) of the Asiago Supernova catalog (ASC; e.g., \\citealp{BBCT99}) we discovered 14 more cases of CCSNe host galaxies classified as type E or S0. If demonstrated, the presence of CCSNe in early type galaxies should be associated with the population of the \\textit{prompt} SN~Ia advocated by \\citealp{MDVP06} to explain the excess of SN~Ia events observed in radio-loud ellipticals. This excess is attributed to recent episodes of massive star formations which are expected to produce CCSNe although in still debated number (\\citealp{GRD08}).  It is known that type Ib/c SNe can be misclassified as SN~Ia because of the similarity of their spectra at some epoch. A recent example is SN2004aw (\\citealp{TPMV+06}). It is not excluded therefore that some less extensively studied SN~Ia reported in early type galaxies, might indeed be misclassified SN~Ib/c. However, since to our knowledge a statistics of these cases is not available, we neglect such possible effect in what follows. In this paper we study in detail all 22 cases with the aim of understanding the reliability of the host galaxy classification. The sample is presented in more detail in Sect.~2. The results of our study of individual host galaxies are presented in Sect.~3, and in Sect.~4 we discuss and summarize our results. Throughout this article we have assumed a value of $H_0=75 \\,\\rm km \\,s^{-1} \\,Mpc^{-1}$ for the Hubble constant. ",
        "conclusions": "\\label{discus}  It is generally believed (e.g., \\citealp{vdBT+91,CET99}) that the hosts of core collapse SNe are objects with young stellar population (generally spiral or irregular galaxies) while the old stellar population of early type galaxies can produce only SNe~Ia. Nevertheless among the morphologically-classified host galaxies of CCSNe, we have noted 22 cases in which the host has been classified as elliptical or S0, in apparent contradiction to this conventional view. Indeed it is well known that the morphological classification of galaxies has margins of subjectivity. For instance \\citet*{vdBLF03} have noted that ``\\ldots late type galaxies of unusually high surface brightness may be misclassified as being of early type'', and mention several additional minor issues with the classifications.  For the 22 cases in which the host galaxies of CCSNe have been reported as early types we have assembled imaging and searched the literature for additional information. We found that 14 hosts clearly ought to be re-classified as spirals, for three galaxies the presence of spiral structure is suspected, one galaxy has been classified as irregular, one as a barred galaxy with an outer ring structure. We share the view with \\citet*{vdBLF03} that  high surface brightness in  galaxies plays a crucial role in misclassification, particularly for largely inclined galaxies. There are no detailed studies on the surface brightness of our sample galaxies but we outline that six out of 22 galaxies are largely inclined objects. Since radial velocity may be another important factor influencing morphological classification we have compared the mean radial velocities of the total sample of CCSNe hosts in ASC with available recession velocity and morphology (910 galaxies) with those of the sample of 22 galaxies. The mean radial velocity of the total sample is $4935\\pm3153\\, \\rm km~s^{-1}$ to be compared with $4996\\pm2434\\, \\rm km~s^{-1}$ of our galaxies. This means that the uncertain classification was not due to the distance effect.  Other parameters, in addition to the morphological type, bar and ring are indicators of young stellar populations. In the sample there is one merger (NGC4589) with elliptical outer isophotes and active nuclear region, which may also have enhanced star formation (e.g., \\citealp{KHT+03}). In fact, the SN in this galaxy is close to the nucleus. Another galaxy (NGC2768) has been confirmed as S0 with a central region showing dust and a disky central gas distribution. This galaxy has a powerful LINER nucleus, and may also be the remnant of a merger. The existence of strong HI and CO emission, and the presence of dust and ionized gas supports the presence of a young stellar population. The core collapse SN in this galaxy is possibly associated with the red star clusters present. Only one galaxy of the 22, NGC2274, remains confirmed with an elliptical morphology without any inner structural peculiarity. This galaxy is a member of an isolated pair, and appears to be in interaction with a close neighbor spiral. NGC2274 may possibly have accreted some gas during a tidal encounter with this gas-rich neighbor, thereby accounting for the presence of massive stars. Despite the low number statistics it is interesting to note the unusually high ratio $\\rm SN~Ibc~/~SN~II=2/1$ found in early type galaxies with respect to other  galaxy types (e.g., Table~5 of \\citealp{BBCT99}).  Our sample of 22 galaxies includes 17 HI sources (more than 77\\% of the total sample). It has been shown many times that the presence of HI emission is correlated with star formation, and the HI is generally found in close proximity to star-forming regions (e.g., \\citealp{Hel07}). A high HI detection rate among our sample galaxies is consistent with the presence of a young stellar population in these objects. Our sample includes also 8 1.4~GHz continuum radio sources. According to the criterion adopted by \\citet{DVPP+05}, among  8 radio sources detected 2 (NGC838, UGC2836) are radio-loud and 6 (IC1429, NGC2768, NGC3720, IC2461, VIIIZw140 and NGC4589) are radio-faint. In two cases (NGC2768, NGC4589) the LINER AGNs might contribute to the radio emission (e.g., \\citealp{FNW+00}). In the 6 remaining cases the radio emission might witness recent star formation (e.g., \\citealp{R-GST+07}). It is interesting to note that two galaxies in the sample are UV excess Markarian galaxies, and three have LINER type AGNs. In one case (NGC1260 with SN2006gy), the SN is located in the inner core of the galaxy characterized by extended $\\rm H\\alpha$ emission, consistent with the supposed large mass of the progenitor.  In summary:  \\begin{enumerate} \\item In 17 cases, the ``early type'' host is a misclassified spiral galaxy. \\item In one case the host galaxy is a misclassified irregular, and in another it is a misclassified ring galaxy. \\item Three galaxies retain the early type classification: NGC4579 a merger remnant with elliptical outer isophotes; NGC2768, a S0 with LINER nucleus, possibly the remnant of a merger; NGC2274 an elliptical with no morphological peculiarity which anyway appears to be in close interaction with a neighbour. \\end{enumerate}  In conclusion, the presence of CCSNe in early type galaxies can be interpreted as an additional indication that residual star formation episodes take place also in E and S0 galaxies due to merging/accretion or interaction with close neighbours. In other words the morphology of a galaxy is the most direct and immediate indicator of its general stellar content and star formation history. Nevertheless the morphology is inadequate to describe local SF episodes occurring at small scale.  The available small statistics is inadequate to perform a detailed analysis of the frequency of CCSNe in early type galaxies, hence to test whether the presence of CCSNe is compatible with the existence of a \\textit{``prompt''} population of SN~Ia as invoked by \\citet{MDVP06}."
    },
    "0806/0806.0074_arXiv.txt": {
        "abstract": "{}{We present the results of a complete spectral survey of the X-ray point sources detected in five XMM-Newton observations along the major axis of M31 but avoiding the central bulge, aimed at establishing the population characteristics of X-ray sources in this galaxy.}{We obtained background subtracted spectra and lightcurves for each of the 335 X-ray point sources detected across the five observations from 2002. We also correlate our source list with those of earlier X-ray surveys and radio, optical and infra-red catalogues.  Sources with more than 50 source counts are individually spectrally fit in order to create the most accurate luminosity functions of M31 to date.}{Based on the spectral fitting of these sources with a power law model, we observe a broad range of best fit photon index.  From this distribution of best fit index, we identify 16 strong high mass X-ray binary system candidates in M31. We show the first cumulative luminosity functions created using the best fit spectral model to each source with more than 50 source counts in the disc of M31. The cumulative luminosity functions show a distinct flattening in the X-ray luminosity L$_\\mathrm{X}$ interval $37.0 \\lesssim \\log \\mathrm{L_X}$ erg s$^{-1} \\lesssim 37.5$.  Such a feature may also be present in the X-ray populations of several other galaxies, but at a much lower statistical significance. We investigate the number of AGN present in our source list and find that above L$_\\mathrm{X}\\sim1.4\\times10^{36}$ erg s$^{-1}$ the observed population is statistically dominated by the point source population of M31.}{} ",
        "introduction": "\\label{sec:intro} The Andromeda Galaxy (M31) is the nearest spiral galaxy to our own, lying at a distance of 760 kpc \\citep{vdBergh00}. The sources in M31 are observed at a nearly uniform distance and through an absorption column significantly lower than for sources in the Galactic plane. Thus M31 is an ideal target for studying the emission from the X-ray point sources in a galaxy similar to the Milky Way.  M31 has been observed with many X-ray observatories since \\emph{Einstein}, when \\citet{VS79} observed 69 point sources above 5$\\times10^{36}$ erg\\ s$^{-1}$. Two ROSAT surveys \\citep{Supper97, Supper01} covered most of the M31 disc and found 560 sources above 5$\\times10^{35}$ erg\\ s$^{-1}$.  There have also been many \\emph{Chandra} \\citep[e.g.][]{Williams04, Kong02, Kong03} and XMM-Newton \\citep[e.g.][]{Trud06, PFH05, Osborne01, Shirey01} surveys of both the disc and central region of M31.  The X-ray emission from M31 is dominated by point sources mostly consisting of X-ray binary systems (XBs). \\citet{Trud06} surveyed 123 sources in the central region of M31 and reported that the majority have X-ray properties reminiscent of Galactic low mass XBs (LMXBs), and labelled 44 sources as XB candidates based on their spectral properties and variability.  Six neighbouring, slightly overlapping XMM-Newton observations along the major axis of M31 were made in January and June 2002.  These observations, along with others taken between 2000 and 2007, form part of a survey of the whole optical D$_{25}$ ellipse of M31. Since the central region of M31 is well studied and the wider survey has not been completed at the time of this work, we have investigated the five remaining major axis observations that exclude the core region. Henceforth we refer to these fields as the M31 disc fields. These observations were long and uninterrupted; together with the unprecedented effective area of XMM-Newton, they yielded up to 40 times the photon counts of the best previous observations. Previous work on these XMM-Newton fields \\citep[see e.g.][]{PFH05, Trud02} has derived only spectral properties for the brightest few sources.  In this paper we re-analyse the five M31 disc fields. For the first time we extend the spectral analysis to sources down to L$_\\mathrm{X} \\gtrsim 10^{36}$ erg s$^{-1}$. We create a new source list, derive the spectral parameters of each source and create spatially resolved cumulative X-ray luminosity functions (CLFs). In Sect.~\\ref{sec:obs} we give details of the observations and data reduction, Sect.~\\ref{sec:analysis} covers the analysis and the results of cross correlations with the \\citet{PFH05} catalogue and catalogues at other wavelengths.  In Sect.~\\ref{sec:results} we give details of the analysis of our spectral fitting including the creation of CLFs and comments on the contamination of the CLF by background AGN.  Finally, Sect.~\\ref{sec:conclusion} summarises our findings. ",
        "conclusions": "\\label{sec:conclusion} We have revisited five archival XMM-Newton observations of the disc of M31.  These data revealed 335 point detections across the 5 fields constituting 329 discrete point sources. All the sources were fit with three spectral models: blackbody, bremsstrahlung and power law and the results of these fits were examined.  Using only the best fit power law model to each source, we investigated the distribution of photon indices of these fits. The broad range of photon index seen in Fig.~\\ref{fig:histo} and the difference in the CLFs seen in Fig.~\\ref{fig:kil_CLF} cast doubt on the validity of assuming the same spectral model for all sources when analysing more distant X-ray point source populations.  Individual spectral fitting has identified the first 18 HMXB candidates in the disc.  The HMXB candidates are all best fit by a power law with a photon index of 0.8-1.2 indicating magnetically acreting neutron stars.  For the first time X-ray point sources in M31 with as few as 50 source counts have been individually spectrally fit, in contrast to previous surveys which have used the same assumed model for all sources except the very brightest. This has led us to create the first CLFs of the M31 disc region created from 240 individually spectrally fit sources shown in Fig.~\\ref{fig:CLF}. It can be seen that the CLFs of the fields are quite similar across the disc of M31 and that there are no obvious changes in the CLF slopes with increasing apparent distance from the core.  The CLFs of both individual fields (Fig.~\\ref{fig:CLF}) and the entire disc both with and without GC sources (Fig.~\\ref{fig:whole_CLF}) show a distinct flattening between 37.0 $\\lesssim \\log \\mathrm{L_X}$ erg s$^{-1} \\lesssim 37.5$.  This flattening could be due to the emission from a mixture of HMXB and LMXB populations or due to a change in the nature of the compact object or the donor star. This prominent flattening in the CLF of M31 sources may also appear at a lower statistical significance in the CLFs of several other galaxies.  Each observation contains not only the point sources in M31 but also some contamination from foreground and background sources. We have estimated that there are around 20-31 background AGN above $10^{36}$ erg s$^{-1}$, in each field observed, and find that above $1.4\\times10^{36}$ erg s$^{-1}$ there are few background AGN according to \\citet{Moretti03}. The CLF here is dominated by the sources intrinsic to the disc of M31 and any foreground interlopers. Following the correlations of the PSPC ROSAT survey with optical catalogues \\citep{Supper01}, we expect 5-10 forground objects above $10^{36}$ erg s$^{-1}$ for each XMM-Newton field \\citep{Shirey01}.  M31 is a prime target for population surveys because of its proximity and similarity to our own Galaxy. We have investigated the X-ray point sources in the disc of M31 in detail and this population challenges theoretical models to explain the features seen in the distribution of sources and in the CLF."
    },
    "0806/0806.0845_arXiv.txt": {
        "abstract": "Black hole mimickers are possible alternatives to black holes, they would look observationally almost like black holes but would have no horizon. The properties in the near-horizon region where gravity is strong can be quite different for both type of objects, but at infinity it could be difficult to discern black holes from their mimickers. To disentangle this possible confusion, we examine the near-horizon properties, and their connection with far away asymptotic properties, of some candidates to black mimickers. We study spherically symmetric uncharged or charged but non-extremal objects, as well as spherically symmetric charged extremal objects. Within the uncharged or charged but non-extremal black hole mimickers, we study non-extremal $\\varepsilon$-wormholes on the threshold of the formation of an event horizon, of which a subclass are called black foils, and gravastars.  Within the charged extremal black hole mimickers we study extremal $ \\varepsilon$-wormholes on the threshold of the formation of an event horizon, quasi-black holes, and wormholes on the basis of quasi-black holes from Bonnor stars. We elucidate, whether or not the objects belonging to these two classes remain regular in the near-horizon limit. The requirement of full regularity, i.e., finite curvature and absence of naked behavior, up to an arbitrary neighborhood of the gravitational radius of the object enables one to rule out potential mimickers in most of the cases. A list ranking the best black hole mimickers up to the worse, both non-extremal and extremal, is as follows: wormholes on the basis of extremal black holes or on the basis of quasi-black holes, quasi-black holes, wormholes on the basis of non-extremal black holes (black foils), and gravastars. Since, in observational astrophysics it is difficult to find extremal configurations (the best mimickers in the ranking), whereas non-extremal configurations are really bad mimickers, the task of distinguishing black holes from their mimickers seems to be less difficult than one could think of it. ",
        "introduction": "In recent years, it has been debated in the literature about possible alternatives to black holes, the black hole mimickers, which would look observationally almost like black holes but would have no horizon. The existence of such objects can, in principle, put in doubt astrophysical data which otherwise are considered as observational confirmation in favor of black holes \\cite{abr}. On one hand, it is clear that the properties in the near-horizon region where gravity is strong can be quite different for both type of objects. On the other hand, the statements about the difficulties in discerning black holes from their mimickers are usually related to measurements at spatial infinity. Thus, one should insist on the question: Can an observer at infinity catch the difference between both type of objects in some indirect way, or even rule out some possible mimicker? In our view, the answer is positive and is connected with key properties, namely, regularity or singularity, of the corresponding geometries. It turns out that the requirement of full regularity up to an arbitrary neighborhood of the gravitational radius of the object enables one to rule out the potential mimickers in most of the cases.  The goal of the present work is to examine the near-horizon properties, and their connection with far away asymptotic properties, of some candidates to black mimickers. We study spherically symmetric configurations, and make, two major divisions, or classes, on those candidates. First, uncharged or charged but non-extremal objects, and second extremal objects. Within the uncharged or charged but non-extremal one can invoke as black hole mimickers, non-extremal $\\varepsilon$-wormholes on the threshold of the formation of an event horizon, some of which are called black foils \\cite{ds} (see \\cite{sm} for the construction with other purposes of $\\varepsilon$-wormholes, which actually can also act as mimickers), and gravastars \\cite{grav}. Within the extremal charged class one can invoke extremal $ \\varepsilon$-wormholes on the threshold of the formation of an event horizon, quasi-black holes \\cite{qbh} (see also \\cite{lemoszanchinjmp}), and wormholes on the basis of quasi-black holes from Bonnor stars, to name a few. We want to elucidate, whether or not the objects belonging to these two classes remain regular in the near-horizon limit. The arguments of \\cite{qbh} which rule out non-extremal limiting configurations as becoming singular do not apply to the wormhole case \\cite{ds}. Thus, we carry out the corresponding analysis anew for both classes of objects. ",
        "conclusions": "We have studied wormhole and other configurations as possible mimickers of black holes. We have separated the configurations into non-extremal and extremal.  For wormholes, we have examined separately two limiting procedures in which the wormhole throat approaches the black hole horizon. In the first procedure, we fix the location of an observer (exactly on the throat). Then, we change the spacetime (making a wormhole on the verge of being a black hole). This is the situation BT. In the second, we change spacetime (making a wormhole on the verge of a black hole), then place the shell outside the throat and move it toward the throat (which coincides with the horizon).  This is the situation TB. This procedure is carried out for non-extremal and extremal configurations separately. In the non-extremal case it turned out that each of the limits under discussion is singular: either the Kretschmann scalar or surface stresses on the throat (or both) diverge. Thus, the limit is singular. In other words, a mimicker of a non-extremal black hole, made from a non-extremal wormhole, including the black foil of \\cite{ds}, is not smooth. We have summarized the results for the non-extremal case in a table.  For the extremal case both the Kretschmann scalar and surface stresses remain finite. This pronounced distinction between properties of the limiting configurations in the non-extremal and extremal cases is one counterpart of the conclusion made in \\cite{qbh} that quasi-black holes can be only extremal. However, one should not forget about some subtleties connected with the fact that singular behavior can, in general, manifest itself not only in the value of the Kretschmann scalar. Even if this scalar is finite, a naked behavior is possible or even inevitable as was shown in in the present paper and in \\cite{qbh} (see, e.g., Sec. V of \\cite{qbh} for a discussion about other subtleties in which the singular features of quasi-black holes are revealed). There is also another candidate for the role of a non-extremal black hole mimicker, a gravastar \\cite{grav}.  However, the corresponding surface stresses grow unbound when the radius approaches the gravitational radius, as we have seen.  {} From an astrophysical viewpoint, the situation BT in the non-extremal case, i.e., a black foil, is more interesting since it implies no necessity of making a shell by hand. It is the case considered in \\cite{ds}. One of the questions raised in \\cite{ds} is whether it is possible or not to distinguish between a true non-extremal black hole, a Schwarzschild black hole say, and a wormhole. The main conclusion of \\cite{ds} is that it is impossible to distinguish for any finite time in the limit under discussion. This conclusion is reached on the basis of considering properties of bodies moving along separate fixed geodesics and emitting signals detected at infinity. However, if from single geodesics we shift our attention to a congruence of geodesics, it turns out that the strong gravity forces, on the near horizon region, leave their imprint on the form of a moving body and, thus, on the properties of signals which an observer at infinity is detecting. If the surface of the body is luminescent, an observer at infinity would see either a finite width instead of a point, a continuous detection instead of separate pulses, and so on. It is essential that in the case discussed in \\cite{ds} the corresponding proper time of deformation tends to zero when the curvature grows unbound, with the time at infinity being finite. Thus, at least in principle, an observer can distinguish between a black hole and an almost singular wormhole. The singular nature of the limit in the non-extremal case makes also questionable the applicability of the membrane paradigm used in \\cite {ds}. The key point in this paradigm consists in boundary conditions according to which a free-falling observer sees a finite value of physical fields on the horizon (see \\cite{th}, Sec. II). However, in the problem under discussion, typically, this observer as well as the geometry itself become ill-defined. Only in some cases (see Model 2, situation TB) the curvature components remain finite in the free-falling frame. But even in such situations the infinite surface stresses on the horizon surface makes the physical meaning of the membrane paradigm unclear since this paradigm relies heavily on the concept of a regular surface.  Objects based on nearly extremal wormholes, although of less interest astrophysically perhaps, have a much better behavior in the sense that both the geometry and surface stresses remain finite. Moreover, typically, there is no naked behavior. In this case, the effect of strong curvature is much less pronounced than in the case of quasi-black holes where a naked behavior is typical \\cite{qbh}. In this sense, a wormhole composed on the basis of two extremal black holes seems to be the best mimicker of an extremal black hole. As by-product, we have obtained a model of a regular black hole!  Thus, if we try to arrange a ranking of black hole mimickers, both non-extremal and extremal, the list looks as follows from top to bottom: wormholes on the basis of extremal black holes or on the basis of quasi-black holes, quasi-black holes, wormholes on the basis of non-extremal black holes (and within these the best are black foils), and gravastars.  Bearing in mind that in observational astrophysics it is difficult to find extremal configurations (the would be best mimickers), whereas non-extremal configurations are really bad mimickers, the task of distinguishing black holes from their mimickers seems to be less difficult than one could think of it.  In the present paper we have restricted ourselves to static spherically symmetric spacetimes. Meanwhile, in a recent work \\cite{cpcce} the status of black hole mimickers is undermined in the rapidly rotating case as well since it is argued that they are unstable. We have also circumscribed our discussion to particular wormholes, gravastars and quasi-black holes \\cite{ds}-\\cite{lemoszanchinjmp}, since these objects are well adapted to our goal of examining the near-horizon properties and their connection with far away asymptotic properties. However, there are many other objects with properties that make them also potential black hole mimickers (see, e.g., \\cite{dars,guz,lib,sken}) and which are worthy of study within our formalism."
    },
    "0806/0806.0418_arXiv.txt": {
        "abstract": "We calculate the prompt neutrino flux from atmospheric charm production by cosmic rays, using the dipole picture in a perturbative QCD framework, which incorporates the parton saturation effects present at high energies.  We compare our results with the next-to-leading order perturbative QCD result and find that saturation effects are large for neutrino energies above $10^6$ GeV, leading to a substantial suppression of the prompt neutrino flux.  We comment on the range of prompt neutrino fluxes due to theoretical uncertainties. ",
        "introduction": "Atmospheric neutrinos are produced in interactions of cosmic rays with Earth's atmosphere.  The observation of low energy ($E_\\nu \\sim$ GeV) atmospheric neutrinos, their flavor-dependent interactions, and their path length dependence~\\cite{SuperK, SNO} has confirmed the existence of neutrino flavor transformation, and therefore the most fundamental property of the neutrinos:\\ that they are not massless.  These observations have provided a remarkable source of information about mass and mixing parameters of neutrinos.  Atmospheric neutrinos are also a background to other sources of neutrinos, such as cosmogenic neutrinos produced in interactions of cosmic rays with the background radiation ~\\cite{ess} and directly from sources such as active galactic nuclei and gamma ray bursts~\\cite{rev}. Observation of neutrinos coming from these distant sources would provide valuable information about the particle production mechanism in astrophysical sources.  Neutrino interactions in the Earth and in the atmosphere could also serve as unique probes of physics beyond the Standard Model \\cite{nonstd} . It has been recently suggested that atmospheric neutrinos, in their interactions in the Earth, could produce supersymmetric particles if the neutrino energies are sufficiently high.  In Ref.~\\cite{ando}, Ando et al.\\ have suggested that the high energy atmospheric neutrino flux may be large enough to produce quasi-stable charged particles that are potentially detectable in the IceCube~\\cite{icecube} neutrino detector. As a background to high energy sources or as a flux to produce exotic particles in the Earth, it is useful to re-evaluate the high energy component of the atmospheric neutrino flux.  The atmospheric fluxes of neutrinos at low energies have been extensively studied~\\cite{agrawal,battistoni,honda,gaisserhonda}. They arise mainly from the products of charged pion and kaon decays. As energies increase, the decay lengths of the mesons become longer than their path lengths in the atmosphere~\\cite{pathlength}, suppressing the production of neutrinos. Other, shorter lived hadrons are also produced at high energies. They too contribute to the neutrino flux, especially from the ``prompt'' decay of charmed mesons. The energy dependence of these prompt neutrinos is less steep than the ``conventional'' neutrino flux from pion and kaon decays. The energy at which the prompt neutrinos become the dominant atmospheric neutrino component depends on the details of the mechanism for charm production in proton--air collisions at high energies. Charm contributions to the atmospheric lepton fluxes have been evaluated analytically and semi-analytically \\cite{volkova,Gondolo:1995fq,Pasquali:1998ji,graciela,Martin:2003us}. There are renewed efforts include charm production with the dual parton model in Monte Carlo simulations of air showers as well \\cite{bat,Berghaus:2007hp}.  For vertical neutrino fluxes, the cross-over between conventional and prompt dominated fluxes occurs in the energy range of $10^5$--$10^6$ GeV for the calculations of Refs. \\cite{Gondolo:1995fq,Pasquali:1998ji,graciela,Martin:2003us}, and the cross-over energy increases with zenith angle. For energies above $\\sim 10^6$~GeV, the dominant contribution to charm production comes from gluons, where saturation effects~\\cite{glr} due to dense, interacting gluons in the nucleus become important. We evaluate the prompt neutrino flux using perturbative QCD in the dipole framework, taking these effects into account.  We study the theoretical uncertainties inherent in this approach and compare with standard next-to-leading order perturbative QCD. The range of QCD-based predictions yields prompt neutrino fluxes that are unlikely to be large enough to produce a detectable number of exotic particles of the type discussed in Ref.~\\cite{ando}.  We begin with a discussion of the cross section for charm production in the dipole picture in Section II. Using the dipole picture results, we discuss the evaluation the prompt neutrino flux from charm decays in Section III. Our results and a comparison with the conventional fluxes are shown in Section IV. We also discuss uncertainties in the QCD approach, and compare our results with the uncertainty band of Ref.~\\cite{ando} in Section IV. ",
        "conclusions": "\\begin{figure} \\begin{center} \\epsfig{file=fig_flux_errorband_both_log.eps, width=\\columnwidth} \\caption{Prompt and conventional $\\nu_\\mu + \\bar\\nu_\\mu$ fluxes in the vertical direction. The shaded band is the theoretical uncertainty band for the prompt flux calculated in this paper with the dipole model. The dashed line shows the conventional flux from Gaisser and Honda (GH) \\cite{gaisserhonda} and the dotted line is the conventional flux calculated in Ref.~\\protect\\cite{Gondolo:1995fq} (TIG). \\label{fig:errorband}} \\end{center} \\end{figure}   Our result for the vertical muon neutrino plus antineutrino flux from atmospheric charm is shown in Figure~\\ref{fig:errorband}, which shows the theoretical uncertainty band for the DM calculation, estimated as described above. For comparison the conventional neutrino fluxes \\cite{gaisserhonda,Gondolo:1995fq} from $\\pi$- and $K$-decays are also shown. We find that the vertical prompt muon neutrino flux becomes dominant over the conventional neutrino flux at energies between $10^5$ GeV and $10^{5.5}$ GeV.  The theoretical uncertainty due to choices of gluon distribution, charm quark mass, factorization scale, and other parameters in the dipole model result in the range of fluxes represented by the shaded area in Figure~\\ref{fig:errorband}. The shape of the prompt neutrinos is only weakly dependent on the choice of parameters, but the overall normalization could vary by up to a factor of two in this model for charm production.    \\begin{figure} \\begin{center} \\epsfig{file=fig_flux_compare_others_lin.eps, width=\\columnwidth} \\caption{Prompt muon neutrino fluxes obtained in perturbative QCD. The shaded area represents the theoretical uncertainty in the prompt neutrino flux evaluated in this paper, and the solid line in the band is our standard result. The dashed curve is the NLO perturbative QCD calculation of Ref. \\protect\\cite{Pasquali:1998ji} (PRS), modified here to include fragmentation, the dotted  curve is the saturation model result of Ref.~\\protect\\cite{Martin:2003us} (MRS), and the dash-dotted curve is the LO perturbative QCD calculation of Ref.\\ \\protect\\cite{Gondolo:1995fq} (TIG). \\label{fig:compare_others}} \\end{center} \\end{figure}   We compare our result to three earlier calculations of the prompt neutrino flux: \\begin{enumerate} \\item Thunman, Ingelman and Gondolo (TIG)~\\cite{Gondolo:1995fq}. This was the first  perturbative QCD calculation and was done at the leading order (LO) in $\\alpha_s$. It takes the fragmentation of charm quarks into account through Monte Carlo simulation using the Lund string model~\\cite{Andersson:1983ia} implemented in the event generator Pythia~\\cite{pythia}. The small-$x$ PDFs are extrapolated with e.g., $xG(x,\\mu^2)\\sim x^{-0.08}$. \\item Pasquali, Reno and Sarcevic (PRS)~\\cite{Pasquali:1998ji}. This result uses the next-to-leading order (NLO) QCD result of~\\cite{NLO} with power law extrapolations of the small-$x$ PDFs. The PRS evaluation does not take fragmentation into account. We have therefore carried out a simplified version of this calculation, taking fragmentation into account in the same way as we did for the DM calculation:\\ we compute the charmed hadron cross section in leading order QCD using KK fragmentation functions~\\cite{KniehlKramer}, and multiply with a K-factor $K=\\sigma(NLO)/\\sigma(LO) \\approx 2$. This reproduces the full NLO calculation of Ref.~\\cite{Pasquali:1998ji} at the parton level to an adequate accuracy. \\item Martin, Ryskin and Sta\\'sto (MRS)~\\cite{Martin:2003us}. This calculation takes fragmentation into account by assigning the neutrino a fixed fraction of the momentum of the mother meson, and is done using the saturation model of Golec-Biernat and W\\\"ustoff~\\cite{GolecBiernat:1998js} described above. \\end{enumerate}  We show the results from these other evaluations of the vertical muon neutrino plus antineutrino flux together with our uncertainty band in Figure~\\ref{fig:compare_others}. The theoretical uncertainties in the standard NLO QCD calculation of the charm cross section are the choice of the renormalization and factorization scales, the charm mass, and the small $x$ behavior of the gluon distribution~\\cite{Vogt:2007aw}. The impact of some of these uncertainties on the neutrino flux has been studied in Ref.\\ \\cite{graciela}.  The MRS curve in Fig. \\ref{fig:compare_others} is at the lower border of our DM uncertainty band. There is approximately a factor of two between the MRS and the central DM results, coming from the different parameterizations of $\\sigma_d$.  The enhancement is also seen in calculations of photoproduction of heavy quarks \\cite{Goncalves:2004dn} comparing the GBW model and the improved DM model of Eq. \\ref{sigmadip1}. The DM cross section for charm pair production in $pp$ collisions lies within the uncertainty band of Ref. \\cite{Vogt:2007aw}.  \\begin{figure}[t] \\begin{center} \\epsfig{file=fig_flux_CGC_NLO_frag_lin.eps, width=\\columnwidth} \\caption{The effect of fragmentation on predicted $\\nu_\\mu + \\bar\\nu_\\mu$ fluxes. The solid lines are DM and NLO QCD (the latter multiplied by two to separate the lines) with Kniehl--Kramer fragmentation. The dashed lines are without fragmentation. \\label{fig:fragmentation}} \\end{center} \\end{figure}  The effect of quark fragmentation on the neutrino fluxes is rather large because fragmentation reduces the energy of the charmed hadrons. For a given hadron energy, fragmentation effects require  higher energy cosmic rays in the steeply falling cosmic ray flux. In Figure~\\ref{fig:fragmentation} we show the effect of including the KK fragmentation functions on both the NLO QCD and DM results. The NLO results are multiplied by a factor of two so that they can be distinguished easily from the DM results. The fragmentation reduces the flux by between 60\\% and 70\\%, and thus it is an important effect to take into account. The Peterson fragmentation function results differ by approximately 10\\% from the results shown in Figure~\\ref{fig:fragmentation}. We include the uncertainty in fragmentation model in our estimate of the theoretical uncertainty; however, we do not consider the result without fragmentation in our uncertainty estimate.   \\begin{figure}[tb] \\begin{center} \\epsfig{file=fig_flux_compare_ABPR_E1_v2.eps, width=\\columnwidth} \\caption{Theoretical uncertainty band (blue, dark band), compared to the uncertainty band from Ref.\\ \\cite{ando} (magenta, light band) with the overlapping region shown as the middle (light blue) band. The flux is scaled by $E$, in units of 1/km$^2$yr\\,sr. \\label{fig:compare_ABPR}} \\end{center} \\end{figure}  Other perturbative QCD calculations are unlikely to give a much larger prompt neutrino flux than our upper limit, if saturation is indeed important. The theoretical expectation is that saturation is important at scales comparable to $\\mu\\sim m_c$ for small $x$ values. If it would turn out that saturation \\emph{does not} occur at the relevant energy scales, the flux is still not expected to be much larger than the PRS result. In Figure~\\ref{fig:compare_ABPR}, we therefore show a comparison of the uncertainty (blue, dark band) compared to the proposed uncertainty range from Ref.~\\cite{ando} (magenta, light band) with their overlapping region (middle, light blue band).  In this plot we take the NLO QCD result as the upper theoretical limit. This gives a larger upper limit than in the earlier plots, which show only the uncertainty in the dipole model result. We stress that, since saturation is expected to be important on theoretical grounds, the uncertainty band in Figure \\ref{fig:errorband} is our main result.  In order to obtain a flux as large as the upper line in the  uncertainty band of Figure~\\ref{fig:compare_ABPR}, we must multiply the upper uncertainty line of our DM result by a factor of 50. A cross section a factor of 50 times larger than the DM evaluation in proton-proton scattering would be incompatible with existing cross section measurements, as illustrated for example by Figure~6 of Ref.~\\cite{Goncalves:2006ch}, which compares the DM result for charm production to fixed-target experimental data.  Measurable stau production rates from prompt atmospheric neutrinos as proposed in Ref.~\\cite{ando} would require the highest fluxes in the lighter band. Our evaluation of the prompt neutrino flux indicates that the upper limit of Ref.\\ \\cite{ando} is unrealistically large. The prompt neutrino flux is unlikely to be large enough for studying stau production from neutrino interactions with Earth and the subsequent detection in neutrino telescopes.  \\begin{figure}[t] \\begin{center} \\epsfig{file=fig_flux_allflavors.eps, width=\\columnwidth} \\caption{Prompt (solid line) and conventional (dashed lines) fluxes of $\\nu_\\mu + \\bar\\nu_\\mu$, $\\nu_e + \\bar\\nu_e$, and $\\mu^+ + \\mu^-$. The conventional fluxes are from Ref.~\\protect\\cite{Gondolo:1995fq}. The three prompt fluxes are approximately equal and are therefore here represented by the $\\nu_\\mu + \\bar\\nu_\\mu$ flux. \\label{fig:flavors}} \\end{center} \\end{figure}  The flavor decomposition of an atmospheric neutrino signal may be an interesting way to explore the prompt contribution. The prompt neutrino fluxes of $\\nu_\\mu + \\bar\\nu_\\mu$ and $\\nu_e + \\bar\\nu_e$ are identical, since the charmed mesons decay equally likely into electrons or muons. The prompt $\\mu^+ + \\mu^-$ flux is approximately equal to the neutrino fluxes. However, this does not hold for the conventional fluxes. Charged pions decay almost exclusively into muons, so the muon neutrino and muon fluxes are much larger than the electron neutrino flux. We show the $\\nu_\\mu + \\bar\\nu_\\mu$ prompt flux together with the corresponding vertical conventional fluxes of muons, muon neutrinos and electron neutrinos (and their antiparticles) in Figure~\\ref{fig:flavors}. If experiments would be able to measure electron neutrino fluxes, the prompt flux will start dominating over the conventional flux for much lower energy $\\sim 10^4$ GeV than for the muon neutrino or muon fluxes.  \\begin{figure}[tb] \\begin{center} \\epsfig{file=fig_flux_nutau.eps, width=\\columnwidth} \\caption{Prompt $\\nu_\\tau + \\bar\\nu_\\tau$ flux (solid line) compared with the prompt $\\nu_\\mu + \\bar\\nu_\\mu$ flux (dashed line). \\label{fig:nutau}} \\end{center} \\end{figure}  \\begin{figure}[t] \\begin{center} \\epsfig{file=fig_flux_direction_log.eps, width=\\columnwidth} \\caption{Dependence on zenith angle of prompt and conventional $\\nu_\\mu + \\bar\\nu_\\mu$ fluxes. The solid  and dashed lines are DM with fragmentation in the vertical and horizontal directions, respectively. The dotted and dash-dotted lines are the Gaisser--Honda conventional fluxes \\cite{gaisserhonda} in the vertical and horizontal directions. \\label{fig:direction}} \\end{center} \\end{figure}   We note that the prompt flux of $\\nu_\\tau + \\bar\\nu_\\tau$ from charm decays is much smaller than the other neutrino flavors \\cite{tau}, since only the $D_s$ meson decays into $\\nu_\\tau$. The $\\nu_\\tau + \\bar\\nu_\\tau$ flux from $D_s$ decays is shown in Figure~\\ref{fig:nutau} together with the prompt $\\nu_\\mu + \\bar\\nu_\\mu$ flux. We have not included the contribution from $B$-meson decays which could give a contribution on the order of 10--20\\%~\\cite{Martin:2003us} since $B$-meson decays to $\\nu_\\tau$ plus tau are kinematically allowed.   The vertical direction is the optimal direction for studying the prompt fluxes. In Figure~\\ref{fig:direction} we show the prompt and conventional $\\nu_\\mu + \\bar\\nu_\\mu$ fluxes in the vertical and horizontal directions. In the horizontal direction the prompt flux does not become larger than the conventional flux until very large energies $\\sim 10^7$ GeV, where the actual number of neutrinos is quite low.   In summary, we have computed prompt neutrino and muon fluxes from cosmic ray interactions in the atmosphere that produce charm pairs. Our evaluation of the fluxes takes parton saturation effects into account via the dipole model, a model with a parametric form guided by QCD and constrained by data. We find that saturation effects in the dipole model decrease the prompt fluxes above $10^5$~GeV.  Our estimate of the theoretical uncertainty in the predicted fluxes in the dipole model is on the order of a factor of two. In comparison to other QCD or dipole model evaluations of the prompt flux, the range of predictions is approximately a factor of 6. Future measurements of the high energy neutrino flux will provide interesting constraints on QCD-based evaluations of the prompt flux of neutrinos, however, the prompt neutrino flux is unlikely to be large enough to probe non-standard model interactions."
    }
}